INFO:
1 Parsec = 3.26 light-years
Blackbody Radiation:
Black-body radiation is the thermal electromagnetic radiation within or surrounding a body in
thermodynamic equilibrium with its environment, emitted by a black body (an idealized opaque, nonreflective body). It has a specific spectrum of wavelengths, inversely related to intensity that depend
only on the body's temperature, which is assumed for the sake of calculations and theory to be
uniform and constant.[1][2][3][4]
The thermal radiation spontaneously emitted by many ordinary objects can be approximated as
black-body radiation. A perfectly insulated enclosure that is in thermal equilibrium internally contains
black-body radiation and will emit it through a hole made in its wall, provided the hole is small
enough to have a negligible effect upon the equilibrium.
In a dark room, a black body at room temperature appears black because most of the energy it
radiates is in the infrared spectrum and cannot be perceived by the human eye. Since the human
eye cannot perceive light waves below the visible frequency, a black body at the lowest just faintly
visible temperature subjectively appears grey, even though its objective physical spectrum peak is in
the infrared range.[5] The human eye essentially does not perceive color at low light levels. When
the object becomes a little hotter, it appears dull red. As its temperature increases further it becomes
bright red, orange, yellow, white, and ultimately blue-white.
Although planets and stars are neither in thermal equilibrium with their surroundings nor perfect
black bodies, black-body radiation is used as a first approximation for the energy they emit.[6] Black
holes are near-perfect black bodies, in the sense that they absorb all the radiation that falls on them.
It has been proposed that they emit black-body radiation (called Hawking radiation), with a
temperature that depends on the mass of the black hole.[7]
The term black body was introduced by Gustav Kirchhoff in 1860.[8] Black-body radiation is also
called thermal radiation, cavity radiation, complete radiation or temperature radiation.
Protostars:
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-mainsequence star, which contracts to later become a main-sequence star at the onset of hydrogen
fusion producing helium.
T Tauri Variables:
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.
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).
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 carbon-rich.[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]
RR Lyrae Variables:
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.
Finding absolute magnitude: M = -2.81 * log(period) -1.43
Carbon Stars:
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.
White Dwarfs:
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.
Planetary Nebulae:
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.
Neutron Stars:
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 (6.2 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 2.14 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 to 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 Xrays 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 so-called 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.
Dwarf Novae:
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]
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.
Recurrent Novae:
A nova (plural novae or novas) is a transient astronomical event that causes the sudden appearance
of a bright, apparently "new" star, that slowly fades over several weeks or many months. Causes of
the dramatic appearance of a nova vary, depending on the circumstances of the two progenitor
stars. All observed novae involve a white dwarf in a close binary system. The main sub-classes of
novae are classical novae, recurrent novae (RNe), and dwarf novae. They are all considered to be
cataclysmic variable stars.
Classical nova eruptions are the most common type. They are likely created in a close binary star
system consisting of a white dwarf and either a main sequence, subgiant, or red giant star. When the
orbital period falls in the range of several days to one day, the white dwarf is close enough to its
companion star to start drawing accreted matter onto the surface of the white dwarf, which creates a
dense but shallow atmosphere. This atmosphere, mostly consisting of hydrogen, is thermally heated
by the hot white dwarf and eventually reaches a critical temperature causing ignition of rapid
runaway fusion.
The sudden increase in energy expels the atmosphere into interstellar space creating the envelope
seen as visible light during the nova event and in past centuries mistaken as a "new" star. A few
novae produce short-lived nova remnants, lasting for perhaps several centuries. Recurrent nova
processes are the same as the classical nova, except that the fusion ignition may be repetitive
because the companion star can again feed the dense atmosphere of the white dwarf.
Novae most often occur in the sky along the path of the Milky Way, especially near the observed
galactic centre in Sagittarius; however, they can appear anywhere in the sky. They occur far more
frequently than galactic supernovae, averaging about ten per year. Most are found telescopically,
perhaps only one every 12–18 months reaching naked-eye visibility. Novae reaching first or second
magnitude occur only several times per century. The last bright nova was V1369 Centauri reaching
3.3 magnitude on 14 December 2013.[1]
Type Ia Supernovae:
A type Ia 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. Somewhat confusingly, this critical mass is often referred to as the Chandrasekhar mass,
despite being 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 its 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]
Magnetic Cataclysmic Variables:
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.
Red Giants:
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 Mclass giant at 88 light-years' distance.
OBJECTS:
HOPS 383:
HOPS 383 is a Class 0 protostar. It is the first class-0 protostar discovered to have had an
outburst,[1] and as of 2020, the youngest protostar known to have had an outburst.[1] The protostar
was discovered by the Herschel Orion Protostar Survey (HOPS) team.[2]
HH 24-26:
HH 24-26 is a molecular cloud and star-forming region containing the Herbig-Haro objects HH 24,
HH 25 and HH 26. This region contains the highest concentration of astrophysical jets known
anywhere in the sky.[2] The molecular cloud is located about 1400 light-years away in the L1630
dark cloud, which is part of the Orion B molecular cloud in the constellation of Orion.[3][4]
The region contains multiple protostars (two class 0 and one class I) and four more evolved IRAS
sources. The three protostars are driving the Herbig-Haro objects in this region.[3]
V1331 Cyg:
V1331 Cygni (also known as V1331 Cyg) is a young star in the constellation Cygnus. V1331 Cyg is
located in the dark nebula LDN 981.[2]
V1331 Cygni is most noted for having an arc-like reflection nebula surrounding it. This circumstellar
disc is a great birthplace for young stars, which form in the cloud.[1]
HBC 672:
Astronomers using a previously captured Hubble imagery spotted a remarkable image of a young
star's unseen, planet-forming disk casting a huge shadow across a more distant cloud in a starforming region. The star is called HBC 672, and the shadow feature was nicknamed the "Bat
Shadow" because it resembles a pair of wings. The nickname turned out to be unexpectedly
appropriate, because now those "wings" appear to be flapping!
Orion Nebula:
The Orion Nebula (also known as Messier 42, M42, or NGC 1976) is a diffuse nebula situated in
the Milky Way, being south of Orion's Belt in the constellation of Orion.[b] It is one of the brightest
nebulae and is visible to the naked eye in the night sky. It is 1,344 ± 20 light-years (412.1 ± 6.1 pc)
away[3][6] and is the closest region of massive star formation to Earth. The M42 nebula is estimated
to be 24 light-years across. It has a mass of about 2,000 times that of the Sun. Older texts frequently
refer to the Orion Nebula as the Great Nebula in Orion or the Great Orion Nebula.[7]
The Orion Nebula is one of the most scrutinized and photographed objects in the night sky and is
among the most intensely studied celestial features.[8] The nebula has revealed much about the
process of how stars and planetary systems are formed from collapsing clouds of gas and dust.
Astronomers have directly observed protoplanetary disks and brown dwarfs within the nebula,
intense and turbulent motions of the gas, and the photo-ionizing effects of massive nearby stars in
the nebula.
Alpha Tauri:
Aldebaran /ælˈdɛbərən/, designated α Tauri (Latinized to Alpha Tauri), is a giant star measured to
be about 65 light-years from the Sun in the zodiac constellation Taurus. It is the brightest star in
Taurus and generally the fourteenth-brightest star in the night sky, though it varies slowly in
brightness between magnitude 0.75 and 0.95. Aldebaran is believed to host a planet several times
the mass of Jupiter, named Aldebaran b.
Aldebaran is cooler than the Sun with a surface temperature of 3,900 K, but its radius is about 44
times the Sun's, so it is over 400 times as luminous. It spins slowly and takes 520 days to complete
a rotation.
The planetary exploration probe Pioneer 10 is heading in the general direction of the star and should
make its closest approach in about two million years.
RR Lyrae:
RR Lyrae is a variable star in the Lyra constellation, figuring in its west near to Cygnus.[10] As the
brightest star in its class,[11] it became the eponym for the RR Lyrae variable class of stars[3] and it
has been extensively studied by astronomers.[7] RR Lyrae variables serve as important standard
candles that are used to measure astronomical distances. The period of pulsation of an RR Lyrae
variable depends on its mass, luminosity and temperature, while the difference between the
measured luminosity and the actual luminosity allows its distance to be determined via the inversesquare law. Hence, understanding the period-luminosity relation for a local set of such stars allows
the distance of more distant stars of this type to be determined.[12]
Omicron Ceti:
Mira (/ˈmaɪrə/), designation Omicron Ceti (ο Ceti, abbreviated Omicron Cet, ο Cet), is a red-giant
star estimated to be 200–400 light-years from the Sun in the constellation Cetus.
ο Ceti is a binary stellar system, consisting of a variable red giant (Mira A) along with a white dwarf
companion (Mira B). Mira A is a pulsating variable star and was the first non-supernova variable star
discovered, with the possible exception of Algol. It is the prototype of the Mira variables.
ESO 577-24:
ESO 577-24, also known as IRAS F13378-1937 or 2MASS J13404134-1952553, resides
approximately 1,400 light-years away from Earth.
It was discovered as part of the National Geographic Society-Palomar Observatory Sky Survey
in the 1950s, and was recorded in the Abell Catalogue of Planetary Nebulae in 1966.
The dazzling nebula is the remains of a dead giant star that has thrown off its outer layers,
leaving behind a small, intensely hot dwarf star.
This diminished remnant will gradually cool and fade, living out its days as the mere ghost of a
once-vast red giant star.
IC 4593:
●
Chandra has found a bubble of ultra-hot gas at the center of a planetary nebula.
●
Planetary nebulas are formed when Sun-like stars run out of fuel, shedding their outer
layers while the star's core shrinks.
●
This image contains X-rays from Chandra (purple) and optical light data from Hubble (pink
and green).
●
IC 4593 is at a distance of about 7,800 light years from Earth, which is the farthest
planetary nebula detected by Chandra.
U Antliae:
U Antliae (U Ant) is a variable star in the constellation Antlia. It is a carbon star surrounded by two
thin shells of dust.
U Antliae is an extremely red C-type carbon star. These cool stars on the asymptotic giant branch
are further reddened by strong mass loss and dust that forms around the star. U Antliae is calculated
to have an effective surface temperature of 2,800 K, although the light that reaches us has an
appearance more like that from a black body with a temperature of 2,300 K surrounded by dust at a
temperature of 72 K.[11] It emits most of its radiation in the infrared and although it is only about 500
times brighter than the sun at visual wavelengths,[9] its bolometric luminosity is 8,000 times higher
than the Sun's.[11]
The visual band light curve of T Antliae, from AAVSO data[12]
U Antliae is an irregular variable star with an apparent magnitude that varies between 5.27 and 6.04.
Approximately 900 light years from Earth, it is surrounded by two shells of dust, thought to have
been ejected 14,000 and 10,000 years ago.[13] The exact origin and structure of the shells is
uncertain, possibly due to enhanced mass loss during thermal pulses, possibly due to interaction of
the stellar wind with interstellar material.[11]
LP 40-365:
LP 40-365 is a low-mass white dwarf star in the constellation Ursa Minor. It travels at high speed
through the Milky Way and has a very unusual elemental composition, lacking hydrogen, helium or
carbon. It may have been produced in a subluminous Type Iax supernova that failed to destroy its
host star totally.[2] [4][5] The "LP" name is derived from the Luyten-Palomar proper motion catalogue
in which it appeared in the 1960s.[6] Another catalog name for this star is "GD 492". [7] The star was
cataloged as a Giclas object with the designation "GD 492" being assigned by Henry Giclas in
1970.[8]
ASASSN-16oh:
●
The fastest growing white dwarf may have been found using data from NASA's Chandra
X-ray Observatory and other telescopes.
●
When stars like our Sun run out of nuclear fuel, they shrink and dim as they become a
"white dwarf" star.
●
Some white dwarf stars have companion stars from which they can siphon off material due
to their strong surface gravity.
●
A new study of ASASSN-16oh in the Small Magellanic Cloud offers a new explanation for
"supersoft" X-rays detected from this white dwarf.
V Sagittae:
V Sagittae or V Sge is a cataclysmic variable binary star system in the constellation Sagitta that is
expected to go nova and briefly become the most luminous point of light in the Milky Way and the
brightest stars in our sky around the year 2083.[5][6] The system is composed of a main sequence
star of about 3.3 solar masses and a white dwarf of about 0.9 solar masses; the fact that the white
dwarf is less massive than its companion is highly unusual,[5] and V Sge is the only super soft X-ray
source nonmagnetic cataclysmic variable found so far.
V Sge has brightened by a factor of 10 over the last century, and based on research reported in
2020, it is expected to continue to brighten and briefly become the brightest star in the night sky
sometime around 2083, plus or minus about 11 years.[7] Over the final few months and days the pair
will coalesce and go nova, eventually becoming a red giant star.
The stars orbit each other about every 0.514 days, and eclipse each other each orbit. The pair is
viewed in the late stages of their spiral in. Their orbital period currently decreases by 4.73 * 10−10
days per cycle, a rate which will accelerate.
Material from the larger star is accreting onto the white dwarf at an exponentially increasing rate,
generating a huge stellar wind. The doubling time for the accretion rate, and hence for the system
luminosity, is about 89 years.[3]
AR Scorpii:
AR Scorpii (AR Sco) is a binary pulsar that consists of a white dwarf and a red dwarf.[2] It is located
close to the ecliptic plane in the constellation Scorpius. Parallax measurements made by Gaia put
the system at a distance of about 380 light-years (120 parsecs).[3]
AR Scorpii is the first, and as of 2021, the only[4] "white dwarf-pulsar" to be discovered.[5] Its unusual
nature was first noticed by amateur astronomers.[6] The 3.56-hour period in AR Scorpii's light curve
caused it to be misclassified as a Delta Scuti variable, but in 2016, this period was found to be the
binary orbital period. In addition, the system shows very strong optical, ultraviolet, and radio
pulsations originating from the red dwarf with a period of just 1.97 minutes, which is a beat period
from the orbital rotation and the white dwarf spin.[2] These pulsations occur when a relativistic beam
from the white dwarf sweeps across the red dwarf, which then reprocesses the beam into the
observed electromagnetic energy. Although the white dwarf shows evidence of accretion in the past,
at present it is not accreting significantly, and the system is powered by the spin-down of the white
dwarf.[5][4] The white dwarf's rotation will slow down on a timescale of 107 years.[4] It has a radius of
about 7×103 km,[4] about the same size as Earth.
SDSS 1035+0551:
Couldn’t find anything
Tycho’s SNR:
SN 1572 (Tycho's Supernova, Tycho's Nova), or B Cassiopeiae (B Cas), was a supernova of Type
Ia in the constellation Cassiopeia, one of eight supernovae visible to the naked eye in historical
records. It appeared in early November 1572 and was independently discovered by many
individuals.
Its supernova remnant has been observed optically but was first detected at radio wavelengths; it is
often known as 3C 10, a radio-source designation, although increasingly as Tycho's supernova
remnant.