Text 20 stellar evolution

There is at present no generally accepted theory of the evolution of stars. However, the following crude picture, based mainly on oversimplified theo­retical considerations, may at least serve as the basis of comparison with ob­servations.

In regions containing gas and dust under «suitable» conditions, conden­sations of various masses form. Such a condensation (initially still with fair­ly low density and temperature) begins to radiate and to contract gravitationally, its temperature, temperature gradient, and luminosity rising. The time taken for such a photo-star to shrink to its dimensions on the main sequ­ence depends on its gravitational energy content and its luminosity throughout the contraction. Presumably a lower limit to this time is obtained by assum­ing a luminosity throughout the contraction equal to the main sequence value. This time is about 2.5 x 10⁷ years for a photo-star of solar mass M, about 10⁶ years for 3M etc. If the photo-star contained an appreciable amount of P, Zi, Be, or B it will burn these elements at different stages of its contrac­tion, increasing the total contraction time slightly.

Once the star reaches its main sequence position, the energy produc­tion from the H→ He conversion keeps pace with the radiative energy loss and the contraction stops. The time taken for the conversion of only 1 per cent of the stellar mass from hydrogen into helium is as long as 10⁹ years for a photo-star of mass M, a few times 10⁷ years for 3M.

Let us assume the star is poorly mixed, which seems likely at least for a considerable fraction of stars. Hydrogen is then converted into helium al­most exclusively in the hottest central regions of the star. Thus a core of different composition (and mean molecular weight) from the rest of the star is formed. When this core contains about 10 per cent of the mass of the star, its configuration ceases to be stable. Its core will contract and heat up, while its envelope expands and cools. At least in the early stages of this develop­ment the star will become much redder and slightly more luminous (red giant or subgiant). Details of this development are not yet known, but it seems likely that in the centre of at least some such stars high enough tem­peratures and densities are reached for the conversion of He into C, O, Ne. It should be remembered that even after leaving the main sequence the bulk of the star still consists of hydrogen and the bulk of the energy production still comes from the H—He conversion in a thin shell just outside the core. But it is rather likely that the onset of further nuclear energy production in the centre of the core can alter the structure of a star. Such processes may be involved in some types of variable stars or supergiants.

It also seems likely that at some late stage in the development of a star, it will become unstable enough to lose mass from its outer layers, either gra­dually or catastrophically. This loss of mass will presumably stop when the mass of the remaining core is less than a certain limit and this core will even­tually become a white dwarf. A white dwarf contains essentially no hydrogen in its interior, has too low temperature for nuclear reactions involving helium or heavier nuclei, and gravitational contraction is prevented by the high pres­sure of the degenerate electron gas. Having no possible sources of energy left, the white dwarf will cool down very slowly and, like old soldiers and some famous generals, merely fade away.

The above outline of stellar evolution at best represents a sort of ave­rage picture, since the evolution of different stars may well depend strongly on various factors. In particular, the structure of a star (as well as the time-scale of the evolution) depends strongly on its mass and its development will be affected by its original rotational velocity, its magnetic field and by the presence or absence of accretion from nearby gas clouds.