Editing Nebular hypothesis (section)

== Formation of planets{{Anchor|Formation of planets}} ==

=== Rocky planets ===
According to the solar nebular disk model, [[rocky planet]]s form in the inner part of the protoplanetary disk, within the [[frost line (astrophysics)|frost line]], where the temperature is high enough to prevent condensation of water ice and other substances into grains.<ref name=Raymond2007>{{cite journal|last=Raymond|first=Sean N.|author2=Quinn, Thomas |author3=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|date=2007|doi=10.1089/ast.2006.06-0126|bibcode=2007AsBio...7...66R|pmid=17407404|issue=1|arxiv = astro-ph/0510285|s2cid=10257401}}</ref> This results in coagulation of purely rocky grains and later in the formation of rocky planetesimals.{{Refn|group=lower-alpha|The [[planetesimal]]s near the outer edge of the terrestrial planet region—2.5 to 4&nbsp;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 />}}<ref name=Raymond2007 /> Such conditions are thought to exist in the inner 3–4&nbsp;AU part of the disk of a Sun-like star.<ref name=Montmerle2006 />

After small planetesimals—about 1&nbsp;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 /> 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&nbsp;years and ends when the largest bodies exceed approximately 1,000&nbsp;km in diameter.<ref name=Kokubo2002>{{cite journal|last=Kokubo|first=Eiichiro|author2=Ida, Shigeru|title=Formation of protoplanet systems and diversity of planetary systems|journal=The Astrophysical Journal|volume=581|issue=1|pages=666–680|date=2002| doi=10.1086/344105|bibcode=2002ApJ...581..666K|doi-access=}}</ref> Slowing of the accretion is caused by gravitational perturbations by large bodies on the remaining planetesimals.<ref name=Kokubo2002 /><ref name=Thommes2003 /> 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 is derived from 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|a(1-e)(M/3M<sub>s</sub>)<sup>1/3</sup>}} is the [[Hill radius]], where a is the [[semimajor axis]], e is the [[orbital eccentricity]], and M<sub>s</sub> is the mass of the central star) 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 isolation mass.<ref name=Thommes2003 /> For the rocky planets it is up to {{Earth mass|0.1}}, 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=Montmerle2006 /><ref name=Kokubo2002 />

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 theory|chaotic]].<ref name=Raymond2006>{{cite journal|last=Raymond|first=Sean N.|author2=Quinn, Thomas |author3=Lunine, Jonathan I. |title=High-resolution simulations of the final assembly of earth-like planets 1: terrestrial accretion and dynamics|journal=Icarus|volume=183|issue=2|pages=265–282|date=2006| doi=10.1016/j.icarus.2006.03.011|bibcode=2006Icar..183..265R|arxiv = astro-ph/0510284|s2cid=119069411}}</ref> During this stage embryos expel remaining planetesimals, and collide with each other. The result of this process, which lasts for 10 to 100&nbsp;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=Montmerle2006 /><ref name=Raymond2006 /><ref name=Bottke2005 /><ref name=Petit2001 /> 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 an equal number of them were thrown out of the Solar System.<ref name=Bottke2005 /> Some of 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 ===
[[File:Fomalhaut Circumstellar Disk.jpg|right|thumb|250px|The dust disk around [[Fomalhaut]]—the brightest star in Piscis Austrinus 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 second possibility is the ''core accretion model'', which is also known as the ''nucleated instability model''.<ref name=Wurchterl2004 /><ref name="ddl2011"/> 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 {{Solar mass|0.1}}).<ref name=ddl2011 /> In this model giant planet formation is divided into two stages: a) accretion of a core of approximately {{Earth mass|10}} and b) accretion of gas from the protoplanetary disk.<ref name=Montmerle2006 /><ref name=Wurchterl2004 /><ref name=dl2018>{{cite book|last=D'Angelo|first=G.|author2=Lissauer, J. J.|chapter=Formation of Giant Planets |bibcode=2018haex.bookE.140D| title=Handbook of Exoplanets |publisher=Springer International Publishing AG, part of Springer Nature| editor=Deeg H., Belmonte J. |pages= 2319–2343|date=2018|arxiv=1806.05649|doi=10.1007/978-3-319-55333-7_140|isbn=978-3-319-55332-0|s2cid=116913980}}</ref> Either method may also lead to the creation of [[brown dwarfs]].<ref name="bodenheimer2013"/><ref name=Janson2011>{{cite journal|last=Janson|first=M.|display-authors=4|author2=Bonavita, M. |author3=Klahr, H. |author4=Lafreniere, D. |author5=Jayawardhana, R. |author6= Zinnecker, H. |title=High-contrast Imaging Search for Planets and Brown Dwarfs around the Most Massive Stars in the Solar Neighborhood|journal=Astrophys. J.|date=2011|volume=736|issue=89|pages=89|arxiv=1105.2577|doi=10.1088/0004-637x/736/2/89 |bibcode = 2011ApJ...736...89J|s2cid=119217803}}</ref> Searches as of 2011 have found that core accretion is likely the dominant formation mechanism.<ref name=Janson2011 />

Giant planet core formation is thought to proceed roughly along the lines of the terrestrial planet formation.<ref name=Kokubo2002 /> It starts with planetesimals that undergo runaway growth, followed by the slower oligarchic stage.<ref name=Thommes2003>{{cite journal|last=Thommes|first=E.W.|author2=Duncan, M.J. |author3=Levison, H.F. |title=Oligarchic growth of giant planets|journal=Icarus|volume=161|issue=2|pages=431–455|date=2003| doi=10.1016/S0019-1035(02)00043-X|bibcode=2003Icar..161..431T|arxiv = astro-ph/0303269|s2cid=16522991}}</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)|frost line]] and consist mainly of ice—the ice to rock ratio is about 4 to 1.<ref name=Inaba2003 /> This enhances the mass of planetesimals fourfold. However, the minimum mass nebula capable of terrestrial planet formation can only form {{Earth mass|1–2}} cores at the distance of Jupiter (5&nbsp;AU) within 10&nbsp;million years.<ref name=Thommes2003 /> The latter number represents the average 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="dangelo2014"/><ref name=Fortier2007>{{cite journal|last=Fortier|first=A.|author2=Benvenuto, A.G.|title=Oligarchic planetesimal accretion and giant planet formation|journal=Astron. Astrophys.|volume=473|issue=1|pages=311–322|date=2007|doi=10.1051/0004-6361:20066729| bibcode=2007A&A...473..311F|arxiv = 0709.1454|s2cid=14812137}}</ref> Some combination of the above-mentioned ideas may explain the formation of the cores of gas giant planets such as [[Jupiter]] and perhaps even [[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&nbsp;AU from the central star.<ref name=Montmerle2006 /> One hypothesis is that they initially accreted in the Jupiter-Saturn region, then were scattered and migrated to their present location.<ref name=Thommes1999>{{cite journal|last=Thommes|first=Edward W.|author2=Duncan, Martin J. |author3=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|date=1999|doi=10.1038/45185|pmid=10604469|issue=6762|bibcode = 1999Natur.402..635T|s2cid=4368864}}</ref> Another possible solution is the growth of the cores of the giant planets via [[pebble accretion]]. In pebble accretion objects between a cm and a meter in diameter falling toward a massive body are slowed enough by gas drag for them to spiral toward it and be accreted. Growth via pebble accretion may be as much as 1000 times faster than by the accretion of planetesimals.<ref name=Lambrechts_Johansen_2012>{{cite journal |title=Rapid growth of gas-giant cores by pebble accretion |journal=Astronomy & Astrophysics |last1=Lambrechts |first1=M. |last2=Johansen |first2=A. |volume=544 |page=A32 |date=August 2012 |doi=10.1051/0004-6361/201219127 |bibcode=2012A&A...544A..32L |arxiv=1205.3030|s2cid=53961588 }}</ref>

Once the cores are of sufficient mass ({{Earth mass|5–10}}), they begin to gather gas from the surrounding disk.<ref name=Montmerle2006 /> Initially it is a slow process, increasing the core masses up to {{Earth mass|30}} in a few million years.<ref name=Inaba2003>{{cite journal|last=Inaba |first=S. |author2=Wetherill, G.W. |author3=Ikoma, M. |title=Formation of gas giant planets: core accretion models with fragmentation and planetary envelope |journal=Icarus |volume=166 |issue=1 |pages=46–62 |date=2003 |doi=10.1016/j.icarus.2003.08.001 |url=http://isotope.colorado.edu/~astr5835/Inaba%20et%20al%202003.pdf |bibcode=2003Icar..166...46I |url-status=dead |archive-url=https://web.archive.org/web/20060912185426/http://isotope.colorado.edu/~astr5835/Inaba%20et%20al%202003.pdf |archive-date=2006-09-12 }}</ref><ref name=Fortier2007 /> After that, the accretion rates increase dramatically and the remaining 90% of the mass is accumulated in approximately 10,000&nbsp;years.<ref name=Fortier2007 /> The accretion of gas stops when the supply from the disk is exhausted.<ref name=dl2018/> This happens gradually, due to the formation of a density gap in the protoplanetary disk and to disk dispersal.<ref name=ddl2011 /><ref name=Papaloizou2007>{{cite encyclopedia |last1=Papaloizou |first1=J. C. B. |last2=Nelson |first2=R. P. |last3=Kley |first3=W. |last4=Masset |first4=F. S. |last5=Artymowicz |first5=P. |display-authors=3 |title=Disk-Planet Interactions During Planet Formation |encyclopedia=Protostars and Planets V |date=2007 |publisher=Arizona Press |editor1=Bo Reipurth |editor2=David Jewitt |editor3=Klaus Keil |bibcode=2007prpl.conf..655P |ref=Papaloizou2007 |pages=655|arxiv = astro-ph/0603196 }}</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 />

During the accretion of gas via streams, a giant planet can be surrounded by a [[circumplanetary disk]]. This circumplanetary disk also carries solids and can form satellites. The [[Galilean moons]] are thought to have formed in such a circumplanetary disk.<ref name="dl2018" />
[[File: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|inclinations]] (see [[Kozai mechanism]]) of planetesimals and embryos in the terrestrial planet region (inside 4&nbsp;AU in the Solar System).<ref name=Bottke2005 /><ref name=Petit2001>{{cite journal|last=Petit|first=Jean-Marc|author2=Morbidelli, Alessandro|title=The Primordial Excitation and Clearing of the Asteroid Belt|journal=Icarus|volume=153|issue=2|pages=338–347|date=2001|doi=10.1006/icar.2001.6702|url=http://www.gps.caltech.edu/classes/ge133/reading/asteroids.pdf|bibcode=2001Icar..153..338P|access-date=2008-03-18|archive-date=2007-02-21|archive-url=https://web.archive.org/web/20070221085835/http://www.gps.caltech.edu/classes/ge133/reading/asteroids.pdf|url-status=dead}}</ref> If giant planets form too early, they can slow or prevent inner planet accretion. 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.|author2=Agnor, Craig |title=The role of giant planets in terrestrial planet formation |journal=The Astronomical Journal|volume=125|issue=5|pages=2692–2713|date=2003|doi=10.1086/374625|url=http://www.boulder.swri.edu/~hal/PDF/tfakess.pdf|bibcode=2003AJ....125.2692L|s2cid=41888579 }}</ref> In addition, the size of the system will shrink, because terrestrial planets will form closer to the central star. The influence of giant planets in the Solar System, particularly that of [[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 the embryos pass close to a giant planet, which may cause them to be ejected from the system.<ref group=lower-alpha>As a variant they may collide with the central star or a giant planet.</ref><ref name=Bottke2005>{{cite journal|last=Bottke|first=William F.|author2=Durda, Daniel D. |author3=Nesvorny, David |display-authors=etal |title=Linking the collisional history of the main asteroid belt to its dynamical excitation and depletion |journal=Icarus|volume=179|issue=1| pages=63–94|date=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|bibcode=2005Icar..179...63B}}</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 are incapable of clearing them all out without the help of embryos. 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 into an [[asteroid belt]], which is a full analog of the asteroid belt in the Solar System, located from 2 to 4&nbsp;AU from the Sun.<ref name=Bottke2005 /><ref name=Petit2001 />

===Exoplanets===

Thousands of exoplanets have been identified in the last twenty years, with, at the very least, billions more, within our observable universe, yet to be discovered.<ref>{{Cite web|url=https://scitechdaily.com/are-we-alone-discovery-of-billions-of-earth-like-planets-may-hold-the-answer/|title=Are We Alone? Discovery of Billions of Earth-Like Planets May Hold the Answer|date=July 8, 2020|website=SciTechDaily}}</ref> The orbits of many of these planets and systems of planets differ significantly from the planets in the Solar System. The exoplanets discovered include hot-Jupiters, warm-Jupiters, super-Earths, and systems of tightly packed inner planets.

The hot-Jupiters and warm-Jupiters are thought to have migrated to their current orbits during or following their formation. A number of possible mechanisms for this migration have been proposed. Type I or Type II migration could smoothly decrease the semimajor axis of the planet's orbit resulting in a warm- or hot-Jupiter. Gravitational scattering by other planets onto eccentric orbits with a perihelion near the star followed by the circularization of its orbit due to tidal interactions with the star can leave a planet on a close orbit. If a massive companion planet or star on an inclined orbit was present an exchange of inclination for eccentricity via the [[Kozai mechanism]] raising eccentricities and lowering perihelion followed by circularization can also result in a close orbit. Many of the Jupiter-sized planets have eccentric orbits which may indicate that gravitational encounters occurred between the planets, although migration while in resonance can also excite eccentricities.<ref name="Baruteau_etal_2014">{{cite book|last1=Baruteau|first1=C.|last2=Crida|first2=A.|last3=Paardekooper|first3=S.-J.|last4=Masset|first4=F.|last5=Guilet|first5=J.|last6=Bitsch|first6=B.|last7=Nelson|first7=R.|last8=Kley|first8=W.|last9=Papaloizou|first9=J.|chapter=Planet-Disk Interactions and Early Evolution of Planetary Systems|title=Protostars and Planets VI|date=2014|pages=667–689|arxiv=1312.4293|bibcode=2014prpl.conf..667B|doi=10.2458/azu_uapress_9780816531240-ch029|isbn=9780816531240|s2cid=67790867}}</ref> The in situ growth of hot Jupiters from closely orbiting super Earths has also been proposed. The cores in this hypothesis could have formed locally or at a greater distance and migrated close to the star.<ref name="Batygin_etal_2016">{{cite journal|last1=Batygin|first1=Konstantin|last2=Bodenheimer|first2=Peter H.|last3=Laughlin|first3=Gregory P.|title=In Situ Formation and Dynamical Evolution of Hot Jupiter Systems|journal=The Astrophysical Journal|date=2016|volume=829|issue=2|page=114|doi=10.3847/0004-637X/829/2/114|arxiv=1511.09157|bibcode = 2016ApJ...829..114B |s2cid=25105765 |doi-access=free }}</ref>

Super-Earths and other closely orbiting planets are thought to have either formed in situ or ex situ, that is, to have migrated inward from their initial locations.<ref name=dangelo_bodenheimer_2016>{{Cite journal|last=D'Angelo|first=G.|author2= Bodenheimer, P. |s2cid=119203398|title=In Situ and Ex Situ Formation Models of Kepler 11 Planets|journal=The Astrophysical Journal|year=2016|volume=828|issue=1|pages=id. 33 (32 pp.)|doi=10.3847/0004-637X/828/1/33|arxiv = 1606.08088 |bibcode = 2016ApJ...828...33D |doi-access=free }}</ref>  The in situ formation of closely orbiting super-Earths would require a massive disk, the migration of planetary embryos followed by collisions and mergers, or the radial drift of small solids from farther out in the disk. The migration of the super-Earths, or the embryos that collided to form them, is likely to have been Type I due to their smaller mass. The resonant orbits of some of the exoplanet systems indicates that some migration occurred in these systems, while the spacing of the orbits in many of the other systems not in resonance indicates that an instability likely occurred in those systems after the dissipation of the gas disk. The absence of Super-Earths and closely orbiting planets in the Solar System may be due to the previous formation of Jupiter blocking their inward migration.<ref name="Morbidelli_Raymond_2016">{{cite journal |last1= Morbidelli |first1= Alessandro |last2= Raymond |first2= Sean |title= Challenges in planet formation |journal= Journal of Geophysical Research: Planets |date=2016 |volume= 121 |issue=10 |pages= 1962–1980 |doi= 10.1002/2016JE005088 |arxiv= 1610.07202 |bibcode = 2016JGRE..121.1962M |s2cid= 119122001 }}</ref>

The amount of gas a super-Earth that formed in situ acquires may depend on when the planetary embryos merged due to giant impacts relative to the dissipation of the gas disk. If the mergers happen after the gas disk dissipates terrestrial planets can form, if in a transition disk a super-Earth with a gas envelope containing a few percent of its mass may form. If the mergers happen too early runaway gas accretion may occur leading to the formation of a gas giant. The mergers begin when the dynamical friction due to the gas disk becomes insufficient to prevent collisions, a process that will begin earlier in a higher metallicity disk.<ref name="Lee_Chiang_2016">{{cite journal|last1=Lee|first1=Eve J.|author1-link=Eve Lee|last2=Chiang|first2=Eugene|title=Breeding Super-Earths and Birthing Super-puffs in Transitional Disks|journal=The Astrophysical Journal|date=2016|volume=817|issue=2|page=90|doi=10.3847/0004-637X/817/2/90|arxiv=1510.08855|bibcode = 2016ApJ...817...90L |s2cid=118456061 |doi-access=free }}</ref> Alternatively gas accretion may be limited due to the envelopes not being in hydrostatic equilibrium, instead gas may flow through the envelope slowing its growth and delaying the onset of runaway gas accretion until the mass of the core reaches 15 Earth masses.<ref name="Lambrechts_Lega">{{cite journal|last1=Lambrechts|first1=Michiel|last2=Lega|first2=Elana|title=Reduced gas accretion on super-Earths and ice giants|journal=Astronomy and Astrophysics|volume=606|pages=A146|date=2017|arxiv=1708.00767|bibcode = 2017A&A...606A.146L |doi=10.1051/0004-6361/201731014|s2cid=118979289}}</ref>