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EQUATION-OF-STATE FOR WHITE DWARF STARS

W HITE dwarfs are a distinct class of stellar objects, that were first encountered in F.W. Bessel’s observation of Sirius in 1836. Bessel examined the wobbling features of the proper motion of Sirius for a decade, and this finally led him to conclude it had an unseen partner—later duped Sirius B. Now we know that Sirius B is a very small star with a mass around 1M⊙, a radius of only 0.008R⊙, and a temperature of 27000◦K. This means that the dwarf have very high mass density and should really be burning its internal storage of hydrogen, but observations tells us that they are devoid of thermonuclear fusion processes and are hence rather ’cold’ and ’dead’ objects made of heavy atoms, mostly helium, carbon and oxygen. White dwarfs are the natural end-point for stellar evolution of smaller stars. Very heavy stars terminates their life in violent supernova processes, whereas stars of masses only up to some solar masses ends their lifes as red giants and expels some of their mass in a planetary nebula, leaving the dense core as a white dwarf progenitor. Since white dwarfs have no effective radiation pressure from thermonuclear reactions, they must have something else to protect them from a gravitational collapse; it turns out that the degenerate electron pressure is able to fully support the star, and this report will go into great details about the nature of this pressure and the equation-of-state (EoS) found in the dwarfs. Adding mass to a dwarf leads to a collapse, where the supporting electron pressure is robbed by inverseβ decays. This leads to the interesting limiting mass as found by S. Chandrasekhar and published in 1931 [1, 2]—this mass will also be examined.