Stellar spectral types

When Payne began her work in the 1920s, stellar spectroscopy was a very active area of research. Numerous elemental and molecular lines had been identified in stellar spectra. The lines observed in each star varied with the inferred temperature of the star, which was understood to mean that the elemental abundances varied with temperature. This body of data was the basis for the spectral typing of stars ( , B, A, F, G, , M, L). However, the power source for stars was not understood and it was not clear why the composition of a star should be related to its temperature. In the 1920s, it was also widely believed that the Sun had the same composition as the Earth models considered the Earth to have formed from the outer layers of the Sun. Payne used the new guantum mechanical understanding of atomic structure to show how and why the spectral lines of the different elements varied as a function of stellar spectral type. She demonstrated how the temperature of the stellar surface controls the spectral lines that are observed. Her analysis led to the conclusion that the chemical... 

This work aims at testing the suggestion of [5] that stellar rotation is faster at lower metallicity by direct measurements, especially in the LMC and SMC, on stars with —3.34 < My < —2.17, i.e. spectral types B0-B6 or masses from 6.7 to 14 M0. This work is complementary to that of [4], which deals with slightly more massive stars. The results are shown on Fig. 1 and commented in the caption. There is an excess of slow rotators in the Galaxy relative to the MCs, but the v sin i distributions of the LMC and the SMC are surprisingly similar. 

Fig. 3.3. Theoretical Hertzsprung-Russell diagram. The right-hand scale shows in absolute bolometric magnitude what the left-hand scale expresses as the logarithm of the intrinsic luminosity in units of the solar intrinsic luminosity (Lq = 4 x 10 erg s ). On the horizontal axis, the logarithm of the effective temperature, i.e. the temperature of the equivalent blackbody, is put into correspondence with the spectral type of the star, as determined by the observer. This temperature-luminosity diagram shows the lifelines of the stars as strands combed out like hair across the graph. With a suitable interpretation, i.e. viewed through the explanatory machinery of nuclear physics, it opens the way to an understanding of stellar evolution and its twin science of nucleosynthesis. (Courtesy of Andre Maeder and co-workers.)...

However, the hot evolved stars do not form a homogeneous group. The most prominent division is into Central Stars of Planetary Nebulae differences exist in spectral type and chemical composition (as deduced from medium and high resolution spectra). Consequently, a unique progenitor for all kinds of hot evolved stars appears to be very unlikely and accurate stellar parameters are needed to compare these stars to predictions of stellar evolution theory. 

This paper summarizes the results of analyses of highly evolved stars with spectral type B or hotter, namely sdB, sdOB and sdO types, CSPN and extremely helium-rich stars. It does not consider white dwarfs since their chemical surface composition is apparently governed by diffusion processes and accretion of interstellar material (Wesemael, 1979 Vauclair et al., 1979 Wesemael and Truran, 1982) and is not linked to their past evolution. Section 2 deals with the positions of the hot evolved stars in the (log Te -log g) plane and their helium to hydrogen ratios. Metal abundances are considered in section 3 and comparisons of stellar evolution calculations with the available data are performed in section 4. 

F star gap The abundance of Li in recent stars has created a very significant puzzle for the understanding of stellar structure. Although Li/H = 2 x io-9 in most recently formed stars, it is about 100 times less in recent spectral type F stars. These Population... 

Figure 2.3 Dust production and gas mass return rate by different stellar types in solar masses per year and kpc-2 in the galaxy at the solar cycle. Stars produce mainly silicate or carbon dust only in some cases is a different kind of dust material formed, probably iron or some iron alloy (peculiar dust). Many additional dust components with much smaller abundance are formed in most cases (Data from Tielens 1999 Zhukovska el al. 2008). Abbreviations of stellar types AGB = asymptotic giant branch stars of spectral types M, S, or C OB = massive stars of spectral types O and B on or close to the main sequence RGB = massive stars on the red giant branch LBV = luminous blue variables WCL = Wolf-Rayet stars from the lower temperature range Novae = mass ejecta from novae SN = mass ejecta from supemovae.

Table 2.3 Observed properties of some young stellar objects and their accretion disks spectral type, effective temperature Teg, luminosity L, estimated stellar mass M, stellar radius Rt, accretion rate M, disk radius T isk as observed by dust emission, inclination of disk with respect to sight line, and disk mass Mdisk estimated from submillimeter dust emission...

Table 8.1 Typical stellar and disk parameters for young stars. The columns show typical spectral types at i Myr, effective temperatures, luminosity, disk masses, and mass accretion rates.

The presence of weak GMF for 21 star with vigorous convection is detected (F9-M3 spectral types and I-V luminosity classes). Furthermore, the variation of global nonaxisymmetric magnetic field ge neral magnetic field - as a function of the stellar rotation has been first determined for two solar-like stars other than the Sun Boo A and 61 Cyg A. 

Column 14 The spectral type of the source. Spectrum type C is a cool dust spectrum (typical Td 30 K), peaking beyond 100 m colour correction factors are of order 1.00, 0.95, 0.99, 1.00. Type W is a warm dust spectrum (typical Td 70 K), peaking between 12 and 100 m colour correction factors are of order 1.03,1.00, 1.00, 1.04. Type S is a stellar spectrum (typical blackbody of about 5000 K) colour correction factors are of order 1.43,1.40,1.32,1.09. A semi-colon indicates that the infrared spectrum is uncertain. Actual, colour-corrected fiux densities can be calculated from the quoted ones by dividing the latter by these colour correction factors (see IRAS, 1989a). 

Our multi-level carbon model atom is adapted from D. Kiselman (private communication), with improved atomic data and better sampling of some absorption lines. The statistical equilibrium code MULTI (Carlsson 1986), together with ID MARCS stellar model atmospheres for a grid of 168 late-type stars with varying Tefj, log g, [Fe/H] and [C/Fe], were used in all Cl non-LTE spectral line formation calculations, to solve radiative-transfer and rate equations and to find the non-LTE solution for the multi-level atom. We put particular attention in the study of the permitted Cl lines around 9100 A, used by Akerman et al. (2004). 

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