the evolution of high-mass stars
In stars with masses greater than about 1.1
M, the central
temperature is high enough for the
CNO cycle to become dominant.
The strong temperature dependence of the CNO cycle means that the
energy generation is much more concentrated at the centre of the
star. The resulting steep temperature gradient is unstable to
and hence such stars have convective cores. Convection has
the effect of mixing the material in the core, bringing fresh hydrogen
into the centre and spreading the newly-produced helium throughout
the core. This keeps the chemical composition of the core uniform,
which means that when the nuclear reactions have used up all of the
hydrogen at the centre, there is no hydrogen left anywhere in the
convectively mixed region and energy production ceases throughout the
core. Just outside the core, hydrogen is still available for burning, but
the temperature is too low for fusion to occur and hydrogen burning
ceases altogether when the core burns out.
The structure of a high-mass star
(M > 1.1 M)
on the sub-giant and red-giant branches of the HR diagram.
The abrupt cessation of hydrogen burning means that
the star now has no nuclear energy source and it is forced into a slow
overall contraction. Part of the gravitational potential energy released
is used to balance the radiation from its surface, but about half of
it goes into heating up the core as the central density and pressure
are increased by the contraction. Eventually, the core is hot enough
that hydrogen can start burning in a thin shell just outside the core
boundary. The star is now similar in structure to a hydrogen-shell-burning
star of lower mass, as shown in
figure 19, but it has a much
thinner shell (because the dependence of energy production on temperature
is more severe). Overall contraction stops, and the energy from the
contraction of the core is now fed into an expansion of the envelope,
just as for lower-mass stars. The
main difference in the evolutionary tracks of higher-mass stars in the
HR diagram is that the
phase of overall contraction causes a hook to the left
before progress to the giant region is resumed once the shell
has ignited, as shown
in figure 20.
The evolution of high-mass stars depicted schematically on an HR
diagram. The track for a solar mass star is also shown for comparison.
The three dots on each track represent, from left to right, the
main-sequence, the onset of core contraction following hydrogen
exhaustion, and the beginning of helium burning, respectively.
The exhausted core now consists of almost pure helium. Although there
is no nuclear energy source, the centre of the star is initially
hotter than the edge of the core and heat flows down this temperature
gradient, cooling the centre and heating up the outside. This
continues until the core becomes isothermal. The isothermal core grows
slowly in mass as the shell gradually burns its way outwards into
fresh fuel, leaving helium 'ash' behind. The core does not need to
contract because it is not losing energy - its temperature is
maintained by the overlying hydrogen-burning shell. This continues
until the core reaches the
Schönberg-Chandrasekhar limit of about 10% of the mass of the
star, at which point an isothermal core cannot support itself against
gravity. This instability only occurs in stars with masses between
about 2 M and
6 M. For
lower mass stars,
the core becomes degenerate before the
Schönberg-Chandrasekhar limit is reached; for
higher mass stars, the central temperature becomes hot enough for
helium fusion to occur before the
Schönberg-Chandrasekhar limit is reached.
In stars in which the Schönberg-Chandrasekhar limit
is reached, the core begins to contract rapidly. The energy released
goes into a rapid expansion of the whole star, and hence a rapid
transition to the giant branch. Because of this rapid transition, very
few stars are observed during this phase, which accounts for the
Hertsprung gap seen in
HR diagrams of galactic clusters.
As lower-mass stars
become red giants less rapidly, this agrees with the
absence of a Hertzsprung gap in the
HR diagrams of globular clusters.
The collapse of the helium core, whether rapid or not, raises the
central temperature to the point where helium ignites via the
triple-alpha process. If the
core is not degenerate (i.e.
M > 2 M),
helium ignites gently and there is no
helium flash. The effect of
helium ignition, whether violent or quiet, is to move a star off
the giant branch towards higher surface temperatures (i.e. to the left
of the HR diagram). There are now two nuclear energy sources, helium-burning
in the core and hydrogen-burning in a shell, and the evolution is much
more complicated than in the core-hydrogen-burning phase.
Nuclear evolution beyond core helium-burning depends on whether or not
the carbon-oxygen core ever becomes hot enough for further fusion
reactions to occur and whether the core becomes degenerate. For stars
with initial masses of
M > 8 M,
the core is non-degenerate and carbon can ignite quietly, burning
first to oxygen and neon. At lower initial masses, a carbon flash
occurs, as the core is degenerate. Further reactions are possible and a series
of burning episodes builds up successive shells of more and more processed
material. Elements produced in these shells include magnesium,
silicon and sulphur. For stellar masses greater than about
11 M, burning can proceed
as far as iron and other elements of comparable nuclear mass, principally
chromium, manganese, cobalt and nickel (the so-called iron-peak
elements). At this point, because iron
is the most stable
element, with the highest binding energy per nucleon, to produce
elements heavier than iron it is necessary to add energy.
The star has thus exhausted all its possible nuclear fuels and it has
an onion skin structure, with successive shells containing the
ashes of the various burning stages, as shown in
The onion-ring structure of a red supergiant (a pre-supernova star).
Note that this diagram is not to scale - the outer hydrogen burning
shell has a radius of order
whereas the star has a radius of order
The evolution of a massive star undergoing these different phases of
core and shell burning beyond helium is very complex.
As the core and shell energy sources vary in relative strength, the star
makes a number of excursions to and fro across the HR diagram. In high-mass
stars, these rightward (core exhaustion) and leftward (core ignition)
excursions, between the red and blue (supergiant) branches respectively,
occur with only a slight systematic increase in luminosity and hence
the evolutionary tracks of high-mass stars occur virtually horizontally
in the HR diagram. In very high-mass stars, the nuclear evolution in the
central regions of the star occurs so quickly that the outer layers
have no time to respond to the successive rounds of core exhaustion and
core ignition, and there is only a relatively steady drift to the
right on the HR diagram before the star arrives at the pre-supernova
state, as shown in figure 20.
It should be noted that the details of
this stage of stellar evolution remain uncertain. In
addition, the time taken to make
these latter excursions in the HR diagram are very small compared to the
duration of the earlier phases. So we do not expect to observe
all the complications of an individual evolutionary track to show up in
a star cluster HR diagram.
With no nuclear energy generation, the iron core becomes degenerate.
As lighter elements continue to burn in shells above it, the iron core
grows in mass until it exceeds the Chandrasekhar limit of
around 1.4 M - the
maximum possible mass of a white dwarf, above which electron degeneracy
pressure is insufficient to prevent gravitational collapse. The core
then begins to collapse. The core's iron nuclei decompose into those
of helium, which then fragment into protons and
neutrons, and the protons then combine with the electrons to form
more neutrons, all at the expense of the star's
gravitational potential energy. In this way, the collapsed core becomes
a neutron star, where it is the degeneracy pressure exerted by
neutrons which prevents continued gravitational collapse.
Meanwhile, the outer layers of the star are still collapsing: they hit
the hard surface of the newly formed neutron star and bounce off,
creating a shockwave which blows off the outer layers of the star in a
Type II supernova explosion, as depicted in the
animation of figure 22.
An animation showing the final stages of a Type II supernova
What remains after a Type II supernova explosion? The expelled
envelope of the star becomes visible as a supernova remnant,
as shown in figure 23, at the centre
of which lies the core of the star. If the core of the star has a
mass below approximately 3 M,
it is a neutron star, those that exhibit pulsed radiation being
known as pulsars.
The crab nebula - the remnant of a supernova which exploded about
900 years ago.
But if the mass of the core exceeds
approximately 3 M,
there is nothing to prevent the core from collapsing to a state of
zero radius and infinite density (within the laws of physics as they
are presently understood). Such an object is known as a
black hole. Isolated black holes cannot be directly
observed, as the escape velocity from the surface exceeds the
speed of light. Evidence that they exist has
been confirmed via observations of binary stars, where the motion
of the visible star is measured in order to determine the mass of
the compact object lying at the heart of the accretion disc in
figure 24. If the mass of the compact
object is measured to be
M > 3 M, the
existence of a black hole is inferred.
An artists impression of an X-ray binary star, some of which are
believed to harbour black holes.
©Vik Dhillon, 27th November 2012