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Stellar Evolution

Stellar evolution

In astrophysics and cosmology , stellar evolution refers to the life history of stars that is driven by the interplay of internal pressure and gravity .

Essentially, throughout the life of a star a tension exists between the compressing force of the star's own gravity and the expanding pressures generated by the nuclear reactions taking place in its core. After cycles of swelling and contraction associated with the burning of progressively heavier nuclear fuels , the star eventually expends its useable nuclear fuel and resumes contraction under the force of its own gravity. There are three possible fates for such a collapsing star. The particular fate for any star is determined by the mass of the star left after blowing away its outer layers.

A star less than 1.44 times the mass of the Sun (termed the Chandrasekhar limit) collapses until the pressure compacted electron clouds exerts enough pressure to balance the compressing force of gravity. Such stars become white dwarfs that are contracted to a radius the size of a planet. This is the fate of most stars.

If a star retains between 1.4 and roughly three times the mass of the Sun, the pressure of the electron clouds is insufficient to stop the gravitational collapse and stars of this mass continue their collapse to become neutron stars. Although neutron stars are only a few miles across, they have enormous density. Within a neutron star the nuclear forces and the repulsion of the compressed atomic nuclei balance the crushing force of gravity.

With the most massive stars, however, there is no known force in the Universe that can stop the final gravitational collapse and such stars collapse to form a singularity a geometrical point of infinite density. As such a star collapses, its gravitational field warps spacetime and a black hole forms around the singularity.

Gravitational collapse is the process which provides the energy required for star formation, which starts with hydrogen fusion in a protostar at a heat of over 15 million K. Gravity, always directed inwards, decreases the radius of interstellar gas clouds, causing them to collapse and form a protostar, the immediate precursor of a star. Interstellar gas is initially cold, but it is heated by the gravitational energy released by the cloud contraction process. The radius of the protostar will continue to shrink under the influence of gravity until enough internal gas pressure, always directed outwards, builds up to stabilize the collapse. At this stage, the protostar is still too cold for hydrogen fusion to be initiated. Protostars can be detected by infrared spectroscopy because the initial warming event releases infrared electromagnetic radiation. If the mass of the protostar is less than 0.08 solar masses, the temperature of its core never reaches the range required for nuclear fusion and the failed star becomes a brown dwarf.

If, however, the mass of the protostar exceeds 0.08 solar masses, hydrogen fusion can proceed and the protostar becomes a main sequence star, with average surface temperatures of 10,800°F (6,000°C) (the internal and coronal temperatures measure in the millions of degrees). Most stars in the Universe are main sequence stars and are found on the diagonal of a Hertzsprung-Russell diagram. The main sequence stage of star evolution is the most stable state a medium-sized star can reach, and it can last for billions of years as such stars undergo very gradual and slow changes in luminosity and temperature. This is because pressure and gravitational forces are in equilibrium and the core has reached the temperature required for the fusion of hydrogen to helium to proceed smoothly. The time spent by a star in the main sequence is a function of its mass. The more massive the star, the less time spent on the main sequence. Although massive stars have large amounts of fuel, hydrogen fusion proceeds so quickly that it is completed within a few hundred thousand years. The fate of such massive stars is to explode violently. Smaller stars fuse their hydrogen at a slower rate. The lightest stars created in the early history of the Universe, for example, are still on the main sequence. The Sun is approximately midway through its main sequence life.

A post-main sequence star has two distinct regions, consisting of a core of helium nuclei and electrons surrounded by an envelope of hydrogen. With two protons in its nucleus, helium requires a higher fusion temperature than the one at which hydrogen fusion is proceeding. Without a source of energy to increase its temperature, the core cannot counter the effect of gravitational collapse and it starts to collapse, heating up as it does. This heat is transferred to the fusing hydrogen layer, which increases the luminosity of the shell and causes it to expand. As it expands, the outer layers cool off. At this point, the star is characterized by expansion and cooling of its shell, which causes it to become redder with increased luminosity. This is termed the red giant phase. When the Sun reaches this stage, it will be large enough to include Mercury in its sphere and hot enough to evaporate Earth's oceans . The core temperature of a red giant is on the order of 100 million K, the threshold temperature for the fusion of helium into carbon . A red giant, however, is initially stable, as pressure and gravity reach equilibrium.

If helium continues to accumulate in the core as the outer portions of the hydrogen envelope continue to fuse, eventually the helium in the core starts fusing into carbon in a violent event referred to as a helium flash, lasting as little as a few seconds. During this phase, the star gradually blows away its outer atmosphere into an expanding shell of gas known as a planetary nebula.

A star takes thousands of years to go through the red giant phase, after which it evolves into a white dwarf. It is then a small, hot star with a surface temperature as high as 100,000 K that makes it glow white. Because of its small size, a white dwarf has a very high density. A white dwarf consists of those elements that were created in its previous evolutionary phases via nucelosynthesis. The original hydrogen was fused into helium then totally or partly fused into carbon. In addition, heavier elements fuse from the carbon. The temperature of a white dwarf is not high enough to initiate a new cycle of fusion. In time, it eventually becomes a black dwarf as it loses its residual heat over billions of years. The size of a white dwarf is limited by a process called electron degeneracy. Electron degeneracy is the stellar equivalent of the Pauli exclusion principle, as is neutron degeneracy. No two electrons can occupy identical states, even under the pressure of a collapsing star of several solar masses. For stellar masses smaller than about 1.44 solar masses, the energy from the gravitational collapse is not sufficient to produce the neutrons of a neutron star, so the collapse is effectively stopped. This maximum mass for a white dwarf is called the Chandrasekhar limit.

When a massive star has fused all of its hydrogen, gravitational collapse is capable of generating sufficient energy so that the core can begin to fuse helium nuclei to form carbon. If the process can go beyond the red giant stage, the star becomes a supergiant. Following fusion and disappearance of the helium, the core can successively burn carbon and other heavier elements until it acquires a core of iron , the heaviest element that can be formed by natural fusion. Another possible fate of white dwarfs is to evolve into novae or another type of supernovae. Novae occur in binary star systems in which one star is a white dwarf. If the companion star evolves into a red giant, it can expand far enough so that gas from its outer shell can be pulled onto the white dwarf. The white dwarf accumulates the additional gas until it reaches nuclear fusion temperatures, at which point the gas ignites explosively into a nova.

Alternatively, a white dwarf may accumulate enough material from its binary star to exceed the Chandrasekhar limit. This results in a sudden and total collapse of the white dwarf, with temperature increases in ranges capable of initiating rapid carbon fusion and subsequent explosion of the white dwarf into a spectacular supernova, that can shine with the brightness of 10 billion suns with a total energy output of 1044 joules, equivalent to the total energy output of the Sun during its entire lifetime.

See also Astronomy; Big Bang theory; Bohr model of the atom; Cosmology; Quantum electrodynamics (QED); Solar sunspot cycles; Solar system; Stellar life cycle

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stellar evolution

stellar evolution, 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.

Phases of Stellar Evolution

Contraction of the Protostar

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.

Mature Stars and the Main Sequence

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.

Old Stars and Death

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.

Validating the Theory of Stellar Evolution

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.


See I. S. Shklovsky, Stars: Their Birth, Life, and Death (1978); D. A. Cooke, The Life and Death of Stars (1985); A. Harpaz, Stellar Evolution (1994); I. Asimov, Star Cycles: The Life and Death of Stars (1995).

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