Final Stages of Stellar Evolution
During the final evolutionary stages the nuclear fuel is consumed, and the star emits radiation owing to cooling. It is relatively cold and has a very high density, while the pressure arises mainly from the matter degeneracy. Chandrasekhar  in 1931 obtained the fundamental result that a star where the pressure is due to degenerate electrons has a maximum mass. At ρ > 1.15 × 109 g cm−3 (for 56Fe, for other nuclei see Table 10.1) the neutronization begins, and stars become unstable. Stable (neutron) stars reappear only at ρ c ≈ 1.5 × 1014 g cm−3 and exist to densities ρ c ≈ 1.15 × 1015 g cm−3, where instability arises from general relativity (GR) effects. Oppenheimer and Volkov established in 1939 the existence of a mass limit for neutron stars , but its value has been recalculated many times for various equations of state. Solving the equilibrium equations for cold stars in Newtonian theory (9.3.13a) and (10.1.1a) and in GR (11.2.3) and (11.2.4) at a given equation of state P(ρ) has allowed derivation of the curve M(ρ c ) demonstrating the existence of two maximum masses and an instability region (falling M(ρ c )). Figure 11.1 represents the curve M(ρ c ) from  for the equation of state corresponding to a minimum energy of matter in full thermodynamic equilibrium.
KeywordsBlack Hole Neutron Star Accretion Disc White Dwarf Stellar Evolution
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