Stellar evolution
From Aetilc
In astrophysics and cosmology, stellar evolution refers to the life history of stars that is driven by the interplay of internal pressure and gravity .
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[edit] Overview
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.
[edit] Protostar
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.
[edit] Main sequence
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.
[edit] Post-main sequence
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.
