ASTRONOMY STUDY SHEET
Part A
Proto Star:
Protostar
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For other uses, see Protostar (disambiguation).
Star formation
Object classes
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Interstellar medium
Molecular cloud
Bok globule
Dark nebula
Young stellar object
Protostar
Pre-main-sequence star
T Tauri star
Herbig Ae/Be star
Herbig–Haro object
Theoretical concepts
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Accretion
Initial mass function
Jeans instability
Kelvin–Helmholtz mechanism
Nebular hypothesis
Planetary migration
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A protostar is a very young star that is still gathering mass from its parent molecular cloud. The
protostellar phase is the earliest one in the process of stellar evolution.[1] For a low-mass star (i.e.
that of the Sun or lower), it lasts about 500,000 years. [2] The phase begins when a molecular
cloud fragment first collapses under the force of self-gravity and an opaque, pressure supported
core forms inside the collapsing fragment. It ends when the infalling gas is depleted, leaving a
pre-main-sequence star, which contracts to later become a main-sequence star at the onset of
hydrogen fusion producing helium.
History[edit]
The modern picture of protostars, summarized above, was first suggested by Chushiro Hayashi
in 1966.[3] In the first models, the size of protostars was greatly overestimated. Subsequent
numerical calculations[4][5][6] clarified the issue, and showed that protostars are only modestly
larger than main-sequence stars of the same mass. This basic theoretical result has been
confirmed by observations, which find that the largest pre-main-sequence stars are also of
modest size.
Protostellar evolution[edit]
Infant star CARMA-7 and its jets are located approximately 1400 light-years from Earth
within the Serpens South star cluster. [7]
Main article: Star formation
Star formation begins in relatively small molecular clouds called dense cores.[8] Each dense core
is initially in balance between self-gravity, which tends to compress the object, and both gas
pressure and magnetic pressure, which tend to inflate it. As the dense core accrues mass from its
larger, surrounding cloud, self-gravity begins to overwhelm pressure, and collapse begins.
Theoretical modeling of an idealized spherical cloud initially supported only by gas pressure
indicates that the collapse process spreads from the inside toward the outside.[9] Spectroscopic
observations of dense cores that do not yet contain stars indicate that contraction indeed occurs.
So far, however, the predicted outward spread of the collapse region has not been observed. [10]
The gas that collapses toward the center of the dense core first builds up a low-mass protostar,
and then a protoplanetary disk orbiting the object. As the collapse continues, an increasing
amount of gas impacts the disk rather than the star, a consequence of angular momentum
conservation. Exactly how material in the disk spirals inward onto the protostar is not yet
understood, despite a great deal of theoretical effort. This problem is illustrative of the larger
issue of accretion disk theory, which plays a role in much of astrophysics.
HBC 1 is a young pre-main-sequence star.[11]
Regardless of the details, the outer surface of a protostar consists at least partially of shocked gas
that has fallen from the inner edge of the disk. The surface is thus very different from the
relatively quiescent photosphere of a pre-main sequence or main-sequence star. Within its deep
interior, the protostar has lower temperature than an ordinary star. At its center, hydrogen-1 is
not yet fusing with itself. Theory predicts, however, that the hydrogen isotope deuterium
(hydrogen-2) fuses with hydrogen-1, creating helium-3. The heat from this fusion reaction tends
to inflate the protostar, and thereby helps determine the size of the youngest observed pre-mainsequence stars.[12]
The energy generated from ordinary stars comes from the nuclear fusion occurring at their
centers. Protostars also generate energy, but it comes from the radiation liberated at the shocks
on its surface and on the surface of its surrounding disk. The radiation thus created must traverse
the interstellar dust in the surrounding dense core. The dust absorbs all impinging photons and
reradiates them at longer wavelengths. Consequently, a protostar is not detectable at optical
wavelengths, and cannot be placed in the Hertzsprung–Russell diagram, unlike the more evolved
pre-main-sequence stars.
The actual radiation emanating from a protostar is predicted to be in the infrared and millimeter
regimes. Point-like sources of such long-wavelength radiation are commonly seen in regions that
are obscured by molecular clouds. It is commonly believed that those conventionally labeled as
Class 0 or Class I sources are protostars. [13][14] However, there is still no definitive evidence for
this identification.
Observed classes of young stars[edit]
For details of observational classification, see Young stellar object.
Class peak emission duration
(Years)
0
submillimeter
104
I
far-infrared
105
II
near-infrared
106
III
visible
107[15]
https://commons.wikimedia.org/w/index.php?title=File%3AA_Young_Star_Flaunts_its_Xray_Spots.ogv
Hertzsprung–Russell
diagram
An observational Hertzsprung–Russell diagram with 22,000 stars plotted from the Hipparcos Catalogue and
1,000 from the Gliese Catalogue of nearby stars. Stars tend to fall only into certain regions of the diagram.
The most prominent is the diagonal, going from the upper-left (hot and bright) to the lower-right (cooler and
less bright), called the main sequence. In the lower-left is where white dwarfs are found, and above the
main sequence are the subgiants, giants and supergiants. The Sun is found on the main sequence at
luminosity 1 (absolute magnitude 4.8) and B− V color index 0.66 (temperature 5780 K, spectral type G2V).
The Hertzsprung–Russell diagram, abbreviated as H–R diagram, HR diagram or
HRD, is a scatter plot of stars showing the relationship between the stars' absolute
magnitudes or luminosities versus their stellar classifications or effective temperatures.
The diagram was created independently in 1911 by Ejnar Hertzsprung and by Henry
Norris Russell in 1913, and represented a major step towards an understanding of
stellar evolution.
Contents
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1
Historical background
2
Forms of diagram
3
Interpretation
○ 3.1
○ The diagram seen by ESA's Gaia mission
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4
Role in the development of stellar physics
5
See also
6
References
○ 6.1
○ Bibliography
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7
External links
Historical background[edit]
In the nineteenth century large-scale photographic spectroscopic surveys of stars were
performed at Harvard College Observatory, producing spectral classifications for tens of
thousands of stars, culminating ultimately in the Henry Draper Catalogue. In one
segment of this work Antonia Maury included divisions of the stars by the width of their
spectral lines.[1] Hertzsprung noted that stars described with narrow lines tended to
have smaller proper motions than the others of the same spectral classification. He took
this as an indication of greater luminosity for the narrow-line stars, and computed
secular parallaxes for several groups of these, allowing him to estimate their absolute
magnitude.[2]
In 1910 Hans Rosenberg published a diagram plotting the apparent magnitude of stars
in the Pleiades cluster against the strengths of the calcium K line and two hydrogen
Balmer lines.[3] These spectral lines serve as a proxy for the temperature of the star, an
early form of spectral classification. The apparent magnitude of stars in the same cluster
is equivalent to their absolute magnitude and so this early diagram was effectively a plot
of luminosity against temperature. The same type of diagram is still used today as a
means of showing the stars in clusters without having to initially know their distance and
luminosity.[4] Hertzsprung had already been working with this type of diagram, but his
first publications showing it were not until 1911. This was also the form of the diagram
using apparent magnitudes of a cluster of stars all at the same distance.[5]
Russell's early (1913) versions of the diagram included Maury's giant stars identified by
Hertzsprung, those nearby stars with parallaxes measured at the time, stars from the
Hyades (a nearby open cluster), and several moving groups, for which the moving
cluster method could be used to derive distances and thereby obtain absolute
magnitudes for those stars.[6]
Forms of diagram[edit]
There are several forms of the Hertzsprung–Russell diagram, and the nomenclature is
not very well defined. All forms share the same general layout: stars of greater
luminosity are toward the top of the diagram, and stars with higher surface temperature
are toward the left side of the diagram.
The original diagram displayed the spectral type of stars on the horizontal axis and the
absolute visual magnitude on the vertical axis. The spectral type is not a numerical
quantity, but the sequence of spectral types is a monotonic series that reflects the
stellar surface temperature. Modern observational versions of the chart replace spectral
type by a color index (in diagrams made in the middle of the 20th Century, most often
the B-V color) of the stars. This type of diagram is what is often called an observational
Hertzsprung–Russell diagram, or specifically a color–magnitude diagram (CMD), and it
is often used by observers.[7] In cases where the stars are known to be at identical
distances such as within a star cluster, a color–magnitude diagram is often used to
describe the stars of the cluster with a plot in which the vertical axis is the apparent
magnitude of the stars. For cluster members, by assumption there is a single additive
constant difference between their apparent and absolute magnitudes, called the
distance modulus, for all of that cluster of stars. Early studies of nearby open clusters
(like the Hyades and Pleiades) by Hertzsprung and Rosenberg produced the first
CMDs, a few years before Russell's influential synthesis of the diagram collecting data
for all stars for which absolute magnitudes could be determined. [3][5]
Another form of the diagram plots the effective surface temperature of the star on one
axis and the luminosity of the star on the other, almost invariably in a log-log plot.
Theoretical calculations of stellar structure and the evolution of stars produce plots that
match those from observations. This type of diagram could be called temperatureluminosity diagram, but this term is hardly ever used; when the distinction is made, this
form is called the theoretical Hertzsprung–Russell diagram instead. A peculiar
characteristic of this form of the H–R diagram is that the temperatures are plotted from
high temperature to low temperature, which aids in comparing this form of the H–R
diagram with the observational form.
Although the two types of diagrams are similar, astronomers make a sharp distinction
between the two. The reason for this distinction is that the exact transformation from
one to the other is not trivial. To go between effective temperature and color requires a
color–temperature relation, and constructing that is difficult; it is known to be a function
of stellar composition and can be affected by other factors like stellar rotation. When
converting luminosity or absolute bolometric magnitude to apparent or absolute visual
magnitude, one requires a bolometric correction, which may or may not come from the
same source as the color–temperature relation. One also needs to know the distance to
the observed objects (i.e., the distance modulus) and the effects of interstellar
obscuration, both in the color (reddening) and in the apparent magnitude (where the
effect is called "extinction"). Color distortion (including reddening) and extinction
(obscuration) are also apparent in stars having significant circumstellar dust. The ideal
of direct comparison of theoretical predictions of stellar evolution to observations thus
has additional uncertainties incurred in the conversions between theoretical quantities
and observations.
Interpretation[edit]
An HR diagram with the instability strip and its components highlighted
Most of the stars occupy the region in the diagram along the line called the main
sequence. During the stage of their lives in which stars are found on the main sequence
line, they are fusing hydrogen in their cores. The next concentration of stars is on the
horizontal branch (helium fusion in the core and hydrogen burning in a shell surrounding
the core). Another prominent feature is the Hertzsprung gap located in the region
between A5 and G0 spectral type and between +1 and − 3 absolute magnitudes (i.e.,
between the top of the main sequence and the giants in the horizontal branch). RR
Lyrae variable stars can be found in the left of this gap on a section of the diagram
called the instability strip. Cepheid variables also fall on the instability strip, at higher
luminosities.
The H-R diagram can be used by scientists to roughly measure how far away a star
cluster or galaxy is from Earth. This can be done by comparing the apparent
magnitudes of the stars in the cluster to the absolute magnitudes of stars with known
distances (or of model stars). The observed group is then shifted in the vertical
direction, until the two main sequences overlap. The difference in magnitude that was
bridged in order to match the two groups is called the distance modulus and is a direct
measure for the distance (ignoring extinction). This technique is known as main
sequence fitting and is a type of spectroscopic parallax. Not only the turn-off in the main
sequence can be used, but also the tip of the red giant branch stars.[8][9]
The diagram seen by ESA's Gaia mission[edit]
Hertzsprung-Russell diagram showing only white dwarfs with data from ESA's Gaia mission
Part of the diagram from ESA's Gaia. The dark line likely represents the transition from partly convective to
fully convective red dwarfs
ESA's Gaia mission showed several features in the diagram that were either not known
or that were suspected to exist. It found a gap in the main sequence that appears for Mdwarfs and that is explained with the transition from a partly convective core to a fully
convective core.[10][11] For white dwarfs the diagram shows several features. Two main
concentrations appear in this diagram following the cooling sequence of white dwarfs
that are explained with the atmospheric composition of white dwarfs, especially
hydrogen versus helium dominated atmospheres of white dwarfs.[12] A third
concentration is explained with core crystallization of the white dwarfs interior. This
releases energy and delays the cooling of white dwarfs. [13][14]
Role in the development of stellar physics[edit]
See also: Stellar nucleosynthesis
HR diagrams for two open clusters, M67 and NGC 188, showing the main-sequence turn-off at different
ages
Contemplation of the diagram led astronomers to speculate that it might demonstrate
stellar evolution, the main suggestion being that stars collapsed from red giants to dwarf
stars, then moving down along the line of the main sequence in the course of their
lifetimes. Stars were thought therefore to radiate energy by converting gravitational
energy into radiation through the Kelvin–Helmholtz mechanism. This mechanism
resulted in an age for the Sun of only tens of millions of years, creating a conflict over
the age of the Solar System between astronomers, and biologists and geologists who
had evidence that the Earth was far older than that. This conflict was only resolved in
the 1930s when nuclear fusion was identified as the source of stellar energy.
Following Russell's presentation of the diagram to a meeting of the Royal Astronomical
Society in 1912, Arthur Eddington was inspired to use it as a basis for developing ideas
on stellar physics. In 1926, in his book The Internal Constitution of the Stars he
explained the physics of how stars fit on the diagram. [15] The paper anticipated the later
discovery of nuclear fusion and correctly proposed that the star's source of power was
the combination of hydrogen into helium, liberating enormous energy. This was a
particularly remarkable intuitive leap, since at that time the source of a star's energy was
still unknown, thermonuclear energy had not been proven to exist, and even that stars
are largely composed of hydrogen (see metallicity), had not yet been discovered.
Eddington managed to sidestep this problem by concentrating on the thermodynamics
of radiative transport of energy in stellar interiors.[16] Eddington predicted that dwarf
stars remain in an essentially static position on the main sequence for most of their
lives. In the 1930s and 1940s, with an understanding of hydrogen fusion, came an
evidence-backed theory of evolution to red giants following which were speculated
cases of explosion and implosion of the remnants to white dwarfs. The term supernova
nucleosynthesis is used to describe the creation of elements during the evolution and
explosion of a pre-supernova star, a concept put forth by Fred Hoyle in 1954.[17] The
pure mathematical quantum mechanics and classical mechanical models of stellar
processes enable the Hertzsprung–Russell diagram to be annotated with known
conventional paths known as stellar sequences—there continue to be added rarer and
more anomalous examples as more stars are analysed and mathematical models
considered.
The Hertzsprung-Russell Diagram
One of the most useful and powerful plots in astrophysics is the Hertzsprung-Russell diagram (hereafter
called the H-R diagram). It originated in 1911 when the Danish astronomer, Ejnar Hertzsprung, plotted the
absolute magnitude of stars against their colour (hence effective temperature). Independently in 1913 the
American astronomer Henry Norris Russell used spectral class against absolute magnitude. Their resultant
plots showed that the relationship between temperature and luminosity of a star was not random but instead
appeared to fall into distinct groups. These are seen in the H-R diagram below. It has a few specific stars
included in the plot but otherwise just shows the main regions.
The majority of stars, including our Sun, are found along a region called the Main Sequence. Main
Sequence stars vary widely in effective temperature but the hotter they are, the more luminous they are,
hence the main sequence tends to follow a band going from the bottom right of the diagram to the top left.
These stars are fusing hydrogen to helium in their cores. Stars spend the bulk of their existence as main
sequence stars. Other major groups of stars found on the H-R diagram are the giants and supergiants;
luminous stars that have evolved off the main sequence, and the white dwarfs. Whilst each of these types
is discussed in detail in later pages we can use their positions on the H-R diagram to infer some of their
properties.
Using the H-R Diagram to Infer Stellar Properties
Let us look at the cool M-class stars as an example. If we look at the H-R diagram below we can see that
in fact there are three main groups of these stars.
At the bottom-right of the diagram we can see two named stars, Proxima Centauri and Barnard's Star.
These are both cool (approximately 2,500 K) and dim (absolute magnitudes of about 13, only about
1/10,000 the luminosity of our Sun). Following the broad band straight up we come across Mira, also cool
but much more luminous. Travelling further up we come across Antares and Betelgeuse. Again these stars
are cool but they are extremely luminous, almost 10,000× as luminous as the Sun. Why do these three
groups differ so much in luminosity?
The answer to this question depends upon the Stefan-Boltzmann relationship. You may recall from equation
4.4 that the energy emitted per unit surface area per second is simply a function of the fourth power of
temperature, that is:
l ≈ σT4 (4.4)
If two stars have the same effective temperature they each have the same power output per square metre
of surface area. As the H-R diagram however shows that one is much more luminous than the other it must
have a greater total power output therefore must have a much greater surface area - the more luminous
star is bigger. We can see this from the full expression for luminosity in equation 4.6:
L ≈ 4πR2σT4 (4.6)
The difference between the three groups of M-class stars is thus a difference in size. This is acknowledged
by the names given to each of the groups. The most luminous ones are called supergiants (luminosity
classes I and II), the luminous ones are called giants (luminosity class III) and the dim ones are part of the
main sequence (luminosity class V) though historically the term dwarf stars was applied to this group.
If we look at the vertical band on the H-R diagram for hotter stars around type A spectral class we see a
similar pattern:
In this case the supergiants Rigel and Deneb have the same effective temperature as Sirius but have
extremely high luminosities. They have large radii than Sirius hence greater surface areas and higher
luminosities. Sirius is a main sequence star but because it is hotter than the red main sequence Barnard's
Star it is much more luminous than it. If you follow the pink band for hot stars down to the bottom of the HR diagram you will notice that it intersects another group of stars that includes Procyon B. These are the
white dwarfs. They are very hot (about 10,000 K or hotter) therefore emit a lot of energy per second for
each square metre of their surface. The fact that they are so dim however, means that they must be
extremely small and have a very low surface area. The terminology of white dwarf must not be confused
with the old-fashioned term of dwarf stars that was applied to main sequence stars. White dwarfs are very
different objects to main sequence stars as we shall see in a later page. Technically they have a luminosity
class of wd. Simple calculations provide a size for white dwarfs roughly that of our Earth, less than 1/100
that of the Sun.
If we compare the dimmest stars on the H-R diagram we can also make some inferences. The following
diagram shows the lower region of the H-R diagram.
Procyon B and Barnard's Star share the same low luminosity with an absolute magnitude of about +13.
Procyon B however is much hotter than Barnard's Star thus emits much more energy per second per unit
surface area. Given that they have the same total power output Procyon B must therefore have less surface
area than Barnard's Star, that is its radius is smaller.
Axes on the H-R Diagram
This points to an interesting and sometimes confusing feature of the H-R diagram - the scales on the axes.
Unlike the height/mass plot earlier in this section, the effective temperature does not increase as it goes
from left to right, it actually decreases, that is the highest temperature is on the left-hand side. If colour
index (B-V) rather than effective temperature is used then it goes from negative (blue) on the left to positive
(red) on the right. A third alternative along the horizontal axis is to use spectral class. Of course, all three
quantities are essentially showing the same thing. The diagram below shows the possible axes for an H-R
diagram.
The vertical axis displays the luminosity of the stars. This is either as a ratio compared with that of the Sun
or as absolute magnitude, M. One point to be careful of when using absolute magnitude is to remember
that the lower or more negative the absolute magnitude, the more luminous the star. The brightest stars
therefore appear at the top of the H-R diagram with the vertical axis having the most negative value of M
at the top.
In some circumstances, such as when plotting stars in a specific open or globular cluster, apparent
magnitude, m, or V, rather than absolute magnitude may be used. This is valid as all the stars in the cluster
are effectively at the same distance away from us hence any differences in apparent magnitude are due to
actual difference in luminosity or M. Diagrams where V is plotted against colour index, B-V, are also known
as colour-magnitude diagrams.
henyey track
The Henyey track is a path taken by pre-main-sequence stars with masses greater than 0.5 solar
masses in the Hertzsprung–Russell diagram after the end of the Hayashi track. The astronomer
Louis G. Henyey and his colleagues in the 1950s showed that the pre-main-sequence star can
remain in radiative equilibrium throughout some period of its contraction to the main sequence.
The Henyey track is characterized by a slow collapse in near hydrostatic equilibrium, approaching
the main sequence almost horizontally in the Hertzsprung–Russell diagram (i.e. the luminosity
remains almost constant).[1]
hayashi track
The Hayashi track is a luminosity–temperature relationship obeyed by infant stars of less than 3 M☉
in the pre-main-sequence phase (PMS phase) of stellar evolution. It is named after Japanese
astrophysicist Chushiro Hayashi. On the Hertzsprung–Russell diagram, which plots luminosity
against temperature, the track is a nearly vertical curve. After a protostar ends its phase of rapid
contraction and becomes a T Tauri star, it is extremely luminous. The star continues to contract, but
much more slowly. While slowly contracting, the star follows the Hayashi track downwards,
becoming several times less luminous but staying at roughly the same surface temperature, until
either a radiative zone develops, at which point the star starts following the Henyey track, or nuclear
fusion begins, marking its entry onto the main sequence.
The shape and position of the Hayashi track on the Hertzsprung–Russell diagram depends on the
star's mass and chemical composition. For solar-mass stars, the track lies at a temperature of
roughly 4000 K. Stars on the track are nearly fully convective and have their opacity dominated by
hydrogen ions. Stars less than 0.5 M☉ are fully convective even on the main sequence, but their
opacity begins to be dominated by Kramers' opacity law after nuclear fusion begins, thus moving
them off the Hayashi track. Stars between 0.5 and 3 M☉ develop a radiative zone prior to reaching
the main sequence. Stars between 3 and 10 M ☉ are fully radiative at the beginning of the pre-mainsequence. Even heavier stars are born onto the main sequence, with no PMS evolution. [1]
At an end of a low- or intermediate-mass star's life, the star follows an analogue of the Hayashi
track, but in reverse—it increases in luminosity, expands, and stays at roughly the same
temperature, eventually becoming a red giant.
T Tauri star
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This article is about the type of variable star. For the particular variable star called "T Tauri",
see T Tauri.
Artist's impression of a T Tauri star with a circumstellar accretion disc
Star formation
Object classes
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Interstellar medium
Molecular cloud
Bok globule
Dark nebula
Young stellar object
Protostar
Pre-main-sequence star
T Tauri star
Herbig Ae/Be star
Herbig–Haro object
Theoretical concepts
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Accretion
Initial mass function
Jeans instability
Kelvin–Helmholtz mechanism
Nebular hypothesis
Planetary migration
v
t
e
T Tauri stars (TTS) are a class of variable stars that are less than about ten million years old. [1]
This class is named after the prototype, T Tauri, a young star in the Taurus star-forming region.
They are found near molecular clouds and identified by their optical variability and strong
chromospheric lines. T Tauri stars are pre-main-sequence stars in the process of contracting to
the main sequence along the Hayashi track, a luminosity–temperature relationship obeyed by
infant stars of less than 3 solar masses (M☉) in the pre-main-sequence phase of stellar evolution.
It ends when a star of 0.5 M☉ or larger develops a radiative zone, or when a smaller star
commences nuclear fusion on the main sequence.
History[edit]
While T Tauri itself was discovered in 1852, the T Tauri class of stars were initially defined by
Alfred Harrison Joy in 1945.[2]
Characteristics[edit]
T Tauri stars comprise the youngest visible F, G, K and M spectral type stars (<2 M☉). Their
surface temperatures are similar to those of main-sequence stars of the same mass, but they are
significantly more luminous because their radii are larger. Their central temperatures are too low
for hydrogen fusion. Instead, they are powered by gravitational energy released as the stars
contract, while moving towards the main sequence, which they reach after about 100 million
years. They typically rotate with a period between one and twelve days, compared to a month for
the Sun, and are very active and variable.
There is evidence of large areas of starspot coverage, and they have intense and variable X-ray
and radio emissions (approximately 1000 times that of the Sun). Many have extremely powerful
stellar winds; some eject gas in high-velocity bipolar jets. Another source of brightness
variability are clumps (protoplanets and planetesimals) in the disk surrounding T Tauri stars.
The ejection of a bubble of hot gas from XZ Tauri, a binary system of T Tauri stars. The scale
is much larger than that of the Solar System.
Their spectra show a higher lithium abundance than the Sun and other main-sequence stars
because lithium is destroyed at temperatures above 2,500,000 K. From a study of lithium
abundances in 53 T Tauri stars, it has been found that lithium depletion varies strongly with size,
suggesting that "lithium burning" by the p-p chain during the last highly convective and unstable
stages during the later pre–main sequence phase of the Hayashi contraction may be one of the
main sources of energy for T Tauri stars. Rapid rotation tends to improve mixing and increase
the transport of lithium into deeper layers where it is destroyed. T Tauri stars generally increase
their rotation rates as they age, through contraction and spin-up, as they conserve angular
momentum. This causes an increased rate of lithium loss with age. Lithium burning will also
increase with higher temperatures and mass, and will last for at most a little over 100 million
years.
The p-p chain for lithium burning is as follows
It will not occur in stars with less than sixty times the mass of Jupiter (MJ). In this way, the rate
of lithium depletion can be used to calculate the age of the star.
Types[edit]
Several types of TTSs exist:[3]
● Classical T Tauri star (CTTS)
● Weak-line T Tauri star (WTTS)
○ Naked T Tauri star (NTTS), which is a subset of WTTS.
Protoplanetary discs in the Orion Nebula
Roughly half of T Tauri stars have circumstellar disks, which in this case are called
protoplanetary discs because they are probably the progenitors of planetary systems like the
Solar System. Circumstellar discs are estimated to dissipate on timescales of up to 10 million
years. Most T Tauri stars are in binary star systems. In various stages of their life, they are called
young stellar object (YSOs). It is thought that the active magnetic fields and strong solar wind of
Alfvén waves of T Tauri stars are one means by which angular momentum gets transferred from
the star to the protoplanetary disc. A T Tauri stage for the Solar System would be one means by
which the angular momentum of the contracting Sun was transferred to the protoplanetary disc
and hence, eventually to the planets.
Analogs of T Tauri stars in the higher mass range (2–8 solar masses)—A and B spectral type
pre–main-sequence stars, are called Herbig Ae/Be-type stars. More massive (>8 solar masses)
stars in pre–main sequence stage are not observed, because they evolve very quickly: when they
become visible (i.e. disperses surrounding circumstellar gas and dust cloud), the hydrogen in the
center is already burning and they are main sequence objects.
Planets[edit]
Planets around T Tauri stars include:
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HD 106906 b around an F-type star
1RXS J160929.1−210524 around a K-type star
Gliese 674 b around an M-type star
V830 Tau b around an M-type star
PDS 70b around a K-type star
Herbig–Haro object
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Hubble Space Telescope images of HH 24 (left) and HH 32 (right; top) – colourful nebulae
are typical of Herbig–Haro objects
Herbig–Haro (HH) objects are bright patches of nebulosity associated with newborn stars.
They are formed when narrow jets of partially ionised gas ejected by stars collide with nearby
clouds of gas and dust at several hundred kilometres per second. Herbig–Haro objects are
commonly found in star-forming regions, and several are often seen around a single star, aligned
with its rotational axis. Most of them lie within about one parsec (3.26 light-years) of the source,
although some have been observed several parsecs away. HH objects are transient phenomena
that last around a few tens of thousands of years. They can change visibly over timescales of a
few years as they move rapidly away from their parent star into the gas clouds of interstellar
space (the interstellar medium or ISM). Hubble Space Telescope observations have revealed the
complex evolution of HH objects over the period of a few years, as parts of the nebula fade while
others brighten as they collide with the clumpy material of the interstellar medium.
First observed in the late 19th century by Sherburne Wesley Burnham, Herbig–Haro objects
were recognised as a distinct type of emission nebula in the 1940s. The first astronomers to study
them in detail were George Herbig and Guillermo Haro, after whom they have been named.
Herbig and Haro were working independently on studies of star formation when they first
analysed the objects, and recognised that they were a by-product of the star formation process.
Although HH objects are a visible wavelength phenomena, many remain invisible at these
wavelengths due to dust and gas, and can only be detected at infrared wavelengths. Such objects,
when observed in near infrared, are called molecular hydrogen emission-line objects (MHOs).
Discovery and history of observations[edit]
The first HH object was observed in the late 19th century by Sherburne Wesley Burnham, when
he observed the star T Tauri with the 36-inch (910 mm) refracting telescope at Lick Observatory
and noted a small patch of nebulosity nearby. [1] It was thought to be an emission nebula, later
becoming known as Burnham's Nebula, and was not recognized as a distinct class of object. [2] T
Tauri was found to be a very young and variable star, and is the prototype of the class of similar
objects known as T Tauri stars which have yet to reach a state of hydrostatic equilibrium
between gravitational collapse and energy generation through nuclear fusion at their centres.[3]
Fifty years after Burnham's discovery, several similar nebulae were discovered with almost starlike appearance. Both Haro and Herbig made independent observations of several of these
objects in the Orion Nebula during the 1940s. Herbig also looked at Burnham's Nebula and
found it displayed an unusual electromagnetic spectrum, with prominent emission lines of
hydrogen, sulfur and oxygen. Haro found that all the objects of this type were invisible in
infrared light.[2]
Following their independent discoveries, Herbig and Haro met at an astronomy conference in
Tucson, Arizona in December 1949. Herbig had initially paid little attention to the objects he had
discovered, being primarily concerned with the nearby stars, but on hearing Haro's findings he
carried out more detailed studies of them. The Soviet astronomer Viktor Ambartsumian gave the
objects their name (Herbig–Haro objects, normally shortened to HH objects), and based on their
occurrence near young stars (a few hundred thousand years old), suggested they might represent
an early stage in the formation of T Tauri stars. [2] Studies of the HH objects showed they were
highly ionised, and early theorists speculated that they were reflection nebulae containing lowluminosity hot stars deep inside. But the absence of infrared radiation from the nebulae meant
there could not be stars within them, as these would have emitted abundant infrared light. In
1975 American astronomer R. D. Schwartz theorized that winds from T Tauri stars produce
shocks in the ambient medium on encounter, resulting in generation of visible light. [2] With the
discovery of the first proto-stellar jet in HH 46/47, it became clear that HH objects are indeed
shock-induced phenomena with shocks being driven by a collimated jet from protostars.[2][4]
Formation[edit]
Main articles: Star formation and Astrophysical jet
HH objects are formed when accreted material is ejected by a protostar as ionized gas along
the star's axis of rotation, as exemplified by HH 34 (right).
Stars form by gravitational collapse of interstellar gas clouds. As the collapse increases the
density, radiative energy loss decreases due to increased opacity. This raises the temperature of
the cloud which prevents further collapse, and a hydrostatic equilibrium is established. Gas
continues to fall towards the core in a rotating disk. The core of this system is called a
protostar.[5] Some of the accreting material is ejected out along the star's axis of rotation in two
jets of partially ionised gas (plasma).[6] The mechanism for producing these collimated bipolar
jets is not entirely understood, but it is believed that interaction between the accretion disk and
the stellar magnetic field accelerates some of the accreting material from within a few
astronomical units of the star away from the disk plane. At these distances the outflow is
divergent, fanning out at an angle in the range of 10−30°, but it becomes increasingly collimated
at distances of tens to hundreds of astronomical units from the source, as its expansion is
constrained.[7][8] The jets also carry away the excess angular momentum resulting from accretion
of material onto the star, which would otherwise cause the star to rotate too rapidly and
disintegrate.[8] When these jets collide with the interstellar medium, they give rise to the small
patches of bright emission which comprise HH objects.[9]
Properties[edit]
Infrared spectrum of HH 46/47 obtained by the Spitzer Space Telescope, showing the medium
in immediate vicinity of the star being silicate-rich
Electromagnetic emission from HH objects is caused when their associated shock waves collide
with the interstellar medium, creating what is called the "terminal working surfaces". [10] The
spectrum is continuous, but also has intense emission lines of neutral and ionized species. [6]
Spectroscopic observations of HH objects' doppler shifts indicate velocities of several hundred
kilometers per second, but the emission lines in those spectra are weaker than what would be
expected from such high-speed collisions. This suggests that some of the material they are
colliding with is also moving along the beam, although at a lower speed. [11][12] Spectroscopic
observations of HH objects show they are moving away from the source stars at speeds of
several hundred kilometres per second. [2][13] In recent years, the high optical resolution of the
Hubble Space Telescope has revealed the proper motion (movement along the sky plane) of
many HH objects in observations spaced several years apart. [14][15] As they move away from the
parent star, HH objects evolve significantly, varying in brightness on timescales of a few years.
Individual compact knots or clumps within an object may brighten and fade or disappear
entirely, while new knots have been seen to appear. [8][10] These arise likely because of the
precession of their jets,[16][17] along with the pulsating and intermittent eruptions from their parent
stars.[9] Faster jets catch up with earlier slower jets, creating the so-called "internal working
surfaces", where streams of gas collide and generate shock waves and consequent emissions.[18]
The total mass being ejected by stars to form typical HH objects is estimated to be of the order of
10−8 to 10−6 M☉ per year,[16] a very small amount of material compared to the mass of the stars
themselves[19] but amounting to about 1–10% of the total mass accreted by the source stars in a
year.[20] Mass loss tends to decrease with increasing age of the source. [21] The temperatures
observed in HH objects are typically about 9,000–12,000 K,[22] similar to those found in other
ionized nebulae such as H II regions and planetary nebulae.[23] Densities, on the other hand, are
higher than in other nebulae, ranging from a few thousand to a few tens of thousands of particles
per cm3,[22] compared to a few thousand particles per cm3 in most H II regions and planetary
nebulae.[23]
Densities also decrease as the source evolves over time. [21] HH objects consist mostly of
hydrogen and helium, which account for about 75% and 24% of their mass respectively. Around
1% of the mass of HH objects is made up of heavier chemical elements, including oxygen,
sulfur, nitrogen, iron, calcium and magnesium. Abundances of these elements, determined from
emission lines of respective ions, are generally similar to their cosmic abundances.[19] Many
chemical compounds found in the surrounding interstellar medium, but not present in the source
material, such as metal hydrides, are believed to have been produced by shock-induced chemical
reactions.[7] Around 20–30% of the gas in HH objects is ionized near the source star, but this
proportion decreases at increasing distances. This implies the material is ionized in the polar jet,
and recombines as it moves away from the star, rather than being ionized by later collisions. [22]
Shocking at the end of the jet can re-ionise some material, giving rise to bright "caps".[6]
Numbers and distribution[edit]
HH 2 (lower right), HH 34 (lower left), and HH 47 (top) were numbered in order of their
discovery; it is estimated that there are up to 150,000 such objects in the Milky Way.
HH objects are named approximately in order of their identification; HH 1/2 being the earliest
such objects to be identified.[24] More than a thousand individual objects are now known. [7] They
are always present in star-forming H II regions, and are often found in large groups. [9] They are
typically observed near Bok globules (dark nebulae which contain very young stars) and often
emanate from them. Several HH objects have been seen near a single energy source, forming a
string of objects along the line of the polar axis of the parent star.[7] The number of known HH
objects has increased rapidly over the last few years, but that is a very small proportion of the
estimated up to 150,000 in the Milky Way, [25] the vast majority of which are too far away to be
resolved. Most HH objects lie within about one parsec of their parent star. Many, however, are
seen several parsecs away.[21][22]
HH 46/47 is located about 450 parsecs (1,500 light-years) away from the Sun and is powered by
a class I protostar binary. The bipolar jet is slamming into the surrounding medium at a velocity
of 300 kilometers per second, producing two emission caps about 2.6 parsecs (8.5 light-years)
apart. Jet outflow is accompanied by a 0.3 parsecs (0.98 light-years) long molecular gas outflow
which is swept up by the jet itself.[7] Infrared studies by Spitzer Space Telescope have revealed a
variety of chemical compounds in the molecular outflow, including water (ice), methanol,
methane, carbon dioxide (dry ice) and various silicates.[7][26] Located around 460 parsecs (1,500
light-years) away in the Orion A molecular cloud, HH 34 is produced by a highly collimated
bipolar jet powered by a class I protostar. Matter in the jet is moving at about 220 kilometers per
second. Two bright bow shocks, separated by about 0.44 parsecs (1.4 light-years), are present on
the opposite sides of the source, followed by series of fainter ones at larger distances, making the
whole complex about 3 parsecs (9.8 light-years) long. The jet is surrounded by a 0.3 parsecs
(0.98 light-years) long weak molecular outflow near the source. [7][27]
Source stars[edit]
Thirteen-year timelapse of material ejecting from a class I protostar, forming the Herbig–Haro
object HH 34
The stars from which HH jets are emitted are all very young stars, a few tens of thousands to
about a million years old. The youngest of these are still protostars in the process of collecting
from their surrounding gases. Astronomers divide these stars into classes 0, I, II and III,
according to how much infrared radiation the stars emit. [28] A greater amount of infrared
radiation implies a larger amount of cooler material surrounding the star, which indicates it is
still coalescing. The numbering of the classes arises because class 0 objects (the youngest) were
not discovered until classes I, II and III had already been defined. [29][28]
Class 0 objects are only a few thousand years old; so young that they are not yet undergoing
nuclear fusion reactions at their centres. Instead, they are powered only by the gravitational
potential energy released as material falls onto them. [30] They mostly contain molecular outflows
with low velocities (less than a hundred kilometres per second) and weak emissions in the
outflows.[17] Nuclear fusion has begun in the cores of Class I objects, but gas and dust are still
falling onto their surfaces from the surrounding nebula, and most of their luminosity is accounted
for by gravitational energy. They are generally still shrouded in dense clouds of dust and gas,
which obscure all their visible light and as a result can only be observed at infrared and radio
wavelengths.[31] Outflows from this class are dominated by ionized species and velocities can
range up to 400 kilometres per second.[17] The in-fall of gas and dust has largely finished in Class
II objects (Classical T Tauri stars), but they are still surrounded by disks of dust and gas, and
produce weak outflows of low luminosity.[17] Class III objects (Weak-line T Tauri stars) have
only trace remnants of their original accretion disk. [28]
About 80% of the stars giving rise to HH objects are binary or multiple systems (two or more
stars orbiting each other), which is a much higher proportion than that found for low mass stars
on the main sequence. This may indicate that binary systems are more likely to generate the jets
which give rise to HH objects, and evidence suggests the largest HH outflows might be formed
when multiple–star systems disintegrate.[32] It is thought that most stars originate from multiple
star systems, but that a sizable fraction of these systems are disrupted before their stars reach the
main sequence due to gravitational interactions with nearby stars and dense clouds of gas. [32][33]
Around proto-brown dwarfs[edit]
The first and currently only (as of May 2017) large-scale Herbig-Haro object around a protobrown dwarf is HH 1165, which is connected to the proto-brown dwarf Mayrit 1701117. HH
1165 has a length of 0.8 light-years (0.26 parsec) and is located in the vicinity of the sigma
Orionis cluster. Previously only small mini-jets (≤0.03 parsec) were found around proto-brown
dwarfs.[34][35]
Infrared counterparts[edit]
HH 49/50 seen in infrared by the Spitzer Space Telescope
HH objects associated with very young stars or very massive protostars are often hidden from
view at optical wavelengths by the cloud of gas and dust from which they form. The intervening
material can diminish the visual magnitude by factors of tens or even hundreds at optical
wavelengths. Such deeply embedded objects can only be observed at infrared or radio
wavelengths,[36] usually in the frequencies of hot molecular hydrogen or warm carbon monoxide
emission.[37] In recent years, infrared images have revealed dozens of examples of "infrared HH
objects". Most look like bow waves (similar to the waves at the head of a ship), and so are
usually referred to as molecular "bow shocks". The physics of infrared bow shocks can be
understood in much the same way as that of HH objects, since these objects are essentially the
same – supersonic shocks driven by collimated jets from the opposite poles of a protostar. [38] It is
only the conditions in the jet and surrounding cloud that are different, causing infrared emission
from molecules rather than optical emission from atoms and ions. [39] In 2009 the acronym
"MHO", for Molecular Hydrogen emission-line Object, was approved for such objects, detected
in near infrared, by the International Astronomical Union Working Group on Designations, and
has been entered into their on-line Reference Dictionary of Nomenclature of Celestial Objects. [38]
The MHO catalog contains over 2000 objects.
See also[edit]
● Protoplanetary disk
Red giant
A red giant is a luminous giant star of low or intermediate mass (roughly 0.3–8 solar
masses (M☉)) in a late phase of stellar evolution. The outer atmosphere is inflated and
tenuous, making the radius large and the surface temperature around 5,000 K (4,700
°C; 8,500 °F) or lower. The appearance of the red giant is from yellow-orange to red,
including the spectral types K and M, but also class S stars and most carbon stars.
Red giants vary in the way by which they generate energy:
● most common red giants are stars on the red-giant branch (RGB) that are still
fusing hydrogen into helium in a shell surrounding an inert helium core
● red-clump stars in the cool half of the horizontal branch, fusing helium into
carbon in their cores via the triple-alpha process
● asymptotic-giant-branch (AGB) stars with a helium burning shell outside a
degenerate carbon–oxygen core, and a hydrogen-burning shell just beyond
that.
Many of the well-known bright stars are red giants, because they are luminous and
moderately common. The K0 RGB star Arcturus is 36 light-years away, and Gamma
Crucis is the nearest M-class giant at 88 light-years' distance.
Characteristics[edit]
Mira, a variable asymptotic giant branch red giant
A red giant is a star that has exhausted the supply of hydrogen in its core and has
begun thermonuclear fusion of hydrogen in a shell surrounding the core. They have
radii tens to hundreds of times larger than that of the Sun. However, their outer
envelope is lower in temperature, giving them a reddish-orange hue. Despite the lower
energy density of their envelope, red giants are many times more luminous than the
Sun because of their great size. Red-giant-branch stars have luminosities up to nearly
three thousand times that of the Sun (L☉), spectral types of K or M, have surface
temperatures of 3,000–4,000 K, and radii up to about 200 times the Sun (R☉). Stars on
the horizontal branch are hotter, with only a small range of luminosities around 75 L ☉.
Asymptotic-giant-branch stars range from similar luminosities as the brighter stars of the
red-giant branch, up to several times more luminous at the end of the thermal pulsing
phase.
Among the asymptotic-giant-branch stars belong the carbon stars of type C-N and late
C-R, produced when carbon and other elements are convected to the surface in what is
called a dredge-up.[1] The first dredge-up occurs during hydrogen shell burning on the
red-giant branch, but does not produce a large carbon abundance at the surface. The
second, and sometimes third, dredge up occurs during helium shell burning on the
asymptotic-giant branch and convects carbon to the surface in sufficiently massive
stars.
The stellar limb of a red giant is not sharply defined, contrary to their depiction in many
illustrations. Rather, due to the very low mass density of the envelope, such stars lack a
well-defined photosphere, and the body of the star gradually transitions into a 'corona'.[2]
The coolest red giants have complex spectra, with molecular lines, emission features,
and sometimes masers, particularly from thermally pulsing AGB stars.[3] Observations
have also provided evidence of a hot chromosphere above the photosphere of red
giants,[4][5][6] where investigating the heating mechanisms for the chromospheres to
form requires 3D simulations of red giants. [7]
Another noteworthy feature of red giants is that, unlike Sun-like stars whose
photospheres have a large number of small convection cells (solar granules), red-giant
photospheres, as well as those of red supergiants, have just a few large cells, the
features of which cause the variations of brightness so common on both types of
stars.[8]
Evolution[edit]
Main article: Stellar evolution § Mid-sized stars
This image tracks the life of a Sun-like star, from its birth on the left side of the frame to its evolution into a
red giant on the right after billions of years
Red giants are evolved from main-sequence stars with masses in the range from about
0.3 M☉ to around 8 M☉.[9] When a star initially forms from a collapsing molecular cloud
in the interstellar medium, it contains primarily hydrogen and helium, with trace amounts
of "metals" (in stellar structure, this simply refers to any element that is not hydrogen or
helium i.e. atomic number greater than 2). These elements are all uniformly mixed
throughout the star. The star reaches the main sequence when the core reaches a
temperature high enough to begin fusing hydrogen (a few million kelvin) and establishes
hydrostatic equilibrium. Over its main sequence life, the star slowly converts the
hydrogen in the core into helium; its main-sequence life ends when nearly all the
hydrogen in the core has been fused. For the Sun, the main-sequence lifetime is
approximately 10 billion years. More-massive stars burn disproportionately faster and so
have a shorter lifetime than less massive stars. [10]
When the star exhausts the hydrogen fuel in its core, nuclear reactions can no longer
continue and so the core begins to contract due to its own gravity. This brings additional
hydrogen into a zone where the temperature and pressure are sufficient to cause fusion
to resume in a shell around the core. The hydrogen-burning shell results in a situation
that has been described as the mirror principle; when the core within the shell contracts,
the layers of the star outside the shell must expand. The detailed physical processes
that cause this are complex, but the behaviour is necessary to satisfy simultaneous
conservation of gravitational and thermal energy in a star with the shell structure. The
core contracts and heats up due to the lack of fusion, and so the outer layers of the star
expand greatly, absorbing most of the extra energy from shell fusion. This process of
cooling and expanding is the subgiant star. When the envelope of the star cools
sufficiently it becomes convective, the star stops expanding, its luminosity starts to
increase, and the star is ascending the red-giant branch of the Hertzsprung–Russell (H–
R) diagram.[10][11]
Mira A is an old star, already shedding its outer layers into space
The evolutionary path the star takes as it moves along the red-giant branch depends on
the mass of the star. For the Sun and stars of less than about 2 M☉[12] the core will
become dense enough that electron degeneracy pressure will prevent it from collapsing
further. Once the core is degenerate, it will continue to heat until it reaches a
temperature of roughly 108 K, hot enough to begin fusing helium to carbon via the triplealpha process. Once the degenerate core reaches this temperature, the entire core will
begin helium fusion nearly simultaneously in a so-called helium flash. In more-massive
stars, the collapsing core will reach 108 K before it is dense enough to be degenerate,
so helium fusion will begin much more smoothly, and produce no helium flash. [10] The
core helium fusing phase of a star's life is called the horizontal branch in metal-poor
stars, so named because these stars lie on a nearly horizontal line in the H–R diagram
of many star clusters. Metal-rich helium-fusing stars instead lie on the so-called red
clump in the H–R diagram.[13]
An analogous process occurs when the central helium is exhausted and the star
collapses once again, causing helium in a shell to begin fusing. At the same time
hydrogen may begin fusion in a shell just outside the burning helium shell. This puts the
star onto the asymptotic giant branch, a second red-giant phase.[14] The helium fusion
results in the build up of a carbon–oxygen core. A star below about 8 M☉ will never start
fusion in its degenerate carbon–oxygen core.[12] Instead, at the end of the asymptoticgiant-branch phase the star will eject its outer layers, forming a planetary nebula with
the core of the star exposed, ultimately becoming a white dwarf. The ejection of the
outer mass and the creation of a planetary nebula finally ends the red-giant phase of the
star's evolution.[10] The red-giant phase typically lasts only around a billion years in total
for a solar mass star, almost all of which is spent on the red-giant branch. The
horizontal-branch and asymptotic-giant-branch phases proceed tens of times faster.
If the star has about 0.2 to 0.5 M☉,[12] it is massive enough to become a red giant but
does not have enough mass to initiate the fusion of helium. [9] These "intermediate" stars
cool somewhat and increase their luminosity but never achieve the tip of the red-giant
branch and helium core flash. When the ascent of the red-giant branch ends they puff
off their outer layers much like a post-asymptotic-giant-branch star and then become a
white dwarf.
Stars that do not become red giants[edit]
Very-low-mass stars are fully convective[15][16] and may continue to fuse hydrogen into
helium for up to a trillion years[17] until only a small fraction of the entire star is
hydrogen. Luminosity and temperature steadily increase during this time, just as for
more-massive main-sequence stars, but the length of time involved means that the
temperature eventually increases by about 50% and the luminosity by around 10 times.
Eventually the level of helium increases to the point where the star ceases to be fully
convective and the remaining hydrogen locked in the core is consumed in only a few
billion more years. Depending on mass, the temperature and luminosity continue to
increase for a time during hydrogen shell burning, the star can become hotter than the
Sun and tens of times more luminous than when it formed although still not as luminous
as the Sun. After some billions more years, they start to become less luminous and
cooler even though hydrogen shell burning continues. These become cool helium white
dwarfs.[9]
Very-high-mass stars develop into supergiants that follow an evolutionary track that
takes them back and forth horizontally over the H–R diagram, at the right end
constituting red supergiants. These usually end their life as a type II supernova. The
most massive stars can become Wolf–Rayet stars without becoming giants or
supergiants at all.[18][19]
Planets[edit]
This section needs to be
updated. The reason given is:
May be outdated. Please help
update this article to reflect
recent events or newly available
information. (April 2015)
Red giants with known planets: the M-type HD 208527, HD 220074 and, as of February
2014, a few tens[20] of known K-giants including Pollux, Gamma Cephei and Iota
Draconis.
Prospects for habitability[edit]
Although traditionally it has been suggested the evolution of a star into a red giant will
render its planetary system, if present, uninhabitable, some research suggests that,
during the evolution of a 1 M☉ star along the red-giant branch, it could harbor a
habitable zone for several billion years at 2 astronomical units (AU) out to around 100
million years at 9 AU out, giving perhaps enough time for life to develop on a suitable
world. After the red-giant stage, there would for such a star be a habitable zone
between 7 and 22 AU for an additional one billion years.[21] Later studies have refined
this scenario, showing how for a 1 M☉ star the habitable zone lasts from 100 million
years for a planet with an orbit similar to that of Mars to 210 million years for one that
orbits at Saturn's distance to the Sun, the maximum time (370 million years)
corresponding for planets orbiting at the distance of Jupiter. However, planets orbiting a
0.5 M☉ star in equivalent orbits to those of Jupiter and Saturn would be in the habitable
zone for 5.8 billion years and 2.1 billion years, respectively; for stars more massive than
the Sun, the times are considerably shorter.[22]
Enlargement of planets[edit]
As of June 2014, fifty giant planets have been discovered around giant stars. However,
these giant planets are more massive than the giant planets found around solar-type
stars. This could be because giant stars are more massive than the Sun (less massive
stars will still be on the main sequence and will not have become giants yet) and more
massive stars are expected to have more massive planets. However, the masses of the
planets that have been found around giant stars do not correlate with the masses of the
stars; therefore, the planets could be growing in mass during the stars' red giant phase.
The growth in planet mass could be partly due to accretion from stellar wind, although a
much larger effect would be Roche lobe overflow causing mass-transfer from the star to
the planet when the giant expands out to the orbital distance of the planet. [23]
Well known examples[edit]
Many of the well-known bright stars are red giants, because they are luminous and
moderately common. The red-giant branch variable star Gamma Crucis is the nearest
M-class giant star at 88 light-years.[24] The K1.5 red-giant branch star Arcturus is 36
light-years away.[25]
Red-giant branch[edit]
● Aldebaran (α Tauri)
● Arcturus (α Bootis)
● Gacrux (γ Crucis)
Red-clump giants[edit]
● Hamal (α Arietis)
● κ Persei
● δ Andromedae[26]
The current size of the Sun (now in the main sequence) compared to its estimated maximum size during its
red-giant phase in the future
Asymptotic giant branch[edit]
● Mira (ο Ceti)
● χ Cygni
● α Herculis
The Sun as a red giant[edit]
Main article: End of the Sun
The Sun will exit the main sequence in approximately 5 billion years and start to turn
into a red giant.[27][28] As a red giant, the Sun will grow so large that it will engulf
Mercury, Venus, and possibly Earth,[29] maybe even Mars and part or all of the asteroid
belt.
Mira variable
Mira, the prototype of the Mira variables
Mira variables /ˈmaɪrə/ (named for the prototype star Mira) are a class of pulsating
stars characterized by very red colours, pulsation periods longer than 100 days, and
amplitudes greater than one magnitude in infrared and 2.5 magnitude at visual
wavelengths.[citation needed] They are red giants in the very late stages of stellar
evolution, on the asymptotic giant branch (AGB), that will expel their outer envelopes as
planetary nebulae and become white dwarfs within a few million years.
Mira variables are stars massive enough that they have undergone helium fusion in
their cores but are less than two solar masses,[1] stars that have already lost about half
their initial mass.[citation needed] However, they can be thousands of times more luminous
than the Sun due to their very large distended envelopes. They are pulsating due to the
entire star expanding and contracting. This produces a change in temperature along
with radius, both of which factors cause the variation in luminosity. The pulsation
depends on the mass and radius of the star and there is a well-defined relationship
between period and luminosity (and colour).[2][3] The very large visual amplitudes are
not due to large luminosity changes, but due to a shifting of energy output between
infra-red and visual wavelengths as the stars change temperature during their
pulsations.[4]
Light curve of χ Cygni.
Early models of Mira stars assumed that the star remained spherically symmetric during
this process (largely to keep the computer modelling simple, rather than for physical
reasons). A recent survey of Mira variable stars found that 75% of the Mira stars which
could be resolved using the IOTA telescope are not spherically symmetric,[5] a result
which is consistent with previous images of individual Mira stars,[6][7][8] so there is now
pressure to do realistic three-dimensional modelling of Mira stars on supercomputers. [9]
Mira variables may be oxygen-rich or carbon-rich. Carbon-rich stars such as R Leporis
arise from a narrow set of conditions that override the normal tendency for AGB stars to
maintain a surplus of oxygen over carbon at their surfaces due to dredge-ups.[10]
Pulsating AGB stars such as Mira variables undergo fusion in alternating hydrogen and
helium shells, which produces periodic deep convection known as dredge-ups. These
dredge-ups bring carbon from the helium burning shell to the surface and would result in
a carbon star. However, in stars above about 4 M☉, hot bottom burning occurs. This is
when the lower regions of the convective region are hot enough for significant CNO
cycle fusion to take place which destroys much of the carbon before it can be
transported to the surface. Thus more massive AGB stars do not become carbonrich.[11]
Mira variables are rapidly losing mass and this material often forms dust shrouds
around the star. In some cases conditions are suitable for the formation of natural
masers.[12]
A small subset of Mira variables appear to change their period over time: the period
increases or decreases by a substantial amount (up to a factor of three) over the course
of several decades to a few centuries. This is believed to be caused by thermal pulses,
where the helium shell reignites the outer hydrogen shell. This changes the structure of
the star, which manifests itself as a change in period. This process is predicted to
happen to all Mira variables, but the relatively short duration of thermal pulses (a few
thousand years at most) over the asymptotic giant branch lifetime of the star (less than
a million years), means we only see it in a few of the several thousand Mira stars
known, possibly in R Hydrae.[13] Most Mira variables do exhibit slight cycle-to-cycle
changes in period, probably caused by nonlinear behaviour in the stellar envelope
including deviations from spherical symmetry. [14][15]
Mira variables are popular targets for amateur astronomers interested in variable star
observations, because of their dramatic changes in brightness. Some Mira variables
(including Mira itself) have reliable observations stretching back well over a century. [16]
List[edit]
The following list contains selected Mira variables. Unless otherwise noted, the given
magnitudes are in the V-band, and distances are from the Gaia DR2 star catalogue.[17]
Star
Mira
Brightest
Dimmest
Period
Distance[citation needed]
magnitude
magnitude
(in
days)
(in parsecs)
2.0
10.1
332
Reference
92+12
[1]
180+45
[2]
[18]
−9
Chi Cygni
3.3
14.2
408
−30
R Hydrae
3.5
10.9
380
224+56
[3]
387+81
[4]
71+5
[5]
497+22
[6]
187+9
[7]
313+40
[8]
−37
R Carinae
3.9
10.5
307
−57
R Leonis
4.4
11.3
310
−4
S Carinae
4.5
9.9
149
−20
R Cassiopeiae
4.7
13.5
430
−8
R Horologii
4.7
14.3
408
−32
R Doradus
4.8
6.3
172
U Orionis
4.8
13.0
377
55±3[18]
[9]
216+19
[10]
277+18
[11]
285+26
[12]
−16
RR Scorpii
5.0
12.4
281
−16
R Serpentis
5.2
14.4
356
−22
T Cephei
5.2
11.3
388
176+13
[13]
320+31
[14]
385+159
[15]
386+48
[16]
933+353
[17]
1078+1137
[18]
238+27
[19]
419+15
[20]
164+25
[21]
−12
R Aquarii
5.2
12.4
387
−26
R Centauri
5.3
11.8
502
[18]
−87
RR Sagittarii
5.4
14
336
−38
R Trianguli
5.4
12.6
267
−201
S Sculptoris
5.5
13.6
367
−366
R Aquilae
5.5
12.0
271
−22
R Leporis
5.5
11.7
445
−14
W Hydrae
5.6
9.6
390
−19
R Andromedae
5.8
15.2
409
242+30
[22]
431+60
[23]
767+34
[24]
215+15
[25]
709+306
[26]
575+48
[27]
1592+1009
[28]
888+47
[29]
1514+1055
[30]
671+109
[31]
−24
S Coronae
Borealis
5.8
14.1
360
−47
U Cygni
5.9
12.1
463
−31
X Ophiuchi
5.9
8.6
338
−13
RS Scorpii
6.0
13.0
319
−164
RT Sagittarii
6.0
14.1
306
−41
RU Sagittarii
6.0
13.8
240
−445
RT Cygni
6.0
13.1
190
−43
R Geminorum
6.0
14.0
370
−441
S Gruis
6.0
15.0
402
−82
V Monocerotis
6.0
13.9
341
426+50
[32]
226+32
[33]
530+28
[34]
674+47
[35]
702+60
[36]
1116+168
[37]
347+653
[38]
729+273
[39]
1553+350
[40]
477+27
[41]
−41
R Cancri
6.1
11.9
357
−25
R Virginis
6.1
12.1
146
−25
R Cygni
6.1
14.4
426
−41
R Boötis
6.2
13.1
223
−52
T Normae
6.2
13.6
244
−129
R Leonis
Minoris
6.3
13.2
372
[18]
−137
S Virginis
6.3
13.2
375
−156
R Reticuli
6.4
14.2
281
−241
S Herculis
6.4
13.8
304
−24
U Herculis
6.4
13.4
404
572+53
[42]
504+46
[43]
574+74
[44]
489+54
[45]
661+65
[46]
581+10000
[47]
−45
R Octantis
6.4
13.2
407
−39
S Pictoris
6.5
14.0
422
−59
R Ursae Majoris
6.5
13.7
302
−44
R Canum
Venaticorum
6.5
12.9
329
−54
R Normae
6.5
12.8
496
[18]
−360
T Ursae Majoris
6.6
13.5
257
1337+218
[48]
227+21
[49]
511+53
[50]
−164
R Aurigae
6.7
13.9
458
−17
RU Herculis
6.7
14.3
486
−44
R Draconis
6.7
13.2
246
662+58
[51]
843+43
[52]
374+37
[53]
353+35
[54]
298+15
[55]
1606+340
[56]
616+81
[57]
654+36
[58]
538+120
[59]
783+48
[60]
−49
V Coronae
Borealis
6.9
12.6
358
−39
T Cassiopeiae
6.9
13.0
445
−31
R Pegasi
6.9
13.8
378
−29
V Cassiopeiae
6.9
13.4
229
−14
T Pavonis
7.0
14.4
244
−239
RS Virginis
7.0
14.6
354
−64
Z Cygni
7.1
14.7
264
−33
S Orionis
7.2
13.1
434
−83
T Draconis
7.2
13.5
422
−43
UV Aurigae
7.3
10.9
394
1107+83
[61]
321+22
[62]
531+23
[63]
633+44
[64]
1117+88
[65]
352+24
[66]
458+36
[67]
318+33
[68]
2589+552
[69]
1126+86
[70]
−72
W Aquilae
7.3
14.3
490
−20
S Cephei
7.4
12.9
487
−21
R Fornacis
7.5
13.0
386
−38
RZ Pegasi
7.6
13.6
437
−76
RT Aquilae
7.6
14.5
327
−21
V Cygni
7.7
13.9
421
−31
RR Aquilae
7.8
14.5
395
−28
S Boötis
7.8
13.8
271
−387
WX Cygni
8.8
13.2
410
−75
W Draconis
8.9
15.4
279
6057+4469
[71]
1407+178
[72]
5669+10000
[73]
−1805
R Capricorni[19]
8.9
14.9
343
−142
UX Cygni
9.0
17.0
569
−2760
LL Pegasi
9.6 K
11.6 K
696
TY Cassiopeiae
10.1
19.0
645
1300[20]
[74]
1328+502
[75]
285+36
[76]
95+22
[77]
333+42
[78]
400+68
[79]
−286
IK Tauri
10.8
16.5
470
−29
CW Leonis
11.0 R
14.8 R
640
[21]
−15
TX
Camelopardalis
11.6 B
17.7 B
557
−33
LP Andromedae
15.1
17.3
614
−51
RR Lyrae variable
The RR Lyrae variable stars fall in a particular area on a Hertzsprung–Russell diagram of color versus
brightness.
RR Lyrae variables are periodic variable stars, commonly found in globular clusters.
They are used as standard candles to measure (extra) galactic distances, assisting with
the cosmic distance ladder. This class is named after the prototype and brightest
example, RR Lyrae.
They are pulsating horizontal branch stars of spectral class A or F, with a mass of
around half the Sun's. They are thought to have shed mass during the red-giant branch
phase, and were once stars of similar or slightly less mass than the Sun, around 0.8
solar masses.
In contemporary astronomy, a period-luminosity relation makes them good standard
candles for relatively nearby targets, especially within the Milky Way and Local Group.
They are also frequent subjects in the studies of globular clusters and the chemistry
(and quantum mechanics) of older stars.
Discovery and recognition[edit]
H-R diagram for globular cluster M5, with the horizontal branch marked in yellow and known RR Lyrae stars
in green
In surveys of globular clusters, these "cluster-type" variables were being rapidly identified in the mid1890s, especially by E. C. Pickering. Probably the first star definitely of RR Lyrae type found outside
a cluster was U Leporis, discovered by J. Kapteyn in 1890. The prototype star RR Lyrae was
discovered prior to 1899 by Williamina Fleming, and reported by Pickering in 1900 as
"indistinguishable from cluster-type variables".[1]
From 1915 to the 1930s, the RR Lyraes became increasingly accepted as a class of star distinct
from the classical Cepheids, due to their shorter periods, differing locations within the galaxy, and
chemical differences. RR Lyrae variables are metal-poor, Population II stars.[1]
RR Lyraes have proven difficult to observe in external galaxies because of their intrinsic faintness.
(In fact, Walter Baade's failure to find them in the Andromeda Galaxy led him to suspect that the
galaxy was much farther away than predicted, to reconsider the calibration of Cepheid variables, and
to propose the concept of stellar populations.[1]) Using the Canada-France-Hawaii Telescope in the
1980s, Pritchet & van den Bergh found RR Lyraes in Andromeda's galactic halo [2] and, more
recently, in its globular clusters.[3]
Classification[edit]
The RR Lyrae stars are conventionally divided into three main types, [1] following classification by S.I.
Bailey based on the shape of the stars' brightness curves:
●
●
●
RRab variables are the most common, making up 91% of all observed RR Lyrae, and
display the steep rises in brightness typical of RR Lyrae
RRc are less common, making up 9% of observed RR Lyrae, and have shorter periods
and more sinusoidal variation
RRd are rare, making up between <1% and 30%[4] of RR Lyrae in a system, and are
double-mode pulsators, unlike RRab and RRc
Distribution[edit]
RR Lyrae-type variable stars close to the galactic center from the VVV ESO public survey
RR Lyrae stars were formerly called "cluster variables" because of their strong (but not exclusive)
association with globular clusters; conversely, over 80% of all variables known in globular clusters
are RR Lyraes.[5] RR Lyrae stars are found at all galactic latitudes, as opposed to classical
Cepheids, which are strongly associated with the galactic plane.
Because of their old age, RR Lyraes are commonly used to trace certain populations in the Milky
Way, including the halo and thick disk.[6]
Several times as many RR Lyraes are known as all Cepheids combined; in the 1980s, about 1900
were known in globular clusters. Some estimates have about 85,000 in the Milky Way. [1]
Though binary star systems are common for typical stars, RR Lyrae are very rarely observed in
pairs.[7]
Properties[edit]
RR Lyrae stars pulse in a manner similar to Cepheid variables, but the nature and histories of these
stars is thought to be rather different. Like all variables on the Cepheid instability strip, pulsations are
caused by the κ-mechanism, when the opacity of ionised helium varies with its temperature.
RR Lyraes are old, relatively low mass, Population II stars, in common with W Virginis and BL
Herculis variables, the type II Cepheids. Classical Cepheid variables are higher mass population I
stars. RR Lyrae variables are much more common than Cepheids, but also much less luminous. The
average absolute magnitude of an RR Lyrae star is about +0.75, only 40 or 50 times brighter than
our Sun.[8] Their period is shorter, typically less than one day, sometimes ranging down to seven
hours. Some RRab stars, including RR Lyrae itself, exhibit the Blazhko effect in which there is a
conspicuous phase and amplitude modulation.[9]
Period-luminosity relationships[edit]
Typical RR Lyrae light curve
Unlike Cepheid variables, RR Lyrae variables do not follow a strict period-luminosity relationship at
visual wavelengths, although they do in the infrared K band.[10] They are normally analysed using a
period-colour-relationship, for example using a Wesenheit function. In this way, they can be used as
standard candles for distance measurements although there are difficulties with the effects of
metallicity, faintness, and blending. The effect of blending can impact RR Lyrae variables sampled
near the cores of globular clusters, which are so dense that in low-resolution observations multiple
(unresolved) stars may appear as a single target. Thus the brightness measured for that seemingly
single star (e.g., an RR Lyrae variable) is erroneously too bright, given those unresolved stars
contributed to the brightness determined. Consequently, the computed distance is wrong, and
certain researchers have argued that the blending effect can introduce a systematic uncertainty into
the cosmic distance ladder, and may bias the estimated age of the Universe and the Hubble
constant.[11][12][13]
Recent developments[edit]
The Hubble Space Telescope has identified several RR Lyrae candidates in globular clusters of the
Andromeda Galaxy[3] and has measured the distance to the prototype star RR Lyrae.[14]
The Kepler space telescope provided accurate photometric coverage of a single field at regular
intervals over an extended period. 37 known RR Lyrae variables lie within the Kepler field, including
RR Lyrae itself, and new phenomena such as period-doubling have been detected.[15]
The Gaia mission mapped 140,784 RR Lyraes, of which 50,220 were not previously known to be
variable, and for which 54,272 interstellar absorption estimates are available.[16]
Carbon star
A carbon star (C-type star) is typically an asymptotic giant branch star, a luminous red
giant, whose atmosphere contains more carbon than oxygen. The two elements
combine in the upper layers of the star, forming carbon monoxide, which consumes all
the oxygen in the atmosphere, leaving carbon atoms free to form other carbon
compounds, giving the star a "sooty" atmosphere and a strikingly ruby red appearance.
There are also some dwarf and supergiant carbon stars, with the more common giant
stars sometimes being called classical carbon stars to distinguish them.
In most stars (such as the Sun), the atmosphere is richer in oxygen than carbon.
Ordinary stars not exhibiting the characteristics of carbon stars but cool enough to form
carbon monoxide are therefore called oxygen-rich stars.
Carbon stars have quite distinctive spectral characteristics, and they were first
recognized by their spectra by Angelo Secchi in the 1860s, a pioneering time in
astronomical spectroscopy
Spectra[edit]
Echelle spectra of the carbon star UU Aurigae.
By definition carbon stars have dominant spectral Swan bands from the molecule C2. Many other
carbon compounds may be present at high levels, such as CH, CN (cyanogen), C3 and SiC2.
Carbon is formed in the core and circulated into its upper layers, dramatically changing the layers'
composition. In addition to carbon, S-process elements such as barium, technetium, and zirconium
are formed in the shell flashes and are "dredged up" to the surface.[1]
When astronomers developed the spectral classification of the carbon stars, they had considerable
difficulty when trying to correlate the spectra to the stars' effective temperatures. The trouble was
with all the atmospheric carbon hiding the absorption lines normally used as temperature indicators
for the stars.
Carbon stars also show a rich spectrum of molecular lines at millimeter wavelengths and
submillimeter wavelengths. In the carbon star CW Leonis more than 50 different circumstellar
molecules have been detected. This star is often used to search for new circumstellar molecules.
Secchi[edit]
Carbon stars were discovered already in the 1860s when spectral classification pioneer Angelo
Secchi erected the Secchi class IV for the carbon stars, which in the late 1890s were reclassified as
N class stars.[2]
Harvard[edit]
Using this new Harvard classification, the N class was later enhanced by an R class for less deeply
red stars sharing the characteristic carbon bands of the spectrum. Later correlation of this R to N
scheme with conventional spectra, showed that the R-N sequence approximately run in parallel with
c:a G7 to M10 with regards to star temperature.[3]
MK-type
R0
R3
R5
R8
Na
Nb
giant equiv.
G7-G8
K1-K2
~K2-K3
K5-M0
~M2-M3
M3-M4
Teff
4300
3900
~3700
3450
---
---
Morgan–Keenan C system[edit]
The later N classes correspond less well to the counterparting M types, because the Harvard
classification was only partially based on temperature, but also carbon abundance; so it soon
became clear that this kind of carbon star classification was incomplete. Instead a new dual number
star class C was erected so to deal with temperature and carbon abundance. Such a spectrum
measured for Y Canum Venaticorum, was determined to be C54, where 5 refers to temperature
dependent features, and 4 to the strength of the C 2 Swan bands in the spectrum. (C54 is very often
alternatively written C5,4).[4] This Morgan–Keenan C system classification replaced the older R-N
classifications from 1960 to 1993.
MK-type
C0
C1
C2
C3
C4
C5
C6
C7
giant equiv.
G4-G6
G7-G8
G9-K0
K1-K2
K3-K4
K5-M0
M1-M2
M3-M4
Teff
4500
4300
4100
3900
3650
3450
---
---
The Revised Morgan–Keenan system[edit]
The two-dimensional Morgan–Keenan C classification failed to fulfill the creators' expectations:
1. it failed to correlate to temperature measurements based on infrared,
2. originally being two-dimensional it was soon enhanced by suffixes, CH, CN, j and
other features making it impractical for en-masse analyses of foreign galaxies' carbon
star populations,
3. and it gradually occurred that the old R and N stars actually were two distinct types of
carbon stars, having real astrophysical significance.
A new revised Morgan–Keenan classification was published in 1993 by Philip Keenan, defining the
classes: C-N, C-R and C-H. Later the classes C-J and C-Hd were added.[5] This constitutes the
established classification system used today.[6]
class
spectrum
population
MV
theory
temperature example(s)
#
known
range (K)[7]
classical carbon stars
C-R: the old Harvard
medium
class R reborn:
disc pop I
are still visible at
the blue end of
the spectrum,
strong isotopic
bands, no
enhanced Ba line
0
red
giants?
5100-2800
S Cam
~25
C-N: the old Harvard
thin disc
class N reborn:
pop I
heavy diffuse
blue absorption,
sometimes
invisible in blue,
s-process
elements
enhanced over
solar abundance,
weak isotopic
bands
-2.2
AGB
3100-2600
R Lep
~90
3900-2800
Y CVn
~20
V Ari, TT
CVn
~20
HD 137613
~7
non-classical carbon stars
C-J:
very strong
isotopic bands of
unknown
unknown unknown
halo pop II
-1.8
C2 and CN
C-H: very strong CH
absorption
bright
5000-4100
giants,
mass
transfer
(all C-H:s
are
binary [8])
CHd:
hydrogen lines
and CH bands
weak or absent
thin disc
pop I
-3.5
unknown
?
Astrophysical mechanisms[edit]
Carbon stars can be explained by more than one astrophysical mechanism. Classical carbon stars
are distinguished from non-classical ones on the grounds of mass, with classical carbon stars being
the more massive.[9]
In the classical carbon stars, those belonging to the modern spectral types C-R and C-N, the
abundance of carbon is thought to be a product of helium fusion, specifically the triple-alpha process
within a star, which giants reach near the end of their lives in the asymptotic giant branch (AGB).
These fusion products have been brought to the stellar surface by episodes of convection (the socalled third dredge-up) after the carbon and other products were made. Normally this kind of AGB
carbon star fuses hydrogen in a hydrogen burning shell, but in episodes separated by 104-105 years,
the star transforms to burning helium in a shell, while the hydrogen fusion temporarily ceases. In this
phase, the star's luminosity rises, and material from the interior of the star (notably carbon) moves
up. Since the luminosity rises, the star expands so that the helium fusion ceases, and the hydrogen
shell burning restarts. During these shell helium flashes, the mass loss from the star is significant,
and after many shell helium flashes, an AGB star is transformed into a hot white dwarf and its
atmosphere becomes material for a planetary nebula.
The non-classical kinds of carbon stars, belonging to the types C-J and C-H, are believed to be
binary stars, where one star is observed to be a giant star (or occasionally a red dwarf) and the other
a white dwarf. The star presently observed to be a giant star accreted carbon-rich material when it
was still a main-sequence star from its companion (that is, the star that is now the white dwarf) when
the latter was still a classical carbon star. That phase of stellar evolution is relatively brief, and most
such stars ultimately end up as white dwarfs. These systems are now being observed a
comparatively long time after the mass transfer event, so the extra carbon observed in the present
red giant was not produced within that star.[9] This scenario is also accepted as the origin of the
barium stars, which are also characterized as having strong spectral features of carbon molecules
and of barium (an s-process element). Sometimes the stars whose excess carbon came from this
mass transfer are called "extrinsic" carbon stars to distinguish them from the "intrinsic" AGB stars
which produce the carbon internally. Many of these extrinsic carbon stars are not luminous or cool
enough to have made their own carbon, which was a puzzle until their binary nature was discovered.
The enigmatic hydrogen deficient carbon stars (HdC), belonging to the spectral class C-Hd, seems
to have some relation to R Coronae Borealis variables (RCB), but are not variable themselves and
lack a certain infrared radiation typical for RCB:s. Only five HdC:s are known, and none is known to
be binary,[10] so the relation to the non-classical carbon stars is not known.
Other less convincing theories, such as CNO cycle unbalancing and core helium flash have also
been proposed as mechanisms for carbon enrichment in the atmospheres of smaller carbon stars.
Other characteristics[edit]
Optical light image of the carbon star VX Andromedae.
Most classical carbon stars are variable stars of the long period variable types.
Observing carbon stars[edit]
Due to the insensitivity of night vision to red and a slow adaption of the red sensitive eye rods to the
light of the stars, astronomers making magnitude estimates of red variable stars, especially carbon
stars, have to know how to deal with the Purkinje effect in order not to underestimate the magnitude
of the observed star.
Generation of interstellar dust[edit]
Owing to its low surface gravity, as much as half (or more) of the total mass of a carbon star may be
lost by way of powerful stellar winds. The star's remnants, carbon-rich "dust" similar to graphite,
therefore become part of the interstellar dust.[11] This dust is believed to be a significant factor in
providing the raw materials for the creation of subsequent generations of stars and their planetary
systems. The material surrounding a carbon star may blanket it to the extent that the dust absorbs
all visible light.
Other classifications[edit]
This section needs expansion.
You can help by adding to it.
(August 2016)
Other types of carbon stars include:
●
●
●
CCS – Cool Carbon Star
CEMP – Carbon-Enhanced Metal-Poor
○ CEMP-no – Carbon-Enhanced Metal-Poor star with no enhancement of
elements produced by the r-process or s-process nucleosynthesis
○ CEMP-r – Carbon-Enhanced Metal-Poor star with an enhancement of
elements produced by r-process nucleosynthesis
○ CEMP-s – Carbon-Enhanced Metal-Poor star with an enhancement of
elements produced by s-process nucleosynthesis
○ CEMP-r/s – Carbon-Enhanced Metal-Poor star with an enhancement of
elements produced by both r-process and s-process nucleosynthesis
CGCS – Cool Galactic Carbon Star
See also[edit]
●
●
●
Barium star – Spectral class G to K giants, whose spectra indicate an overabundance of
s-process elements
S-type star – Cool giant with approximately equal quantities of carbon and oxygen in its
atmosphere
Technetium star – Star whose stellar spectrum contains absorption lines of technetium
●
Marc Aaronson – American astronomer (1950–1987), American astronomer and noted
researcher of carbon stars
Specimens:
●
●
●
R Leporis, Hind's Crimson Star: an example of a carbon star
IRC +10216, CW Leonis: the most studied carbon star, and also the brightest star in the
sky at N-band
La Superba, Y Canum Venaticorum: one of the brighter carbon stars
White dwarf
For other uses, see White dwarf (disambiguation).
"Degenerate dwarf" redirects here. Not to be confused with Degenerate star.
Image of Sirius A and Sirius B taken by the Hubble Space Telescope. Sirius B, which is a white dwarf, can
be seen as a faint point of light to the lower left of the much brighter Sirius A.
A white dwarf, also called a degenerate dwarf, is a stellar core remnant composed
mostly of electron-degenerate matter. A white dwarf is very dense: its mass is
comparable to that of the Sun, while its volume is comparable to that of Earth. A white
dwarf's faint luminosity comes from the emission of residual thermal energy; no fusion
takes place in a white dwarf.[1] The nearest known white dwarf is Sirius B, at 8.6 light
years, the smaller component of the Sirius binary star. There are currently thought to be
eight white dwarfs among the hundred star systems nearest the Sun.[2] The unusual
faintness of white dwarfs was first recognized in 1910.[3]: 1 The name white dwarf was
coined by Willem Luyten in 1922.
White dwarfs are thought to be the final evolutionary state of stars whose mass is not
high enough to become a neutron star or black hole. This includes over 97% of the
other stars in the Milky Way.[4]: §1 After the hydrogen-fusing period of a main-sequence
star of low or medium mass ends, such a star will expand to a red giant during which it
fuses helium to carbon and oxygen in its core by the triple-alpha process. If a red giant
has insufficient mass to generate the core temperatures required to fuse carbon (around
1 billion K), an inert mass of carbon and oxygen will build up at its center. After such a
star sheds its outer layers and forms a planetary nebula, it will leave behind a core,
which is the remnant white dwarf.[5] Usually, white dwarfs are composed of carbon and
oxygen (CO white dwarf). If the mass of the progenitor is between 8 and 10.5 solar
masses (M☉), the core temperature will be sufficient to fuse carbon but not neon, in
which case an oxygen–neon–magnesium (ONeMg or ONe) white dwarf may form.[6]
Stars of very low mass will be unable to fuse helium; hence, a helium white dwarf [7][8]
may form by mass loss in binary systems.
The material in a white dwarf no longer undergoes fusion reactions, so the star has no
source of energy. As a result, it cannot support itself by the heat generated by fusion
against gravitational collapse, but is supported only by electron degeneracy pressure,
causing it to be extremely dense. The physics of degeneracy yields a maximum mass
for a non-rotating white dwarf, the Chandrasekhar limit—approximately 1.44 times M☉—
beyond which it cannot be supported by electron degeneracy pressure. A carbon–
oxygen white dwarf that approaches this mass limit, typically by mass transfer from a
companion star, may explode as a type Ia supernova via a process known as carbon
detonation;[1][5] SN 1006 is thought to be a famous example.
A white dwarf is very hot when it forms, but because it has no source of energy, it will
gradually cool as it radiates its energy away. This means that its radiation, which initially
has a high color temperature, will lessen and redden with time. Over a very long time, a
white dwarf will cool and its material will begin to crystallize, starting with the core. The
star's low temperature means it will no longer emit significant heat or light, and it will
become a cold black dwarf.[5] Because the length of time it takes for a white dwarf to
reach this state is calculated to be longer than the current age of the known universe
(approximately 13.8 billion years),[9] it is thought that no black dwarfs yet exist. [1][4] The
oldest known white dwarfs still radiate at temperatures of a few thousand kelvins, which
establishes an observational limit on the maximum possible age of the universe.[10]
Discovery[edit]
See also: List of white dwarfs
The first white dwarf discovered was in the triple star system of 40 Eridani, which contains the
relatively bright main sequence star 40 Eridani A, orbited at a distance by the closer binary system of
the white dwarf 40 Eridani B and the main sequence red dwarf 40 Eridani C. The pair 40 Eridani B/C
was discovered by William Herschel on 31 January 1783.[11] In 1910, Henry Norris Russell, Edward
Charles Pickering and Williamina Fleming discovered that, despite being a dim star, 40 Eridani B
was of spectral type A, or white.[12] In 1939, Russell looked back on the discovery:[3]: 1
I was visiting my friend and generous benefactor, Prof. Edward C. Pickering. With characteristic
kindness, he had volunteered to have the spectra observed for all the stars – including comparison
stars – which had been observed in the observations for stellar parallax which Hinks and I made at
Cambridge, and I discussed. This piece of apparently routine work proved very fruitful – it led to the
discovery that all the stars of very faint absolute magnitude were of spectral class M. In conversation
on this subject (as I recall it), I asked Pickering about certain other faint stars, not on my list,
mentioning in particular 40 Eridani B. Characteristically, he sent a note to the Observatory office and
before long the answer came (I think from Mrs. Fleming) that the spectrum of this star was A. I knew
enough about it, even in these paleozoic days, to realize at once that there was an extreme
inconsistency between what we would then have called "possible" values of the surface brightness
and density. I must have shown that I was not only puzzled but crestfallen, at this exception to what
looked like a very pretty rule of stellar characteristics; but Pickering smiled upon me, and said: "It is
just these exceptions that lead to an advance in our knowledge", and so the white dwarfs entered
the realm of study!
The spectral type of 40 Eridani B was officially described in 1914 by Walter Adams.[13]
The white dwarf companion of Sirius, Sirius B, was next to be discovered. During the nineteenth
century, positional measurements of some stars became precise enough to measure small changes
in their location. Friedrich Bessel used position measurements to determine that the stars Sirius (α
Canis Majoris) and Procyon (α Canis Minoris) were changing their positions periodically. In 1844 he
predicted that both stars had unseen companions:[14]
If we were to regard Sirius and Procyon as double stars, the change of their motions would not
surprise us; we should acknowledge them as necessary, and have only to investigate their amount
by observation. But light is no real property of mass. The existence of numberless visible stars can
prove nothing against the existence of numberless invisible ones.
Bessel roughly estimated the period of the companion of Sirius to be about half a century; [14] C.A.F.
Peters computed an orbit for it in 1851.[15] It was not until 31 January 1862 that Alvan Graham Clark
observed a previously unseen star close to Sirius, later identified as the predicted companion. [15]
Walter Adams announced in 1915 that he had found the spectrum of Sirius B to be similar to that of
Sirius.[16]
In 1917, Adriaan van Maanen discovered Van Maanen's Star, an isolated white dwarf.[17] These
three white dwarfs, the first discovered, are the so-called classical white dwarfs.[3]: 2 Eventually,
many faint white stars were found which had high proper motion, indicating that they could be
suspected to be low-luminosity stars close to the Earth, and hence white dwarfs. Willem Luyten
appears to have been the first to use the term white dwarf when he examined this class of stars in
1922;[12][18][19][20][21] the term was later popularized by Arthur Stanley Eddington.[12][22] Despite
these suspicions, the first non-classical white dwarf was not definitely identified until the 1930s. 18
white dwarfs had been discovered by 1939.[3]: 3 Luyten and others continued to search for white
dwarfs in the 1940s. By 1950, over a hundred were known,[23] and by 1999, over 2,000 were
known.[24] Since then the Sloan Digital Sky Survey has found over 9,000 white dwarfs, mostly
new.[25]
Composition and structure[edit]
Although white dwarfs are known with estimated masses as low as 0.17 M ☉[26] and as high as 1.33
M☉,[27] the mass distribution is strongly peaked at 0.6 M ☉, and the majority lie between 0.5 and 0.7
M☉.[27] The estimated radii of observed white dwarfs are typically 0.8–2% the radius of the Sun;[28]
this is comparable to the Earth's radius of approximately 0.9% solar radius. A white dwarf, then,
packs mass comparable to the Sun's into a volume that is typically a million times smaller than the
Sun's; the average density of matter in a white dwarf must therefore be, very roughly, 1,000,000
times greater than the average density of the Sun, or approximately 106 g/cm 3, or 1 tonne per cubic
centimetre.[1] A typical white dwarf has a density of between 104 and 107 g/cm 3. White dwarfs are
composed of one of the densest forms of matter known, surpassed only by other compact stars such
as neutron stars, quark stars (hypothetical),[29] and black holes.
White dwarfs were found to be extremely dense soon after their discovery. If a star is in a binary
system, as is the case for Sirius B or 40 Eridani B, it is possible to estimate its mass from
observations of the binary orbit. This was done for Sirius B by 1910,[30] yielding a mass estimate of
0.94 M☉, which compares well with a more modern estimate of 1.00 M ☉.[31] Since hotter bodies
radiate more energy than colder ones, a star's surface brightness can be estimated from its effective
surface temperature, and that from its spectrum. If the star's distance is known, its absolute
luminosity can also be estimated. From the absolute luminosity and distance, the star's surface area
and its radius can be calculated. Reasoning of this sort led to the realization, puzzling to
astronomers at the time, that due to their relatively high temperature and relatively low absolute
luminosity, Sirius B and 40 Eridani B must be very dense. When Ernst Öpik estimated the density of
a number of visual binary stars in 1916, he found that 40 Eridani B had a density of over 25,000
times the Sun's, which was so high that he called it "impossible".[32] As A.S. Eddington put it later, in
1927:[33]: 50
We learn about the stars by receiving and interpreting the messages which their light brings to us.
The message of the companion of Sirius when it was decoded ran: "I am composed of material
3,000 times denser than anything you have ever come across; a ton of my material would be a little
nugget that you could put in a matchbox." What reply can one make to such a message? The reply
which most of us made in 1914 was — "Shut up. Don't talk nonsense."
As Eddington pointed out in 1924, densities of this order implied that, according to the theory of
general relativity, the light from Sirius B should be gravitationally redshifted.[22] This was confirmed
when Adams measured this redshift in 1925.[34]
Material
Density in
Notes
kg/m3
Supermassiv
e black hole
c. 1,000[35]
Critical density of a
black hole of around
108 solar masses.
Water (fresh)
1,000
At STP
Osmium
22,610
Near room
temperature
The core of
the Sun
c. 150,000
White dwarf
1 × 109[1]
Atomic nuclei
2.3 × 1017[3
6]
Does not depend
strongly on size of
nucleus
Neutron star
core
8.4 × 1016
– 1 × 10
Small black
hole
18
2 × 1030[37]
Critical density of an
Earth-mass black
hole.
Such densities are possible because white dwarf material is not composed of atoms joined by
chemical bonds, but rather consists of a plasma of unbound nuclei and electrons. There is therefore
no obstacle to placing nuclei closer than normally allowed by electron orbitals limited by normal
matter.[22] Eddington wondered what would happen when this plasma cooled and the energy to
keep the atoms ionized was no longer sufficient.[38] This paradox was resolved by R. H. Fowler in
1926 by an application of the newly devised quantum mechanics. Since electrons obey the Pauli
exclusion principle, no two electrons can occupy the same state, and they must obey Fermi–Dirac
statistics, also introduced in 1926 to determine the statistical distribution of particles which satisfy the
Pauli exclusion principle.[39] At zero temperature, therefore, electrons can not all occupy the lowestenergy, or ground, state; some of them would have to occupy higher-energy states, forming a band
of lowest-available energy states, the Fermi sea. This state of the electrons, called degenerate,
meant that a white dwarf could cool to zero temperature and still possess high energy. [38][40]
Compression of a white dwarf will increase the number of electrons in a given volume. Applying the
Pauli exclusion principle, this will increase the kinetic energy of the electrons, thereby increasing the
pressure.[38][41] This electron degeneracy pressure supports a white dwarf against gravitational
collapse. The pressure depends only on density and not on temperature. Degenerate matter is
relatively compressible; this means that the density of a high-mass white dwarf is much greater than
that of a low-mass white dwarf and that the radius of a white dwarf decreases as its mass
increases.[1]
The existence of a limiting mass that no white dwarf can exceed without collapsing to a neutron star
is another consequence of being supported by electron degeneracy pressure. Such limiting masses
were calculated for cases of an idealized, constant density star in 1929 by Wilhelm Anderson[42] and
in 1930 by Edmund C. Stoner.[43] This value was corrected by considering hydrostatic equilibrium for
the density profile, and the presently known value of the limit was first published in 1931 by
Subrahmanyan Chandrasekhar in his paper "The Maximum Mass of Ideal White Dwarfs".[44] For a
2
non-rotating white dwarf, it is equal to approximately 5.7M☉/μe , where μe is the average molecular
weight per electron of the star.[45]: eqn.(63) As the carbon-12 and oxygen-16 which predominantly
compose a carbon–oxygen white dwarf both have atomic number equal to half their atomic weight,
one should take μe equal to 2 for such a star,[40] leading to the commonly quoted value of 1.4 M ☉.
(Near the beginning of the 20th century, there was reason to believe that stars were composed
chiefly of heavy elements,[43]: 955 so, in his 1931 paper, Chandrasekhar set the average molecular
weight per electron, μe, equal to 2.5, giving a limit of 0.91 M☉.) Together with William Alfred Fowler,
Chandrasekhar received the Nobel prize for this and other work in 1983.[46] The limiting mass is now
called the Chandrasekhar limit.
If a white dwarf were to exceed the Chandrasekhar limit, and nuclear reactions did not take place,
the pressure exerted by electrons would no longer be able to balance the force of gravity, and it
would collapse into a denser object called a neutron star.[47] Carbon–oxygen white dwarfs accreting
mass from a neighboring star undergo a runaway nuclear fusion reaction, which leads to a Type Ia
supernova explosion in which the white dwarf may be destroyed, before it reaches the limiting
mass.[48]
New research indicates that many white dwarfs – at least in certain types of galaxies – may not
approach that limit by way of accretion. It has been postulated that at least some of the white dwarfs
that become supernovae attain the necessary mass by colliding with one another. It may be that in
elliptical galaxies such collisions are the major source of supernovae. This hypothesis is based on
the fact that the X-rays produced by those galaxies are 30 to 50 times less than what is expected to
be produced by type Ia supernovas of that galaxy as matter accretes on the white dwarf from its
encircling companion. It has been concluded that no more than 5 percent of the supernovae in such
galaxies could be created by the process of accretion onto white dwarfs. The significance of this
finding is that there could be two types of supernovae, which could mean that the Chandrasekhar
limit might not always apply in determining when a white dwarf goes supernova, given that two
colliding white dwarfs could have a range of masses. This in turn would confuse efforts to use
exploding white dwarfs as standard candles in determining distances.[49]
White dwarfs have low luminosity and therefore occupy a strip at the bottom of the Hertzsprung–
Russell diagram, a graph of stellar luminosity versus color or temperature. They should not be
confused with low-luminosity objects at the low-mass end of the main sequence, such as the
hydrogen-fusing red dwarfs, whose cores are supported in part by thermal pressure,[50] or the even
lower-temperature brown dwarfs.[51]
Mass–radius relationship and mass limit[edit]
The relationship between the mass and radius of white dwarfs can be derived using an energy
minimization argument. The energy of the white dwarf can be approximated by taking it to be the
sum of its gravitational potential energy and kinetic energy. The gravitational potential energy of a
unit mass piece of white dwarf, Eg, will be on the order of −G M / R, where G is the gravitational
constant, M is the mass of the white dwarf, and R is its radius.
The kinetic energy of the unit mass, Ek, will primarily come from the motion of electrons, so it will be
2
approximately N p / 2m, where p is the average electron momentum, m is the electron mass, and
N is the number of electrons per unit mass. Since the electrons are degenerate, we can estimate p
to be on the order of the uncertainty in momentum, Δp, given by the uncertainty principle, which
says that Δp Δx is on the order of the reduced Planck constant, ħ. Δx will be on the order of the
−1/3
average distance between electrons, which will be approximately n
, i.e., the reciprocal of the
cube root of the number density, n, of electrons per unit volume. Since there are N·M electrons in
3
the white dwarf, where M is the star's mass and its volume is on the order of R , n will be on the
3
order of N M / R .[40]
Solving for the kinetic energy per unit mass, Ek, we find that
The white dwarf will be at equilibrium when its total energy, Eg + Ek, is minimized. At this point, the
kinetic and gravitational potential energies should be comparable, so we may derive a rough massradius relationship by equating their magnitudes:
Solving this for the radius, R, gives[40]
Dropping N, which depends only on the composition of the white dwarf, and the universal constants
leaves us with a relationship between mass and radius:
i.e., the radius of a white dwarf is inversely proportional to the cube root of its mass.
2
Since this analysis uses the non-relativistic formula p / 2m for the kinetic energy, it is nonrelativistic. If we wish to analyze the situation where the electron velocity in a white dwarf is close to
2
the speed of light, c, we should replace p / 2m by the extreme relativistic approximation p c for the
kinetic energy. With this substitution, we find
If we equate this to the magnitude of Eg, we find that R drops out and the mass, M, is forced to
be[40]
Radius–mass relations for a model white dwarf. Mlimit is denoted as MCh
To interpret this result, observe that as we add mass to a white dwarf, its radius will decrease, so, by
the uncertainty principle, the momentum, and hence the velocity, of its electrons will increase. As
this velocity approaches c, the extreme relativistic analysis becomes more exact, meaning that the
mass M of the white dwarf must approach a limiting mass of Mlimit. Therefore, no white dwarf can
be heavier than the limiting mass Mlimit, or 1.4 M☉.
For a more accurate computation of the mass-radius relationship and limiting mass of a white dwarf,
one must compute the equation of state which describes the relationship between density and
pressure in the white dwarf material. If the density and pressure are both set equal to functions of
the radius from the center of the star, the system of equations consisting of the hydrostatic equation
together with the equation of state can then be solved to find the structure of the white dwarf at
equilibrium. In the non-relativistic case, we will still find that the radius is inversely proportional to the
cube root of the mass.[45]: eqn.(80) Relativistic corrections will alter the result so that the radius
becomes zero at a finite value of the mass. This is the limiting value of the mass – called the
Chandrasekhar limit – at which the white dwarf can no longer be supported by electron degeneracy
pressure. The graph on the right shows the result of such a computation. It shows how radius varies
with mass for non-relativistic (blue curve) and relativistic (green curve) models of a white dwarf. Both
models treat the white dwarf as a cold Fermi gas in hydrostatic equilibrium. The average molecular
weight per electron, μe, has been set equal to 2. Radius is measured in standard solar radii and
mass in standard solar masses.[45][52]
These computations all assume that the white dwarf is non-rotating. If the white dwarf is rotating, the
equation of hydrostatic equilibrium must be modified to take into account the centrifugal pseudoforce arising from working in a rotating frame.[53] For a uniformly rotating white dwarf, the limiting
mass increases only slightly. If the star is allowed to rotate nonuniformly, and viscosity is neglected,
then, as was pointed out by Fred Hoyle in 1947,[54] there is no limit to the mass for which it is
possible for a model white dwarf to be in static equilibrium. Not all of these model stars will be
dynamically stable.[55]
Radiation and cooling[edit]
The degenerate matter that makes up the bulk of a white dwarf has a very low opacity, because any
absorption of a photon requires that an electron must transition to a higher empty state, which may
not be possible as the energy of the photon may not be a match for the possible quantum states
available to that electron, hence radiative heat transfer within a white dwarf is low; it does, however,
have a high thermal conductivity. As a result, the interior of the white dwarf maintains a uniform
temperature, approximately 107 K. An outer shell of non-degenerate matter cools from
approximately 107 K to 104 K. This matter radiates roughly as a black body. A white dwarf remains
visible for a long time, as its tenuous outer atmosphere of normal matter begins to radiate at about
107 K, upon formation, while its greater interior mass is at 107 K but cannot radiate through its
normal matter shell.[56]
The visible radiation emitted by white dwarfs varies over a wide color range, from the blue-white
color of an O-type main sequence star to the red of an M-type red dwarf.[57] White dwarf effective
surface temperatures extend from over 150,000 K[24] to barely under 4,000 K.[58][59] In accordance
with the Stefan–Boltzmann law, luminosity increases with increasing surface temperature; this
surface temperature range corresponds to a luminosity from over 100 times the Sun's to under
1
⁄10,000 that of the Sun's.[59] Hot white dwarfs, with surface temperatures in excess of 30,000 K,
have been observed to be sources of soft (i.e., lower-energy) X-rays. This enables the composition
and structure of their atmospheres to be studied by soft X-ray and extreme ultraviolet
observations.[60]
White dwarfs also radiate neutrinos through the Urca process.[61]
A comparison between the white dwarf IK Pegasi B (center), its A-class companion IK Pegasi A (left) and
the Sun (right). This white dwarf has a surface temperature of 35,500 K.
As was explained by Leon Mestel in 1952, unless the white dwarf accretes matter from a companion
star or other source, its radiation comes from its stored heat, which is not replenished. [62][63]: §2.1
White dwarfs have an extremely small surface area to radiate this heat from, so they cool gradually,
remaining hot for a long time.[5] As a white dwarf cools, its surface temperature decreases, the
radiation which it emits reddens, and its luminosity decreases. Since the white dwarf has no energy
sink other than radiation, it follows that its cooling slows with time. The rate of cooling has been
estimated for a carbon white dwarf of 0.59 M☉ with a hydrogen atmosphere. After initially taking
approximately 1.5 billion years to cool to a surface temperature of 7,140 K, cooling approximately
500 more kelvins to 6,590 K takes around 0.3 billion years, but the next two steps of around 500
kelvins (to 6,030 K and 5,550 K) take first 0.4 and then 1.1 billion years.[64]: Table 2
Most observed white dwarfs have relatively high surface temperatures, between 8,000 K and 40,000
K.[25][65] A white dwarf, though, spends more of its lifetime at cooler temperatures than at hotter
temperatures, so we should expect that there are more cool white dwarfs than hot white dwarfs.
Once we adjust for the selection effect that hotter, more luminous white dwarfs are easier to
observe, we do find that decreasing the temperature range examined results in finding more white
dwarfs.[66] This trend stops when we reach extremely cool white dwarfs; few white dwarfs are
observed with surface temperatures below 4,000 K,[67] and one of the coolest so far observed, WD
0346+246, has a surface temperature of approximately 3,900 K.[58] The reason for this is that the
Universe's age is finite;[68][69] there has not been enough time for white dwarfs to cool below this
temperature. The white dwarf luminosity function can therefore be used to find the time when stars
started to form in a region; an estimate for the age of our Galactic disk found in this way is 8 billion
years.[66] A white dwarf will eventually, in many trillions of years, cool and become a non-radiating
black dwarf in approximate thermal equilibrium with its surroundings and with the cosmic
background radiation. No black dwarfs are thought to exist yet.[1]
The white dwarf cooling sequence seen by ESA's Gaia mission
Although white dwarf material is initially plasma – a fluid composed of nuclei and electrons – it was
theoretically predicted in the 1960s that at a late stage of cooling, it should crystallize, starting at its
center.[70] The crystal structure is thought to be a body-centered cubic lattice.[4][71] In 1995 it was
suggested that asteroseismological observations of pulsating white dwarfs yielded a potential test of
the crystallization theory,[72] and in 2004, observations were made that suggested approximately
90% of the mass of BPM 37093 had crystallized.[70][73][74] Other work gives a crystallized mass
fraction of between 32% and 82%.[75] As a white dwarf core undergoes crystallization into a solid
phase, latent heat is released which provides a source of thermal energy that delays its cooling. [76]
This effect was first confirmed in 2019 after the identification of a pile up in the cooling sequence of
more than 15,000 white dwarfs observed with the Gaia satellite.[77]
Low-mass helium white dwarfs (mass < 0.20 M☉), often referred to as "extremely low-mass white
dwarfs, ELM WDs" are formed in binary systems. As a result of their hydrogen-rich envelopes,
residual hydrogen burning via the CNO cycle may keep these white dwarfs hot on a long timescale.
In addition, they remain in a bloated proto-white dwarf stage for up to 2 Gyr before they reach the
cooling track.[78]
Atmosphere and spectra[edit]
Artist's impression of the WD J0914+1914 system.[79]
Although most white dwarfs are thought to be composed of carbon and oxygen, spectroscopy
typically shows that their emitted light comes from an atmosphere which is observed to be either
hydrogen or helium dominated. The dominant element is usually at least 1,000 times more abundant
than all other elements. As explained by Schatzman in the 1940s, the high surface gravity is thought
to cause this purity by gravitationally separating the atmosphere so that heavy elements are below
and the lighter above.[80][81]: §§5–6 This atmosphere, the only part of the white dwarf visible to us, is
thought to be the top of an envelope which is a residue of the star's envelope in the AGB phase and
may also contain material accreted from the interstellar medium. The envelope is believed to consist
of a helium-rich layer with mass no more than 1⁄100 of the star's total mass, which, if the atmosphere
is hydrogen-dominated, is overlain by a hydrogen-rich layer with mass approximately 1⁄10,000 of the
stars total mass.[59][82]: §§4–5
Although thin, these outer layers determine the thermal evolution of the white dwarf. The degenerate
electrons in the bulk of a white dwarf conduct heat well. Most of a white dwarf's mass is therefore at
almost the same temperature (isothermal), and it is also hot: a white dwarf with surface temperature
between 8,000 K and 16,000 K will have a core temperature between approximately 5,000,000 K
and 20,000,000 K. The white dwarf is kept from cooling very quickly only by its outer layers' opacity
to radiation.[59]
Primary and secondary features
A H lines present
B He I lines
C Continuous spectrum; no lines
O He II lines, accompanied by He I or H lines
Z
Metal lines
Q Carbon lines present
X Unclear or unclassifiable spectrum
Secondary features only
P Magnetic white dwarf with detectable polarization
H Magnetic white dwarf without detectable polarization
E Emission lines present
V Variable
The first attempt to classify white dwarf spectra appears to have been by G. P. Kuiper in 1941,[57][83]
and various classification schemes have been proposed and used since then. [84][85] The system
currently in use was introduced by Edward M. Sion, Jesse L. Greenstein and their coauthors in 1983
and has been subsequently revised several times. It classifies a spectrum by a symbol which
consists of an initial D, a letter describing the primary feature of the spectrum followed by an optional
sequence of letters describing secondary features of the spectrum (as shown in the adjacent table),
and a temperature index number, computed by dividing 50,400 K by the effective temperature. For
example:
●
●
A white dwarf with only He I lines in its spectrum and an effective temperature of 15,000
K could be given the classification of DB3, or, if warranted by the precision of the
temperature measurement, DB3.5.
A white dwarf with a polarized magnetic field, an effective temperature of 17,000 K, and
a spectrum dominated by He I lines which also had hydrogen features could be given the
classification of DBAP3.
The symbols "?" and ":" may also be used if the correct classification is uncertain.[24][57]
White dwarfs whose primary spectral classification is DA have hydrogen-dominated atmospheres.
They make up the majority, approximately 80%, of all observed white dwarfs. [59] The next class in
number is of DBs, approximately 16%.[86] The hot, above 15,000 K, DQ class (roughly 0.1%) have
carbon-dominated atmospheres.[87] Those classified as DB, DC, DO, DZ, and cool DQ have heliumdominated atmospheres. Assuming that carbon and metals are not present, which spectral
classification is seen depends on the effective temperature. Between approximately 100,000 K to
45,000 K, the spectrum will be classified DO, dominated by singly ionized helium. From 30,000 K to
12,000 K, the spectrum will be DB, showing neutral helium lines, and below about 12,000 K, the
spectrum will be featureless and classified DC.[82]: §2.4 [59]
Molecular hydrogen (H2) has been detected in spectra of the atmospheres of some white dwarfs.[88]
Metal-rich white dwarfs[edit]
Around 25–33% of white dwarfs have metal lines in their spectra, which is notable because any
heavy elements in a white dwarf should sink into the star's interior in just a small fraction of the star's
lifetime.[89] The prevailing explanation for metal-rich white dwarfs is that they have recently accreted
rocky planetesimals.[89] The bulk composition of the accreted object can be measured from the
strengths of the metal lines. For example, a 2015 study of the white dwarf Ton 345 concluded that its
metal abundances were consistent with those of a differentiated, rocky planet whose mantle had
been eroded by the host star's wind during its asymptotic giant branch phase.[90]
Magnetic field[edit]
Magnetic fields in white dwarfs with a strength at the surface of c. 1 million gauss (100 teslas) were
predicted by P. M. S. Blackett in 1947 as a consequence of a physical law he had proposed which
stated that an uncharged, rotating body should generate a magnetic field proportional to its angular
momentum.[91] This putative law, sometimes called the Blackett effect, was never generally
accepted, and by the 1950s even Blackett felt it had been refuted. [92]: 39–43 In the 1960s, it was
proposed that white dwarfs might have magnetic fields due to conservation of total surface magnetic
flux that existed in its progenitor star phase.[93] A surface magnetic field of c. 100 gauss (0.01 T) in
the progenitor star would thus become a surface magnetic field of c. 100·100 2 = 1 million gauss (100
T) once the star's radius had shrunk by a factor of 100.[81]: §8 [94]: 484 The first magnetic white dwarf
to be discovered was GJ 742 (also known as GRW +70 8247) which was identified by James Kemp,
John Swedlund, John Landstreet and Roger Angel in 1970 to host a magnetic field by its emission of
circularly polarized light.[95] It is thought to have a surface field of approximately 300 million gauss
(30 kT).[81]: §8
Since 1970 magnetic fields have been discovered in well over 200 white dwarfs, ranging from
2 × 103 to 109 gauss (0.2 T to 100 kT).[96] The large number of presently known magnetic white
dwarfs is due to the fact that most white dwarfs are identified by low-resolution spectroscopy, which
is able to reveal the presence of a magnetic field of 1 megagauss or more. Thus the basic
identification process also sometimes results in discovery of magnetic fields. [97] It has been
estimated that at least 10% of white dwarfs have fields in excess of 1 million gauss (100 T). [98][99]
The highly magnetized white dwarf in the binary system AR Scorpii was identified in 2016 as the first
pulsar in which the compact object is a white dwarf instead of a neutron star. [100]
Chemical bonds[edit]
The magnetic fields in a white dwarf may allow for the existence of a new type of chemical bond,
perpendicular paramagnetic bonding, in addition to ionic and covalent bonds, resulting in what has
been initially described as "magnetized matter" in research published in 2012. [101]
Variability[edit]
Main article: Pulsating white dwarf
See also: Cataclysmic variables
DAV (GCVS: ZZA)
DA spectral type, having only hydrogen absorption lines in its
spectrum
DBV (GCVS: ZZB)
DB spectral type, having only helium absorption lines in its spectrum
GW Vir (GCVS:
ZZO)
Atmosphere mostly C, He and O; may be divided into DOV and PNNV
stars
Early calculations suggested that there might be white dwarfs whose luminosity varied with a period
of around 10 seconds, but searches in the 1960s failed to observe this. [81]: §7.1.1 [104] The first
variable white dwarf found was HL Tau 76; in 1965 and 1966, and was observed to vary with a
period of approximately 12.5 minutes.[105] The reason for this period being longer than predicted is
that the variability of HL Tau 76, like that of the other pulsating variable white dwarfs known, arises
from non-radial gravity wave pulsations.[81]: §7 Known types of pulsating white dwarf include the
DAV, or ZZ Ceti, stars, including HL Tau 76, with hydrogen-dominated atmospheres and the spectral
type DA;[81]: 891, 895 DBV, or V777 Her, stars, with helium-dominated atmospheres and the spectral
type DB;[59]: 3525 and GW Vir stars, sometimes subdivided into DOV and PNNV stars, with
atmospheres dominated by helium, carbon, and oxygen.[103][106] GW Vir stars are not, strictly
speaking, white dwarfs, but are stars which are in a position on the Hertzsprung-Russell diagram
between the asymptotic giant branch and the white dwarf region. They may be called pre-white
dwarfs.[103][107] These variables all exhibit small (1%–30%) variations in light output, arising from a
superposition of vibrational modes with periods of hundreds to thousands of seconds. Observation of
these variations gives asteroseismological evidence about the interiors of white dwarfs.[108]
Formation[edit]
White dwarfs are thought to represent the end point of stellar evolution for main-sequence stars with
masses from about 0.07 to 10 M☉.[4][109] The composition of the white dwarf produced will depend
on the initial mass of the star. Current galactic models suggest the Milky Way galaxy currently
contains about ten billion white dwarfs.[110]
Stars with very low mass[edit]
If the mass of a main-sequence star is lower than approximately half a solar mass, it will never
become hot enough to fuse helium in its core. It is thought that, over a lifespan that considerably
exceeds the age of the Universe (c. 13.8 billion years),[9] such a star will eventually burn all its
hydrogen, for a while becoming a blue dwarf, and end its evolution as a helium white dwarf
composed chiefly of helium-4 nuclei.[111] Due to the very long time this process takes, it is not
thought to be the origin of the observed helium white dwarfs. Rather, they are thought to be the
product of mass loss in binary systems[5][7][8][112][113][114] or mass loss due to a large planetary
companion.[115][116]
Stars with low to medium mass[edit]
If the mass of a main-sequence star is between 0.5 and 8 M☉ like our sun, its core will become
sufficiently hot to fuse helium into carbon and oxygen via the triple-alpha process, but it will never
become sufficiently hot to fuse carbon into neon. Near the end of the period in which it undergoes
fusion reactions, such a star will have a carbon–oxygen core which does not undergo fusion
reactions, surrounded by an inner helium-burning shell and an outer hydrogen-burning shell. On the
Hertzsprung–Russell diagram, it will be found on the asymptotic giant branch. It will then expel most
of its outer material, creating a planetary nebula, until only the carbon–oxygen core is left. This
process is responsible for the carbon–oxygen white dwarfs which form the vast majority of observed
white dwarfs.[112][117][118]
Stars with medium to high mass[edit]
If a star is massive enough, its core will eventually become sufficiently hot to fuse carbon to neon,
and then to fuse neon to iron. Such a star will not become a white dwarf, because the mass of its
central, non-fusing core, initially supported by electron degeneracy pressure, will eventually exceed
the largest possible mass supportable by degeneracy pressure. At this point the core of the star will
collapse and it will explode in a core-collapse supernova which will leave behind a remnant neutron
star, black hole, or possibly a more exotic form of compact star.[109][119] Some main-sequence stars,
of perhaps 8 to 10 M☉, although sufficiently massive to fuse carbon to neon and magnesium, may be
insufficiently massive to fuse neon. Such a star may leave a remnant white dwarf composed chiefly
of oxygen, neon, and magnesium, provided that its core does not collapse, and provided that fusion
does not proceed so violently as to blow apart the star in a supernova.[120][121] Although a few white
dwarfs have been identified which may be of this type, most evidence for the existence of such
comes from the novae called ONeMg or neon novae. The spectra of these novae exhibit
abundances of neon, magnesium, and other intermediate-mass elements which appear to be only
explicable by the accretion of material onto an oxygen-neon-magnesium white dwarf.[6][122][123]
Type Iax supernova[edit]
Type Iax supernova, that involve helium accretion by a white dwarf, have been proposed to be a
channel for transformation of this type of stellar remnant. In this scenario, the carbon detonation
produced in a Type Ia supernova is too weak to destroy the white dwarf, expelling just a small part of
its mass as ejecta, but produces an asymmetric explosion that kicks the star, often known as a
zombie star, to high speeds of a hypervelocity star. The matter processed in the failed detonation is
re-accreted by the white dwarf with the heaviest elements such as iron falling to its core where it
accumulates.[124] These iron-core white dwarfs would be smaller than the carbon–oxygen kind of
similar mass and would cool and crystallize faster than those.[125]
Fate[edit]
Artist's concept of white dwarf aging
Internal structures of white dwarfs. To the left is a newly formed white dwarf, in the center is a cooling and
crystallizing white dwarf, and the right is a black dwarf.
A white dwarf is stable once formed and will continue to cool almost indefinitely, eventually to
become a black dwarf. Assuming that the Universe continues to expand, it is thought that in 1019 to
1020 years, the galaxies will evaporate as their stars escape into intergalactic space.[126]: §IIIA White
dwarfs should generally survive galactic dispersion, although an occasional collision between white
dwarfs may produce a new fusing star or a super-Chandrasekhar mass white dwarf which will
explode in a Type Ia supernova.[126]: §§IIIC, IV The subsequent lifetime of white dwarfs is thought to
be on the order of the hypothetical lifetime of the proton, known to be at least 1034–1035 years.
Some grand unified theories predict a proton lifetime between 1030 and 1036 years. If these theories
are not valid, the proton might still decay by complicated nuclear reactions or through quantum
gravitational processes involving virtual black holes; in these cases, the lifetime is estimated to be no
more than 10200 years. If protons do decay, the mass of a white dwarf will decrease very slowly with
time as its nuclei decay, until it loses enough mass to become a nondegenerate lump of matter, and
finally disappears completely.[126]: §IV
A white dwarf can also be cannibalized or evaporated by a companion star, causing the white dwarf
to lose so much mass that it becomes a planetary mass object. The resultant object, orbiting the
former companion, now host star, could be a helium planet or diamond planet.[127][128]
Debris disks and planets[edit]
Artist's impression of debris around a white dwarf [129]
Comet falling into white dwarf (artist's impression)[130]
A white dwarf's stellar and planetary system is inherited from its progenitor star and may interact with
the white dwarf in various ways. Infrared spectroscopic observations made by NASA's Spitzer Space
Telescope of the central star of the Helix Nebula suggest the presence of a dust cloud, which may
be caused by cometary collisions. It is possible that infalling material from this may cause X-ray
emission from the central star.[131][132] Similarly, observations made in 2004 indicated the presence
of a dust cloud around the young (estimated to have formed from its AGB progenitor about 500
million years ago) white dwarf G29-38, which may have been created by tidal disruption of a comet
passing close to the white dwarf.[133] Some estimations based on the metal content of the
atmospheres of the white dwarfs consider that at least 15% of them may be orbited by planets
and/or asteroids, or at least their debris.[134] Another suggested idea is that white dwarfs could be
orbited by the stripped cores of rocky planets, that would have survived the red giant phase of their
star but losing their outer layers and, given those planetary remnants would likely be made of
metals, to attempt to detect them looking for the signatures of their interaction with the white dwarf's
magnetic field.[135] Other suggested ideas of how white dwarfs are polluted with dust involve the
scattering of asteroids by planets[136][137][138] or via planet-planet scattering.[139] Liberation of
exomoons from their host planet could cause white dwarf pollution with dust. Either the liberation
could cause asteroids to be scattered towards the white dwarf or the exomoon could be scattered
into the Roche-Radius of the white dwarf.[140] The mechanism behind the pollution of white dwarfs in
binaries was also explored as these systems are more likely to lack a major planet, but this idea
cannot explain the presence of dust around single white dwarfs.[141] While old white dwarfs show
evidence of dust accretion, white dwarfs older than ~1 billion years or >7000 K with dusty infrared
excess were not detected[142] until the discovery of LSPM J0207+3331 in 2018, which has a cooling
age of ~3 billion years. The white dwarf shows two dusty components that are being explained with
two rings with different temperatures.[143]
Exoplanet orbits WD 1856+534
(NASA; video; 2:10)
There is a planet in the white dwarf–pulsar binary system PSR B1620-26.
There are two circumbinary planets around the white dwarf–red dwarf binary NN Serpentis.
The metal-rich white dwarf WD 1145+017 is the first white dwarf observed with a disintegrating
minor planet which transits the star.[144][145] The disintegration of the planetesimal generates a
debris cloud which passes in front of the star every 4.5 hours, causing a 5-minute-long fade in the
star's optical brightness.[145] The depth of the transit is highly variable.[145]
The white dwarf WD 0145+234 shows brightening in the mid-infrared, seen in NEOWISE data. The
brightening is not seen before 2018. It is interpreted as the tidal disruption of an exoasteroid, the first
time such an event has been observed.[146]
WD 0806-661 has a Y-dwarf that orbits the white dwarf in a wide orbit with a projected distance of
2500 astronomical units. Considering the low mass and the wide orbit of this object, WD 0806-661 B
can be interpreted as either a sub-brown dwarf or a directly imaged exoplanet.
WD J0914+1914 is the first single white dwarf star found to have a giant planet orbiting it. The giant
planet is being evaporated by the strong ultraviolet radiation of the hot white dwarf. Part of the
evaporated material is being accreted in a gaseous disk around the white dwarf. The weak hydrogen
line as well as other lines in the spectrum of the white dwarf revealed the presence of the giant
planet.[147]
In September 2020, astronomers reported the discovery, for the first time, of a very massive Jupitersized planet, named WD 1856 b, closely orbiting, every 36 hours, a white dwarf, named WD
1856+534.[148][149][150]
Habitability[edit]
It has been proposed that white dwarfs with surface temperatures of less than 10,000 Kelvins could
harbor a habitable zone at a distance of c. 0.005 to 0.02 AU that would last upwards of 3 billion
years. This is so close that any habitable planets would be tidally locked. The goal is to search for
transits of hypothetical Earth-like planets that could have migrated inward and/or formed there. As a
white dwarf has a size similar to that of a planet, these kinds of transits would produce strong
eclipses.[151] Newer research casts some doubts on this idea, given that the close orbits of those
hypothetical planets around their parent stars would subject them to strong tidal forces that could
render them uninhabitable by triggering a greenhouse effect.[152] Another suggested constraint to
this idea is the origin of those planets. Leaving aside formation from the accretion disk surrounding
the white dwarf, there are two ways a planet could end in a close orbit around stars of this kind: by
surviving being engulfed by the star during its red giant phase, and then spiralling inward, or inward
migration after the white dwarf has formed. The former case is implausible for low-mass bodies, as
they are unlikely to survive being absorbed by their stars. In the latter case, the planets would have
to expel so much orbital energy as heat, through tidal interactions with the white dwarf, that they
would likely end as uninhabitable embers.[153]
Binary stars and novae[edit]
The merger process of two co-orbiting white dwarfs produces gravitational waves
If a white dwarf is in a binary star system and is accreting matter from its companion, a variety of
phenomena may occur, including novae and Type Ia supernovae. It may also be a super-soft x-ray
source if it is able to take material from its companion fast enough to sustain fusion on its
surface.[154] On the other hand, phenomena in binary systems such as tidal interaction and star-disc
interaction, moderated by magnetic fields or not, act on the rotation of accreting white dwarfs. In fact,
the fastest-spinning, securely known white dwarfs, are members of binary systems (being the white
dwarf in CTCV J2056-3014 the fastest one). [155] A close binary system of two white dwarfs can
radiate energy in the form of gravitational waves, causing their mutual orbit to steadily shrink until the
stars merge.[156][157]
Type Ia supernovae[edit]
Main article: Type Ia supernova
The mass of an isolated, nonrotating white dwarf cannot exceed the Chandrasekhar limit of ~1.4 M☉.
This limit may increase if the white dwarf is rotating rapidly and nonuniformly. [158] White dwarfs in
binary systems can accrete material from a companion star, increasing both their mass and their
density. As their mass approaches the Chandrasekhar limit, this could theoretically lead to either the
explosive ignition of fusion in the white dwarf or its collapse into a neutron star.[47]
Accretion provides the currently favored mechanism called the single-degenerate model for Type Ia
supernovae. In this model, a carbon–oxygen white dwarf accretes mass and compresses its core by
pulling mass from a companion star.[48]: 14 It is believed that compressional heating of the core leads
to ignition of carbon fusion as the mass approaches the Chandrasekhar limit.[48] Because the white
dwarf is supported against gravity by quantum degeneracy pressure instead of by thermal pressure,
adding heat to the star's interior increases its temperature but not its pressure, so the white dwarf
does not expand and cool in response. Rather, the increased temperature accelerates the rate of the
fusion reaction, in a runaway process that feeds on itself. The thermonuclear flame consumes much
of the white dwarf in a few seconds, causing a Type Ia supernova explosion that obliterates the
star.[1][48][159] In another possible mechanism for Type Ia supernovae, the double-degenerate
model, two carbon–oxygen white dwarfs in a binary system merge, creating an object with mass
greater than the Chandrasekhar limit in which carbon fusion is then ignited.[48]: 14
Observations have failed to note signs of accretion leading up to Type Ia supernovae, and this is
now thought to be because the star is first loaded up to above the Chandrasekhar limit while also
being spun up to a very high rate by the same process. Once the accretion stops the star gradually
slows until the spin is no longer enough to prevent the explosion. [160]
The historical bright SN 1006 is thought to have been a type Ia supernova from a white dwarf,
possibly the merger of two white dwarfs.[161] Tycho's Supernova of 1572 was also a type Ia
supernova, and its remnant has been detected.[162]
Post-common envelope binary[edit]
Main article: Post common envelope binary
A post-common envelope binary (PCEB) is a binary consisting of a white dwarf and a closely tidallylocked red dwarf (in other cases this might be a brown dwarf instead of a red dwarf). These binaries
form when the red dwarf is engulfed in the red giant phase. As the red dwarf orbits inside the
common envelope, it is slowed down in the denser environment. This slowed orbital speed is
compensated with a decrease of the orbital distance between the red dwarf and the core of the red
giant. The red dwarf spirals inwards towards the core and might merge with the core. If this does not
happen and instead the common envelope is ejected, then the binary ends up in a close orbit,
consisting of a white dwarf and a red dwarf. This type of binary is called a post-common envelope
binary. The evolution of the PCEB continues as the two dwarf stars orbit closer and closer due to
magnetic braking and by releasing gravitational waves. The binary might evolve at some point into a
cataclysmic variable, and therefore post-common envelope binaries are sometimes called precataclysmic variables.
Cataclysmic variables[edit]
Main article: Cataclysmic variable star
Before accretion of material pushes a white dwarf close to the Chandrasekhar limit, accreted
hydrogen-rich material on the surface may ignite in a less destructive type of thermonuclear
explosion powered by hydrogen fusion. These surface explosions can be repeated as long as the
white dwarf's core remains intact. This weaker kind of repetitive cataclysmic phenomenon is called a
(classical) nova. Astronomers have also observed dwarf novae, which have smaller, more frequent
luminosity peaks than the classical novae. These are thought to be caused by the release of
gravitational potential energy when part of the accretion disc collapses onto the star, rather than
through a release of energy due to fusion. In general, binary systems with a white dwarf accreting
matter from a stellar companion are called cataclysmic variables. As well as novae and dwarf novae,
several other classes of these variables are known, including polars and intermediate polars, both of
which feature highly magnetic white dwarfs.[1][48][163][164] Both fusion- and accretion-powered
cataclysmic variables have been observed to be X-ray sources.[164]
Other non-pre-supernova binaries[edit]
Other non-pro-supernova binaries include binaries that consist of a main sequence star (or giant)
and a white dwarf. The binary Sirius AB is probably the most famous example. White dwarfs can
also exist as binaries or multiple star systems that only consist of white dwarfs. An example of a
resolved triple white dwarf system is WD J1953-1019, discovered with Gaia DR2 data. One
interesting field is the study of remnant planetary systems around white dwarfs. While stars are
bright and often outshine the exoplanets and brown dwarfs that orbit them, the white dwarfs are
faint. This allows astronomers to study these brown dwarfs or exoplanets in more detail. The subbrown dwarf around the white dwarf WD 0806−661 is one such example.
Nearest[edit]
Identifier
WD
Number
Distance Type
(ly)
Absolute
Mass Luminosity
magnitude
(M☉)
(L☉)
Age Objects
(Gyr)
in
system
Sirius B
0642–166
8.66
DA
11.18
0.98
0.0295
0.10
2
Procyon B
0736+053
11.46
DQZ
13.20
0.63
0.00049
1.37
2
Van
Maanen 2
0046+051
14.07
DZ
14.09
0.68
0.00017
3.30
1
LP 145141
1142–645
15.12
DQ
12.77
0.61
0.00054
1.29
1
40 Eridani
B
0413-077
16.39
DA
11.27
0.59
0.0141
0.12
3
Stein 2051
B
0426+588
17.99
DC
13.43
0.69
0.00030
2.02
2
G 240-72
1748+708
20.26
DQ
15.23
0.81
0.000085
5.69
1
Gliese
223.2
0552–041
21.01
DZ
15.29
0.82
0.000062
7.89
1
Gliese
3991
1708+437
24.23
D??
>15
0.5
<0.000086
>6
2
B[166]
Gallery[edit]
●
Illustration of rocky debris around a white dwarf
[167]
Planetary nebula
X-ray/optical composite image of the Cat's Eye Nebula (NGC 6543)
NGC 6326, a planetary nebula with glowing wisps of outpouring gas that are lit up by a binary[1] central star
A planetary nebula (PN, plural PNe), is a type of emission nebula consisting of an
expanding, glowing shell of ionized gas ejected from red giant stars late in their lives.[2]
The term "planetary nebula" is a misnomer because they are unrelated to planets. The
term originates from the planet-like round shape of these nebulae observed by
astronomers through early telescopes. The first usage may have occurred during the
1780s with the English astronomer William Herschel who described these nebulae as
resembling planets; however, as early as January 1779, the French astronomer Antoine
Darquier de Pellepoix described in his observations of the Ring Nebula, "very dim but
perfectly outlined; it is as large as Jupiter and resembles a fading planet". [3][4][5] Though
the modern interpretation is different, the old term is still used.
All planetary nebulae form at the end of the life of a star of intermediate mass, about 1-8
solar masses. It is expected that the Sun will form a planetary nebula at the end of its
life cycle.[6] They are relatively short-lived phenomena, lasting perhaps a few tens of
millennia, compared to considerably longer phases of stellar evolution.[7] Once all of the
red giant's atmosphere has been dissipated, energetic ultraviolet radiation from the
exposed hot luminous core, called a planetary nebula nucleus (P.N.N.), ionizes the
ejected material.[2] Absorbed ultraviolet light then energizes the shell of nebulous gas
around the central star, causing it to appear as a brightly coloured planetary nebula.
Planetary nebulae probably play a crucial role in the chemical evolution of the Milky
Way by expelling elements into the interstellar medium from stars where those elements
were created. Planetary nebulae are observed in more distant galaxies, yielding useful
information about their chemical abundances.
Starting from the 1990s, Hubble Space Telescope images revealed that many planetary
nebulae have extremely complex and varied morphologies. About one-fifth are roughly
spherical, but the majority are not spherically symmetric. The mechanisms that produce
such a wide variety of shapes and features are not yet well understood, but binary
central stars, stellar winds and magnetic fields may play a role.
Observations[edit]
NGC 7293, the Helix Nebula.
NGC 2392, the Eskimo Nebula.
Discovery[edit]
The first planetary nebula discovered (though not yet termed as such) was the Dumbbell Nebula in
the constellation of Vulpecula. It was observed by Charles Messier on July 12, 1764 and listed as
M27 in his catalogue of nebulous objects.[8] To early observers with low-resolution telescopes, M27
and subsequently discovered planetary nebulae resembled the giant planets like Uranus. As early as
January 1779, the French astronomer Antoine Darquier de Pellepoix described in his observations of
the Ring Nebula, "a very dull nebula, but perfectly outlined; as large as Jupiter and looks like a
fading planet".[3][4][5]
The nature of these objects remained unclear. In 1782, William Herschel, discoverer of Uranus,
found the Saturn Nebula (NGC 7009) and described it as "A curious nebula, or what else to call it I
do not know". He later described these objects as seeming to be planets "of the starry kind". [9] As
noted by Darquier before him, Herschel found that the disk resembled a planet but it was too faint to
be one. In 1785, Herschel wrote to Jérôme Lalande:
These are celestial bodies of which as yet we have no clear idea and which are perhaps of a type
quite different from those that we are familiar with in the heavens. I have already found four that
have a visible diameter of between 15 and 30 seconds. These bodies appear to have a disk that is
rather like a planet, that is to say, of equal brightness all over, round or somewhat oval, and about as
well defined in outline as the disk of the planets, of a light strong enough to be visible with an
ordinary telescope of only one foot, yet they have only the appearance of a star of about ninth
magnitude.[10]
He assigned these to Class IV of his catalogue of "nebulae", eventually listing 78 "planetary
nebulae", most of which are in fact galaxies.[11]
Herschel used the term "planetary nebulae" for these objects. The origin of this term not
known.[8][12] The label "planetary nebula" became ingrained in the terminology used by astronomers
to categorize these types of nebulae, and is still in use by astronomers today. [13][14]
Spectra[edit]
The nature of planetary nebulae remained unknown until the first spectroscopic observations were
made in the mid-19th century. Using a prism to disperse their light, William Huggins was one of the
earliest astronomers to study the optical spectra of astronomical objects.[12]
On August 29, 1864, Huggins was the first to analyze the spectrum of a planetary nebula when he
observed Cat's Eye Nebula.[8] His observations of stars had shown that their spectra consisted of a
continuum of radiation with many dark lines superimposed. He found that many nebulous objects
such as the Andromeda Nebula (as it was then known) had spectra that were quite similar. However,
when Huggins looked at the Cat's Eye Nebula, he found a very different spectrum. Rather than a
strong continuum with absorption lines superimposed, the Cat's Eye Nebula and other similar
objects showed a number of emission lines.[12] Brightest of these was at a wavelength of 500.7
nanometres, which did not correspond with a line of any known element.[15]
At first, it was hypothesized that the line might be due to an unknown element, which was named
nebulium. A similar idea had led to the discovery of helium through analysis of the Sun's spectrum in
1868.[8] While helium was isolated on Earth soon after its discovery in the spectrum of the Sun,
"nebulium" was not. In the early 20th century, Henry Norris Russell proposed that, rather than being
a new element, the line at 500.7 nm was due to a familiar element in unfamiliar conditions. [8]
Physicists showed in the 1920s that in gas at extremely low densities, electrons can occupy excited
metastable energy levels in atoms and ions that would otherwise be de-excited by collisions that
would occur at higher densities.[16] Electron transitions from these levels in nitrogen and oxygen ions
(O+, O2+ (a.k.a. O iii), and N+) give rise to the 500.7 nm emission line and others.[8] These spectral
lines, which can only be seen in very low density gases, are called forbidden lines. Spectroscopic
observations thus showed that nebulae were made of extremely rarefied gas. [17]
Planetary nebula NGC 3699 is distinguished by an irregular mottled appearance and a dark rift.[18]
Central stars[edit]
The central stars of planetary nebulae are very hot.[2] Only when a star has exhausted most of its
nuclear fuel can it collapse to a small size. Planetary nebulae are understood as a final stage of
stellar evolution. Spectroscopic observations show that all planetary nebulae are expanding. This led
to the idea that planetary nebulae were caused by a star's outer layers being thrown into space at
the end of its life.[8]
Modern observations[edit]
Towards the end of the 20th century, technological improvements helped to further the study of
planetary nebulae.[19] Space telescopes allowed astronomers to study light wavelengths outside
those that the Earth's atmosphere transmits. Infrared and ultraviolet studies of planetary nebulae
allowed much more accurate determinations of nebular temperatures, densities and elemental
abundances.[20][21] Charge-coupled device technology allowed much fainter spectral lines to be
measured accurately than had previously been possible. The Hubble Space Telescope also showed
that while many nebulae appear to have simple and regular structures when observed from the
ground, the very high optical resolution achievable by telescopes above the Earth's atmosphere
reveals extremely complex structures.[22][23]
Under the Morgan-Keenan spectral classification scheme, planetary nebulae are classified as TypeP, although this notation is seldom used in practice.[24]
Origins[edit]
Computer simulation of the formation of a planetary nebula from a star with a warped disk, showing the
complexity which can result from a small initial asymmetry.
Stars greater than 8 solar masses (M⊙) will probably end their lives in dramatic supernovae
explosions, while planetary nebulae seemingly only occur at the end of the lives of intermediate and
low mass stars between 0.8 M⊙ to 8.0 M⊙.[25] Progenitor stars that form planetary nebulae will
spend most of their lifetimes converting their hydrogen into helium in the star's core by nuclear fusion
at about 15 million K. This generated energy creates outward pressure from fusion reactions in the
core, balancing the crushing inward pressures of the star's gravity.[26] This state of equilibrium is
known as the main sequence, which can last for tens of millions to billions of years, depending on
the mass.
When the hydrogen source in the core starts to diminish, gravity starts compressing the core,
causing a rise in temperature to about 100 million K.[27] Such higher core temperatures then make
the star's cooler outer layers expand to create much larger red giant stars. This end phase causes a
dramatic rise in stellar luminosity, where the released energy is distributed over a much larger
surface area, which in fact causes the average surface temperature to be lower. In stellar evolution
terms, stars undergoing such increases in luminosity are known as asymptotic giant branch stars
(AGB).[27] During this phase, the star can lose 50 to 70% of its total mass from its stellar wind.[28]
For the more massive asymptotic giant branch stars that form planetary nebulae, whose progenitors
exceed about 3M⊙, their cores will continue to contract. When temperatures reach about 100 million
K, the available helium nuclei fuse into carbon and oxygen, so that the star again resumes radiating
energy, temporarily stopping the core's contraction. This new helium burning phase (fusion of helium
nuclei) forms a growing inner core of inert carbon and oxygen. Above it is a thin helium-burning
shell, surrounded in turn by a hydrogen-burning shell. However, this new phase lasts only 20,000
years or so, a very short period compared to the entire lifetime of the star.
The venting of atmosphere continues unabated into interstellar space, but when the outer surface of
the exposed core reaches temperatures exceeding about 30,000 K, there are enough emitted
ultraviolet photons to ionize the ejected atmosphere, causing the gas to shine as a planetary
nebula.[27]
Lifetime[edit]
The Necklace Nebula consists of a bright ring, measuring about two light-years across, dotted with dense,
bright knots of gas that resemble diamonds in a necklace. The knots glow brightly due to absorption of
ultraviolet light from the central stars.[29]
After a star passes through the asymptotic giant branch (AGB) phase, the short planetary nebula
phase of stellar evolution begins[19] as gases blow away from the central star at speeds of a few
kilometers per second. The central star is the remnant of its AGB progenitor, an electron-degenerate
carbon-oxygen core that has lost most of its hydrogen envelope due to mass loss on the AGB. [19] As
the gases expand, the central star undergoes a two-stage evolution, first growing hotter as it
continues to contract and hydrogen fusion reactions occur in the shell around the core and then
slowly cooling when the hydrogen shell is exhausted through fusion and mass loss.[19] In the second
phase, it radiates away its energy and fusion reactions cease, as the central star is not heavy
enough to generate the core temperatures required for carbon and oxygen to fuse. [8][19] During the
first phase, the central star maintains constant luminosity,[19] while at the same time it grows ever
hotter, eventually reaching temperatures around 100,000 K. In the second phase, it cools so much
that it does not give off enough ultraviolet radiation to ionize the increasingly distant gas cloud. The
star becomes a white dwarf, and the expanding gas cloud becomes invisible to us, ending the
planetary nebula phase of evolution.[19] For a typical planetary nebula, about 10,000 years [19]
passes between its formation and recombination of the resulting plasma.[8]
Role in galactic enrichment[edit]
ESO 455-10 is a planetary nebula located in the constellation of Scorpius (The Scorpion).[30]
Planetary nebulae may play a very important role in galactic evolution. Newly born stars consist
almost entirely of hydrogen and helium,[31] but as stars evolve through the asymptotic giant branch
phase,[32] they create heavier elements via nuclear fusion which are eventually expelled by strong
stellar winds.[33] Planetary nebulae usually contain larger proportions of elements such as carbon,
nitrogen and oxygen, and these are recycled into the interstellar medium via these powerful winds.
In this way, planetary nebulae greatly enrich the Milky Way and their nebulae with these heavier
elements – collectively known by astronomers as metals and specifically referred to by the metallicity
parameter Z.[34]
Subsequent generations of stars formed from such nebulae also tend to have higher metallicities.
Although these metals are present in stars in relatively tiny amounts, they have marked effects on
stellar evolution and fusion reactions. When stars formed earlier in the universe they theoretically
contained smaller quantities of heavier elements.[35] Known examples are the metal poor Population
II stars. (See Stellar population.)[36][37] Identification of stellar metallicity content is found by
spectroscopy.
Characteristics[edit]
Physical characteristics[edit]
NGC 6720, the Ring Nebula
Lemon slice nebula (IC 3568).
A typical planetary nebula is roughly one light year across, and consists of extremely rarefied gas,
with a density generally from 100 to 10,000 particles per cm 3.[38] (The Earth's atmosphere, by
comparison, contains 2.5×1019 particles per cm3.) Young planetary nebulae have the highest
densities, sometimes as high as 106 particles per cm3. As nebulae age, their expansion causes their
density to decrease. The masses of planetary nebulae range from 0.1 to 1 solar masses.[38]
Radiation from the central star heats the gases to temperatures of about 10,000 K.[39] The gas
temperature in central regions is usually much higher than at the periphery reaching 16,000–25,000
K.[40] The volume in the vicinity of the central star is often filled with a very hot (coronal) gas having
the temperature of about 1,000,000 K. This gas originates from the surface of the central star in the
form of the fast stellar wind.[41]
Nebulae may be described as matter bounded or radiation bounded. In the former case, there is not
enough matter in the nebula to absorb all the UV photons emitted by the star, and the visible nebula
is fully ionized. In the latter case, there are not enough UV photons being emitted by the central star
to ionize all the surrounding gas, and an ionization front propagates outward into the circumstellar
envelope of neutral atoms.[42]
Numbers and distribution[edit]
About 3000 planetary nebulae are now known to exist in our galaxy, [43] out of 200 billion stars. Their
very short lifetime compared to total stellar lifetime accounts for their rarity. They are found mostly
near the plane of the Milky Way, with the greatest concentration near the galactic center.[44]
Morphology[edit]
This animation shows how the two stars at the heart of a planetary nebula like Fleming 1 can control the
creation of the spectacular jets of material ejected from the object.
Only about 20% of planetary nebulae are spherically symmetric (for example, see Abell 39).[45] A
wide variety of shapes exist with some very complex forms seen. Planetary nebulae are classified by
different authors into: stellar, disk, ring, irregular, helical, bipolar, quadrupolar,[46] and other types,[47]
although the majority of them belong to just three types: spherical, elliptical and bipolar. Bipolar
nebulae are concentrated in the galactic plane, probably produced by relatively young massive
progenitor stars; and bipolars in the galactic bulge appear to prefer orienting their orbital axes
parallel to the galactic plane.[48] On the other hand, spherical nebulae are probably produced by old
stars similar to the Sun.[41]
The huge variety of the shapes is partially the projection effect—the same nebula when viewed
under different angles will appear different.[49] Nevertheless, the reason for the huge variety of
physical shapes is not fully understood.[47] Gravitational interactions with companion stars if the
central stars are binary stars may be one cause. Another possibility is that planets disrupt the flow of
material away from the star as the nebula forms. It has been determined that the more massive stars
produce more irregularly shaped nebulae.[50] In January 2005, astronomers announced the first
detection of magnetic fields around the central stars of two planetary nebulae, and hypothesized that
the fields might be partly or wholly responsible for their remarkable shapes. [51][52]
Membership in clusters[edit]
Abell 78, 24 inch telescope on Mt. Lemmon, AZ. Courtesy of Joseph D. Schulman.
Planetary nebulae have been detected as members in four Galactic globular clusters: Messier 15,
Messier 22, NGC 6441 and Palomar 6. Evidence also points to the potential discovery of planetary
nebulae in globular clusters in the galaxy M31.[53] However, there is currently only one case of a
planetary nebula discovered in an open cluster that is agreed upon by independent
researchers.[54][55][56] That case pertains to the planetary nebula PHR 1315-6555 and the open
cluster Andrews-Lindsay 1. Indeed, through cluster membership, PHR 1315-6555 possesses among
the most precise distances established for a planetary nebula (i.e., a 4% distance solution). The
cases of NGC 2818 and NGC 2348 in Messier 46, exhibit mismatched velocities between the
planetary nebulae and the clusters, which indicates they are line-of-sight coincidences.[44][57][58] A
subsample of tentative cases that may potentially be cluster/PN pairs includes Abell 8 and Bica
6,[59][60] and He 2-86 and NGC 4463.[61]
Theoretical models predict that planetary nebulae can form from main-sequence stars of between
one and eight solar masses, which puts the progenitor star's age at greater than 40 million years.
Although there are a few hundred known open clusters within that age range, a variety of reasons
limit the chances of finding a planetary nebula within.[44] For one reason, the planetary nebula phase
for more massive stars is on the order of millennia, which is a blink of the eye in astronomic terms.
Also, partly because of their small total mass, open clusters have relatively poor gravitational
cohesion and tend to disperse after a relatively short time, typically from 100 to 600 million years. [62]
Current issues in planetary nebula studies[edit]
The distances to planetary nebulae are generally poorly determined,[63] but the Gaia mission is now
measuring direct parallactic distances between their central stars and neighboring stars.[64] It is also
possible to determine distances to nearby planetary nebula by measuring their expansion rates.
High resolution observations taken several years apart will show the expansion of the nebula
perpendicular to the line of sight, while spectroscopic observations of the Doppler shift will reveal the
velocity of expansion in the line of sight. Comparing the angular expansion with the derived velocity
of expansion will reveal the distance to the nebula.[22]
The issue of how such a diverse range of nebular shapes can be produced is a debatable topic. It is
theorised that interactions between material moving away from the star at different speeds gives rise
to most observed shapes.[47] However, some astronomers postulate that close binary central stars
might be responsible for the more complex and extreme planetary nebulae. [65] Several have been
shown to exhibit strong magnetic fields,[66] and their interactions with ionized gas could explain
some planetary nebulae shapes.[52]
There are two main methods of determining metal abundances in nebulae. These rely on
recombination lines and collisionally excited lines. Large discrepancies are sometimes seen between
the results derived from the two methods. This may be explained by the presence of small
temperature fluctuations within planetary nebulae. The discrepancies may be too large to be caused
by temperature effects, and some hypothesize the existence of cold knots containing very little
hydrogen to explain the observations. However, such knots have yet to be observed.[67]
Neutron star
From Wikipedia, the free encyclopedia
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For the story by Larry Niven, see Neutron Star (short story).
Simulated view of a Neutron star with accretion disk. The disk appears distorted near the star due to
extreme gravitational lensing
Radiation from the rapidly spinning pulsar PSR B1509-58 makes nearby gas emit X-rays (gold) and
illuminates the rest of the nebula, here seen in infrared (blue and red).
A neutron star is the collapsed core of a massive supergiant star, which had a total mass of
between 10 and 25 solar masses, possibly more if the star was especially metal-rich.[1] Except
for black holes, and some hypothetical objects (e.g. white holes, quark stars, and strange stars),
neutron stars are the smallest and densest currently known class of stellar objects. [2] Neutron
stars have a radius on the order of 10 kilometres (6 mi) and a mass of about 1.4 solar
masses.[3] They result from the supernova explosion of a massive star, combined with
gravitational collapse, that compresses the core past white dwarf star density to that of atomic
nuclei.
Once formed, they no longer actively generate heat, and cool over time; however, they may still
evolve further through collision or accretion. Most of the basic models for these objects imply
that neutron stars are composed almost entirely of neutrons (subatomic particles with no net
electrical charge and with slightly larger mass than protons); the electrons and protons present
in normal matter combine to produce neutrons at the conditions in a neutron star. Neutron stars
are partially supported against further collapse by neutron degeneracy pressure, a phenomenon
described by the Pauli exclusion principle, just as white dwarfs are supported against collapse
by electron degeneracy pressure. However, neutron degeneracy pressure is not by itself
sufficient to hold up an object beyond 0.7M☉[4][5] and repulsive nuclear forces play a larger role
in supporting more massive neutron stars. [6][7] If the remnant star has a mass exceeding the
Tolman–Oppenheimer–Volkoff limit of around 2 solar masses, the combination of degeneracy
pressure and nuclear forces is insufficient to support the neutron star and it continues collapsing
to form a black hole. The most massive neutron star detected so far, PSR J0740+6620, is
estimated to be 2.08 +/- 0.07 solar masses.
Neutron stars that can be observed are very hot and typically have a surface temperature of
around 600000 K.[8][9][10][11][a] They are so dense that a normal-sized matchbox containing
neutron-star material would have a weight of approximately 3 billion tonnes, the same weight as
a 0.5 cubic kilometre chunk of the Earth (a cube with edges of about 800 metres) from Earth's
surface.[12][13] Their magnetic fields are between 108 and 1015 (100 million and 1 quadrillion)
times stronger than Earth's magnetic field. The gravitational field at the neutron star's surface is
about 2×1011 (200 billion) times that of Earth's gravitational field.
As the star's core collapses, its rotation rate increases as a result of conservation of angular
momentum, and newly formed neutron stars hence rotate at up to several hundred times per
second. Some neutron stars emit beams of electromagnetic radiation that make them
detectable as pulsars. Indeed, the discovery of pulsars by Jocelyn Bell Burnell and Antony
Hewish in 1967 was the first observational suggestion that neutron stars exist. The radiation
from pulsars is thought to be primarily emitted from regions near their magnetic poles. If the
magnetic poles do not coincide with the rotational axis of the neutron star, the emission beam
will sweep the sky, and when seen from a distance, if the observer is somewhere in the path of
the beam, it will appear as pulses of radiation coming from a fixed point in space (the so-called
"lighthouse effect"). The fastest-spinning neutron star known is PSR J1748-2446ad, rotating at a
rate of 716 times a second[14][15] or 43,000 revolutions per minute, giving a linear speed at the
surface on the order of 0.24 c (i.e., nearly a quarter the speed of light).
There are thought to be around one billion neutron stars in the Milky Way,[16] and at a minimum
several hundred million, a figure obtained by estimating the number of stars that have
undergone supernova explosions.[17] However, most are old and cold and radiate very little;
most neutron stars that have been detected occur only in certain situations in which they do
radiate, such as if they are a pulsar or part of a binary system. Slow-rotating and non-accreting
neutron stars are almost undetectable; however, since the Hubble Space Telescope detection
of RX J185635−3754 in the 1990s, a few nearby neutron stars that appear to emit only thermal
radiation have been detected. Soft gamma repeaters are conjectured to be a type of neutron
star with very strong magnetic fields, known as magnetars, or alternatively, neutron stars with
fossil disks around them.[18]
Neutron stars in binary systems can undergo accretion which typically makes the system bright
in X-rays while the material falling onto the neutron star can form hotspots that rotate in and out
of view in identified X-ray pulsar systems. Additionally, such accretion can "recycle" old pulsars
and potentially cause them to gain mass and spin-up to very fast rotation rates, forming the socalled millisecond pulsars. These binary systems will continue to evolve, and eventually the
companions can become compact objects such as white dwarfs or neutron stars themselves,
though other possibilities include a complete destruction of the companion through ablation or
merger. The merger of binary neutron stars may be the source of short-duration gamma-ray
bursts and are likely strong sources of gravitational waves. In 2017, a direct detection
(GW170817) of the gravitational waves from such an event was observed,[19] and gravitational
waves have also been indirectly observed in a system where two neutron stars orbit each other.
Formation[edit]
Simplified representation of the formation of neutron stars.
Any main-sequence star with an initial mass of above 8 times the mass of the sun (8 M☉) has the
potential to produce a neutron star. As the star evolves away from the main sequence, subsequent
nuclear burning produces an iron-rich core. When all nuclear fuel in the core has been exhausted,
the core must be supported by degeneracy pressure alone. Further deposits of mass from shell
burning cause the core to exceed the Chandrasekhar limit. Electron-degeneracy pressure is
overcome and the core collapses further, sending temperatures soaring to over 5×109 K. At these
temperatures, photodisintegration (the breaking up of iron nuclei into alpha particles by high-energy
gamma rays) occurs. As the temperature climbs even higher, electrons and protons combine to form
neutrons via electron capture, releasing a flood of neutrinos. When densities reach nuclear density
of 4×1017 kg/m 3, a combination of strong force repulsion and neutron degeneracy pressure halts the
contraction.[20] The infalling outer envelope of the star is halted and flung outwards by a flux of
neutrinos produced in the creation of the neutrons, becoming a supernova. The remnant left is a
neutron star. If the remnant has a mass greater than about 3 M ☉, it collapses further to become a
black hole.[21]
As the core of a massive star is compressed during a Type II supernova or a Type Ib or Type Ic
supernova, and collapses into a neutron star, it retains most of its angular momentum. But, because
it has only a tiny fraction of its parent's radius (and therefore its moment of inertia is sharply
reduced), a neutron star is formed with very high rotation speed, and then over a very long period it
slows. Neutron stars are known that have rotation periods from about 1.4 ms to 30 s. The neutron
star's density also gives it very high surface gravity, with typical values ranging from 1012 to 1013
m/s2 (more than 1011 times that of Earth).[11] One measure of such immense gravity is the fact that
neutron stars have an escape velocity ranging from 100,000 km/s to 150,000 km/s, that is, from a
third to half the speed of light. The neutron star's gravity accelerates infalling matter to tremendous
speed. The force of its impact would likely destroy the object's component atoms, rendering all the
matter identical, in most respects, to the rest of the neutron star.[citation needed]
Properties[edit]
Mass and temperature[edit]
A neutron star has a mass of at least 1.1 solar masses (M☉). The upper limit of mass for a neutron
star is called the Tolman–Oppenheimer–Volkoff limit and is generally held to be around 2.1
M☉,[22][23] but a recent estimate puts the upper limit at 2.16 M☉.[24] The maximum observed mass of
neutron stars is about 2.14 M☉ for PSR J0740+6620 discovered in September, 2019.[25] Compact
stars below the Chandrasekhar limit of 1.39 M☉ are generally white dwarfs whereas compact stars
with a mass between 1.4 M☉ and 2.16 M☉ are expected to be neutron stars, but there is an interval
of a few tenths of a solar mass where the masses of low-mass neutron stars and high-mass white
dwarfs can overlap. It is thought that beyond 2.16 M ☉ the stellar remnant will overcome the strong
force repulsion and neutron degeneracy pressure so that gravitational collapse will occur to produce
a black hole, but the smallest observed mass of a stellar black hole is about 5 M☉.[b] Between 2.16
M☉ and 5 M☉, hypothetical intermediate-mass stars such as quark stars and electroweak stars have
been proposed, but none have been shown to exist.[b]
The temperature inside a newly formed neutron star is from around 1011 to 1012 kelvins.[27]
However, the huge number of neutrinos it emits carry away so much energy that the temperature of
an isolated neutron star falls within a few years to around 106 kelvins.[27] At this lower temperature,
most of the light generated by a neutron star is in X-rays.
Some researchers have proposed a neutron star classification system using Roman numerals (not
to be confused with the Yerkes luminosity classes for non-degenerate stars) to sort neutron stars by
their mass and cooling rates: type I for neutron stars with low mass and cooling rates, type II for
neutron stars with higher mass and cooling rates, and a proposed type III for neutron stars with even
higher mass, approaching 2 M☉, and with higher cooling rates and possibly candidates for exotic
stars.[28]
Density and pressure[edit]
Neutron stars have overall densities of 3.7×1017 to 5.9×1017 kg/m 3 (2.6×1014 to 4.1×1014 times the
density of the Sun),[c] which is comparable to the approximate density of an atomic nucleus of
3×1017 kg/m 3.[29] The neutron star's density varies from about 1×109 kg/m 3 in the crust—increasing
with depth—to about 6×1017 or 8×1017 kg/m 3 (denser than an atomic nucleus) deeper inside.[27] A
neutron star is so dense that one teaspoon (5 milliliters) of its material would have a mass over
5.5×1012 kg, about 900 times the mass of the Great Pyramid of Giza. In the enormous gravitational
field of a neutron star, that teaspoon of material would weigh 1.1×1025 N, which is 15 times what the
Moon would weigh if it were placed on the surface of the Earth.[d] The entire mass of the Earth at
neutron star density would fit into a sphere of 305 m in diameter (the size of the Arecibo Telescope).
The pressure increases from 3.2×1031 to 1.6×1034 Pa from the inner crust to the center.[30]
The equation of state of matter at such high densities is not precisely known because of the
theoretical difficulties associated with extrapolating the likely behavior of quantum chromodynamics,
superconductivity, and superfluidity of matter in such states. The problem is exacerbated by the
empirical difficulties of observing the characteristics of any object that is hundreds of parsecs away,
or farther.
A neutron star has some of the properties of an atomic nucleus, including density (within an order of
magnitude) and being composed of nucleons. In popular scientific writing, neutron stars are
therefore sometimes described as "giant nuclei". However, in other respects, neutron stars and
atomic nuclei are quite different. A nucleus is held together by the strong interaction, whereas a
neutron star is held together by gravity. The density of a nucleus is uniform, while neutron stars are
predicted to consist of multiple layers with varying compositions and densities.
Magnetic field[edit]
The magnetic field strength on the surface of neutron stars ranges from c. 10 4 to 1011 tesla.[31]
These are orders of magnitude higher than in any other object: For comparison, a continuous 16 T
field has been achieved in the laboratory and is sufficient to levitate a living frog due to diamagnetic
levitation. Variations in magnetic field strengths are most likely the main factor that allows different
types of neutron stars to be distinguished by their spectra, and explains the periodicity of pulsars. [31]
The neutron stars known as magnetars have the strongest magnetic fields, in the range of 108 to
1011 tesla,[32] and have become the widely accepted hypothesis for neutron star types soft gamma
repeaters (SGRs)[33] and anomalous X-ray pulsars (AXPs).[34] The magnetic energy density of a
108 T field is extreme, greatly exceeding the mass-energy density of ordinary matter.[e] Fields of this
strength are able to polarize the vacuum to the point that the vacuum becomes birefringent. Photons
can merge or split in two, and virtual particle-antiparticle pairs are produced. The field changes
electron energy levels and atoms are forced into thin cylinders. Unlike in an ordinary pulsar,
magnetar spin-down can be directly powered by its magnetic field, and the magnetic field is strong
enough to stress the crust to the point of fracture. Fractures of the crust cause starquakes, observed
as extremely luminous millisecond hard gamma ray bursts. The fireball is trapped by the magnetic
field, and comes in and out of view when the star rotates, which is observed as a periodic soft
gamma repeater (SGR) emission with a period of 5–8 seconds and which lasts for a few minutes.[36]
The origins of the strong magnetic field are as yet unclear. [31] One hypothesis is that of "flux
freezing", or conservation of the original magnetic flux during the formation of the neutron star.[31] If
an object has a certain magnetic flux over its surface area, and that area shrinks to a smaller area,
but the magnetic flux is conserved, then the magnetic field would correspondingly increase.
Likewise, a collapsing star begins with a much larger surface area than the resulting neutron star,
and conservation of magnetic flux would result in a far stronger magnetic field. However, this simple
explanation does not fully explain magnetic field strengths of neutron stars. [31]
Gravity and equation of state[edit]
Gravitational light deflection at a neutron star. Due to relativistic light deflection over half the surface is
visible (each grid patch represents 30 by 30 degrees).[37] In natural units, this star's mass is 1 and its radius
is 4, or twice its Schwarzschild radius.[37]
The gravitational field at a neutron star's surface is about 2×1011 times stronger than on Earth, at
around 2.0×1012 m/s2.[38] Such a strong gravitational field acts as a gravitational lens and bends the
radiation emitted by the neutron star such that parts of the normally invisible rear surface become
visible.[37] If the radius of the neutron star is 3GM/c2 or less, then the photons may be trapped in an
orbit, thus making the whole surface of that neutron star visible from a single vantage point, along
with destabilizing photon orbits at or below the 1 radius distance of the star.
A fraction of the mass of a star that collapses to form a neutron star is released in the supernova
explosion from which it forms (from the law of mass–energy equivalence, E = mc2). The energy
comes from the gravitational binding energy of a neutron star.
Hence, the gravitational force of a typical neutron star is huge. If an object were to fall from a height
of one meter on a neutron star 12 kilometers in radius, it would reach the ground at around 1400
kilometers per second.[39] However, even before impact, the tidal force would cause
spaghettification, breaking any sort of an ordinary object into a stream of material.
Because of the enormous gravity, time dilation between a neutron star and Earth is significant. For
example, eight years could pass on the surface of a neutron star, yet ten years would have passed
on Earth, not including the time-dilation effect of the star's very rapid rotation.[40]
Neutron star relativistic equations of state describe the relation of radius vs. mass for various
models.[41] The most likely radii for a given neutron star mass are bracketed by models AP4
(smallest radius) and MS2 (largest radius). EB is the ratio of gravitational binding energy mass
equivalent to the observed neutron star gravitational mass of M kilograms with radius R meters,[42]
A 2 M☉ neutron star would not be more compact than 10,970 meters radius (AP4 model). Its
mass fraction gravitational binding energy would then be 0.187, − 18.7% (exothermic). This is not
near 0.6/2 = 0.3, − 30%.
The equation of state for a neutron star is not yet known. It is assumed that it differs significantly
from that of a white dwarf, whose equation of state is that of a degenerate gas that can be
described in close agreement with special relativity. However, with a neutron star the increased
effects of general relativity can no longer be ignored. Several equations of state have been
proposed (FPS, UU, APR, L, SLy, and others) and current research is still attempting to
constrain the theories to make predictions of neutron star matter.[11][44] This means that the
relation between density and mass is not fully known, and this causes uncertainties in radius
estimates. For example, a 1.5 M☉ neutron star could have a radius of 10.7, 11.1, 12.1 or 15.1
kilometers (for EOS FPS, UU, APR or L respectively).[44]
Structure
Cross-section of neutron star. Densities are in terms of ρ0 the saturation nuclear matter density,
where nucleons begin to touch.
Current understanding of the structure of neutron stars is defined by existing mathematical
models, but it might be possible to infer some details through studies of neutron-star
oscillations. Asteroseismology, a study applied to ordinary stars, can reveal the inner structure
of neutron stars by analyzing observed spectra of stellar oscillations.[11]
Current models indicate that matter at the surface of a neutron star is composed of ordinary
atomic nuclei crushed into a solid lattice with a sea of electrons flowing through the gaps
between them. It is possible that the nuclei at the surface are iron, due to iron's high binding
energy per nucleon.[45] It is also possible that heavy elements, such as iron, simply sink
beneath the surface, leaving only light nuclei like helium and hydrogen.[45] If the surface
temperature exceeds 106 kelvins (as in the case of a young pulsar), the surface should be fluid
instead of the solid phase that might exist in cooler neutron stars (temperature <106
kelvins).[45]
The "atmosphere" of a neutron star is hypothesized to be at most several micrometres thick,
and its dynamics are fully controlled by the neutron star's magnetic field. Below the atmosphere
one encounters a solid "crust". This crust is extremely hard and very smooth (with maximum
surface irregularities on the order of millimetres or less), due to the extreme gravitational
field.[46][47]
Proceeding inward, one encounters nuclei with ever-increasing numbers of neutrons; such
nuclei would decay quickly on Earth, but are kept stable by tremendous pressures. As this
process continues at increasing depths, the neutron drip becomes overwhelming, and the
concentration of free neutrons increases rapidly. In that region, there are nuclei, free electrons,
and free neutrons. The nuclei become increasingly small (gravity and pressure overwhelming
the strong force) until the core is reached, by definition the point where mostly neutrons exist.
The expected hierarchy of phases of nuclear matter in the inner crust has been characterized as
"nuclear pasta", with fewer voids and larger structures towards higher pressures.[48] The
composition of the superdense matter in the core remains uncertain. One model describes the
core as superfluid neutron-degenerate matter (mostly neutrons, with some protons and
electrons). More exotic forms of matter are possible, including degenerate strange matter
(containing strange quarks in addition to up and down quarks), matter containing high-energy
pions and kaons in addition to neutrons,[11] or ultra-dense quark-degenerate matter.
Radiation
Computer renders of a neutron star with accretion disk, with magnetic field lines projected,
showing bursts of powerful X-rays and Radio Waves. The simulations are taken from 2017 data
from NASA's NuSTAR and Swift, and ESA's XMM-Newto observatories
File:Pulsar anim.ogv
Animation of a rotating pulsar. The sphere in the middle represents the neutron star, the curves
indicate the magnetic field lines and the protruding cones represent the emission zones.
Pulsars
Main article: Pulsar
Neutron stars are detected from their electromagnetic radiation. Neutron stars are usually
observed to pulse radio waves and other electromagnetic radiation, and neutron stars observed
with pulses are called pulsars.
Pulsars' radiation is thought to be caused by particle acceleration near their magnetic poles,
which need not be aligned with the rotational axis of the neutron star. It is thought that a large
electrostatic field builds up near the magnetic poles, leading to electron emission.[49] These
electrons are magnetically accelerated along the field lines, leading to curvature radiation, with
the radiation being strongly polarized towards the plane of curvature.[49] In addition, high
energy photons can interact with lower energy photons and the magnetic field for
electron− positron pair production, which through electron–positron annihilation leads to further
high energy photons.[49]
The radiation emanating from the magnetic poles of neutron stars can be described as
magnetospheric radiation, in reference to the magnetosphere of the neutron star.[50] It is not to
be confused with magnetic dipole radiation, which is emitted because the magnetic axis is not
aligned with the rotational axis, with a radiation frequency the same as the neutron star's
rotational frequency.[49]
If the axis of rotation of the neutron star is different to the magnetic axis, external viewers will
only see these beams of radiation whenever the magnetic axis point towards them during the
neutron star rotation. Therefore, periodic pulses are observed, at the same rate as the rotation
of the neutron star.
Non-pulsating neutron stars
In addition to pulsars, non-pulsating neutron stars have also been identified, although they may
have minor periodic variation in luminosity.[51][52] This seems to be a characteristic of the X-ray
sources known as Central Compact Objects in Supernova remnants (CCOs in SNRs), which are
thought to be young, radio-quiet isolated neutron stars.[51]
Spectra
In addition to radio emissions, neutron stars have also been identified in other parts of the
electromagnetic spectrum. This includes visible light, near infrared, ultraviolet, X-rays, and
gamma rays.[50] Pulsars observed in X-rays are known as X-ray pulsars if accretion-powered,
while those identified in visible light are known as optical pulsars. The majority of neutron stars
detected, including those identified in optical, X-ray, and gamma rays, also emit radio
waves;[53] the Crab Pulsar produces electromagnetic emissions across the spectrum.[53]
However, there exist neutron stars called radio-quiet neutron stars, with no radio emissions
detected.[54]
Rotation
Neutron stars rotate extremely rapidly after their formation due to the conservation of angular
momentum; in analogy to spinning ice skaters pulling in their arms, the slow rotation of the
original star's core speeds up as it shrinks. A newborn neutron star can rotate many times a
second.
Spin down
P–P-dot diagram for known rotation-powered pulsars (red), anomalous X-ray pulsars (green),
high-energy emission pulsars (blue) and binary pulsars (pink)
Over time, neutron stars slow, as their rotating magnetic fields in effect radiate energy
associated with the rotation; older neutron stars may take several seconds for each revolution.
This is called spin down. The rate at which a neutron star slows its rotation is usually constant
and very small.
The periodic time (P) is the rotational period, the time for one rotation of a neutron star. The
spin-down rate, the rate of slowing of rotation, is then given the symbol {\displaystyle {\dot
{P}}}{\dot {P}} (P-dot), the derivative of P with respect to time. It is defined as periodic time
increase per unit time; it is a dimensionless quantity, but can be given the units of s⋅s− 1 (seconds
per second).[49]
The spin-down rate (P-dot) of neutron stars usually falls within the range of 10− 22 to 10− 9 s⋅s− 1,
with the shorter period (or faster rotating) observable neutron stars usually having smaller P-dot.
As a neutron star ages, its rotation slows (as P increases); eventually, the rate of rotation will
become too slow to power the radio-emission mechanism, and the neutron star can no longer
be detected.[49]
P and P-dot allow minimum magnetic fields of neutron stars to be estimated.[49] P and P-dot
can be also used to calculate the characteristic age of a pulsar, but gives an estimate which is
somewhat larger than the true age when it is applied to young pulsars.[49]
P and P-dot can also be combined with neutron star's moment of inertia to estimate a quantity
called spin-down luminosity, which is given the symbol {\displaystyle {\dot {E}}}{\displaystyle
{\dot {E}}} (E-dot). It is not the measured luminosity, but rather the calculated loss rate of
rotational energy that would manifest itself as radiation. For neutron stars where the spin-down
luminosity is comparable to the actual luminosity, the neutron stars are said to be "rotation
powered".[49][50] The observed luminosity of the Crab Pulsar is comparable to the spin-down
luminosity, supporting the model that rotational kinetic energy powers the radiation from it.[49]
With neutron stars such as magnetars, where the actual luminosity exceeds the spin-down
luminosity by about a factor of one hundred, it is assumed that the luminosity is powered by
magnetic dissipation, rather than being rotation powered.[55]
P and P-dot can also be plotted for neutron stars to create a P–P-dot diagram. It encodes a
tremendous amount of information about the pulsar population and its properties, and has been
likened to the Hertzsprung–Russell diagram in its importance for neutron stars.[49]
Spin up
Neutron star rotational speeds can increase, a process known as spin up. Sometimes neutron
stars absorb orbiting matter from companion stars, increasing the rotation rate and reshaping
the neutron star into an oblate spheroid. This causes an increase in the rate of rotation of the
neutron star of over a hundred times per second in the case of millisecond pulsars.
The most rapidly rotating neutron star currently known, PSR J1748-2446ad, rotates at 716
revolutions per second.[56] A 2007 paper reported the detection of an X-ray burst oscillation,
which provides an indirect measure of spin, of 1122 Hz from the neutron star XTE J1739285,[57] suggesting 1122 rotations a second. However, at present, this signal has only been
seen once, and should be regarded as tentative until confirmed in another burst from that star.
Glitches and starquakes
NASA artist's conception of a "starquake", or "stellar quake".
Sometimes a neutron star will undergo a glitch, a sudden small increase of its rotational speed
or spin up. Glitches are thought to be the effect of a starquake—as the rotation of the neutron
star slows, its shape becomes more spherical. Due to the stiffness of the "neutron" crust, this
happens as discrete events when the crust ruptures, creating a starquake similar to
earthquakes. After the starquake, the star will have a smaller equatorial radius, and because
angular momentum is conserved, its rotational speed has increased.
Starquakes occurring in magnetars, with a resulting glitch, is the leading hypothesis for the
gamma-ray sources known as soft gamma repeaters.[58]
Recent work, however, suggests that a starquake would not release sufficient energy for a
neutron star glitch; it has been suggested that glitches may instead be caused by transitions of
vortices in the theoretical superfluid core of the neutron star from one metastable energy state
to a lower one, thereby releasing energy that appears as an increase in the rotation rate.[59]
"Anti-glitches"
An "anti-glitch", a sudden small decrease in rotational speed, or spin down, of a neutron star
has also been reported.[60] It occurred in the magnetar 1E 2259+586, that in one case
produced an X-ray luminosity increase of a factor of 20, and a significant spin-down rate
change. Current neutron star models do not predict this behavior. If the cause was internal, it
suggests differential rotation of solid outer crust and the superfluid component of the magnetar's
inner structure.[60]
Population and distances
Central neutron star at the heart of the Crab Nebula.[61]
At present, there are about 2,000 known neutron stars in the Milky Way and the Magellanic
Clouds, the majority of which have been detected as radio pulsars. Neutron stars are mostly
concentrated along the disk of the Milky Way, although the spread perpendicular to the disk is
large because the supernova explosion process can impart high translational speeds (400 km/s)
to the newly formed neutron star.
Some of the closest known neutron stars are RX J1856.5− 3754, which is about 400 light-years
from Earth, and PSR J0108− 1431 about 424 light years.[62] RX J1856.5-3754 is a member of a
close group of neutron stars called The Magnificent Seven. Another nearby neutron star that was
detected transiting the backdrop of the constellation Ursa Minor has been nicknamed Calvera by
its Canadian and American discoverers, after the villain in the 1960 film The Magnificent Seven.
This rapidly moving object was discovered using the ROSAT/Bright Source Catalog.
Neutron stars are only detectable with modern technology during the earliest stages of their
lives (almost always less than 1 million years) and are vastly outnumbered by older neutron
stars that would only be detectable through their blackbody radiation and gravitational effects on
other stars.
Binary neutron star systems
Circinus X-1: X-ray light rings from a binary neutron star (24 June 2015; Chandra X-ray
Observatory)
About 5% of all known neutron stars are members of a binary system. The formation and
evolution of binary neutron stars[63] and double neutron stars[64] can be a complex process.
Neutron stars have been observed in binaries with ordinary main-sequence stars, red giants,
white dwarfs, or other neutron stars. According to modern theories of binary evolution, it is
expected that neutron stars also exist in binary systems with black hole companions. The
merger of binaries containing two neutron stars, or a neutron star and a black hole, has been
observed through the emission of gravitational waves.[65][66]
X-ray binaries
Main article: X-ray binary
Binary systems containing neutron stars often emit X-rays, which are emitted by hot gas as it
falls towards the surface of the neutron star. The source of the gas is the companion star, the
outer layers of which can be stripped off by the gravitational force of the neutron star if the two
stars are sufficiently close. As the neutron star accretes this gas, its mass can increase; if
enough mass is accreted, the neutron star may collapse into a black hole.[67]
Neutron star binary mergers and nucleosynthesis
Main article: Stellar collision
The distance between two neutron stars in a close binary system is observed to shrink as
gravitational waves are emitted.[68] Ultimately, the neutron stars will come into contact and
coalesce. The coalescence of binary neutron stars is one of the leading models for the origin of
short gamma-ray bursts. Strong evidence for this model came from the observation of a
kilonova associated with the short-duration gamma-ray burst GRB 130603B,[69] and finally
confirmed by detection of gravitational wave GW170817 and short GRB 170817A by LIGO,
Virgo, and 70 observatories covering the electromagnetic spectrum observing the
event.[70][71][72][73] The light emitted in the kilonova is believed to come from the radioactive
decay of material ejected in the merger of the two neutron stars. This material may be
responsible for the production of many of the chemical elements beyond iron,[74] as opposed to
the supernova nucleosynthesis theory.
Planets
An artist's conception of the pulsar planet PSR B1257+12 C, with bright aurorae.
Main article: Pulsar planet
Neutron stars can host exoplanets. These can be original, circumbinary, captured, or the result of
a second round of planet formation. Pulsars can also strip the atmosphere off from a star, leaving
a planetary-mass remnant, which may be understood as a chthonian planet or a stellar object
depending on interpretation. For pulsars, such pulsar planets can be detected with the pulsar
timing method, which allows for high precision and detection of much smaller planets than with
other methods. Two systems have been definitively confirmed. The first exoplanets ever to be
detected were the three planets Draugr, Poltergeist and Phobetor around PSR B1257+12,
discovered in 1992–1994. Of these, Draugr is the smallest exoplanet ever detected, at a mass of
twice that of the Moon. Another system is PSR B1620− 26, where a circumbinary planet orbits a
neutron star-white dwarf binary system. Also, there are several unconfirmed candidates. Pulsar
planets receive little visible light, but massive amounts of ionizing radiation and high-energy
stellar wind, which makes them rather hostile environments.
History of discoveries
The first direct observation of a neutron star in visible light. The neutron star is RX J1856.5− 3754.
At the meeting of the American Physical Society in December 1933 (the proceedings were
published in January 1934), Walter Baade and Fritz Zwicky proposed the existence of neutron
stars,[75][f] less than two years after the discovery of the neutron by James Chadwick.[78] In
seeking an explanation for the origin of a supernova, they tentatively proposed that in
supernova explosions ordinary stars are turned into stars that consist of extremely closely
packed neutrons that they called neutron stars. Baade and Zwicky correctly proposed at that
time that the release of the gravitational binding energy of the neutron stars powers the
supernova: "In the supernova process, mass in bulk is annihilated". Neutron stars were thought
to be too faint to be detectable and little work was done on them until November 1967, when
Franco Pacini pointed out that if the neutron stars were spinning and had large magnetic fields,
then electromagnetic waves would be emitted. Unbeknown to him, radio astronomer Antony
Hewish and his research assistant Jocelyn Bell at Cambridge were shortly to detect radio pulses
from stars that are now believed to be highly magnetized, rapidly spinning neutron stars, known
as pulsars.
In 1965, Antony Hewish and Samuel Okoye discovered "an unusual source of high radio
brightness temperature in the Crab Nebula".[79] This source turned out to be the Crab Pulsar
that resulted from the great supernova of 1054.
In 1967, Iosif Shklovsky examined the X-ray and optical observations of Scorpius X-1 and
correctly concluded that the radiation comes from a neutron star at the stage of accretion.[80]
In 1967, Jocelyn Bell Burnell and Antony Hewish discovered regular radio pulses from PSR
B1919+21. This pulsar was later interpreted as an isolated, rotating neutron star. The energy
source of the pulsar is the rotational energy of the neutron star. The majority of known neutron
stars (about 2000, as of 2010) have been discovered as pulsars, emitting regular radio pulses.
In 1968, Richard V. E. Lovelace and collaborators discovered period {\displaystyle P\!\approx
33}{\displaystyle P\!\approx 33} ms of the Crab pulsar using Arecibo Observatory.[81][82] After
this discovery, scientists concluded that pulsars were rotating neutron stars.[83] Before that,
many scientists believed that pulsars were pulsating white dwarfs.
In 1971, Riccardo Giacconi, Herbert Gursky, Ed Kellogg, R. Levinson, E. Schreier, and H.
Tananbaum discovered 4.8 second pulsations in an X-ray source in the constellation Centaurus,
Cen X-3.[84] They interpreted this as resulting from a rotating hot neutron star. The energy
source is gravitational and results from a rain of gas falling onto the surface of the neutron star
from a companion star or the interstellar medium.
In 1974, Antony Hewish was awarded the Nobel Prize in Physics "for his decisive role in the
discovery of pulsars" without Jocelyn Bell who shared in the discovery.[85]
In 1974, Joseph Taylor and Russell Hulse discovered the first binary pulsar, PSR B1913+16,
which consists of two neutron stars (one seen as a pulsar) orbiting around their center of mass.
Albert Einstein's general theory of relativity predicts that massive objects in short binary orbits
should emit gravitational waves, and thus that their orbit should decay with time. This was
indeed observed, precisely as general relativity predicts, and in 1993, Taylor and Hulse were
awarded the Nobel Prize in Physics for this discovery.[86]
In 1982, Don Backer and colleagues discovered the first millisecond pulsar, PSR B1937+21.[87]
This object spins 642 times per second, a value that placed fundamental constraints on the
mass and radius of neutron stars. Many millisecond pulsars were later discovered, but PSR
B1937+21 remained the fastest-spinning known pulsar for 24 years, until PSR J1748-2446ad
(which spins more than 700 times a second) was discovered.
In 2003, Marta Burgay and colleagues discovered the first double neutron star system where
both components are detectable as pulsars, PSR J0737− 3039.[88] The discovery of this system
allows a total of 5 different tests of general relativity, some of these with unprecedented
precision.
In 2010, Paul Demorest and colleagues measured the mass of the millisecond pulsar PSR
J1614− 2230 to be 1.97±0.04 M☉, using Shapiro delay.[89] This was substantially higher than any
previously measured neutron star mass (1.67 M☉, see PSR J1903+0327), and places strong
constraints on the interior composition of neutron stars.
In 2013, John Antoniadis and colleagues measured the mass of PSR J0348+0432 to be
2.01±0.04 M☉, using white dwarf spectroscopy.[90] This confirmed the existence of such
massive stars using a different method. Furthermore, this allowed, for the first time, a test of
general relativity using such a massive neutron star.
In August 2017, LIGO and Virgo made first detection of gravitational waves produced by
colliding neutron stars.[91]
In October 2018, astronomers reported that GRB 150101B, a gamma-ray burst event detected
in 2015, may be directly related to the historic GW170817 and associated with the merger of
two neutron stars. The similarities between the two events, in terms of gamma ray, optical and
x-ray emissions, as well as to the nature of the associated host galaxies, are "striking",
suggesting the two separate events may both be the result of the merger of neutron stars, and
both may be a kilonova, which may be more common in the universe than previously
understood, according to the researchers.[92][93][94][95]
In July 2019, astronomers reported that a new method to determine the Hubble constant, and
resolve the discrepancy of earlier methods, has been proposed based on the mergers of pairs
of neutron stars, following the detection of the neutron star merger of GW170817.[96][97] Their
measurement of the Hubble constant is 70.3+5.3
− 5.0 (km/s)/Mpc.[98]
A 2020 study by University of Southampton PhD student Fabian Gittins suggested that surface
irregularities ("mountains") may only be fractions of a millimeter tall (about 0.000003% of the
neutron star's diameter), hundreds of times smaller than previously predicted, a result bearing
implications for the non-detection of gravitational waves from spinning neutron
stars.[47][99][100]
Subtypes table
Different Types of Neutron Stars (24 June 2020)
Neutron star
Isolated neutron star (INS):[50][51][101][102] not in a binary system.
Rotation-powered pulsar (RPP or "radio pulsar"):[51] neutron stars that emit directed pulses of
radiation towards us at regular intervals (due to their strong magnetic fields).
Rotating radio transient (RRATs):[51] are thought to be pulsars which emit more sporadically
and/or with higher pulse-to-pulse variability than the bulk of the known pulsars.
Magnetar: a neutron star with an extremely strong magnetic field (1000 times more than a
regular neutron star), and long rotation periods (5 to 12 seconds).
Soft gamma repeater (SGR).[50]
Anomalous X-ray pulsar (AXP).[50]
Radio-quiet neutron stars.
X-ray dim isolated neutron stars.[51]
Central compact objects in supernova remnants (CCOs in SNRs): young, radio-quiet nonpulsating X-ray sources, thought to be Isolated Neutron Stars surrounded by supernova
remnants.[51]
X-ray pulsars or "accretion-powered pulsars": a class of X-ray binaries.
Low-mass X-ray binary pulsars: a class of low-mass X-ray binaries (LMXB), a pulsar with a
main sequence star, white dwarf or red giant.
Millisecond pulsar (MSP) ("recycled pulsar").
"Spider Pulsar", a pulsar where their companion is a semi-degenerate star.[103]
"Black Widow" pulsar, a pulsar that falls under the "Spider Pulsar" if the companion has
extremely low mass (less than 0.1 solar masses).
"Redback" pulsar, are if the companion is more massive.
Sub-millisecond pulsar.[104]
X-ray burster: a neutron star with a low mass binary companion from which matter is accreted
resulting in irregular bursts of energy from the surface of the neutron star.
Intermediate-mass X-ray binary pulsars: a class of intermediate-mass X-ray binaries (IMXB), a
pulsar with an intermediate mass star.
High-mass X-ray binary pulsars: a class of high-mass X-ray binaries (HMXB), a pulsar with a
massive star.
Binary pulsars: a pulsar with a binary companion, often a white dwarf or neutron star.
X-ray tertiary (theorized).[105]
Theorized compact stars with similar properties.
Protoneutron star (PNS), theorized.[106]
Exotic star
Thorne–Żytkow object: currently a hypothetical merger of a neutron star into a red giant star.
Quark star: currently a hypothetical type of neutron star composed of quark matter, or strange
matter. As of 2018, there are three candidates.
Electroweak star: currently a hypothetical type of extremely heavy neutron star, in which the
quarks are converted to leptons through the electroweak force, but the gravitational collapse of
the neutron star is prevented by radiation pressure. As of 2018, there is no evidence for their
existence.
Preon star: currently a hypothetical type of neutron star composed of preon matter. As of 2018,
there is no evidence for the existence of preons.
Dwarf and Recurrent
Novae/Novas
Dwarf novae undergo multiple eruptions in which their brightness increases by about 2 – 5
magnitudes. Each eruption lasts between 2 and 20 days with the nova duration related to the
interval between outbursts. This time interval is quasi-periodic and ranges from days to
decades.
While classical and recurrent novae result from runaway thermonuclear reactions on the
surface of a white dwarf, dwarf novae originate via a different mechanism. Two possible
scenarios have been put forward to account for the outbursts, with the Disk Instability
Model, currently popular amongst astronomers, able to reproduce many of the
characteristics of dwarf nova eruptions.
In this model, the accretion disk surrounding the white dwarf is able to accumulate a certain
amount of gas at a steady rate before becoming unstable. These instabilities cause material to
be dumped more rapidly onto the white dwarf resulting in the release of large amounts of
gravitational energy — the nova outburst. Once the accretion disk has lost enough mass, it
once again becomes stable, the increased transfer of material to the white dwarf returns to
normal levels, and the system returns to its quiescent state. The cycle is repeated once the
disk again accumulates sufficient mass for instabilities within the accretion disk to arise.
The alternate model is the Mass Transfer Burst Model, where a sudden increase in the
transfer of material from the companion star to the accretion disk causes the disk to collapse.
The result is that matter is suddenly dumped onto the white dwarf releasing large amounts of
gravitational energy — the nova outburst.
Due to the different origin for the outbursts, one of the features which distinguishes dwarf
novae from classical and recurrent novae is that they do not eject a shell of material with the
outburst. In addition, the increase in their brightness is much less than for other novae, hence
the term dwarf in the name. They are the most abundant type of nova, with several hundred
dwarf novae known.
Recurrent novae are thought to arise in the same way as classical novae, through a
white dwarf in a close binary system accreting a surface layer of hydrogen from a
main sequence companion. Once the temperature at the bottom of this hydrogen layer
reaches about 10 million Kelvin, a runaway thermonuclear reaction takes place which
ejects the unburnt hydrogen into a rapidly-expanding shell around the white dwarf. This
is the nova outburst.
While classical novae have only been seen in outburst once, recurrent novae have
undergone at least two outbursts over the past century (since astronomers started
taking notice!). The time interval between outbursts varies from 10 to 100 years, and
astronomers propose that classical novae will be seen as recurrent novae given enough
time.
The 8 recurrent novae astronomers know about tend to be slightly brighter than
classical novae in their quiescent state. In outburst, they have the same brightness or
are slightly fainter than classical novae, with the brightest maxima occuring for those
novae with the shortest time interval between outbursts.
A U Geminorum-type variable star, or dwarf nova (pl. novae) is one of several types of
cataclysmic variable star, consisting of a close binary star system in which one of the components is
a white dwarf that accretes matter from its companion. Dwarf novae are dimmer and repeat more
frequently than "classical" novae.[1]
Overview[edit]
The first one to be observed was U Geminorum in 1855; however, the mechanism was not known till
1974, when Brian Warner showed that the nova is due to the increase of the luminosity of the
accretion disk.[2] They are similar to classical novae in that the white dwarf is involved in periodic
outbursts, but the mechanisms are different. Classical novae result from the fusion and detonation of
accreted hydrogen on the primary's surface. Current theory suggests that dwarf novae result from
instability in the accretion disk, when gas in the disk reaches a critical temperature that causes a
change in viscosity, resulting in a temporary increase in mass flow through the disc, which heats the
whole disc and hence increases its luminosity. The mass transfer from the donor star is less than
this increased flow through the disc, so the disc will eventually drop back below the critical
temperature and revert to a cooler, duller mode.[3][4]
Dwarf novae are distinct from classical novae in other ways; their luminosity is lower, and they are
typically recurrent on a scale from days to decades.[3] The luminosity of the outburst increases with
the recurrence interval as well as the orbital period; recent research with the Hubble Space
Telescope suggests that the latter relationship could make dwarf novae useful standard candles for
measuring cosmic distances.[3][4]
There are three subtypes of U Geminorum star (UG):[5]
●
●
●
SS Cygni stars (UGSS), which increase in brightness by 2-6 mag in V in 1–2 days, and
return to their original brightnesses in several subsequent days.
SU Ursae Majoris stars (UGSU), which have brighter and longer "supermaxima"
outbursts, or "super-outbursts," in addition to normal outbursts. Varieties of SU Ursae
Majoris star include ER Ursae Majoris stars and WZ Sagittae stars (UGWZ).[6]
Z Camelopardalis stars (UGZ), which temporarily "halt" at a particular brightness below
their peak.
In addition to the large outbursts, some dwarf novae show periodic brightening known as
“superhumps”. They are caused by deformations of the accretion disk when its rotation is in
resonance with the orbital period of the binary.
●
AAVSO light curve of U Geminorum (SS Cygni type)
●
Light curve of eclipsing dwarf nova HT Cassiopeia during outburst, showing eclipses and
superhumps (SU Ursae Majoris type)
●
Light curve of Z Camelopardalis (Z Camelopardalis type)
^DWARF AND RECURRENT SUPER NOVA IMAGE
Type Ia supernova
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At the core of a planetary nebula, Henize 2-428, two white dwarf stars slightly under one solar mass each
are expected to merge and create a Type Ia supernova destroying both in about 700 million years (artist's
impression).
A type 1a supernova (read: "type one-A") is a type of supernova that occurs in binary systems (two
stars orbiting one another) in which one of the stars is a white dwarf. The other star can be anything
from a giant star to an even smaller white dwarf.[1]
Physically, carbon–oxygen white dwarfs with a low rate of rotation are limited to below 1.44 solar
masses (M☉).[2][3] Beyond this "critical mass", they reignite and in some cases trigger a supernova
explosion; this critical mass is often referred to as the Chandrasekhar mass, but is marginally
different from the absolute Chandrasekhar limit, where electron degeneracy pressure is unable to
prevent catastrophic collapse. If a white dwarf gradually accretes mass from a binary companion, or
merges with a second white dwarf, the general hypothesis is that a white dwarf's core will reach the
ignition temperature for carbon fusion as it approaches the Chandrasekhar mass. Within a few
seconds of initiation of nuclear fusion, a substantial fraction of the matter in the white dwarf
undergoes a runaway reaction, releasing enough energy (1–2×1044 J)[4] to unbind the star in a
supernova explosion.[5]
The type Ia category of supernova produces a fairly consistent peak luminosity because of this fixed
critical mass at which a white dwarf will explode. Their consistent peak luminosity allows these
explosions to be used as standard candles to measure the distance to their host galaxies: the visual
magnitude of a type Ia supernova, as observed from Earth, indicates its distance from Earth.
In May 2015, NASA reported that the Kepler space observatory observed KSN 2011b, a type Ia
supernova in the process of exploding. Details of the pre-nova moments may help scientists better
judge the quality of Type Ia supernovae as standard candles, which is an important link in the
argument for dark energy.[6]
In September 2021, astronomers reported that the Hubble Space Telescope had taken three images
of a Type Ia supernova through a gravitational lens. This supernova appeared at three different times
in the evolution of its brightness due to the differing path length of the light in the three images; at
− 24, 92, and 107 days from peak luminosity. A fourth image will appear in 2037 allowing observation
of the entire luminosity cycle of the supernova.[7]
Consensus model[edit]
Spectrum of SN 1998aq, a type Ia supernova, one day after maximum light in the B band[8]
The Type Ia supernova is a subcategory in the Minkowski–Zwicky supernova classification scheme,
which was devised by German-American astronomer Rudolph Minkowski and Swiss astronomer
Fritz Zwicky.[9] There are several means by which a supernova of this type can form, but they share
a common underlying mechanism. Theoretical astronomers long believed the progenitor star for this
type of supernova is a white dwarf, and empirical evidence for this was found in 2014 when a Type
Ia supernova was observed in the galaxy Messier 82.[10] When a slowly-rotating[2] carbon–oxygen
white dwarf accretes matter from a companion, it can exceed the Chandrasekhar limit of about 1.44
M☉, beyond which it can no longer support its weight with electron degeneracy pressure. [11] In the
absence of a countervailing process, the white dwarf would collapse to form a neutron star, in an
accretion-induced non-ejective process,[12] as normally occurs in the case of a white dwarf that is
primarily composed of magnesium, neon, and oxygen.[13]
The current view among astronomers who model Type Ia supernova explosions, however, is that
this limit is never actually attained and collapse is never initiated. Instead, the increase in pressure
and density due to the increasing weight raises the temperature of the core, [3] and as the white
dwarf approaches about 99% of the limit,[14] a period of convection ensues, lasting approximately
1,000 years.[15] At some point in this simmering phase, a deflagration flame front is born, powered
by carbon fusion. The details of the ignition are still unknown, including the location and number of
points where the flame begins.[16] Oxygen fusion is initiated shortly thereafter, but this fuel is not
consumed as completely as carbon.[17]
G299 Type Ia supernova remnant.
Once fusion begins, the temperature of the white dwarf increases. A main sequence star supported
by thermal pressure can expand and cool which automatically regulates the increase in thermal
energy. However, degeneracy pressure is independent of temperature; white dwarfs are unable to
regulate temperature in the manner of normal stars, so they are vulnerable to runaway fusion
reactions. The flare accelerates dramatically, in part due to the Rayleigh–Taylor instability and
interactions with turbulence. It is still a matter of considerable debate whether this flare transforms
into a supersonic detonation from a subsonic deflagration.[15][18]
Regardless of the exact details of how the supernova ignites, it is generally accepted that a
substantial fraction of the carbon and oxygen in the white dwarf fuses into heavier elements within a
period of only a few seconds,[17] with the accompanying release of energy increasing the internal
temperature to billions of degrees. The energy released (1–2×1044 J)[4] is more than sufficient to
unbind the star; that is, the individual particles making up the white dwarf gain enough kinetic energy
to fly apart from each other. The star explodes violently and releases a shock wave in which matter
is typically ejected at speeds on the order of 5,000–20,000 km/s, roughly 6% of the speed of light.
The energy released in the explosion also causes an extreme increase in luminosity. The typical
visual absolute magnitude of Type Ia supernovae is Mv = − 19.3 (about 5 billion times brighter than
the Sun), with little variation.[15]
The theory of this type of supernova is similar to that of novae, in which a white dwarf accretes
matter more slowly and does not approach the Chandrasekhar limit. In the case of a nova, the
infalling matter causes a hydrogen fusion surface explosion that does not disrupt the star. [15]
Type Ia supernovae differ from Type II supernovae, which are caused by the cataclysmic explosion
of the outer layers of a massive star as its core collapses, powered by release of gravitational
potential energy via neutrino emission.[19]
Formation[edit]
Formation process
An accretion disc forms around a compact body (such as a white dwarf) stripping gas from a companion
giant star. NASA image
Supercomputer simulation of the explosion phase of the deflagration-to-detonation model of supernova
formation.
Single degenerate progenitors[edit]
One model for the formation of this category of supernova is a close binary star system. The
progenitor binary system consists of main sequence stars, with the primary possessing more mass
than the secondary. Being greater in mass, the primary is the first of the pair to evolve onto the
asymptotic giant branch, where the star's envelope expands considerably. If the two stars share a
common envelope then the system can lose significant amounts of mass, reducing the angular
momentum, orbital radius and period. After the primary has degenerated into a white dwarf, the
secondary star later evolves into a red giant and the stage is set for mass accretion onto the primary.
During this final shared-envelope phase, the two stars spiral in closer together as angular
momentum is lost. The resulting orbit can have a period as brief as a few hours. [20][21] If the
accretion continues long enough, the white dwarf may eventually approach the Chandrasekhar limit.
The white dwarf companion could also accrete matter from other types of companions, including a
subgiant or (if the orbit is sufficiently close) even a main sequence star. The actual evolutionary
process during this accretion stage remains uncertain, as it can depend both on the rate of accretion
and the transfer of angular momentum to the white dwarf companion. [22]
It has been estimated that single degenerate progenitors account for no more than 20% of all Type
Ia supernovae.[23]
Double degenerate progenitors[edit]
A second possible mechanism for triggering a Type Ia supernova is the merger of two white dwarfs
whose combined mass exceeds the Chandrasekhar limit. The resulting merger is called a superChandrasekhar mass white dwarf.[24][25] In such a case, the total mass would not be constrained by
the Chandrasekhar limit.
Collisions of solitary stars within the Milky Way occur only once every 10 7 to 1013 years; far less
frequently than the appearance of novae.[26] Collisions occur with greater frequency in the dense
core regions of globular clusters[27] (cf. blue stragglers). A likely scenario is a collision with a binary
star system, or between two binary systems containing white dwarfs. This collision can leave behind
a close binary system of two white dwarfs. Their orbit decays and they merge through their shared
envelope.[28] A study based on SDSS spectra found 15 double systems of the 4,000 white dwarfs
tested, implying a double white dwarf merger every 100 years in the Milky Way: this rate matches
the number of Type Ia supernovae detected in our neighborhood.[29]
A double degenerate scenario is one of several explanations proposed for the anomalously massive
(2 M☉) progenitor of SN 2003fg.[30][31] It is the only possible explanation for SNR 0509-67.5, as all
possible models with only one white dwarf have been ruled out. [32] It has also been strongly
suggested for SN 1006, given that no companion star remnant has been found there.[23]
Observations made with NASA's Swift space telescope ruled out existing supergiant or giant
companion stars of every Type Ia supernova studied. The supergiant companion's blown out outer
shell should emit X-rays, but this glow was not detected by Swift's XRT (X-ray telescope) in the 53
closest supernova remnants. For 12 Type Ia supernovae observed within 10 days of the explosion,
the satellite's UVOT (ultraviolet/optical telescope) showed no ultraviolet radiation originating from the
heated companion star's surface hit by the supernova shock wave, meaning there were no red
giants or larger stars orbiting those supernova progenitors. In the case of SN 2011fe, the companion
star must have been smaller than the Sun, if it existed.[33] The Chandra X-ray Observatory revealed
that the X-ray radiation of five elliptical galaxies and the bulge of the Andromeda Galaxy is 30–50
times fainter than expected. X-ray radiation should be emitted by the accretion discs of Type Ia
supernova progenitors. The missing radiation indicates that few white dwarfs possess accretion
discs, ruling out the common, accretion-based model of Ia supernovae.[34] Inward spiraling white
dwarf pairs are strongly-inferred candidate sources of gravitational waves, although they have not
been directly observed.
Double degenerate scenarios raise questions about the applicability of Type Ia supernovae as
standard candles, since total mass of the two merging white dwarfs varies significantly, meaning
luminosity also varies.
Type Iax[edit]
It has been proposed that a group of sub-luminous supernovae that occur when helium accretes
onto a white dwarf should be classified as Type Iax.[35][36] This type of supernova may not always
completely destroy the white dwarf progenitor, but instead leave behind a zombie star.[37]
Observation[edit]
Supernova remnant N103B taken by the Hubble Space Telescope.[38]
Unlike the other types of supernovae, Type Ia supernovae generally occur in all types of galaxies,
including ellipticals. They show no preference for regions of current stellar formation. [39] As white
dwarf stars form at the end of a star's main sequence evolutionary period, such a long-lived star
system may have wandered far from the region where it originally formed. Thereafter a close binary
system may spend another million years in the mass transfer stage (possibly forming persistent nova
outbursts) before the conditions are ripe for a Type Ia supernova to occur. [40]
A long-standing problem in astronomy has been the identification of supernova progenitors. Direct
observation of a progenitor would provide useful constraints on supernova models. As of 2006, the
search for such a progenitor had been ongoing for longer than a century. [41] Observation of the
supernova SN 2011fe has provided useful constraints. Previous observations with the Hubble Space
Telescope did not show a star at the position of the event, thereby excluding a red giant as the
source. The expanding plasma from the explosion was found to contain carbon and oxygen, making
it likely the progenitor was a white dwarf primarily composed of these elements. [42] Similarly,
observations of the nearby SN PTF 11kx,[43] discovered January 16, 2011 (UT) by the Palomar
Transient Factory (PTF), lead to the conclusion that this explosion arises from single-degenerate
progenitor, with a red giant companion, thus suggesting there is no single progenitor path to SN Ia.
Direct observations of the progenitor of PTF 11kx were reported in the August 24 edition of Science
and support this conclusion, and also show that the progenitor star experienced periodic nova
eruptions before the supernova – another surprising discovery. [43][44] However, later analysis
revealed that the circumstellar material is too massive for the single-degenerate scenario, and fits
better the core-degenerate scenario.[45]
Light curve[edit]
This plot of luminosity (relative to the Sun, L0) versus time shows the characteristic light curve for a Type Ia
supernova. The peak is primarily due to the decay of nickel (Ni), while the later stage is powered by cobalt
(Co).
Light curve for type Ia SN 2018gv
Type Ia supernovae have a characteristic light curve, their graph of luminosity as a function of time
after the explosion. Near the time of maximal luminosity, the spectrum contains lines of intermediate-
mass elements from oxygen to calcium; these are the main constituents of the outer layers of the
star. Months after the explosion, when the outer layers have expanded to the point of transparency,
the spectrum is dominated by light emitted by material near the core of the star, heavy elements
synthesized during the explosion; most prominently isotopes close to the mass of iron (iron-peak
elements). The radioactive decay of nickel-56 through cobalt-56 to iron-56 produces high-energy
photons, which dominate the energy output of the ejecta at intermediate to late times. [15]
The use of Type Ia supernovae to measure precise distances was pioneered by a collaboration of
Chilean and US astronomers, the Calán/Tololo Supernova Survey.[46] In a series of papers in the
1990s the survey showed that while Type Ia supernovae do not all reach the same peak luminosity,
a single parameter measured from the light curve can be used to correct unreddened Type Ia
supernovae to standard candle values. The original correction to standard candle value is known as
the Phillips relationship[47] and was shown by this group to be able to measure relative distances to
7% accuracy.[48] The cause of this uniformity in peak brightness is related to the amount of nickel-56
produced in white dwarfs presumably exploding near the Chandrasekhar limit. [49]
The similarity in the absolute luminosity profiles of nearly all known Type Ia supernovae has led to
their use as a secondary standard candle in extragalactic astronomy. [50] Improved calibrations of the
Cepheid variable distance scale[51] and direct geometric distance measurements to NGC 4258 from
the dynamics of maser emission[52] when combined with the Hubble diagram of the Type Ia
supernova distances have led to an improved value of the Hubble constant.
In 1998, observations of distant Type Ia supernovae indicated the unexpected result that the
universe seems to undergo an accelerating expansion.[53][54] Three members from two teams were
subsequently awarded Nobel Prizes for this discovery.[55]
Subtypes[edit]
Supernova remnant SNR 0454-67.2 is likely the result of a Type Ia supernova explosion.[56]
There is significant diversity within the class of Type Ia supernovae. Reflecting this, a plethora of
sub-classes have been identified. Two prominent and well-studied examples include 1991T-likes, an
overluminous
subclass that exhibits particularly strong iron absorption lines and
abnormally small silicon features,[57] and 1991bg-likes, an exceptionally dim
subclass characterized by strong early titanium absorption features and rapid photometric and
spectral evolution.[58] Despite their abnormal luminosities, members of both peculiar groups can be
standardized by use of the Phillips relation to determine distance.[59]
See also[edit]
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Carbon detonation
Cosmic distance ladder
History of supernova observation
List of supernova remnants
Supernova remnant
MAGNETIC ACCRETION IN CATACLYSMIC
VARIABLES: SEARCH FOR NEW OBJECTS AND
SIMULTANEOUS MODELLING USING X-RAY
EMISSION AND OPTICAL POLARIZATION
PAB305
Isabel Lima, (National Institute of Space Research, Brazil)
January 22, 2019 @ 12:00 PM - 1:00 PM
Magnetic cataclysmic variables are compact binary systems in which mass transfer occurs
from a low-mass star onto a magnetic white dwarf. In polars and intermediate polars
subclasses, the magnetic field of the white dwarf is strong enough to disrupt the inner
accretion disk or even completely prevent disk formation. SW Sextantis stars are a type of
novalike cataclysmic variable, which show observational characteristics that indicate
matter outside the orbital plane is distributed asymmetrically in azimuth. Many scenarios
have been proposed to explain the SW Sex behavior, one of them is magnetic accretion. As
a first part of the project, we intend to search for signs of magnetic accretion in SW Sex
systems such as circular polarization and/or coherent photometric and/or polarimetric
variability that can be associated with the white dwarf spin period. In this project, we aim for
the discovery of magnetic cataclysmic variables from large surveys (such as ZTF, CRTS).
Magnetic cataclysmic variables candidates or confirmed objects can be further studied
using the CYCLOPS code.This code performs multi-wavelength fitting of the accretion
column flux using simultaneous photometric and polarimetric light curves and X-ray spectra
and light curves. It considers cyclotron and free-free emission from a 3D post-shock region,
which is non-homogeneous in terms of density, temperature and magnetic field. This kind
of research is inline with my PhD, which is about modeling the optical and X-ray data of IPs.
Cataclysmic variable stars (CV) are stars which irregularly increase in brightness by a large factor,
then drop back down to a quiescent state. They were initially called novae, from the Latin 'new',
since ones with an outburst brightness visible to the naked eye and an invisible quiescent brightness
appeared as new stars in the sky.
Cataclysmic variable stars are binary stars that consist of two components; a white dwarf primary,
and a mass transferring secondary. The stars are so close to each other that the gravity of the white
dwarf distorts the secondary, and the white dwarf accretes matter from the companion. Therefore,
the secondary is often referred to as the donor star. The infalling matter, which is usually rich in
hydrogen, forms in most cases an accretion disk around the white dwarf. Strong UV and X-ray
emission is often seen from the accretion disc, powered by the loss of gravitational potential energy
from the infalling material.[citation needed]
Material at the inner edge of disc falls onto the surface of the white dwarf primary. A classical nova
outburst occurs when the density and temperature at the bottom of the accumulated hydrogen layer
rise high enough to ignite runaway hydrogen fusion reactions, which rapidly convert the hydrogen
layer to helium. If the accretion process continues long enough to bring the white dwarf close to the
Chandrasekhar limit, the increasing interior density may ignite runaway carbon fusion and trigger a
Type Ia supernova explosion, which would completely destroy the white dwarf.
The accretion disc may be prone to an instability leading to dwarf nova outbursts, when the outer
portion of the disc changes from a cool, dull mode to a hotter, brighter mode for a time, before
reverting to the cool mode. Dwarf novae can recur on a timescale of days to decades.
Contents
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1
Classification
2
Discovery
3
Superhumps
4
References
5
External links
Classification[edit]
Cataclysmic variables are subdivided into several smaller groups, often named after a bright
prototype star characteristic of the class. In some cases the magnetic field of the white dwarf is
strong enough to disrupt the inner accretion disk or even prevent disk formation altogether. Magnetic
systems often show strong and variable polarization in their optical light, and are therefore
sometimes called polars; these often exhibit small-amplitude brightness fluctuations at what is
presumed to be the period of rotation of the white dwarf.
Supernovae
These are classed as cataclysmic variables and have extremely
large outbursts that destroy the progenitor star. Some result from
white dwarfs in binary systems but others are very massive stars.
(Classical) novae
These cataclysmic variables have very large outbursts, of 6 to 19
magnitudes, caused by thermonuclear fusion of material accreted
onto the white dwarf.
Recurrent novae
These have outbursts of about 4 to 9 magnitudes, repeating every
10 to 80 years.[1] Examples include T Pyxidis and RS Ophiuchi.
Dwarf novae
Dwarf novae, or U Geminorum stars, are cataclysmic variables which are
observed to brighten repeatedly, though by a smaller amount than
classical novae.
Z Camelopardalis starsTemporarily "halt" at a particular brightness below
their peak
SU Ursae Majoris starsHave "superoutbursts" which are brighter than the
average
SS Cygni stars
Luminous red novae
Have outbursts of two distinct lengths
These are stellar mergers that become very red after outburst.
Polars
AM Herculis stars are binaries in which the magnetic field of the white
dwarf has synchronized the latter's rotational period with the binary
orbital period. Matter from the donor star is magnetically channeled onto
the white dwarf rather than forming a disc.
DQ Herculis, also called 'intermediate polars', have a slightly weaker
magnetic field than AM Herculis stars; there is an accretion disc, but
substructure in it is created by the field.
VY Sculptoris
These are stars which occasionally drop in brightness by more
than one magnitude, with very occasional dwarf-nova-type
outbursts during the dim state. They may be a subclass of
polars.[2]
AM Canum Venaticorum
These are cataclysmic variables both of whose components are
white dwarfs; the accretion disc is composed primarily of helium,
and they are of interest as sources of gravitational waves.
SW Sextantis
These are like dwarf novae but have the accretion disc in a steady
state, so do not show outbursts; the disc emits non-uniformly.
They are usually also eclipsing variables, though this appears to
be a selection artefact.[3]
Z Andromedae
(symbiotic variables)
These are close binaries with a large cool component losing mass
to a hotter compact component and accretion disc.
There are over 1600 known CV systems.[4] The catalog was frozen as of 1 February 2006 though
more are discovered each year.
Discovery[edit]
Cataclysmic variables are among the classes of astronomical objects most commonly found by
amateurs, since a cataclysmic variable in its outburst phase is bright enough to be detectable with
very modest instruments, and the only celestial objects easily confused with them are bright
asteroids whose movement from night to night is clear.
Verifying that an object is a cataclysmic variable is also fairly straightforward: they are usually quite
blue objects, they exhibit rapid and strong variability, and they tend to have peculiar emission lines.
They emit in the ultraviolet and X-ray ranges; they are expected also to emit gamma rays, from
annihilation of positrons from proton-rich nuclei produced in the fusion explosion, but this has not yet
been detected.[5]
Around six galactic novae (i.e. in our own galaxy) are discovered each year, whilst models based on
observations in other galaxies suggest that the rate of occurrence ought to be between 20 and 50; [6]
this discrepancy is due partly to obscuration by interstellar dust, and partly to a lack of observers in
the southern hemisphere and to the difficulties of observing while the Sun is up and at full moon.
Superhumps[edit]
Main article: Superhump
Some cataclysmic variables experience periodic brightenings caused by deformations of the
accretion disk when its rotation is in resonance with the orbital period of the binary.
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Advanced
Basic
Spectra and What They Can Tell Us
A rainbow rises over a misty forest. (Credit: U.S. Fish and Wildlife Service)
A spectrum is simply a chart or a graph that shows the intensity of light being emitted over a
range of energies. Have you ever seen a spectrum before? Probably. Nature makes beautiful
ones we call rainbows. Sunlight sent through raindrops is spread out to display its various colors
(the different colors are just the way our eyes perceive radiation with slightly different energies).
Spectroscopy can be very useful in helping scientists understand how an object like a black
hole, neutron star, or active galaxy produces light, how fast it is moving, and what elements it
is composed of. Spectra can be produced for any energy of light, from low-energy radio waves
to very high-energy gamma rays.
Each spectrum holds a wide variety of information. For instance, there are many different
mechanisms by which an object, like a star, can produce light. Each of these mechanisms has a
characteristic spectrum.
The Electromagnetic Spectrum
White light (what we call visible or optical light) can be split up into its constituent colors easily
and with a familiar result: the rainbow. All we have to do is use a slit to focus a narrow beam of
the light at a prism. This setup is actually a basic spectrometer.
Spectrum of white light
The resultant rainbow is really a continuous spectrum that shows us the different energies of
light (from red to blue) present in visible light. But the electromagnetic spectrum encompasses
more than just optical light. It covers all energies of light, extending from low-energy radio
waves, to microwaves, to infrared, to optical light, to ultraviolet, to very high-energy X-rays and
gamma rays.
The full electromagnetic spectrum. (Credit: NASA's Imagine the Universe)
Tell Me More About the Electromagnetic Spectrum!
What Can Scientists Learn From a Spectrum?
Continuous
Emission or Bright Line
Absorption or Dark Line
Three types of spectra: continuous, emission line and absorption.(Credit: NASA's Imagine the
Universe)
Each element in the periodic table can appear in gaseous form and will produce a series of
bright lines unique to that element. Hydrogen will not look like helium which will not look like
carbon which will not look like iron... and so on. Thus, astronomers can identify what kinds of
stuff are in stars from the lines they find in the star's spectrum. This type of study is called
spectroscopy.
The science of spectroscopy is quite sophisticated. From spectral lines astronomers can
determine not only the element, but the temperature and density of that element in the star. The
spectral line also can tell us about any magnetic field of the star. The width of the line can tell
us how fast the material is moving. We can learn about winds in stars from this. If the lines shift
back and forth we can learn that the star may be orbiting another star. We can estimate the
mass and size of the star from this. If the lines grow and fade in strength we can learn about the
physical changes in the star. Spectral information can also tell us about material around stars.
This material may be falling onto the star from a doughnut-shaped disk around the star called an
accretion disk. These disks often form around a neutron star or black hole. The light from the
stuff between the stars allows astronomers to study the interstellar medium (ISM). This tells
us what type of stuff fills the space between the stars. Space is not empty! There is lots of gas
and dust between the stars. Spectroscopy is one of the fundamental tools which scientists use
to study the Universe.
Spectral Analysis
In a star, there are many elements present. We can tell which ones are there by looking at the
spectrum of the star. The science of spectroscopy is quite sophisticated. From spectral lines
astronomers can determine not only the element, but the temperature and density of that
element in the star. The lines can also tell us about the magnetic field of the star. The width of
the line can tell us how fast the material is moving, giving us information about stellar wind. If
the lines shift back and forth, it means that the star may be orbiting another star - the spectrum
will give the information necessary to estimating the mass and size of the star system and the
companion star. If the lines grow and fade in strength we can learn about the physical changes
in the star.
Spectral information, particularly from energies of light other than optical, can tell us about
material around stars. This material may have been pulled from a companion star by a black
hole or a neutron star, where it will form an orbiting disk. Around a compact object (black hole,
neutron star), the material in this accretion disk is heated to the point that it gives off X-rays,
and the material eventually falls onto the black hole or neutron star. It is by looking at the
spectrum of X-rays being emitted by that object and its surrounding disk that we can learn about
the nature of these objects.
As discussed in the introduction to spectra, a spectrum is simply a chart or graph of the intensity
of light of a source as a function of the energy of that light. Each spectrum holds a wide variety
of information. For instance, there are many different mechanisms by which an object, like a
star, can produce light. Each of these mechanisms has a characteristic spectrum.
Look at the sample spectrum below.
X-ray spectrum of supernova remnant Cas A from ASCA data. (Credit: Holt et al., PASJ 1994)
This spectrum was taken by the ASCA satellite. It is an X-ray spectrum of a supernova
remnant (SNR), which is a huge cloud of gaseous matter swept up from the explosion of a
massive star. The x-axis shows the range of energy of light that is being emitted (the units are
keV, or kilo-electron Volts, which are the common units astronomers use to talk about X-ray
light). The y-axis of the graph shows the intensity of the light being emitted by the SNR — that
is, how many photons of light the SNR is giving off at each energy.
There are two main types of spectra in this graph – a continuum and emission lines. Notice that
the overall curve decreases with higher energies. This decreasing curve is called a continuum,
and is produced by a process that emits photons at all energies. That curve, though, has bumps
on it, which are marked by the symbols for various elements like neon (Ne), calcium (Ca), and
iron (Fe). These bumps are called line emission and occur at very specific energies. These two
different types of X-ray emission (continuum and line) are produced by different processes and
each tells us different things about the source emitting them.
Below we'll go into more detail about line and continuum emission, such as what mechanisms
cause them, and what they can tell us about the light-emitting object. However, you might first
want to understand a little more about atoms and about how spectra are represented
graphically. Use the links below to read more about those topics.
Tell Me More About Understanding the Atom!
Tell Me More About Graphical Representations of Spectra!
Line Emission
Remember that if we use a spectrometer on white light (like an ordinary light bulb), we'll see a
rainbow. However, what if instead we used our spectrometer to look at a tube of pure gas, like
hydrogen? First, we need to heat the hydrogen to a very high temperature, or give the atoms of
hydrogen energy by running an electric current through the tube. This would cause the gas to
glow, or in other words, to emit radiation.
A hydrogen emission lamp. (Credit: Photo by Wilco Oelen via Wikimedia Commons
If we looked at the spectrum of light given off by the hydrogen gas with our spectroscope,
instead of seeing a continuum of colors, we would just see a few bright lines. Below we see the
spectrum, which shows the unique fingerprint of hydrogen.
The spectrum of hydrogen. The bright lights are emission lines, and the pattern is unique to
hydrogen. Other elements show different patterns of emission lines, making this spectrum a
"fingerprint" of the element hydrogen.
The bright lines are called emission lines. Remember how we heated the hydrogen to give the
atoms energy? By doing that, we excited the electrons in the atom. When the electrons fell
back to their ground state, they gave off photons of light at hydrogen's characteristic energies. If
we altered the amount or abundance of hydrogen gas we have, we could change the intensity
and brightness of the lines, because more photons would be produced. But we couldn't change
their color, because no matter how much or how little hydrogen gas was present, the pattern of
lines would be the same. Hydrogen's pattern of emission lines is unique to it. The brightness of
the emission lines can give us a great deal of information about the abundance of hydrogen
present. This is particularly useful in a star, where there are many elements mixed together.
Each element in the periodic table can appear in gaseous form and will each produce a series
of bright emission lines unique to that element. The spectrum of hydrogen will not look like the
spectrum of helium, or the spectrum of carbon, or of any other element.
Hydrogen
Helium
Carbon
We know that the continuum of the electromagnetic spectrum extends from low-energy radio
waves, to microwaves, to infrared, to optical light, to ultraviolet, to X-rays and gamma rays. In
the same way, hydrogen's unique spectrum extends over a range, as do the spectra of the other
elements. The above spectra are in the optical range of light. Line emission can actually occur
at any energy of light (i.e. visible, UV, etc. ) and with any type of atom. However, not all atoms
have line emission at all wavelengths. The difference in energy between levels in the atom is
not great enough for the emission to be X-rays in atoms of lighter elements, for example.
Continuum Emission
We'll focus this discussion of continuum emission on those processes that produce X-rays. Just
like visible light, with its range of energies from red to blue, X-rays have a continuum, or a range
of energies associated with it. X-rays usually range in energy from around 0.5 keV up to around
1000 keV.
Like line emission, continuum X-ray emission involves charged particles. Continuum emission is
a result of the acceleration of a population of charged particles. All X-ray sources contain such
particles. These particles must be at least partially ionized, which means their electrons need to
be unbound from their nuclei to be free to zip around when they are heated to extreme
temperatures. For an electron to radiate X-rays, the gas containing the electron must have
extreme conditions, such as temperatures of millions of degrees, super-strong magnetic fields,
or the electrons themselves must be moving at nearly the speed of light. Extreme conditions
can be found in disks of matter orbiting black holes or in supernova remnants. Strong magnetic
fields, like those created in the wake of a supernova explosion, can also accelerate fast-moving
ions in spirals around the field lines to the point of X-ray emission. Electrons can be accelerated
to nearly the speed of light in the shockwave created by a supernova explosion.
There are three mechanisms that will produce continuum X-ray emission. They are
Synchrotron Radiation, Bremsstrahlung, and Compton Scattering. Because the
populations of electrons have a continuous range of energies, and they can be accelerated
through a range of energies, the radiation produced is continuous, and not at the discreet
energies of line emission.
Illustration of the process of synchrotron radiation. (Credit: NASA's Imagine the Universe)
Synchrotron radiation is emitted when a fast electron interacts with a magnetic field. A
magnetic field in an area where an electron is traveling will cause the electron to change
direction by exerting a force on it perpendicular to the direction the electron is moving. As a
result, the electron will be accelerated, causing it to radiate electromagnetic energy. This is
called magnetic bremsstrahlung or synchrotron radiation (after radiation observed from
particle accelerators by that name). If the electrons and the magnetic field are energetic enough,
the emitted radiation can be in the form of X-rays.
Gas that is at about 1 million to 10 million degrees, such as the gas heated by a supernova
explosion, produces most of its emission in X-rays from thermal bremsstrahlung. Gas can be
heated to these temperatures by the resulting shockwave of a supernova explosion, or in an
accretion disk around a black hole or neutron star.
Illustration of the process of Bremsstralung radiation. (Credit: NASA's Imagine the Universe)
Bremsstrahlung occurs when an electron passes close to a positive ion, and the strong electric
forces cause its trajectory to change. The acceleration of the electron in this way causes it to
radiate electromagnetic energy. This radiation is called bremsstrahlung, which in German
translates literally to "braking radiation." Thermal bremsstrahlung occurs in a hot gas, where
many electrons are stripped from their nuclei, leaving a population of electrons and positive
ions. If the gas is hot enough (millions of degrees Celsius), this kind of radiation will primarily
take the form of X-rays.
Supernova remnants (SNRs) have strong magnetic fields where ions can be accelerated by the
shock wave of the supernova to high energies, producing X-ray synchrotron radiation. X-rays
produced by SNR require electrons with energies of about 10 4 GeV each. We would have to
heat an electron to a temperature of about ten trillion degrees for it to have this much energy!
Video illustration of Bremsstralung radiation
Credit:NASA
Comptonization is when a photon collides with an electron - the photon will either give up
energy to or gain energy from the electron, changing the electron's velocity as a result.
Compton scattered radiation and synchrotron radiation are major components of the diffuse Xray background and emission from active galaxies.
Illustration of the process of Compton radiation. (Credit: NASA's Imagine the Universe)
Light Curves and What They Can Tell Us
Images show a scientist where in an object light is emitted. Another piece of information we
have about light is when it reaches the detector. Astronomers use this "timing" information to
create light curves and perform timing analysis.
Tell me more about the history of timing in astronomy
Light curves are graphs that show the brightness of an object over a period of time. In the study
of objects which change their brightness over time, such as novae, supernovae, and variable
stars, the light curve is a simple but valuable tool to a scientist.
This video compares the X-ray 'heartbeats', or light curves, of GRS 1915 and IGR J17091, two
black holes that ingest gas from companion stars. (Credit: NASA's Goddard Space Flight
Center)
Watch on YouTube.com
If we had the following information about a particular source:
Brightness
(Magnitude)
Date
Brightness
(Magnitude)
Date
April 21
9.2
June 20
8.7
April 27
9.3
June 26
8.3
May 3
9.7
July 2
8.6
May 9
9.9
July 8
9.1
May 15
9.6
July 14
9.1
May 21
9.8
July 20
9.2
May 27
9.9
July 26
9.5
June 2
9.7
Aug 1
9.9
June 8
9.1
Aug 7
9.7
June 14
8.8
Aug 13
9.7
then we might make a light curve, and it would look like this:
A simple light curve made from the data in the table above
The plot shows the brightness of a certain astronomical object viewed through a telescope
every 6 days over the course of a few months. This gives us a light curve of the object we have
measured. But light curves can be generated for any measure of brightness which is measured
over time. So, if we measured the number of X-rays being emitted by a star during every
second for an hour, we could generate a light curve from our observations.
What can we learn from light curves?
The record of changes in brightness that a light curve provides can help astronomers
understand processes at work within the object they are studying and identify specific
categories (or classes) of stellar events. We know generally what light curves look like for a set
of objects, so when we plot a new light curve, we can compare it to those standard light curves
to possibly identify the type of object we're observing.
If the light curve we measured looked like the graph below, we would identify the object as an
eclipsing binary star. The light curve also shows us that it takes 10 days for one of the stars in
the binary to orbit completely around the other. Astronomers would say that the binary system
has an orbital period of 10 days.
Light curve of an eclipsing binary system
If, instead, the light curve we measured looked like the one below, we would know that this
object was the death of a star by a massive explosion called a supernova!
Light curve of a supernova
Multiwavelength Astronomy
The sky as seen in different wavelengths. From the top: radio, infrared, optical, X-ray, and
gamma-ray. (Click image for a larger version. Credits: radio: Haslam et al. 1982; infrared:
NASA; optical: ESO/S. Brunier; X-ray: Max Planck Institute for Extraterrestrial Physics and S.
L. Snowden; gamma-ray: NASA/DOE/Fermi LAT Collaboration)
The night sky has always served as a source of wonder and mystery to people. However, it has
only been in the past few decades that we have truly begun to 'see' the Universe in all its glory.
This is because we have only recently been able to look at the Universe over the entire
electromagnetic spectrum. Our Universe contains objects that produce a vast range of
radiation with wavelengths either too short or too long for our eyes to see.
Instruments that examine all parts of the electromagnetic (EM) spectrum have been available to
us only in the 20th century, since the rocket age was required to get instruments sensitive to the
infrared, ultraviolet, X-rays, and gamma-ray wavelengths above the Earth. We had to wait
until the rocket age to get a view from space, which is extremely important, since radiation in
these wavelengths is absorbed by the Earth's atmosphere).
Tell me more about the electromagnetic spectrum!
Some astronomical objects emit mostly infrared radiation, others mostly visible light, and still
others mostly ultraviolet radiation. What determines the type of electromagnetic radiation
emitted by astronomical objects? The simple answer is temperature.
A solid contains molecules and atoms that are in continuous vibration. Inside a gas are
molecules that are flying about freely at high rates, continually bumping into each other and
surrounding matter. The energy of the motions of the molecules or atoms is called heat. The
hotter the solid or gas, the more rapid the motion of the molecules or atoms. And temperature is
just a measure of the average energy of those particles.
What is the relationship between temperature and electromagnetic radiation? And what objects
do we see at different temperatures or in different regions of the electromagnetic spectrum?
Check out the chart below to see.
Type Of
Radiation
Gamma-rays
Radiated by
Objects at this
Temperature
> 108 Kelvin (K)
Typical Sources
accretion disks around black holes
X-rays
106-108 K
gas in clusters of galaxies;
supernova remnants; stellar corona
Ultraviolet
104-106 K
supernova remnants; very hot stars
Visible
103-104 K
planets, stars, some satellites
Infrared
10-103 K
cool clouds of dust and gas; planets
Microwave
1-10 K
cool clouds of gas, including those
around newly formed stars; the
cosmic microwave background
Radio
<1K
radio emission produced by electrons
moving in magnetic fields
Black body[edit]
Main article: Black body
All normal (baryonic) matter emits electromagnetic radiation when it has a temperature above
absolute zero. The radiation represents a conversion of a body's internal energy into electromagnetic
energy, and is therefore called thermal radiation. It is a spontaneous process of radiative distribution
of entropy.
Color of a black body from 800 K to 12200 K. This range of colors approximates the range of colors of stars
of different temperatures, as seen or photographed in the night sky.
Conversely, all normal matter absorbs electromagnetic radiation to some degree. An object that
absorbs all radiation falling on it, at all wavelengths, is called a black body. When a black body is at
a uniform temperature, its emission has a characteristic frequency distribution that depends on the
temperature. Its emission is called black-body radiation.
The concept of the black body is an idealization, as perfect black bodies do not exist in nature.[17]
Graphite and lamp black, with emissivities greater than 0.95, however, are good approximations to a
black material. Experimentally, black-body radiation may be established best as the ultimately stable
steady state equilibrium radiation in a cavity in a rigid body, at a uniform temperature, that is entirely
opaque and is only partly reflective.[17] A closed box with walls of graphite at a constant temperature
with a small hole on one side produces a good approximation to ideal black-body radiation
emanating from the opening.[18][19]
Black-body radiation has the unique absolutely stable distribution of radiative intensity that can
persist in thermodynamic equilibrium in a cavity.[17] In equilibrium, for each frequency the total
intensity of radiation that is emitted and reflected from a body (that is, the net amount of radiation
leaving its surface, called the spectral radiance) is determined solely by the equilibrium temperature
and does not depend upon the shape, material or structure of the body. [20] For a black body (a
perfect absorber) there is no reflected radiation, and so the spectral radiance is entirely due to
emission. In addition, a black body is a diffuse emitter (its emission is independent of direction).
Consequently, black-body radiation may be viewed as the radiation from a black body at thermal
equilibrium.
Black-body radiation becomes a visible glow of light if the temperature of the object is high
enough.[21] The Draper point is the temperature at which all solids glow a dim red, about 798 K. [22]
At 1000 K, a small opening in the wall of a large uniformly heated opaque-walled cavity (such as an
oven), viewed from outside, looks red; at 6000 K, it looks white. No matter how the oven is
constructed, or of what material, as long as it is built so that almost all light entering is absorbed by
its walls, it will contain a good approximation to black-body radiation. The spectrum, and therefore
color, of the light that comes out will be a function of the cavity temperature alone. A graph of the
amount of energy inside the oven per unit volume and per unit frequency interval plotted versus
frequency is called the black-body curve. Different curves are obtained by varying the temperature.
The temperature of a Pāhoehoe lava flow can be estimated by observing its color. The result agrees well
with other measurements of temperatures of lava flows at about 1,000 to 1,200 °C (1,830 to 2,190 °F).
Two bodies that are at the same temperature stay in mutual thermal equilibrium, so a body at
temperature T surrounded by a cloud of light at temperature T on average will emit as much light into
the cloud as it absorbs, following Prevost's exchange principle, which refers to radiative equilibrium.
The principle of detailed balance says that in thermodynamic equilibrium every elementary process
works equally in its forward and backward sense.[23][24] Prevost also showed that the emission from
a body is logically determined solely by its own internal state. The causal effect of thermodynamic
absorption on thermodynamic (spontaneous) emission is not direct, but is only indirect as it affects
the internal state of the body. This means that at thermodynamic equilibrium the amount of every
wavelength in every direction of thermal radiation emitted by a body at temperature T, black or not,
is equal to the corresponding amount that the body absorbs because it is surrounded by light at
temperature T.[25]
When the body is black, the absorption is obvious: the amount of light absorbed is all the light that
hits the surface. For a black body much bigger than the wavelength, the light energy absorbed at
any wavelength λ per unit time is strictly proportional to the black-body curve. This means that the
black-body curve is the amount of light energy emitted by a black body, which justifies the name.
This is the condition for the applicability of Kirchhoff's law of thermal radiation: the black-body curve
is characteristic of thermal light, which depends only on the temperature of the walls of the cavity,
provided that the walls of the cavity are completely opaque and are not very reflective, and that the
cavity is in thermodynamic equilibrium.[26] When the black body is small, so that its size is
comparable to the wavelength of light, the absorption is modified, because a small object is not an
efficient absorber of light of long wavelength, but the principle of strict equality of emission and
absorption is always upheld in a condition of thermodynamic equilibrium.
In the laboratory, black-body radiation is approximated by the radiation from a small hole in a large
cavity, a hohlraum, in an entirely opaque body that is only partly reflective, that is maintained at a
constant temperature. (This technique leads to the alternative term cavity radiation.) Any light
entering the hole would have to reflect off the walls of the cavity multiple times before it escaped, in
which process it is nearly certain to be absorbed. Absorption occurs regardless of the wavelength of
the radiation entering (as long as it is small compared to the hole). The hole, then, is a close
approximation of a theoretical black body and, if the cavity is heated, the spectrum of the hole's
radiation (i.e., the amount of light emitted from the hole at each wavelength) will be continuous, and
will depend only on the temperature and the fact that the walls are opaque and at least partly
absorptive, but not on the particular material of which they are built nor on the material in the cavity
(compare with emission spectrum).
The radiance or observed intensity is not a function of direction. Therefore, a black body is a perfect
Lambertian radiator.
Real objects never behave as full-ideal black bodies, and instead the emitted radiation at a given
frequency is a fraction of what the ideal emission would be. The emissivity of a material specifies
how well a real body radiates energy as compared with a black body. This emissivity depends on
factors such as temperature, emission angle, and wavelength. However, it is typical in engineering to
assume that a surface's spectral emissivity and absorptivity do not depend on wavelength so that the
emissivity is a constant. This is known as the gray body assumption.
9-year WMAP image (2012) of the cosmic microwave background radiation across the universe.[27][28]
With non-black surfaces, the deviations from ideal black-body behavior are determined by both the
surface structure, such as roughness or granularity, and the chemical composition. On a "per
wavelength" basis, real objects in states of local thermodynamic equilibrium still follow Kirchhoff's
Law: emissivity equals absorptivity, so that an object that does not absorb all incident light will also
emit less radiation than an ideal black body; the incomplete absorption can be due to some of the
incident light being transmitted through the body or to some of it being reflected at the surface of the
body.
In astronomy, objects such as stars are frequently regarded as black bodies, though this is often a
poor approximation. An almost perfect black-body spectrum is exhibited by the cosmic microwave
background radiation. Hawking radiation is the hypothetical black-body radiation emitted by black
holes, at a temperature that depends on the mass, charge, and spin of the hole. If this prediction is
correct, black holes will very gradually shrink and evaporate over time as they lose mass by the
emission of photons and other particles.
A black body radiates energy at all frequencies, but its intensity rapidly tends to zero at high
frequencies (short wavelengths). For example, a black body at room temperature (300 K) with one
square meter of surface area will emit a photon in the visible range (390–750 nm) at an average rate
of one photon every 41 seconds, meaning that, for most practical purposes, such a black body does
not emit in the visible range.[29]
The study of the laws of black bodies and the failure of classical physics to describe them helped
establish the foundations of quantum mechanics.
Further explanation[edit]
According to the Classical Theory of Radiation, if each Fourier mode of the equilibrium radiation (in
an otherwise empty cavity with perfectly reflective walls) is considered as a degree of freedom
capable of exchanging energy, then, according to the equipartition theorem of classical physics,
there would be an equal amount of energy in each mode. Since there are an infinite number of
modes, this would imply infinite heat capacity, as well as a nonphysical spectrum of emitted radiation
that grows without bound with increasing frequency, a problem known as the ultraviolet catastrophe.
In the longer wavelengths this deviation is not so noticeable, as
are very small. In the shorter wavelengths of the ultraviolet range, however, classical theory predicts
the energy emitted tends to infinity, hence the ultraviolet catastrophe. The theory even predicted that
all bodies would emit most of their energy in the ultraviolet range, clearly contradicted by the
experimental data which showed a different peak wavelength at different temperatures (see also
Wien's law).
As the temperature increases, the peak of the emitted black-body radiation curve moves to higher
intensities and shorter wavelengths.[30] The black-body radiation graph is also compared with the classical
model of Rayleigh and Jeans.
Instead, in the quantum treatment of this problem, the numbers of the energy modes are quantized,
attenuating the spectrum at high frequency in agreement with experimental observation and
resolving the catastrophe. The modes that had more energy than the thermal energy of the
substance itself were not considered, and because of quantization modes having infinitesimally little
energy were excluded.
Thus for shorter wavelengths very few modes (having energy more than
{\displaystyle h\nu }
) were allowed, supporting the data that the energy emitted is reduced for wavelengths less than
the wavelength of the observed peak of emission.
Notice that there are two factors responsible for the shape of the graph. Firstly, longer wavelengths
have a larger number of modes associated with them. Secondly, shorter wavelengths have more
energy associated per mode. The two factors combined give the characteristic maximum
wavelength.
Calculating the black-body curve was a major challenge in theoretical physics during the late
nineteenth century. The problem was solved in 1901 by Max Planck in the formalism now known as
Planck's law of black-body radiation.[31] By making changes to Wien's radiation law (not to be
confused with Wien's displacement law) consistent with thermodynamics and electromagnetism, he
found a mathematical expression fitting the experimental data satisfactorily. Planck had to assume
that the energy of the oscillators in the cavity was quantized, i.e., it existed in integer multiples of
some quantity. Einstein built on this idea and proposed the quantization of electromagnetic radiation
itself in 1905 to explain the photoelectric effect. These theoretical advances eventually resulted in
the superseding of classical electromagnetism by quantum electrodynamics. These quanta were
called photons and the black-body cavity was thought of as containing a gas of photons. In addition,
it led to the development of quantum probability distributions, called Fermi–Dirac statistics and
Bose–Einstein statistics, each applicable to a different class of particles, fermions and bosons.
The wavelength at which the radiation is strongest is given by Wien's displacement law, and the
overall power emitted per unit area is given by the Stefan–Boltzmann law. So, as temperature
increases, the glow color changes from red to yellow to white to blue. Even as the peak wavelength
moves into the ultra-violet, enough radiation continues to be emitted in the blue wavelengths that the
body will continue to appear blue. It will never become invisible—indeed, the radiation of visible light
increases monotonically with temperature.[32] The Stefan–Boltzmann law also says that the total
radiant heat energy emitted from a surface is proportional to the fourth power of its absolute
temperature. The law was formulated by Josef Stefan in 1879 and later derived by Ludwig
Boltzmann. The formula E = σT4 is given, where E is the radiant heat emitted from a unit of area per
unit time, T is the absolute temperature, and σ = 5.670367×10− 8 W·m − 2⋅K− 4 is the Stefan–
Boltzmann constant.[33]
Color index
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For the colorant reference database, see Colour Index International. For the term in geology,
see Color index (geology).In astronomy, the color index is a simple numerical expression that
determines the color of an object, which in the case of a star gives its temperature. The smaller
the color index, the more blue (or hotter) the object is. Conversely, the larger the color index, the
more red (or cooler) the object is. This is a consequence of the logarithmic magnitude scale, in
which brighter objects have smaller (more negative) magnitudes than dimmer ones. For
comparison, the yellowish Sun has a B− V index of 0.656 ± 0.005,[2] whereas the bluish Rigel has
a B− V of − 0.03 (its B magnitude is 0.09 and its V magnitude is 0.12, B− V = − 0.03).[3]
Traditionally, the color index uses Vega as a zero point.
To measure the index, one observes the magnitude of an object successively through two
different filters, such as U and B, or B and V, where U is sensitive to ultraviolet rays, B is
sensitive to blue light, and V is sensitive to visible (green-yellow) light (see also: UBV system).
The set of passbands or filters is called a photometric system. The difference in magnitudes
found with these filters is called the U− B or B− V color index respectively.
In principle, the temperature of a star can be calculated directly from the B− V index, and there
are several formulae to make this connection.[4] A good approximation can be obtained by
considering stars as black bodies, using Ballesteros' formula[5] (also implemented in the
PyAstronomy package for Python):[6]
Color indices of distant objects are usually affected by interstellar extinction, that is, they are redder
than those of closer stars. The amount of reddening is characterized by color excess, defined as the
difference between the observed color index and the normal color index (or intrinsic color
index), the hypothetical true color index of the star, unaffected by extinction. For example, in the UBV
photometric system we can write it for the B− V color:
The passbands most optical astronomers use are the UBVRI filters, where the U, B, and V filters are
as mentioned above, the R filter passes red light, and the I filter passes infrared light. This system of
filters is sometimes called the Johnson–Cousins filter system, named after the originators of the
system (see references). These filters were specified as particular combinations of glass filters and
photomultiplier tubes. M. S. Bessell specified a set of filter transmissions for a flat response detector,
thus quantifying the calculation of the color indices.[7] For precision, appropriate pairs of filters are
chosen depending on the object's color temperature: B− V are for mid-range objects, U− V for hotter
objects, and R− I for cool ones.
Stellar Evolution
Stellar evolution is a description of the way that stars change with time. On human
timescales, most stars do not appear to change at all, but if we were to look for billions
of years, we would see how stars are born, how they age, and finally how they die.
The primary factor determining how a star evolves is its mass as it reaches the main
sequence. The following is a brief outline tracing the evolution of a low-mass and a
high-mass star.
The life of a star
Stars are born out of the gravitational collapse of cool, dense molecular
clouds. As the cloud collapses, it fragments into smaller regions, which
themselves contract to form stellar cores. These protostars rotate faster and
increase in temperature as they condense, and are surrounded by a
protoplanetary disk out of which planets may later form.
The central temperature of the contracting protostar increases to the point
where nuclear reactions begin. At this point, hydrogen is converted into
helium in the core and the star is born onto the main sequence. For about 90%
of its life, the star will continue to burn hydrogen into helium and will remain a
main sequence star.
Once the hydrogen in the core has all been burned to helium, energy generation
stops and the core begins to contract. This raises the internal temperature of
the star and ignites a shell of hydrogen burning around the inert core.
Meanwhile, the helium core continues to contract and increase in temperature,
which leads to an increased energy generation rate in the hydrogen shell. This
causes the star to expand enormously and increase in luminosity – the star
becomes a red giant.
Eventually, the core reaches temperatures high enough to burn helium into
carbon. If the mass of the star is less than about 2.2 solar masses, the entire
core ignites suddenly in a helium core flash. If the star is more massive than
this, the ignition of the core is more gentle. At the same time, the star
continues to burn hydrogen in a shell around the core.
The star burns helium into carbon in its core for a much shorter time than it
burned hydrogen. Once the helium has all been converted, the inert carbon core
begins to contract and increase in temperature. This ignites a helium burning
shell just above the core, which in turn is surrounded by a hydrogen burning
shell.
What happens next depends on the mass of the star
Stars less than 8 solar masses
The inert carbon core continues to contract but never reaches
temperatures sufficient to initiate carbon burning. However, the
existence of two burning shells leads to a thermally unstable situation
in which hydrogen and helium burning occur out of phase with each
other. This thermal pulsing is characteristic of asymptotic giant branch
stars.
The carbon core continues to contract until it is supported by electron
degeneracy pressure. No further contraction is possible (the core is
now supported by the pressure of electrons, not gas pressure), and the
core has formed a white dwarf. Meanwhile, each thermal pulse causes
the outer layers of the star to expand, resulting in a period of mass
loss. Eventually, the outer layers of the star are ejected completely and
ionised by the white dwarf to form a planetary nebula.
Stars greater than 8 solar masses
The contracting core will reach the temperature for carbon ignition, and
begin to burn to neon. This process of core burning followed by core
contraction and shell burning, is repeated in a series of nuclear reactions
producing successively heavier elements until iron is formed in the core.
Iron cannot be burned to heavier elements as this reaction does not
generate energy – it requires an input of energy to proceed. The star
has therefore finally run out of fuel and collapses under its own gravity.
The mass of the core of the star dictates what happens next. If the core
has a mass less than about 3 times that of our Sun, the collapse of
the core may be halted by the pressure of neutrons (this is an even
more extreme state than the electron pressure that supports white
dwarfs!). In this case, the core becomes a neutron star. The sudden
halt in the contraction of the core produces a shock wave which
propagates back out through the outer layers of the star, blowing it
apart in a core-collapse supernova explosion. If the core has a mass
greater than about 3 solar masses, even neutron pressure is not
sufficient to withstand gravity, and it will collapse further into a stellar
black hole.
The ejected gas expands into the interstellar medium, enriching it with
all the elements synthesised during the star’s lifetime and in the
explosion itself. These supernova remnants are the chemical distribution
centres of the Universe.
An important tool in the study of stellar evolution is the Hertzsprung-Russell
diagram (HR diagram), which plots the absolute magnitudes of stars against their
spectral type (or alternatively, stellar luminosity versus effective temperature). As
a star evolves, it moves to specific regions in the HR diagram, following a characteristic
path that depends on the star’s mass and chemical composition.
See also: black hole.
Stellar classification
In astronomy, stellar classification is the classification of stars based on their spectral
characteristics. Electromagnetic radiation from the star is analyzed by splitting it with a prism or
diffraction grating into a spectrum exhibiting the rainbow of colors interspersed with spectral lines.
Each line indicates a particular chemical element or molecule, with the line strength indicating the
abundance of that element. The strengths of the different spectral lines vary mainly due to the
temperature of the photosphere, although in some cases there are true abundance differences. The
spectral class of a star is a short code primarily summarizing the ionization state, giving an objective
measure of the photosphere's temperature.
Most stars are currently classified under the Morgan–Keenan (MK) system using the letters O, B, A,
F, G, K, and M, a sequence from the hottest (O type) to the coolest (M type). Each letter class is
then subdivided using a numeric digit with 0 being hottest and 9 being coolest (e.g., A8, A9, F0, and
F1 form a sequence from hotter to cooler). The sequence has been expanded with classes for other
stars and star-like objects that do not fit in the classical system, such as class D for white dwarfs and
classes S and C for carbon stars.
In the MK system, a luminosity class is added to the spectral class using Roman numerals. This is
based on the width of certain absorption lines in the star's spectrum, which vary with the density of
the atmosphere and so distinguish giant stars from dwarfs. Luminosity class 0 or Ia+ is used for
hypergiants, class I for supergiants, class II for bright giants, class III for regular giants, class IV for
subgiants, class V for main-sequence stars, class sd (or VI) for subdwarfs, and class D (or VII) for
white dwarfs. The full spectral class for the Sun is then G2V, indicating a main-sequence star with a
surface temperature around 5,800 K.
Conventional colour description[edit]
Main article: Green star (astronomy)
Just-saturated RGB-camera discs
The conventional colour description takes into account only the peak of the stellar spectrum. In
actuality, however, stars radiate in all parts of the spectrum. Because all spectral colours combined
appear white, the actual apparent colours the human eye would observe are far lighter than the
conventional colour descriptions would suggest. This characteristic of 'lightness' indicates that the
simplified assignment of colours within the spectrum can be misleading. Excluding colour-contrast
effects in dim light, in typical viewing conditions there are no green, indigo, or violet stars. Red
dwarfs are a deep shade of orange, and brown dwarfs do not literally appear brown, but
hypothetically would appear dim grey to a nearby observer.
Modern classification[edit]
Main-sequence stars arranged from O to M Harvard classes
The modern classification system is known as the Morgan–Keenan (MK) classification. Each star is
assigned a spectral class from the older Harvard spectral classification and a luminosity class using
Roman numerals as explained below, forming the star's spectral type.
Other modern stellar classification systems, such as the UBV system, are based on color indices—
the measured differences in three or more color magnitudes. Those numbers are given labels such as
"U− V" or "B− V", which represent the colors passed by two standard filters (e.g. Ultraviolet, Blue and
Visual).
Harvard spectral classification[edit]
The Harvard system is a one-dimensional classification scheme by astronomer Annie Jump Cannon,
who re-ordered and simplified the prior alphabetical system by Draper (see #History). Stars are
grouped according to their spectral characteristics by single letters of the alphabet, optionally with
numeric subdivisions. Main-sequence stars vary in surface temperature from approximately 2,000 to
50,000 K, whereas more-evolved stars can have temperatures above 100,000 K. Physically, the
classes indicate the temperature of the star's atmosphere and are normally listed from hottest to
coldest.
A common mnemonic for remembering the order of the spectral type letters, from hottest to coolest,
is "Oh, Be A Fine Guy/Girl: Kiss Me!".[9]
The spectral classes O through M, as well as other more specialized classes discussed later, are
subdivided by Arabic numerals (0–9), where 0 denotes the hottest stars of a given class. For
example, A0 denotes the hottest stars in class A and A9 denotes the coolest ones. Fractional
numbers are allowed; for example, the star Mu Normae is classified as O9.7.[10] The Sun is
classified as G2.[11]
Conventional color descriptions are traditional in astronomy, and represent colors relative to the
mean color of an A class star, which is considered to be white. The apparent color[5] descriptions are
what the observer would see if trying to describe the stars under a dark sky without aid to the eye, or
with binoculars. However, most stars in the sky, except the brightest ones, appear white or bluish
white to the unaided eye because they are too dim for color vision to work. Red supergiants are
cooler and redder than dwarfs of the same spectral type, and stars with particular spectral features
such as carbon stars may be far redder than any black body.[dubious – discuss]
The fact that the Harvard classification of a star indicated its surface or photospheric temperature (or
more precisely, its effective temperature) was not fully understood until after its development, though
by the time the first Hertzsprung–Russell diagram was formulated (by 1914), this was generally
suspected to be true.[12] In the 1920s, the Indian physicist Meghnad Saha derived a theory of
ionization by extending well-known ideas in physical chemistry pertaining to the dissociation of
molecules to the ionization of atoms. First he applied it to the solar chromosphere, then to stellar
spectra.[13]
Harvard astronomer Cecilia Payne then demonstrated that the O-B-A-F-G-K-M spectral sequence is
actually a sequence in temperature.[14] Because the classification sequence predates our
understanding that it is a temperature sequence, the placement of a spectrum into a given subtype,
such as B3 or A7, depends upon (largely subjective) estimates of the strengths of absorption
features in stellar spectra. As a result, these subtypes are not evenly divided into any sort of
mathematically representable intervals.
Yerkes spectral classification[edit]
Montage of false color spectra for main-sequence stars[15]
The Yerkes spectral classification, also called the MKK system from the authors' initials, is a system
of stellar spectral classification introduced in 1943 by William Wilson Morgan, Philip C. Keenan, and
Edith Kellman from Yerkes Observatory.[16] This two-dimensional (temperature and luminosity)
classification scheme is based on spectral lines sensitive to stellar temperature and surface gravity,
which is related to luminosity (whilst the Harvard classification is based on just surface temperature).
Later, in 1953, after some revisions of list of standard stars and classification criteria, the scheme
was named the Morgan–Keenan classification, or MK,[17] and this system remains in use.
Denser stars with higher surface gravity exhibit greater pressure broadening of spectral lines. The
gravity, and hence the pressure, on the surface of a giant star is much lower than for a dwarf star
because the radius of the giant is much greater than a dwarf of similar mass. Therefore, differences
in the spectrum can be interpreted as luminosity effects and a luminosity class can be assigned
purely from examination of the spectrum.
A number of different luminosity classes are distinguished, as listed in the table below.[18]
Luminosity
class
Description
Examples
0 or Ia+
hypergiants or extremely luminous
supergiants
Cygnus OB2#12 – B3-4Ia+[19]
Ia
luminous supergiants
Eta Canis Majoris – B5Ia[20]
Iab
intermediate-size luminous supergiants
Gamma Cygni – F8Iab[21]
Ib
less luminous supergiants
Zeta Persei – B1Ib[22]
II
bright giants
Beta Leporis – G0II[23]
III
normal giants
Arcturus – K0III[24]
IV
subgiants
Gamma Cassiopeiae –
B0.5IVpe[25]
V
main-sequence stars (dwarfs)
Achernar – B6Vep[22]
sd (prefix) or VI
subdwarfs
HD 149382 – sdB5 or B5VI[26]
D (prefix) or VII
white dwarfs[c]
van Maanen 2 – DZ8[27]
Marginal cases are allowed; for example, a star may be either a supergiant or a bright giant, or may
be in between the subgiant and main-sequence classifications. In these cases, two special symbols
are used:
●
●
A slash (/) means that a star is either one class or the other.
A dash (-) means that the star is in between the two classes.
For example, a star classified as A3-4III/IV would be in between spectral types A3 and A4, while
being either a giant star or a subgiant.
Sub-dwarf classes have also been used: VI for sub-dwarfs (stars slightly less luminous than the
main sequence).
Nominal luminosity class VII (and sometimes higher numerals) is now rarely used for white dwarf or
"hot sub-dwarf" classes, since the temperature-letters of the main sequence and giant stars no
longer apply to white dwarfs.
Occasionally, letters a and b are applied to luminosity classes other than supergiants; for example, a
giant star slightly less luminous than typical may be given a luminosity class of IIIb, while a
luminosity class IIIa indicates a star slightly brighter than a typical giant. [28]
A sample of extreme V stars with strong absorption in He II λ4686 spectral lines have been given the
Vz designation. An example star is HD 93129 B.[29]
Spectral peculiarities[edit]
Additional nomenclature, in the form of lower-case letters, can follow the spectral type to indicate
peculiar features of the spectrum.[30]
Code
Spectral peculiarities for stars
:
uncertain spectral value[18]
...
Undescribed spectral peculiarities exist
!
comp
e
Special peculiarity
Composite spectrum [31]
Emission lines present[31]
[e]
"Forbidden" emission lines present
er
"Reversed" center of emission lines weaker than edges
eq
Emission lines with P Cygni profile
f
N III and He II emission[18]
f*
N IV λ4058Å is stronger than the N III λ4634Å, λ4640Å, & λ4642Å lines[32]
f+
Si IV λ4089Å & λ4116Å are emitted, in addition to the N III line[32]
(f)
N III emission, absence or weak absorption of He II
(f+)
[33]
((f))
Displays strong He II absorption accompanied by weak N III emissions [34]
((f*))
[33]
h
WR stars with hydrogen emission lines.[35]
ha
He wk
WR stars with hydrogen seen in both absorption and emission.[35]
Weak Helium lines
k
Spectra with interstellar absorption features
m
Enhanced metal features[31]
n
Broad ("nebulous") absorption due to spinning [31]
nn
Very broad absorption features[18]
neb
A nebula's spectrum mixed in[31]
p
Unspecified peculiarity, peculiar star.[d][31]
pq
Peculiar spectrum, similar to the spectra of novae
q
P Cygni profiles
s
Narrow ("sharp") absorption lines[31]
ss
Very narrow lines
sh
Shell star features[31]
var
Variable spectral feature[31] (sometimes abbreviated to "v")
wl
Weak lines[31] (also "w" & "wk")
Element Abnormally strong spectral lines of the specified element(s)[31]
symbol
For example, 59 Cygni is listed as spectral type B1.5Vnne,[36] indicating a spectrum with the general
classification B1.5V, as well as very broad absorption lines and certain emission lines.
History[edit]
The reason for the odd arrangement of letters in the Harvard classification is historical, having
evolved from the earlier Secchi classes and been progressively modified as understanding
improved.
Secchi classes[edit]
During the 1860s and 1870s, pioneering stellar spectroscopist Angelo Secchi created the Secchi
classes in order to classify observed spectra. By 1866, he had developed three classes of stellar
spectra, shown in the table below.[37][38][39]
In the late 1890s, this classification began to be superseded by the Harvard classification, which is
discussed in the remainder of this article.[40][41][42]
Class number
Secchi class I
Secchi class description
White and blue stars with broad heavy hydrogen lines, such as Vega and
Altair. This includes the modern class A and early class F.
Secchi class I
A subtype of Secchi class I with narrow lines in place of wide bands, such
as Rigel and Bellatrix. In modern terms, this corresponds to early B-type
(Orion subtype) stars
Secchi class II
Yellow stars – hydrogen less strong, but evident metallic lines, such as the
Sun, Arcturus, and Capella. This includes the modern classes G and K as
well as late class F.
Secchi class III
Orange to red stars with complex band spectra, such as Betelgeuse and
Antares.
This corresponds to the modern class M.
Secchi class IV In 1868, he discovered carbon stars, which he put into a distinct group:[43]
Red stars with significant carbon bands and lines, corresponding to modern
classes C and S.
Secchi class V
In 1877, he added a fifth class:[44]
Emission-line stars, such as Gamma Cassiopeiae and Sheliak, which are in
modern class Be. In 1891, Edward Charles Pickering proposed that class V
should correspond to the modern class O (which then included Wolf–Rayet
stars) and stars within planetary nebulae.[45]
The Roman numerals used for Secchi classes should not be confused with the completely unrelated
Roman numerals used for Yerkes luminosity classes and the proposed neutron star classes.
Draper system[edit]
Secchi
Draper
Comment
I
A, B, C, D
II
E, F, G, H, I, K, L
III
M
IV
N
Did not appear in the catalogue
V
O
Included Wolf–Rayet spectra with bright lines
V
P
Planetary nebulae
Hydrogen lines dominant
Q
Other spectra
Classes carried through into the MK system are in bold.
In the 1880s, the astronomer Edward C. Pickering began to make a survey of stellar spectra at the
Harvard College Observatory, using the objective-prism method. A first result of this work was the
Draper Catalogue of Stellar Spectra, published in 1890. Williamina Fleming classified most of the
spectra in this catalogue and was credited with classifying over 10,000 featured stars and
discovering 10 novae and more than 200 variable stars.[48] With the help of the Harvard computers,
especially Williamina Fleming, the first iteration of the Henry Draper catalogue was devised to
replace the Roman-numeral scheme established by Angelo Secchi.[49]
The catalogue used a scheme in which the previously used Secchi classes (I to V) were subdivided
into more specific classes, given letters from A to P. Also, the letter Q was used for stars not fitting
into any other class.[46][47] Fleming worked with Pickering to differentiate 17 different classes based
on the intensity of hydrogen spectral lines, which causes variation in the wavelengths emanated
from stars and results in variation in color appearance. The spectra in class A tended to produce the
strongest hydrogen absorption lines while spectra in class O produced virtually no visible lines. The
lettering system displayed the gradual decrease in hydrogen absorption in the spectral classes when
moving down the alphabet. This classification system was later modified by Annie Jump Cannon and
Antonia Maury to produce the Harvard spectral classification scheme.[48][50]
The old Harvard system (1897)[edit]
In 1897, another astronomer at Harvard, Antonia Maury, placed the Orion subtype of Secchi class I
ahead of the remainder of Secchi class I, thus placing the modern type B ahead of the modern type
A. She was the first to do so, although she did not use lettered spectral types, but rather a series of
twenty-two types numbered from I–XXII.[51][52]
Gro
ups
Summary
I− V
included ‘Orion type’ stars that displayed an increasing strength in
hydrogen absorption lines from group I to group V
VI
acted as an intermediate between the ‘Orion type’ and Secchi type I group
VII−
were Secchi’s type 1 stars, with decreasing strength in hydrogen absorption
XI
lines from groups VII− XI
XIII−
included Secchi type 2 stars with decreasing hydrogen absorption lines and
increasing solar-type metallic lines
XVI
included Secchi type 3 stars with increasing spectral lines
XVII
− XX
XXI
included Secchi type 4 stars
XXII
included Wolf–Rayet stars
Because the 22 Roman numeral groupings didn't account for additional variations in spectra, three
additional divisions were made to further specify differences: Lowercase letters were added to
differentiate relative line appearance in spectra; the lines were defined as[53]
(
a
)
average
width
(
b
)
hazy
(
c
)
sharp
Antonia Maury published her own stellar classification catalogue in 1897 called "Spectra of Bright
Stars Photographed with the 11 inch Draper Telescope as Part of the Henry Draper Memorial",
which included 4,800 photographs and Maury's analyses of 681 bright northern stars. This was the
first instance in which a woman was credited for an observatory publication. [54]
The current Harvard system (1912)[edit]
In 1901, Annie Jump Cannon returned to the lettered types, but dropped all letters except O, B, A, F,
G, K, M, and N used in that order, as well as P for planetary nebulae and Q for some peculiar
spectra. She also used types such as B5A for stars halfway between types B and A, F2G for stars
one fifth of the way from F to G, and so on.[55][56]
Finally, by 1912, Cannon had changed the types B, A, B5A, F2G, etc. to B0, A0, B5, F2, etc.[57][58]
This is essentially the modern form of the Harvard classification system. This system was developed
through the analysis of spectra on photographic plates, which could convert light emanated from
stars into a readable spectra.[59]
Mount Wilson classes[edit]
Proper motion of stars of early type in ± 200,000 years
A luminosity classification known as the Mount Wilson system was used to distinguish between stars
of different luminosities.[60][61][62] This notation system is still sometimes seen on modern
spectra.[63]
Class
Meaning
sd
Subdwarf
d
Dwarf
sg
Subgiant
g
Giant
c
Supergiant
The movement of stars of late type around the apex (left) and antapex (right) in ± 200,000 years
Spectral types[edit]
The stellar classification system is taxonomic, based on type specimens, similar to classification of
species in biology: The categories are defined by one or more standard stars for each category and
sub-category, with an associated description of the distinguishing features. [64]
"Early" and "late" nomenclature[edit]
Stars are often referred to as early or late types. "Early" is a synonym for hotter, while "late" is a
synonym for cooler.
Depending on the context, "early" and "late" may be absolute or relative terms. "Early" as an
absolute term would therefore refer to O or B, and possibly A stars. As a relative reference it relates
to stars hotter than others, such as "early K" being perhaps K0, K1, K2 and K3.
"Late" is used in the same way, with an unqualified use of the term indicating stars with spectral
types such as K and M, but it can also be used for stars that are cool relative to other stars, as in
using "late G" to refer to G7, G8, and G9.
In the relative sense, "early" means a lower Arabic numeral following the class letter, and "late"
means a higher number.
This obscure terminology is a hold-over from a late nineteenth century model of stellar evolution,
which supposed that stars were powered by gravitational contraction via the Kelvin–Helmholtz
mechanism, which is now known to not apply to main-sequence stars. If that were true, then stars
would start their lives as very hot "early-type" stars and then gradually cool down into "late-type"
stars. This mechanism provided ages of the Sun that were much smaller than what is observed in
the geologic record, and was rendered obsolete by the discovery that stars are powered by nuclear
fusion.[65] The terms "early" and "late" were carried over, beyond the demise of the model they were
based on.
Class O[edit]
Main article: O-type star
See also: O-type main-sequence star, Blue giant, and Blue supergiant
The spectrum of an O5V star
O-type stars are very hot and extremely luminous, with most of their radiated output in the ultraviolet
range. These are the rarest of all main-sequence stars. About 1 in 3,000,000 (0.00003%) of the
main-sequence stars in the solar neighborhood are O-type stars.[e][8] Some of the most massive
stars lie within this spectral class. O-type stars frequently have complicated surroundings that make
measurement of their spectra difficult.
O-type spectra formerly were defined by the ratio of the strength of the He II λ4541 relative to that of
He I λ4471, where λ is the radiation wavelength. Spectral type O7 was defined to be the point at
which the two intensities are equal, with the He I line weakening towards earlier types. Type O3 was,
by definition, the point at which said line disappears altogether, although it can be seen very faintly
with modern technology. Due to this, the modern definition uses the ratio of the nitrogen line N IV
λ4058 to N III λλ4634-40-42.[66]
O-type stars have dominant lines of absorption and sometimes emission for He II lines, prominent
ionized (Si IV, O III, N III, and C III) and neutral helium lines, strengthening from O5 to O9, and
prominent hydrogen Balmer lines, although not as strong as in later types. Higher-mass O-type stars
do not retain extensive atmospheres due to the extreme velocity of their stellar wind, which may
reach 2,000 km/s. Because they are so massive, O-type stars have very hot cores and burn through
their hydrogen fuel very quickly, so they are the first stars to leave the main sequence.
When the MKK classification scheme was first described in 1943, the only subtypes of class O used
were O5 to O9.5.[67] The MKK scheme was extended to O9.7 in 1971[68] and O4 in 1978,[69] and
new classification schemes that add types O2, O3, and O3.5 have subsequently been introduced.[70]
Spectral standards:[64]
●
●
O7V – S Monocerotis
O9V – 10 Lacertae
Class B[edit]
See also: B-type main-sequence star, Blue giant, and Blue supergiant
B-class stars in the Jewel Box cluster (Credit: ESO VLT)
B-type stars are very luminous and blue. Their spectra have neutral helium lines, which are most
prominent at the B2 subclass, and moderate hydrogen lines. As O- and B-type stars are so
energetic, they only live for a relatively short time. Thus, due to the low probability of kinematic
interaction during their lifetime, they are unable to stray far from the area in which they formed, apart
from runaway stars.
The transition from class O to class B was originally defined to be the point at which the He II λ4541
disappears. However, with modern equipment, the line is still apparent in the early B-type stars.
Today for main-sequence stars, the B class is instead defined by the intensity of the He I violet
spectrum, with the maximum intensity corresponding to class B2. For supergiants, lines of silicon are
used instead; the Si IV λ4089 and Si III λ4552 lines are indicative of early B. At mid-B, the intensity
of the latter relative to that of Si II λλ4128-30 is the defining characteristic, while for late B, it is the
intensity of Mg II λ4481 relative to that of He I λ4471. [66]
These stars tend to be found in their originating OB associations, which are associated with giant
molecular clouds. The Orion OB1 association occupies a large portion of a spiral arm of the Milky
Way and contains many of the brighter stars of the constellation Orion. About 1 in 800 (0.125%) of
the main-sequence stars in the solar neighborhood are B-type main-sequence stars.[e][8]
Massive yet non-supergiant entities known as "Be stars" are main-sequence stars that notably have,
or had at some time, one or more Balmer lines in emission, with the hydrogen-related
electromagnetic radiation series projected out by the stars being of particular interest. Be stars are
generally thought to feature unusually strong stellar winds, high surface temperatures, and
significant attrition of stellar mass as the objects rotate at a curiously rapid rate.[71] Objects known as
"B(e)" or "B[e]" stars possess distinctive neutral or low ionisation emission lines that are considered
to have 'forbidden mechanisms', undergoing processes not normally allowed under current
understandings of quantum mechanics.
Spectral standards:[64]
●
●
B0V – Upsilon Orionis
B0Ia – Alnilam
●
●
●
●
B2Ia – Chi2 Orionis
B2Ib – 9 Cephei
B3V – Eta Ursae Majoris
B3V – Eta Aurigae
●
●
●
B3Ia – Omicron2 Canis Majoris
B5Ia – Eta Canis Majoris
B8Ia – Rigel
Class A[edit]
See also: A-type main-sequence star
Class A Vega (left) compared to the Sun (right)
A-type stars are among the more common naked eye stars, and are white or bluish-white. They
have strong hydrogen lines, at a maximum by A0, and also lines of ionized metals (Fe II, Mg II, Si II)
at a maximum at A5. The presence of Ca II lines is notably strengthening by this point. About 1 in
160 (0.625%) of the main-sequence stars in the solar neighborhood are A-type stars.[e][8][72]
Spectral standards:[64]
●
●
●
●
●
●
●
A0Van – Gamma Ursae Majoris
A0Va – Vega
A0Ib – Eta Leonis
A0Ia – HD 21389
A1V – Sirius A
A2Ia – Deneb
A3Va – Fomalhaut
Class F[edit]
See also: F-type main-sequence star
Canopus, an F-type supergiant and the second-brightest star in the night sky
F-type stars have strengthening spectral lines H and K of Ca II. Neutral metals (Fe I, Cr I) beginning
to gain on ionized metal lines by late F. Their spectra are characterized by the weaker hydrogen
lines and ionized metals. Their color is white. About 1 in 33 (3.03%) of the main-sequence stars in
the solar neighborhood are F-type stars.[e][8]
Spectral standards:[64]
●
●
●
F0IIIa – Zeta Leonis
F0Ib – Alpha Leporis
F2V – 78 Ursae Majoris
Class G[edit]
"G star" redirects here. For other uses, see G star (disambiguation).
See also: G-type main-sequence star, Yellow supergiant, and Yellow hypergiant
The Sun, a G2 main-sequence star, with dark sunspots
G-type stars, including the Sun,[11] have prominent spectral lines H and K of Ca II, which are most
pronounced at G2. They have even weaker hydrogen lines than F, but along with the ionized metals,
they have neutral metals. There is a prominent spike in the G band of CN molecules. Class G main-
sequence stars make up about 7.5%, nearly one in thirteen, of the main-sequence stars in the solar
neighborhood.[e][8]
Class G contains the "Yellow Evolutionary Void".[73] Supergiant stars often swing between O or B
(blue) and K or M (red). While they do this, they do not stay for long in the unstable yellow
supergiant class.
Spectral standards:[64]
●
●
●
●
●
●
●
●
●
●
●
●
G0V – Beta Canum Venaticorum
G0IV – Eta Boötis
G0Ib – Beta Aquarii
G2V – Sun
G5V – Kappa1 Ceti
G5IV – Mu Herculis
G5Ib – 9 Pegasi
G8V – 61 Ursae Majoris
G8IV – Beta Aquilae
G8IIIa – Kappa Geminorum
G8IIIab – Epsilon Virginis
G8Ib – Epsilon Geminorum
Class K[edit]
See also: K-type main-sequence star
"K-type star" redirects here. For the Korean nuclear fusion project, see KSTAR.
Arcturus, a K1.5 giant compared to the Sun and Antares
K-type stars are orangish stars that are slightly cooler than the Sun. They make up about 12% of the
main-sequence stars in the solar neighborhood.[e][8] There are also giant K-type stars, which range
from hypergiants like RW Cephei, to giants and supergiants, such as Arcturus, whereas orange
dwarfs, like Alpha Centauri B, are main-sequence stars.
They have extremely weak hydrogen lines, if those are present at all, and mostly neutral metals (Mn
I, Fe I, Si I). By late K, molecular bands of titanium oxide become present. Mainstream theories
(those rooted in lower harmful radioactivity and star longevity) would thus suggest such stars have
the optimal chances of heavily evolved life developing on orbiting planets (if such life is directly
analogous to earth's) due to a broad habitable zone yet much lower harmful periods of emission
compared to those with the broadest such zones.[74][75]
Spectral standards:[64]
●
●
●
●
●
●
●
●
K0V – Sigma Draconis
K0III – Pollux
K0III – Epsilon Cygni
K2V – Epsilon Eridani
K2III – Kappa Ophiuchi
K3III – Rho Boötis
K5V – 61 Cygni A
K5III – Gamma Draconis
Class M[edit]
See also: Red dwarf, Red giant, and Red supergiant
Class M stars are by far the most common. About 76% of the main-sequence stars in the solar
neighborhood are class M stars.[e][f][8] However, class M main-sequence stars (red dwarfs) have
such low luminosities that none are bright enough to be seen with the unaided eye, unless under
exceptional conditions. The brightest-known M class main-sequence star is Lacaille 8760, class
M0V, with magnitude 6.7 (the limiting magnitude for typical naked-eye visibility under good
conditions is typically quoted as 6.5), and it is extremely unlikely that any brighter examples will be
found.
Although most class M stars are red dwarfs, most of the largest-ever supergiant stars in the Milky
Way are M stars, such as VV Cephei, Antares, and Betelgeuse, which are also class M.
Furthermore, the larger, hotter brown dwarfs are late class M, usually in the range of M6.5 to M9.5.
The spectrum of a class M star contains lines from oxide molecules (in the visible spectrum,
especially TiO) and all neutral metals, but absorption lines of hydrogen are usually absent. TiO
bands can be strong in class M stars, usually dominating their visible spectrum by about M5.
Vanadium(II) oxide bands become present by late M.
Spectral standards:[64]
●
●
●
●
M0IIIa – Beta Andromedae
M2III – Chi Pegasi
M1-M2Ia-Iab – Betelgeuse
M2Ia – Mu Cephei ("Herschel’s garnet")
Extended spectral types[edit]
A number of new spectral types have been taken into use from newly discovered types of stars. [76]
Hot blue emission star classes[edit]
UGC 5797, an emission-line galaxy where massive bright blue stars are formed[77]
Spectra of some very hot and bluish stars exhibit marked emission lines from carbon or nitrogen, or
sometimes oxygen.
Class W: Wolf–Rayet[edit]
Main article: Wolf–Rayet star
Hubble Space Telescope image of the nebula M1-67 and the Wolf–Rayet star WR 124 in the center
Once included as type O stars, the Wolf–Rayet stars of class W or WR are notable for spectra
lacking hydrogen lines. Instead their spectra are dominated by broad emission lines of highly ionized
helium, nitrogen, carbon, and sometimes oxygen. They are thought to mostly be dying supergiants
with their hydrogen layers blown away by stellar winds, thereby directly exposing their hot helium
shells. Class W is further divided into subclasses according to the relative strength of nitrogen and
carbon emission lines in their spectra (and outer layers). [35]
WR spectra range is listed below:[78][79]
●
WN[35] – spectrum dominated by N III-V and He I-II lines
○ WNE (WN2 to WN5 with some WN6) – hotter or "early"
○ WNL (WN7 to WN9 with some WN6) – cooler or "late"
○
Extended WN classes WN10 and WN11 sometimes used for the Ofpe/WN9
○
stars[35]
h tag used (e.g. WN9h) for WR with hydrogen emission and ha (e.g. WN6ha)
for both hydrogen emission and absorption
●
WN/C – WN stars plus strong C IV lines, intermediate between WN and WC stars[35]
●
WC[35] – spectrum with strong C II-IV lines
○ WCE (WC4 to WC6) – hotter or "early"
○ WCL (WC7 to WC9) – cooler or "late"
WO (WO1 to WO4) – strong O VI lines, extremely rare, extension of the WCE class into
incredibly hot temperatures (up to 200 kK or more)
●
Although the central stars of most planetary nebulae (CSPNe) show O-type spectra,[80] around 10%
are hydrogen-deficient and show WR spectra.[81] These are low-mass stars and to distinguish them
from the massive Wolf–Rayet stars, their spectra are enclosed in square brackets: e.g. [WC]. Most
of these show [WC] spectra, some [WO], and very rarely [WN].
The "Slash" stars[edit]
Main article: Slash star
The slash stars are O-type stars with WN-like lines in their spectra. The name "slash" comes from
their printed spectral type having a slash in it (e.g. "Of/WNL"[66]).
There is a secondary group found with this spectra, a cooler, "intermediate" group designated
"Ofpe/WN9".[66] These stars have also been referred to as WN10 or WN11, but that has become
less popular with the realisation of the evolutionary difference from other Wolf–Rayet stars. Recent
discoveries of even rarer stars have extended the range of slash stars as far as O2-3.5If*/WN5-7,
which are even hotter than the original "slash" stars.[82]
The magnetic O stars[edit]
They are O stars with strong magnetic fields. Designation is Of?p. [66]
Cool red and brown dwarf classes[edit]
Main articles: Brown dwarf and Red dwarf
The new spectral types L, T, and Y were created to classify infrared spectra of cool stars. This
includes both red dwarfs and brown dwarfs that are very faint in the visible spectrum.[83]
Brown dwarfs, stars that do not undergo hydrogen fusion, cool as they age and so progress to later
spectral types. Brown dwarfs start their lives with M-type spectra and will cool through the L, T, and
Y spectral classes, faster the less massive they are; the highest-mass brown dwarfs cannot have
cooled to Y or even T dwarfs within the age of the universe. Because this leads to an unresolvable
overlap between spectral types' effective temperature and luminosity for some masses and ages of
different L-T-Y types, no distinct temperature or luminosity values can be given.[7]
Class L[edit]
Artist's impression of an L-dwarf
Class L dwarfs get their designation because they are cooler than M stars and L is the remaining
letter alphabetically closest to M. Some of these objects have masses large enough to support
hydrogen fusion and are therefore stars, but most are of substellar mass and are therefore brown
dwarfs. They are a very dark red in color and brightest in infrared. Their atmosphere is cool enough
to allow metal hydrides and alkali metals to be prominent in their spectra.[84][85][86]
Due to low surface gravity in giant stars, TiO- and VO-bearing condensates never form. Thus, L-type
stars larger than dwarfs can never form in an isolated environment. However, it may be possible for
these L-type supergiants to form through stellar collisions, an example of which is V838 Monocerotis
while in the height of its luminous red nova eruption.
Class T: methane dwarfs[edit]
Artist's impression of a T-dwarf
Class T dwarfs are cool brown dwarfs with surface temperatures between approximately 550 and
1,300 K (277 and 1,027 °C; 530 and 1,880 °F). Their emission peaks in the infrared. Methane is
prominent in their spectra.[84][85]
Study of the number of proplyds (protoplanetary disks, clumps of gas in nebulae from which stars
and planetary systems are formed) indicates that the number of stars in the galaxy should be several
orders of magnitude higher than what was previously conjectured. It is theorized that these proplyds
are in a race with each other. The first one to form will become a protostar, which are very violent
objects and will disrupt other proplyds in the vicinity, stripping them of their gas. The victim proplyds
will then probably go on to become main-sequence stars or brown dwarfs of the L and T classes,
which are quite invisible to us.[citation needed][87]
Class Y[edit]
See also: Sub-brown dwarf and Substellar object
Artist's impression of a Y-dwarf
Brown dwarfs of spectral class Y are cooler than those of spectral class T and have qualitatively
different spectra from them. A total of 17 objects have been placed in class Y as of August 2013. [88]
Although such dwarfs have been modelled[89] and detected within forty light-years by the Wide-field
Infrared Survey Explorer (WISE)[76][90][91][92][93] there is no well-defined spectral sequence yet and
no prototypes. Nevertheless, several objects have been proposed as spectral classes Y0, Y1, and
Y2.[94]
The spectra of these prospective Y objects display absorption around 1.55 micrometers.[95] Delorme
et al. have suggested that this feature is due to absorption from ammonia, and that this should be
taken as the indicative feature for the T-Y transition.[95][96] In fact, this ammonia-absorption feature
is the main criterion that has been adopted to define this class. [94] However, this feature is difficult to
distinguish from absorption by water and methane,[95] and other authors have stated that the
assignment of class Y0 is premature.[97]
The latest brown dwarf proposed for the Y spectral type, WISE 1828+2650, is a > Y2 dwarf with an
effective temperature originally estimated around 300 K, the temperature of the human
body.[90][91][98] Parallax measurements have, however, since shown that its luminosity is
inconsistent with it being colder than ~400 K. The coolest Y dwarf currently known is WISE
0855−0714 with an approximate temperature of 250 K.[99]
The mass range for Y dwarfs is 9–25 Jupiter masses, but young objects might reach below one
Jupiter mass, which means that Y class objects straddle the 13 Jupiter mass deuterium-fusion limit
that marks the current IAU division between brown dwarfs and planets.[94]
Peculiar brown dwarfs[edit]
Symbols used for peculiar brown dwarfs
pec This suffix stands for "peculiar" (e.g. L2pec).[100]
sd
This prefix (e.g. sdL0) stands for subdwarf and
indicates a low metallicity and blue color [101]
β
Objects with the beta (β) suffix (e.g. L4β) have an
intermediate surface gravity.[102]
γ
Objects with the gamma (γ) suffix (e.g. L5γ) have a
low surface gravity.[102]
red
The red suffix (e.g. L0red) indicates objects without
signs of youth, but high dust content.[103]
blue The blue suffix (e.g. L3blue) indicates unusual blue
near-infrared colors for L-dwarfs without obvious
low metallicity.[104]
Young brown dwarfs have low surface gravities because they have larger radii and lower masses
compared to the field stars of similar spectral type. These sources are marked by a letter beta ( β) for
intermediate surface gravity and gamma ( γ) for low surface gravity. Indication for low surface gravity
are weak CaH, KI and NaI lines, as well as strong VO line.[102] Alpha (α) stands for normal surface
gravity and is usually dropped. Sometimes an extremely low surface gravity is denoted by a delta
(δ).[104] The suffix "pec" stands for peculiar. The peculiar suffix is still used for other features that are
unusual and summarizes different properties, indicative of low surface gravity, subdwarfs and
unresolved binaries.[105] The prefix sd stands for subdwarf and only includes cool subdwarfs. This
prefix indicates a low metallicity and kinematic properties that are more similar to halo stars than to
disk stars.[101] Subdwarfs appear bluer than disk objects.[106] The red suffix describes objects with
red color, but an older age. This is not interpreted as low surface gravity, but as a high dust
content.[103][104] The blue suffix describes objects with blue near-infrared colors that cannot be
explained with low metallicity. Some are explained as L+T binaries, others are not binaries, such as
2MASS J11263991−5003550 and are explained with thin and/or large-grained clouds.[104]
Late giant carbon-star classes[edit]
Carbon-stars are stars whose spectra indicate production of carbon – a byproduct of triple-alpha
helium fusion. With increased carbon abundance, and some parallel s-process heavy element
production, the spectra of these stars become increasingly deviant from the usual late spectral
classes G, K, and M. Equivalent classes for carbon-rich stars are S and C.
The giants among those stars are presumed to produce this carbon themselves, but some stars in
this class are double stars, whose odd atmosphere is suspected of having been transferred from a
companion that is now a white dwarf, when the companion was a carbon-star.
Class C: carbon stars[edit]
Main article: Carbon star
Image of the carbon star R Sculptoris and its striking spiral structure
Originally classified as R and N stars, these are also known as carbon stars. These are red giants,
near the end of their lives, in which there is an excess of carbon in the atmosphere. The old R and N
classes ran parallel to the normal classification system from roughly mid-G to late M. These have
more recently been remapped into a unified carbon classifier C with N0 starting at roughly C6.
Another subset of cool carbon stars are the C–J-type stars, which are characterized by the strong
presence of molecules of 13CN in addition to those of 12CN.[107] A few main-sequence carbon stars
are known, but the overwhelming majority of known carbon stars are giants or supergiants. There
are several subclasses:
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C-R – Formerly its own class (R) representing the carbon star equivalent of late G- to
early K-type stars.
C-N – Formerly its own class representing the carbon star equivalent of late K- to M-type
stars.
C-J – A subtype of cool C stars with a high content of 13C.
C-H – Population II analogues of the C-R stars.
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C-Hd – Hydrogen-deficient carbon stars, similar to late G supergiants with CH and C2
bands added.
Class S[edit]
Main article: S-type star
Class S stars form a continuum between class M stars and carbon stars. Those most similar to class
M stars have strong ZrO absorption bands analogous to the TiO bands of class M stars, whereas
those most similar to carbon stars have strong sodium D lines and weak C2 bands.[108] Class S
stars have excess amounts of zirconium and other elements produced by the s-process, and have
more similar carbon and oxygen abundances than class M or carbon stars. Like carbon stars, nearly
all known class S stars are asymptotic-giant-branch stars.
The spectral type is formed by the letter S and a number between zero and ten. This number
corresponds to the temperature of the star and approximately follows the temperature scale used for
class M giants. The most common types are S3 to S5. The non-standard designation S10 has only
been used for the star Chi Cygni when at an extreme minimum.
The basic classification is usually followed by an abundance indication, following one of several
schemes: S2,5; S2/5; S2 Zr4 Ti2; or S2*5. A number following a comma is a scale between 1 and 9
based on the ratio of ZrO and TiO. A number following a slash is a more-recent but less-common
scheme designed to represent the ratio of carbon to oxygen on a scale of 1 to 10, where a 0 would
be an MS star. Intensities of zirconium and titanium may be indicated explicitly. Also occasionally
seen is a number following an asterisk, which represents the strength of the ZrO bands on a scale
from 1 to 5.
Classes MS and SC: Intermediate carbon-related classes[edit]
In between the M and S classes, border cases are named MS stars. In a similar way, border cases
between the S and C-N classes are named SC or CS. The sequence M → MS → S → SC → C-N is
hypothesized to be a sequence of increased carbon abundance with age for carbon stars in the
asymptotic giant branch.
White dwarf classifications[edit]
Main article: White dwarf spectroscopy
The class D (for Degenerate) is the modern classification used for white dwarfs—low-mass stars that
are no longer undergoing nuclear fusion and have shrunk to planetary size, slowly cooling down.
Class D is further divided into spectral types DA, DB, DC, DO, DQ, DX, and DZ. The letters are not
related to the letters used in the classification of other stars, but instead indicate the composition of
the white dwarf's visible outer layer or atmosphere.
The white dwarf types are as follows:[109][110]
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DA – a hydrogen-rich atmosphere or outer layer, indicated by strong Balmer hydrogen
spectral lines.
DB – a helium-rich atmosphere, indicated by neutral helium, He I, spectral lines.
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DO – a helium-rich atmosphere, indicated by ionized helium, He II, spectral lines.
DQ – a carbon-rich atmosphere, indicated by atomic or molecular carbon lines.
DZ – a metal-rich atmosphere, indicated by metal spectral lines (a merger of the obsolete
white dwarf spectral types, DG, DK, and DM).
DC – no strong spectral lines indicating one of the above categories.
DX – spectral lines are insufficiently clear to classify into one of the above categories.
The type is followed by a number giving the white dwarf's surface temperature. This number is a
rounded form of 50400/Teff, where Teff is the effective surface temperature, measured in kelvins.
Originally, this number was rounded to one of the digits 1 through 9, but more recently fractional
values have started to be used, as well as values below 1 and above 9.[109][111]
Two or more of the type letters may be used to indicate a white dwarf that displays more than one of
the spectral features above.[109]
Extended white dwarf spectral types[edit]
Sirius A and B (a white dwarf of type DA2) resolved by Hubble
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DAB – a hydrogen- and helium-rich white dwarf displaying neutral helium lines
DAO – a hydrogen- and helium-rich white dwarf displaying ionized helium lines
DAZ – a hydrogen-rich metallic white dwarf
DBZ – a helium-rich metallic white dwarf
A different set of spectral peculiarity symbols are used for white dwarfs than for other types of
stars:[109]
Code
P
Spectral peculiarities for stars
Magnetic white dwarf with detectable
polarization
E
Emission lines present
H
Magnetic white dwarf without detectable
polarization
V
Variable
PEC
Spectral peculiarities exist
Non-stellar spectral types: Classes P and Q[edit]
Finally, the classes P and Q are left over from the system developed by Cannon for the Henry
Draper Catalogue. They are occasionally used for certain non-stellar objects: Type P objects are
stars within planetary nebulae (typically young white dwarfs or hydrogen-poor M giants); type Q
objects are novae.[citation needed]
Stellar remnants[edit]
Main articles: Neutron star, Black hole, and Exotic star
Stellar remnants are objects associated with the death of stars. Included in the category are white
dwarfs, and as can be seen from the radically different classification scheme for class D, non-stellar
objects are difficult to fit into the MK system.
The Hertzsprung–Russell diagram, which the MK system is based on, is observational in nature so
these remnants cannot easily be plotted on the diagram, or cannot be placed at all. Old neutron
stars are relatively small and cold, and would fall on the far right side of the diagram. Planetary
nebulae are dynamic and tend to quickly fade in brightness as the progenitor star transitions to the
white dwarf branch. If shown, a planetary nebula would be plotted to the right of the diagram's upper
right quadrant. A black hole emits no visible light of its own, and therefore would not appear on the
diagram.[112]
A classification system for neutron stars using Roman numerals has been proposed: type I for less
massive neutron stars with low cooling rates, type II for more massive neutron stars with higher
cooling rates, and a proposed type III for more massive neutron stars (possible exotic star
candidates) with higher cooling rates.[113] The more massive a neutron star is, the higher neutrino
flux it carries. These neutrinos carry away so much heat energy that after only a few years the
temperature of an isolated neutron star falls from the order of billions to only around a million Kelvin.
This proposed neutron star classification system is not to be confused with the earlier Secchi
spectral classes and the Yerkes luminosity classes.
Replaced spectral classes[edit]
Several spectral types, all previously used for non-standard stars in the mid-20th century, have been
replaced during revisions of the stellar classification system. They may still be found in old editions
of star catalogs: R and N have been subsumed into the new C class as C-R and C-N.