Helium is the second most abundant element in the visible universe and there is accordingly a vast amount of information about its distribution from optical and radio emission lines in nebulae, optical absorption lines in spectra of hot stars (effective temperatures above 104 K), scale heights of the atmospheres of the major planets and the influence of initial helium content on stellar structure and evolution and pulsation. However, primordial helium has been enhanced by astration to varying degrees in different objects, so that the observed mass fraction Y is in general only an upper limit to the primordial mass fraction Yp, and, conversely, Y may be reduced in planetary and certain stellar atmospheres by gravitational settling. Thus in the 1960's a certain amount of confusion reigned as to whether there was really a universal lower limit to Y as required by SBBN (cf. Pagel 1982). The existence of hot subdwarfs with weak helium lines was one source of confusion until Sargent & Searle (1967) showed that other abundance anomalies in these stars implied processes similar to those operating in the chemically peculiar A-type stars of Population I. Christy (1966) had already shown that the pulsational characteristics of RR Lyrae stars, many of which occur in halo globular clusters, demand a substantial helium content in the envelope and Iben (1968) likewise deduced a substantial initial helium content from the relative numbers (i.e. lifetimes), in globular clusters, of core helium-burning stars on the horizontal branch of the luminosity-effective temperature diagram (the Hertzsprung-Russell diagram) and hydrogen shell-burning stars on the red giant branch. Searle & Sargent (1972) finally clinched the matter by taking spectra of two blue compact galaxies discovered by Zwicky, I Zw 18 and II Zw 40, and showing that their light is dominated by H II regions with very low abundances of observable heavy elements (N, O, Ne and S), but nearly normal helium.
A thorough account of all the data available up to 1983 has been given in the proceedings of the ESO Workshop (Shaver, Kunth & Kjar 1983) and table 1 gives a slightly updated version of the results presented there, which strongly support the view that there is a universal floor corresponding to a primordial abundance 0.20 < Yp < 0.25, but for reasons briefly indicated in the last column of the table it is difficult to reach the better than 5 per cent precision (i.e. less than ±.01) that is needed in order to constrain the SBBN model significantly. Fig. 1 shows that, with the lower limit to derived from D + 3He, and standard values of N and 1/2, Yp needs to be at least 0.235 for consistency with SBBN and that a firm upper limit to Yp may give a tighter upper limit to than can be derived from other elements. However, the Sun and hot stars give only a rather imprecise upper limit and in subdwarfs and halo globular clusters, which are likely to approximate pristine material quite closely, the relevant stars are too cool to show helium lines in their spectra and one has to rely on systematics of the HR diagram: location of the zero-age main sequence (ZAMS, needing knowledge of distance); effective temperature at the blue edge of the instability strip where RR Lyrae variables are found; the difference in magnitude between the horizontal branch (HB) and a well-defined point on the ZAMS; and the relative numbers of stars in the HB and red giant stages. These estimates are dependent on a full understanding of stellar opacities, reaction rates, instability mechanisms, convection, semi-convection etc. (cf. Caputo 1985), so that the systematic uncertainties are difficult to judge.
|Sun||< .28±.02||Struc.; oscill.||Turck-Chieze 88||Eq. of state|
|< .28±.05||Promiinence em.lines||Heasley 78||Self-abs.|
|Hot stars||< .28±.04||Absorption lines||Wolff 85||Precision|
|Subdwarf||.19±.05||Lum. and eff. temp||Carney 83||Parallax; convection|
|Glob.clust.||.23:||Variables and lum.evol.||Caputo 83||Physical basis|
|.20±.03 :||N(HB)/N(RG)||Cole 83||of stellar structure|
|.23±.02||'' : ''||Buzzoni 83||and evolution|
|.24±01||'' : ''||Caputo 87||(see Caputo 1985)|
|Plan.nebula||.22±.02||opt.em.lines||Peimbert 83||Self-and Gal. enrichment|
|Gal. HII reg||.22:||opt.and radio em.lines||Mezger 83||He0 and Gal. enrichment|
|Extragal.||.233±.005||opt. em.lines||Lequeux 79||(See text)|
|H II region||< .243±.010||'' : ''||Kunth 83|