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A white dwarf is a stellar core remnant composed mostly of electron-degenerate matter. A white dwarf is very dense: its mass is comparable to the Sun's, while its volume is comparable to Earth's. A white dwarf's low luminosity comes from the emission of residual thermal energy; no fusion takes place in a white dwarf.[1] The nearest known white dwarf is Sirius B, at 8.6 light years, the smaller component of the Sirius binary star. There are currently thought to be eight white dwarfs among the hundred star systems nearest the Sun.[2] The unusual faintness of white dwarfs was first recognized in 1910.[3]: 1 The name white dwarf was coined by Willem Luyten in 1922.
White dwarfs are thought to be the final evolutionary state of stars whose mass is not high enough to become a neutron star or black hole. This includes over 97% of the stars in the Milky Way.[4]: §1 After the hydrogen-fusing period of a main-sequence star of low or medium mass ends, such a star will expand to a red giant during which it fuses helium to carbon and oxygen in its core by the triple-alpha process. If a red giant has insufficient mass to generate the core temperatures required to fuse carbon (around 1 billion K), an inert mass of carbon and oxygen will build up at its center. After such a star sheds its outer layers and forms a planetary nebula, it will leave behind a core, which is the remnant white dwarf.[5] Usually, white dwarfs are composed of carbon and oxygen (CO white dwarf). If the mass of the progenitor is between 7 and 9 solar masses (M☉), the core temperature will be sufficient to fuse carbon but not neon, in which case an oxygen–neon–magnesium (ONeMg or ONe) white dwarf may form.[6] Stars of very low mass will be unable to fuse helium; hence, a helium white dwarf[7][8] may form by mass loss in binary systems.
The material in a white dwarf no longer undergoes fusion reactions, so the star has no source of energy. As a result, it cannot support itself by the heat generated by fusion against gravitational collapse, but is supported only by electron degeneracy pressure, causing it to be extremely dense. The physics of degeneracy yields a maximum mass for a non-rotating white dwarf, the Chandrasekhar limit — approximately 1.44 times M☉ — beyond which it cannot be supported by electron degeneracy pressure. A carbon–oxygen white dwarf that approaches this mass limit, typically by mass transfer from a companion star, may explode as a type Ia supernova via a process known as carbon detonation;[1][5] SN 1006 is thought to be a famous example.
A white dwarf is very hot when it forms, but because it has no source of energy, it will gradually cool as it radiates its energy away. This means that its radiation, which initially has a high color temperature, will lessen and redden with time. Over a very long time, a white dwarf will cool and its material will begin to crystallize, starting with the core. The star's low temperature means it will no longer emit significant heat or light, and it will become a cold black dwarf.[5] Because the length of time it takes for a white dwarf to reach this state is calculated to be longer than the current age of the known universe (approximately 13.8 billion years),[9] it is thought that no black dwarfs yet exist.[1][4] The oldest known white dwarfs still radiate at temperatures of a few thousand kelvins, which establishes an observational limit on the maximum possible age of the universe.[10]
- ^ a b c Johnson, J. (2007). "Extreme stars: White dwarfs & neutron stars" (Lecture notes). Astronomy 162. Ohio State University. Archived from the original on 31 March 2012. Retrieved 17 October 2011.
- ^ Henry, T.J. (1 January 2009). "The one hundred nearest star systems". Research Consortium on Nearby Stars. Archived from the original on 12 November 2007. Retrieved 21 July 2010.
- ^ Cite error: The named reference
schatzmanwas invoked but never defined (see the help page). - ^ a b Fontaine, G.; Brassard, P.; Bergeron, P. (2001). "The potential of white dwarf cosmochronology". Publications of the Astronomical Society of the Pacific. 113 (782): 409–435. Bibcode:2001PASP..113..409F. doi:10.1086/319535.
- ^ a b c Richmond, M. "Late stages of evolution for low-mass stars". Lecture notes, Physics 230. Rochester Institute of Technology. Archived from the original on 4 September 2017. Retrieved 3 May 2007.
- ^ Werner, K.; Hammer, N.J.; Nagel, T.; Rauch, T.; Dreizler, S. (2005). On possible oxygen / neon white dwarfs: H1504+65 and the white dwarf donors in ultracompact X-ray binaries. 14th European Workshop on White Dwarfs. Vol. 334. p. 165. arXiv:astro-ph/0410690. Bibcode:2005ASPC..334..165W.
- ^ Liebert, James; Bergeron, P.; Eisenstein, D.; Harris, H. C.; Kleinman, S. J.; Nitta, A.; Krzesinski, J. (2004). "A helium white dwarf of extremely low mass". The Astrophysical Journal. 606 (2): L147. arXiv:astro-ph/0404291. Bibcode:2004ApJ...606L.147L. doi:10.1086/421462. S2CID 118894713.
- ^ "Cosmic weight loss: The lowest mass white dwarf" (Press release). Harvard-Smithsonian Center for Astrophysics. 17 April 2007. Archived from the original on 22 April 2007. Retrieved 20 April 2007.
- ^ Spergel, D.N.; Bean, R.; Doré, O.; Nolta, M.R.; Bennett, C.L.; Dunkley, J.; et al. (2007). "Wilkinson Microwave Anisotropy Probe (WMAP) three year results: Implications for cosmology". The Astrophysical Journal Supplement Series. 170 (2): 377–408. arXiv:astro-ph/0603449. Bibcode:2007ApJS..170..377S. doi:10.1086/513700. S2CID 1386346.
- ^ §3, Heger, A.; Fryer, C.L.; Woosley, S.E.; Langer, N.; Hartmann, D.H. (2003). "How massive single stars end their life". Astrophysical Journal. 591 (1): 288–300. arXiv:astro-ph/0212469. Bibcode:2003ApJ...591..288H. doi:10.1086/375341. S2CID 59065632.