Nebular hypothesis
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223153727
2008-07-02T20:18:59Z
Trevor MacInnis
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rm Category:FA-Class articles,
{{Star formation}}
In [[cosmogony]], the '''nebular hypothesis''' is the most widely accepted model explaining the [[formation and evolution of the Solar System]]. It was first proposed in 1734 by [[Emanuel Swedenborg]].<ref name=Swedenborg1734/> Originally applied only to our own [[Solar System]], this method of planetary system formation is now thought to be at work throughout the [[universe]].<ref name=Montmerle2006/> The widely accepted modern variant of the nebular hypothesis is '''Solar Nebular Disk Model''' (SNDM) or simply '''Solar Nebular Model'''.<ref name=Woolfson1993/>
According to SNDM stars form in massive and dense clouds of [[molecular hydrogen]]—[[giant molecular cloud]]s (GMC). They are gravitationally unstable, and matter coalesces to smaller denser clumps within, which then proceed to collapse and form stars. Star formation is a complex process, which always produces a gaseous [[protoplanetary disk]] around the young star. This may give birth to planets in certain circumstances, which are not well known. Thus the formation of planetary systems is thought to be a natural result of star formation. A sun-like star usually takes around 100 million years to form.<ref name=Montmerle2006/>
The protoplanetary disk is an [[accretion disk]] which continues to feed the central star. Initially very hot, the disk later cools in what is known as the [[T tauri star]] stage; here, formation of small [[dust]] grains made of [[rock (geology)|rock]]s and ices is possible. The grains may eventually [[coagulation|coagulate]] into [[kilometer]] sized [[planetesimal]]s. If the disk is massive enough the runaway accretions begin, resulting in the rapid—100,000 to 300,000 years—formation of Moon- to Mars-sized [[planetary embryo]]s. Near the star, the planetary embryos go through a stage of violent mergers, producing a few [[terrestrial planet]]s. The last stage takes around 100 million to a billion years.<ref name=Montmerle2006/>
The formation of [[giant planet]]s is a more complicated process. It is thought to occur beyond the so called [[Frost line (astrophysics)|snow line]], where planetary embryos are mainly made of various ices. As a result they are several times more massive than in the inner part of the protoplanetary disk. What follows after the embryo formation is not completely clear. However, some embryos appear to continue to grow and eventually reach 5–10 Earth masses—the threshold value, which is necessary to begin accretion of the [[hydrogen]]–[[helium]] gas from the disk. The accumulation of gas by the core is initially a slow process, which continues for several million years, but after the forming protoplanet reaches about 30 Earth masses it accelerates and proceeds in a runaway manner. The [[Jupiter (planet)|Jupiter]] and [[Saturn (planet)|Saturn]]–like planets are thought to accumulate the bulk of their mass during only 10,000 years. The accretion stops when the gas is exhausted. The formed planets can migrate over long distances during or after their formation. The [[ice giant]]s like [[Uranus]] and [[Neptune (planet)|Neptune]] are thought to be failed cores, which formed too late when the disk had almost disappeared.<ref name=Montmerle2006/>
==History==
{{main article|History of Solar System formation and evolution hypotheses}}
The nebular hypothesis was first proposed in 1734 by [[Emanuel Swedenborg]].<ref name=Swedenborg1734>{{cite book|author=Swedenborg, Emanuel|year=1734|title=(Principia) Latin: Opera Philosophica et Mineralia (English: Philosophical and Mineralogical Works)|volume=I}}</ref> [[Immanuel Kant]], who was familiar with Swedenborg's work, developed the theory further in 1755.<ref name=Woolfson1993>{{cite journal|last=Woolfson|first=M.M.|title=Solar System – its origin and evolution|journal=Q. J. R. astr. Soc.|volume=34| pages=1–20|year=1993|url=http://adsabs.harvard.edu/abs/1993QJRAS..34....1W}}</ref> He argued that gaseous clouds—[[nebulae]], which slowly rotate, gradually collapse and flatten due to [[gravity]] and eventually form [[star]]s and [[planet]]s. A similar model was proposed in 1796 by [[Pierre-Simon Laplace]].<ref name=Woolfson1993/> It featured a contracting and cooling protosolar cloud—the protosolar nebula. As the nebular contracted, it flattened and shed rings of material, which later collapsed into the planets.<ref name=Woolfson1993/> While the Laplacian nebular model dominated in the 19th century, it encountered a number of difficulties. The main problem was [[angular momentum]] distribution between the Sun and planets. The planets have 99% of the momentum, and this fact could not be explained by the nebular model.<ref name=Woolfson1993/> As a result this theory of planet formation was largely abandoned at the beginning of the 20th century.
The fall of the Laplacian model stimulated scientists to find a replacement for it. During the 20th century many theories were proposed including the ''plantesimal theory'' of [[Thomas Chrowder Chamberlin|Thomas Chamberlin]] and [[Forest Ray Moulton|Forest Moulton]] (1901), ''tidal model'' of Jeans (1917), ''accretion model'' of [[Otto Schmidt]] (1944), ''protoplanet theory'' of [[William McCrea (astronomer)|William McCrea]] (1960) and finally ''capture theory'' of [[Michael Woolfson]].<ref name=Woolfson1993/> In 1978 [[Andrew Prentice]] resurrected the initial Laplacan ideas about planet formation and developed the ''modern Laplacian theory''.<ref name=Woolfson1993/> None of these attempts was completely successful and many of the proposed theories were descriptive.
The birth of the modern widely accepted theory of planetary formation—Solar Nebular Disk Model (SNDM)—can be traced to the works of Soviet astronomer [[Victor Safronov]].<ref name=NewScientist>{{cite web|url=http://space.newscientist.com/channel/solar-system/comets-asteroids/mg13117837.100|title=Birth of the planets: The Earth and its fellow planets may be survivors from a time when planets ricocheted around the Sun like ball bearings on a pinball table|publisher=New Scientist|author=Henbest, Nigel |year=1991|accessdate=2008-04-18}}</ref> His book ''Evolution of the protoplanetary cloud and formation of the Earth and the planets'',<ref name=Safronov1972>{{cite book |first=Viktor Sergeevich|last=Safronov|title=Evolution of the Protoplanetary Cloud and Formation of the Earth and the Planets|isbn= 0706512251|year=1972| publisher=Israel Program for Scientific Translations }}</ref> which was translated to English in 1972, had a long lasting effect on the way scientists think about the formation of the planets.<ref name=Safronov>{{cite journal|url=http://adsabs.harvard.edu/full/1989Metic..24..347W|title=Leonard Medal Citation for Victor Sergeevich Safronov|journal=Meteoritics|author=Wetherill, George W. |year=1989|volume=24|pages=347}}</ref> In this book almost all major problems of the planetary formation process were formulated and some of them solved. The Safronov's ideas were further developed in the works of [[George Wetherill]], who discovered ''runaway accretion''.<ref name=Woolfson1993/> While originally applied only to our own [[Solar System]], the SNDM was subsequently thought by theorists to be at work throughout the [[universe]]; over 280 [[extrasolar planet]]s have since been discovered in our [[galaxy]].
==Solar Nebular Model: achievements and problems==
===Achievements===
The star formation process naturally results in the appearance of [[accretion disk]]s around young stellar objects.<ref name=Andre1994/> At the age of about 1 million years, 100% of stars may have such disks.<ref name=Haisch2001/> This conclusion is supported by the discovery of the gaseous and dusty disks around [[protostar]]s and [[T Tauri star]]s as well as by theoretical considerations.<ref name=Padgett1999/> The observations of the disks show that the [[dust|dust grain]]s inside them grow in size on the short time scale (over thousands of years) producing 1 [[cm]] sized particles.<ref name=Kessler-Silacci2006/>
The accretion process, by which 1 km [[planetesimal]]s grow into 1,000 km sized bodies, is well understood now.<ref name=Kokubo2002/> This process develops inside any disk, where the number density of planetesimals is sufficiently high, and proceeds in a runaway manner. Growth later slows and continues as the oligarchic accretion. The end result is formation of [[planetary embryo]]s of varying sizes, which depend on the distance from the star.<ref name=Kokubo2002/> Various simulations have demonstrated that the merger of embryos in the inner part of the protoplanetary disk leads to the formation of a few Earth sized bodies. Thus the origin of [[terrestrial planet]]s is now considered to be an almost solved problem.<ref name=Raymond2006/>
===Problems===
The physics of accretion disks encounters some problems.<ref name=Wurchterl2004/> The most important one is how the material, which is accreted by the protostar, loses its [[angular momentum]]. The momentum is probably transported to the outer parts of the disk, but the precise mechanism of this transport is not well understood. The process or processes responsible for the disappearance of the disks are also poorly known.<ref name=Klahr2003/><ref name=Nakamoto1995/>
The formation of planetesimals is the biggest unsolved problem in the Nebular Disk Model. How 1 cm sized particles coalesce into 1 km planetesimals is a mystery. This mechanism appears to be the key to the question as to why some stars have planets, while others have nothing around them, even [[debris disk|dust belt]]s.<ref name=Youdin2002/>
The formation of [[giant planet]]s is another unsolved problem. Current theories are unable to explain how their cores can form fast enough to accumulate significant amounts of gas from the quickly disappearing protoplanetary disk.<ref name=Kokubo2002/><ref name=Inaba2003/> The mean lifetime of the disks, which are less than 10<sup>7</sup> years, appears to be shorter than the time necessary for the core formation.<ref name=Haisch2001/> Another problem of giant planet formation is their migration. Some calculations show that interaction with the disk can cause rapid inward migration, which, if not stopped, will result in the planets plunging into the star.<ref name= Papaloizou2007/>
==Formation of stars and protoplanetary disks==
===Protostars===
{{main|protostar}}
[[Image:Ssc2005-02b.jpg|right|thumb|300px|The visible-light (left) and infrared (right) views of the [[Trifid Nebula]]—a giant star-forming cloud of gas and dust located 5,400 light-years away in the constellation Sagittarius]]
[[Star]]s are thought to form inside [[molecular cloud|giant clouds]] of cold [[molecular hydrogen]]—[[giant molecular cloud]]s roughly 300,000 times the mass of the Sun and 20 [[parsec]]s in diameter.<ref name=Pudritz2002>{{cite journal|last=Pudritz|first=Ralph E.|title=Clustered Star Formation and the Origin of Stellar Masses|journal=Science|volume=295| pages=68–75|year=2002|doi=10.1126/science.1068298|url=http://www.sciencemag.org/cgi/content/full/295/5552/68}}</ref><ref name=Montmerle2006>{{cite journal|last=Montmerle|first=Thierry|coauthors=Augereau, Jean-Charles; Chaussidon, Marc et.al|title=Solar System Formation and Early Evolution: the First 100 Million Years|journal=Earth, Moon, and Planets|volume=98|publisher=Spinger|pages=39–95|year=2006|doi=10.1007/s11038-006-9087-5| url=http://adsabs.harvard.edu/abs/2006EM%26P...98...39M}}</ref> Over millions of years giant molecular clouds are prone to [[gravitational collapse|collapse]] and fragmentation.<ref name=Clark2005>{{cite journal|last=Clark|first=Paul C.|coauthors=Bonnell, Ian A.|title=The onset of collapse in turbulently supported molecular clouds|journal=Mon.Not.R.Astron.Soc.|volume=361| pages=2–16|year=2005|doi=10.1111/j.1365-2966.2005.09105.x|url=http://adsabs.harvard.edu/abs/2005MNRAS.361....2C}}</ref> These fragments then form small, dense cores which in turn collapse into stars.<ref name=Pudritz2002/> The cores range in mass from a fraction to several times that of the Sun and are called protostellar (protosolar) nebulae.<ref name=Montmerle2006/> They possess diameters of 0.01–01 pc (2,000–20,000 AU) and a particle number density of roughly 10,000 to 100,000 cm<sup>−3</sup>.{{Ref_label|A|a|none}}<ref name=Pudritz2002/><ref name=Motte1998>{{cite journal|last=Motte|first=F.|coauthors=Andre, P.; Neri, R.|title=The initial conditions of star formation in the ρ Ophiuchi main cloud: wide-field millimeter continuum mapping|journal=Astron. Astrophys.|volume=336|pages=150–172|year=1998| url=http://adsabs.harvard.edu/abs/1998A%26A...336..150M}}</ref>
The initial collapse of a solar-mass protostellar nebula takes around 100,000 years.<ref name=Pudritz2002/><ref name=Montmerle2006/> Every nebula begins it with a certain amount of [[angular momentum]]. Gas in the central part of the nebula, whose angular momentum is relatively low, undergoes fast compression and forms a hot [[hydrostatic]] (not contracting) core containing a small fraction of the mass of the original nebula.<ref name=Stahler1980/> This core forms the seed of what will become a star.<ref name=Stahler1980/><ref name=Montmerle2006/> As the collapse continues, conservation of angular momentum means that the rotation of the infalling envelop accelerates,<ref name=Nakamoto1995/><ref name=Yorke1999/> which largely prevents the gas from directly [[accretion (astrophysics)|accreting]] onto the central core. The gas is instead forced to spread outwards near its equatorial plane, forming a [[accretion disk|disk]], which in turn accretes onto the core.<ref name=Montmerle2006/><ref name=Nakamoto1995>{{cite journal|last=Nakamoto|first=Taishi|coauthors=Nakagawa, Yushitsugu|title=Formation, early evolution, and gravitational stability of protoplanetary disks|journal=The Astrophysical Journal|volume=421|pages=640–650|year=1994|doi=10.1086/173678| url=http://adsabs.harvard.edu/abs/1994ApJ...421..640N }}</ref><ref name=Yorke1999>{{cite journal|last=Yorke|first=Harold W.|coauthors=Bodenheimer, Peter|title=The formation of protostellar disks. III. The influence of gravitationally induced angular momentum transport on disk structure and appearance|journal=The Astrophysical Journal|volume=525|pages=330–342|year=1999|doi=10.1086/307867| url=http://adsabs.harvard.edu/abs/1999ApJ...525..330Y}}</ref> The core gradually grows in mass until it becomes a young hot [[protostar]].<ref name=Stahler1980/> At this stage, the protostar and its disk are heavily obscured by the infalling envelope and are not directly observable.<ref name=Andre1994/> In fact the remaining envelope's [[opacity]] is so high that even [[millimeter-wave]] radiation has trouble escaping from inside it.<ref name=Andre1994/><ref name=Montmerle2006/> Such objects are observed as very bright condensations, which emit mainly millimeter-wave and [[Terahertz radiation|submillimeter-wave]] radiation.<ref name=Motte1998/> They are classified as spectral Class 0 protostars.<ref name=Andre1994>{{cite journal|last=Andre|first=Philippe|coauthors=Montmerle, Thierry|title=From T Tauri stars protostars: circumstellar material and young stellar objects in the ρ Ophiuchi cloud|journal=The Astrophysical Journal|volume=420| pages=837–862|year=1994|doi=10.1086/173608| url=http://adsabs.harvard.edu/abs/1994ApJ...420..837A}}</ref> The collapse is often accompanied by [[bipolar outflow]]s—[[jet (gas)|jet]]s, which emanate along the [[rotation|rotational axis]] of the inferred disk. The jets are frequently observed in star-forming regions (see [[Herbig-Haro object|Herbig-Haro (HH) object]]s).<ref name=Lee2000>{{cite journal|last=Lee|first=Chin-Fei|coauthors=Mundy, Lee G.; Reipurth, Bo et.al.|title=CO outflows from young stars: confronting the jet and wind models|journal=The Astrophysical Journal|volume=542|pages=925–945|year=2000|doi=10.1086/317056|url=http://adsabs.harvard.edu/abs/2000ApJ...542..925L}}</ref> The luminosity of the Class 0 protostars is high— a protostar of the solar mass may radiate at up to 100 solar luminosities.<ref name=Andre1994/> Their main source of energy is [[gravitational collapse]]; at this stage the protostars do not fuse hydrogen.<ref name=Stahler1980>{{cite journal|last=Stahler|first=Steven W.|coauthors=Shu, Frank H.; Taam, Ronald E.|title=The evolution of protostars: II The hydrostatic core|journal=The Astrophysical Journal|volume=242|pages=226–241|year=1980|url=http://adsabs.harvard.edu/abs/1980ApJ...242..226S}}</ref><ref name=Stahler1988/>
[[Image:Ssc2003-06f.jpg|left|thumb|250px|Infrared image of the molecular outflow from an otherwise hidden newborn star HH 46/47]]
As the envelope's material continues to infall onto the disk, it eventually becomes thin and transparent and the young stellar object (YSO) becomes observable; initially in [[far-infrared]] light and later in the visible.<ref name=Motte1998/> Around this time the protostar begins to [[nuclear fusion|fuse]] [[deuterium]] and then ordinary hydrogen.<ref name=Stahler1988>{{cite journal|last=Stahler|first=Steven W.|title=Deuterium and the Stellar Birthline|journal=The Astrophysical Journal|volume=332|pages=804–825|year=1988|url=http://adsabs.harvard.edu/abs/1988ApJ...332..804S}}</ref> This birth of a new star occurs at approximately 100,000 years after the collapse has begun.<ref name=Montmerle2006 /> The external appearance of the YSO at this stage corresponds to the spectral class I protostars,<ref name=Andre1994/> which are also called young [[T Tauri star]]s or evolved protostars.<ref name=Andre1994/> By this time the forming star has already accreted much of its mass: the total mass of the disk and remaining envelope does not exceed 10–20% of the mass of the central YSO.<ref name=Motte1998/>
At the next stage the envelope completely disappears, having been gathered up by the disk, and the protostar becomes a classical T Tauri star.{{Ref_label|B|b|none}} This happens after about 1 million years.<ref name=Montmerle2006/> The mass of the disk around a classical T Tauri star is about 1–3% of the stellar mass, and it is accreted at the rate of between a 10 millionth to 1 billionth a solar mass per year.<ref name=Hartmann1998>{{cite journal|last=Hartmann|first=Lee|coauthors=Calvet, Nuria; Gullbring, Eric; D’Alessio, Paula|title=Accretion and the evolution of T Tauri disks|journal=The Astrophysical Journal|volume=495| pages=385–400|year=1998|doi=10.1086/305277|url=http://adsabs.harvard.edu/abs/1998ApJ...495..385H}}</ref> A pair of bipolar jets is usually present as well.<ref name=Shu1997/> The accretion explains all peculiar properties of classical T Tauri stars: strong [[flux]] in the [[emission line]]s (up to 100% of the intrinsic [[luminosity]] of the star), [[magnetic]] [[activity]], [[photometry (astronomy)|photometric]] [[variability]] and jets.<ref name=Muzerolle2001>{{cite journal|last=Muzerolle|first=James|coauthors=Calvet, Nuria; Hartmann, Lee|title= Emission-line diagnostics of T Tauri magnetospheric accretion. II. Improved model tests and insights into accretion physics|journal=The Astrophysical Journal|volume=550|pages=944–961|year=2001|doi=10.1086/319779| url=http://adsabs.harvard.edu/abs/2001ApJ...550..944M}}</ref> The emission lines actually form as the accreted gas hits the "surface" of the star, which happens around its [[magnetic pole]]s.<ref name=Muzerolle2001/> The jets are byproducts of accretion: they carry away excessive angular momentum. The classical T Tauri stage lasts about 10 million years.<ref name=Montmerle2006/> The disk eventually disappears due to accretion onto central star, planet formation, ejection by jets and [[photoevaporation]] by UV-radiation from the central star and nearby stars.<ref name=Adams2004>{{cite journal|last=Adams|first=Fred C.|coauthors=Hollenbach, David; Laughlin, Gregory; Gorti, Uma|title=Photoevaporation of circumstellar disks due to external far-ultraviolet radiation in stellar aggregates|journal=The Astrophysical Journal|volume=611|pages=360–379|year=2004|doi=10.1086/421989| url=http://adsabs.harvard.edu/abs/2004ApJ...611..360A}}</ref> As a result the young star becomes a weakly lined T Tauri star, which slowly, over timeframe of hundreds of millions of years, evolves into an ordinary sun-like star.<ref name=Stahler1980/>
===Protoplanetary disks===
{{See also|Protoplanetary disk|planetesimal}}
[[Image:M42proplyds.jpg|right|thumb|250px|A protoplanetary disk forming in the [[Orion Nebula]]]]
Under certain circumstances the disk, which can now be called protoplanetary, may give birth to a [[planetary system]].<ref name=Montmerle2006/> The [[protoplanetary disk]]s are ubiquitous around all Sun-like stars.<ref name=Megeath2005>{{cite journal|last=Megeath|first=S.T.|coauthors=Hartmann, L.; Luhmann, K.L.; Fazio, G.G.|title=Spitzer/IRAC photometry of the ρ Chameleontis association|journal=The Astrophysical Journal|volume=634|pages=L113–L116|year=2005|doi=10.1086/498503| url=http://adsabs.harvard.edu/abs/2005ApJ...634L.113M}}</ref><ref name=Haisch2001>{{cite journal|last=Haisch|first=Karl E.|coauthors=Lada, Elizabeth A.; Lada, Charles J.|title=Disk frequencies and lifetimes in young clusters|journal=The Astrophysical Journal|volume=553|pages=L153–L156|year=2001|doi=10.1086/320685| url=http://adsabs.harvard.edu/abs/2001ApJ...553L.153H}}</ref> They exist from the beginning of a star's formation, but at the earliest stages are unobservable due to the [[opacity]] of the surrounding envelope.<ref name=Andre1994/> The disk of a Class 0 [[protostar]] is thought to be massive and hot. It is an [[accretion (astrophysics)|accretion disk]], which feeds the central protostar.<ref name=Nakamoto1995/><ref name=Yorke1999/> The temperature can easily exceed 400 [[Kelvin|K]] inside 5 AU and 1,000 K inside 1 AU.<ref name=Chick1997>{{cite journal|last=Chick|first=Kenneth M.|coauthors=Cassen, Patrick|title=Thermal processing of interstellar dust grains in the primitive solar environment|journal=The Astrophysical Journal|volume=477|pages=398–409|year=1997|doi=10.1086/303700|url=http://adsabs.harvard.edu/abs/1997ApJ...477..398C}}</ref> The heating of the disk is primarily caused by the [[viscosity|viscous]] [[dissipation]] of [[turbulence]] in it and by the infall of the gas from the nebula.<ref name=Nakamoto1995/><ref name=Yorke1999/> The high [[temperature]] in the inner disk causes most of the [[volatiles|volatile]] material—water, organics, and some [[rock (geology)|rocks]] to evaporate, leaving only the most [[refractory]] elements like [[iron]]. The ice can survive only in the outer part of the disk.<ref name=Chick1997/>
The main problem in the physics of accretion disks is the generation of turbulence and the mechanism responsible for the high [[viscosity|effective viscosity]].<ref name=Montmerle2006/> The turbulent viscosity is thought to be responsible for the [[transport phenomena|transport]] of the mass to the central protostar and momentum to the periphery of the disk. This is vital for accretion, because the gas can be accreted by the central protostar only if it losses most of its angular momentum, which must be carried away by the small part of the gas drifting outwards.<ref name=Nakamoto1995/><ref name=Klahr2003>{{cite journal|last=Klahr|first=H.H.|coauthors=Bodenheimer, P.|title=Turbulence in accretion disks: vorticity generation and angular momentum transport via the global baroclinic instability|journal=The Astrophysical Journal|volume=582|pages=869–892| year=2003|doi=10.1086/344743|url=http://adsabs.harvard.edu/abs/2003ApJ...582..869K}}</ref> The result of this process is the growth of both the protostar and of the disk [[radius]], which can reach 1,000 AU if the initial angular momentum of the nebula is large enough.<ref name=Yorke1999/> Large disks are routinely observed in many star-forming regions such as the [[Orion nebula]].<ref name=Padgett1999>{{cite journal|last=Padgett|first=Deborah L.|coauthors=Brandner, Wolfgang; Stapelfeldt, Karl L. et.al.|title=Hubble space telescope/nicmos imaging of disks and envelopes around very young stars|journal=The Astronomical Journal|volume=117|pages=1490–1504|year=1999|doi=10.1086/300781| url=http://adsabs.harvard.edu/abs/1999AJ....117.1490P}}</ref>
The lifespan of the accretion disks is about 10 million years.<ref name=Haisch2001/> By the time the star reaches the classical T-Tauri stage, the disk becomes thinner and cools.<ref name=Hartmann1998/> Less volatile materials start to [[condensation|condense]] in the inner part of it forming dust grains, which have size of 0.1–1 μm and contain [[crystalline]] [[silicate]]s.<ref name=Kessler-Silacci2006/> The transport of the material from the outer disk can mix these newly formed [[cosmic dust|dust grain]]s with [[primordial]] ones, which contain organic matter and other volatiles. This mixing can explain some peculiarities in the composition of solar system bodies such as the presence of [[interstellar]] grains in the primitive [[meteorite]]s and refractory inclusions in comets.<ref name=Chick1997/>
Dust particles tend to stick to each other in the dense disk environment, leading to the formation of larger particles up to several centimeters in size.<ref name=Michikoshi2006>{{cite journal|last=Michikoshi|first=Shugo|coauthors=Inutsuka, Shu-ichiro|title=A two-fluid analysis of the kelvin-helmholtz instability in the dusty layer of a protoplanetary disk: a possible path toward planetesimal formation through gravitational instability|journal=The Astrophysical Journal|volume=641|pages=1131–1147|year=2006|doi=10.1086/499799| url=http://adsabs.harvard.edu/abs/2006ApJ...641.1131M}}</ref> The signatures of the dust processing and [[coagulation]] are observed in the infrared spectra of the young disks.<ref name=Kessler-Silacci2006>{{cite journal|last=Kessler-Silacci|first=Jacqueline|coauthors=Augereau, Jean-Charles; Dullemond, Cornelis P. et.al.|title= c2d SPITZER IRS spectra of disks around T Tauri stars. I. Silicate emission and grain growth |journal=The Astrophysical Journal|volume=639|pages=275–291|year=2006|doi=10.1086/300781|url=http://adsabs.harvard.edu/abs/1999AJ....117.1490P}}</ref> Further aggregation can lead to the formation of [[planetesimal]]s measuring 1 km across or larger, which are the building blocks of [[planet]]s.<ref name=Michikoshi2006/><ref name=Montmerle2006/> Planetesimal formation is another unsolved problem of disk physics, as simple sticking becomes ineffective as dust particles grow larger.<ref name=Youdin2002/> The favorite hypothesis is formation by the [[Jeans instability|gravitational instability]]. Particles several centimeters in size or larger slowly settle near the middle plane of the disk, forming a very thin—less than 100 km—and dense layer. This layer is gravitationally unstable and may fragment into numerous clumps, which in turn collapse into planetesimals.<ref name=Youdin2002>{{cite journal|last=Youdin|first=Andrew N.|coauthors=Shu, Frank N.|title=Planetesimal formation by gravitational instability|journal=The Astrophysical Journal|volume=580|pages=494–505|year=2002|doi=10.1086/343109| url=http://adsabs.harvard.edu/abs/2002ApJ...580..494Y}}</ref><ref name=Montmerle2006/>
Planetary formation can also be triggered by gravitational instability within the disk itself, which leads to its [[fragmentation]] into clumps. Some them, if they are dense enough, will [[gravitational collapse|collapse]],<ref name=Klahr2003/> which can lead to rapid formation of [[gas giant]] planets and even [[brown dwarf]]s at the timescale of 1,000 years.<ref name=Boss2003>{{cite journal|last=Boss|first=Alan P.|title=Rapid formation of outer giant planets by disk instability|journal=The Astrophysical Journal|volume=599|pages=577–581|year=2003|doi=10.1086/379163|url=http://adsabs.harvard.edu/abs/2003ApJ...599..577B}}</ref> However it is only possible in massive disks—more massive than 0.3 solar masses. In comparison typical disk masses are 0.01–0.03 solar masses. Because the massive disks are rare, this mechanism of the planet formation is thought to be infrequent.<ref name=Wurchterl2004>{{cite encyclopedia|last=Wurchterl|first=G.|title=Planet Formation Towards Estimating Galactic Habitability|encyclopedia=Astrobiology:Future Perspectives|year=2004|publisher=Kluwer Academic Publishers|editor=P. Ehrenfreund et al.|pages=67–96| url=http://www.springerlink.com/content/pr4rj4240383l585/}}</ref><ref name=Montmerle2006/>
The ultimate [[dissipation]] of protoplanetary disks is triggered by a number of different mechanisms. The inner part of the disk is either accreted by the star or ejected by the [[bipolar outflow|bipolar jet]]s,<ref name=Hartmann1998/><ref name=Shu1997>{{cite journal|last=Shu|first=Frank H.|coauthors=Shang, Hsian; Glassgold, Alfred E.; Lee, Typhoon|title=X-rays and Fluctuating X-Winds from Protostars|journal=Science |volume=277|pages=1475–1479|year=1997|doi=10.1126/science.277.5331.1475 |url=http://www.sciencemag.org/cgi/content/full/277/5331/1475}}</ref> whereas the outer part can evaporate under the star's powerful [[ultraviolet|UV]] [[radiation]] during the T Tauri stage<ref name=Font2004>{{cite journal|last=Font|first=Andreea S.|coauthors=McCarthy, Ian G.; Johnstone, Doug; Ballantyne, David R.|title=Photoevaporation of circumstellar disks around young stars|journal=The Astrophysical Journal|volume=607|pages=890–903|year=2004|doi=10.1086/383518| url=http://adsabs.harvard.edu/abs/2004ApJ...607..890F}}</ref> or by nearby stars.<ref name=Adams2004/> The gas in the central part can either be accreted or ejected by the growing planets, while the small dust particles are ejected by the [[radiation pressure]] of the central star. What is finally left is either a planetary system, a remnant disk of dust without planets, or nothing, if planetesimals failed to form.<ref name=Montmerle2006/>
Because planetesimals are so numerous, and spread throughout the protoplanetary disk, some survive the formation of a planetary system. [[Asteroid]]s are understood to be left-over planetesimals, gradually grinding each other down into smaller and smaller bits, while [[comet]]s are typically planetesimals from the farther reaches of a planetary system. Meteorites are samples of planetesimals that reach a planetary surface, and provide a great deal of information about the formation of our Solar System. Primitive-type meteorites are chunks of shattered low-mass planetesimals, where no thermal [[Planetary differentiation|differentiation]] took place, while processed-type meteorites are chunks from shattered massive planetesimals.<ref name=Bottke2005/>
==Formation of planets==
===Rocky planets ===
According to SNDM [[rocky planet]]s form in the inner part of the protoplanetary disk, where temperature is high enough to prevent [[condensation]] of water and other ices.<ref name=Raymond2007>{{cite journal|last=Raymond|first=Sean N.|coauthors=Quinn, Thomas; Lunine, Jonathan I.|title=High-resolution simulations of the final assembly of Earth-like planets 2: water delivery and planetary habitability|journal=Astrobiology|volume=7|pages=66–84|year=2007|doi=10.1089/ast.2006.06-0126|url=http://adsabs.harvard.edu/abs/2007AsBio...7...66R}}</ref> This results in coagulation of purely rocky grains and later in the formation of rocky planetesimals.{{Ref_label|C|c|none}}<ref name=Raymond2007/> Such conditions are thought to exist in the inner 3–4 AU part of the disk of a sun-like star.<ref name=Montmerle2006/>
After small planetesimals—about 1 km in diameter—have formed by one way or another, ''runaway accretion'' begins.<ref name=Kokubo2002/> It is called runaway because the mass growth rate is proportional to {{nowrap|R<sup>4</sup>~M<sup>4/3</sup>}}, where R and M are the radius and mass of the growing body, respectively.<ref name=Thommes2003/> It is obvious that the specific (divided by mass) growth accelerates as the mass increases. This leads to the preferential growth of larger bodies at the expense of smaller ones.<ref name=Kokubo2002/> The runaway accretion lasts between 10,000 and 100,000 years and ends when the largest bodies exceed approximately 1,000 km in diameter.<ref name=Kokubo2002>{{cite journal|last=Kokubo|first=Eiichiro|coauthors=Ida, Shigeru|title=Formation of protoplanet systems and diversity of planetary systems|journal=The Astrophysical Journal|volume=581|pages=666–680|year=2002| doi=10.1086/344105|url=http://adsabs.harvard.edu/abs/2002ApJ...581..666K}}</ref> Slowing of the accretion is caused by gravitational perturbations by large bodies on the remaining planetesimals.<ref name=Thommes2003/><ref name=Kokubo2002/> In addition, the influence of larger bodies stops further growth of smaller bodies.<ref name=Kokubo2002/>
The next stage is called ''oligarchic accretion''.<ref name=Kokubo2002/> It is characterized by the dominance of several hundred of the largest bodies—oligarchs, which continue to slowly accrete planetesimals.<ref name=Kokubo2002/> No body other than the oligarchs can grow.<ref name=Thommes2003/> At this stage the rate of accretion is proportional to R<sup>2</sup>, which the geometrical [[Cross section (geometry)|cross-section]] of an oligarch.<ref name=Thommes2003/> The specific accretion rate is proportional to {{nowrap|M<sup>−1/3</sup>}}; and it declines with the mass of the body. This allows smaller oligarchs to catch up to larger ones. The oligarchs are kept at the distance of about {{nowrap|10·H<sub>r</sub>}} ({{nowrap|H<sub>r</sub>}}={{nowrap|(M/3M<sub>s</sub>)<sup>1/3</sup>}} is [[Hill radius]] and M<sub>s</sub> is the mass of Sun) from each other by the influence of the remaining planetesimals.<ref name=Kokubo2002/> Their orbital eccentricities and inclinations remain small. The oligarchs continue to accrete until planetesimals are exhausted in the disk around them.<ref name=Kokubo2002/> Sometimes nearby oligarchs merge. The final mass of an oligarch depends on the distance from the star and surface density of planetesimals and is called the isolations mass.<ref name=Thommes2003/> For the rocky planets it is up to 0.1 of the Earth mass, or one [[Mars]] mass.<ref name=Montmerle2006/> The final result of the oligarchic stage is the formation of about 100 [[Moon]]- to Mars-sized planetary embryos uniformly spaced at about {{nowrap|10·H<sub>r</sub>}}.<ref name=Raymond2006/> They are thought to reside inside gaps in the disk and to be separated by rings of remaining planetesimals. This stage is thought to last a few hundred thousand years.<ref name=Kokubo2002/><ref name=Montmerle2006/>
The last stage of rocky planet formation is the ''merger stage''.<ref name=Montmerle2006/> It begins when only a small number of planetesimals remains and embryos become massive enough to perturb each other, which causes their orbits to become [[chaos|chaotic]].<ref name=Raymond2006>{{cite journal|last=Raymond|first=Sean N.|coauthors=Quinn, Thomas; Lunine, Jonathan I.|title=High-resolution simulations of the final assembly of earth-like planets 1: terrestrial accretion and dynamics|journal=Icarus|volume=183|pages=265–282|year=2006| doi=10.1016/j.icarus.2006.03.011|url=http://adsabs.harvard.edu/abs/2006Icar..183..265R}}</ref> During this stage embryos expel remaining planetesimals, and collide with each other. The result of this process, which lasts for 10 to 100 million years, is the formation of a limited number of Earth sized bodies. Simulations show that the number of surviving planets is on average from 2 to 5.<ref name=Petit2001/><ref name=Bottke2005/><ref name=Raymond2006/><ref name=Montmerle2006/> In the Solar System they may be represented by Earth and [[Venus]].<ref name=Raymond2006/> Formation of both planets required merging of approximately 10–20 embryos, while the equal number of them were thrown out of the Solar System.<ref name=Bottke2005/> Some the embryos, which originated in the [[asteroid belt]], are thought to have brought water to Earth.<ref name=Raymond2007/> Mars and [[Mercury (planet)|Mercury]] may be regarded as remaining embryos that survived that rivalry.<ref name=Bottke2005/> Rocky planets, which have managed to coalesce, settle eventually into more or less stable orbits, explaining why planetary systems are generally packed to the limit; or, in other words, why they always appear to be at the brink of instability.<ref name=Raymond2006/>
===Giant planets===
[[Image:Ssc2003-06i.jpg|right|thumb|250px|The dust disk around [[Fomalhaut]]—the brightest star in Piscis Austrini constellation. Asymmetry of the disk may be caused by a giant planet (or planets) orbiting the star.]]
The formation of [[giant planet]]s is an outstanding problem in the [[planetary science]]s.<ref name=Wurchterl2004/> In the framework of the Solar Nebular Model two theories for their formation exist. The first one is the ''disk instability model'', where giant planets form in the massive protoplanetary disks as a result of its [[gravity|gravitational]] fragmentation (see above).<ref name=Boss2003/> The disk instability may also lead to the formation of [[brown dwarf]]s, which are usually classified as stars. The second possibility is the ''core accretion model'', which is also known as the ''nucleated instability model''.<ref name=Wurchterl2004/> The latter scenario is thought to be the most promising one, because it can explain the formation of the giant planets in relatively low mass disks (less than 0.1 solar masses). In this model giant planet formation is divided into two stages: a) accretion of a core of approximately 10 Earth masses and b) accretion of gas from the protoplanetary disk.<ref name=Wurchterl2004/><ref name=Montmerle2006/>
Giant planet core formation is thought to proceed roughly along the lines of the terrestrial planet formation.<ref name=Kokubo2002/> It starts with planetesimals, which then undergo the runaway growth followed by the slower oligarchic stage.<ref name=Thommes2003>{{cite journal|last=Thommes|first=E.W.|coauthors=Duncan, M.J.; Levison, H.F.|title=Oligarchic growth of giant planets|journal=Icarus|volume=161|pages=431–455|year=2003| doi=10.1016/S0019-1035(02)00043-X|url=http://adsabs.harvard.edu/abs/2003Icar..161..431T}}</ref> Hypotheses do not predict a merger stage, due to the low probability of collisions between planetary embryos in the outer part of planetary systems.<ref name=Thommes2003/> An additional difference is the composition of the [[planetesimal]]s, which in the case of giant planets form beyond the so called [[Frost line (astrophysics)|snow line]] and consist mainly of ice—ice to rock ratio is about 4 to 1.<ref name=Inaba2003/> This enhances the mass of planetesimals four-fold. However the minimum mass nebular, which is capable of terrestrial planet formation, can only form 1–2 Earth mass cores at the distance of Jupiter (5 AU) within 10 million years.<ref name=Thommes2003/> The latter number represents an averages lifetime of gaseous disks around sun-like stars.<ref name=Haisch2001/> The proposed solutions include enhanced mass of the disk—a tenfold increase would suffice;<ref name=Thommes2003/> protoplanet migration, which allows the embryo to accrete more planetesimals;<ref name=Inaba2003/> and finally accretion enhancement due to [[drag (physics)|gas drag]] in the gaseous envelopes of the embryos.<ref name=Inaba2003/><ref name=Fortier2007>{{cite journal|last=Fortier|first=A.|coauthors=Benvenuto, A.G.|title=Oligarchic planetesimal accretion and giant planet formation|journal=Astron.Astrophys.|volume=473|pages=311–322|year=2007|doi=10.1051/0004-6361:20066729| url=http://adsabs.harvard.edu/abs/2007A%26A...473..311F}}</ref> Some combination of the above-mentioned ideas may explain the formation of the cores of gas giant planets such as [[Jupiter (planet)|Jupiter]] and perhaps even [[Saturn (planet)|Saturn]].<ref name=Wurchterl2004/> The formation of planets like [[Uranus]] and [[Neptune]] is more problematic, since no theory has been capable of providing for the in situ formation of their cores at the distance of 20–30 AU from the central star.<ref name=Montmerle2006/> To resolve this issue an idea has been brought forward that they initially accreted in the Jupiter-Saturn region and then were scattered and migrated to their present location.<ref name=Thommes1999>{{cite journal|last=Thommes|first=Edward W.|coauthors=Duncan, Martin J.; Levison, Harold F.|title=The formation of Uranus and Neptune in the Jupiter-Saturn region of the Solar System|journal=Nature|volume=402|pages=635–638| url=http://www.boulder.swri.edu/~hal/PDF/un-scat_nature.pdf|year=1999|doi=10.1038/45185|format=pdf}}</ref>
Once the cores are of sufficient mass (5–10 Earth masses), they begin to gather gas from the surrounding disk.<ref name=Montmerle2006/> Initially it is a slow process, which can increase the core masses up to 30 Earth masses in a few million years.<ref name=Inaba2003>{{cite journal|last=Inaba|first=S.|coauthors=Wetherill, G.W.; Ikoma, M.|title=Formation of gas giant planets: core accretion models with fragmentation and planetary envelope|journal=Icarus|volume=166|pages=46–62|year=2003|doi=10.1016/j.icarus.2003.08.001| url=http://isotope.colorado.edu/~astr5835/Inaba%20et%20al%202003.pdf|format=pdf}}</ref><ref name=Fortier2007/> After that the accretion rates increase dramatically and the remaining 90% of the mass is accumulated in approximately 10,000 years.<ref name=Fortier2007/> The accretion of the gas stops, when it is exhausted. This happens when a gap opens in the protoplanetary disk.<ref name=Papaloizou2007>{{cite encyclopedia|last=Papaloizou|first=J.C.B.|coauthors=Nelson, R.P.; Kley, W. et.al.|title=Disk-Planet Interactions During Planet Formation|encyclopedia=Protostars and Planets V|year=2007|publisher=Arizona Press|editor=Bo Reipurth; David Jewitt; Klaus Keil|url=http://adsabs.harvard.edu/abs/2007prpl.conf..655P}}</ref> In this model ice giants—Uranus and Neptune are failed cores that began gas accretion too late, when almost all gas had already disappeared. The post runaway gas accretion stage is characterized by migration of the newly formed giant planets and continued slow gas accretion.<ref name=Papaloizou2007/> Migration is caused by the interaction of the planet sitting in the gap with the remaining disk. It stops, when the protoplanetary disk disappears or when the end of the disk is attained. The latter case corresponds to the so called [[hot Jupiters]], which are likely to have stopped their migration, when they reached the inner hole in the protoplanetary disk.<ref name=Papaloizou2007/>
[[Image:Planet formation.jpg|left|thumb|250px|In this artist's conception, a planet spins through a clearing (gap) in a nearby star's dusty, planet-forming disc.]]
Giant planets can significantly influence [[terrestrial planet]] formation. The presence of giants tends to increase [[Orbital eccentricity| eccentricities]] and [[orbital inclination|inclination]]s of planetesimals and embryos in the terrestrial planet region (inside 4 AU in the Solar System).<ref name=Bottke2005/><ref name=Petit2001>{{cite journal|last=Petit|first=Jean-Marc|coauthors=Morbidelli, Alessandro|title=The Primordial Excitation and Clearing of the Asteroid Belt|journal=Icarus|volume=153| pages=338–347|year=2001|doi=10.1006/icar.2001.6702| url=http://www.gps.caltech.edu/classes/ge133/reading/asteroids.pdf|format=pdf}}</ref> On the one hand, if giant planets form too early they can slow or prevent inner planet accretion. On the other hand, if they form near the end of the oligarchic stage, as is thought to have happened in the Solar System, they will influence the merges of planetary embryos making them more violent.<ref name=Bottke2005/> As a result the number of terrestrial planets will decrease and they will be more massive.<ref name=Levinson2003>{{cite journal|last=Levison|first=Harold F.|coauthors=Agnor, Craig |title=The role of giant planets in terrestrial planet formation |journal=The Astronomical Journal|volume=125|pages=2692–2713|year=2003|doi=10.1086/374625|url=http://www.boulder.swri.edu/~hal/PDF/tfakess.pdf|format=pdf }}</ref> In addition, the size of the system will shrink, because terrestrial planets will form closer to the central star. In the Solar System the influence of giant planets, particularly that of [[Jupiter (planet)|Jupiter]], is thought to have been limited because they are relatively remote from the terrestrial planets.<ref name=Levinson2003/>
The region of a planetary system adjacent to the giant planets will be influenced in a different way.<ref name=Petit2001/> In such a region eccentricities of embryos may become so large that they may pass close to a giant planet. As a result they may and probably will be thrown out of the planetary system.{{Ref_label|D|d|none}}<ref name=Bottke2005>{{cite journal|last=Bottke|first=William F.|coauthors=Durda, Daniel D.; Nesvorny, David et.al.|title=Linking the collisional history of the main asteroid belt to its dynamical excitation and depletion |journal=Icarus|volume=179| pages=63–94|year=2005|doi=10.1016/j.icarus.2005.05.017 |url=http://www.boulder.swri.edu/~bottke/Reprints/Bottke_Icarus_2005_179_63-94_Linking_Collision_Dynamics_MB.pdf|format=pdf }}</ref><ref name=Petit2001/> If all embryos are removed then no planets will form in this region.<ref name=Petit2001/> An additional consequence is that a huge number of small planetesimals will remain, because giant planets along without help from embryos are incapable of clearing them all out. The total mass of remaining planetesimals will be small, because cumulative action of the embryos before their ejection and giant planets is still strong enough to remove 99% of the small bodies.<ref name=Bottke2005/> Such a region will eventually evolve in an [[asteroid belt]], which is a full analog of the main asteroid belt in the Solar System located at the distance 2 to 4 AU from the Sun.<ref name=Bottke2005/><ref name=Petit2001/>
== Meaning of ''accretion'' ==
Use of the term [[accretion disk]] for the [[protoplanetary disk]] leads to confusion over the planetary accretion process.
The protoplanetary disk is sometimes referred to as an accretion disk, because while the young [[T Tauri star|T Tauri]]-like protosun is still contracting, gaseous material may still be falling onto it, accreting on its surface from the disk's inner edge.<ref name=Yorke1999/>
However, that meaning should not be confused with the process of accretion forming the planets. In this context, accretion refers to the process of cooled, solidified grains of dust and ice orbiting the [[protostar]] in the protoplanetary disk, colliding and sticking together and gradually growing, up to and including the high energy collisions between sizable [[planetesimal]]s.<ref name=Kokubo2002/>
In addition, the [[giant planet]]s probably had accretion disks of their own, in the first meaning of the word. The clouds of captured hydrogen and helium gas contracted, spun up, flattened, and deposited gas onto the surface of each giant [[protoplanet]], while solid bodies within that disk accreted into the giant planet's regular moons.<ref name=Canup2002>{{cite journal|last=Canup|first=Robin M.|coauthors=Ward, William R.|title=Formation of the Galilean Satellites: Conditions of Accretion|journal=The Astronomical Journal|year=2002|volume=124|pages=3404–3423|doi=10.1086/344684| url=http://www.boulder.swri.edu/~robin/cw02final.pdf|format=pdf}}</ref>
==See also==
*[[Formation and evolution of the Solar System]]
*[[History of Earth]]
*[[Asteroid Belt]], [[Kuiper Belt]], and [[Oort Cloud]]
*[[Bok globule]], [[Herbig-Haro object]]
*[[T Tauri star]]
== Notes ==
<div class="references-small">
<ol type="a">
<li>{{Note_label|A|a|none}} Compare it with the particle number density of the air at the sea level—<math>2.8\;\times\; e^{19} </math> cm<sup>−3</sup>.
<li>{{Note_label|B|b|none}} The T Tauri stars are young stars with mass less than about 2.5 solar masses showing a heightened level of activity. They are divided into two classes: weakly lined and classical T Tauri stars.<ref name=Mohanty2005>{{cite journal|last=Mohanty|first=Subhanjoy|coauthors=Jayawardhana, Ray; Basri, Gibor|title=The T Tauri Phase down to Nearly Planetary Masses: Echelle Spectra of 82 Very Low Mass Stars and Brown Dwarfs|journal=The Astrophysical Journal|volume=626|pages=498–522|year=2005|doi=10.1086/429794|url=http://adsabs.harvard.edu/abs/2005ApJ...626..498M}}</ref> The latter have accretion disks and continue to accrete hot gas, which manifests itself by strong emission lines in their spectrum. The former do not possess accretion disks. Classical T Tauri stars evolve into weakly lined T Tauri stars.<ref name=Martin1994>{{cite journal|last=Martin|first=E.L.|coauthors=Rebolo, R.; Magazzu, A.; Pavlenko Ya.V.|title=Pre-main sequence lithium burning|journal=Astron. Astrophys.|volume=282|pages=503–517|year=1994| url=http://adsabs.harvard.edu/abs/1994A%26A...282..503M}}</ref>
<li>{{Note_label|C|c|none}} The [[planetesimal]]s near the outer edge of the terrestrial planet region—2.5 to 4 AU from the Sun—may accumulate some amount of ice. However the rocks will still dominate, like in the [[asteroid belt|outer main belt]] in the Solar System.<ref name=Raymond2007/>
<li>{{Note_label|D|d|none}} As a variant they may collide with the central star or a giant planet.<ref name=Petit2001/></li>
</ol>
</div>
==References==
{{reflist|2}}
[[Category:Sun]]
[[Category:Solar System]]
[[Category:Circumstellar discs]]
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