WHITE DWARFS

White dwarf stars, also called degenerate dwarfs, represent the endpoint of the evolution of stars with initial masses ranging from about 0.08 to about 8 solar masses. This large range encompasses the vast majority of stars present in our Galaxy and thus white dwarf stars represent the most common endpoint of Stellar Evolution. The defining characteristic of these objects is the fact that their mass is typically of the order of half that of the Sun, while their size is more akin to that of a planet. Their compact nature gives rise to large average densities and surface gravities.

White dwarfs are conceived to be the final evolutionary state of all stars whose mass is not high enough for supernova. It is believed that over 95% of the stars of our Galaxy will eventually end up as white dwarfs. After the hydrogen fusion stage of a main-sequence star of low or medium mass ends, it will expand to a red giant which fuses helium to carbon and oxygen in its core by the triple-alpha process. If a red giant does not have sufficient mass to generate the core temperatures required to fuse carbon, an inert mass of carbon and oxygen will build up at its center. Once it sheds its outer layers to form a planetary nebula, it will leave behind this core, which forms the remnant white dwarf. As a result, white dwarfs are composed of carbon and oxygen.

The first glimpse of the existence of these objects came in 1844, with the study carried out by F Bessel, the great German mathematician and astronomer, of the proper motion of Sirius, the brightest star in the night sky. The irregularities in the apparent motion of Sirius on the celestial sphere led Bessel to suggest the presence of an unobserved, solar-mass companion orbiting the bright primary. The companion, Sirius B, was first observed by A G Clark in 1862 and represents, together with 40 Eri B, the first known examples of white dwarf stars.

However, how could such stars withstand the tendency to collapse onto themselves under the influence of their gravitational field? The answer to that question came in 1925, when R H Fowler first applied the newly-developed principles of Quantum Mechanics to stars. He showed that, in white dwarf stars, the density is high enough for the gas of free electrons to become degenerate. Electrons are said to be degenerate when a majority of them occupy the lowest possible energy states available to them. This occurs, at fixed temperature, when the electrons are packed sufficiently close to each other. Because of the Pauli Exclusion Principle that no more than two electrons with oppositely directed spins may occupy the same energy state, the electrons retain kinetic motions even when cooled to zero temperature. The amplitude of this kinetic activity increases with increasing density, when electrons become more degenerate. In a white dwarf, the pressure generated by this kinetic motion, clearly of quantum mechanical origin, prevents the gravitational collapse of the star.

The first detailed stellar models appropriate to white dwarfs were calculated in the 1930s by S Chandrasekhar, who received the 1983 Nobel Prize in Physics for his achievements.

Questions to Ponder

• If a white dwarf is a dead star, why is it so hot?
• Can our Sun become a White Dwarf?