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[[Image:HRDiagram.png|thumb|280px|A Hertzsprung–Russell diagram plots the actual brightness (or [[absolute magnitude]]) of a star against its [[color index]] (represented as B-V). The main sequence is visible as a prominent diagonal band that runs from the upper left to the lower right. This plot shows 22,000 stars from the [[Hipparcos Catalogue]] together with 1,000 low-luminosity stars (red and white dwarfs) from the [[Gliese Catalogue of Nearby Stars]].]]
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In [[astronomy]], the '''main sequence''' is a continuous and distinctive band of [[star]]s that appears on plots of stellar [[Color index|color]] versus [[Absolute magnitude|brightness]]. These color-magnitude plots are known as [[Hertzsprung–Russell diagram]]s after their co-developers, [[Ejnar Hertzsprung]] and [[Henry Norris Russell]]. Stars on this band are known as '''main-sequence stars''' or "dwarf" stars.<ref name=smith91/><ref name=powell06/>
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After a star has formed, it generates thermal energy in the dense core region through the [[nuclear fusion]] of [[hydrogen]] atoms into [[helium]]. During this stage of the star's lifetime, it is located along the main sequence at a position determined primarily by its mass, but also based upon its chemical composition and other factors. All main-sequence stars are in [[hydrostatic equilibrium]], where outward thermal pressure from the hot core is balanced by the inward gravitational pressure from the overlying layers. The strong dependence of the rate of energy generation in the core on the temperature and pressure helps to sustain this balance. Energy generated at the core makes its way to the surface and is radiated away at the [[photosphere]]. The energy is carried by either [[radiation]] or [[convection]], with the latter occurring in regions with steeper temperature gradients, higher opacity or both.
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The main sequence is sometimes divided into upper and lower parts, based on the dominant process that a star uses to generate energy. Stars below about 1.5 times the [[Solar mass|mass of the Sun]] (or 1.5 solar masses) primarily fuse hydrogen atoms together in a series of stages to form helium, a sequence called the [[proton–proton chain]]. Above this mass, in the upper main sequence, the nuclear fusion process mainly uses atoms of [[carbon]], [[nitrogen]] and [[oxygen]] as intermediaries in the [[CNO cycle]] that produces helium from hydrogen atoms. Main-sequence stars with more than two solar masses undergo convection in their core regions, which acts to stir up the newly created helium and maintain the proportion of fuel needed for fusion to occur. Below this mass, stars have cores that are entirely radiative with convective zones near the surface. With decreasing stellar mass, the proportion of the star forming a convective envelope steadily increases, while main-sequence stars below 0.4 solar masses undergo convection throughout their mass. When core convection does not occur, a helium-rich core develops surrounded by an outer layer of hydrogen.
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In general, the more massive a star is, the shorter its lifespan on the main sequence. After the hydrogen fuel at the core has been consumed, the star [[stellar evolution|evolves]] away from the main sequence on the HR diagram. The behavior of a star now depends on its mass, with stars below 0.23 [[solar mass]]es becoming [[white dwarf]]s directly, while stars with up to ten solar masses pass through a [[red giant]] stage.<ref name=romp69 /> More massive stars can explode as a [[supernova]],<ref name=science304/> or collapse directly into a [[black hole]].
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[[File:Hot and brilliant O stars in star-forming regions.jpg|thumb|Hot and brilliant [[O-type main-sequence star]]s in star-forming regions. These are all regions of star formation that contain many hot young stars including several bright stars of spectral type O.<ref>{{cite news|title=The Brightest Stars Don't Live Alone|url=http://www.eso.org/public/news/eso1230/|accessdate=27 July 2012|newspaper=ESO Press Release}}</ref> ]]
 
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In the early part of the 20th century, information about the types and distances of [[star]]s became more readily available. The [[spectrum|spectra]] of stars were shown to have distinctive features, which allowed them to be categorized. [[Annie Jump Cannon]] and [[Edward C. Pickering]] at [[Harvard College Observatory]] developed a method of categorization that became known as the [[Stellar classification|Harvard Classification Scheme]], published in the ''Harvard Annals'' in 1901.<ref name=longair06/>
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In [[Potsdam]] in 1906, the Danish astronomer [[Ejnar Hertzsprung]] noticed that the reddest stars—classified as K and M in the Harvard scheme—could be divided into two distinct groups. These stars are either much brighter than the Sun, or much fainter. To distinguish these groups, he called them "giant" and "dwarf" stars. The following year he began studying [[star cluster]]s; large groupings of stars that are co-located at approximately the same distance. He published the first plots of color versus [[luminosity]] for these stars. These plots showed a prominent and continuous sequence of stars, which he named the Main Sequence.<ref name=brown/>
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At [[Princeton University]],.[[Henry Norris Russell]] was following a similar course of research. He was studying the relationship between the spectral classification of stars and their actual brightness as corrected for distance—their [[absolute magnitude]]. For this purpose he used a set of stars that had reliable [[parallax]]es and many of which had been categorized at Harvard. When he plotted the spectral types of these stars against their absolute magnitude, he found that dwarf stars followed a distinct relationship. This allowed the real brightness of a dwarf star to be predicted with reasonable accuracy.<ref name=obs36/>
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Of the red stars observed by Hertzsprung, the dwarf stars also followed the spectra-luminosity relationship discovered by Russell. However, the giant stars are much brighter than dwarfs and so, do not follow the same relationship. Russell proposed that the "giant stars must have low density or great surface-brightness, and the reverse is true of dwarf stars". The same curve also showed that there were very few faint white stars.<ref name=obs36/>
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In 1933, [[Bengt Strömgren]] introduced the term Hertzsprung–Russell diagram to denote a luminosity-spectral class diagram.<ref name=zfa7/> This name reflected the parallel development of this technique by both Hertzsprung and Russell earlier in the century.<ref name=brown/>
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As evolutionary models of stars were developed during the 1930s, it was shown that, for stars of a uniform chemical composition, a relationship exists between a star's mass and its luminosity and radius. That is, for a given mass and composition, there is a unique solution for determining the star's radius and luminosity. This became known as the [[Vogt-Russell theorem]]; named after Heinrich Vogt and Henry Norris Russell. By this theorem, once a star's chemical composition and its position on the main sequence is known, so too is the star's mass and radius. (However, it was subsequently discovered that the theorem breaks down somewhat for stars of non-uniform composition.)<ref name=schatzman33/>
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A refined scheme for [[stellar classification]] was published in 1943 by W. W. Morgan and P. C. Keenan.<ref name=keenan_morgan43/> The MK classification assigned each star a spectral type—based on the Harvard classification—and a luminosity class.  The Harvard classification had been developed by assigning a different letter to each star based on the strength of the hydrogen spectra line, before the relationship between spectra and temperature was known.  When ordered by temperature and when duplicate classes were removed,  the [[spectral type]]s of stars followed, in order of decreasing temperature with colors ranging from blue to red, the sequence O, B, A, F, G, K and M. (A popular [[mnemonic]] for memorizing this sequence of stellar classes is "Oh Be A Fine Girl/Guy, Kiss Me".) The luminosity class ranged from I to V, in order of decreasing luminosity. Stars of luminosity class V belonged to the main sequence.<ref name=tnc/>
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{{Main|Star formation}}
 
When a [[protostar]] is formed from the [[Jeans instability|collapse]] of a [[giant molecular cloud]] of gas and dust in the local [[interstellar medium]], the initial composition is homogeneous throughout, consisting of about 70% hydrogen, 28% helium and trace amounts of other elements, by mass.<ref name=asr34_1/> The initial mass of the star depends on the local conditions within the cloud. (The mass distribution of newly formed stars is described empirically by the [[initial mass function]].)<ref name=science295_5552/> During the initial collapse, this [[pre-main-sequence star]] generates energy through gravitational contraction. Upon reaching a suitable density, energy generation is begun at the core using an [[exothermic]] [[nuclear fusion]] process that converts hydrogen into helium.<ref name=tnc/>
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{{Star nav}}
  <li>[http://www.thomassankara.net/spip.php?article605&date=2008-05/ http://www.thomassankara.net/spip.php?article605&date=2008-05/]</li>
 
 
Once nuclear fusion of hydrogen becomes the dominant energy production process and the excess energy gained from gravitational contraction has been lost,<ref name=science293_5538/> the star lies along a [[curve]] on the [[Hertzsprung–Russell diagram]] (or HR diagram) called the standard main sequence. Astronomers will sometimes refer to this stage as "zero age main sequence", or ZAMS.<ref name=zams_sao/> The ZAMS curve can be calculated using computer models of stellar properties at the point when stars begin hydrogen fusion. From this point, the brightness and surface temperature of stars typically increase with age.<ref name="clayton83"/>
  <li>[http://www.mihmusicbizservices.com/ http://www.mihmusicbizservices.com/]</li>
 
 
A star remains near its initial position on the main sequence until a significant amount of hydrogen in the core has been consumed, then begins to evolve into a more luminous star. (On the HR diagram, the evolving star moves up and to the right of the main sequence.) Thus the main sequence represents the primary hydrogen-burning stage of a star's lifetime.<ref name=tnc/>
  <li>[http://bbs.lmzol.com/forum.php?mod=viewthread&tid=4452098 http://bbs.lmzol.com/forum.php?mod=viewthread&tid=4452098]</li>
 
 
== Properties ==
  <li>[http://www.nnzqm.com/forum.php?mod=viewthread&tid=689988 http://www.nnzqm.com/forum.php?mod=viewthread&tid=689988]</li>
The majority of stars on a typical HR diagram lie along the main-sequence curve. This line is pronounced because both the [[stellar classification|spectral type]] and the [[luminosity]] depend only on a star's mass, at least to [[orders of approximation|zeroth-order approximation]], as long as it is fusing hydrogen at its core—and that is what almost all stars spend most of their "active" lives doing.<ref name=mss_atoe/>
 
 
</ul>
The temperature of a star determines its [[spectral type]] via its effect on the physical properties of [[Plasma (physics)|plasma]] in its [[photosphere]]. A star's energy emission as a function of wavelength is influenced by both its temperature and composition. A key indicator of this energy distribution is given by the [[color index]], ''B''&nbsp;−&nbsp;''V'', which measures the star's [[apparent magnitude|magnitude]] in blue (''B'') and green-yellow (''V'') light by means of filters.<ref group=note>By measuring the difference between these values, this eliminates the need to correct the magnitudes for distance. However, see [[Extinction (astronomy)|extinction]].</ref> This difference in magnitude provides a measure of a star's temperature.
 
==Dwarf terminology==
Main-sequence stars are called dwarf stars, but this terminology is partly historical and can be somewhat confusing. For the cooler stars, dwarfs such as [[red dwarf]]s, [[orange dwarf]]s, and [[yellow dwarf]]s are indeed much smaller and dimmer than other stars of those colors. However, for hotter blue and white stars, the size and brightness difference between so-called ''dwarf'' stars that are on the main sequence and the so-called ''giant'' stars that are not becomes smaller; for the hottest stars it is not directly observable. For those stars the terms ''dwarf'' and ''giant'' refer to differences in [[spectral line]]s which indicate if a star is on the main sequence or off it.  Nevertheless, very hot main-sequence stars are still sometimes called dwarfs, even though they have roughly the same size and brightness as the "giant" stars of that temperature.<ref name=moore06/>
 
The common use of ''dwarf'' to mean main sequence is confusing in another way, because there are dwarf stars which are not main-sequence stars. For example, a [[white dwarf]] is a different kind of star that is much smaller than a main-sequence star—roughly the size of the [[Earth]]. These represent the final evolutionary stage of many main-sequence stars.<ref name=wd_sao/>
 
==Parameters==
By treating the star as an idealized energy radiator known as a [[black body]], the luminosity ''L'' and radius ''R'' can be related to the [[effective temperature]] <math>T_{\rm eff}</math> by the [[Stefan–Boltzmann law]]:
 
:''L'' = 4πσ''R''<sup>2</sup>''T<sub>eff</sub>''<sup>4</sup>
 
where ''σ'' is the [[Stefan–Boltzmann constant]]. As the position of a star on the HR diagram shows its approximate luminosity, this relation can be used to estimate its radius.<ref name=ohrd/>
 
The mass, radius and luminosity of a star are closely interlinked, and their respective values can be approximated by three relations. First is the [[Stefan–Boltzmann law]], which relates the luminosity ''L'', the radius ''R'' and the surface temperature ''T<sub>eff</sub>''. Second is the [[mass–luminosity relation]], which relates the luminosity ''L'' and the mass ''M''. Finally, the relationship between ''M'' and ''R'' is close to linear. The ratio of ''M'' to ''R'' increases by a factor of only three over 2.5 [[orders of magnitude]] of ''M''. This relation is roughly proportional to the star's inner temperature ''T<sub>I</sub>'', and its extremely slow increase reflects the fact that the rate of energy generation in the core strongly depends on this temperature, while it has to fit the mass–luminosity relation. Thus, a too high or too low temperature will result in stellar instability.
 
A better approximation is to take <math>\epsilon = L / M</math>, the energy generation rate per unit mass, as ε is proportional to ''T<sub>I</sub>''<sup>15</sup>, where ''T<sub>I</sub>'' is the core temperature. This is suitable for stars at least as massive as the Sun, exhibiting the [[CNO cycle]], and gives the better fit ''R'' ∝ ''M''<sup>0.78</sup>.<ref>{{cite web
| title=A course on stars' physical properties, formation and evolution
| publisher=University of St. Andrews
| url=http://www-star.st-and.ac.uk/~kw25/teaching/stars/STRUC4.pdf
| accessdate=2010-05-18 }}</ref>
 
=== Sample parameters ===
The table below shows typical values for stars along the main sequence. The values of [[luminosity]] (''L''), [[radius]] (''R'') and [[mass]] (''M'') are relative to the Sun—a dwarf star with a spectral classification of G2 V. The actual values for a star may vary by as much as 20–30% from the values listed below.<ref name=siess00/>
 
<!-- Please include a solid reference if you add additional values on this table. -->
:{| class="wikitable" style="text-align: center;" border="1" cellspacing="0" cellpadding="4"
|+ Table of main-sequence stellar parameters<ref name=zombeck/>
|- bgcolor="#FFFFCC"
!rowspan="2" style="font-size: smaller;"|[[Stellar classification|Stellar<br />Class]]
!style="font-size: smaller;"|[[Radius]]
!style="font-size: smaller;"|Mass
!style="font-size: smaller;"|Luminosity
!style="font-size: smaller;"|Temperature
!rowspan="2"|Examples<ref name=simbad/>
|- bgcolor="#FFFFEE"
|R/[[solar radius|R<sub>☉</sub>]]
|M/[[solar mass|M<sub>☉</sub>]]
|L/[[solar luminosity|L<sub>☉</sub>]]
|[[kelvin|K]]
|-
| O6 || 18  || 40  || 500,000 || 38,000
|style="text-align: left;"|[[Theta1 Orionis C]]
|-
| B0 || 7.4  || 18  || 20,000  || 30,000
|style="text-align: left;"|[[Phi Orionis|Phi<sup>1</sup> Orionis]]
|-
| B5 || 3.8  || 6.5  || 800    || 16,400
|style="text-align: left;"|[[Pi Andromedae|Pi Andromedae A]]
|-
| A0 || 2.5  || 3.2  || 80      || 10,800
|style="text-align: left;"|[[Alpha Coronae Borealis|Alpha Coronae Borealis A]]
|-
| A5 || 1.7  || 2.1  || 20      || 8,620
|style="text-align: left;"|[[Beta Pictoris]]
|-
| F0 || 1.3  || 1.7  || 6      || 7,240
|style="text-align: left;"|[[Gamma Virginis]]
|-
| F5 || 1.2  || 1.3 || 2.5    || 6,540
|style="text-align: left;"|[[Eta Arietis]]
|-
| G0 || 1.05 || 1.10 || 1.26    || 5,920
|style="text-align: left;"|[[Beta Comae Berenices]]
|-
| G2 || 1.00 || 1.00 || 1.00    || 5,780
|style="text-align: left;"|[[Sun]]<ref name=bydef group=note>The Sun is a typical type G2V star.</ref>
|-
| G5 || 0.93 || 0.93 || 0.79    || 5,610
|style="text-align: left;"|[[Alpha Mensae]]
|-
| K0 || 0.85 || 0.78 || 0.40    || 5,240
|style="text-align: left;"|[[70 Ophiuchi|70 Ophiuchi A]]
|-
| K5 || 0.74 || 0.69 || 0.16    || 4,410
|style="text-align: left;"|[[61 Cygni|61 Cygni A]]<ref name=apj129/>
|-
| M0 || 0.63 || 0.47 || 0.063  || 3,920
|style="text-align: left;"|[[Gliese 185]]<ref name=simbad_ltt2151/>
|-
| M5 || 0.32 || 0.21 || 0.0079  || 3,120
|style="text-align: left;"|[[EZ Aquarii|EZ Aquarii A]]
|-
| M8 || 0.13 || 0.10 || 0.0008  || 2,660
|style="text-align: left;"|[[VB 10|Van Biesbroeck's star]]<ref name=recons/>
|}
 
== Energy generation ==
{{See also|Stellar nucleosynthesis}}
All main-sequence stars have a core region where energy is generated by nuclear fusion. The temperature and density of this core are at the levels necessary to sustain the energy production that will support the remainder of the star. A reduction of energy production would cause the overlaying mass to compress the core, resulting in an increase in the fusion rate because of higher temperature and pressure. Likewise an increase in energy production would cause the star to expand, lowering the pressure at the core. Thus the star forms a self-regulating system in [[hydrostatic equilibrium]] that is stable over the course of its main sequence lifetime.<ref name=brainerd/>
 
[[Image:Nuclear energy generation.svg|right|280px|thumb|This graph shows the [[logarithm]] of the relative energy output (&epsilon;) for the [[Proton-proton chain reaction|proton-proton]] (PP), [[CNO cycle|CNO]] and [[Triple-alpha process|triple-α]] fusion processes at different temperatures. The dashed line shows the combined energy generation of the PP and CNO processes within a star. At the Sun's core temperature, the PP process is more efficient.]]
Main-sequence stars employ two types of hydrogen fusion processes, and the rate of energy generation from each type depends on the temperature in the core region. Astronomers divide the main sequence into upper and lower parts, based on which of the two is the dominant fusion process. In the lower main sequence, energy is primarily generated as the result of the [[proton-proton chain]], which directly fuses hydrogen together in a series of stages to produce helium.<ref name=hannu/> Stars in the upper main sequence have sufficiently high core temperatures to efficiently use the [[CNO cycle]]. (See the chart.) This process uses atoms of [[carbon]], [[nitrogen]] and [[oxygen]] as intermediaries in the process of fusing hydrogen into helium.
 
At a stellar core temperature of 18 million [[kelvin]]s, the PP process and CNO cycle are equally efficient, and each type generates half of the star's net luminosity. As this is the core temperature of a star with about 1.5 [[solar mass]]es, the upper main sequence consists of stars above this mass. Thus, roughly speaking, stars of spectral class F or cooler belong to the lower main sequence, while class A stars or hotter are upper main-sequence stars.<ref name="clayton83"/> The transition in primary energy production from one form to the other spans a range difference of less than a single solar mass. In the Sun, a one solar mass star, only 1.5% of the energy is generated by the CNO cycle.<ref name=apj555/> By contrast, stars with 1.8 solar masses or above generate almost their entire energy output through the CNO cycle.<ref name=maurizio05/>
 
The observed upper limit for a main-sequence star is 120–200 solar masses.<ref name=apj620_1/> The theoretical explanation for this limit is that stars above this mass can not radiate energy fast enough to remain stable, so any additional mass will be ejected in a series of pulsations until the star reaches a stable limit.<ref name=apj162/> The lower limit for sustained proton–proton nuclear fusion is about 0.08 solar masses.<ref name=hannu/> Below this threshold are sub-stellar objects that can not sustain hydrogen fusion, known as [[brown dwarf]]s.<ref name=apj406_1/>
 
== Structure ==
{{Main|Stellar structure}}
[[Image:Solar internal structure.svg|right|280px|thumb|This diagram shows a cross-section of a Sun-like star, showing the internal structure.]]
 
Because there is a temperature difference between the core and the surface, or [[photosphere]], energy is transported outward. The two modes for transporting this energy are [[radiation]] and [[convection]]. A [[radiation zone]], where energy is transported by [[radiation]], is stable against convection and there is very little mixing of the plasma. By contrast, in a [[convection zone]] the energy is transported by bulk movement of plasma, with hotter material rising and cooler material descending. Convection is a more efficient mode for carrying energy than radiation, but it will only occur under conditions that create a steep temperature gradient.<ref name=brainerd/><ref name=aller91/>
 
In massive stars (above 10 solar masses)<ref name=aaa102_1/> the rate of energy generation by the CNO cycle is very sensitive to temperature, so the fusion is highly concentrated at the core. Consequently, there is a high temperature gradient in the core region, which results in a convection zone for more efficient energy transport.<ref name=hannu/> This mixing of material around the core removes the helium ash from the hydrogen-burning region, allowing more of the hydrogen in the star to be consumed during the main-sequence lifetime. The outer regions of a massive star transport energy by radiation, with little or no convection.<ref name=brainerd/>
 
Intermediate mass stars such as [[Sirius]] may transport energy primarily by radiation, with a small core convection region.<ref name=lockner06/> Medium-sized, low mass stars like the Sun have a core region that is stable against convection, with a convection zone near the surface that mixes the outer layers. This results in a steady buildup of a helium-rich core, surrounded by a hydrogen-rich outer region. By contrast, cool, very low-mass stars (below 0.4 solar masses) are convective throughout.<ref name=science295_5552/> Thus the helium produced at the core is distributed across the star, producing a relatively uniform atmosphere and a proportionately longer main sequence lifespan.<ref name=brainerd/>
 
== Luminosity-color variation ==
As non-fusing helium ash accumulates in the core of a main-sequence star, the reduction in the abundance of hydrogen per unit mass results in a gradual lowering of the fusion rate within that mass. Since it is the outflow of fusion-supplied energy that supports the higher layers of the star, the core is compressed, producing higher temperatures and pressures. Both factors increase the rate of fusion thus moving the equilibrium towards a smaller, denser, hotter core producing more energy whose increased outflow pushes the higher layers further out. Thus there is a steady increase in the luminosity and radius of the star over time.<ref name="clayton83"/> For example, the luminosity of the early Sun was only about 70% of its current value.<ref name=sp74/> As a star ages this luminosity increase changes its position on the HR diagram. This effect results in a broadening of the main sequence band because stars are observed at random stages in their lifetime. That is, the main sequence band develops a thickness on the HR diagram; it is not simply a narrow line.<ref name=padmanabhan01/>
 
Other factors that broaden the main sequence band on the HR diagram include uncertainty in the distance to stars and the presence of unresolved [[binary star]]s that can alter the observed stellar parameters. However, even perfect observation would show a fuzzy main sequence because mass is not the only parameter that affects a star's color and luminosity. Variations in chemical composition caused by the initial abundances, the star's [[Stellar evolution|evolutionary status]],<ref name=apj128_3/> interaction with a [[Binary star|close companion]],<ref name=tayler94/> [[Stellar rotation|rapid rotation]],<ref name=mnras113/> or a [[Stellar magnetic field|magnetic field]] can all slightly change a main-sequence star's HR diagram position, to name just a few factors. As an example, there are [[Metallicity|metal-poor stars]] (with a very low abundance of elements with higher atomic numbers than helium) that lie just below the main sequence and are known as [[subdwarf]]s. These stars are fusing hydrogen in their cores and so they mark the lower edge of main sequence fuzziness caused by variance in chemical composition.<ref name=cwcs13/>
 
A nearly vertical region of the HR diagram, known as the [[instability strip]], is occupied by pulsating [[variable star]]s known as [[Cepheid variable]]s. These stars vary in magnitude at regular intervals, giving them a pulsating appearance. The strip intersects the upper part of the main sequence in the region of class ''A'' and ''F'' stars, which are between one and two solar masses. Pulsating stars in this part of the instability strip that intersects the upper part of the main sequence are called [[Delta Scuti variable]]s. Main-sequence stars in this region experience only small changes in magnitude and so this variation is difficult to detect.<ref name=green04/> Other classes of unstable main-sequence stars, like [[Beta Cephei variable]]s, are unrelated to this instability strip.
 
== Lifetime ==
[[Image:Isochrone ZAMS Z2pct.png|360px|right|thumb|This plot gives an example of the mass-luminosity relationship for zero-age main-sequence stars. The mass and luminosity are relative to the present-day Sun.]]
The total amount of energy that a star can generate through nuclear fusion of hydrogen is limited by the amount of hydrogen fuel that can be consumed at the core. For a star in equilibrium, the energy generated at the core must be at least equal to the energy radiated at the surface. Since the luminosity gives the amount of energy radiated per unit time, the total life span can be estimated, to [[Orders of approximation|first approximation]], as the total energy produced divided by the star's luminosity.<ref name=rit_ms/>
 
For a star with at least 0.5 solar masses, once the hydrogen supply in its core is exhausted and it expands to become a [[red giant]], it can start to fuse [[helium]] atoms to form [[carbon]]. The energy output of the helium fusion process per unit mass is only about a tenth the energy output of the hydrogen process, and the luminosity of the star increases.<ref name="prialnik00"/> This results in a much shorter length of time in this stage compared to the main sequence lifetime. (For example, the Sun is predicted to spend {{nowrap|130 million years}} burning helium, compared to about 12 billion years burning hydrogen.)<ref name=mnras386_1/> Thus, about 90% of the observed stars above 0.5 solar masses will be on the main sequence.<ref name=arnett96/> On average, main-sequence stars are known to follow an empirical [[Mass–luminosity relation|mass-luminosity relationship]].<ref name=lecchini07/> The luminosity (''L'') of the star is roughly proportional to the total mass (''M'') as the following [[power law]]:
 
:<math>\begin{smallmatrix}L\ \propto\ M^{3.5}\end{smallmatrix}</math>
 
This relationship applies to main-sequence stars in the range 0.1–50 solar masses.<ref name=rolfs_rodney88/>
 
The amount of fuel available for nuclear fusion is proportional to the mass of the star. Thus, the lifetime of a star on the main sequence can be estimated by comparing it to solar evolutionary models. The [[Sun]] has been a main-sequence star for about 4.5 billion years and it will become a red giant in 6.5 billion years,<ref name=apj418>{{cite journal
| last=Sackmann | first=I.-Juliana
| coauthors=Boothroyd, Arnold I.; Kraemer, Kathleen E.
| title=Our Sun. III. Present and Future
| journal=Astrophysical Journal |date=November 1993 | volume=418 | pages=457–468
| doi=10.1086/173407 | bibcode=1993ApJ...418..457S}}</ref> for a total main sequence lifetime of roughly 10<sup>10</sup> years. Hence:<ref name=hansen_kawaler94>{{cite book
| first=Carl J. | last=Hansen
| coauthors=Kawaler, Steven D. | year=1994
| title=Stellar Interiors: Physical Principles, Structure, and Evolution | page=28 | publisher=Birkhäuser
| isbn=0-387-94138-X }}</ref>
 
:<math>\begin{smallmatrix} \tau_{\rm MS}\ \approx \ 10^{10} \text{years} \cdot \left[ \frac{M}{M_{\bigodot}} \right] \cdot \left[ \frac{L_{\bigodot}}{L} \right]\ =\ 10^{10} \text{years} \cdot \left[ \frac{M}{M_{\bigodot}} \right]^{-2.5} \end{smallmatrix}</math>
 
where ''M'' and ''L'' are the mass and luminosity of the star, respectively, <math>\begin{smallmatrix}M_{\bigodot}\end{smallmatrix}</math> is a solar mass, <math>\begin{smallmatrix}L_{\bigodot}\end{smallmatrix}</math> is the [[solar luminosity]] and <math>\tau_{\rm MS}</math> is the star's estimated main sequence lifetime.
 
Although more massive stars have more fuel to burn and might be expected to last longer, they also must radiate a proportionately greater amount with increased mass. Thus, the most massive stars may remain on the main sequence for only a few million years, while stars with less than a tenth of a solar mass may last for over a trillion years.<ref name=apj482>{{cite journal
| last=Laughlin | first=Gregory
| coauthors=Bodenheimer, Peter; Adams, Fred C.
| title=The End of the Main Sequence
| journal=The Astrophysical Journal
| year=1997 | volume=482
| issue=1 | pages=420–432
| doi= 10.1086/304125 | bibcode=1997ApJ...482..420L}}</ref>
 
The exact mass-luminosity relationship depends on how efficiently energy can be transported from the core to the surface. A higher [[Opacity (optics)|opacity]] has an insulating effect that retains more energy at the core, so the star does not need to produce as much energy to remain in [[hydrostatic equilibrium]]. By contrast, a lower opacity means energy escapes more rapidly and the star must burn more fuel to remain in equilibrium.<ref name=imamura07>{{cite web
| last=Imamura | first=James N. | date=1995-02-07
| url=http://zebu.uoregon.edu/~imamura/208/feb6/mass.html
| title=Mass-Luminosity Relationship
| publisher=University of Oregon
| accessdate=2007-01-08 | archiveurl = http://web.archive.org/web/20061214065335/http://zebu.uoregon.edu/~imamura/208/feb6/mass.html| archivedate = December 14, 2006}}</ref> Note, however, that a sufficiently high opacity can result in energy transport via [[convection]], which changes the conditions needed to remain in equilibrium.<ref name="clayton83">{{cite book
| first=Donald D. | last=Clayton | year=1983
| title=Principles of Stellar Evolution and Nucleosynthesis | publisher=University of Chicago Press
| isbn=0-226-10953-4 }}</ref>
 
In high-mass main-sequence stars, the opacity is dominated by [[electron scattering]], which is nearly constant with increasing temperature. Thus the luminosity only increases as the cube of the star's mass.<ref name="prialnik00"/> For stars below 10 times the solar mass, the opacity becomes dependent on temperature, resulting in the luminosity varying approximately as the fourth power of the star's mass.<ref name=rolfs_rodney88>{{cite book
| first=Claus E. | last=Rolfs
| coauthors=Rodney, William S. | year=1988
| title=Cauldrons in the Cosmos: Nuclear Astrophysics
| publisher=University of Chicago Press
| isbn=0-226-72457-3 }}</ref> For very low mass stars, molecules in the atmosphere also contribute to the opacity. Below about 0.5 solar masses, the luminosity of the star varies as the mass to the power of 2.3, producing a flattening of the slope on a graph of mass versus luminosity. Even these refinements are only an approximation, however, and the mass-luminosity relation can vary depending on a star's composition.<ref name=science295_5552>{{cite journal
| last=Kroupa | first=Pavel
| title=The Initial Mass Function of Stars: Evidence for Uniformity in Variable Systems
| journal=Science | year=2002 | volume=295
| issue=5552 | pages=82–91 | url=http://www.sciencemag.org/cgi/content/full/295/5552/82
| accessdate=2007-12-03 | doi=10.1126/science.1067524
| pmid=11778039 |arxiv = astro-ph/0201098 |bibcode = 2002Sci...295...82K }}</ref>
 
== Evolutionary tracks ==
{{See also|Stellar evolution}}
[[Image:Open cluster HR diagram ages.gif|right|thumb|250px|This shows the [[Hertzsprung–Russell diagram]]s for two open clusters. [[NGC 188]] (blue) is older, and shows a lower turn off from the main sequence than that seen in [[Messier 67|M67]] (yellow). The dots outside the two sequences are mostly foreground and background stars with no relation to the clusters.]]
 
Once a main-sequence star consumes the hydrogen at its core, the loss of energy generation causes its gravitational collapse to resume. Stars with less than 0.23 solar masses,<ref name=romp69>{{cite journal
| author=Adams, Fred C.; Laughlin, Gregory
| title=A Dying Universe: The Long Term Fate and Evolution of Astrophysical Objects
| journal=Reviews of Modern Physics |date=April 1997 | volume=69 | issue=2 | pages=337–372
| doi=10.1103/RevModPhys.69.337 | bibcode=1997RvMP...69..337A|arxiv = astro-ph/9701131 }}</ref> are predicted to directly become [[white dwarf]]s once energy generation by nuclear fusion of hydrogen at their core comes to a halt. In stars between this threshold and 10 solar masses, the hydrogen surrounding the helium core reaches sufficient temperature and pressure to undergo fusion, forming a hydrogen-burning shell. In consequence of this change, the outer envelope of the star expands and decreases in temperature, turning it into a [[red giant]]. At this point the star is evolving off the main sequence and entering the giant branch. The path which the star now follows across the HR diagram, to the upper right of the main sequence, is called an evolutionary track.
 
The helium core of a red giant continues to collapse until it is entirely supported by [[electron degeneracy pressure]]—a [[quantum mechanics|quantum mechanical]] effect that restricts how closely matter can be compacted. For stars of more than about 0.5 solar masses,<ref name=nature433>{{cite journal
| author= Fynbo, Hans O. U. ''et al.''
| title=Revised rates for the stellar triple-α process from measurement of 12C nuclear resonances
| journal=Nature | year=2004 | volume=433 | pages=136–139
| doi=10.1038/nature03219
| pmid= 15650733
| issue= 7022 }}</ref>
the core eventually reaches a temperature where it becomes hot enough to burn helium into carbon via the [[triple alpha process]].<ref name=sitko00>{{cite web
| last=Sitko | first=Michael L. | date=2000-03-24
| url=http://www.physics.uc.edu/~sitko/Spring00/4-Starevol/starevol.html
| title=Stellar Structure and Evolution
| publisher=University of Cincinnati
| accessdate=2007-12-05
}}</ref><ref name=pmss_atoe>{{cite web
| author=Staff | date=2006-10-12 | url=http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_postmain.html
| title=Post-Main Sequence Stars
| publisher=Australia Telescope Outreach and Education
| accessdate=2008-01-08 }}</ref>
Stars with more than 5–7.5 solar masses can additionally fuse elements with higher atomic numbers.<ref name=aaas141>{{cite journal
| author=Girardi, L.; Bressan, A.; Bertelli, G.; Chiosi, C.
| title=Evolutionary tracks and isochrones for low- and intermediate-mass stars: From 0.15 to 7 M<sub>sun</sub>, and from Z=0.0004 to 0.03
| journal=Astronomy and Astrophysics Supplement
| year=2000 | volume=141
| issue=3 | pages=371–383
| doi=10.1051/aas:2000126
|arxiv = astro-ph/9910164 |bibcode = 2000A&AS..141..371G }}</ref><ref name=apj675_1>{{cite journal | author=Poelarends, A. J. T.; Herwig, F.; Langer, N.; Heger, A.
| title=The Supernova Channel of Super-AGB Stars
| journal=The Astrophysical Journal |date=March 2008 | volume=675 | issue=1 | pages=614–625
| doi=10.1086/520872 | bibcode=2008ApJ...675..614P|arxiv = 0705.4643 }}</ref>
For stars with ten or more solar masses, this process can lead to an increasingly dense core that finally collapses, ejecting the star's overlying layers in a [[Type II supernova]] explosion,<ref name="science304" /> [[Type Ib supernova]] or [[Type Ic supernova]].
 
When a [[star cluster|cluster of stars]] is formed at about the same time, the life span of these stars will depend on their individual masses. The most massive stars will leave the main sequence first, followed steadily in sequence by stars of ever lower masses. Thus the stars will evolve in order of their position on the main sequence, proceeding from the most massive at the left toward the right of the HR diagram. The current position where stars in this cluster are leaving the main sequence is known as the turn-off point. By knowing the main sequence lifespan of stars at this point, it becomes possible to estimate the age of the cluster.<ref name=science299_5603>{{cite journal
| last=Krauss | first=Lawrence M.
| coauthors=Chaboyer, Brian
| title=Age Estimates of Globular Clusters in the Milky Way: Constraints on Cosmology
| journal=Science
| year=2003 | volume=299
| issue=5603 | pages=65–69
| doi= 10.1126/science.1075631
| pmid=12511641 |bibcode = 2003Sci...299...65K }}</ref>
 
== See also ==
{{Portal|Star}}
* [[Hertzsprung–Russell diagram]]
* [[Hydrogen-burning process]]
 
== Notes ==
<div class="references-small">
<references group=note/>
</div>
 
==References==
{{Reflist|2|refs=
 
<ref name=smith91>{{Cite web | url=http://cass.ucsd.edu/public/tutorial/HR.html | title=The Hertzsprung-Russell Diagram | author=Harding E. Smith | date=1999-04-21 | work=Gene Smith's Astronomy Tutorial | publisher=Center for Astrophysics & Space Sciences, University of California, San Diego | accessdate=2009-10-29 }}</ref>
 
<ref name=powell06>{{Cite web | url=http://www.atlasoftheuniverse.com/hr.html | title=The Hertzsprung Russell Diagram | author=Richard Powell | year=2006 | work=An Atlas of the Universe | accessdate=2009-10-29 }}</ref>
 
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<ref name=schatzman33>{{cite book | first=Evry L. | last=Schatzman | year=1993 | coauthors=Praderie, Francoise | title=The Stars | pages=96–97 | publisher=Springer | isbn=3-540-54196-9 }}</ref>
 
<ref name=keenan_morgan43>{{cite book | first=W. W. | last=Morgan | coauthors=Keenan, P. C.; Kellman, E. | year=1943 | title=An atlas of stellar spectra, with an outline of spectral classification | publisher=The University of Chicago press | location=Chicago, Illinois | url=http://nedwww.ipac.caltech.edu/level5/ASS_Atlas/MK_contents.html | accessdate=2008-08-12 }}</ref>
 
<ref name=tnc>{{cite book | first=Albrecht | last=Unsöld | year=1969 | title=The New Cosmos | page=268 | publisher=Springer-Verlag New York Inc | isbn=0-387-90886-2 }}</ref>
 
<ref name=asr34_1>{{cite journal | last=Gloeckler | first=George | coauthors=Geiss, Johannes | title=Composition of the local interstellar medium as diagnosed with pickup ions  | journal=Advances in Space Research | year=2004 | volume=34 | issue=1 | pages=53–60 | bibcode=2004AdSpR..34...53G | doi=10.1016/j.asr.2003.02.054 }}</ref>
 
<ref name=science295_5552>{{cite journal | last=Kroupa | first=Pavel | title=The Initial Mass Function of Stars: Evidence for Uniformity in Variable Systems | journal=Science | date=2002-01-04 | volume=295 | issue=5552 | pages=82–91 | url=http://www.sciencemag.org/cgi/content/abstract/295/5552/82 | accessdate=2008-12-08 | pmid=11778039 | doi=10.1126/science.1067524 |arxiv = astro-ph/0201098 |bibcode = 2002Sci...295...82K }}</ref>
 
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<ref name=zams_sao>{{cite web | url=http://astronomy.swin.edu.au/cms/astro/cosmos/Z/Zero+Age+Main+Sequence | title=Zero Age Main Sequence | work=The SAO Encyclopedia of Astronomy | publisher=Swinburne University | accessdate=2007-12-09 }}</ref>
 
<ref name=mss_atoe>{{cite web | url=http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_mainsequence.html | title=Main Sequence Stars | publisher=Australia Telescope Outreach and Education | accessdate=2007-12-04 }}</ref>
 
<ref name=moore06>{{cite book | first=Patrick | last=Moore | authorlink=Patrick Moore | year=2006 | title=The Amateur Astronomer | publisher=Springer | isbn=1-85233-878-4 }}</ref>
 
<ref name=wd_sao>{{cite web | url=http://astronomy.swin.edu.au/cosmos/W/White+Dwarf | title=White Dwarf | work=COSMOS—The SAO Encyclopedia of Astronomy | publisher=Swinburne University | accessdate=2007-12-04 }}</ref>
 
<ref name=siess00>{{cite web | last=Siess | first=Lionel | year=2000 | url=http://www.astro.ulb.ac.be/~siess/WWWTools/Isochrones | title=Computation of Isochrones | publisher=Institut d'Astronomie et d'Astrophysique, Université libre de Bruxelles | accessdate=2007-12-06 }}—Compare, for example, the model isochrones generated for a ZAMS of 1.1 solar masses. This is listed in the table as 1.26 times the [[solar luminosity]]. At metallicity Z=0.01 the luminosity is 1.34 times solar luminosity. At metallicity Z=0.04 the luminosity is 0.89 times the solar luminosity.</ref>
 
<ref name=ohrd>{{cite web | url=http://astro.unl.edu/naap/hr/hr_background3.html | title=Origin of the Hertzsprung-Russell Diagram | publisher=University of Nebraska | accessdate=2007-12-06 }}</ref>
 
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<ref name=simbad>{{cite web | url=http://simbad.u-strasbg.fr/simbad/ | title=SIMBAD Astronomical Database | publisher=Centre de Données astronomiques de Strasbourg | accessdate=2008-11-21 }}</ref>
 
<ref name=apj129>{{cite journal | author=Luck, R. Earle; Heiter, Ulrike | title=Stars within 15 Parsecs: Abundances for a Northern Sample | journal=The Astronomical Journal | year=2005 | volume=129 | issue=2 | pages=1063–1083 | bibcode=2005AJ....129.1063L  | doi=10.1086/427250 }}</ref>
 
<ref name=recons>{{cite web | author=Staff | date=2008-01-01 | url=http://www.chara.gsu.edu/RECONS/TOP100.posted.htm | title=List of the Nearest Hundred Nearest Star Systems | publisher=Research Consortium on Nearby Stars | accessdate=2008-08-12 }}</ref>
 
<ref name=brainerd>{{cite web | last=Brainerd | first=Jerome James | date=2005-02-16 | url=http://www.astrophysicsspectator.com/topics/stars/MainSequence.html | title=Main-Sequence Stars | publisher=The Astrophysics Spectator | accessdate=2007-12-04 }}</ref>
 
<ref name=hannu>{{cite book | first=Hannu | last=Karttunen | year=2003 | title=Fundamental Astronomy | publisher=Springer | isbn=3-540-00179-4 }}</ref>
 
<ref name=clayton83>{{cite book | first=Donald D. | last=Clayton | year=1983 | title=Principles of Stellar Evolution and Nucleosynthesis | publisher=University of Chicago Press | isbn=0-226-10953-4 }}</ref>
 
<ref name=apj555>{{cite journal | author=Bahcall, John N.; Pinsonneault, M. H.; Basu, Sarbani | title=Solar Models: Current Epoch and Time Dependences, Neutrinos, and Helioseismological Properties | journal=The Astrophysical Journal | date=2001-07-10 | volume=555 | issue=2 | pages=990–1012 | bibcode=2003PhRvL..90m1301B | doi=10.1086/321493 |arxiv = astro-ph/0212331 }}</ref>
 
<ref name=maurizio05>{{cite book | first=Maurizio | last=Salaris | coauthors=Cassisi, Santi | year=2005 | title=Evolution of Stars and Stellar Populations | page=128 | publisher=John Wiley and Sons | isbn=0-470-09220-3 }}</ref>
 
<ref name=apj620_1>{{cite journal | last=Oey | first=M. S. | coauthors=Clarke, C. J. | title=Statistical Confirmation of a Stellar Upper Mass Limit | journal=The Astrophysical Journal | year=2005 | volume=620 | issue=1 | pages=L43–L46 | bibcode=2005ApJ...620L..43O | doi=10.1086/428396 |arxiv = astro-ph/0501135 }}</ref>
 
<ref name=apj162>{{cite journal | last=Ziebarth | first=Kenneth | title=On the Upper Mass Limit for Main-Sequence Stars | journal=Astrophysical Journal | year=1970 | volume=162 | pages=947–962 | bibcode=1970ApJ...162..947Z | doi=10.1086/150726 }}</ref>
 
<ref name=apj406_1>{{cite journal | author=Burrows, A.; Hubbard, W. B.; Saumon, D.; Lunine, J. I. | title=An expanded set of brown dwarf and very low mass star models | journal=Astrophysical Journal, Part 1 |date=March 1993 | volume=406 | issue=1 | pages=158–171 | doi=10.1086/172427 | bibcode=1993ApJ...406..158B}}</ref>
 
<ref name=aller91>{{cite book | first=Lawrence H. | last=Aller | year=1991 | title=Atoms, Stars, and Nebulae | publisher=Cambridge University Press | isbn=0-521-31040-7 }}</ref>
 
<ref name=aaa102_1>{{cite journal | author=Bressan, A. G.; Chiosi, C.; Bertelli, G. | title=Mass loss and overshooting in massive stars | journal=Astronomy and Astrophysics | year=1981 | volume=102 | issue=1 | pages=25–30 | bibcode=1981A&A...102...25B }}</ref>
 
<ref name=lockner06>{{cite web | last=Lochner | first=Jim | coauthors=Gibb, Meredith; Newman, Phil | date=2006-09-06 | url=http://imagine.gsfc.nasa.gov/docs/science/know_l2/stars.html | title=Stars | publisher=NASA | accessdate=2007-12-05 }}</ref>
 
<ref name=science295_5552>{{cite journal | last=Kroupa | first=Pavel | title=The Initial Mass Function of Stars: Evidence for Uniformity in Variable Systems | journal=Science | date=2002-01-04 | volume=295 | issue=5552 | pages=82–91 | url=http://www.sciencemag.org/cgi/content/full/295/5552/82 | accessdate=2008-11-21 | doi=10.1126/science.1067524 | pmid=11778039 |arxiv = astro-ph/0201098 |bibcode = 2002Sci...295...82K }}</ref>
 
<ref name=sp74>{{cite journal | last=Gough | first=D. O. | title=Solar interior structure and luminosity variations | journal=Solar Physics | year=1981 | volume=74 | issue=1 | pages=21–34 | bibcode=1981SoPh...74...21G | doi=10.1007/BF00151270 }}</ref>
 
<ref name=padmanabhan01>{{cite book  | first=Thanu | last=Padmanabhan | year=2001 | title=Theoretical Astrophysics | publisher=Cambridge University Press | isbn=0-521-56241-4 }}</ref>
 
<ref name=apj128_3>{{cite journal | last=Wright | first=J. T. | title=Do We Know of Any Maunder Minimum Stars? | journal=The Astronomical Journal | year=2004 | volume=128 | issue=3 | pages=1273–1278 | url=http://adsabs.harvard.edu/cgi-bin/bib_query?arXiv:astro-ph/0406338 | accessdate=2007-12-06 | doi=10.1086/423221 | bibcode=2004AJ....128.1273W|arxiv = astro-ph/0406338 }}</ref>
 
<ref name=tayler94>{{cite book | first=Roger John | last=Tayler | year=1994 | title=The Stars: Their Structure and Evolution | publisher=Cambridge University Press | isbn=0-521-45885-4 }}</ref>
 
<ref name=mnras113>{{cite journal | last=Sweet | coauthors=Roy, A. E. | first=I. P. A. | title=The structure of rotating stars | journal=Monthly Notices of the Royal Astronomical Society | year=1953 | volume=113 | pages=701–715 | bibcode=1953MNRAS.113..701S }}</ref>
 
<ref name=cwcs13>{{cite conference | last=Burgasser | first=Adam J. | coauthors=Kirkpatrick, J. Davy; Lepine, Sebastien | title=Spitzer Studies of Ultracool Subdwarfs: Metal-poor Late-type M, L and T Dwarfs | booktitle=Proceedings of the 13th Cambridge Workshop on Cool Stars, Stellar Systems and the Sun | pages=237 | publisher=Dordrecht, D. Reidel Publishing Co | date=July 5–9, 2004 | location=Hamburg, Germany | url=http://adsabs.harvard.edu/cgi-bin/bib_query?arXiv:astro-ph/0409178 | accessdate=2007-12-06 }}</ref>
 
<ref name=green04>{{cite book | first=S. F. | last=Green | coauthors=Jones, Mark Henry; Burnell, S. Jocelyn | year=2004 | title=An Introduction to the Sun and Stars | publisher=Cambridge University Press | isbn=0-521-54622-2 }}</ref>
 
<ref name=rit_ms>{{cite web | last=Richmond | first=Michael W. | date=2004-11-10 | url=http://spiff.rit.edu/classes/phys230/lectures/star_age/star_age.html | title=Stellar evolution on the main sequence  | publisher=Rochester Institute of Technology | accessdate=2007-12-03 }}</ref>
 
<ref name="prialnik00">{{cite book | first=Dina | last=Prialnik | year=2000 | title=An Introduction to the Theory of Stellar Structure and Evolution | publisher=Cambridge University Press | isbn=0-521-65937-X }}</ref>
 
<ref name=mnras386_1>{{cite journal | author=Schröder, K.-P.; Connon Smith, Robert | title=Distant future of the Sun and Earth revisited | journal=Monthly Notices of the Royal Astronomical Society |date=May 2008 | volume=386 | issue=1 | pages=155–163 | bibcode=2008MNRAS.386..155S | doi=10.1111/j.1365-2966.2008.13022.x |arxiv = 0801.4031 }}</ref>
 
<ref name=arnett96>{{cite book | first=David | last=Arnett | year=1996 | title=Supernovae and Nucleosynthesis: An Investigation of the History of Matter, from the Big Bang to the Present | publisher=Princeton University Press | isbn=0-691-01147-8 }}—Hydrogen fusion produces 8×10<sup>18</sup>&nbsp;[[erg]]/[[gram|g]] while helium fusion produces 8×10<sup>17</sup>&nbsp;erg/g.</ref>
 
<ref name=lecchini07>For a detailed historical reconstruction of the theoretical derivation of this relationship by Eddington in 1924, see: {{cite book | first=Stefano | last=Lecchini | year=2007 | title=How Dwarfs Became Giants. The Discovery of the Mass-Luminosity Relation | url=http://www.amazon.de/Dwarfs-Giants-Discovery-Mass-Luminosity-Relation/dp/3952288268 | publisher=Bern Studies in the History and Philosophy of Science | isbn=3-9522882-6-8}}</ref>
 
<ref name=rolfs_rodney88>{{cite book  | first=Claus E. | last=Rolfs | coauthors=Rodney, William S. | year=1988 | title=Cauldrons in the Cosmos: Nuclear Astrophysics | page=46 | publisher=University of Chicago Press | isbn=0-226-72457-3 }}</ref>
 
<ref name=simbad_ltt2151>{{cite web | url=http://simbad.u-strasbg.fr/simbad/sim-basic?Ident=Gliese | title=LTT 2151 – High proper-motion Star | publisher=Centre de Données astronomiques de Strasbourg | accessdate=2008-08-12 }}</ref>
 
}}
 
==External links==
* [http://astrosun2.astro.cornell.edu/academics/courses/astro101/herter/java/evolve/evolve.htm A java based applet for stellar evolution.]
* {{cite web
| last=Charity
| first=Mitchell
| date=2001-06-04
| url=http://www.vendian.org/mncharity/dir3/starcolor/
| title=What color are the stars?
| publisher=Vendian Systems
| accessdate=2008-11-26 }}
 
{{Star}}
 
{{Featured article}}
 
{{DEFAULTSORT:Main Sequence}}
[[Category:Hertzsprung–Russell classifications]]
[[Category:Main-sequence stars|*]]
[[Category:Star types]]
[[Category:Stellar evolution]]

Latest revision as of 09:30, 20 November 2014

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