Editing Nebular hypothesis (section)

=== 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 />