The
way in which a star evolves during its lifetime
depends principally on its MASS.
In this article, an outline is given of the development of
a solar mass star (comparable to the Sun) and a much more massive star,
10 times the mass of the Sun (eg. Rigel in the constellation of Orion).
The interior of a nebula is a dynamic environment, subjected to the gravitational and radiative effects of nearby stars. If an adjacent star were to suffer a supernova, for example, the compression wave generated in the nebula might cause part of it to collapse. A region of the nebula with a mass of, say, 1000 solar masses could start to contract. Contraction causes the cloud to warm up, and, if there was any initial rotational motion, conservation of angular momentum produces a large increase in the angular velocity of the cloud constituents.
Such a rotating mass is very unstable and tends to fragment into smaller "clumps", each of which may contract further under gravity to form a single star. The group of contracting fragments constitutes the basis of a cluster of stars (eg. the Pleiades star cluster in Taurus).
After
several thousand years of contraction the cloud temperature may still only
be of the order of 100K, but as contraction increases the core temperature
of each fragment rises. Each fragment is becoming a protostar. Such
cool objects radiate in the infra-red region of the electromagnetic spectrum
and may be observed using telescopes sensitive
to these wavelengths. Many protostars are currently forming in the aforementioned
nebula, above.
At this stage, a protostar will be well off to the lower right (low T, low luminosity) in the Hertzsprung-Russell (HR) diagram.
Given a few million years of contraction, and as the loss of gravitational potential energy is converted into kinetic energy and then to heat, the protostar core temperature will reach a few million degrees Kelvin. At this point nuclear fusion reactions start. These provide an outward radiation pressure which opposes gravity and halts the stars collapse. A dynamic equilibrium is established. At a surface temperature of 2000-3000K, the protostar is (according to Wien’s Law) still emitting radiation predominantly in the IR.
For a
solar mass star, the core temperature stabilises at around 107
K. At this temperature, the main fusion process generating energy in the
core of the star is the proton-proton
chain (hydrogen "burning").
The surface temperature of the star is now about 6000K, emitting radiation
predominantly in the visible region. Such a star would be of spectral
class G, like our Sun. The star is now shining on the Main Sequence
of the HR diagram.
Even
though a star is a massive object, its supply of hydrogen is not unlimited.
Eventually, the amount of H in the core becomes depleted, much of the initial
H having been converted into He. As the number of p-p fusions diminishes,
the equilibrium is upset and the core contracts under gravity. As this
happens the core temperature rises.
The temperature needed to initiate He fusion (by the "triple alpha" process) is of the order of 108K. In the early stages of core collapse the temperature is still far too low for He nuclei to fuse. Instead, H in the outer layers of the star starts to heat up and undergo p-p reactions. Clearly, the structure of the star undergoes radical changes: the core becomes smaller and heats up and the outside of the star expands and cools.
The star swells to become a red giant, with 100 times its main sequence luminosity. The transition from the main sequence to the red giant phase may take 500 million years.
Meanwhile, the core temperature eventually becomes hot enough to allow He fusion to occur. Fusion of He halts the core collapse and further increases the stars luminosity.
The red
giant stage can be observed above and to the right of the main sequence
on the HR diagram.