Editing Nebular hypothesis (section)

== Formation of stars and protoplanetary disks ==

=== Protostars ===
{{Main|Protostar}}
[[File: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&nbsp;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&nbsp;times the mass of the Sun ({{Solar mass|link=y}}) and 20&nbsp;[[parsec]]s in diameter.<ref name=Montmerle2006>{{cite journal|last=Montmerle|first=Thierry|author2=Augereau, Jean-Charles |author3=Chaussidon, Marc |display-authors=etal |title=Solar System Formation and Early Evolution: the First 100 Million Years|journal=Earth, Moon, and Planets|volume=98|issue=1–4|pages=39–95|date=2006|doi=10.1007/s11038-006-9087-5| bibcode=2006EM&P...98...39M|s2cid=120504344}}</ref><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|date=2002|doi=10.1126/science.1068298|url=http://www.sciencemag.org/cgi/content/full/295/5552/68|pmid=11778037|issue=5552|bibcode = 2002Sci...295...68P|s2cid=33585808}}</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.|author2=Bonnell, Ian A.|title=The onset of collapse in turbulently supported molecular clouds|journal=Mon. Not. R. Astron. Soc.|volume=361|issue=1| pages=2–16|date=2005|doi=10.1111/j.1365-2966.2005.09105.x|bibcode=2005MNRAS.361....2C|doi-access=free}}</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 0.01–0.1&nbsp;pc (2,000–20,000&nbsp;AU) and a [[particle number density]] of roughly 10,000 to 100,000&nbsp;cm<sup>−3</sup>.<ref group=lower-alpha>Compare it with the particle number density of the air at the sea level—{{val|2.8|e=19|u=cm<sup>−3</sup>}}.</ref><ref name=Pudritz2002 /><ref name=Motte1998>{{cite journal|last1=Motte|first1=F.|last2=Andre|first2=P.|last3=Neri|first3=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|date=1998| bibcode=1998A&A...336..150M}}</ref>

The initial collapse of a solar-mass protostellar nebula takes around 100,000&nbsp;years.<ref name=Montmerle2006 /><ref name=Pudritz2002 /> Every nebula begins with a certain amount of [[angular momentum]]. Gas in the central part of the nebula, with relatively low angular momentum, 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=Montmerle2006 /><ref name=Stahler1980 /> As the collapse continues, conservation of angular momentum means that the rotation of the infalling envelope 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|author2=Nakagawa, Yushitsugu|title=Formation, early evolution, and gravitational stability of protoplanetary disks|journal=The Astrophysical Journal|volume=421|pages=640–650|date=1994|doi=10.1086/173678| bibcode=1994ApJ...421..640N|doi-access=}}</ref><ref name=Yorke1999>{{cite journal|last=Yorke|first=Harold W.|author2=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|issue=1|pages=330–342|date=1999|doi=10.1086/307867| bibcode=1999ApJ...525..330Y|doi-access=free}}</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 (optics)|opacity]] is so high that even [[millimeter-wave]] radiation has trouble escaping from inside it.<ref name=Montmerle2006 /><ref name=Andre1994 /> 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&nbsp;0 protostars.<ref name=Andre1994>{{cite journal|last=Andre|first=Philippe|author2=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|date=1994|doi=10.1086/173608| bibcode=1994ApJ...420..837A|doi-access=free}}</ref> The collapse is often accompanied by [[bipolar outflow]]s—[[jet (gas)|jets]]—that emanate along the [[rotation]]al axis of the inferred disk. The jets are frequently observed in star-forming regions (see [[Herbig–Haro object|Herbig–Haro (HH) objects]]).<ref name=Lee2000>{{cite journal|last=Lee|first=Chin-Fei|author2=Mundy, Lee G. |author3=Reipurth, Bo |display-authors=etal |title=CO outflows from young stars: confronting the jet and wind models|journal=The Astrophysical Journal|volume=542|issue=2|pages=925–945|date=2000|doi=10.1086/317056|bibcode=2000ApJ...542..925L|s2cid=118351543 |doi-access=free}}</ref> The luminosity of the Class&nbsp;0 protostars is high&nbsp;— a solar-mass protostar may radiate at up to 100 solar luminosities.<ref name=Andre1994 /> The source of this energy is [[gravitational collapse]], as their cores are not yet hot enough to begin [[nuclear fusion]].<ref name=Stahler1980>{{cite journal|last=Stahler|first=Steven W.|author2=Shu, Frank H. |author3=Taam, Ronald E. |title=The evolution of protostars: II The hydrostatic core|journal=The Astrophysical Journal|volume=242|pages=226–241|date=1980|bibcode=1980ApJ...242..226S|doi=10.1086/158459|doi-access=free}}</ref><ref name=Stahler1988 />

[[File:Embedded Outflow in Herbig-Haro object HH 46 47.jpg|left|thumb|250px|Infrared image of the molecular outflow from an otherwise hidden newborn star HH 46/47]]
As the infall of its material onto the disk continues, the envelope 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]]. If the protostar is sufficiently massive (above 80 Jupiter masses ({{Jupiter mass|link=y}})), hydrogen fusion follows. Otherwise, if its mass is too low, the object becomes a [[brown dwarf]].<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|date=1988|bibcode=1988ApJ...332..804S|doi=10.1086/166694}}</ref> This birth of a new star occurs approximately 100,000&nbsp;years after the collapse begins.<ref name=Montmerle2006 /> Objects at this stage are known as Class I protostars,<ref name=Andre1994 /> which are also called young [[T Tauri star]]s, evolved protostars, or young stellar objects.<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.{{Refn|group=lower-alpha|The T Tauri stars are young stars with mass less than about {{Solar mass|2.5}} 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|author2=Jayawardhana, Ray |author3=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|issue=1|pages=498–522|date=2005|doi=10.1086/429794|bibcode=2005ApJ...626..498M|arxiv = astro-ph/0502155|s2cid=8462683}}</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 |last1=Martin |first1=E. L. |last2=Rebolo |first2=R. |last3=Magazzu |first3=A. |last4=Pavlenko |first4=Ya. V. |title=Pre-main sequence lithium burning |journal=Astron. Astrophys. |volume=282 |pages=503–517 |date=1994 |bibcode=1994A&A...282..503M |arxiv=astro-ph/9308047}}</ref>}} This happens after about 1&nbsp;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 a rate of 10<sup>−7</sup> to {{Solar mass|10<sup>−9</sup>}} per year.<ref name=Hartmann1998>{{cite journal|last=Hartmann|first=Lee|author2=Calvet, Nuria|author2-link=Nuria Calvet|author3=Gullbring, Eric|author4=D’Alessio, Paula|author4-link=Paola D'Alessio|title=Accretion and the evolution of T Tauri disks|journal=The Astrophysical Journal|volume=495|issue=1| pages=385–400|date=1998|doi=10.1086/305277|bibcode=1998ApJ...495..385H|doi-access=free}}</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]] [[variable star|variability]] and jets.<ref name=Muzerolle2001>{{cite journal |last1=Muzerolle |first1=James |last2=Calvet |first2=Nuria|author2-link=Nuria Calvet |last3=Hartmann |first3=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 |issue=2 |pages=944–961 |date=2001 |doi=10.1086/319779 |bibcode=2001ApJ...550..944M|doi-access=free }}</ref> The emission lines actually form as the accreted gas hits the "surface" of the star, which happens around its [[Poles of astronomical bodies#Magnetic poles|magnetic poles]].<ref name=Muzerolle2001 /> The jets are byproducts of accretion: they carry away excessive angular momentum. The classical T Tauri stage lasts about 10&nbsp;million years.<ref name=Montmerle2006 /> The disk eventually disappears due to accretion onto the 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.|author2=Hollenbach, David |author3=Laughlin, Gregory |author4= Gorti, Uma |title=Photoevaporation of circumstellar disks due to external far-ultraviolet radiation in stellar aggregates|journal=The Astrophysical Journal|volume=611|issue=1|pages=360–379|date=2004|doi=10.1086/421989| bibcode=2004ApJ...611..360A|arxiv = astro-ph/0404383|s2cid=16093937}}</ref> As a result, the young star becomes a [[Weak-lined T Tauri star|weakly lined T Tauri star]], which slowly, over hundreds of millions of years, evolves into an ordinary Sun-like star.<ref name=Stahler1980 />

=== Protoplanetary disks ===
{{See also|Protoplanetary disk|planetesimal}}
[[File:NASA-14114-HubbleSpaceTelescope-DebrisDisks-20140424.jpg|thumb|right|250px|[[Debris disks]] detected in [[Hubble Space Telescope|HST]] archival images of young stars, HD 141943 and HD 191089, using improved imaging processes (24 April 2014).<ref name="NASA-20140424">{{cite web |last1=Harrington |first1=J.D. |last2=Villard |first2=Ray |title=RELEASE 14–114 Astronomical Forensics Uncover Planetary Disks in NASA's Hubble Archive |url=http://www.nasa.gov/press/2014/april/astronomical-forensics-uncover-planetary-disks-in-nasas-hubble-archive |date=24 April 2014 |work=[[NASA]] |url-status=live |archive-date=2014-04-25 |archive-url=https://web.archive.org/web/20140425125432/http://www.nasa.gov/press/2014/april/astronomical-forensics-uncover-planetary-disks-in-nasas-hubble-archive/ |access-date=2014-04-25}}</ref>]]
Under certain circumstances the disk, which can now be called protoplanetary, may give birth to a [[planetary system]].<ref name=Montmerle2006 /> Protoplanetary disks have been observed around a very high fraction of stars in young [[star clusters]].<ref name=Haisch2001>{{cite journal|last=Haisch|first=Karl E.|author2=Lada, Elizabeth A. |author3=Lada, Charles J. |title=Disk frequencies and lifetimes in young clusters|journal=The Astrophysical Journal|volume=553|issue=2|pages=L153–L156|date=2001|doi=10.1086/320685| bibcode=2001ApJ...553L.153H|arxiv = astro-ph/0104347|s2cid=16480998}}</ref><ref name=Megeath2005>{{cite journal|last=Megeath|first=S.T.|author2=Hartmann, L. |author3=Luhmann, K.L. |author4= Fazio, G.G. |title=Spitzer/IRAC photometry of the ρ Chameleontis association|journal=The Astrophysical Journal|volume=634|issue=1|pages=L113–L116|date=2005|doi=10.1086/498503| bibcode=2005ApJ...634L.113M|arxiv = astro-ph/0511314|s2cid=119007015}}</ref> They exist from the beginning of a star's formation, but at the earliest stages are unobservable due to the [[Opacity (optics)|opacity]] of the surrounding envelope.<ref name=Andre1994 /> The disk of a Class&nbsp;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&nbsp;[[Kelvin|K]] inside 5&nbsp;AU and 1,000&nbsp;K inside 1&nbsp;AU.<ref name=Chick1997>{{cite journal|last=Chick|first=Kenneth M.|author2=Cassen, Patrick|title=Thermal processing of interstellar dust grains in the primitive solar environment|journal=The Astrophysical Journal|volume=477|issue=1|pages=398–409|date=1997|doi=10.1086/303700|bibcode=1997ApJ...477..398C|doi-access=free}}</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 [[Volatile (astrogeology)|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 />
[[File:M42proplyds.jpg|thumb|left|250px|A protoplanetary disk forming in the [[Orion Nebula]]]]
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 loses 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.|author2=Bodenheimer, P.|title=Turbulence in accretion disks: vorticity generation and angular momentum transport via the global baroclinic instability|journal=The Astrophysical Journal|volume=582|issue=2|pages=869–892| date=2003|doi=10.1086/344743|bibcode=2003ApJ...582..869K|arxiv = astro-ph/0211629|s2cid=119362731}}</ref> The result of this process is the growth of both the protostar and of the disk [[radius]], which can reach 1,000&nbsp;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.|author2=Brandner, Wolfgang|author3=Stapelfeldt, Karl L.|display-authors=etal|title=Hubble space telescope/nicmos imaging of disks and envelopes around very young stars|journal=The Astronomical Journal|volume=117|issue=3|pages=1490–1504|date=1999|doi=10.1086/300781| bibcode=1999AJ....117.1490P|arxiv = astro-ph/9902101|s2cid=16498360}}</ref>
[[File:Artist’s impression of the disc and gas streams around HD 142527 (Animation).ogv|thumb|300px|Artist's impression of the disc and gas streams around young star [[HD 142527]].<ref>{{cite news|title=ALMA Sheds Light on Planet-Forming Gas Streams|url=http://www.eso.org/public/news/eso1301/|access-date=10 January 2013|newspaper=ESO Press Release}}</ref> ]]
The lifespan of the accretion disks is about 10&nbsp;million&nbsp;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]] close to its center, forming 0.1–1&nbsp;μm dust grains that 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 grains]] with [[primordial elements|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 dust|interstellar]] grains in primitive [[meteorite]]s and refractory inclusions in comets.<ref name=Chick1997 />
[[File:NASA-ExocometsAroundBetaPictoris-ArtistView.jpg|thumb|350px|left|Various [[planet formation]] processes, including [[exocomets]] and other [[planetesimal]]s, around [[Beta Pictoris]], a very young type [[A V star]] ([[NASA]] artist's conception).]]
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|author2=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|issue=2|pages=1131–1147|date=2006|doi=10.1086/499799| bibcode=2006ApJ...641.1131M|arxiv=astro-ph/0412643|s2cid=15477674}}</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|author2=Augereau, Jean-Charles|author3=Dullemond, Cornelis P.|display-authors=etal|title= c2d SPITZER IRS spectra of disks around T Tauri stars. I. Silicate emission and grain growth |journal=The Astrophysical Journal|volume=639|issue=3|pages=275–291|date=2006|doi=10.1086/499330 |arxiv = astro-ph/0511092 |bibcode = 2006ApJ...639..275K|s2cid=118938125}}</ref> Further aggregation can lead to the formation of [[planetesimal]]s measuring 1&nbsp;km across or larger, which are the building blocks of [[planet]]s.<ref name=Montmerle2006 /><ref name=Michikoshi2006 /> Planetesimal formation is another unsolved problem of disk physics, as simple sticking becomes ineffective as dust particles grow larger.<ref name=Youdin2002 />

One hypothesis is formation by [[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&nbsp;km—and dense layer. This layer is gravitationally unstable and may fragment into numerous clumps, which in turn collapse into planetesimals.<ref name=Montmerle2006 /><ref name=Youdin2002>{{cite journal|last=Youdin|first=Andrew N.|author2=Shu, Frank N.|title=Planetesimal formation by gravitational instability|journal=The Astrophysical Journal|volume=580|issue=1|pages=494–505|date=2002|doi=10.1086/343109| bibcode=2002ApJ...580..494Y|arxiv = astro-ph/0207536|s2cid=299829}}</ref> However, the differing velocities of the gas disk and the solids near the mid-plane can generate turbulence which prevents the layer from becoming thin enough to fragment due to gravitational instability.<ref name="Johansen_etal_2006">{{cite journal|last1=Johansen|first1=Anders|last2=Henning|first2=Thomas|last3=Klahr|first3=Hubert|title=Dust Sedimentation and Self-sustained Kelvin-Helmholtz Turbulence in Protoplanetary Disk Midplanes|journal=The Astrophysical Journal|date=2006|volume=643|issue=2|pages=1219–1232|doi=10.1086/502968|arxiv=astro-ph/0512272|bibcode = 2006ApJ...643.1219J|s2cid=15999094}}</ref> This may limit the formation of planetesimals via gravitational instabilities to specific locations in the disk where the concentration of solids is enhanced.<ref name="Protostars_and Planets_2014">{{cite book |last1=Johansen |first1=A. |last2=Blum |first2=J. |last3=Tanaka |first3=H. |last4=Ormel |first4=C. |last5=Bizzarro |first5=M. |last6=Rickman |first6=H. |title=Protostars and Planets VI |date=2014 |chapter=The Multifaceted Planetesimal Formation Process |editor1-last=Beuther |editor1-first=H. |editor2-last=Klessen |editor2-first=R. S. |editor3-last=Dullemond |editor3-first=C. P. |editor4-last=Henning |editor4-first=T. |pages=547–570 |publisher=University of Arizona Press |arxiv=1402.1344 |bibcode=2014prpl.conf..547J |doi=10.2458/azu_uapress_9780816531240-ch024 |isbn=978-0-8165-3124-0|s2cid=119300087 }}</ref>

Another possible mechanism for the formation of planetesimals is the [[streaming instability]] in which the drag felt by particles orbiting through gas creates a feedback effect causing the growth of local concentrations. These local concentrations push back on the gas creating a region where the headwind felt by the particles is smaller. The concentration is thus able to orbit faster and undergoes less radial drift. Isolated particles join these concentrations as they are overtaken or as they drift inward causing it to grow in mass. Eventually these concentrations form massive filaments which fragment and undergo gravitational collapse forming planetesimals the size of the larger asteroids.<ref name=Johansen_Jacquet_2015>{{cite book |last1=Johansen |first1=A. |last2=Jacquet |first2=E. |last3=Cuzzi |first3=J. N. |last4=Morbidelli |first4=A. |last5=Gounelle |first5=M. |date=2015 |chapter=New Paradigms For Asteroid Formation |editor1-last=Michel |editor1-first=P. |editor2-last=DeMeo |editor2-first=F. |editor3-last=Bottke |editor3-first=W. |title=Asteroids IV |pages=471 |publisher=University of Arizona Press |series=Space Science Series |arxiv=1505.02941 |bibcode=2015aste.book..471J |isbn=978-0-8165-3213-1|doi=10.2458/azu_uapress_9780816532131-ch025 |s2cid=118709894 }}</ref>

Planetary formation can also be triggered by gravitational instability within the disk itself, which leads to its fragmentation into clumps. Some of 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 on the timescale of 1,000&nbsp;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|issue=1|pages=577–581|date=2003|doi=10.1086/379163|bibcode=2003ApJ...599..577B|doi-access=free}}</ref> If these clumps migrate inward as the collapse proceeds tidal forces from the star can result in a significant [[tidal downsizing|mass loss]] leaving behind a smaller body.<ref name=Nayaksin_2010>{{cite journal|last1=Nayakshin|first1=Sergie|title=Formation of planets by tidal downsizing of giant planet embryos|journal=Monthly Notices of the Royal Astronomical Society Letters |date=2010|volume=408|issue=1|page=L36–L40|doi=10.1111/j.1745-3933.2010.00923.x|doi-access=free |arxiv=1007.4159|bibcode=2010MNRAS.408L..36N|s2cid=53409577}}</ref> However it is only possible in massive disks—more massive than {{Solar mass|0.3}}. In comparison, typical disk masses are {{Solar mass|0.01–0.03}}. Because the massive disks are rare, this mechanism of planet formation is thought to be infrequent.<ref name=Montmerle2006 /><ref name=Wurchterl2004>{{cite encyclopedia|last=Wurchterl|first=G.|title=Planet Formation Towards Estimating Galactic Habitability|encyclopedia=Astrobiology:Future Perspectives|date=2004|publisher=Kluwer Academic Publishers|editor=P. Ehrenfreund |display-editors=etal|pages=67–96|isbn=9781402023040 |doi=10.1007/1-4020-2305-7|chapter=Planet Formation|series=Astrophysics and Space Science Library|chapter-url=https://cds.cern.ch/record/1338806}}</ref> On the other hand, it may play a major role in the formation of [[brown dwarf]]s.<ref name=Stamatellosetal2007>{{cite journal|last=Stamatellos|first=Dimitris|author2=Hubber, David A. |author3=Whitworth, Anthony P. |title=Brown dwarf formation by gravitational fragmentation of massive, extended protostellar discs|journal=[[Monthly Notices of the Royal Astronomical Society Letters]]|volume=382|issue=1|pages=L30–L34|date=2007|doi=10.1111/j.1745-3933.2007.00383.x |doi-access=free |bibcode = 2007MNRAS.382L..30S |arxiv = 0708.2827|s2cid=17139868}}</ref>
[[File:PIA18469-AsteroidCollision-NearStarNGC2547-ID8-2013.jpg|thumb|right|300px|Asteroid collision—building planets (artist concept).]]
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 jets]],<ref name=Hartmann1998 /><ref name=Shu1997>{{cite journal|last=Shu|first=Frank H.|author2=Shang, Hsian |author3=Glassgold, Alfred E. |author4= Lee, Typhoon |title=X-rays and Fluctuating X-Winds from Protostars|journal=Science |volume=277|issue=5331|pages=1475–1479|date=1997|doi=10.1126/science.277.5331.1475 |url=http://www.sciencemag.org/cgi/content/full/277/5331/1475|bibcode = 1997Sci...277.1475S}}</ref> whereas the outer part can [[photoevaporation|evaporate]] under the star's powerful [[ultraviolet|UV]] [[radiation]] during the T Tauri stage<ref name=Font2004>{{cite journal|last=Font|first=Andreea S.|author2=McCarthy, Ian G. |author3=Johnstone, Doug |author4= Ballantyne, David R. |title=Photoevaporation of circumstellar disks around young stars|journal=The Astrophysical Journal|volume=607|issue=2|pages=890–903|date=2004|doi=10.1086/383518| bibcode=2004ApJ...607..890F|arxiv = astro-ph/0402241|s2cid=15928892}}</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 comets 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 the 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 /> Interstellar objects could have been captured, and become part of the young Solar system.<ref>{{Cite journal|last1=Grishin|first1=Evgeni|last2=Perets|first2=Hagai B.|last3=Avni|first3=Yael|date=2019-08-11|title=Planet seeding through gas-assisted capture of interstellar objects|journal=Monthly Notices of the Royal Astronomical Society|language=en|volume=487|issue=3|pages=3324–3332|doi=10.1093/mnras/stz1505|doi-access=free |issn=0035-8711|arxiv=1804.09716|s2cid=119066860}}</ref>