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The Evolution of High-Mass Stars

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Planets, Stars and Stellar Systems

Abstract

The evolution of stars more massive than 8 M⊙ is discussed in this chapter. On the main sequence, these stars have spectral types of B2 or earlier, but depending on their mass can evolve into red supergiants, blue supergiants, Cepheids, Wolf–Rayet stars, Of stars, or luminous blue variables before ending their evolution as core collapse supernovae and neutron stars or black holes. The chapter begins with a general discussion of the energy production in the interior of a massive star as it evolves. The main fusion reactions that generate the star’s energy are listed. Some observed properties of the O and early B main-sequence stars and their evolved products are discussed including the best determinations of their masses. The computation of contemporary evolutionary tracks that include stellar rotation and magnetic fields is detailed. The equations of stellar structure including those for energy conservation, momentum transfer, mass conservation, and energy transport are listed. The discussion includes the meridional circulation in the interior of a rotating massive star and its effect on the transport of nuclear-processed material to the surface and the impact of rotation, mass loss, and metallicity on the evolutionary tracks. Recent evolutionary tracks from the Geneva group are presented. Finally the newest evolutionary tracks and the surface abundances predicted by the calculations are compared with recent observations.

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Notes

  1. 1.

    The ratio of the number of 20 M⊙ stars to one solar mass stars in the Milky Way galaxy is about 10− 5; for every 100 M⊙ star, there are a more than a million suns (Massey 2003).

  2. 2.

    Thermonuclear reaction chains are frequently written in a compact form showing the by-products in parentheses (e.g., the CNO bi-cycle is \({\mathrm{O}}^{16}(\mathrm{p},\gamma ){\mathrm{F}}^{17}{(\beta }^{+}\nu _{\mathrm{e}}){O}^{17}(\mathrm{p},\alpha ){\mathrm{N}}^{14}\), where p,  β +,  ν e, and α are the proton, positron, neutrino, and helium nucleus.)

  3. 3.

    In the jargon of nucleosynthesis, the term “ash” is used to mean the nuclear product of a series of thermonuclear reactions.

  4. 4.

    Historically such thermonuclear reactions have been called “burning” as the fusion reactions consume a nuclear “fuel,” but this is a misnomer as the term implies that the reactions are chemical rather than nuclear. Nevertheless, the term is occasionally used in this chapter.

  5. 5.

     4He comes from the photo-disintegration of 28Si.

  6. 6.

    The far ultraviolet spectra of B stars as cool as B0.5 IV (e.g., HR 1887) reveal violet asymmetries in the resonance lines from abundant ions such as Si IV, and such features most certainly contain a weak wind component.

  7. 7.

    The Eddington limit is the maximum radiative luminosity a star can reach and still be in hydrostatic equilibrium. \(\mathrm{L}_{\mathrm{Ed}} = 4\pi \mathrm{cGM}/\bar{\kappa }\), where c the speed of light, G is the gravitational constant, M the star’s mass, and \(\bar{\kappa }\) the mean opacity in the surface layers, which can include electron scattering and a temperature-dependent component due to lines from the Fe group elements.

  8. 8.

    See Porter and Rivinius (2003) for historical information on these interesting objects and Neiner (2011) for a recent update on their properties, evolution, and activity.

  9. 9.

    γ Cas, (Sechhi 1867).

  10. 10.

    The carbon lines were typically weaker than those computed with the Lanz–Hubeny model atmospheres and solar abundances. This could mean a reduced carbon abundance or emission filling from the circumstellar disk.

  11. 11.

    The duration of the corresponding WR phase is not very sensitive to the choice of this limit if it remains in the range 0.3–0.4 as mentioned by Meynet and Maeder (2003).

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Acknowledgment

GJP is grateful for the support from several NASA grants including NNX07AF89G, NNX07AH56G, and NNX10AD66G and the Women in Science and Engineering Program (WiSE) at USC. RH acknowledges support from the World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan, and from ESF-EuroGENESIS program. RH also expresses his gratitude to the groups of Professors Maeder and Meynet in Geneva and Professor Thielemann in Basel for very fruitful collaborations and support and for the use of results and figures in this chapter.

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  1. Space Sciences Center and Department of Physics and Astronomy, University of Southern California, 835 W. 37th St., SHS-274, 90089-1341, Los Angeles, CA, USA

    Prof. Geraldine J. Peters

  2. Astrophysics Group, EPSAM Institute, Keele University, Keele, UK

    Dr. Raphael Hirschi

  3. Kavli IPMU, University of Tokyo, 277-8583, Kashiwa, Chiba, Japan

    Dr. Raphael Hirschi

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  1. Department of Physics & Space Sciences, Florida Institute of Technology, University Boulevard, Melbourne, FL, USA

    Terry D. Oswalt

  2. Department of Physics and Astronomy, University of Leicester, University Road, Leicester, UK

    Martin A. Barstow

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Peters, G.J., Hirschi, R. (2013). The Evolution of High-Mass Stars. In: Oswalt, T.D., Barstow, M.A. (eds) Planets, Stars and Stellar Systems. Springer, Dordrecht. https://doi.org/10.1007/978-94-007-5615-1_9

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