Stellar evolution is the series of phases that a star passes through between its birth and its death. The following article describes the evolution of typical stars.

The space between stars contains gas and dust at a very low density. This interstellar matter tends to gather into clouds. Sometimes the density becomes high enough so that gravity causes contraction, leading to the formation of a protostar. As a protostar slowly contracts, its pressure and temperature increase, the temperature rise being the result of the release of gravitational energy. Any hot object radiates energy, and the protostar eventually becomes hot enough to shine, although temperatures are not yet great enough to sustain nuclear reactions. The pressure builds up almost enough to balance gravity, but the radiation emitted drains energy and inhibits the ability of the internal pressure to complete the balance. Therefore the contraction (and heating) slowly continue.

The temperature at the center of the protostar finally becomes high enough to initiate nuclear reactions and the subsequent release of nuclear energy. Hydrogen is the most abundant element, and hydrogen-burning reactions, in which hydrogen is converted to helium with an accompanying release of large amounts of energy, are initially the most important reactions. When the nuclear energy released exactly balances the radiation energy lost into space, the protostar finally enters a state of balance, and contraction ceases. At this point the object becomes a true star.

Hydrogen-Burning and Helium-Burning Stages
A star that is in balance and burning hydrogen in its core is called a main sequence star. On a plot of luminosity versus temperature, known as a Hertzsprung-Russell diagram, main-sequence stars fall along a diagonal line. All stars begin their careers in the main-sequence phase. If a main-sequence star has a large mass, it will have a high surface temperature and will therefore be very luminous. If it has a small mass, it will be rather cool and faint. The Sun is a main-sequence star somewhat above average in mass, surface temperature, and radiant-energy output.

When the hydrogen fuel in the core is used up, the star loses its main-sequence status. This can happen in less than a million years for the most luminous stars but takes many trillions of years for the faintest. The Sun has a main-sequence lifetime of about 10 billion years, of which half is over.

When all the core hydrogen has been converted to helium through nuclear reactions, the release of nuclear energy stops. The star falls out of energy balance, and the central portions contract further under gravity and grow still hotter. The fact that nuclear reaction ceases, however, does not mean that a star will no longer radiate energy into space. Stars shine because they are hot, and a post-main-sequence star is still quite hot.

As the central parts of a star get hotter, nuclear reactions resume, either in the form of hydrogen-burning in the regions just outside the helium core or in the form of helium-burning reactions in the core itself. Any nucleus can undergo nuclear reactions if conditions are violent enough. Hydrogen-burning takes place at temperatures of about 10 million K, but temperatures of about 100 million K are necessary to ignite helium. Hydrogen-burning produces helium, while helium-burning produces carbon, oxygen, and other rather heavy nuclei. The heavier the nucleus, the higher the temperature required to bring it into nuclear reaction.

The later nuclear reactions are brought about by the further contraction and heating of the inner parts of the star, after the hydrogen has been exhausted in the core. During this phase the outer layers of the star actually expand and cool, and the luminosity can become quite high. The star, no longer on the main sequence, is now what astronomers call a giant or, if the luminosity is extremely great, a supergiant.

Old Age and Death
In old age, with further contraction and heating of the inner parts of a star, heavier particles are ignited. The burning of helium is followed by the burning of carbon, oxygen, silicon, and so on. As the giant or supergiant star ages, it builds up layers of successively heavier elements in its interior, with the heaviest materials in the core and lighter materials in shells around the core. This process cannot go on indefinitely, however. In smaller stars the material can become so dense that it resists further contractions, a state known as "degeneracy." The star then slowly radiates away what heat energy is available and ends its life as a cold, dark body. Stars of this type that are observable are known as white dwarfs. They have not yet cooled to the point at which they fade from view and become black dwarfs.

The densities needed to produce electron degeneracy are quite large. White dwarfs themselves have densities that are typically tons per cubic centimeter. The electrons in a star cannot become degenerate if the star has a mass greater than about 1.4 times the mass of the Sun. This mass, known as the Chandrasekhar limit, is an upper limit for the mass of a white dwarf. Stars with masses greater than this limit may undergo violent explosions that eject much of their material into space. These explosions are called supernovas. The core of such a star apparently collapses with violence, driving off the outer layers and breaking down the core particles into neutrons. In several cases the remnants of supernova explosions have been detected and are known as neutron stars. Rotating neutron stars are called pulsars. Neutron stars, which cannot contract further, end as cold, dark bodies.

If a star never sheds enough mass to become degenerate, there is nothing to stop continued contraction when the nuclear sources have been completely used up. The star will keep getting smaller until it becomes an object known as a black hole.

Thomas L. Swihart

Bibliography: Cohen, Martin, In Darkness Born: The Story of Star Formation (1988); Harpaz, Amos, Stellar Evolution (1994); Hartmann, Lee, Accretion Processes in Star Formation (2001); Kaler, James, Extreme Stars: At the Edge of Creation (2001); Kippenhahn, Rudolf, and Weigert, Alfred, Stellar Structure and Evolution (1990); Lada, C. J., and Kylafis, N. D., eds., The Physics of Star Formation and Early Stellar Evolution (1991); Swihart, Thomas L., Quantitative Astronomy (1991).