Why Are Some Stars Not on the Main Sequence? Again

Changes to a star over its lifespan

Representative lifetimes of stars as a function of their masses

The change in size with time of a Dominicus-like star

Artist's depiction of the life cycle of a Sun-like star, starting as a primary-sequence star at lower left and then expanding through the subgiant and giant phases, until its outer envelope is expelled to form a planetary nebula at upper right

Chart of stellar evolution

Stellar development is the process past which a star changes over the course of time. Depending on the mass of the star, its lifetime tin range from a few million years for the virtually massive to trillions of years for the least massive, which is considerably longer than the age of the universe. The table shows the lifetimes of stars as a function of their masses.^[1] All stars are formed from collapsing clouds of gas and dust, ofttimes called nebulae or molecular clouds. Over the form of millions of years, these protostars settle down into a state of equilibrium, becoming what is known equally a chief-sequence star.

Nuclear fusion powers a star for most of its existence. Initially the energy is generated by the fusion of hydrogen atoms at the core of the main-sequence star. Later, as the preponderance of atoms at the core becomes helium, stars like the Sun begin to fuse hydrogen forth a spherical beat surrounding the core. This procedure causes the star to gradually grow in size, passing through the subgiant phase until it reaches the ruddy-giant phase. Stars with at to the lowest degree one-half the mass of the Sun can also begin to generate energy through the fusion of helium at their core, whereas more than-massive stars tin can fuse heavier elements along a series of concentric shells. Once a star similar the Sun has exhausted its nuclear fuel, its core collapses into a dense white dwarf and the outer layers are expelled as a planetary nebula. Stars with around 10 or more times the mass of the Lord's day tin explode in a supernova as their inert fe cores collapse into an extremely dense neutron star or black hole. Although the universe is not old enough for any of the smallest red dwarfs to have reached the end of their existence, stellar models advise they will slowly become brighter and hotter earlier running out of hydrogen fuel and becoming low-mass white dwarfs.^[2]

Stellar evolution is non studied by observing the life of a unmarried star, as almost stellar changes occur also slowly to be detected, fifty-fifty over many centuries. Instead, astrophysicists come to understand how stars evolve by observing numerous stars at various points in their lifetime, and by simulating stellar structure using calculator models.

Star formation [edit]

Simplistic representation of the stages of stellar evolution

Protostar [edit]

Schematic of stellar evolution

Stellar evolution starts with the gravitational collapse of a behemothic molecular cloud. Typical giant molecular clouds are roughly 100 light-years (nine.5×10¹⁴ km) beyond and comprise up to 6,000,000 solar masses (one.two×10³⁷ kg). Every bit it collapses, a behemothic molecular cloud breaks into smaller and smaller pieces. In each of these fragments, the collapsing gas releases gravitational potential energy as rut. As its temperature and pressure increase, a fragment condenses into a rotating brawl of superhot gas known equally a protostar.^[3] Filamentary structures are truly ubiquitous in the molecular deject. Dumbo molecular filaments will fragment into gravitationally bound cores, which are the precursors of stars. Continuous accretion of gas, geometrical angle, and magnetic fields may control the detailed fragmentation manner of the filaments. In supercritical filaments, observations take revealed quasi-periodic chains of dumbo cores with spacing comparable to the filament inner width, and embedded 2 protostars with gas outflows.^[four]

A protostar continues to grow past accretion of gas and dust from the molecular cloud, becoming a pre-chief-sequence star as it reaches its final mass. Further development is determined past its mass. Mass is typically compared to the mass of the Dominicus: 1.0M _☉ (2.0×ten³⁰ kg) ways 1 solar mass.

Protostars are encompassed in dust, and are thus more than readily visible at infrared wavelengths. Observations from the Wide-field Infrared Survey Explorer (WISE) have been especially of import for unveiling numerous galactic protostars and their parent star clusters.^{[five] [6]}

Brownish dwarfs and sub-stellar objects [edit]

Protostars with masses less than roughly 0.08M _☉ (1.6×ten²⁹ kg) never reach temperatures loftier enough for nuclear fusion of hydrogen to begin. These are known as dark-brown dwarfs. The International Astronomical Union defines chocolate-brown dwarfs as stars massive enough to fuse deuterium at some betoken in their lives (thirteen Jupiter masses (M _J), two.v × 10²⁸ kg, or 0.0125M _☉). Objects smaller than 13M _J are classified equally sub-brown dwarfs (but if they orbit around some other stellar object they are classified as planets).^[7] Both types, deuterium-burning and not, shine dimly and fade away slowly, cooling gradually over hundreds of millions of years.

Principal sequence stellar mass objects [edit]

For a more than-massive protostar, the cadre temperature volition eventually reach ten million kelvin, initiating the proton–proton concatenation reaction and assuasive hydrogen to fuse, kickoff to deuterium and then to helium. In stars of slightly over iOne thousand _☉ (two.0×ten³⁰ kg), the carbon–nitrogen–oxygen fusion reaction (CNO bike) contributes a large portion of the energy generation. The onset of nuclear fusion leads relatively apace to a hydrostatic equilibrium in which free energy released by the cadre maintains a loftier gas pressure, balancing the weight of the star's affair and preventing further gravitational collapse. The star thus evolves rapidly to a stable country, starting time the master-sequence phase of its evolution.

A new star will sit at a specific bespeak on the main sequence of the Hertzsprung–Russell diagram, with the main-sequence spectral type depending upon the mass of the star. Small, relatively cold, depression-mass red dwarfs fuse hydrogen slowly and will remain on the main sequence for hundreds of billions of years or longer, whereas massive, hot O-type stars will get out the master sequence after just a few 1000000 years. A mid-sized yellow dwarf star, like the Sun, volition remain on the principal sequence for about x billion years. The Sun is thought to be in the centre of its main sequence lifespan.

Mature stars [edit]

Internal structures of main-sequence stars, convection zones with arrowed cycles and radiative zones with ruby flashes. To the left a low-mass cerise dwarf, in the center a mid-sized yellow dwarf and at the correct a massive blue-white main-sequence star.

Eventually the star'south core exhausts its supply of hydrogen and the star begins to evolve off the principal sequence. Without the outward radiation pressure level generated by the fusion of hydrogen to counteract the force of gravity the core contracts until either electron degeneracy pressure becomes sufficient to oppose gravity or the core becomes hot enough (effectually 100 MK) for helium fusion to begin. Which of these happens first depends upon the star's mass.

Depression-mass stars [edit]

What happens after a low-mass star ceases to produce free energy through fusion has not been direct observed; the universe is effectually xiii.8 billion years old, which is less time (by several orders of magnitude, in some cases) than it takes for fusion to finish in such stars.

Recent astrophysical models suggest that red dwarfs of 0.1Yard _☉ may stay on the primary sequence for some six to twelve trillion years, gradually increasing in both temperature and luminosity, and have several hundred billion years more to collapse, slowly, into a white dwarf.^{[9] [10]} Such stars will not become red giants as the whole star is a convection zone and it volition non develop a degenerate helium core with a trounce burning hydrogen. Instead, hydrogen fusion volition go on until most the whole star is helium.

Slightly more than massive stars do expand into crimson giants, but their helium cores are not massive enough to reach the temperatures required for helium fusion so they never attain the tip of the red-giant branch. When hydrogen vanquish burning finishes, these stars move directly off the red-giant co-operative like a post-asymptotic-giant-branch (AGB) star, just at lower luminosity, to get a white dwarf.^[2] A star with an initial mass virtually 0.half dozenChiliad _☉ will exist able to reach temperatures high enough to fuse helium, and these "mid-sized" stars keep to farther stages of evolution beyond the cherry-red-giant branch.^[xi]

Mid-sized stars [edit]

The evolutionary rail of a solar mass, solar metallicity, star from principal sequence to post-AGB

Stars of roughly 0.6–10M _☉ become carmine giants, which are big not-main-sequence stars of stellar nomenclature K or Grand. Red giants lie along the right edge of the Hertzsprung–Russell diagram due to their red colour and big luminosity. Examples include Aldebaran in the constellation Taurus and Arcturus in the constellation of Boötes.

Mid-sized stars are red giants during two different phases of their post-main-sequence evolution: ruddy-giant-branch stars, with inert cores made of helium and hydrogen-burning shells, and asymptotic-giant-branch stars, with inert cores made of carbon and helium-burning shells within the hydrogen-burning shells.^[12] Between these two phases, stars spend a period on the horizontal branch with a helium-fusing core. Many of these helium-fusing stars cluster towards the cool finish of the horizontal branch as K-type giants and are referred to as reddish dodder giants.

Subgiant phase [edit]

When a star exhausts the hydrogen in its core, it leaves the main sequence and begins to fuse hydrogen in a shell outside the cadre. The core increases in mass equally the shell produces more helium. Depending on the mass of the helium core, this continues for several million to 1 or two billion years, with the star expanding and cooling at a similar or slightly lower luminosity to its principal sequence state. Somewhen either the core becomes degenerate, in stars around the mass of the sun, or the outer layers cool sufficiently to become opaque, in more massive stars. Either of these changes cause the hydrogen shell to increase in temperature and the luminosity of the star to increase, at which bespeak the star expands onto the red-giant co-operative.^[13]

Crimson-giant-branch phase [edit]

The expanding outer layers of the star are convective, with the textile being mixed by turbulence from nigh the fusing regions upwardly to the surface of the star. For all simply the everyman-mass stars, the fused material has remained deep in the stellar interior prior to this betoken, so the convecting envelope makes fusion products visible at the star's surface for the first time. At this stage of evolution, the results are subtle, with the largest furnishings, alterations to the isotopes of hydrogen and helium, beingness unobservable. The effects of the CNO wheel appear at the surface during the first dredge-up, with lower ¹²C/¹³C ratios and altered proportions of carbon and nitrogen. These are detectable with spectroscopy and have been measured for many evolved stars.

The helium core continues to grow on the ruddy-giant co-operative. It is no longer in thermal equilibrium, either degenerate or above the Schönberg–Chandrasekhar limit, and so it increases in temperature which causes the rate of fusion in the hydrogen shell to increase. The star increases in luminosity towards the tip of the scarlet-giant branch. Scarlet-giant-branch stars with a degenerate helium core all achieve the tip with very similar cadre masses and very similar luminosities, although the more massive of the ruby giants become hot plenty to ignite helium fusion before that point.

Horizontal co-operative [edit]

In the helium cores of stars in the 0.6 to two.0 solar mass range, which are largely supported by electron degeneracy pressure, helium fusion will ignite on a timescale of days in a helium flash. In the nondegenerate cores of more than massive stars, the ignition of helium fusion occurs relatively slowly with no flash.^[xiv] The nuclear power released during the helium flash is very large, on the order of x^eight times the luminosity of the Sun for a few days^[thirteen] and ten¹¹ times the luminosity of the Sun (roughly the luminosity of the Milky way Galaxy) for a few seconds.^[fifteen] However, the energy is consumed by the thermal expansion of the initially degenerate core and thus cannot be seen from exterior the star.^{[13] [15] [16]} Due to the expansion of the core, the hydrogen fusion in the overlying layers slows and full energy generation decreases. The star contracts, although not all the fashion to the principal sequence, and information technology migrates to the horizontal branch on the Hertzsprung–Russell diagram, gradually shrinking in radius and increasing its surface temperature.

Core helium flash stars evolve to the red end of the horizontal co-operative simply do not migrate to higher temperatures earlier they gain a degenerate carbon-oxygen cadre and start helium shell burning. These stars are often observed every bit a blood-red clump of stars in the colour-magnitude diagram of a cluster, hotter and less luminous than the red giants. Higher-mass stars with larger helium cores move along the horizontal branch to higher temperatures, some becoming unstable pulsating stars in the xanthous instability strip (RR Lyrae variables), whereas some become fifty-fifty hotter and can form a blue tail or blue hook to the horizontal branch. The morphology of the horizontal branch depends on parameters such as metallicity, age, and helium content, only the exact details are even so being modelled.^[17]

Asymptotic-giant-branch phase [edit]

After a star has consumed the helium at the cadre, hydrogen and helium fusion continues in shells around a hot core of carbon and oxygen. The star follows the asymptotic giant co-operative on the Hertzsprung–Russell diagram, paralleling the original red-giant evolution, only with even faster energy generation (which lasts for a shorter time).^[xviii] Although helium is beingness burnt in a shell, the majority of the free energy is produced by hydrogen burning in a shell further from the core of the star. Helium from these hydrogen burning shells drops towards the centre of the star and periodically the energy output from the helium beat increases dramatically. This is known as a thermal pulse and they occur towards the end of the asymptotic-behemothic-branch stage, sometimes even into the post-asymptotic-giant-branch phase. Depending on mass and composition, there may be several to hundreds of thermal pulses.

There is a stage on the ascent of the asymptotic-behemothic-branch where a deep convective zone forms and can bring carbon from the core to the surface. This is known as the 2nd dredge upwardly, and in some stars there may even exist a third dredge up. In this way a carbon star is formed, very cool and strongly reddened stars showing stiff carbon lines in their spectra. A process known every bit hot bottom burning may catechumen carbon into oxygen and nitrogen before it can be dredged to the surface, and the interaction betwixt these processes determines the observed luminosities and spectra of carbon stars in item clusters.^[nineteen]

Another well known class of asymptotic-giant-co-operative stars is the Mira variables, which pulsate with well-defined periods of tens to hundreds of days and large amplitudes upward to well-nigh 10 magnitudes (in the visual, total luminosity changes by a much smaller amount). In more-massive stars the stars get more luminous and the pulsation period is longer, leading to enhanced mass loss, and the stars become heavily obscured at visual wavelengths. These stars tin can be observed equally OH/IR stars, pulsating in the infrared and showing OH maser activity. These stars are clearly oxygen rich, in contrast to the carbon stars, just both must be produced by dredge ups.

Mail service-AGB [edit]

These mid-range stars ultimately reach the tip of the asymptotic-giant-branch and run out of fuel for shell called-for. They are not sufficiently massive to start total-scale carbon fusion, so they contract over again, going through a period of postal service-asymptotic-giant-branch superwind to produce a planetary nebula with an extremely hot central star. The central star and then cools to a white dwarf. The expelled gas is relatively rich in heavy elements created within the star and may be particularly oxygen or carbon enriched, depending on the type of the star. The gas builds upward in an expanding shell chosen a circumstellar envelope and cools every bit information technology moves away from the star, allowing grit particles and molecules to form. With the high infrared energy input from the central star, platonic conditions are formed in these circumstellar envelopes for maser excitation.

Information technology is possible for thermal pulses to be produced in one case mail service-asymptotic-giant-branch evolution has begun, producing a diverseness of unusual and poorly understood stars known as born-once again asymptotic-giant-branch stars.^[20] These may result in farthermost horizontal-branch stars (subdwarf B stars), hydrogen deficient post-asymptotic-giant-branch stars, variable planetary nebula central stars, and R Coronae Borealis variables.

Massive stars [edit]

Reconstructed paradigm of Antares, a red supergiant

In massive stars, the core is already large plenty at the onset of the hydrogen burning beat out that helium ignition will occur before electron degeneracy pressure has a chance to become prevalent. Thus, when these stars aggrandize and absurd, they do not brighten as dramatically every bit lower-mass stars; however, they were more luminous on the main sequence and they evolve to highly luminous supergiants. Their cores become massive plenty that they cannot support themselves by electron degeneracy and will eventually plummet to produce a neutron star or blackness hole.^{[ citation needed ]}

Supergiant evolution [edit]

Extremely massive stars (more than than approximately 40G _☉), which are very luminous and thus have very rapid stellar winds, lose mass so speedily due to radiations pressure that they tend to strip off their own envelopes before they can expand to become red supergiants, and thus retain extremely high surface temperatures (and blueish-white color) from their chief-sequence time onwards. The largest stars of the current generation are about 100-150M _☉ considering the outer layers would be expelled past the extreme radiation. Although lower-mass stars normally do non fire off their outer layers and so rapidly, they can too avoid becoming ruddy giants or reddish supergiants if they are in binary systems close enough so that the companion star strips off the envelope equally it expands, or if they rotate rapidly enough so that convection extends all the way from the core to the surface, resulting in the absenteeism of a separate core and envelope due to thorough mixing.^[21]

The onion-like layers of a massive, evolved star only earlier cadre plummet (not to scale)

The core of a massive star, defined every bit the region depleted of hydrogen, grows hotter and denser every bit it accretes material from the fusion of hydrogen exterior the core. In sufficiently massive stars, the cadre reaches temperatures and densities loftier enough to fuse carbon and heavier elements via the blastoff process. At the end of helium fusion, the core of a star consists primarily of carbon and oxygen. In stars heavier than nigh 8Thousand _☉, the carbon ignites and fuses to class neon, sodium, and magnesium. Stars somewhat less massive may partially ignite carbon, but they are unable to fully fuse the carbon before electron degeneracy sets in, and these stars volition eventually get out an oxygen-neon-magnesium white dwarf.^{[22] [23]}

The exact mass limit for full carbon burning depends on several factors such as metallicity and the detailed mass lost on the asymptotic giant branch, but is approximately 8-9M _☉.^[22] Subsequently carbon burning is complete, the core of these stars reaches about 2.5M _☉ and becomes hot enough for heavier elements to fuse. Before oxygen starts to fuse, neon begins to capture electrons which triggers neon burning. For a range of stars of approximately 8-12One thousand _☉, this process is unstable and creates delinquent fusion resulting in an electron capture supernova.^{[24] [23]}

In more massive stars, the fusion of neon gain without a runaway deflagration. This is followed in turn by complete oxygen burning and silicon burning, producing a cadre consisting largely of atomic number 26-superlative elements. Surrounding the core are shells of lighter elements still undergoing fusion. The timescale for consummate fusion of a carbon core to an iron core is so short, only a few hundred years, that the outer layers of the star are unable to react and the appearance of the star is largely unchanged. The fe cadre grows until it reaches an effective Chandrasekhar mass, higher than the formal Chandrasekhar mass due to various corrections for the relativistic effects, entropy, charge, and the surrounding envelope. The effective Chandrasekhar mass for an iron core varies from nigh 1.34M _☉ in the least massive carmine supergiants to more than 1.eight1000 _☉ in more massive stars. In one case this mass is reached, electrons begin to be captured into the fe-peak nuclei and the core becomes unable to support itself. The core collapses and the star is destroyed, either in a supernova or straight collapse to a black hole.^[23]

Supernova [edit]

The Crab Nebula, the shattered remnants of a star which exploded every bit a supernova visible in 1054 Advert

When the core of a massive star collapses, it will form a neutron star, or in the case of cores that exceed the Tolman–Oppenheimer–Volkoff limit, a blackness hole. Through a process that is not completely understood, some of the gravitational potential energy released past this cadre collapse is converted into a Blazon Ib, Type Ic, or Blazon II supernova. Information technology is known that the core collapse produces a massive surge of neutrinos, equally observed with supernova SN 1987A. The extremely energetic neutrinos fragment some nuclei; some of their energy is consumed in releasing nucleons, including neutrons, and some of their energy is transformed into heat and kinetic energy, thus augmenting the stupor moving ridge started by rebound of some of the infalling material from the collapse of the cadre. Electron capture in very dense parts of the infalling thing may produce additional neutrons. Considering some of the rebounding matter is bombarded by the neutrons, some of its nuclei capture them, creating a spectrum of heavier-than-iron textile including the radioactive elements upward to (and probable beyond) uranium.^[25] Although non-exploding cherry-red giants can produce significant quantities of elements heavier than iron using neutrons released in side reactions of before nuclear reactions, the abundance of elements heavier than atomic number 26 (and in particular, of certain isotopes of elements that have multiple stable or long-lived isotopes) produced in such reactions is quite different from that produced in a supernova. Neither affluence alone matches that found in the Solar System, so both supernovae and ejection of elements from red giants are required to explain the observed abundance of heavy elements and isotopes thereof.

The energy transferred from plummet of the cadre to rebounding material not just generates heavy elements, but provides for their acceleration well across escape velocity, thus causing a Type Ib, Type Ic, or Blazon 2 supernova. Current agreement of this energy transfer is still not satisfactory; although current estimator models of Type Ib, Type Ic, and Type 2 supernovae account for part of the energy transfer, they are not able to business relationship for plenty free energy transfer to produce the observed ejection of cloth.^[26] Yet, neutrino oscillations may play an of import function in the energy transfer problem as they not merely touch on the free energy available in a particular flavour of neutrinos simply too through other general-relativistic furnishings on neutrinos.^{[27] [28]}

Some evidence gained from assay of the mass and orbital parameters of binary neutron stars (which require two such supernovae) hints that the plummet of an oxygen-neon-magnesium core may produce a supernova that differs observably (in ways other than size) from a supernova produced by the plummet of an iron core.^[29]

The about massive stars that exist today may be completely destroyed past a supernova with an energy greatly exceeding its gravitational bounden free energy. This rare outcome, caused by pair-instability, leaves backside no black pigsty remnant.^[thirty] In the past history of the universe, some stars were even larger than the largest that exists today, and they would immediately collapse into a blackness hole at the end of their lives, due to photodisintegration.

Stellar evolution of low-mass (left bicycle) and high-mass (correct bike) stars, with examples in italics

Stellar remnants [edit]

After a star has burned out its fuel supply, its remnants tin can take one of iii forms, depending on the mass during its lifetime.

White and black dwarfs [edit]

For a star of 1Yard _☉, the resulting white dwarf is of well-nigh 0.sixOne thousand _☉, compressed into approximately the volume of the Earth. White dwarfs are stable considering the in pull of gravity is balanced past the degeneracy pressure of the star's electrons, a consequence of the Pauli exclusion principle. Electron degeneracy pressure provides a rather soft limit against further pinch; therefore, for a given chemical limerick, white dwarfs of college mass have a smaller volume. With no fuel left to burn, the star radiates its remaining oestrus into space for billions of years.

A white dwarf is very hot when it commencement forms, more than 100,000 One thousand at the surface and even hotter in its interior. Information technology is so hot that a lot of its energy is lost in the form of neutrinos for the kickoff ten one thousand thousand years of its existence, but will accept lost most of its energy later a billion years.^[31]

The chemical composition of the white dwarf depends upon its mass. A star of a few solar masses will ignite carbon fusion to course magnesium, neon, and smaller amounts of other elements, resulting in a white dwarf equanimous importantly of oxygen, neon, and magnesium, provided that information technology can lose enough mass to get below the Chandrasekhar limit (see beneath), and provided that the ignition of carbon is not so violent as to blow the star autonomously in a supernova.^[32] A star of mass on the guild of magnitude of the Sun volition be unable to ignite carbon fusion, and will produce a white dwarf composed importantly of carbon and oxygen, and of mass too low to collapse unless matter is added to it later (see below). A star of less than about half the mass of the Dominicus will exist unable to ignite helium fusion (equally noted earlier), and will produce a white dwarf equanimous chiefly of helium.

In the finish, all that remains is a cold dark mass sometimes called a black dwarf. Notwithstanding, the universe is not one-time enough for any blackness dwarfs to exist yet.

If the white dwarf's mass increases above the Chandrasekhar limit, which is ane.4M _☉ for a white dwarf equanimous chiefly of carbon, oxygen, neon, and/or magnesium, and then electron degeneracy force per unit area fails due to electron capture and the star collapses. Depending upon the chemical limerick and pre-collapse temperature in the middle, this volition lead either to collapse into a neutron star or runaway ignition of carbon and oxygen. Heavier elements favor continued core plummet, because they require a higher temperature to ignite, considering electron capture onto these elements and their fusion products is easier; higher core temperatures favor runaway nuclear reaction, which halts core collapse and leads to a Type Ia supernova.^[33] These supernovae may be many times brighter than the Type II supernova marker the death of a massive star, even though the latter has the greater total energy release. This instability to plummet means that no white dwarf more than massive than approximately 1.4One thousand _☉ can be (with a possible minor exception for very quickly spinning white dwarfs, whose centrifugal force due to rotation partially counteracts the weight of their thing). Mass transfer in a binary system may cause an initially stable white dwarf to surpass the Chandrasekhar limit.

If a white dwarf forms a close binary system with another star, hydrogen from the larger companion may accrete around and onto a white dwarf until information technology gets hot plenty to fuse in a runaway reaction at its surface, although the white dwarf remains beneath the Chandrasekhar limit. Such an explosion is termed a nova.

Neutron stars [edit]

Bubble-similar shock wave even so expanding from a supernova explosion fifteen,000 years agone

Ordinarily, atoms are generally electron clouds by volume, with very meaty nuclei at the center (proportionally, if atoms were the size of a football game stadium, their nuclei would exist the size of dust mites). When a stellar core collapses, the pressure causes electrons and protons to fuse by electron capture. Without electrons, which keep nuclei apart, the neutrons collapse into a dense ball (in some ways similar a giant atomic nucleus), with a sparse overlying layer of degenerate matter (chiefly iron unless matter of different limerick is added later). The neutrons resist farther pinch by the Pauli exclusion principle, in a way coordinating to electron degeneracy pressure, but stronger.

These stars, known as neutron stars, are extremely minor—on the society of radius 10 km, no bigger than the size of a big city—and are phenomenally dumbo. Their period of rotation shortens dramatically as the stars compress (due to conservation of angular momentum); observed rotational periods of neutron stars range from about one.5 milliseconds (over 600 revolutions per 2d) to several seconds.^[34] When these rapidly rotating stars' magnetic poles are aligned with the Earth, we detect a pulse of radiation each revolution. Such neutron stars are called pulsars, and were the outset neutron stars to be discovered. Though electromagnetic radiation detected from pulsars is most oftentimes in the form of radio waves, pulsars have also been detected at visible, Ten-ray, and gamma ray wavelengths.^[35]

Black holes [edit]

If the mass of the stellar remnant is high enough, the neutron degeneracy pressure volition be insufficient to preclude plummet beneath the Schwarzschild radius. The stellar remnant thus becomes a black pigsty. The mass at which this occurs is not known with certainty, but is currently estimated at between 2 and 3M _☉.

Black holes are predicted by the theory of general relativity. According to classical full general relativity, no matter or data can menstruation from the interior of a blackness pigsty to an outside observer, although quantum furnishings may allow deviations from this strict rule. The existence of black holes in the universe is well supported, both theoretically and by astronomical ascertainment.

Considering the cadre-collapse machinery of a supernova is, at nowadays, only partially understood, information technology is still non known whether information technology is possible for a star to plummet directly to a black hole without producing a visible supernova, or whether some supernovae initially form unstable neutron stars which then plummet into black holes; the exact relation between the initial mass of the star and the last remnant is also not completely sure. Resolution of these uncertainties requires the analysis of more supernovae and supernova remnants.

Models [edit]

A stellar evolutionary model is a mathematical model that can exist used to compute the evolutionary phases of a star from its formation until it becomes a remnant. The mass and chemical limerick of the star are used as the inputs, and the luminosity and surface temperature are the simply constraints. The model formulae are based upon the physical agreement of the star, unremarkably under the supposition of hydrostatic equilibrium. All-encompassing computer calculations are so run to determine the changing land of the star over fourth dimension, yielding a table of data that can be used to determine the evolutionary runway of the star across the Hertzsprung–Russell diagram, forth with other evolving properties.^[36] Authentic models can be used to estimate the current historic period of a star by comparing its physical properties with those of stars along a matching evolutionary track.^[37]

See too [edit]

• Galaxy formation and development – From a homogeneous beginning, the formation of the first galaxies, the way galaxies modify over time
• Chronology of the universe – History and future of the universe
• Template:Nature timeline
• Nucleosynthesis – Process that creates new atomic nuclei from pre-existing nucleons, primarily protons and neutrons
• Standard solar model
• Stellar population – Grouping of stars past similar metallicity (metallicity)
• Stellar rotation § After formation – Angular motion of a star about its axis – Rotations slow equally stars age
• Timeline of stellar astronomy

References [edit]

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8. ^ Prialnik (2000, Fig. 8.19, p. 174)
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13. ^ ^{a b c} Ryan & Norton (2010, p. 115)
14. ^ Ryan & Norton (2010, p. 125)
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• Hansen, Carl J.; Kawaler, Steven D.; Trimble, Virginia (2004). Stellar interiors: concrete principles, structure, and evolution (2nd ed.). Springer-Verlag. ISBN0-387-20089-4.
• Prialnik, Dina (2000). An Introduction to the Theory of Stellar Structure and Evolution. Cambridge University Printing. ISBN0-521-65065-8.
• Ryan, Sean G.; Norton, Andrew J. (2010). Stellar Evolution and Nucleosynthesis. Cambridge University Press. ISBN978-0-521-13320-three.

Farther reading [edit]

• Astronomy 606 (Stellar Structure and Evolution) lecture notes, Cole Miller, Department of Astronomy, University of Maryland
• Astronomy 162, Unit 2 (The Structure & Evolution of Stars) lecture notes, Richard W. Pogge, Department of Astronomy, Ohio State University

External links [edit]

• Stellar evolution simulator
• Pisa Stellar Models
• MESA stellar evolution codes (Modules for Experiments in Stellar Astrophysics)
• "The Life of Stars", BBC Radio four discussion with Paul Murdin, Janna Levin and Phil Charles (In Our Fourth dimension, Mar. 27, 2003)

silviathelly.blogspot.com

Source: https://en.wikipedia.org/wiki/Stellar_evolution

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