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1.4 Brief Outline of Stellar Evolution


An interstellar cloud collapses forming a generation of stars. The masses of the stars are spread over a range from maybe 0.01 Msun or less to maybe 100Msun or so. The distribution function of stellar masses at birth, known as the ``initial mass function'' (IMF), has more small stars than big ones.

Figure 1.6

Fig. 1.6. HR diagram for nearby stars with well-known paraflaxes (i.e. distances) from the HIPPARCOS astrometric satellite mission, after M.A.C. Perryman et al., Astr. Astrophys., 304, 69 (1995). The abscissa, B (blue) - V (visual) magnitude is a measure of ``redness'' or inverse surface temperature ranging from about 2 x 104 K on the left to about 3000 K on the right, while the ordinate measures luminosity in the visual band expressed by absolute magnitude MHp, in the HIPPARCOS photometric system. Luminosity decreases by a factor of 100 for every 5 units increase in M and the Sun's MV appeq MHp is 4.8. The main sequence forms a band going from top left to bottom right in the diagram, with the ZAMS forming a lower boundary. Subgiants and red giants go upward and to the right from the MS and a few white dwarfs can be seen at the lower left. Courtesy Michael Perryman.


The young stars, which appear as variable stars with emission lines known as T Tauri stars and related classes, initially derive their energy from gravitational contraction, which leads to a steady increase in their internal temperature (see Chapter 5). Eventually the central temperature becomes high enough (~ 107 K or 1 keV) to switch on hydrogen burning and the star lies on the ``zero-age main sequence'' (ZAMS) of the Hertzsprung-Russell (HR) diagram in which luminosity is plotted against surface temperature (Fig. 1.6). Stars spend most of their lives (about 80 per cent) in a main-sequence band stretching slightly upwards from the ZAMS; the corresponding time is short (a few x 106 years) for the most massive and luminous stars and very long (> 1010 years) for stars smaller then the Sun, because the luminosity varies as a high power of the mass and so bigger stars use up their nuclear fuel supplies faster.


When hydrogen in a central core occupying about 10 per cent of the total mass is exhausted, there is an energy crisis. The core, now consisting of helium, contracts gravitationally, heating a surrounding hydrogen shell, which consequently ignites to form helium and gradually eats its way outwards (speaking in terms of the mass coordinate). At the same time, the envelope expands, making the star a red giant in the upper right part of the diagram.


Eventually the contracting core becomes hot enough to ignite helium (3alpha -> 12C; 12C + 4He -> 16O) and the core contraction is halted.


In sufficiently big stars (>~ 10Msun) this process repeats; successive stages of gravitational contraction and heating permit the ashes of the previous burning stage to be ignited leading to C, Ne, O and Si burning in the centre with less advanced burning stages in surrounding shells leading to an onion-like structure with hydrogen-rich material on the outside. Silicon burning leads to a core rich in iron-group elements and with a temperature of the order of 109 K, i.e. about 100 keV.


The next stage of contraction is catastrophic, partly because all nuclear energy supplies have been used up when the iron group is reached, and partly because the core, having reached nearly the Chandrasekhar limiting mass for a white dwarf supported by electron degeneracy pressure, is close to instability and also suffers loss of pressure due to neutronization by inverse beta-decays. Further contraction leads to photodisintegration, which absorbs energy, and this leads to dynamical collapse of the core which continues until it reaches nuclear density and forms a neutron star. (If the mass of collapsing material is too large, then a black hole will probably form instead.) The stiff equation of state of nuclear matter leads to a bounce which sends a shock out into the surrounding layers. This heats them momentarily to high temperatures, maybe 5 x 109 K in the silicon layer, leading to explosive nucleosynthesis of iron-peak elements, mainly 56Ni (which later decays by electron capture and beta+ to 56Fe); more external layers are heated to lower temperatures resulting in milder changes. Assisted by high-energy neutrinos, the shock expels the outer layers in a supernova explosion (Type II and related classes); the ejecta eventually feed the products into the ISM which thus becomes enriched in ``metals'' in course of time. This scenario was first put forward in essentials by Hoyle (1946), and modern versions give a fairly good fit to the local abundances of elements from oxygen to calcium. The iron yield is uncertain because it depends on the mass cut between expelled and infalling material in the silicon layer, but can be parameterized to fit observational data, e.g. for SN 1987A in the Large Magellanic Cloud (LMC). The upshot is that iron-group elements are probably underproduced relative to local abundances, but the deficit is plausibly made up by contributions from supernovae of Type Ia. A subset of Type II supernovae may also be the site of the r-process (see Chapter 6).


For intermediate mass stars (IMS), ~ 1Msun leq M leq~ 8Msun, stages (i) to (iii) are much as before, bitt these never reach the stage of carbon burning because the carbon-oxygen core becomes degenerate first by virtue of high density, and later evolution is limited by extensive mass loss from the surface. After core helium exhaustion, these stars re-ascend the giant branch along the so-called asymptotic giant branch (AGB) track (see Fig. 5.14) with a double shell source: helium-burning outside the CO core and hydrogen-burning outside the He core. This is an unstable situation giving rise to thermal pulses or ``helium shell flashes'' in which the two sources alternately switch on and off driving inner and outer convection zones (in which mixing takes place) during their active phases (see Chapter 5). The helium-burning shell generates 12C and neutrons, either from 22Ne(alpha, n)25Mg or from 13C(alpha, n)16O, leading to s-processing, and the products are subsequently brought up to the surface in what is known as the third dredge-up process. This process leads to observable abundance anomalies in the spectra of AGB stars, carbon and S stars; see Figs. 1.7, 1.8). The products are then ejected into the ISM by mass loss in the form of stellar winds and planetary nebulae (PN), leaving a white dwarf as the final remnant. If the white dwarf is a member of a close binary system, it can occasionally be ``rejuvenated'' by accreting material from its companion. giving rise to cataclysmic variables, novae and supernovae of Type Ia (cf. Chapter 5).

Figure 1.7

Fig. 1.7. Photographic (negative) spectra of stars showing various aspects of nucleosynthesis. Top: (a) carbon star X Cancri with 12C/133C appeq 4 from H-burning by the CNO cycle and a suggestion of enhanced Zr; (b) peculiar carbon star HD 137613 without 13C bands in which hydrogen is apparently weak (H-deficient carbon star); (c) carbon star HD 52432, with 12C/13C ~ 4. Middle: (a) Normal carbon star HD 156074, showing the CH band and Hgamma; (b) Peculiar (H-poor) carbon star HD 182040 showing C2 but weak Hgamma and CH. Bottom: (a) normal F-type star xi Pegasi (slightly hotter than the Sun); (b) old Galactic halo population star HD 19445 with similar temperature (shown by Hgamma) but very low metal abundance (about 1/100 solar). After E.M. and OR. Burbidge, WA. Fowler & V Hoyle 1957, Rev. Mod. Phys., 29, 547 (B2FH). Courtesy Margaret and Geoffrey Burbidge.

Figure 1.8

Fig. 1.8. Spectra showing effects of s-process. Top: (a) normal G-type giant kappa Gem (similar temperature to the Sun); (b) Ba II star HD 46407, probably not an AGB star but affected by a companion which was. Middle: (c) M-type giant 56 Leo. showing TiO bands; (d) S-type AGB star R Andromedae showing ZrO bands, partly due to enhanced Zr abundance. Bottom: another spectral region of the same two stars showing Tc features in R And which indicate s-processing and dredge-up within a few half-lives of technetium (2 x 105 yrs) before the present. After B2FH (1957). Courtesy Margaret and Geoffrey Burbidge.

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