Stellar Structure and
Evolution
The Life of the
Stars
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From observations of
stars and from comparison between observations and theoretical
models of stars (huge computer programs written by astronomers in an
attempt to understand the physics of the stars), it is now well
established that the life of a star is governed almost exclusively
by its birth mass. Massive stars burn up their fuel much faster than
less massive ones, and the enormous energy production at the centre
of the more massive stars makes their inner structure somewhat
different from the structure of the less massive ones, as we will
come back to in a little while.
Since the heaviest stars
are about 1000 times heavier than the lightest, we expect to observe
large differences in stellar properties over the entire mass range.
Furthermore, the nuclear burning of hydrogen into helium - the very
process that makes most stars shine - takes place near the stellar
cores, and the process changes the stellar properties over time (see
below). We therefore observe differences between stars of different
age, even if their mass is about the same. By the way, nuclear
burning is not an actual "burning", but just a term astronomers uses
for the fusion process of, for example, forming one helium nucleus
out of four hydrogen nuclei. This process is described in more
detail below.
It is now known that our Universe is
approximately 14 billion years old, and that the oldest stars are
nearly as old as this. At the same time, new stars are still being
born, which means that we can observe stars in all stages of life,
as well as we can observe stars with different masses. The result is
that we can learn about both the internal structure and the
evolution of stars, by studying many stars of different mass and
age.
In the following, we will describe the life cycle of a
star, but in order to limit the length of our description, we will
focus on stars which have masses similar to that of the
Sun.
Sun-like Stars
The Sun, as
other stars, is fuelled by nuclear reactions in the central regions,
where the temperature, pressure and density are so high that nuclear
fusion of hydrogen into helium can take place. In the Sun, this is
happening inside the inner 25% in radius, corresponding to the
central 1.5% in volume. However, since stars are much denser in the
core compared to the outer regions, this region actually contains
about 30% of the total mass of the Sun.
You can read
more about the conditions in these central regions of a star like
the Sun here.
The Sun will spend
about 90% of its life in this relatively quite phase, slowly
converting hydrogen into helium and shining steadily. Although the
nuclear reactions change the chemical composition of the central
regions, as the amount of helium gradually increases while the
hydrogen content decreases as the star gets older, the star will not
change its luminosity and size very much in this phase of life. This
stage of evolution is called the "main sequence" and all stars spend
most of their lifetime in this phase.
In total, the Sun will
spend about 10 billion years (of which 5 have already passed) on the
main sequence, while a star 10 times as massive will spend only 10
million years in this phase. A star half as massive as the Sun will
spend about 140 billion years on the main sequence. The huge
differences between these time-scales are due to the rate of nuclear
burning - the more massive stars are much more luminous and use up
their fuel quicker, underlining that the evolution of stars is
indeed governed by the mass.
We will now look into
what happens in the central region of the star as time passes by and
the star evolves away from the main sequence and into the so-called
post-main sequence phases that follow life on the main
sequence.
As the star gets older, the amount of hydrogen
available for nuclear reactions naturally gets smaller (it is
converted into helium). But the core also gets hotter which keeps
the energy production up to level. Once hydrogen is completely
depleted in the centre, the temperature in the surrounding parts is
so high that fusion takes place in a shell around the centre. This
marks the end of the main-sequence phase and brings about rapid
changes for the star.
Energy is now only being
produced in a thin, hot shell around the core, which consists only
of helium. But because of the high temperature, the energy
production actually increases and so does the luminosity. The
surface temperature, on the other hand, is decreased slightly and
the diameter of the star grows rapidly: the star becomes a so-called
giant. In this phase, the star consists of a very dense core of
approximately half the mass of the star, and a huge, but very thin,
outer envelope.
Since there is no energy production in the
centre, this part of the star contracts and heats up. Ultimately,
the temperature in the core gets so high that another nuclear fusion
reaction, where helium nuclei are transformed into carbon, sets in.
For a star like the Sun, this phase of helium burning will only last
for about 100-200 million years.
When all the helium is used
up, the star will swell even more, before finally expelling the
outer layers completely. This exposes the carbon core, which is now
very small (about the size of the Earth), very dense (one cubic
centimeter weighs 1000 kg) and very hot: a white dwarf is
born.
The white dwarf does not produce energy but shines
because it is hot. This means that white dwarfs can be observed from
Earth, at least in sufficiently large telescopes, despite the fact
that the white dwarf does not produce any energy. Over time, it will
radiate its energy away and become cooler and cooler, and less and
less luminous. Finally, it will end up in the stellar graveyard, as
a dark, but still massive object, which can no longer be seen even
with the largest telescopes.
The diagram above,
showing luminosity versus temperature for stars, is called an
Hertzsprung-Russell diagram and will be described in more detail
later in the text.
Energy Transport inside a
Star
This was a very brief description of the life
of a star like the Sun. However, a very important point for
understanding both the internal structure as well as the evolution
of stars is to understand the way stars transport energy from the
hot central regions, where it is created, to the cooler surface,
from where it is radiated into space.
Except in very dense
objects such as white dwarfs, the energy can be transported up
through the star in two well-known ways: by radiation or by
convection. In the dense objects, energy can also be transported by
conduction, but this can be neglected in "normal" stars.
Energy
transportation by both radiation and convection are well-known
everyday phenomena. The Earth receives energy from the Sun in the
form of radiation, and the atmosphere is transparent to radiation -
we feel a difference between being in the sunlight and being in the
shade. Similarly, convection can be seen as the hot air rising over
a heater, or over the asphalt on a hot day.
For stars, the
preferred way of energy transportation from the interior to the
surface depends on many factors, such as temperature, pressure and
the detailed chemical composition of the stellar matter. In the
outer layers of the Sun, energy is primarily transported by
convection, while radiation is the dominant form of energy transport
in the interior regions.
Stars with a mass only slightly
larger than that of the Sun (~10% larger) have convective cores and
only a shallow outer convective layer. These stars have convective
cores because the energy production is so large, that all the energy
simply cannot be transported by radiation alone. Convection - hot
gas rising up in higher, slightly cooler layers and cold gas sinking
- must therefore set in.
For stars even more massive, only
the core is convective, while stars less massive than the Sun have a
very deep outer convection zone. Again, the mass is the dominant
factor in a star's life.
These differences are very important
for the stellar structure and evolution, but especially energy
transport by convection is, in fact, very poorly understood - even
in the convection phenomenon we observe here on Earth.
For
the theoretical modeling of stars, convection is one of the most
difficult processes to describe. In fact, although we understand
stars pretty well, there are still many aspects of the physics of
stars that are unknown or poorly understood. One of the reasons of
this is that we are generally limited to observe only the surfaces
of stars, including the sun. But with asteroseismology we can use
stellar oscillations - or star quakes - to look beyond the surface
and learn about the stellar structure in much more detail. This is
the topic of the next chapter.
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