Helium: abundance

Opacity effects are also important. This can refer to differences in the treatment of interior opacities or to the effects of uncertain stellar compositions on the opacities. An increase in opacity makes temperature gradients larger, keeps the star convective for longer, raises Tf,cz once the radiative core develops and so leads to enhanced Li depletion. Opacity is increased by an increase in overall metallicity or a decrease in the Helium abundance. Changes of only 0.1 dex in metallicity can lead to an order of magnitude change in Li depletion (e.g. see Fig. 2 of [37]). 

We adopted as the present chemical composition the O/H, C/O, N/O, and Fe/H abundances from the H II region Hubble V (Peimbert et al. 2005) and from A-type supergiants (Venn et al. 2001). With these abundances and assuming the solar abundances by Asplund et al. (2005), we have determined its metallicity (Z = 0.6Zq). Since Venn et al. find no metallicity gradient, we have assumed that at present the ISM is well mixed. We have obtained the amount of gaseous mass, Mgas = 2 x 108Mq, based on the H I measurement inside r < 5 kpc, the present helium abundance, and an estimation of M(H2). 

Abstract. New results of the Primordial Helium abundance (Yp) measurement by radio recombination line (RRL) observations from five galactic HII regions are presented. The RRL observations were carried out with two telescopes RT32 (22.4 and 8.3 GHz, Medicina, Italy) and RT22 (36.5 and 22.4 GHz, Pushchino, Russia). The results of the first run of the low frequency RRL observations (408 MHz) with the Croce del Nord radiotelescope (Medicina Observatory, Italy) are also presented. 

Apart from improved calculations of the primordial helium abundance by Alpher, Follin and Herman (1953) and in a prescient paper by Hoyle and Tayler (1964). 

Helium is the second most abundant element in the visible Universe and accordingly there is a mass of data from optical and radio emission lines in nebulae, optical emission lines from the solar chromosphere and prominences and absorption lines in spectra of hot stars. Further estimates are derived more indirectly by applying theories of stellar structure, evolution and pulsation. However, because of the relative insensitivity of Tp to cosmological parameters, combined with the need to allow for additional helium from stellar nucleosynthesis in most objects, the requirements for accuracy are very severe better than 5 per cent to place cosmological limits on Nv and better still to place interesting constraints on t] or One can, however, assert with confidence that there is a universal floor to the helium abundance in observed objects corresponding to 0.23 < Fp < 0.25. 

The B stars, with 104 K < Teff < 3 x 104 K, show He I lines in absorption, but again these are young objects and the inferred helium abundances are not yet very precise. 

Table 4.3, but the errors are hard to quantify because many features of stellar evolution are sensitive to uncertainties in the input physics and the helium abundance is generally only one of many factors influencing stellar structure at different stages of evolution (see Chapter 5). 

In recent years, intensive efforts have been made to improve the determination of the pre-Galactic helium abundance (generally assumed to be the same as FP) and AF/AZ from observations of extragalactic H II regions. Advantages of this method are ... 

The two error terms refer to Yp and the regression slope respectively. In contrast to some earlier work, based on less homogeneous data sets and apparently affected by underlying absorption lines, notably in I Zw 18, this result, together with a similar one by Peimbert, Luridiana and Peimbert (2007), gives a primordial helium abundance in excellent agreement with the one predicted theoretically on the basis of the microwave background fluctuations and the lower estimates of deuterium abundance (see Fig. 4.3), a comparatively small value of about 2 for AT/AZ and no... 

They carry the stamp of the Big Bang, whence their great interest for cosmology. Their lithium content in particular is a precious clue as to the nucleonic density of the Universe, combined with deuterium and helium abundances measured in extremely metal-poor media (see Appendix 1). 

The bulk of the analyzed CSPN is. hotter than 60000K, covering a wide range in gravity. They are clearly separated from the classical sdOs and the sdB/sdOB stars. Most of them have normal helium abundances. However, some cases of intermediate and extreme enrichments as well as depletion of helium have been found too. (Note that the helium-poor CSPN are almost at a DA white dwarf stage). 

It is remarkable that He-rich stars appear to have a higher rate of mass loss than stars with solar-type atmospheres with the same Teff and L-values The Wolf-Rayet stars have, on the average, Al-values that are 140 times larger than the values for corresponding 0 and B type stars (De Jager et al., 1987). This may be due to the fact that WR stars, with their large Helium abundance, are relatively closer to their Eddington limit than the most luminous 0-type stars. 

Saio, Kato, and Nomoto (1988) recently examined under what conditions a massive star undergoes a blue-red-blue evolution. The evolution of a star of initial mass 20 M0 star in the HR diagram is shown in Figure 1 from the zero-age main-sequence through carbon ignition at the center. The metallicity in the envelope was assumed to be Z = 0.005 and the Schwarzschild criterion was adopted. The star shows the three types of evolutionary path (A, B, C) depending on the mass loss, metallicity, and the change in the helium abundance Y in the envelope. 

Figure 2b Envelope solutions for different metallicity. For log Teff = 4.2, the lines for the constant total mass (M = Menv + 5.45 M = 20 and 16 M ) are given as a function of the luminosity L and the mass fraction of heavy elements Z. The solid and dashed lines are for helium abundance Y = 0.255 — Z and 0.4 —Z, respectively.

Another constraint on Menv is obtained from the presupernova evolutionary model where the star evolves from the blue to the red when the mass of the hydrogen-rich envelope significantly decreases and comes back from the red to the blue if the helium abundance is sufficiently enhanced by mass loss and mixing (Saio, Kato, and Nomoto 1988). The observed N/C ratio, which is 40 times the solar ratio (Panagia 1988), is consistent with Menv 7 - 11 M0. 

Figure 2. The helium abundance profile left at the end of the main-sequence phase for the two sequences in Fig. 1.

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