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<document>
<title>White dwarf</title>
<authors>Wikipedia</authors>
<publisher>wikipedia.org</publisher>
<publishedplace>network</publishedplace>
<publishedyear>2010</publishedyear>
<edition></edition>
<category>Star types</category>
<section id="7719" fatherid="0" title="" lv="" priority="-1" level="1">
A white dwarf, also called a degenerate dwarf, is a small star composed mostly of electron-degenerate matter.  They are very dense; a white dwarf's mass is comparable to that of the Sun and its volume is comparable to that of the Earth.  Its faint luminosity comes from the emission of stored thermal energy.  White dwarfs comprise roughly 6% of all known stars in the solar neighborhood.The One Hundred Nearest Star Systems, Todd J. Henry, RECONS, April 11, 2007.  Accessed on line May 4, 2007.  The unusual faintness of white dwarfs was first recognized in 1910 by Henry Norris Russell, Edward Charles Pickering, and Williamina Fleming;, p. 1 the name white dwarf was coined by Willem Luyten in 1922.
White dwarfs are thought to be the final evolutionary state of all stars whose mass is not too high—over 97% of the stars in our galaxy., §1.   After the hydrogen–fusing lifetime of a main-sequence star of low or medium mass ends, it will expand to a red giant which 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, an inert mass of carbon and oxygen will build up at its center.  After shedding its outer layers to form a planetary nebula, it will leave behind this core, which forms the remnant white dwarf.Late stages of evolution for low-mass stars, Michael Richmond, lecture notes, Physics 230, Rochester Institute of Technology.  Accessed on line May 3, 2007.  Usually, therefore, white dwarfs are composed of carbon and oxygen.  It is also possible that core temperatures suffice to fuse carbon but not neon, in which case an oxygen-neon–magnesium white dwarf may be formed.On Possible Oxygen/Neon White Dwarfs: H1504+65 and the White Dwarf Donors in Ultracompact X-ray Binaries, K. Werner, N. J. Hammer, T. Nagel, T. Rauch, and S. Dreizler, pp. 165 ff. in 14th European Workshop on White Dwarfs; Proceedings of a meeting held at Kiel, July 19–23, 2004, edited by D. Koester and S. Moehler, San Francisco: Astronomical Society of the Pacific, 2005.  Also, some helium white dwarfsA Helium White Dwarf of Extremely Low Mass, James Liebert, P. Bergeron, Daniel Eisenstein, H.C. Harris, S.J. Kleinman, Atsuko Nitta, and Jurek Krzesinski, The Astrophysical Journal 606, #2 (May 2004), pp. L147-L149.  Accessed on line March 5, 2007.Cosmic weight loss: The lowest mass white dwarf, press release, Harvard-Smithsonian Center for Astrophysics, April 17, 2007. appear to have been formed 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, nor is it supported against gravitational collapse by the heat generated by fusion.  It is supported only by electron degeneracy pressure, causing it to be extremely dense.  The physics of degeneracy yields a maximum mass for a nonrotating white dwarf, the Chandrasekhar limit-approximately 1.4 solar masses-beyond which it cannot be supported by 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.Extreme Stars: White Dwarfs -anddocsymbol- Neutron Stars, Jennifer Johnson, lecture notes, Astronomy 162, Ohio State University.  Accessed on line May 3, 2007. (SN 1006 is thought to be a famous example.)
A white dwarf is very hot when it is formed but since it has no source of energy, it will gradually radiate away its energy and cool down.  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 to temperatures at which it will no longer be visible, and become a cold black dwarf.  However, since no white dwarf can be older than the age of the Universe (approximately 13.7 billion years),Wilkinson Microwave Anisotropy Probe (WMAP) Three Year Results: Implications for Cosmology, D. N. Spergel, R. Bean, O. Doré, M. R. Nolta, C. L. Bennett, J. Dunkley, G. Hinshaw, N. Jarosik, E. Komatsu, L. Page, H. V. Peiris, L. Verde, M. Halpern, R. S. Hill, A. Kogut, M. Limon, S. S. Meyer, N. Odegard, G. S. Tucker, J. L. Weiland, E. Wollack, and E. L. W, arXiv:astro-ph/0603449v2, February 27, 2007. even the oldest white dwarfs still radiate at temperatures of a few thousand kelvins, and no black dwarfs are thought to exist yet.The Potential of White Dwarf Cosmochronology, G. Fontaine, P. Brassard, and P. Bergeron, Publications of the Astronomical Society of the Pacific 113, #782 (April 2001), pp. 409-435.

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<section id="7720" fatherid="0" title="Discovery" lv="" priority="2" level="1">
The first white dwarf discovered was in the triple star system of 40 Eridani, which contains the relatively b main sequence star 40 Eridani A, orbited at a distance by the closer binary system of the white dwarf 40 Eridani B and the main sequence red dwarf 40 Eridani C.  The pair 40 Eridani B/C was discovered by Friedrich Wilhelm Herschel on January 31, 1783;Catalogue of Double Stars, William Herschel, Philosophical Transactions of the Royal Society of London 75 (1785), pp. 40-126, p. 73 it was again observed by Friedrich Georg Wilhelm Struve in 1825 and by Otto Wilhelm von Struve in 1851.The orbit and the masses of 40 Eridani BC, W. H. van den Bos, Bulletin of the Astronomical Institutes of the Netherlands 3, #98 (July 8, 1926), pp. 128-132.Astrometric study of four visual binaries, W. D. Heintz, Astronomical Journal 79, #7 (July 1974), pp. 819-825.  In 1910, it was discovered by Henry Norris Russell, Edward Charles Pickering and Williamina Fleming that despite being a dim star, 40 Eridani B was of spectral type A, or white.How Degenerate Stars Came to be Known as White Dwarfs, J. B. Holberg, Bulletin of the American Astronomical Society 37 (December 2005),  p. 1503.  In 1939, Russell looked back on the discovery:White Dwarfs, E. Schatzman, Amsterdam: North-Holland, 1958., p. 1
I was visiting my friend and generous benefactor, Prof. Edward C. Pickering.  With characteristic kindness, he had volunteered to have the spectra observed for all the stars-including comparison stars-which had been observed in the observations for stellar parallax which Hinks and I made at Cambridge, and I discussed.  This piece of apparently routine work proved very fruitful-it led to the discovery that all the stars of very faint absolute magnitude were of spectral class M.  In conversation on this subject (as I recall it), I asked Pickering about certain other faint stars, not on my list, mentioning in particular 40 Eridani B.  Characteristically, he sent a note to the Observatory office and before long the answer came (I think from Mrs Fleming) that the spectrum of this star was A.  I knew enough about it, even in these paleozoic days, to realize at once that there was an extreme inconsistency between what we would then have called "possible" values of the surface bness and density.  I must have shown that I was not only puzzled but crestfallen, at this exception to what looked like a very pretty rule of stellar characteristics; but Pickering smiled upon me, and said: "It is just these exceptions that lead to an advance in our knowledge", and so the white dwarfs entered the realm of study!
The spectral type of 40 Eridani B was officially described in 1914 by Walter Adams.An A-Type Star of Very Low Luminosity, Walter S. Adams, Publications of the Astronomical Society of the Pacific 26, #155 (October 1914), p. 198.
The companion of Sirius, Sirius B, was next to be discovered.  During the nineteenth century, positional measurements of some stars became precise enough to measure small changes in their location. Friedrich Bessel used just such precise measurements to determine that the stars Sirius (α Canis Majoris) and Procyon (α Canis Minoris) were changing their positions.  In 1844 he predicted that both stars had unseen companions:On the Variations of the Proper Motions of Procyon and Sirius, F. W. Bessel, communicated by J. F. W. Herschel, Monthly Notices of the Royal Astronomical Society 6 (December 1844), pp. 136-141.
If we were to regard Sirius and Procyon as double stars, the change of their motions would not surprise us; we should acknowledge them as necessary, and have only to investigate their amount by observation.  But light is no real property of mass.  The existence of numberless visible stars can prove nothing against the existence of numberless invisible ones.
Bessel roughly estimated the period of the companion of Sirius to be about half a century; C. H. F. Peters computed an orbit for it in 1851.The Companion of Sirius, Camille Flammarion, The Astronomical Register 15, #176 (August 1877), pp. 186-189.  It was not until January 31, 1862 that Alvan Graham Clark observed a previously unseen star close to Sirius, later identified as the predicted companion.  Walter Adams announced in 1915 that he had found the spectrum of Sirius B to be similar to that of Sirius.The Spectrum of the Companion of Sirius, W. S. Adams, Publications of the Astronomical Society of the Pacific 27, #161 (December 1915), pp. 236-237.
In 1917, Adriaan Van Maanen discovered Van Maanen's Star, an isolated white dwarf.Two Faint Stars with Large Proper Motion, A. van Maanen, Publications of the Astronomical Society of the Pacific 29, #172 (December 1917), pp. 258-259.   These three white dwarfs, the first discovered, are the so-called classical white dwarfs., p. 2  Eventually, many faint white stars were found which had high proper motion, indicating that they could be suspected to be low-luminosity stars close to the Earth, and hence white dwarfs.  Willem Luyten appears to have been the first to use the term white dwarf when he examined this class of stars in 1922;The Mean Parallax of Early-Type Stars of Determined Proper Motion and Apparent Magnitude, Willem J. Luyten, Publications of the Astronomical Society of the Pacific 34, #199 (June 1922), pp. 156-160.Note on Some Faint Early Type Stars with Large Proper Motions, Willem J. Luyten, Publications of the Astronomical Society of the Pacific 34, #197 (February 1922), pp. 54-55.Additional Note on Faint Early-Type Stars with Large Proper-Motions, Willem J. Luyten, Publications of the Astronomical Society of the Pacific 34, #198 (April 1922), p. 132.Third Note on Faint Early Type Stars with Large Proper Motion, Willem J. Luyten, Publications of the Astronomical Society of the Pacific 34, #202 (December 1922),  pp. 356-357.  the term was later popularized by Arthur Stanley Eddington.  Despite these suspicions, the first non-classical white dwarf was not definitely identified until the 1930s.  18 white dwarfs had been discovered by 1939., p. 3  Luyten and others continued to search for white dwarfs in the 1940s.  By 1950, over a hundred were known,The search for white dwarfs, W. J. Luyten, Astronomical Journal 55, #1183 (April 1950), pp. 86-89. and by 1999, over 2,000 were known.A Catalog of Spectroscopically Identified White Dwarfs, George P. McCook and Edward M. Sion, The Astrophysical Journal Supplement Series 121, #1 (March 1999), pp. 1-130.  Since then the Sloan Digital Sky Survey has found over 9,000 white dwarfs, mostly new.A Catalog of Spectroscopically Confirmed White Dwarfs from the Sloan Digital Sky Survey Data Release 4, Daniel J. Eisenstein, James Liebert, Hugh C. Harris, S. J. Kleinman, Atsuko Nitta, Nicole Silvestri, Scott A. Anderson, J. C. Barentine, Howard J. Brewington, J. Brinkmann, Michael Harvanek, Jurek Krzesiński, Eric H. Neilsen, Jr., Dan Long, Donald P. Schneider, and Stephanie A. Snedden, The Astrophysical Journal Supplement Series 167, #1 (November 2006), pp. 40-58.

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<section id="7721" fatherid="0" title="Composition and structure" lv="" priority="3" level="1">
<section id="7722" fatherid="7721" title="" lv="" priority="-1" level="2">
Although white dwarfs are known with estimated masses as low as 0.17The Lowest Mass White Dwarf, Mukremin Kulic, Carlos Allende Prieto, Warren R. Brown, and D. Koester, The Astrophysical Journal 660, #2 (May 2007), pp. 1451-1461. and as high as 1.33White dwarf mass distribution in the SDSS, S. O. Kepler, S. J. Kleinman, A. Nitta, D. Koester, B. G. Castanheira, O. Giovannini, A. F. M. Costa, and L. Althaus, Monthly Notices of the Royal Astronomical Society 375, #4 (March 2007), pp. 1315-1324. solar masses, the mass distribution is strongly peaked at 0.6 solar mass, and the majority lie between 0.5 to 0.7 solar mass.  The estimated radii of observed white dwarfs, however, are typically between 0.008 and 0.02 times the radius of the Sun;Masses and radii of white-dwarf stars. III - Results for 110 hydrogen-rich and 28 helium-rich stars, H. L. Shipman, The Astrophysical Journal 228 (February 15, 1979), pp. 240-256. this is comparable to the Earth's radius of approximately 0.009 solar radius.  A white dwarf, then, packs mass comparable to the Sun's into a volume that is typically a million times smaller than the Sun's; the average density of matter in a white dwarf must therefore be, very roughly, 1,000,000 times greater than the average density of the Sun, or approximately 106 grams (1 tonne) per cubic centimeter.  White dwarfs are composed of one of the densest forms of matter known, surpassed only by other compact stars such as neutron stars, black holes and, hypothetically, quark stars.Exotic Phases of Matter in Compact Stars, Fredrik Sandin, licentiate thesis, Luleå University of Technology, May 8, 2005.
White dwarfs were found to be extremely dense soon after their discovery.  If a star is in a binary system, as is the case for Sirius B and 40 Eridani B, it is possible to estimate its mass from observations of the binary orbit.  This was done for Sirius B by 1910,Preliminary General Catalogue, L. Boss, Washington, D.C.: Carnegie Institution, 1910. yielding a mass estimate of 0.94 solar mass.  (A more modern estimate is 1.00 solar mass.)The Age and Progenitor Mass of Sirius B, James Liebert, Patrick A. Young, David Arnett, J. B. Holberg, and Kurtis A. Williams, The Astrophysical Journal 630, #1 (September 2005) pp. L69-L72.  Since hotter bodies radiate more than colder ones, a star's surface bness can be estimated from its effective surface temperature, and hence from its spectrum.   If the star's distance is known, its overall luminosity can also be estimated.  Comparison of the two figures yields the star's radius.  Reasoning of this sort led to the realization, puzzling to astronomers at the time, that Sirius B and 40 Eridani B must be very dense.  For example, when Ernst Öpik estimated the density of a number of visual binary stars in 1916, he found that 40 Eridani B had a density of over 25,000 times the Sun's, which was so high that he called it "impossible".The Densities of Visual Binary Stars, E. Öpik, The Astrophysical Journal 44 (December 1916), pp. 292-302.  As Arthur Stanley Eddington put it later in 1927:Stars and Atoms, A. S. Eddington, Oxford: Clarendon Press, 1927., p. 50
We learn about the stars by receiving and interpreting the messages which their light brings to us.  The message of the Companion of Sirius when it was decoded ran: "I am composed of material 3,000 times denser than anything you have ever come across; a ton of my material would be a little nugget that you could put in a matchbox."  What reply can one make to such a message?  The reply which most of us made in 1914 was-"Shut up.  Don't talk nonsense."
As Eddington pointed out in 1924, densities of this order implied that, according to the theory of general relativity, the light from Sirius B should be gravitationally redshifted.On the relation between the masses and luminosities of the stars, A. S. Eddington, Monthly Notices of the Royal Astronomical Society 84 (March 1924), pp. 308-332.  This was confirmed when Adams measured this redshift in 1925.The Relativity Displacement of the Spectral Lines in the Companion of Sirius, Walter S. Adams, Proceedings of the National Academy of Sciences of the United States of America 11, #7 (July 1925), pp. 382-387.
Such densities are possible because white dwarf material is not composed of atoms bound by chemical bonds, but rather consists of a plasma of unbound nuclei and electrons.  There is therefore no obstacle to placing nuclei closer to each other than electron orbitals-the regions occupied by electrons bound to an atom-would normally allow.   Eddington, however, wondered what would happen when this plasma cooled and the energy which kept the atoms ionized was no longer present.On Dense Matter, R. H. Fowler, Monthly Notices of the Royal Astronomical Society 87 (1926), pp. 114-122.  This paradox was resolved by R. H. Fowler in 1926 by an application of the newly devised quantum mechanics.  Since electrons obey the Pauli exclusion principle, no two electrons can occupy the same state, and they must obey Fermi-Dirac statistics, also introduced in 1926 to determine the statistical distribution of particles which satisfy the Pauli exclusion principle.The Development of the Quantum Mechanical Electron Theory of Metals: 1900-28, Lillian H. Hoddeson and G. Baym, Proceedings of the Royal Society of London, Series A, Mathematical and Physical Sciences 371, #1744 (June 10, 1980), pp. 8-23.  At zero temperature, therefore, electrons could not all occupy the lowest-energy, or ground, state; some of them had to occupy higher-energy states, forming a band of lowest-available energy states, the Fermi sea.   This state of the electrons, called degenerate, meant that a white dwarf could cool to zero temperature and still possess high energy.  Another way of deriving this result is by use of the uncertainty principle: the high density of electrons in a white dwarf means that their positions are relatively localized, creating a corresponding uncertainty in their momenta.  This means that some electrons must have high momentum and hence high kinetic energy.
Compression of a white dwarf will increase the number of electrons in a given volume.  Applying either the Pauli exclusion principle or the uncertainty principle, we can see that this will increase the kinetic energy of the electrons, causing pressure.Lecture 12 - Degeneracy pressure, Rachel Bean, lecture notes, Astronomy 211, Cornell University.  Accessed on line September 21, 2007.  This electron degeneracy pressure is what supports a white dwarf against gravitational collapse.  It depends only on density and not on temperature.  Degenerate matter is relatively compressible; this means that the density of a high-mass white dwarf is so much greater than that of a low-mass white dwarf that the radius of a white dwarf decreases as its mass increases.
The existence of a limiting mass that no white dwarf can exceed is another consequence of being supported by electron degeneracy pressure. These masses were first published in 1929 by Wilhelm Anderson-anddocsymbol-Uuml;ber die Grenzdichte der Materie und der Energie, Wilhelm Anderson, Zeitschrift f-anddocsymbol-uuml;r Physik 56, #11-12 (November 1929), pp. 851-856. and in 1930 by Edmund C. Stoner.The Equilibrium of Dense Stars, Edmund C. Stoner, Philosophical Magazine (7th series) 9 (1930), pp. 944-963.  The modern value of the limit was first published in 1931 by Subrahmanyan Chandrasekhar in his paper "The Maximum Mass of Ideal White Dwarfs".The Maximum Mass of Ideal White Dwarfs, S. Chandrasekhar, The Astrophysical Journal 74, #1 (July 1931), pp. 81-82.  For a nonrotating white dwarf, it is equal to approximately 5.7/μe2 solar masses, where μe is the average molecular weight per electron of the star.The Highly Collapsed Configurations of a Stellar Mass (second paper), S. Chandrasekhar, Monthly Notices of the Royal Astronomical Society, 95 (1935), pp. 207-225., eq. (63)  As the carbon-12 and oxygen-16 which predominantly compose a carbon-oxygen white dwarf both have atomic number equal to half their atomic weight, one should take μe equal to 2 for such a star, leading to the commonly quoted value of 1.4 solar masses.  (Near the beginning of the 20th century, there was reason to believe that stars were composed chiefly of heavy elements,, p. 955 so, in his 1931 paper, Chandrasekhar set the average molecular weight per electron, μe, equal to 2.5, giving a limit of 0.91 solar mass.)  Together with William Alfred Fowler, Chandrasekhar received the Nobel prize for this and other work in 1983.The Nobel Prize in Physics 1983, Nobel Foundation.  Accessed on line May 4, 2007.  The limiting mass is now called the Chandrasekhar limit.
If a white dwarf were to exceed the Chandrasekhar limit, and nuclear reactions did not take place, the pressure exerted by electrons would no longer be able to balance the force of gravity, and it would collapse into a denser object such as a neutron star.The Possible White Dwarf-Neutron Star Connection, R. Canal and J. Gutierrez, arXiv:astro-ph/9701225v1, January 29, 1997.  However, carbon-oxygen white dwarfs accreting mass from a neighboring star undergo a runaway nuclear fusion reaction, which leads to a Type Ia supernova explosion in which the white dwarf is destroyed, just before reaching the limiting mass.Type IA Supernova Explosion Models, Wolfgang Hillebrandt and Jens C. Niemeyer, Annual Review of Astronomy and Astrophysics 38 (2000), pp. 191-230.
White dwarfs have low luminosity and therefore occupy a strip at the bottom of the Hertzsprung-Russell diagram, a graph of stellar luminosity versus color (or temperature). They should not be confused with low-luminosity objects at the low-mass end of the main sequence, such as the hydrogen-fusing red dwarfs, whose cores are supported in part by thermal pressure,Theory of Low-Mass Stars and Substellar Objects, Gilles Chabrier and Isabelle Baraffe, Annual Review of Astronomy and Astrophysics 38 (2000), pp. 337-377. or the even lower-temperature brown dwarfs.The Hertzsprung-Russell (HR) diagram, Jim Kaler, online article.  Accessed on line May 5, 2007.

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<section id="7723" fatherid="7721" title="Mass-radius relationship and mass limit" lv="" priority="2" level="2">
It is simple to derive a rough relationship between the mass and radii of white dwarfs using an energy minimization argument.  The energy of the white dwarf can be approximated by taking it to be the sum of its gravitational potential energy and kinetic energy.  The gravitational potential energy of a unit mass piece of white dwarf, Eg, will be on the order of −GM/R, where G is the gravitational constant, M is the mass of the white dwarf, and R is its radius.  The kinetic energy of the unit mass, Ek, will primarily come from the motion of electrons, so it will be approximately N p2/2m, where p is the average electron momentum, m is the electron mass, and N is the number of electrons per unit mass.  Since the electrons are degenerate, we can estimate p to be on the order of the uncertainty in momentum, Δp, given by the uncertainty principle, which says that Δp Δx is on the order of the reduced Planck constant, ħ.  Δx will be on the order of the average distance between electrons, which will be approximately n−1/3, i.e., the reciprocal of the cube root of the number density, n, of electrons per unit volume.  Since there are N M electrons in the white dwarf and its volume is on the order of R3, n will be on the order of N M / R3.Estimating Stellar Parameters from Energy Equipartition, ScienceBits.  Accessed on line May 9, 2007.
Solving for the kinetic energy per unit mass, Ek, we find that
:E_k approx frac{N (Delta p)^2}{2m} approx frac{N hbar^2 n^{2/3}}{2m} approx frac{M^{2/3} N^{5/3} hbar^2}{2m R^2}.
The white dwarf will be at equilibrium when its total energy, Eg + Ek, is minimized.  At this point, the kinetic and gravitational potential energies should be comparable, so we may derive a rough mass-radius relationship by equating their magnitudes:
:|E_g|approxfrac{GM}{R} = E_kapproxfrac{M^{2/3} N^{5/3} hbar^2}{2m R^2}.
Solving this for the radius, R, gives
: R approx frac{N^{5/3} hbar^2}{2m GM^{1/3}}.
Dropping N, which depends only on the composition of the white dwarf, and the universal constants leaves us with a relationship between mass and radius:
:R sim frac{1}{M^{1/3}}, ,
i.e., the radius of a white dwarf is inversely proportional to the cube root of its mass.
Since this analysis uses the non-relativistic formula p2/2m for the kinetic energy, it is non-relativistic.  If we wish to analyze the situation where the electron velocity in a white dwarf is close to the speed of light, c, we should replace p2/2m by the extreme relativistic approximation p c for the kinetic energy.  With this substitution, we find
:E_{k {
m relativistic}} approx frac{M^{1/3} N^{4/3} hbar c}{R}.
If we equate this to the magnitude of Eg, we find that R drops out and the mass, M, is forced to be
:M_{
m limit} approx N^2 (frac{hbar c}{G}
ight)^{3/2}.
To interpret this result, observe that as we add mass to a white dwarf, its radius will decrease, so, by the uncertainty principle, the momentum, and hence the velocity, of its electrons will increase.  As this velocity approaches c, the extreme relativistic analysis becomes more exact, meaning that the mass M of the white dwarf must approach Mlimit.  Therefore, no white dwarf can be heavier than the limiting mass Mlimit.
For a more accurate computation of the mass-radius relationship and limiting mass of a white dwarf, one must compute the equation of state which describes the relationship between density and pressure in the white dwarf material.  If the density and pressure are both set equal to functions of the radius from the center of the star, the system of equations consisting of the hydrostatic equation together with the equation of state can then be solved to find the structure of the white dwarf at equilibrium.  In the non-relativistic case, we will still find that the radius is inversely proportional to the cube root of the mass., eq. (80)  Relativistic corrections will alter the result so that the radius becomes zero at a finite value of the mass.  This is the limiting value of the mass-called the Chandrasekhar limit-at which the white dwarf can no longer be supported by electron degeneracy pressure.  The graph on the  shows the result of such a computation.  It shows how radius varies with mass for non-relativistic (blue curve) and relativistic (green curve) models of a white dwarf.  Both models treat the white dwarf as a cold Fermi gas in hydrostatic equilibrium.  The average molecular weight per electron, μe, has been set equal to 2.  Radius is measured in standard solar radii and mass in standard solar masses.Standards for Astronomical Catalogues, Version 2.0, section 3.2.2.  Accessed on line January 12, 2007.
These computations all assume that the white dwarf is nonrotating.  If the white dwarf is rotating, the equation of hydrostatic equilibrium must be modified to take into account the centrifugal pseudo-force arising from working in a rotating frame.The Structure, Stability, and Dynamics of Self-Gravitating Systems, Joel E. Tohline, online book.  Accessed on line May 30, 2007.  For a uniformly rotating white dwarf, the limiting mass increases only slightly.  However, if the star is allowed to rotate nonuniformly, and viscosity is neglected, then, as was pointed out by Fred Hoyle in 1947,Note on equilibrium configurations for rotating white dwarfs, F. Hoyle, Monthly Notices of the Royal Astronomical Society 107 (1947), pp. 231-236. there is no limit to the mass for which it is possible for a model white dwarf to be in static equilibrium.  Not all of these model stars, however, will be dynamically stable.Rapidly Rotating Stars. II. Massive White Dwarfs, Jeremiah P. Ostriker and Peter Bodenheimer, The Astrophysical Journal 151 (March 1968), pp. 1089-1098.

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<section id="7724" fatherid="7721" title="Radiation and cooling" lv="" priority="3" level="2">
The visible radiation emitted by white dwarfs varies over a wide color range, from the blue-white color of an O-type main sequence star to the red of a M-type red dwarf.A proposed new white dwarf spectral classification system, E. M. Sion, J. L. Greenstein, J. D. Landstreet, J. Liebert, H. L. Shipman, and G. A. Wegner, The Astrophysical Journal 269, #1 (June 1, 1983), pp. 253-257.  White dwarf effective surface temperatures extend from over 150,000 K to under 4,000 K.White dwarfs, Gilles Fontaine and Fran-anddocsymbol-ccedil;ois Wesemael, in Encyclopedia of Astronomy and Astrophysics, edited by Paul Murdin, Bristol and Philadelphia: Institute of Physics Publishing and London, New York and Tokyo: Nature Publishing Group, 2001.  ISBN 0333750888.  In accordance with the Stefan-Boltzmann law, luminosity increases with increasing surface temperature; this surface temperature range corresponds to a luminosity from over 100 times the Sun's to under 1/10,000th that of the Sun's.  Hot white dwarfs, with surface temperatures in excess of 30,000 K, have been observed to be sources of soft (i.e., lower-energy) X-rays.  This enables the composition and structure of their atmospheres to be studied by soft X-ray and extreme ultraviolet observations.X-ray emission from isolated hot white dwarfs, J. Heise, Space Science Reviews 40 (February 1985), pp. 79-90.
As was explained by Leon Mestel in 1952, unless the white dwarf accretes matter from a companion star or other source, its radiation comes from its stored heat, which is not replenished.On the theory of white dwarf stars. I. The energy sources of white dwarfs, L. Mestel, Monthly Notices of the Royal Astronomical Society 112 (1952), pp. 583–597.White Dwarf Stars and the Hubble Deep Field, S. D. Kawaler, pp. 252–271 in The Hubble Deep Field: Proceedings of the Space Telescope Science Institute Symposium, held in Baltimore, Maryland, May 6–9, 1997, edited by Mario Livio, S. Michael Fall, and Piero Madau, Space Telescope Science Institute symposium series, 11, New York: Cambridge University Press, 1998, ISBN 0521630975., §2.1.  White dwarfs have an extremely small surface area to radiate this heat from, so they cool gradually, remaining hot for a long time.  As a white dwarf cools, its surface temperature decreases, the radiation which it emits reddens, and its luminosity decreases.  Since the white dwarf has no energy sink other than radiation, it follows that its cooling slows with time.   Bergeron, Ruiz, and Leggett, for example, estimate that after a carbon white dwarf of 0.59 solar mass with a hydrogen atmosphere has cooled to a surface temperature of 7,140 K, taking approximately 1.5 billion years, cooling approximately 500 more kelvins to 6,590 K takes around 0.3 billion years, but the next two steps of around 500 kelvins (to 6,030 K and 5,550 K) take first 0.4 and then 1.1 billion years.The Chemical Evolution of Cool White Dwarfs and the Age of the Local Galactic Disk, P. Bergeron, Maria Teresa Ruiz, and S. K. Leggett, The Astrophysical Journal Supplement Series 108, #1 (January 1997), pp. 339-387., Table 2.  Although white dwarf material is initially plasma-a fluid composed of nuclei and electrons-it was theoretically predicted in the 1960s that at a late stage of cooling, it should crystallize, starting at the center of the star.Testing White Dwarf Crystallization Theory with Asteroseismology of the Massive Pulsating DA Star BPM 37093, T. S. Metcalfe, M. H. Montgomery, and A. Kanaan, The Astrophysical Journal 605, #2 (April 2004), pp. L133-L136.  The crystal structure is thought to be a body-centered cubic lattice.Crystallization of carbon-oxygen mixtures in white dwarfs, J. L. Barrat, J. P. Hansen, and R. Mochkovitch, Astronomy and Astrophysics 199, #1-2 (June 1988), pp. L15-L18. In 1995 it was pointed out that asteroseismological observations of pulsating white dwarfs yielded a potential test of the crystallization theory,The Status of White Dwarf Asteroseismology and a Glimpse of the Road Ahead, D. E. Winget, Baltic Astronomy 4 (1995), pp. 129-136.  and in 2004, Travis Metcalfe and a team of researchers at the Harvard-Smithsonian Center for Astrophysics estimated, on the basis of such observations, that approximately 90% of the mass of BPM 37093 had crystallized.Diamond star thrills astronomers, David Whitehouse, BBC News, February 16, 2004.  Accessed on line January 6, 2007.Press release, Harvard-Smithsonian Center for Astrophysics, 2004.Whole Earth Telescope observations of BPM 37093: a seismological test of crystallization theory in white dwarfs, A. Kanaan, A. Nitta, D. E. Winget, S. O. Kepler, M. H. Montgomery, T. S. Metcalfe, et al., arXiv:astro-ph/0411199v1, November 8, 2004.    Other work gives a crystallized mass fraction of between 32% and 82%.Asteroseismology of the Crystallized ZZ Ceti Star BPM 37093: A Different View, P. Brassard and G. Fontaine, The Astrophysical Journal 622, #1 (March 2005), pp. 572-576.
Most observed white dwarfs have relatively high surface temperatures, between 8,000 K and 40,000 K.III/235A: A Catalogue of Spectroscopically Identified White Dwarfs, G.P. McCook and E.M. Sion, on line at the Centre de Données astronomiques de Strasbourg.  Accessed on line May 9, 2007.  A white dwarf, though, spends more of its lifetime at cooler temperatures than at hotter temperatures, so we should expect that there are more cool white dwarfs than hot white dwarfs.  Once we adjust for the selection effect that hotter, more luminous white dwarfs are easier to observe, we do find that decreasing the temperature range examined results in finding more white dwarfs.The Cool White Dwarf Luminosity Function and the Age of the Galactic Disk, S. K. Leggett, Maria Teresa Ruiz, and P. Bergeron, The Astrophysical Journal 497 (April 1998), pp. 294-302.  This trend stops when we reach extremely cool white dwarfs; few white dwarfs are observed with surface temperatures below 4,000 K,Discovery of New Ultracool White Dwarfs in the Sloan Digital Sky Survey, Evalyn Gates, Geza Gyuk, Hugh C. Harris, Mark Subbarao, Scott Anderson, S. J. Kleinman, James Liebert, Howard Brewington, J. Brinkmann, Michael Harvanek, Jurek Krzesinski, Don Q. Lamb, Dan Long, Eric H. Neilsen, Jr., Peter R. Newman, Atsuko Nitta, and Stephanie A. Snedden, The Astrophysical Journal 612, #2 (September 2004), pp. L129-L132. and one of the coolest so far observed, WD 0346+246, has a surface temperature of approximately 3,900 K.WD 0346+246: A Very Low Luminosity, Cool Degenerate in Taurus, N. C. Hambly, S. J. Smartt, and S. Hodgkin, The Astrophysical Journal 489 (November 1997), pp. L157-L160.  The reason for this is that, as the Universe's age is finite,The Moment of Creation: Big Bang Physics from Before the First Millisecond to the Present Universe, James S. Trefil, Mineola, New York: Dover Publications, 2004.  ISBN 0486438139. there has not been time for white dwarfs to cool down below this temperature.  The white dwarf luminosity function can therefore be used to find the time when stars started to form in a region; an estimate for the age of the Galactic disk found in this way is 8 billion years.
A white dwarf will eventually cool and become a non-radiating black dwarf in approximate thermal equilibrium with its surroundings and with the cosmic background radiation.  However, no black dwarfs are thought to exist yet.

</section>
<section id="7725" fatherid="7721" title="Atmosphere and spectra" lv="" priority="4" level="2">
Although most white dwarfs are thought to be composed of carbon and oxygen, spectroscopy typically shows that their emitted light comes from an atmosphere which is observed to be either hydrogen-dominated or helium-dominated.  The dominant element is usually at least 1,000 times more abundant than all other elements.  As explained by Schatzman in the 1940s, the high surface gravity is thought to cause this purity by gravitationally separating the atmosphere so that heavy elements are on the bottom and lighter ones on top.Théorie du débit d'énergie des naines blanches, Evry Schatzman, Annales d'Astrophysique 8 (January 1945), pp. 143-209.Physics of white dwarf stars, D. Koester and G. Chanmugam, Reports on Progress in Physics 53 (1990), pp. 837-915., §5-6  This atmosphere, the only part of the white dwarf visible to us, is thought to be the top of an envelope which is a residue of the star's envelope in the AGB phase and may also contain material accreted from the interstellar medium.  The envelope is believed to consist of a helium-rich layer with mass no more than 1/100th of the star's total mass, which, if the atmosphere is hydrogen-dominated, is overlain by a hydrogen-rich layer with mass approximately 1/10,000th of the stars total mass.White Dwarf Stars, Steven D. Kawaler, in Stellar remnants, S. D. Kawaler, I. Novikov, and G. Srinivasan, edited by Georges Meynet and Daniel Schaerer, Berlin: Springer, 1997.  Lecture notes for Saas-Fee advanced course number 25.  ISBN 3540615202., §4-5.
Although thin, these outer layers determine the thermal evolution of the white dwarf.  The degenerate electrons in the bulk of a white dwarf conduct heat well.  Most of a white dwarf's mass is therefore almost isothermal, and it is also hot: a white dwarf with surface temperature between 8,000 K and 16,000 K will have a core temperature between approximately 5,000,000 K and 20,000,000 K.  The white dwarf is kept from cooling very quickly only by its outer layers' opacity to radiation.
The first attempt to classify white dwarf spectra appears to have been by G. P. Kuiper in 1941,List of Known White Dwarfs, Gerard P. Kuiper, Publications of the Astronomical Society of the Pacific 53, #314 (August 1941), pp. 248-252. and various classification schemes have been proposed and used since then.The Spectra and Luminosities of White Dwarfs, Willem J. Luyten, Astrophysical Journal 116 (September 1952), pp. 283-290.Stellar atmospheres, Jesse Leonard Greenstein, in Stars and Stellar Systems, vol. 6, Stellar Atmospheres, edited by J. L. Greenstein, Chicago: University of Chicago Press, 1960.  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 ), and a temperature index number, computed by dividing 50,400 K by the effective temperature.  For example:
The symbols ? and : may also be used if the correct classification is uncertain.
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.  A small fraction (roughly 0.1%) have carbon-dominated atmospheres, the hot (above 15,000 K) DQ class.White dwarf stars with carbon atmospheres, Patrick Dufour, James Liebert, G. Fontaine, and N. Behara, Nature 450, #7169 (November 2007), pp. 522–524, ,   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.,§ 2.4  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.

</section>
<section id="7726" fatherid="7721" title="Magnetic field" lv="" priority="5" level="2">
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.The magnetic field of massive rotating bodies, P. M. S. Blackett, Nature 159, #4046 (May 17, 1947), pp. 658-666.  This putative law, sometimes called the Blackett effect, was never generally accepted, and by the 1950s even Blackett felt it had been refuted.Patrick Maynard Stuart Blackett, Baron Blackett, of Chelsea, 18 November 1897-13 July 1974, Bernard Lovell, 
Biographical Memoirs of Fellows of the Royal Society 21 (November 1975), pp. 1-115., pp. 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., §8;Coherent Mechanisms of Radio Emission and Magnetic Models of Pulsars, V. L. Ginzburg, V. V. Zheleznyakov, and V. V. Zaitsev, Astrophysics and Space Science 4 (1969), pp. 464-504., p. 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.Discovery of Circularly Polarized Light from a White Dwarf, James C. Kemp, John B. Swedlund, J. D. Landstreet, and  J. R. P. Angel, The Astrophysical Journal 161 (August 1970), pp. L77-L79.  It is thought to have a surface field of approximately 300 million gauss (30 kT)., §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% of white dwarfs have fields in excess of 1 million gauss (100 T).The fraction of DA white dwarfs with kilo-Gauss magnetic fields, S. Jordan, R. Aznar Cuadrado, R. Napiwotzki, H. M. Schmid, and S. K. Solanki, Astronomy and Astrophysics 462, #3 (February 11, 2007), pp. 1097-1101.The True Incidence of Magnetism Among Field White Dwarfs, James Liebert, P. Bergeron, and J. B. Holberg, Astronomical Journal 125, #1 (January 2003), pp. 348-353.

</section>
</section>
<section id="7727" fatherid="0" title="Variability" lv="" priority="4" level="1">
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., § 7.1.1;Ultrashort-Period Stellar Oscillations. I. Results from White Dwarfs, Old Novae, Central Stars of Planetary Nebulae, 3C 273, and Scorpius XR-1, George M. Lawrence, Jeremiah P. Ostriker, and James E. Hesser, The Astrophysical Journal 148, #3 (June 1967), pp. L161-L163.  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.A New Short-Period Blue Variable, Arlo U. Landolt, The Astrophysical Journal 153, #1 (July 1968), pp. 151-164.  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., § 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;, pp. 891, 895 DBV, or V777 Her, stars, with helium-dominated atmospheres and the spectral type DB;White dwarfs, Gilles Fontaine and Fran-anddocsymbol-ccedil;ois Wesemael, in Encyclopedia of Astronomy and Astrophysics, ed. Paul Murdin, Bristol and Philadelphia: Institute of Physics Publishing and London, New York and Tokyo: Nature Publishing Group, 2001.  ISBN 0333750888., p. 3525 and GW Vir stars (sometimes subdivided into DOV and PNNV stars), with atmospheres dominated by helium, carbon, and oxygen.Mapping the Instability Domains of GW Vir Stars in the Effective Temperature-Surface Gravity Diagram, P.-O. Quirion, G. Fontaine, and P. Brassard, The Astrophysical Journal Supplement Series 171, #1 (July 2007), pp. 219-248.,§1.1, 1.2;Detection of non-radial g-mode pulsations in the newly discovered PG 1159 star HE 1429-1209, T. Nagel and K. Werner, Astronomy and Astrophysics 426 (2004), pp. L45-L48.,§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., § 1.1;The Extent and Cause of the Pre-White Dwarf Instability Strip, M. S. O'Brien, The Astrophysical Journal 532, #2 (April 2000), pp. 1078-1088.  These variables all exhibit small (1%-30%) 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.Asteroseismology of white dwarf stars, D. E. Winget, Journal of Physics: Condensed Matter 10, #49 (December 14, 1998), pp. 11247-11261.  DOI 10.1088/0953-8984/10/49/014.

</section>
<section id="7728" fatherid="0" title="Formation" lv="" priority="5" level="1">
<section id="7729" fatherid="7728" title="" lv="" priority="-1" level="2">
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.How Massive Single Stars End Their Life, A. Heger, C. L. Fryer, S. E. Woosley, N. Langer, and D. H. Hartmann, The Astrophysical Journal 591, #1 (2003), pp. 288-300.  The composition of the white dwarf produced will differ depending on the initial mass of the star.
</section>
<section id="7730" fatherid="7728" title="Stars with very low mass" lv="" priority="2" level="2">
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) 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.The End of the Main Sequence, Gregory Laughlin, Peter Bodenheimer,  and Fred C. Adams, Astrophysical Journal 482, #1 (June 10, 1997), pp. 420-432.  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 systems Stars Beyond Maturity, Simon Jeffery, online article.  Accessed on line May 3, 2007.Helium core white dwarf evolution-including white dwarf companions to neutron stars, M. J. Sarna, E. Ergma, and J. Gerskevits, Astronomische Nachrichten 322, #5/6 (December 2001), pp. 405-410.The formation of helium white dwarfs in close binary systems - II, O. G. Benvenuto, M. A. De Vito, Monthly Notices of the Royal Astronomical Society 362, #3 (September 2005), pp. 891-905. or mass loss due to a large planetary companion.Formation of undermassive single white dwarfs and the influence of planets on late stellar evolution, G. Nelemans and T. M. Tauris, Astronomy and Astrophysics 335 (July 1998), pp. L85–L88.
</section>
<section id="7731" fatherid="7728" title="Stars with low to medium mass" lv="" priority="3" level="2">
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 .  This process is responsible for the carbon-oxygen white dwarfs which form the vast majority of observed white dwarfs.the evolution of low-mass stars, Vik Dhillon, lecture notes, Physics 213, University of Sheffield.  Accessed on line May 3, 2007.the evolution of high-mass stars, Vik Dhillon, lecture notes, Physics 213, University of Sheffield.  Accessed on line May 3, 2007.
</section>
<section id="7732" fatherid="7728" title="Stars with medium to high mass" lv="" priority="4" level="2">
If a star is massive enough, 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, because 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.Strange quark matter in stars: a general overview, J-anddocsymbol-uuml;rgen Schaffner-Bielich, Journal of Physics G: Nuclear and Particle Physics 31, #6 (2005), pp. S651-S657; also arXiv:astro-ph/0412215v1.  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.Evolution of 8-10 solar mass stars toward electron capture supernovae. I - Formation of electron-degenerate O + Ne + Mg cores, Ken'ichi Nomoto, The Astrophysical Journal 277 (February 15, 1984), pp. 791-805.The evolution and explosion of massive stars, S. E. Woosley, A. Heger, and T. A. Weaver, Reviews of Modern Physics 74, #4 (October 2002), pp. 1015-1071.  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.Chandra and FUSE spectroscopy of the hot bare stellar core H 1504+65, K. Werner, T. Rauch, M. A. Barstow, and J. W. Kruk, Astronomy and Astrophysics 421 (2004), pp. 1169-1183.On the interpretation and implications of nova abundances: an abundance of riches or an overabundance of enrichments, Mario Livio and James W. Truran, The Astrophysical Journal 425, #2 (April 1994), pp. 797-801.
</section>
</section>
<section id="7733" fatherid="0" title="Fate" lv="" priority="6" level="1">
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.A dying universe: the long-term fate and evolution of astrophysical objects, Fred C. Adams and Gregory Laughlin, Reviews of Modern Physics 69, #2 (April 1997), pp. 337-372., §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., §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 enough mass to become a nondegenerate lump of matter, and finally disappears completely., §IV.
</section>
<section id="7734" fatherid="0" title="Stellar system" lv="" priority="7" level="1">
<section id="7735" fatherid="7734" title="" lv="" priority="-1" level="2">
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.Comet clash kicks up dusty haze, BBC News, February 13, 2007.  Accessed on line September 20, 2007.A Debris Disk around the Central Star of the Helix Nebula?, K. Y. L. Su, Y.-H. Chu, G. H. Rieke, P. J. Huggins, R. Gruendl, R. Napiwotzki, T. Rauch, W. B. Latter, and K. Volk, The Astrophysical Journal 657, #1 (March 2007), pp. L41-L45.  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.The Dust Cloud around the White Dwarf G29-38, William T. Reach, Marc J. Kuchner, Ted von Hippel, Adam Burrows, Fergal Mullally, Mukremin Kilic, and D. E. Winget, The Astrophysical Journal 635, #2 (December 2005), pp. L161-L164.   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.
</section>
<section id="7736" fatherid="7734" title="Type Ia supernovae" lv="" priority="2" level="2">
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.)Presupernova Evolution of Accreting White Dwarfs with Rotation, S.-C. Yoon and N. Langer, Astronomy and Astrophysics 419, #2 (May 2004), pp. 623-644.  Accessed on line May 30, 2007.  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.
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,, 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.  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.Theoretical light curves for deflagration models of type Ia supernova, S. I. Blinnikov, F. K. Röpke, E. I. Sorokina, M. Gieseler, M. Reinecke, C. Travaglio, W. Hillebrandt, and M. Stritzinger, Astronomy and Astrophysics 453, #1 (July 2006), pp.229-240.  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., p. 14.

</section>
<section id="7737" fatherid="7734" title="Cataclysmic variables" lv="" priority="3" level="2">
Before accretion of material pushes a white dwarf close to the Chandrasekhar limit, accreted hydrogen-rich material on the surface may ignite in a less destructive type of thermonuclear explosion powered by hydrogen fusion. 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 be caused by the release of gravitational potential energy when part of the accretion disc collapses onto the star, rather than by fusion.  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.Imagine the Universe! Cataclysmic Variables, fact sheet at NASA Goddard.  Accessed on line May 4, 2007.Introduction to Cataclysmic Variables (CVs), fact sheet at NASA Goddard. Accessed on line May 4, 2007.  Both fusion- and accretion-powered cataclysmic variables have been observed to be X-ray sources.
</section>
</section>
<section id="7738" fatherid="0" title="See also" lv="" priority="-8" level="1">
</section>
<section id="7739" fatherid="0" title="References" lv="" priority="-9" level="1">
</section>
<section id="7740" fatherid="0" title="External links and further reading" lv="" priority="-10" level="1">
<section id="7741" fatherid="7740" title="General" lv="" priority="-1" level="2">
</section>
<section id="7742" fatherid="7740" title="Physics" lv="" priority="-2" level="2">
</section>
<section id="7743" fatherid="7740" title="Variability" lv="" priority="-3" level="2">
</section>
<section id="7744" fatherid="7740" title="Magnetic field" lv="" priority="-4" level="2">
</section>
<section id="7745" fatherid="7740" title="Frequency" lv="" priority="-5" level="2">
</section>
<section id="7746" fatherid="7740" title="Observational" lv="" priority="-6" level="2">
Category:Dark matter
Category:Star types
Category:Stellar evolution
Category:Stellar phenomena
Category:Exotic matter




mzn:اسپه کوتوله

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