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The first attempt to classify white dwarf spectra appears to have been by Gerard P. Kuiper in 1941,5271 and various classification schemes have been proposed and used since then.7273 The system currently in use was introduced by Edward M. Sion and his coauthors in 1983 and has been subsequently revised several times. It classifies a spectrum by a symbol which consists of an initial D, a letter describing the primary feature of the spectrum followed by an optional sequence of letters describing secondary features of the spectrum (as shown in the table to the right), and a temperature index number, computed by dividing 50,400 K by the effective temperature. For example:

  • A white dwarf with only He I lines in its spectrum and an effective temperature of 15,000 K could be given the classification of DB3, or, if warranted by the precision of the temperature measurement, DB3.5.
  • A white dwarf with a polarized magnetic field, an effective temperature of 17,000 K, and a spectrum dominated by He I lines which also had hydrogen features could be given the classification of DBAP3.

The symbols ? and : may also be used if the correct classification is uncertain.5224

White dwarfs whose primary spectral classification is DA have hydrogen-dominated atmospheres. They make up the majority (approximately three-quarters) of all observed white dwarfs.54 A small fraction (roughly 0.1 percent) have carbon-dominated atmospheres, the hot (above 15,000 K) DQ class.74 The classifiable remainder (DB, DC, DO, DZ, and cool DQ) have helium-dominated atmospheres. Assuming that carbon and metals are not present, which spectral classification is seen depends on the effective temperature. Between approximately 100,000 K to 45,000 K, the spectrum will be classified DO, dominated by singly ionized helium. From 30,000 K to 12,000 K, the spectrum will be DB, showing neutral helium lines, and below about 12,000 K, the spectrum will be featureless and classified DC.70,§ 2.454 The reason for the absence of white dwarfs with helium-dominated atmospheres and effective temperatures between 30,000 K and 45,000 K, called the DB gap, is not clear. It is suspected to be due to competing atmospheric evolutionary processes, such as gravitational separation and convective mixing.54

Magnetic field

Magnetic fields in white dwarfs with a strength at the surface of ~1 million gauss (100 teslas) were predicted by P. M. S. Blackett in 1947 as a consequence of a physical law he had proposed which stated that an uncharged, rotating body should generate a magnetic field proportional to its angular momentum.75 This putative law, sometimes called the Blackett effect, was never generally accepted, and by the 1950s even Blackett felt it had been refuted.76, 39-43 In the 1960s, it was proposed that white dwarfs might have magnetic fields because of conservation of total surface magnetic flux during the evolution of a non-degenerate star to a white dwarf. A surface magnetic field of ~100 gauss (0.01 T) in the progenitor star would thus become a surface magnetic field of ~100•1002=1 million gauss (100 T) once the star's radius had shrunk by a factor of 100.69, §8;77, 484 The first magnetic white dwarf to be observed was GJ 742, which was detected to have a magnetic field in 1970 by its emission of circularly polarized light.78 It is thought to have a surface field of approximately 300 million gauss (30 kT).69, §8 Since then magnetic fields have been discovered in well over 100 white dwarfs, ranging from 2×103 to 109 gauss (0.2 T to 100 kT). Only a small number of white dwarfs have been examined for fields, and it has been estimated that at least 10 percent of white dwarfs have fields in excess of 1 million gauss (100 T).7980

Variability

DAV (GCVS: ZZA)DA spectral type, having only hydrogen absorption lines in its spectrumDBV (GCVS: ZZB)DB spectral type, having only helium absorption lines in its spectrumGW Vir (GCVS: ZZO)Atmosphere mostly C, He and O;
may be divided into DOV and PNNV stars Types of pulsating white dwarf8182, §1.1, 1.2.See also: Cataclysmic variables

Early calculations suggested that there might be white dwarfs whose luminosity varied with a period of around 10 seconds, but searches in the 1960s failed to observe this.69, § 7.1.1;83 The first variable white dwarf found was HL Tau 76; in 1965 and 1966, Arlo U. Landolt observed it to vary with a period of approximately 12.5 minutes.84 The reason for this period being longer than predicted is that the variability of HL Tau 76, like that of the other pulsating variable white dwarfs known, arises from non-radial gravity wave pulsations.69, § 7. Known types of pulsating white dwarf include the DAV, or ZZ Ceti, stars, including HL Tau 76, with hydrogen-dominated atmospheres and the spectral type DA;69, 891, 895 DBV, or V777 Her, stars, with helium-dominated atmospheres and the spectral type DB;54, 3525 and GW Vir stars (sometimes subdivided into DOV and PNNV stars), with atmospheres dominated by helium, carbon, and oxygen.82,§1.1, 1.2;85,§1. GW Vir stars are not, strictly speaking, white dwarfs, but are stars which are in a position on the Hertzsprung-Russell diagram between the asymptotic giant branch and the white dwarf region. They may be called pre-white dwarfs.82, § 1.1;86 These variables all exhibit small (1 percent-30 percent) variations in light output, arising from a superposition of vibrational modes with periods of hundreds to thousands of seconds. Observation of these variations gives asteroseismological evidence about the interiors of white dwarfs.87

Formation

White dwarfs are thought to represent the end point of stellar evolution for main-sequence stars with masses from about 0.07 to 10 solar masses.887 The composition of the white dwarf produced will differ depending on the initial mass of the star.

Stars with very low mass

If the mass of a main-sequence star is lower than approximately half a solar mass, it will never become hot enough to fuse helium at its core. It is thought that, over a lifespan exceeding the age (~13.7 billion years)6 of the Universe, such a star will eventually burn all its hydrogen and end its evolution as a helium white dwarf composed chiefly of helium-4 nuclei. Owing to the time this process takes, it is not thought to be the origin of observed helium white dwarfs. Rather, they are thought to be the product of mass loss in binary systems348990911 or mass loss due to a large planetary companion.92

Stars with low to medium mass

If the mass of a main-sequence star is between approximately 0.5 and 8 solar masses, its core will become sufficiently hot to fuse helium into carbon and oxygen via the triple-alpha process, but it will never become sufficiently hot to fuse carbon into neon. Near the end of the period in which it undergoes fusion reactions, such a star will have a carbon-oxygen core which does not undergo fusion reactions, surrounded by an inner helium-burning shell and an outer hydrogen-burning shell. On the Hertzsprung-Russell diagram, it will be found on the asymptotic giant branch. It will then expel most of its outer material, creating a planetary nebula, until only the carbon-oxygen core is left. This process is responsible for the carbon-oxygen white dwarfs which form the vast majority of observed white dwarfs.899394

Stars with medium to high mass

If a star is sufficiently massive, its core will eventually become sufficiently hot to fuse carbon to neon, and then to fuse neon to iron. Such a star will not become a white dwarf as the mass of its central, non-fusing, core, supported by electron degeneracy pressure, will eventually exceed the largest possible mass supportable by degeneracy pressure. At this point the core of the star will collapse and it will explode in a core-collapse supernova which will leave behind a remnant neutron star, black hole, or possibly a more exotic form of compact star.8895 Some main-sequence stars, of perhaps 8 to 10 solar masses, although sufficiently massive to fuse carbon to neon and magnesium, may be insufficiently massive to fuse neon. Such a star may leave a remnant white dwarf composed chiefly of oxygen, neon, and magnesium, provided that its core does not collapse, and provided that fusion does not proceed so violently as to blow apart the star in a supernova.9697 Although some isolated white dwarfs have been identified which may be of this type, most evidence for the existence of such stars comes from the novae called ONeMg or neon novae. The spectra of these novae exhibit abundances of neon, magnesium, and other intermediate-mass elements which appear to be only explicable by the accretion of material onto an oxygen-neon-magnesium white dwarf.29899

Fate

A white dwarf is stable once formed and will continue to cool almost indefinitely; eventually, it will become a black white dwarf, also called a black dwarf. Assuming that the Universe continues to expand, it is thought that in 1019 to 1020 years, the galaxies will evaporate as their stars escape into intergalactic space.100, §IIIA. White dwarfs should generally survive this, although an occasional collision between white dwarfs may produce a new fusing star or a super-Chandrasekhar mass white dwarf which will explode in a type Ia supernova.100, §IIIC, IV. The subsequent lifetime of white dwarfs is thought to be on the order of the lifetime of the proton, known to be at least 1032 years. Some simple grand unified theories predict a proton lifetime of no more than 1049 years. If these theories are not valid, the proton may decay by more complicated nuclear processes, or by quantum gravitational processes involving a virtual black hole; in these cases, the lifetime is estimated to be no more than 10200 years. If protons do decay, the mass of a white dwarf will decrease very slowly with time as its nuclei decay, until it loses so much mass as to become a nondegenerate lump of matter, and finally disappears completely.100, §IV.

Stellar system

A white dwarf's stellar and planetary system is inherited from its progenitor star and may interact with the white dwarf in various ways. Infrared spectroscopic observations made by NASA's Spitzer Space Telescope of the central star of the Helix Nebula suggest the presence of a dust cloud, which may be caused by cometary collisions. It is possible that infalling material from this may cause X-ray emission from the central star.101102 Similarly, observations made in 2004 indicated the presence of a dust cloud around the young white dwarf star G29-38 (estimated to have formed from its AGB progenitor about 500 million years ago), which may have been created by tidal disruption of a comet passing close to the white dwarf.103 If a white dwarf is in a binary system with a stellar companion, a variety of phenomena may occur, including novae and Type Ia supernovae. It may also be a super-soft x-ray source if it is able to take material from its companion fast enough to sustain fusion on its surface.

Type Ia supernovae

Multiwavelength X-ray image of SN 1572 or Tycho's Nova, the remnant of a Type Ia supernova.

The mass of an isolated, nonrotating white dwarf cannot exceed the Chandrasekhar limit of ~1.4 solar masses. (This limit may increase if the white dwarf is rotating rapidly and nonuniformly.)104 White dwarfs in binary systems, however, can accrete material from a companion star, increasing both their mass and their density. As their mass approaches the Chandrasekhar limit, this could theoretically lead to either the explosive ignition of fusion in the white dwarf or its collapse into a neutron star.44

Accretion provides the currently favored mechanism, the single-degenerate model, for type Ia supernovae. In this model, a carbon-oxygen white dwarf accretes material from a companion star,45, p. 14. increasing its mass and compressing its core. It is believed that compressional heating of the core leads to ignition of carbon fusion as the mass approaches the Chandrasekhar limit.45 Because the white dwarf is supported against gravity by quantum degeneracy pressure instead of by thermal pressure, adding heat to the star's interior increases its temperature but not its pressure, so the white dwarf does not expand and cool in response. Rather, the increased temperature accelerates the rate of the fusion reaction, in a runaway process that feeds on itself. The thermonuclear flame consumes much of the white dwarf in a few seconds, causing a type Ia supernova explosion that obliterates the star.545105 In another possible mechanism for type Ia supernovae, the double-degenerate model, two carbon-oxygen white dwarfs in a binary system merge, creating an object with mass greater than the Chandrasekhar limit in which carbon fusion is then ignited.45, 14.

Cataclysmic variables

When accretion of material does not push a white dwarf close to the Chandrasekhar limit, accreted hydrogen-rich material on the surface may still ignite in a thermonuclear explosion. Since the white dwarf's core remains intact, these surface explosions can be repeated as long as accretion continues. This weaker kind of repetitive cataclysmic phenomenon is called a (classical) nova. Astronomers have also observed dwarf novae, which have smaller, more frequent luminosity peaks than classical novae. These are thought to not be caused by fusion but rather by the release of gravitational potential energy during accretion. In general, binary systems with a white dwarf accreting matter from a stellar companion are called cataclysmic variables. As well as novae and dwarf novae, several other classes of these variables are known.545106107 Both fusion- and accretion-powered cataclysmic variables have been observed to be X-ray sources.107

See also

  • Brown dwarf
  • Hertzsprung-Russell diagram
  • Neutron star
  • Planetary nebula
  • Red dwarf
  • Supernova

Notes

  1. 1.0 1.1 1.2 1.3 1.4 Michael Richmond, Late stages of evolution for low-mass stars. Lecture notes, Physics 230, Rochester Institute of Technology. Retrieved November 7, 2008.
  2. 2.0 2.1 K. Werner, N.J. Hammer, T. Nagel, T. Rauch, and S. Dreizler. 2004. On Possible Oxygen/Neon White Dwarfs: H1504+65 and the White Dwarf Donors in Ultracompact X-ray Binaries. in D. Koester and S. Moehler. 2005. 14th European Workshop on White Dwarfs; Proceedings of a meeting held at Kiel, July 19-23, 2004. (San Francisco, CA: Astronomical Society of the Pacific. ISBN 9781583811979), 165. Retrieved November 7, 2008.
  3. 3.0 3.1 James Liebert, P. Bergeron, Daniel Eisenstein, H.C. Harris, S.J. Kleinman, Atsuko Nitta, and Jurek Krzesinski. 2004. A Helium White Dwarf of Extremely Low Mass. The Astrophysical Journal. 606(2):L147-L149. Retrieved November 7, 2008.
  4. 4.0 4.1 Press release. 2007. Cosmic weight loss: The lowest mass white dwarf. Harvard-Smithsonian Center for Astrophysics. Retrieved November 7, 2008.
  5. 5.0 5.1 5.2 5.3 5.4 5.5 5.6 Jennifer Johnson, Extreme Stars: White Dwarfs & Neutron Stars. Lecture notes, Astronomy 162. Ohio State University. Retrieved November 7, 2008.
  6. 6.0 6.1 D.N. Spergel, et al. 2007. Wilkinson Microwave Anisotropy Probe (WMAP) Three Year Results: Implications for Cosmology. ApJS. 170:377. Retrieved November 7, 2008.
  7. 7.0 7.1 7.2 G. Fontaine, P. Brassard, and P. Bergeron. 2001. The Potential of White Dwarf Cosmochronology. Publications of the Astronomical Society of the Pacific 113 (782):409-435. Retrieved November 7, 2008.
  8. 8.0 8.1 8.2 8.3 E. Schatzman, 1958. White Dwarfs. (Amsterdam, NL: North-Holland.)
  9. 9.0 9.1 9.2 9.3 J.B. Holberg, 2005. How Degenerate Stars Came to be Known as White Dwarfs. Bulletin of the American Astronomical Society 37:1503. Retrieved November 7, 2008.
  10. ↑ William Herschel, (1785). Catalogue of Double Stars. Philosophical Transactions of the Royal Society of London 75:40-126. Retrieved November 7, 2008.
  11. ↑ W.H. van den Bos. (1926). The orbit and the masses of 40 Eridani BC. Bulletin of the Astronomical Institutes of the Netherlands 3(98): 128-132. Retrieved November 7, 2008.
  12. ↑ W.D. Heintz, 1974. Astrometric study of four visual binaries. Astronomical Journal 79(7):819-825. Retrieved November 7, 2008.
  13. ↑ Walter S. Adams, (1914) An A-Type Star of Very Low Luminosity. Publications of the Astronomical Society of the Pacific 26(155):198. Retrieved November 7, 2008.
  14. 14.0 14.1 F.W. Bessel, J.F.W. Herschel. (1844). On the Variations of the Proper Motions of Procyon and Sirius. Monthly Notices of the Royal Astronomical Society 6:136-141. Retrieved November 7, 2008.
  15. 15.0 15.1 Camille Flammarion, (1877). The Companion of Sirius. The Astronomical Register 15(176):186-189. Retrieved November 7, 2008.
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  19. ↑ Willem J. Luyten, 1922. Note on Some Faint Early Type Stars with Large Proper Motions. Publications of the Astronomical Society of the Pacific 34(197):54-55. Retrieved November 7, 2008.
  20. ↑ Luyten, Willem J. 1922. Additional Note on Faint Early-Type Stars with Large Proper-Motions. Publications of the Astronomical Society of the Pacific 34(198):132. Retrieved November 7, 2008.
  21. ↑ Willem J. Luyten, 1922. Third Note on Faint Early Type Stars with Large Proper Motion. Publications of the Astronomical Society of the Pacific 34(202):356-357. Retrieved November 7, 2008.
  22. 22.0 22.1 22.2 A.S. Eddington, 1924. On the relation between the masses and luminosities of the stars. Monthly Notices of the Royal Astronomical Society. 84:308-332. Retrieved November 7, 2008.
  23. ↑ Luyten, Willem J. 1950. The search for white dwarfs. Astronomical Journal. 55(1183):86-89. Retrieved November 7, 2008.
  24. 24.0 24.1 24.2 24.3 George P. McCook and Edward M. Sion. 1999. A Catalog of Spectroscopically Identified White Dwarfs. The Astrophysical Journal Supplement Series 121(1):1-130. Retrieved November 7, 2008.
  25. 25.0 25.1 Daniel J. Eisenstein, et al. 2006. A Catalog of Spectroscopically Confirmed White Dwarfs from the Sloan Digital Sky Survey Data Release 4. The Astrophysical Journal Supplement Series 167(1):40-58. Retrieved November 7, 2008.
  26. ↑ Mukremin Kulic, Carlos Allende Prieto, Warren R. Brown, and D. Koester. 2007. The Lowest Mass White Dwarf. The Astrophysical Journal 660(2):1451-1461. Retrieved November 7, 2008.
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  39. ↑ Wilhelm Anderson, 1929. Über die Grenzdichte der Materie und der Energie. Zeitschrift für Physik. 56(11-12):851-856.
  40. 40.0 40.1 Edmund C. Stoner, 1930. The Equilibrium of Dense Stars. Philosophical Magazine (7th series) 9:944-963.
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  56. ↑ P. Bergeron, Maria Teresa Ruiz, and S.K. Leggett. The Chemical Evolution of Cool White Dwarfs and the Age of the Local Galactic Disk. The Astrophysical Journal Supplement Series 108(1):339-387. Retrieved November 7, 2008.
  57. 57.0 57.1 T.S. Metcalfe, M.H. Montgomery, and A. Kanaan. 2004. Testing White Dwarf Crystallization Theory with Asteroseismology of the Massive Pulsating DA Star BPM 37093. The Astrophysical Journal 605():L133-L136. Retrieved November 7, 2008.
  58. ↑ J.L. Barrat, J.P. Hansen, and R. Mochkovitch. 1988. Crystallization of carbon-oxygen mixtures in white dwarfs. Astronomy and Astrophysics 199(1-2):L15-L18. Retrieved November 7, 2008.
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  60. ↑ David Whitehouse, 2004. Diamond star thrills astronomers. BBC News. Retrieved November 7, 2008.
  61. ↑ Press release. Harvard-Smithsonian Center for Astrophysics. Retrieved November 7, 2008.
  62. ↑ A. Kanaan, et al. 2004. Whole Earth Telescope observations of BPM 37093: a seismological test of crystallization theory in white dwarfs. arXiv:astro-ph/0411199v1. Retrieved November 7, 2008.
  63. ↑ P. Brassard, and G. Fontaine. 2005. Asteroseismology of the Crystallized ZZ Ceti Star BPM 37093: A Different View. The Astrophysical Journal 622(1):572-576. Retrieved November 7, 2008.
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  67. ↑ James S. Trefil, (1983) 2004. The Moment of Creation: Big Bang Physics from Before the First Millisecond to the Present Universe. (reprint ed. Mineola, NY: Dover Publications. ISBN 0486438139).
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  70. 70.0 70.1 Steven D. Kawaler, "White Dwarf Stars" in S.D. Kawaler, I. Nov

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