life history of a star, beginning with its condensation out of the interstellar gas (see interstellar matter) and ending, sometimes catastrophically, when the star has exhausted its nuclear fuel or can no longer adjust itself to a stable configuration. Because a star's total energy reserve is finite, a star shining today cannot continue to produce its present luminosity steadily into the indefinite future, nor can it have done so from the indefinite past. Thus, stellar evolution is a necessary consequence of the physical theory of stellar structure, which requires that the luminosity, temperature, and size of a star must change as its chemical composition changes because of thermonuclear reactions.
The initial phase of stellar evolution is contraction of the protostar from the interstellar gas, which consists of mostly hydrogen, some helium, and traces of heavier elements. In this stage, which typically lasts millions of years, half the gravitational potential energy released by the collapsing protostar is radiated away and half goes into increasing the temperature of the forming star. Eventually the temperature becomes high enough for thermonuclear reactions to begin; if the mass of the protostar is too small to raise the temperature to the ignition point for the thermonuclear reaction, the result is a brown dwarf, or “failed star.” In these thermonuclear reactions, loosely called “hydrogen burning,” four hydrogen nuclei are fused to form a helium nucleus (see nucleosynthesis). This point in time is conventionally called age zero.
Many protostar contractions have been observed in isolated gas clouds; that is, where one cloud contracted to form one star. However, in 1995, the first example of a star-forming region was found in the Eagle Nebula, some 7,000 light-years from the earth. In this region, stars are being formed at the tips of long, fingerlike columns stretching from a huge cloud of interstellar gas and dust; the columns are being eroded by radiation (a process called photoevaporation) from stars in the vicinity, leaving scattered knots of matter that contract into stars.
Once formed, a star settles into a long “middle age” during which it shines steadily as it converts its hydrogen supply into helium. For stars of a given chemical composition, the mass alone determines the luminosity, surface temperature, and size of the star. The luminosity increases very sharply with an increase in the mass; doubling the mass (which is proportional to the energy supply) increases the luminosity (which is proportional to the rate of using energy) more than 10 times. Hence the more massive and luminous a star is, the faster it depletes its hydrogen and the faster it evolves.
Because the middle age of a star is the longest period in stellar evolution, one would expect most of the observed stars to be at this stage and to show a strong correlation of luminosity with color (color is a measure of stellar temperature). This prediction is confirmed by plotting stars on a Hertzsprung-Russell diagram, in which the majority of stars fall along a diagonal line called the main sequence. The main sequence is most heavily populated at the low luminosity end; these are the stars that evolve most slowly and so remain longest on the main sequence.
As a star's hydrogen is converted into helium, its chemical composition becomes inhomogeneous: helium-rich in the core, where the nuclear reactions occur, and more nearly pure hydrogen in the surrounding envelope. The hydrogen near the center of the core is consumed first. As this is depleted, the site of the nuclear reactions moves out from the center of the core and fusion occurs in successive concentric shells. Finally fusion occurs only in a thin, outer shell of the core, the only place where both the hydrogen content and the temperature are high enough to sustain the reactions.
As the helium content of the star's core builds up, the core contracts and releases gravitational energy, which heats up the core and actually increases the rates of the nuclear reactions. Thus the rate of hydrogen consumption rises as the hydrogen is used up. To accommodate the higher luminosity resulting from the increased reaction rates, the envelope must expand to allow an increased flow of energy to the surface of the star. As the outer regions of the star expand, they cool.
The star now consists of a dense, helium rich core surrounded by a huge, tenuous envelope of relatively cool gas; the star has become a red giant. Eventually, the contracting stellar core will reach temperatures in excess of 100 million degrees Kelvin. At this point, helium burning sets in. With the ignition of that process, the expansion of the envelope is halted and then reversed; the star retreats from the red giant phase, shrinking in size and luminosity, and reapproaches the main sequence. The exact course of evolution is uncertain, but as the star recrosses the main sequence, it will probably become unstable. The star may eject some of its mass or become an exploding nova or supernova star; at the very least, it will become a pulsating variable star, possibly a Cepheid variable.
In the later stages of evolution, further contraction and elevation of temperature open up new thermonuclear reactions. It is believed that the heavier elements in the universe, up to iron, were synthesized in the interiors of stars by a variety of intricate nuclear reactions, many involving neutron absorption. Elements heavier than iron are made in supernova explosions. As a result of the nuclear reactions, the chemical composition of the late-stage star becomes highly inhomogeneous; its structure is fractionated into a number of concentric shells consisting of different elements around an iron core.
The final outcome of stellar evolution depends critically on the remaining mass of the old star. The vast majority of stars do not develop iron cores. If the mass is not greater than the Chandrasekhar mass limit (1.5 times the sun's mass), the star will become a white dwarf, glowing feebly for billions of years by radiating away its remaining heat energy until it becomes a black dwarf, a totally dead star. If the star is too massive to become a stable white dwarf, contraction will continue until the temperature reaches about 5 billion degrees Kelvin. At this temperature the iron nuclei in the core begin to absorb electrons; this creates neutron-rich isotopes and simultaneously deprives the core of its pressure. With further collapse and increase in density, the core becomes a special kind of rigid solid. At still higher density, the solid “evaporates” as the nuclei break up into free neutrons. The resulting neutron fluid forms the core of a new astrophysical body, called a neutron star, of which pulsars are examples. If the stellar mass is too great to be stable even as a neutron star, complete gravitational collapse will ensue and a black hole will form.
Because the computed lifetimes of stars range from millions to billions of years, one cannot follow an individual star through its life history observationally, or even observe significant changes in the whole span of human history, except from the violent events of nova and supernova explosions. However, new stars are continually being formed and hence stars of all ages exist at the present epoch; examples of the various stages of stellar evolution can be found in different stars. The age of a star is not a directly observable characteristic but must be inferred from the very evolutionary theory one is trying to validate. Confidence in this circular reasoning results from its self-consistency and its ability to draw together into a unified picture a wide variety of observational data on individual stars, clusters of stars, and galaxies.
See cosmology; star clusters.