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PLASMA
Biography
of a Star
Nuclear fusion is a
virtually inexhaustible
source of energy, and
for decades now scientists
exploiting it. A process
that continues to present
difficulties in laboratories
on Earth has been running
smoothly in stars like
our own Sun for billions
of years. But how do the
stars work? How are
they born? How do they
die? ACHIM WEISS
at the MAX PLANCK
INSTITUTE
PHYSICS
FOR
ASTRO-
in Garching
tracks the life cycle of the
cosmic plasma spheres –
not with a telescope,
but by using computer
model calculations.
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or a glimpse of the world’s largest laboratory, you need only
look into the clear night sky far
from the bright city lights. And if,
at the same time, you take a deep
breath of fresh country air, you will
be supplying your body with the
very substances that are produced
in this laboratory. Elements such as
nitrogen, oxygen and carbon originate in nurseries that have sparkled
on the terrestrial firmament since
time immemorial: sometimes brighter, sometimes less bright; sometimes
white, sometimes in shades of yellow, blue or red.
The stars have always fascinated
man. As recently as the 1850s, however, researchers were still speculating over the nature of these flickering lights. “We do not know what
the stars are, and never will,” one
professor is reported to have answered when asked by a young physics student whether there might not
perhaps be some way of learning
more about the universe than merely
the position, distance and brightness
of the Sun, moon and stars. The student’s name was Karl Friedrich Zöllner, and he was by no means satisfied with his professor’s answer.
Undeterred, he continued his studies
and became one of the first astrophysicists – a profession that he
played a part in shaping.
Achim Weiss shares the same profession, and works, appropriately, at
the Max Planck Institute for Astrophysics. He has a surprisingly uncomplicated answer to Zöllner’s
question: “Stars are simple plasma
spheres that are subject to their own
gravity.” A plasma is a gas consisting of ions, electrons and neutral
particles; over 99 percent of the vis-
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ible matter in the universe is in this
state. For its part, gravitation is the
dominating force in space, acting
upon all objects that are substantially larger than molecules. Little
else is needed in the way of parts to
build a star. Ingredients such as
magnetic fields, vibration or electrical phenomena are rarely significant – either in nature or in the
computer in Garching on which
Weiss models stars.
JUMP-START FROM
AN EXPLODED SUN
In space, the birth of a star begins
with a giant gas cloud. The mass of
this cloud must be so great that
gravity prevails against the internal
pressure and the turbulence that
would drive the filigree structure
apart. For its birth to proceed, the
star presumably needs a little gentle
help from outside, such as the pressure wave of a nearby supernova,
that is, an exploded sun (see the box
“Furious Finale”).
At some point, the cloud breaks up
into smaller lumps, each of which
collapses. Shackled by gravitation,
the particles within such a fragment
bunch up. “If this were to continue
indefinitely, the star’s birth would
end in a black hole,” says Achim
Weiss. How does the inside of the
emerging gas sphere withstand the
growing gravitational pressure?
What stops the stellar embryo from
breaking up?
The compressive work of gravity
generates heat and pressure. The
heat causes the electrons to separate
from the cores of their atoms – a
plasma is produced. And the pressure enables the gas to build up a
“counter-force” against the gravita-
P HOTO : NASA, ESA AND THE H UBBLE H ERITAGE
(STS C I/AURA) – ESA/H UBBLE C OLLABORATION
have been working on
tion: at any given distance from the
sphere’s center, the pressure is exactly equal to the weight of the gas
masses lying above it. The star has
become a stable structure. Or as an
astrophysicist would put it: it is in a
state of hydrostatic equilibrium.
Such a state can be reproduced by
a simple experiment: carefully press
in a bicycle pump, then block off the
outlet with your finger. Since air in
the pump is no longer able to flow
out, pressure builds up in the tube
and prevents the piston from moving. If the right amount of pressure is
applied to the piston, it remains stationary in the tube of the pump and a
form of equilibrium is produced.
“What happens next in the star’s
life depends entirely on its mass,”
Plasma laboratory on the firmament: Dozens of young
stars shine in the NGC 3603 nebula. Today, astrophysicists
replay their births and biographies on computers.
says Achim Weiss. The mass is therefore the decisive parameter in the
model calculations. In a perfectly
normal, average star like our own
Sun (mass: 1.989 x 1030 kg), an event
with far-reaching consequences occurs after its birth, which lasts a few
hundred thousand years. In the cen-
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ter, the gas – primarily hydrogen –
heats up to a temperature of over ten
million degrees Celsius. At this astronomically high temperature, a fusion reactor ignites, and nucleosynthesis begins: four hydrogen nuclei
(protons) combine to form a nucleus
of helium-4.
Only now has the cosmic gas
sphere become a full member of the
star family. The reason is that stars
have another property that differen22
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a deuterium nucleus, consisting of
one proton and one neutron, to form
a helium-3 nucleus. However, lightweights such as these among the
stars never reach the stage of steady
hydrogen burning. (The term “burning” is used for historical reasons
and is usual in astrophysics; it actually refers to “fusion” and is unrelated to chemical combustion.)
These “black sheep” of the star
family are called brown dwarfs. Their
lives are fairly unspectacular: owing
to their low core temperature, the
gas pressure is not sufficient to keep
By the turn of the 20th century, astrophysicists had already gathered extensive data.
For example, they had determined the masses of many stars; they knew their apparent
and absolute (actual, independent of distance) brightnesses, their colors and surface
temperatures, and their spectral types. In 1913, American Henry N. Russell (1877 to
1957) had the idea to check whether the different characteristics were in some way connected. Some years earlier, the Dane Ejnar Hertzsprung (1873 to 1967) had been wondering much the same thing. Hertzsprung had examined many stars closely on paper and
compared their luminosity and spectral types.
Russell drew a simple system of coordinates with x- and y-axes, familiar to schoolchildren from geometry lessons. On this diagram, he entered the stars according to their
spectral type or temperature (x-axis) and absolute brightness or luminosity (y-axis).
Contrary to what might have been expected, the resulting distribution was far from random. Instead, the majority of stars were located on or close to a line running diagonally
from the top left corner to the bottom right corner: the main sequence.
Some stars, however, were located at the bottom left of the diagram. They are extremely hot and belong to the spectral types O, B, A and F. At the same time, their absolute
brightness is low – that is, they shine weakly. Something that radiates hot but appears
only as a weak light must have a relatively small surface area. These miniature balls of
gas are thus located on the dwarf sequence. Finally, there is a second large group outside the main sequence. Stars in this group are all located in the upper area of the Hertzsprung-Russell diagram with very high absolute brightness values. Some of them have
relatively cool surfaces. Nevertheless, they are brighter than the stars in the main sequence in the spectral classes K or M. They must therefore have very large radii. They
are located in the giant sequence.
The color-magnitude diagram (CMD) is equivalent in practice to the Hertzsprung-Russell
diagram (HRD). The spectral type is replaced by the “color index,” which describes the
brightness of a star in two different spectral ranges and is a measure for the color temperature of the star’s surface. The rule is that the higher the color index, the redder the
star’s appearance. Since the members of a cluster are all located the same distance from
the Earth, their apparent brightness in the sky reflects the actual differences in their
luminosity. The apparent brightness is therefore often substituted for the absolute
brightness in color-magnitude diagrams.
MATTER RAINING FROM
THE PARENT CLOUD
As the brown dwarfs shrink and cool
down, however, the properties of the
gases composed of free electrons
change: they degenerate, as physicists say. This state has a peculiar
feature: the temperature becomes decoupled from the pressure and density, and the star is able to cool down
without the pressure dropping. The
star remains stabilized, and therefore
does not vanish as a small black
hole; instead, it becomes progressively colder and darker.
But back to stars of normal weight.
A few million years after birth, the
young star checks the deluge of matter from its parent cloud by means
of increasingly intense radiation and
a rising wind of charged particles
that it spits off its surface into space.
With these mechanisms, the star
avoids a further increase in mass
and reaches the nuclear fusion
phase. At this point, it enters the
main sequence in the HertzsprungRussell diagram (see the box “Stellar
Class Society”).
A star might be expected to respect its place in this society forever,
according to its initial mass. But this
is by no means the case. The population density in the HertzsprungRussell diagram (HRD) reflects the
relative frequency with which individual star types occur at a particular point in time. If, however, the
data from the same stars were to be
entered in an HRD every couple
hundred thousand years, and the
measurements repeated over a period of several billion years, we would
notice movement: in the resulting
Temperature
25,000 10,000
6,000
3,000
-10
-5
104
Supergiants
102
0
Giants
Main-sequence stars
+5
10
Luminosity
tiates them crucially from planets:
they shine, because they derive energy from nucleosynthesis. The fusion reactor also ensures that the gas
remains hot and delivers sufficient
pressure to maintain the hydrostatic
equilibrium.
Some stars, however, do not possess sufficient substance at birth. If
their mass is less than 75 times that
of the planet Jupiter, or in other
words less than 8 percent of the mass
of the Sun, fusion reactions may still
occur on a limited scale within them;
a proton, for example, may fuse with
STELLAR CLASS SOCIETY
Absolute brightness
Gas, dust and radiation are the scenic background in
the stars’ nursery – as in the Great Orion nebula, which
can be observed even with a pair of binoculars.
the gas spheres in equilibrium in the
long term. Ultimately, gravity gains
the upper hand. The brown dwarfs
shrink and convert their gravitational energy into heat. Incidentally, this
process, known as the Kelvin-Helmholtz contraction, was discussed by
astronomers as one of the possible
sources of stars’ energy, before they
solved the riddle in the 20th century
with the aid of nuclear fusion.
I LLUSTRATION : C HRISTOPH S CHNEIDER
P HOTO : NASA, ESA, T. M EGEATH (U NIVERSITY
OF T OLEDO ) AND
M. R OBBERTO (STS C I)
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10-2
+10
Dwarf stars
10-4
+15
O
B
A
F
G
K
M
Spectral classes
The Hertzsprung-Russell diagram (HRD) clearly shows the four most important
families between which the stars may migrate in the course of their development.
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At this point, the Sun draws its energy from two sources. Whereas in
the shell, the hydrogen is fusing to
form helium, the triple-alpha process
is taking place in the core: a carbon
nucleus is created from each set of
three helium nuclei (alpha particles).
This takes place in a roundabout
way, however. The fusion of two helium nuclei first produces an unstable beryllium nucleus with a half-life
of only 10-16 seconds.
Only when, during its extremely
brief existence, this helium nucleus
collides with another helium nucleus
is stable carbon produced. The capture of further helium nuclei may
also cause oxygen and neon nuclei
to form. In order to ignite the helium, the core contracts, as already
1
H
Triple-alpha process
4
He
1
H
8
Be
Electron
(e-)
Gamma
radiation
Proton-proton
chain reaction
12
C
Positron
(e+)
4
He
Neutrino
1
H
1
H
4
Hee
1
H
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1
H
3
He
-
e
(e+)
4
He
1
1
H
H
I LLUSTRATIONS : C HRISTOPH S CHNEIDER
Fusion reactions take place on a large scale in
the interior of the stars. One of the most frequent
of these is the proton-proton chain reaction
(right): two hydrogen nuclei (protons, orange)
first fuse, causing one of the two protons to be
transformed into a neutron (blue) and creating
a deuterium nucleus. In the next step, a further
proton attaches itself to the deuterium nucleus
and forms a helium-3 nucleus. Finally, two of
these helium-3 nuclei fuse to form a stable helium-4 nucleus. During the reactions, energy is
released in the form of gamma rays. The triplealpha process (above) first produces beryllium-8
from two helium-4 nuclei (alpha particles); this
then fuses immediately with a further helium-4
nucleus to form a stable carbon-12 nucleus.
Weiss solves this time problem by
calculating a further model for a
point later on in the Sun’s life, for
example a million years from now.
“Approximately 10,000 individual
models are needed in order to describe the entire life of a star,” says
Weiss. The time interval between
these models must not be too great,
however, particularly at an advanced
stage in the star’s life: at the giant
stage, events follow in quick succession – once the helium in the core
has transformed completely into carbon and oxygen. The core is then
surrounded by two shells: in the inner shell, helium burns to form carbon; in the outer shell, hydrogen
burns to form helium.
In the space of a few tens of thousands of years, a star goes through a
wild phase. First, the carbon/oxygen
core contracts, while at the same
time, the envelope expands. This
process does not take place evenly,
however, but rather in bursts of
greater or lesser regularity during
which the star inflates, once again
increasing strongly in size and luminosity. During this process, the two
outer shells do not burn simultaneously, but alternately.
CONVECTION THOROUGHLY
MIXES THE GAS
And an astonishing process takes
place within the star: “The complicated interplay of forces creates the
conditions for the nucleosynthesis of
heavy elements,” explains Achim
Weiss, “and violent convection flows
are generated within the star.” These
flows use particles to transport energy, and thoroughly mix the gas. The
heat given off by a radiator is transported in the same way: hot air rises,
while cool air falls. You need only
hold your hand above a hot radiator
to experience this phenomenon for
yourself.
The resulting “eddies” in the star
cause a certain amount of hydrogen
from the outer layer to reach the he-
OF
NUCLEAR FUSION
TAKES A DETOUR
mentioned, becoming hotter in the
process. At the same time, however,
the outer shell greatly expands, causing the surface temperature to fall
from values of some 6,000 degrees
to around 3,000 degrees Celsius. The
Sun has increased its radius a hundredfold, and shines with a reddish
light up to 5,000 times as brightly as
it does at present: it has become a
red giant. Accordingly, it migrates in
the Hertzsprung-Russell diagram to
the giant sequence.
“Recording such a biography requires numerical programs that describe the star as an ideal gas sphere,”
says Achim Weiss. In principle, the
task is to divide the star mathematically into “onion skins,” and to determine the chemical composition,
physical structure (mass, temperature, density, energy flow) and type
of nuclear reaction for each of them.
In order for a star to be analyzed for
a particular point in time, Weiss and
his colleagues typically require a
thousand layers. The result is a snapshot of the stellar glass sphere: a
model of a star.
In the second step, Weiss then calculates the changes that take place
in this model, for example as a result
of the nuclear fusion processes, over
a given time. He then generates the
next, slightly older model. In this
way, the researcher tracks the development of a star in the computer.
In order to test the calculations in
practice, some kind of initial model
is first required. For this purpose,
Weiss uses the measurable state parameters of an actual, undeveloped
star as approximate values – so its
mass, luminosity and radius. He then
sets these state parameters to zero
for the center and begins to calculate
in stages from the inside out. “Only
once we have found a solution in
this way for the initial model do we
begin the actual calculations,” says
the astrophysicist.
What is the subsequent fate of a
star with the mass of our Sun? Achim
K.J. B ORKOWSKI (U NIVERSITY
and heats up so much that, ultimately, the helium ignites.
AND
heart of the Sun until it consists entirely of helium, something that will
happen in around six billion years’
time. Since the Sun is already four
and a half billion years old, it will
have had a fairly stable life of ten
billion years by that point.
When hydrogen burning at the
Sun’s center ceases, the star has a
problem. It loses energy, but tries to
maintain the hydrostatic equilibrium.
Fusion in the interior no longer delivers energy. The Sun uses a trick to
compensate for this deficit: the core
begins to contract, and converts
gravitational energy into heat. In the
process, it heats up, becoming so hot
that the layers outside the burnt-out
core reach a sufficiently high temperature to maintain the hydrogen
fusion. Calculations show that this
burning of the shell eats its way progressively outward over time. And
something is also happening on the
inside: the core contracts further still
P HOTO : J.P. H ARRINGTON
time-lapse movie, some stars would
enter the main sequence and remain
in it for a long time, only to leave it
very quickly toward the giant sequence, finally “crashing” into the
dwarfs. In other words, stars are by
no means static plasma spheres –
they develop. “I am interested in
these differences in stars’ biographies for my calculations,” says Max
Planck researcher Weiss.
Let us consider a star of the same
type as our Sun. Nuclear fusion functions smoothly only when the external conditions such as pressure, density and temperature are right, and
sufficient fuel is also available. At
this point, the Sun has consumed
about half of the hydrogen at its core
by nuclear fusion; around 70 percent
of its mass lies within half the solar
radius of 350,000 kilometers. Over
time, the hydrogen reserves are completely exhausted, and increasing
quantities of helium collect at the
M ARYLAND ),
AND
NASA/ESA
FOCUS
Astronomers refer to this planetary cloud as the
“Cat’s Eye” nebula. It bears witness to the slow
death of a star with the mass of our own Sun.
lium that is burning in the shell beneath it. There, the protons are able
to react with the carbon, resulting in
neutrons being released. The neutrons are captured by the iron particles that were present in the star in
small quantities from the beginning,
resulting in the formation of neutron-rich iron isotopes.
If too many neutrons accumulate,
radioactive beta decay occurs, which
in turn creates stable cobalt nuclei.
The neutrons are thus captured progressively by the atomic nuclei,
which then become progressively
heavier. This “s-process” (s for slow)
produces all elements up to and including lead. According to Achim
Weiss, “one day, the Sun will produce barium and other rare earths
such as lanthanum.”
At any rate, the star’s death is
now imminent. In the final phase, it
loses several tenths of its mass within the space of a few tens of thousands of years, at the end of which
99 percent of its mass is accounted
for by its carbon/oxygen core and
only half a percent each by the thin
hydrogen envelope and the helium
shell. The carbon/oxygen core is
effectively blasted clear in much
the same way that the desert wind
blasts a stone free of sand. The material that is carried off forms an
expanding envelope surrounding
the star; it is lit by the star, and it
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P HOTO : NASA/JPL-C ALTECH /O. K RAUSE (S TEWARD O BSERVATORY )
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Despite its name, a supernova is not a new star; it is an explosion
caused by a sun that has already been in existence for a long
time. Astrophysicists differentiate between several different types
of supernova according to the development of their brightness
over time and their spectral properties. The mechanisms of the
explosions also differ, but are not yet understood in detail. At the
Max Planck Institute for Astrophysics in Garching, the group
headed by Director Wolfgang Hillebrandt, Hans-Thomas Janka and
Friedrich Röpke is seeking explanations for the background to
such cosmic catastrophes. Broadly speaking, the scenarios can be
described as follows:
In the case of Type Ia supernovae, a pair of stars orbit each other
closely: a white dwarf, the corpse of an old star, greedily draws
matter from its partner in a binary system. Ingestion of this “power nutrition” brings the white dwarf back to life. If it is overfed, it
reaches critical mass (the Chandrasekhar limit). At this point, it
becomes unstable and begins to contract. This, in turn, releases
gravitational energy, which causes the star to heat up. As a result,
carbon and oxygen in its core ignite, and silicon and nickel are
produced in nuclear burning reactions. Finally, a burning wave in
the form of a detonation or deflagration front passes through the
gas sphere: the star “explodes.”
1
H fuses in the core to form He
Non-burning envelope
assumes the most diverse shapes,
such as rings, spheres or asymmetrical structures. In the “hard core,”
the fusion processes ultimately
grind to a complete halt.
The star’s meager remains have a
temperature of a few tens of thousands of degrees, and are now only
as large as the Earth. The star now
appears in the Hertzsprung-Russell
diagram as a white dwarf: at first
still hot and bright, but in the absence of nuclear fusion, cooling
down and becoming dark, first
quickly, then more and more slowly
– just like the brown dwarfs. When
the computer has churned out the
state parameters for such a white
dwarf – endless columns of figures
for values such as the density, radius,
mass and temperature – Weiss’ work
is normally over; a white dwarf is
the final stage of a star of low or
medium mass.
Death of a star in the sky: A supernova exploded in Cassiopeia
in 1680. The figure shows the expanding gas envelope.
neutrons. This causes the pressure within the core to drop; the core
then collapses within a fraction of a second to form an object that
is 10,000 times as dense: a neutron star. The matter in the center
of the neutron star presents great resistance to further compression. The stellar matter, which continues to fall from further outside, slams into this hard neutron star at ultrasonic velocity. Before
long, the inevitable happens: a shock wave runs from the inside
outward, tearing the gas envelope with it. The star bursts, suddenly
shining billions of times more brightly than before.
The other types of supernova, of which Type II is the most common, are the result of the explosion of a single star of at least
eight solar masses. Once this star has consumed its main supply of
hydrogen and helium at the end of its life, carbon – the ash of
the helium fire – ignites in its core. The temperatures rise to a
billion degrees Celsius. Neutrinos are produced in large numbers.
Finally, over a period of just a few years, the star produces elements of increasing atomic weight: neon, oxygen, silicon, and finally iron. Iron is the last of these, since iron atoms cannot undergo further fusion. The reactor is extinguished.
In the matter that is projected outward, atomic nuclear reactions
produce large quantities of radioactive material (primarily nickel),
as well as isotopes of cobalt and titanium. Extreme conditions also
prevail in the supernova explosion; under these conditions, heavy
elements such as gold, lead and uranium are produced from atomic nuclei of the iron group as a result of successive capture of alpha particles (helium nuclei) and free neutrons and protons. Owing to the extremely rapid attachment of nucleons to existing
atomic nuclei, these forms of nucleosynthesis are termed the rprocess (r for rapid) in the case of neutron entrapment, and rpprocess in that of proton entrapment.
Death of a star in the computer: In a Type Ia supernova, a thermonuclear flame ignites in the white dwarf, tearing it apart.
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A STROPHYSICS
FOR
A CHIM W EISS – MPI
BASED ON WORK BY
C HRISTOPH S CHNEIDER
C HRISTOPH S CHNEIDER ,
In space, the explosion debris of supernovae form luminous, in
some cases bizarre gas clouds, enriched with heavy elements.
Over time, they become mixed with the interstellar matter, from
which new stars may again be born, and the cycle of the elements
begins once more.
TOP :
The death of a star with over 30 solar masses leaves behind it an
even more extreme object. The mass of the stellar core is so great
that its collapse can no longer be prevented: the burned-out neutron sphere cannot be stabilized, and it collapses under its own
gravity. The gravitation of such a structure is so immense that
even light is incapable of leaving it: a black hole has emerged.
I LLUSTRATION ,
I LLUSTRATION ,
I MAGE : MPI
FOR
A STROPHYSICS - F RIEDRICH R ÖPKE
Deep inside a supernova is the neutron star, a compact object
with a diameter of just 20 to 30 kilometers and a mass one and
a half times that of our Sun. Since the angular momentum is
conserved when the rotating stellar core collapses, the neutron
star rotates extremely quickly. Particles are continually emitted
from its surface and accelerated in its strong magnetic field.
In the process, they emit “synchrotron radiation” in two cones.
If this cosmic lighthouse beam crosses the Earth, the star flashes
at an interval ranging from milliseconds to seconds – and astronomers observe a pulsar.
BOTTOM :
By this time, the core of the star possesses a mass close to the
Chandrasekhar limit, and its density has risen to 10,000 tons per
cubic centimeter. Electrons are squashed into the protons and form
AN ENERGY CRISIS
AMONG THE HEAVYWEIGHTS
Life for the heavyweights is faster
and more dramatic: while a star such
as the Sun remains on the main sequence of the Hertzsprung-Russell
diagram for 10 billion years, a star
with ten times its mass stays there
for only 20 million years. It is much
more wasteful of its fuel reserves,
and ultimately fuses elements in its
core up to and including iron. Should
it experience an energy crisis, it
bursts. At the Max Planck Institute
for Astrophysics, a dedicated research group is studying the simulation of supernovae of this kind (see
the box “Furious Finale”).
What connection exists between a
white dwarf and the star from which
it developed? This is one of the
problems that Achim Weiss is studying with the aid of his models. For
this purpose, the researcher obtains
from catalogs the data of suns belonging to a cluster. Clusters are collections of several hundreds or thousands of suns that were born almost
simultaneously
many
millions of years ago.
Since they were not all endowed with the same mass at
birth, their lives have taken
different paths. Their ages
can be determined from the
“population density” at
various points in the Hertzsprung-Russell diagram.
Let us assume that a cluster is 500 million years old,
and that Weiss finds within it
a white dwarf with a cooling
age of 100 million years. The
cooling age is the time that
has elapsed since the star
developed into a white
dwarf. In this example,
this means that the star
had previously lived
normally for 400 million years. “The problem
to be solved now, says
Weiss, is: What star takes
400 million years to develop
into a white dwarf?” In this case,
it could be a star with approximately three solar masses. Using his
models, the researcher examines this
“initial-final mass relationship” and
obtains results that are sometimes
confusing.
All stars with the same initial mass
would normally be assumed to have
the same final mass as well. As an
example, however, the final masses
of the white dwarfs in the Beehive
Cluster differ by a factor of two. “I
have no idea why this is the case,”
says Achim Weiss. The form of energy transport within the gas spheres
and the mass loss from the surfaces
are evidently decisive factors: “Models with greater mass, which have
large convective cores, deliver clearer results.” Achim Weiss intends to
continue the search for an answer to
this question. By no means do we already know everything in astrophysics – even if we do now have a pretty
good idea of “what the stars are.”
HELMUT HORNUNG
He fuses to form C
H fuses to form He
Non-burning envelope
2
3
C fuses to form Na, Ne and Mg
He fuses to form C
H fuses to form He
Non-burning envelope
4
Degenerate iron core
S and Si fuse to form Fe
O fuses to form S and Si
Ne fuses to form O and Mg
C fuses to form Na, Ne and Mg
He fuses to form C
H fuses to form He
Non-burning envelope
Stellar element cuisine:
From simple hydrogen fusion (1) in the
core of a star, the process passes through
the various stages of shell burning (2, 3),
ending in the creation of heavy elements
up to and including iron (4).
6
4
Luminosity log
(logarithmic)
(L/L???)
10
Furious Finale
2
0
-2
4.6
4.4
4.2
4.0
3.8
3.6
T????
Temperaturelog(logarithmic)
10
Modeled paths of three stars with one solar mass (red) and with three (green) and
ten (blue) solar masses in the HertzsprungRussell diagram. The diagram shows the
development from the main sequence to
the giant stage.
2/2008 MAXPL
A NCK
R
ESEARCH
27
3.4