惑星物理学2019：惑星形成に関する概論

2019೥4݄11೔ɹߨٛʮ࿭੕෺ཧֶʯ ࿭੕ܗ੒ʹؔ͢Δ֓࿦ େֶӃཧֶݚڀՊ Ӊ஦෺ཧֶڭࣨɹࠤʑ໦وڭ

ຊ೔ͷ಺༰ ✤ ݱࡏͷଠཅܥͷ࢟  ɹଟछଟ༷ͳ࿭੕ɾӴ੕ɾখఱମ ✤ ଠཅܥܗ੒࿦ͷ؆୯ͳϨϏϡʔ  ɹݹయతඪ४Ϟσϧʮژ౎Ϟσϧʯͷ֓ཁ ✤ ܥ֎࿭੕ͷൃݟͱ൚࿭੕ܗ੒ཧ࿦ͷߏங  ɹҟܗͷ࿭੕ͨͪͷ࡞Γํ

ݱࡏͷଠཅܥͷ࢟

ଠཅܥͷߏ੒ϝϯόʔ ஍ٿܕ࿭੕ ɹɹਫ੕ ɹɹۚ੕ ɹɹ஍ٿ ɹɹՐ੕ ڊେΨε࿭੕ ɹɹɹ໦੕ ɹɹɹ౔੕ ڊେණ࿭੕
ɹɹఱԦ੕ ɹɹւԦ੕ (c) ikachi.org

֤ఱମͷيಓ ᯁੴͷ฼ఱମ ୹पظክ੕ͷ૥ (c) wakatsuki

͞Βʹԕ͘·Ͱ޿͕Δଠཅܥ ʮΦʔϧτͷӢʯ ௕पظክ੕ͷ૥ ʹ ఱจ୯Ґʢ"6ʣ ଠཅ͔Β஍ٿ·Ͱͷڑ཭ ʢ໿ԯສLNʣ (c) http://kahaku.go.jp

ਫ੕ ۚ੕ ஍ٿ Ր੕ يಓ௕൒ܘ [AU] 0.39 0.72 1 1.52
ެసपظ [೥] 0.241 0.615 1 1.881 ࣭ྔ [஍ٿ = 1] 0.055 0.82 1 0.11 ൒ܘ [km] 2440 6052 6378 3396 ີ౓ [kg/m3] 5430 5240 5520 3930 Ӵ੕ͷ਺ 0 0 1 2 ஍ٿܕ࿭੕ͷੑ࣭

໦੕ ౔੕ ఱԦ੕ ւԦ੕ يಓ௕൒ܘ [AU] 5.2 9.6 19.2 30.1
ެసपظ [೥] 11.86 29.46 84.02 164.7 ࣭ྔ [஍ٿ = 1] 317.8 95.2 14.5 17.2 ൒ܘ [km] 71490 60270 25560 24760 ີ౓ [kg/m3] 1330 690 1270 1640 Ӵ੕ͷ਺ 79 62 27 14 ڊେΨε࿭੕ɾණ࿭੕ͷੑ࣭

஍ٿܕ࿭੕ͷ಺෦ߏ଄ ਫ੕ ஍֪ Ϛϯτϧ ίΞ Ր੕ ஍ٿ ݄ ۚ੕ ݻମͷ಺֩
ӷମͷ֎֩ ίΞʁ (c) Wikipedia

ڊେΨε࿭੕ɾණ࿭੕ͷ಺෦ߏ଄ ໦੕ ౔੕ ఱԦ੕ ւԦ੕ ஍ٿ ਫૉ෼ࢠ ۚଐਫૉ ਫૉɾϔϦ΢ϜɾϝλϯΨε ϚϯτϧʢਫɾΞϯϞχΞɾϝλϯණʣ
ίΞʢؠੴɾණʣ (c) Wikipedia

ଠཅܥͷ୅දతͳӴ੕ͨͪ (c) Wikipedia

ଠཅܥܗ੒࿦ͷ؆୯ͳϨϏϡʔ

ݪ࢝ଠཅܥԁ൫ͷ̎ͭͷϞσϧ
©Newton Press ژ౎ϞσϧʢྛϞσϧʣ ΩϟϝϩϯϞσϧ ஍ٿܕ࿭੕ɾڊେΨε࿭੕ɾڊେණ࿭੕ͷ࡞Γ෼͚ ɹˠଠཅܥܗ੒ʹؔͯ͠͸ɺژ౎Ϟσϧͷํʹ܉഑ (c) YouTube

ژ౎Ϟσϧͷجຊ֓೦ ԁ൫Ծઆ ɾ࿭੕ܥ͸ݪ࢝࿭੕ܥԁ൫͔Βܗ੒͞ΕΔ ɾԁ൫͸Ψεͱμετ͔Βߏ੒͞ΕΔ ඍ࿭੕Ծઆ ɾμετͷूੵʹΑͬͯඍ࿭੕͕ܗ੒͞ΕΔ ɾඍ࿭੕ͷूੵʹΑͬͯݻମ࿭੕͕ܗ੒͞ΕΔ ɾݻମ࿭੕ʹΨε͕߱Γੵ΋Δ͜ͱʹΑͬͯ ɹΨε࿭੕͕ܗ੒͞ΕΔ ɹɹɹɹɹɹɹɹɹɹɹɹ<ྛ஧࢛࿠ଞ
>

෼ࢠӢ͔Βओܥྻ੕΁ͷਐԽ (c) Yusuke Tsukamoto ʢ̍ʣ੕ؒ෼ࢠӢͷऩॖͱίΞͷܗ੒ ʢ̎ʣݪ࢝੕ͷܗ੒ͱ੒௕ ʢ̏ʣओܥྻ੕΁ͷਐԽ ੕ܗ੒ͷ ̏ஈ֊

ݪ࢝࿭੕ܥԁ൫ ෼ࢠӢίΞͷऩॖ ɹॏྗͱԕ৺ྗͷͭΓ͍͋ ݪ࢝࿭੕ܥԁ൫͕ܗ੒ (c) NASA ෼ࢠӢίΞ ݪ࢝࿭੕ܥԁ൫

ݪ࢝ଠཅܥԁ൫ͷ૊੒ Ұൠʹԁ൫࣭ྔͷ͸ΨεʢਫૉɾϔϦ΢Ϝʣ ࢒Γͷ͕μετʢݻମ੒෼ʣ ɾݱࡏͷଠཅܥͷ࿭੕ͷݻମ੒෼ʢ໿.ଠཅʣ ɹɹˠ͢ΓͭͿͯ͠ԁ൫ঢ়ʹͳΒ͢ ɾݻମ੒෼ͷ໿ഒͷ࣭ྔͷΨε੒෼ΛՃ͑Δ ࠷খ࣭ྔԁ൫Ϟσϧʢژ౎Ϟσϧʣ ݪ࢝ଠཅܥԁ൫ͷॳظ࣭ྔ͸໿.ଠཅ ॏྗͱԕ৺ྗͷ௼Γ߹͍͔Β൒ܘ͸໿"6 4OPXMJOFҎԕͰ͸ਫ͕ڽ݁͠ݻମ໘ີ౓্͕ঢ

ݪ࢝ଠཅܥԁ൫ͷ࣭ྔ෼෍ ଠཅ͔Βͷڑ཭ʢ"6ʣ ໘ີ౓ʢHDNʣ Ψε੒෼ ݻମ੒෼ Ψε੒෼ɿਫૉɾϔϦ΢Ϝ ݻମ੒෼ɿμετʢؠੴɾۚଐమɾණʣ "6ҎԕͰ͸ ਫৠؾ͕ڽॖ ɹɹɹˣ
ණμετͷ෼͚ͩ ໘ີ౓্͕ঢ͢Δ TOPXMJOF

ݪ࢝࿭੕ܥԁ൫ͷ؍ଌ ࣮ࡍʹ༷ʑͳܗͷԁ൫͕؍ଌ͞Ε͍ͯΔ ɹˠݪ࢝࿭੕ܥԁ൫͸͔֬ʹଘࡏ͢Δʂ Subaru ALMA

"-."๬ԕڸʹΑΔ؍ଌ (੕ܗ੒θϛɿ2019/01/24) DSHARP Paper I (Andrews et al. 2018, ApJL
869, L41) Munetake MOMOSE (Ibaraki U.) Figure 3. Gallery of 240 GHz (1.25 mm) continuum emission images for the disks in the DSHARP sample. Beam sizes and 10 au scalebars are shown in the lower left and right corners of each panel, respectively. All images are shown with an asinh stretch to reduce the dynamic range (accentuate fainter details without over- saturating the bright emission peaks). For more quantitative details regarding the image dimensions and intensity scales, see Huang et al. (2018a) and Kurtovic et al. The Astrophysical Journal Letters, 869:L41 (15pp), 2018 December 20 Andrews et al. <"OESFXTFUBM > %JTL4VCTUSVDUVSFT BU)JHI"OHVMBS 3FTPMVUJPO1SPKFDU %4)"31 ݪ࢝࿭੕ܥԁ൫͸ଟ༷ͳ ߏ଄Λ࣋ͭ͜ͱ͕໌Β͔ ʹͳ͖ͬͯͨ

ଠཅܥܗ੒ඪ४ཧ࿦ʢژ౎Ϟσϧʣ
©Newton Press 巨大氷惑星形成

ඍ࿭੕ܗ੒γφϦΦͱ༷ʑͳࠔ೉ μετ(≲µm) ࿭੕(≳103km) ݪ࢝࿭੕ܥԁ൫ ඍ࿭੕(≳km) ࿭੕ܗ੒ͱඍ࿭੕ܗ੒ !5 ice +rock ice+rock
rock Itokawa (~0.5km) rock ॏྗूੵ ෼ࢠؒྗूੵ + μετ૚ͷࣗݾ ॏྗෆ҆ఆ? εϊʔϥΠϯ (~1-3AU?) ௚઀߹ମ੒௕ μετͷࣗݾॏྗෆ҆ఆ μετ ඍ࿭੕ ЖN NN N LN ੩ి൓ൃোน ௓ͶฦΓোน த৺੕མԼোน িಥഁյোน ཚྲྀোน ☓☓ ☓ ☓☓

!10-4 # " $ [g/cm3] roll
imp 2 3 ~ 40 . 2 ~ E E N dN d f f E imp = - p dV V 0 V Suyama et al. submitted to ApJ μετͷ߹ମ੒௕ μετͷ߹ମ੒௕ ɹˠඍ࿭੕ܗ੒ ඍ࿭੕ͷԁ൫͕ܗ੒ (c) Toru Sayama ණඍ࿭੕͸ܗ੒Մೳ ؠੴඍ࿭੕͸ഁյ͕ ୎ӽ͠ܗ੒ෆՄೳʁ

ଠཅܥܗ੒ඪ४ཧ࿦ʢژ౎Ϟσϧʣ
©Newton Press 巨大氷惑星形成

ඍ࿭੕ͷ߹ମ੒௕ ਺LNαΠζͷ ඍ࿭੕͕ܗ੒ ޓ͍ʹিಥɾ߹ମ Λ܁Γฦ͠੒௕ ˣ ๫૸త੒௕ ɹେཻ͖͍ࢠ΄Ͳ੒௕͕଎͍ டংత੒௕ ɹશͯͷཻࢠ͕ಉ͡଎౓Ͱ੒௕
(c) Kouji KANBA

ଟମ໰୊ઐ༻ܭࢉػ(3"1& ଟମʢඍ࿭੕ʣͷॏྗܭࢉ ɹˠܭࢉྔ͕๲େʹͳΔ ཻࢠؒ૬ޓ࡞༻ͷ෦෼͚ͩΛ ઐ༻ܭࢉػͰܭࢉ͍ͨ͠ ɹˠ(3"1&஀ੜʂ (3"1&ͱ຀໺३Ұ࿠ڭत (c) Junichiro Makino

KOKUBO AND IDA FIG. 4. Time evolution of the maximum
mass (solid curve) and the mean mass (dashed curve) of the system. thanthisrangearenotstatisticallyvalidsinceeachmassbinoften has only a few bodies. First, the distribution tends to relax to a ๫૸త੒௕ͷ༷ࢠ ฏۉ஋ ࠷େͷఱମ ඍ࿭੕ͷ๫૸త੒௕ ɹˠݪ࢝࿭੕͕஀ੜ͢Δ 20 KOKUBO AND IDA FIG. 3. Snapshots of a planetesimal system on the a–e plane. The circles represent planetesimals and their radii are proportional to the radii of planetesimals. The system initially consists of 3000 equal-mass (1023 g) planetesimals. FIG. 4. Time evolution of the maximum mass (solid curve) and the mean mass (dashed curve) of the system. thanthisrangearenotstatisticallyvalidsinceeachmassbinoften has only a few bodies. First, the distribution tends to relax to a decreasing function of mass through dynamical friction among (energy equipartition of) bodies (t = 50,000, 100,000 years). Second, the distributions tend to ﬂatten (t = 200,000 years). This is because as a runaway body grows, the system is mainly heated by the runaway body (Ida and Makino 1993). In this case, the eccentricity and inclination of planetesimals are scaled by the يಓ௕൒ܘ<"6> يಓ཭৺཰ ࣭ྔ<H> ࣌ؒ<೥> <,PLVCP*EB >

Չ઎త੒௕ͷ༷ࢠ FORMATION OF PROTOPLANETS FROM PLANETESIMALS 23 FIG. 7. Snapshots
of a planetesimal system on the a–e plane. The circles represent planetesimals and their radii are proportional to the radii of planetesimals. The system initially consists of 4000 planetesimals whose to- tal mass is 1.3 × 1027 g. The initial mass distribution is given by the power- FIG. 8. The number of bodies in linear mass bins is plotted for t = 100,000, 200,000, 300,000, 400,000, and 500,000 years. In Fig. 10, we plot the maximum mass and the mean mass of يಓ཭৺཰ ֤৔ॴͰඍ࿭੕͕๫૸త੒௕ ɹˠ౳αΠζͷݪ࢝࿭੕͕ฒͿ Չ઎త੒௕ͱΑͿ ʹ ֤يಓͰͷݪ࢝࿭੕ ࣭ྔ [kg] ܗ੒࣌ؒ [yr] ஍ٿيಓ 1×1024 7×105 ໦੕يಓ 3×1025 4×107 ఱԦ੕يಓ 8×1025 2×109 يಓ௕൒ܘ<"6> <,PLVCP*EB >

ݪ࢝࿭੕͔Β࿭੕΁ )-/ 2 .3 ( )
-/(% -/ ) 0 # " 4+ ݪ࢝࿭੕ͷ࣭ྔ<஍ٿ࣭ྔ> يಓ௕൒ܘ<"6> ஍ٿܕ࿭੕ ɹݪ࢝࿭੕ಉ࢜ͷ߹ମ ڊେΨε࿭੕ ɹݪ࢝࿭੕ͷΨεั֫ ڊେණ࿭੕ ɹݪ࢝࿭੕ͦͷ·· TOPXMJOF (c) Eiichiro Kokubo

δϟΠΞϯτΠϯύΫτ ݪ࢝࿭੕ಉ࢜ͷڊେఱମিಥΛ܁Γฦ͠ ݱࡏͷ࿭੕΁ (c) NASA

δϟΠΞϯτΠϯύΫτ يಓ௕൒ܘ<"6> يಓ཭৺཰ planets is hnM i ’ 2:0 Æ
0:6, which means that the typical result- ing system consists of two Earth-sized planets and a smaller planet. In this model, we obtain hna i ’ 1:8 Æ 0:7. In other words, one or two planets tend to form outside the initial distribution of protoplanets. In most runs, these planets are smaller scattered planets. Thus we obtain a high efﬁciency of h fa i ¼ 0:79 Æ 0:15. The accretion timescale is hTacc i ¼ 1:05 Æ 0:58 ð Þ ; 108 yr. These results are consistent with Agnor et al. (1999), whose initial con- Fig. 2.—Snapshots of the system on the a-e (left) and a-i (right) planes at t ¼ 0, 1 are proportional to the physical sizes of the planets. KOKUBO, KOMIN 1134 ௕͍࣌ؒΛ͔͚ͯݪ࢝࿭੕ಉ࢜ͷيಓ͕ཚΕΔ ɹˠޓ͍ʹিಥɾ߹ମͯ͠ΑΓେ͖ͳఱମʹ੒௕ <,PLVCP*EB > (c) Hidenori Genda

ڊେఱମিಥʹΑΔԁ൫ܗ੒ ݪ࢝஍ٿʹՐ੕αΠζͷ ݪ࢝࿭੕͕িಥ ඈͼࢄͬͨഁย͕஍ٿͷ पғʹԁ൫Λܗ੒ (c) Robin Canup [Canap &
Asphaug, 2001]

ԁ൫தͰͷ݄ܗ੒ [Ida et al., 1997] प஍ٿԁ൫಺Ͱ݄ܗ੒ ݱࡏͷ݄ͷҐஔΑΓ΋ ͔ͳΓ಺ଆʢ໿ʣ

Moon Formation Simulations ਺ϲ݄ʙ਺೥Ͱɺͻͱͭͷ݄͕Ͱ͖Δ

ڊେΨε࿭੕ͷܗ੒ ݪ࢝࿭੕ʹԁ൫Ψε͕๫૸తʹྲྀೖˠΨε࿭੕΁ (c) NVIDIA

Ψεั֫ʹΑΔڊେΨε࿭੕ܗ੒ ݪ࢝࿭੕͸ॏྗʹΑΓपғͷԁ൫ΨεΛั֫ ɾ஍ٿ࣭ྔҎԼˠେؾѹͰࢧ͑ΒΕͯ҆ఆʹଘࡏ ɾ஍ٿ࣭ྔҎ্ˠେؾ่͕յɾ๫૸తʹΨεั֫ يಓ෇ۙʹ࢒͍ͬͯΔΨεΛશͯՃ଎౓తʹั֫ ɹˠٸܹʹ࣭ྔΛ૿͠໦੕ɾ౔੕΁ͱ੒௕͢Δ (c) Nagoya U

ΨεͷྲྀΕͷ਺஋γϛϡϨʔγϣϯ (c) Takayuki Tanigawa

ڊେΨε࿭੕ͷܗ੒ͷ༷ࢠ Fig. 4.—Structure around the Hill sphere for model M04
on the midplane (left) and in three dimensions, shown in bird’s-eye view (right). The gas stream MACHIDA ET AL. 1226 V 1.—Time sequence for model M04. The density (color scale) and velocity distributions (arrows) on the cross section in the ˜ z ¼ 0 plane are plotted. The bottom ¼ 3) are 4 times the spatial magniﬁcation of the top panels (l ¼ 1). Three levels of grids are shown in each top (l ¼ 1, 2, and 3) and bottom (l ¼ 3, 4, and 5) panel. l of the outermost grid is denoted in the top left corner of each panel. The elapsed time ˜ tp and the central density ˜ c on the midplane are denoted above each of the ls. The velocity scale in units of the sound speed is denoted below each panel. पғͷԁ൫Ψε͕ݪ࢝࿭੕ͷॏྗݍ಺ʹั֫͞ΕΔ (c) Takayuki Tanigawa

प࿭੕ԁ൫಺ͰӴ੕͕ܗ੒ ΠΦ Τ΢ϩύ Ψχϝσ ΧϦετ λΠλϯ (c) Wikipedia (c) Nagoya
U

ܥ֎࿭੕ͷൃݟͱ ɹ൚࿭੕ܗ੒ཧ࿦ͷߏங

.BZPS2VFMP[ εΠεͷ؍ଌνʔϜ ਓྨॳͷܥ֎࿭੕ݕग़ʂ ϖΨαε࠲൪੕ͷपΓʹ)PU+VQJUFS͕ଘࡏʂ ೥݄ ਓྨॳͷܥ֎࿭੕ൃݟ (c) NASA (c)
Michel Mayor

ଠཅܥ֎࿭੕͕ଓʑͱݟ͔ͭΔ ೥݄೔ݱࡏ ݸͷܥ֎࿭੕͕֬ೝ

όϥΤςΟʹ෋Ήܥ֎࿭੕ܥ ඪ४తͳ࿭੕ܗ੒γφϦΦʹΑͬͯઆ໌Մೳ͔ʁ (c) NASA (c) Wikipedia (c) NASA (c) picshype.com

࿭੕ܥͷଟ༷ੑΛੜΈग़͢ཁૉ ɾݪ࢝࿭੕ܥԁ൫ͷ࣭ྔͷҧ͍ ɹˠΨε࿭੕ͷݸ਺΍Ґஔͷҧ͍ΛੜΉʁ ɾيಓෆ҆ఆʹΑΔ࿭੕ܥͷมԽ ɹˠ௕͍࣌ؒΛ͔͚ͯҟͳΔ࿭੕ܥ΁Ҡߦʁ ɾܗ੒தͷ࿭੕ͷத৺੕ํ޲΁ͷམԼ ʢλΠϓ*࿭੕མԼˍλΠϓ**࿭੕མԼʣ ɹˠ࠷ऴతͳ࿭੕ͷҐஔͷҧ͍ΛੜΉʁ ɾ࿭੕ͷҠಈʹ൐͏࿭੕ܥͷมԽ ɹˠΑΓଟ༷ͳ࿭੕ܥ͕ܗ੒͞ΕΔʁ

ଟ༷ͳݪ࢝࿭੕ܥԁ൫ Դڇ࠲ ΁ͼ͔͍ͭ࠲ ԁ൫ͷ࣭ྔ<ଠཅ࣭ྔ> ൃ
ݟ ਺ ଠཅܥ෮ݩԁ൫ Ӊ஦ʹ͸༷ʑͳ࣭ྔΛ࣋ͭݪ࢝࿭੕ܥԁ൫͕ଘࡏ ɹˠԁ൫ͷ࣭ྔͷҧ͍͕ଟ༷ͳ࿭੕ܥΛੜΈग़͢ʂʁ

ଟ༷ͳԁ൫͔Βੜ·ΕΔଟ༷ͳ࿭੕ ԁ൫ͷ࣭ྔͷҧ͍ˠΨε࿭੕ͷ਺ͱҐஔͷҧ͍ the escape velocity of protoplanets. This high random
velocity makes the accretion process slow and inefficient and thus Tgrow longer. This accretion inefficiency is a severe problem On the ot in circular o HD 192263 with Æ1e 1 for in situ f case. It is d slingshot m circular orb the magnet may be wea disks may b Terrestria Jovian plan planetary a key process systems. We confir holds in Æsolid ¼ Æ1 ð ¼ 1=2; 3= tions. We d systems dep disk profile growth tim and (17), re a Mdisk T <T grow disk T <T cont disk Fig. 13.—Schematic illustration of the diversity of planetary systems against the initial disk mass for < 2. The left large circles stand for central stars. The double circles (cores with envelopes) are Jovian planets, and the others are terrestrial and Uranian planets. [ See the electronic edition of the Journal for a color version of this figure.] ݪ࢝࿭੕ܥԁ൫ͷ࣭ྔ يಓ௕൒ܘ த৺੕͔Βͷڑ཭ <,PLVCP*EB >

Weidenschilling &,Marzari (1996),,Lin,& a GM a GM a
GM a GM a GM * * 3 * 2 * 1 * $&,60 45 '"#, (t >~ 1My) /% , "# 2-+0'! 3. *) 1."#3. 00 a1 0 . 00 a1(. final e يಓෆ҆ఆʹΑΔ࿭੕ܥͷมԽ ࿭੕ؒͷॏྗͷӨڹ͕ ੵΈॏͳͬͯ࠷ऴతʹ ޓ͍ͷيಓ͕ෆ҆ఆԽ ҟͳΔ࿭੕ܥ΁ ˣ </BHBTBXBFUBM > ᶃ ᶃ ᶄ ᶄ ᶅ ᶅ

λΠϓ*࿭੕མԼ ݄࣭ྔʙ஍ٿ࣭ྔͷఱମʹޮ͘ϝΧχζϜ ఱମ͕ԁ൫ʹཱͯͨີ౓೾ʹΑΓ֯ӡಈྔΛࣦ͏ (c) Frédéric Masset

λΠϓ**࿭੕མԼ ஍ٿ࣭ྔҎ্ͷఱମʹޮ͘ϝΧχζϜ ఱମ͕ԁ൫ʹߔΛ࡞Γԁ൫ͱͱ΋ʹத৺੕ʹམԼ (c) Frédéric Masset

࿭੕ͷҠಈʹ൐͏࿭੕ܥͷมԽ earing continues through scattering. After 00 million years the
inner disk is composed the collection of planetesimals at 0.06 AU, a M] planet at 0.12 AU, the hot Jupiter at 0.21 U, and a 3 M] planet at 0.91 AU. Previous sults have shown that these planets are likely be stable for billion-year time scales (15). Many bodies remain in the outer disk, and ac- orbital time scales and high inclinations. Two of the four simulations from Fig. 2 contain a 90.3 M] planet on a low-eccentricity orbit in the habitable zone, where the temper- ature is adequate for water to exist as liquid on a planet_s surface (23). We adopt 0.3 M] as a lower limit for habitability, including long-term climate stabilization via plate tectonics (24). three categories: (i) hot Earth analogs interior to the giant planet; (ii) Bnormal[ terrestrial planets between the giant planet and 2.5 AU; and (iii) outer planets beyond 2.5 AU, whose accretion has not completed by the end of the simulation. Properties of simulated planets are segregated (Table 1): hot Earths have very low eccentric- ities and inclinations and high masses because g. 1. Snapshots in time of the evolution of one simulation. Each panel ots the orbital eccentricity versus semimajor axis for each surviving body. he size of each body is proportional to its physical size (except for the ant planet, shown in black). The vertical ‘‘error bars’’ represent the sine of each body’s inclination on the y-axis scale. The color of each dot corresponds to its water content (as per the color bar), and the dark inner dot represents the relative size of its iron core. For scale, the Earth’s water content is roughly 10j3 (28). λΠϓ* **࿭੕མԼʹ ΑΓ࿭੕ܥͷيಓ͕େ͖ ͔͖͘ཚ͞ΕΔ they accrete on the migration time scale (105 years), so there is a large amount of damping during their formation. These planets are remi- niscent of the recently discovered, close-in 7.5 M] planet around GJ 876 (25), whose formation is also attributed to migrating resonances (26). ଟ༷ͳ࿭੕ܥܗ੒ <3BZNPOEFUBM >

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惑星物理学2019：惑星形成に関する概論

惑星物理学2019：惑星形成に関する概論

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