When a main-sequence star has consumed the hydrogen at its core, the loss of energy generation causes its gravitational collapse to resume and the star evolves off the main sequence. The path which the star follows across the HR diagram is called an evolutionary track.
If the star has a mass of less than 0.23Mʘ, it is predicted to directly become a white dwarf when the energy generation by nuclear fusion of hydrogen at its core comes to a halt.
In stars more massive than 0.23Mʘ, the hydrogen surrounding the helium core reaches sufficient temperature and pressure to undergo fusion, forming a hydrogen-burning shell and causing the outer layers of the star to expand and cool. The stage as these stars move away from the main sequence is known as the subgiant branch which is relatively brief and appears as a gap in the evolutionary track since few stars are observed at that point.
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When the helium core of low-mass stars becomes degenerate, or the outer layers of intermediate-mass stars cool sufficiently to become opaque, their hydrogen shells increase in temperature and the stars start to become more luminous. This is known as the red giant branch, which is a relatively long-lived stage and it appears prominently in H-R diagrams. These stars will eventually end their lives as white dwarfs.
The most massive stars do not become red giants, instead, their cores quickly become hot enough to fuse helium and eventually heavier elements and they are known as supergiants. They follow approximately horizontal evolutionary tracks from the main sequence across the top of the H-R diagram.
Supergiants are relatively rare and do not show prominently on most H-R diagrams. Their cores will eventually collapse, usually leading to a supernova and leaving behind either a neutron star or black hole.
Any main-sequence star with an initial mass of above 8 times the mass of the Sun has the potential to produce a neutron star.
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The neutron was discovered in 1932. In 1934, Walter Baade and Fritz Zwicky suggested that Supernovae involve a collapse of a massive star, resulting in a neutron star.
As the star evolves away from the main sequence, subsequent nuclear burning produces an iron-rich core.
When all the nuclear fuel in the core has been exhausted, the core must be supported by degeneracy pressure alone.
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Exceeding the Chandrashekar limit
Further deposits of mass from the shell burning cause the core to exceed the Chandrashekar limit.
Electron degeneracy pressure is overcome and the core collapses further, sending temperatures soaring to over 5×109K.
At these temperatures, photodisintegration (the breaking up of iron nuclei into alpha particles by high-energy gamma rays) occurs.
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As the temperature climbs even higher, electrons and protons combine to form neutrons via electron capture, releasing a flood of neutrinos.
e– + p → n + υe
When densities reach a nuclear density of 4×1017kg/m3, a combination of strong force repulsion and neutron degeneracy pressure halts the contraction.
The infalling outer envelope of the star is halted and flung outwards by a flux of neutrinos produced in the creation of the neutrons, becoming a Supernova.
The remnant left is a neutron star.
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