formationofsolarsystemanaloguesii: discsvian-body simulations

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MNRAS 000 1ndash24 (2018) Preprint 5 July 2018 Compiled using MNRAS LATEX style file v30

Formation of Solar system analogues II

post-gas phase growth and water accretion in extended

discs via N-body simulations

M P Ronco1234⋆ G C de Elıa12dagger1Instituto de Astrofısica de La Plata CCT La Plata - CONICET UNLP Paseo del Bosque SN (1900) La Plata Argentina2Facultad de Ciencias Astronomicas y Geofısicas Universidad Nacional de La Plata Paseo del Bosque SN (1900) La Plata Argentina3Instituto de Astrofısica Pontificia Universidad Catolica de Chile Santiago Chile4Nucleo Milenio de Formacion Planetaria (NPF) Chile

Accepted XXX Received YYY in original form ZZZ

ABSTRACT

This work is the second part of a project that attempts to analyze the formationof Solar system analogues (SSAs) from the gaseous to the post-gas phase in a self-consistently way In the first paper (PI) we presented our model of planet formationduring the gaseous phase which provided us with embryo distributions planetesimalsurface density eccentricity and inclination profiles of SSAs considering different plan-etesimal sizes and type I migration rates at the time the gas dissipates In this secondwork we focus on the late accretion stage of SSAs using the results obtained in PI asinitial conditions to carry out N-body simulations One of our interests is to analyzethe formation of rocky planets and their final water contents within the habitable zoneOur results show that the formation of potentially habitable planets (PHPs) seems tobe a common process in this kind of scenarios However the efficiency in forming PHPsis directly related to the size of the planetesimals The smaller the planetesimals thegreater the efficiency in forming PHPs We also analyze the sensitivity of our resultsto scenarios with type I migration rates and gap-opening giants finding that bothphenomena act in a similar way These effects seem to favor the formation of PHPsfor small planetesimal scenarios and to be detrimental for scenarios formed from bigplanetesimals Finally another interesting result is that the formation of water-richPHPs seems to be more common than the formation of dry PHPs

Key words planets and satellites dynamical evolution and stability - planets andsatellites formation - methods numerical

1 INTRODUCTION

During the past years improvements in the exoplanetdetection techniques and the increase in the quantityand quality of the missions have led to the discoveryof exoplanets of all kinds To date the number of con-firmed exoplanets rises to 3798 and continues growing(httpexoplaneteu Schneider et al 2011) In particu-lar the number of rocky planets has grown the most dur-ing the last years discovered mainly by the Kepler mis-sion (httpswwwnasagovmission_pageskepler) Thenumber of gas giant exoplanets in orbits with semimajor-axis greater than sim 1 au has also grown and many of theserocky and gas giant planets are part of multiple planetary

⋆ E-mail mroncoastropuccldagger E-mail gdeeliafcaglpunlpeduar

systems 633 until now Moreover multiple planetary sys-tems with more than one gas giant planet have also beendetected From this many of these multiple planetary sys-tems present a central star similar to our sun However todate we have not yet been able to detect planetary systemssimilar to our own Solar system This is planetary systemsformed simultaneously by at least a giant planet with a massof the order of Saturn or Jupiter in the outer zones of thedisc and by rocky planets in the inner zone particularly byrocky planets within the habitable zone

Several works studied the formation of terrestrial plan-ets and water accretion process in different dynamical sce-narios or within the framework of the formation of ourSolar system For example Raymond et al (2004) andRaymond et al (2006) analyzed the effects produced by aJovian planet located in the outer zones of the disc on theformation and dynamics of terrestrial planets around sun-

copy 2018 The Authors

2 M P Ronco et al

like stars Later Mandell et al (2007) explored the forma-tion of terrestrial-type planets in planetary systems with onemigrating gas giant around solar-type stars They also con-sidered planetary systems with one inner migrating gas giantand one outer non-migrating giant planet Fogg amp Nelson(2009) examined the effect of giant planet migration onthe formation of inner terrestrial planet systems and foundthat this phenomena does not prevent terrestrial planet for-mation Salvador Zain et al (2017) focused on the forma-tion and water delivery within the habitable zone in dif-ferent dynamical environments and also around solar-typestars These authors considered planetary systems that har-bor Jupiter or Saturn analogues around the location of thesnowline and found that the formation of water-rich planetswithin the habitable zones of those systems seems to be acommon process More recently Raymond amp Izidoro (2017)showed that the inner Solar systemrsquos water could be a re-sult of the formation of the giant planets These authors findthat when the gas disc starts to dissipate a population ofscattered eccentric planetesimals cross the terrestrial planetregion delivering water to the growing Earth

During our first work (Ronco et al 2017 hereafter PI)we improved our model of planet formation which calculatesthe evolution of a planetary system during the gaseous phaseto find via a population synthesis analysis suitable scenar-ios and physical parameters of the disc to form Solar systemanalogues (SSAs from now on) We were specially interestedin the planet distributions and in the surface density ec-centricity and inclination profiles for the planetesimal popu-lation at the end of the gas-phase obtained considering dif-ferent formation scenarios with different planetesimal sizesand different type I migration rates

The aim of this second work is to analyze the formationof SSAs during the post-gas phase through the developmentof N-body simulations focusing on the formation of rockyplanets and their water contents in the inner regions of thedisc particularly on the habitable zone To do so and inorder to analyze the formation of SSAs in a self-consistentlyway we make use of the results obtained in PI as the initialconditions for our N-body simulations and form final plane-tary systems from different planetesimal sizes Many worksin the literature that analyze the post-gas phase formation ofplanetary systems in particular the formation of rocky plan-ets in the inner regions of the disc considered arbitrary ini-tial condition (Chambers 2001 Raymond et al 2004 2006OrsquoBrien et al 2006 Mandell et al 2007 Raymond et al2009 Walsh et al 2011 Ronco amp de Elıa 2014) Howeveras it was already shown by Ronco et al (2015) more re-alistic initial conditions obtained from a model of planetformation lead to different accretion histories of the planetsthat remain within the habitable zone than considering ad

hoc initial conditions We also analyze the differences in theresults of planetary systems that were affected by the type Imigration regime (Tanaka et al 2002) during the gas-phaseand by planetary systems with gas giant planets that wereable to open a gap in the disc and were affected by type IImigration

This work is then organized as follow In Section 2 wesummarize our model of planet formation explained in de-tailed in PI in Section 3 we describe the initial conditionsfor the development of our N-body simulations obtainedfrom the results of PI in Section 4 we describe the most

important properties of the N-body code and then show ourmain results in Section 5 Then we compare our results withthe observed potentially habitable planets and with the cur-rent gas giant population in Section 6 Finally discussionsand conclusions are presented in Section 7 and Section 8respectively

2 SUMMARY OF THE GAS-PHASE PLANET

FORMATION MODEL

In PI we improved a model of planet formation based on pre-vious works (Guilera et al 2010 2011 2014) which we nowcalled PLANETALP This code describes the formation andevolution in time of a planetary system during the gaseousphase taking into account many relevant phenomena forthe planetary system formation The model presents a pro-toplanetary disc which is characterized by a gaseous compo-nent evolving due to an α-viscosity driven accretion (Pringle1981) and photoevaporation (DrsquoAngelo amp Marzari 2012)processes and a solid component represented by a plan-etesimal and an embryo population The planetesimal pop-ulation is being subject to accretion ejection (Laakso et al2006) and scattering (Ida amp Lin 2004 Alibert et al 2005)by the embryos and radial drift due to gas drag in-cluding the Epstein Stokes and quadratic regimes Theembryo population grows by accretion of planetesimals(Inaba et al 2001) gas (Ida amp Lin 2004 Miguel et al 2011Mordasini et al 2009 Tanigawa amp Ikoma 2007) and due tothe fusion between them taking into account their atmo-spheres (Inamdar amp Schlichting 2015) while they migratevia type I (Tanaka et al 2002) or type II (Crida et al 2006)migration mechanisms We also consider that the evolu-tion of the eccentricities and inclinations of the planetesi-mals are due to two principal processes the embryo gravi-tational excitation (Ohtsuki et al 2002) and the dampingdue to the nebular gas drag (Rafikov 2004 Chambers 2008)The particularities and the detailed description of the im-plementation of each phenomena can be found in PI and inGuilera et al (2014)

In PI we also described the method initial parametersand initial formation scenarios considered to synthesize apopulation of different planetary systems around a solar-mass star with the particular aim of forming SSAs Follow-ing Andrews et al (2010) the gas and solid surface densityprofiles are modeled as

Σg(R) = Σ0g

(R

Rc

)minusγeminus(

R

Rc

)2minusγ

(1)

Σp(R) = Σ0pηice

(R

Rc

)minusγeminus(

R

Rc

)2minusγ

(2)

where Σ0p and z0Σ

0g are normalization constants with z0 =

00153 the primordial abundance of heavy elements in thesun (Lodders et al 2009) and ηice is a parameter that rep-resents the increase in the amount of solid material due tothe condensation of water beyond the snowline given by 1if R ge Rice and given by 1β if R lt Rice where Rice is theposition of the snowline at 27 au (Hayashi 1981)

The variable parameters which were taken randomlyfrom range values obtained from the same observationalworks (Andrews et al 2010) are the mass of the disc Md

MNRAS 000 1ndash24 (2018)

Solar system analogues II 3

between 001M⊙ and 015M⊙ the characteristic radius Rc

between 20 au and 50 au and the density profile gradientγ between 05 and 15 The factor β in the solid surfacedensity profile is also taken randomly between 11 and 3The viscosity parameter α for the disc and the rate of EUVionizing photons f41 have a uniform distribution in log be-tween 10

minus4 and 10minus2 and between 10

minus1 and 104 respectively

(DrsquoAngelo amp Marzari 2012)

The formation scenarios considered include differentplanetesimal sizes (rp = 100 m 1 km 10 km and 100 km)but we consider a single size per simulation The type Imigration rates for isothermal-discs were slowed down in1 ( fmigI = 001) and 10 ( fmigI = 01) We also consid-ered the extreme cases in where type I migration was nottaken into account ( fmigI = 0) or was not slowed down at all( fmigI = 1)

We also constrained the disc lifetime τ between 1 Myrand 12 Myr following Pfalzner et al (2014) and consideringonly those protoplanetary discs whose combinations betweenα and f41 satisfied this condition

Finally the random combinations of all the parametersgave rise to a great diversity of planetary systems which wereanalyzed in PI From this diversity only 43 of the totalof the developed simulations formed planetary systems thatat the end of the gaseous phase were SSAs Recalling fromPI we define SSAs as those planetary systems that presentat the end of the gaseous phase only rocky planets in theinner region of the disc and at least one gas giant planet likeJupiter or Saturn beyond 15 au The great majority of theseSSAs were found in formation scenarios with small planetes-imals and lownull type I migration rates These SSAs pre-sented different properties and architectures for exampledifferent number of gas giant planets and icy giant planetsand different final planetesimal profiles Moreover from arange of disc masses between 001M⊙ and 015M⊙ SSAs wereformed only within discs more massive than Md ge 004Moplus

since low-mass discs have not enough solid material to formmassive cores before the disc dissipation timescales How-ever these SSAs were not sensitive to the different values ofRc γ β α and τ since we found them throughout the wholerange of values considered for each of these parameters Themean values for the disc parameters that formed SSAs werelt Md gt= 01M⊙ lt Rc gt= 328 au lt γ gt= 095 lt β gt= 207lt α = 34 times 10

minus3 and lt τ gt= 645 Myr

It is important to remark that the planetary systemsobtained at the end of the gaseous phase and which aregoing to be consider as initial conditions for the post-gasgrowth were formed within the framework of a planetaryformation code that although presents important physicalphenomena for the formation of a planetary system dur-ing the gaseous phase it also presents certain limitationsSome of these limitations are for example the use of a ra-dial 1D model to compute the evolution of a classical α-disc the consideration of isothermal discs and the corre-sponding type I migration rates for such discs and the lackof the calculation of the gravitational interactions betweenembryos Then the initial conditions for the N-body simula-tions carried out in this work are dependent on the physicsof our planetary formation model and more detailed treat-ments of the physics could affect the final embryo and plan-etesimal configurations of planetary systems at the end ofthe gaseous phase For example recent works which con-

sider non-isothermal discs have shown that type I migra-tion rates could be very different with respect to previousisothermal type I migration rates (Paardekooper et al 20102011) Moreover for non-isothermal discs migration con-vergent zones are generated acting as traps for migratingplanets (Hellary amp Nelson 2012 Bitsch et al 2015) On theother hand the presence of a dead zone along the disc couldalso affect both the planet migration and the planetesimalradial drifts Guilera amp Sandor (2017) showed that even inisothermal discs the presence of a dead zone in the disc leadsto pressure maxima at the edges of such zones acting as trapsfor migrating planets and also as convergent zones for drift-ing planetesimals Finally alternative models for α-discs aswind-discs have been proposed to model the evolution ofthe gaseous component of protoplanetary discs showing dif-ferent final disc structures (Suzuki et al 2010 2016) Thesedifferences are important and could affect models of plan-etesimal formation and planet migration

3 INITIAL CONDITIONS

The goal of this second work is to study and analyze theevolution of the post-gas phase of planetary systems thatresulted in SSAs at the end of the gaseous phase From thisFig 8 of PI showed some of the embryo and planetesimaldistributions of the different scenarios of SSAs at the end ofthe gaseous phase

Some of the final configurations obtained after the gas iscompletely dissipated are used in this work as initial condi-tions to develop N-body simulations Since one of our maingoals is to analyze the formation of rocky planets within thehabitable zone (HZ) of SSAs we focuse on the long-termevolution of planetary systems between 05 au and 30 auWe are not interested in the formation of close-in planetsvery near the central star Thus we do not consider plane-tary embryos neither planetesimals inside 05 au

Table 1 and Fig 1 show the main properties of the cho-sen outcomes of the results of PI that are going to be usedto analyze the late-accretion stage Due to the high com-putational cost of N-body simulations of this kind it is notviable to develop simulations for each SSA found Thus wehave to chose the most representative one for each forma-tion scenario Therefore the chosen outcomes shown in Fig1 that represent the embryo distributions and the planetesi-mal surface density profiles of SSAs at the end of the gaseousphase were carefully selected from each formation scenariothis is taking into account the different planetesimal sizesand different type I migration rates considering that it wasthe most representative of the whole group

As it was explained in PI a water radial distribution forthe embryos and planetesimals was included in our model ofplanet formation The colorscale in Fig 1 shows the watercontents in embryos at the end of the gaseous phase As thefraction per unit mass of water in embryos and planetesi-mals at the beginning of the gaseous phase was given by

fH2O(R) =

0 if R le Rice

β minus 1

β if R gt Rice

(3)


4 M P Ronco et al

Table 1 Chosen configurations of SSAs at the end of the gas-stage to be used as initial conditions for the development of N-bodysimulations The total mass in embryos (without considering the mass of the Jupiter or Saturn analogues) and planetesimals is calculatedbetween 05 au and 30 au except for the lowast cases which represent scenarios that had been cutted of until 15 au in order to use 1000planetesimals The mass in embryos inside and outside represent the mass in embryos inside and outside the giants location The lastfour columns of this table show characteristics concerning water contents β is a variable factor that represents the jump in the solidsurface density due to the water condensation beyond the snowline The maximum water content by mass is represented in

Formation scenario Total mass in Mass in embryos Initial water contents

Scenario Planetesimal fmigI Embryos Planetesimals Inside Outside β Max water Dry Water-richsize contents embryos embryos

S1 100 m 0 6665Moplus 25 times 10minus4Moplus 6665Moplus 0Moplus 212 528 72 14

S2 100 m 001 8633Moplus 828Moplus 7517Moplus 1116oplus 211 526 49 15S3 100 m 01 3513Moplus 1378Moplus 2654Moplus 859Moplus 175 428 62 10S4 1 km 0 11436Moplus 17450Moplus 6147Moplus 5289Moplus 214 533 28 25S5 1 km 001 9361Moplus 15655Moplus 4847Moplus 4515Moplus 227 559 25 28S6 10 km 0 8902Moplus 19530Moplus 4259Moplus 4643Moplus 137 270 23 23S7 10 km 001 12091Moplus 22643Moplus 2341Moplus 975Moplus 266 624 23 20S8 10 km 01 13466Moplus 27184Moplus 1887Moplus 11579Moplus 275 636 23 34Slowast

9100 km 0 11644Moplus 16440Moplus 6622Moplus 5022Moplus 149 328 22 24

Slowast10


10minus5

10minus4

10minus3

10minus2

10minus1

100

101100m minus fmigI=0

Fraction of water by mass

S1

Mas

s [M

J]

0 01 02 03 04 05 06



S2

0 01 02 03 04 05 06

100m minus fmigI=01


S3

0 01 02 03 04 05 06

1km minus fmigI=0


S4

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1

1001km minus fmigI=001


S5

Σ p [g

cm

minus2]

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1

100

101

1 10

10km minus fmigI=0


S6

Mas

s [M

J]

Semimajorminusaxis [au]

0 01 02 03 04 05 06

1 10



S7


0 01 02 03 04 05 06

1 10

10km minus fmigI=01


S8


0 01 02 03 04 05 06

1 10

100km minus fmigI=0


S9


0 01 02 03 04 05 06

1 1010minus5

10minus4

10minus3

10minus2

10minus1


S10

Σ p [g

cm

minus2]


0 01 02 03 04 05 06

Figure 1 Chosen SSAs configurations at the end of the gaseous phase to develop N-body simulations in a mass vs semimajor-axisvs planetesimal surface density plane Each colored point represents a planetary embryo the black and grey big dots represent Jupiterand Saturn analogues respectively while the colored points with black borders represent icy giants like Neptune-like planets The solidblack line in all the planetary systems represents the planetesimal density profile when the gas is already gone The colorscale representsthe final amounts of water by mass in each planet and the size of each planet is represented in logarithmic scale

the water contents in embryos at the end of the gaseousphase which represent the initial water contents for thepost-gas phase evolution vary on the one hand accordingto the β value considered for each scenario and on the otherhand due to the growth of each embryo during the gas phaseAs we mentioned before β takes random values from 11 (which represents a 9 of water by mass) to 3 (which repre-sents a 66 of water by mass) The planetesimal populationremains always the same regarding water contents planetes-imals inside 27 au remain dry and planetesimals beyond27 au contain the amount of water by mass given by Eq3 As we already mentioned the last four columns of ta-ble 1 show the main properties related to water contents ineach formation scenario considered to develop N-body sim-ulations

It is important to note that at the end of the gaseous

phase we are able to distinguish which is the origin of thewater content of each planet This is we can know the per-centage of water by mass at the end of the gaseous phasedue to the accretion of planetesimals and the percentage ofwater by mass due to the fusion between embryos

An interesting difference between some of the forma-tion scenarios considered is that due to the position of theformed giant planets during the gaseous phase not all theSSAs present a water-rich embryo population inside the gi-ant locations This topic will play an important role in theanalysis of the water contents in planets within the HZ atthe end of the N-body simulations and it will be attendedin the following sections



4 PROPERTIES OF THE N-BODY

SIMULATIONS

Our N-body simulations begin when the gas is already dissi-pated in the protoplanetary disc The numerical code used tocarry out our N-body simulations is the MERCURY code de-veloped by Chambers (1999) In particular we adopted thehybrid integrator which uses a second-order mixed variablesymplectic algorithm for the treatment of the interaction be-tween objects with separations greater than 3 Hill radii anda Burlisch-Stoer method for resolving closer encounters

The MERCURY code calculates the orbits of gas gi-ants planetary embryos and planetesimals When these or-bits cross each other close encounters and collisions betweenthe bodies may occur and are registered by the MERCURYcode It is important to mention that in our simulations col-lisions between all the bodies are treated as inelastic merg-ers which conserve mass and water content and because ofthat our model does not consider water loss during impactsAs a consequence the final mass and water contents of theresulting planets should be considered as upper limits

Given the high numerical cost of the N-body simula-tions and in order to perform them in a viable CPU timewe only consider gravitational between planets (embryos andgiant planets) and between planets and planetesimals butplanetesimals are not self-interacting This is planetesimalsperturb and interact with embryos and giant planets butthey ignore one another completely and they cannot collidewith each other (Raymond et al 2006) It is important toremark that these planetesimals are in fact superplanetes-

imalsWhile the size of the planetesimals is only relevantduring the gas phase when planetesimals interact with thegaseous component of the disc due to the drift of the neb-ular gas (see PI) the size ceases to be relevant in the post-gas phase during the N-body simulations Thus due to thenumerical limitations of N-body simulations regarding thelimited number of bodies that has to be used it is not possi-ble to consider in a consistently way realistic planetesimalsof the proposed sizes in PI this is planetesimals of 100 m1 km 10 km and 100 km Instead we have to consider highermass superplanetesimals that represent an idealized swarmof a larger number of real size planetesimals Thus from nowon when we speak about planetesimals we are referring tosuperplanetesimals

As we already mentioned we are interested in the post-gas evolution of planetary systems similar to our own andthis is why we consider planetary embryos and planetes-imals in extended discs between 05 au and 30 au Thenumber of embryos varies between 40 and 90 according toeach scenario and present physical densities of 3 gr cmminus3 Weconsider 1000 planetesimals in all the developed simulationswith physical densities of 15 gr cmminus3 Since the remainingmass in planetesimals at the end of the gaseous phase insimulations with planetesimals of 100 km and type I migra-tion rates reduced to a 01 are too high we cut the discuntil 15 au in order to be able to consider 1000 planetesi-mals The choice of considering only 1000 planetesimals allwith the same mass for each scenario has to do with the factthat considering more planetesimals would be unviable tak-ing into account the time needed to run these simulationsand since this is a number that will allow us to study in

a moderately detailed way the evolution of this populationthroughout the whole disc

Since terrestrial planets in our Solar system may haveformed within 100 Myr - 200 Myr (Touboul et al 2007Dauphas amp Pourmand 2011 Jacobson et al 2014) we in-tegrate each simulation for 200 Myr To compute the innerorbit with enough precision we use a maximum time step ofsix days which is shorter than 120 th of the orbital periodof the innermost body in all the simulations The reliabil-ity of all our simulations is evaluated taking into accountthe criterion of energy |dEE | lt Et where Et is an energythreshold and it is considered as 10

minus3 However in order toconserve the energy to at least one part in 10

3 we carry outsimulations assuming shorter time steps of four and up totwo days depending on the requirement of each simulationThose simulations that do not meet this energy thresholdvalue are directly discarded

To avoid any numerical error for small-perihelion or-bits we use a non-realistic size for the sun radius of 01 au(Raymond et al 2009) The ejection distance considered inall our simulations was 1000 au This value is realistic enoughfor the study of extended planetary systems and is also avalue that allows us to reduce the total number of bod-ies of our simulations as time passes Higuchi et al (2007)found that planetesimals with semimajor-axis gt 1000 au in-crease their perihelion distances outside the planetary region(100 au) due to the fact that they would already be affectedby galactic tides In addition Veras amp Evans (2013) showedthat wide-orbit planets with semimajor-axis greater than1000 au are substantially affected by galactic tides whilecloser-in planets at sim 100 au are only affected if the hoststar is near the galactic center within the inner hundredparsecs

The semimayor-axes of the planetesimals are generatedusing an acceptance-rejection method taking into accountthe surface density profile of planetesimals shown with blackcurves in each formation scenario of Fig 1 As it was ex-plained before the orbital eccentricities and inclinations forthe planetesimal population are results of the gas-phase evo-lution (see PI) Each planetesimal surface density profile hasan eccentricity and inclination profile that are also used asinitial conditions for this population Figure 2 shows the ec-centricity profile (black curve) of the planetesimal surfacedensity profile (red curve) of the results of S8 as an exam-ple For embryos and giant planets the initial eccentricitiesand inclinations are taken randomly considering values lowerthan 002 and 05 respectively Recalling from PI embryosand giant planets evolve in circular and coplanar orbits dur-ing the gaseous phase Thus to be consistent with our modelof planet formation we initiate the N-body simulations withlow values for the eccentricities and inclinations It is wothmentioning that the model developed in PI does not yetconsider the gravitational self-stirring of embryos and giantplanets Therefore this limitation in our model could changethe configurations of our SSAs at the end of the gaseousstage

The rest of the orbital parameters such as the argumentof pericenter ω longitude of ascending node Ω and the meananomaly M are also taken randomly between 0

and 360 for

the gas giants the embryos and the planetesimals

Due to the stochastic nature of the accretion process


6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]


e profilePlanetesimals

Planetesimal density

Zoom

Figure 2 Eccentricity profile (black curve) for the planetesimalsurface density profile (red curve) of the end of the gas-phase evo-

lution of scenario S8 of Fig 1 The oscillations in the eccentricityprofiles of planetesimals are caused during the gaseous phase (seePI) by the embryo gravitational excitation (Ohtsuki et al 2002)and by the damping due to the nebular gas drag (Rafikov 2004Chambers 2008) The zoom shows some planetesimals in blackpoints All of them the 1000 planetesimals represent the mass inplanetesimals available between 05 au and 30 au located on topof the eccentricity profile

we perform 10 simulations with different random numberseeds per each formation scenario (see Table 1)

5 RESULTS

In this section we describe the main results of our N-bodysimulations for the formation of planetary systems fromSSAs configurations It is worth remembering that the con-sidered architecture of a SSA consists on a planetary systemwith rocky planets in the inner region of the disc and atleast one giant planet a Jupiter or Saturn analogue beyond15 au The main goal of our work is to study the formationprocess of the whole system focusing on the formation ofrocky planets and the water delivery within the HZ As wehave developed a great number of N-body simulations takinginto account different planetesimal sizes and different migra-tion rates and in order to show our results in an organizedway we first describe the results of scenarios formed fromplanetesimals of 100 km 10 km 1 km and 100 m that didnot undergo type I migration and that present gas giantsthat had not managed to open a gap during the gas phaseas the reference scenarios

The sensitivity of the results to different type I migra-tion rates and to giant planets that open a gap in the gasdisc are going to be exposed in the subsequent sections

51 The habitable zone

One of the goals of this work is to analyze the formationof rocky planets within the HZ and their final water con-tents The HZ was first defined by Kasting et al (1993)as the circumstellar region inside where a planet can re-tain liquid water on its surface For a solar-mass star thisregion remains between 08 au and 15 au (Kasting et al

1993 Selsis et al 2007) Later and particularly for the Solarsystem Kopparapu et al (2013ab) proposed that an opti-mistic HZ is set between 075 au and 177 au These limitswere estimated based on the belief that Venus has not hadliquid water on its surface for the last 1 billion years andthat Mars was warm enough to maintain liquid water onits surface The authors also defined a conservative HZ de-termined by loss of water and by the maximum greenhouseeffect provided by a CO2 atmosphere between 099 au and167 au However locating a planet within these regions doesnot guarantee the potential habitability taking into accountthat if the orbits of planets found inside these regions aretoo eccentric they could be most of the time outside themand this could avoid long times of permanent liquid wateron their surfaces In order to prevent this situation in pre-vious works that also analyze the formation of rocky planetswithin the HZ we considered that a planet is in the HZand can hold liquid water if it has a perihelion q ge 08 auand a aphelion Q le 15 au (Ronco amp de Elıa 2014) or hasconservative or optimistic chances to be habitable if it hasa perihelion q ge 075 au and a aphelion Q le 177 au and ifit has a perihelion q ge 099 au and a aphelion Q le 167 aurespectively (Ronco et al 2015)

Williams amp Pollard (2002) showed that the planetaryhabitability would not be compromised if the planet doesnot stay within the HZ during all the integration If an oceanis present to act as a heat capacitor or the atmosphere pro-tects the planet from seasonal changes it is primarily thetime-averaged flux that affects the habitability over eccen-tric orbits In this sense Williams amp Pollard (2002) deter-mined that the mean flux approximation could be valid forall eccentricities However Bolmont et al (2016) exploredthe limits of the mean flux approximation varying the lumi-nosity of the central star and the eccentricity of the planetand found that planets orbiting a 1L⊙ star with eccentric-ities higher than 06 would not be able to maintain liquidwater throughout the orbit

Kopparapu et al (2014) improved the analysis of theHZ boundaries founding a dependence between the limits ofthe HZ with the masses of the planets Assuming H2O (in-ner HZ) and CO2 (outer HZ) atmospheres their results indi-cated that larger planets have wider HZs than smaller onesParticularly Kopparapu et al (2014) provided with para-metric equations to calculate the corresponding HZ bound-aries of planets of 01Moplus 1Moplus and 5Moplus They determinedthat for a 1L⊙ the inner conservative boundary of a planetof 01Moplus is 1005 au for a planet of 1Moplus is 095 au and fora planet of 5Moplus is 0917 au while the outer limits remainthe same for the three of them and is 167 au

In this work on the one hand we consider that thoseplanets with masses lower than 01Moplus present the same in-ner HZ boundary of a planet of 01Moplus planets with massesbetween 01Moplus and 1Moplus present the same inner HZ bound-aries of a planet of 1Moplus and planets with masses between1Moplus and 5Moplus present the same inner HZ boundaries of aplanet of 5Moplus Planets with masses greater than 5Moplus alsopresent the inner HZ boundaries of a planet of 5Moplus

On the other hand we also calculate the time-averagedincident stellar flux for the inner and outer HZ limits of aneccentric planet in terms of the solar flux following

Flowasteff=

Feff

(1 minus e2)12 (4)



where Feff is the effective flux from a circular orbit Thenfollowing Williams amp Pollard (2002) we can draw constantflux curves in the semimajor-axis vs eccentricity plane forthe HZ boundaries

Therefore given the previous mentioned works and ourconsiderations the adopted criteria for the definition of theHZ and the class of potentially habitable planets PHPs fromnow on is as follows

(i) If a planet presents at the end of the N-body simula-tions a semimajor-axis between its HZ boundaries and hasa perihelion q greater and a aphelion Q lower than the limitvalues for their HZ boundaries we consider this planet is aPHP class A

(ii) If a planet presents a (a e) pair that remains insidethe constant flux curves associated to the inner and outerHZ edges and has an eccentricity lower or equal than 06we consider this planet is a PHP class B

(iii) Finally if a planet remains within a 10 of the inneror the outer constant flux curve and has an eccentricitylower or equal than 06 we consider this planet is a PHPclass C

Although planets of the three classes will be consideredas PHPs the most important ones taking into account theinner and outer HZ boundaries and the mean flux definitionare those of class A and B Planets of class C for beingin a region that extends in a 10 on both sides of eachboundary are considered as marginally habitable but PHPsat last Figure 3 shows as an example the different HZregions for a planet of 5Moplus These regions are reduced if weconsider less massive planets since larger planets have widerHZs than smaller ones

Regarding final water contents we consider that to bea PHP of real interest a planet has to harbor at least somewater content better if this content is of the order of the es-timations of water on Earth The mass of water on the Earthsurface is around 28 times 10

minus4Moplus which is defined as 1 Earthocean However as the actual mass of water in the mantleis still uncertain the total mass of water on the Earth isnot yet very well known However Lecuyer amp Gillet (1998)estimated that the mantle water is between 08 times 10

minus4Moplus

and 8 times 10minus4Moplus and Marty (2012) estimated a value of

2 times 10minus3Moplus Thus taking into account the water reservoirs

that should be in the Earthrsquos mantle and the mass of wa-ter on the surface the present-day water content on Earthshould be around 01 to 02 by mass which representsbetween 36 to 71 Earth oceans Moreover Abe et al (2000)suggested that the primitive mantle of the Earth could havecontained between 10 and 50 Earth oceans However otherauthors assumed that the current content is not very dif-ferent from the one that was accreted during the formationprocess (see Raymond et al 2004 and references therein)Althought to be a PHP of real interest the planet has topresent at least some water content we do not consider aminimum for this value since we could not ensure the po-tential non-habitability for lower values

52 Description of the reference scenarios

Here we describe the most important characteristics of thefour reference scenarios planetary systems formed from

Figure 3 Habitable zone class A B and C for a planet of 5Moplus

planetesimals of 100 km 10 km 1 km and 100 m and with-out suffering migration regimes during the gaseous phase

521 Scenarios with planetesimals of 100 km

Main perturber A gas giant planet similar to a Saturn-like1 planet acts as the main perturber of the systems Thisgiant is initially located at 297 au and then migrates inwarddue to an embryo-driven migration until it ends up locatedbetween 157 au and 197 au at the end of the evolutionwith a mass ranging between 1714Moplus and 190Moplus (this is18MS and 2MS) taking into account the whole group ofsimulations

General dynamical evolution The dynamical evo-lution of all the simulations performed with planetesimals of100 km is very similar Figure 4 six snapshots in time of thesemimajor-axis eccentricity plane of SIM4 as an example ofthe whole group The initial conditions correspond to Slowast

9in

Fig 1 From the very beginning the inner embryo popula-tion which does not coexist with planetesimals gets higheccentricities due to their own mutual gravitational pertur-bations and particularly with the giant planet The plane-tary embryos evolved through collisions with planetesimalsand other embryos Particularly the collisional events be-tween the inner embryos happened within the first millionyear of evolution where at that time only one planet re-mains in the inner zone of the disc

Mass remotion Between a 30 and a 40 of the ini-tial embryos and between a 6 and a 14 of the initialplanetesimals survive at the end of the simulations Con-sidering mean values for all the simulations a 276 of theinitial embryos collide with the central star while a 247is accreted and a 39 of is ejected from the systems The

1 From PI we remember that our consideration of a Saturn-likeplanet is a giant planet with an envelope that is greater than themass of the core but its total mass is less than 200Moplus


8 M P Ronco et al

Figure 4 Time evolution of SIM4 of the scenario with planetesimals of 100 km and without considering type I migration The shadedareas represent the HZ regions already described in Section 1 From now on all the HZ regions plotted are the widest possible whichcorrespond to a planet of 5Moplus Planets are plotted as colored circles planetesimals are plotted with black dots and the giant planetwhich is a Saturn-like planet in this case is represented in grey The size of the Saturn-like planet is fixed and is not in the same scalefor the rest of the planets The same happens with the planetesimal population each planetesimal adopts a size of half the size of thesmallest embryo independently of its mass Then the size for the rest of the planets is represented by a lineal relation between the massand the size where the size for the smallest embryo is the double of the size of the planetesimals and the size of the biggest planet issmaller that the size for the Saturn-like planet The color scale represents the fraction of water of the planets with respect to their totalmasses This scenario always presents a Mega-Earth in the inner region per simulation Half of them are completely dry as it can be seenin this case This scenario does not form PHPs within the HZ

rest of them remains in the inner zone2 as a Mega-Earth3

and in the outer zones of the disc coexisting simultaneously

2 From now on when we talk about planets or planetesimals fromthe inner zone we are referring to those planets or planetesimalsthat are located inside the giants location at the beginning ofthe N-body simulations Similarly planets or planetesimals fromthe outer zone are those planets or planetesimals that are locatedbeyond the giants location at the beginning of the N-body simu-lations3 We classify the formed rocky planets as Mega-Earths if theypresent masses between 8Moplus and 30Moplus Super-Earths if theypresent masses between 2Moplus and 8Moplus and Earth-like planets iftheir masses are lower than 2Moplus and greater than half an Earth

with a planetesimal population The planetesimals suffer asimilar process They are located behind the giant and areperturbed by the giant and the surrounding planets Con-sidering all the simulations the great majority (sim 843)are ejected from the systems and only a few collide withthe central star (31) or are accreted (25) by the sur-rounding planets Following these results the accretion ofplanetesimals is an inefficient process in this scenario Therest of the planetesimals (sim 10) presents high eccentricitiesat the end of the evolution suggesting that if we extendedthe simulations by a few million years more this populationwould probably end up being completely ejected from thesystems



Inner region Inside the final location of the giantwe always find a Mega-Earth with final mass in a rangeof 145Moplus minus 26Moplus located between 037 au and 079 auThe seeds of these Mega-Earths where initially located al-ways inside the snowline between 056 au and 237 au withan initial mass between 170Moplus and 676Moplus During theevolution of the systems these seeds also migrated inwardsfinally settling between 037 au and 079 au Half of theseMega-Earths are completely dry at the end of the evolutionThe other half present amounts of water by mass due tothe acretion of very few water-rich planetesimals Howevernone of them lie within the HZ This is why this scenariois not interesting from an astrobiological point of view theefficiency in forming PHPs is completely null

Outer region Beyond the giant a mixture of dry andwater-rich planets and a planetesimal population with highmean eccentricities can be found until sim 400 au This mixof material was found in all the developed simulations andthe average number of dry planets found in the outermostregions is sim 4 The gravitational interactions between thegas giant and the embryo population cause that some of theinternal embryos are injected into the outer zones of the dischowever no water-rich embryos from beyond the snowline areinjected in the inner regions

Finally after 200 Myr of evolution the planetary sys-tems resulting from Slowast

9in figure 1 are composed by a dry

Mega-Earth at around 055 au in half of the cases a Saturn-like planet near or in the outer edge of the HZ and of severalplanets beyond the giant planet


Main perturbers In these systems two Jupiter-like plan-ets with masses of 079MJ and 064MJ and initially locatedat 373 au and 588 au respectively act as the main per-turbers After 200 Myr of evolution the innermost giant islocated at around 285 au and the outermost giant is locatedat around 578 au

General dynamical evolution From the beginingthe inner embryos are quickly excited by the two Jupiter-likegravitational perturbations and by the perturbations witheach other The planetesimal population which remains be-hind the gas giants is also perturbed However very fewplanetesimals are able to cross the giant barrier Thereforeinner planetary embryos grow almost only by the accretionof other embryos within the first Myr of evolution As an ex-ample of this scenario figure 5 shows six snapshots in timeof the semimajor-axis eccentricity plane of SIM1

Mass remotion During the evolution of the systemsa sim 8 of the initial embryo population collides with thecentral star a 198 is accreted and a sim 377 is ejectedfrom the system considering the whole group of simulationsThe planetesimals are mostly ejected from the systems (sim771) due to the interactions with the giants and the outerplanets Only a few collide with the central star (17) or areaccreted (13) by the surrounding planets In this scenarioas in the previous one the planetesimal accretion is also avery inefficient process along the entire length of the disc

Inner region Inside the orbit of the giant planets wefind between 1 and 2 planets Considering the 10 developedsimulations the trend shows that is more likely to find 2planets instead of 1 These planets are completely dry Super-

Table 2 Principal properties of the planets that remain withinthe HZ in scenarios with planetesimals of 10 km and without typeI migration The row in bold corresponds to the results of Fig 5

SIM ai [au] af [au] Mass [Moplus] W [] Habitableclass

SIM1 185 130 1080 0 A

SIM2 mdash mdash mdash mdash mdashSIM3 mdash mdash mdash mdash mdashSIM4 122 101 1441 0 BSIM5 mdash mdash mdash mdash mdashSIM6 mdash mdash mdash mdash mdashSIM7 mdash mdash mdash mdash mdashSIM8 199 117 686 0 ASIM9 215 145 1157 0 ASIM10 057 094 356 0 C

Earths or dry Mega-Earths with masses in a range of 356Moplus

to 1596Moplus However the trend is to find planets more mas-sive than 10Moplus In this scenario 5 of 10 simulations formedone PHP Although most of them are class A all of themare completely dry because their accretion seeds come fromthe region inside the snowline Their masses range between356Moplus and 1157Moplus The properties of these PHPs can beseen in table 2

Outer region Beyond the gas giants planets andplanetesimals survived until 400 au - 500 au In these sys-tems the mix of material is less efficient than for planetes-imals of 100 km We only find one dry planet in the outerzone beyond the giants in half of the simulationsbull It seems that the gravitational interactions between thegas giants and the embryos in this case are stronger thanwith planetesimals of 100 km and instead of injecting dryembryos into the outer zone they directly eject them out ofthe system

Finally at the end of the evolution the planetary sys-tems resulting from S6 in figure 1 are formed mostly by twodry Super or Mega-Earths two Jupiter-like planets betweensim 285 au and 578 au and by several planets beyond thesegiants In 5 of 10 simulations we find a dry PHP


Main perturbers A Saturn-like planet and a Jupiter-likeplanet act as the main perturbers of these systems TheSaturn-like planet is the gas giant which is closer to the cen-tral star and is initially located at 4 au This planet finisheslocated between 168 au and 297 au due to an embryo-drivenmigration after 200 Myr of evolution and presents a meanmass of 46Moplus (048MS) This inward migration is significantbut it is not enough to disperse all the solid material of theinner zone of the disc as it does for example the Saturn-like planet of the scenario of 100km The Jupiter-like planetoriginally located at 567 au finishes always located aroundsim 46 au with a mass of 16MJbull Icy giants Unlike the previous scenarios this onepresents two Neptune-like planets4 between the rocky em-

4 From PI the adopted definition for an icy giant planet is that ofa planet which presents a mass of the envelope greater than 1 ofthe total mass but lower than the mass of the core Particularly if


10 M P Ronco et al

Figure 5 Time evolution of a simulation of the scenario with planetesimals of 10 km and without considering type I migration Theshaded areas represent the HZ regions already described in Section 1 Planets are plotted as colored circles planetesimals are plotted withblack dots and the giant planets which are Jupiter-like planets in this case are represented in black The size of the Jupiter-like planetsis fixed and is not in the same scale for the rest of the planets The same happens with the planetesimal population each planetesimaladopts a size of half the size of the smallest embryo independently of its mass Then the size for the rest of the planets is represented

by a lineal relation between the mass and the size where the size for the smallest embryo is the double of the size of the planetesimalsand the size of the biggest planet is smaller that the size for a Saturn-like planet like in the previous scenario

The color scale represents the fraction of water of the planets with respect to their total masses This scenario form dry PHPs withinthe HZ

bryos and the gas giant planets at the beginning of thesimulations (see S4 in Fig 1) with masses of sim 15Moplus and sim

17Moplus General dynamical evolution From the beginning

the embryos are excited due to the mutual perturbations be-tween them and with both the icy and the gas giants Thefirst most common result found in 6 of 10 simulations isthat one of the two Neptunes is ejected from the systemswithin the first sim 01 Myr The second most common resultwhich happens in 5 of the 6 simulations mentioned before

the mass of the envelope is greater than 1 of the total mass andlower than 30 of the total mass we consider it is a Neptune-likeplanet

is that the other Neptune is scattered towards the outerzone of the disc presenting encounters with an excited plan-etesimal population Here it can be seen the importance ofthe dynamical friction produced by the planetesimals thatdampens the eccentricity and inclination of the Neptune-like planet The mass available in planetesimals to producethis effect is during this process high enough to cause theplanet to settle in the mid-plane of the disc in less than sim

1 Myr We can see in figure 6 six snapshots in time of thesemimajor-axis eccentricity plane of SIM10 as an exampleof this scenario Although is not the most common resultwe find that in 3 of the 10 simulations one of the Neptune-like planets survives inside the giants locations and finisheswithin the HZ while the other one is ejected

Mass remotion During the simulations a 158 of



Table 3 Principal properties of the planets that remain withinthe HZ in scenarios with planetesimals of 1 km and without typeI migration The planets with lowast are the Neptune-like planets Therow in bold corresponds to the results of Fig 6


SIM1 mdash mdash mdash mdash mdashSIM2 mdash mdash mdash mdash mdashSIM3 292 139 1537

lowast 3910 ASIM4 343 114 1732

lowast 3891 ASIM5 mdash mdash mdash mdash mdashSIM6 mdash mdash mdash mdash mdashSIM7 107 097 255 0 B


lowast 3891 A

the initial embryo population collides with the central stara 127 is accreted and a 4673 is ejected from the sys-tem While only between 1 and 2 planets remain inside thegiants the rest of them are located beyond them The plan-etesimals perturbed by the giants and outermost planetsare mainly ejected from the system (sim 844) Only a fewcollide with the central star (06) or are accreted (08)by the surrounding planets Only a sim 140 of the initialplanetesimals survives at the end of the simulations withsemimajor-axis almost greater than 20 au and mean higheccentricities Here again the planetesimal accretion is aninefficient process above all in the inner region of the discsince no planet originally located in the inner zone has suc-ceeded in accreting planetesimals coming from beyond thegiants

Inner region The inner zone of the disc in generalpresents one or two planets The trend shows that is morelikely to find one which in almost all the cases is a Super-Earth However in 3 of the 10 simulations one of innerplanets is a Neptune-like planet With the exception of theNeptune-like planets the inner planets are completely dryand present masses in a range between 11Moplus and 618Moplus In general this scenario achieved to form one PHP in 4 of 10simulations but as we mentioned before 3 of that 4 PHPsare Neptune-like planets which present high water contentsThe other PHP is a dry Super-Earth of 255Moplus The prop-erties of these PHPs can be seen in table 3

Outer region Beyond the giants there is a planet anda planetesimal population until sim 400 au In most of thesimulations the planets in this outer region of the disc arewater-rich planets and we do not see planets from the innerzone Moreover in most of the cases one of these planets isone of the Neptune-like planets originally located between26 au and 29 au

Finally at the end of the simulations the planetarysystems resulting from S4 in figure 1 are formed by one ortwo inner planets which are dry Super-Earths or water-richNeptune-like planets A Saturn-like is followed by a Jupiter-like planet and other several planets and planetesimals arebeyond them We find PHPs in 4 of 10 simulations One ofthem is a dry Super-Earth and the other 3 are Neptune-likeplanets that were not scattered The common result regard-ing these Neptunes in the HZ is that they did not suffer

collisions either from embryos or from planetesimals reasonwhy they could have conserved their primordial envelopesduring all the evolution Thus although their final locationsdefine them as PHPs it is in fact very likely that they arenot Therefore this scenario is not really interesting

524 Scenarios with planetesimals of 100 m

Main perturbers These simulations present two Jupiter-like planets of 149MJ and 5MJ located at around sim 99 auand sim 193 au during all the evolution respectively as themain perturbers of the systembull Icy giants Two Neptune-like planets of sim 13Moplus and sim

15Moplus are located between the inner embryo distribution andthe Jupiter analogues at 682 au and at 780 au respectivelyat the beginning of the simulations (see S1 in Fig 1)

General dynamical evolution As from the begin-ing the inner embryos get excited very quickly acquiringvery high eccentricities This causes that the orbits of theseobjects cross each other leading to accretion collisions ejec-tions and collisions with the central starbull From the beginning these collisions may occur betweendry and water-rich embryos since these two different popula-tions are not separated by a big perturber as in the previousscenarios

The fate of the Neptune-like planets is always the sameThe trend shows that one of them is ejected from the systemwithin the first million year of evolution while the other onewhich in most of the cases the inner Neptune stays in theinner zone The most common result among the simulationsis that this Neptune-like planet do not suffer collisions ofany kind Thus it could retain its primitive atmosphereWe show in Fig 7 six snapshots in time of the evolution ofSIM10 as an example of this scenario The initial conditionscorrespond to S1 in Fig 1 Although in these scenarios theplanetesimal population resides in the inner zone of the discthey are initially located between 5 au and 7 au just beyondone of the Neptunes Thus this population does not directlyinteract with the inner embryo population between 05 auand 4 au

Inner region The inner region of the disc presentsbetween 2 and 3 planets and in all the simulations one ofthese inner planet is one of the Neptune-like planets Theothers are dry or water-rich Earths Super-Earths or Mega-Earths with masses in a range between 102Moplus and 910oplus bull This scenario forms 6 PHPs 5 of them are unlike the onesin previous scenarios Super-Earths with very high watercontents what indicates that their accretion seeds come frombeyond the snowlineThe masses of these PHPs are between 264Moplus and 714Moplus The most interesting characteristics of these PHPs can beseen in table 4

Outer region This scenario does not present any kindof objects nor planets nor planetesimals beyond the giantsafter 200 Myr of evolutionbull Planetary systems in this case extend only up to the po-sition of the outermost Jupiter analogue at sim 20 au Thisis a direct consequence of the initial conditions which donot present an embryo and planetesimal population beyondthe gas giants Besides the strong gravitational interactionsbetween the two gas giants with the inner embryo and plan-etesimal population only manage to eject bodies outside the


12 M P Ronco et al

Figure 6 Time evolution of SIM7 of the scenario with planetesimals of 1 km and without considering type I migration The shadedareas represent the HZ regions already described in Section 1 Planets are plotted as colored circles planetesimals are plotted with blackdots and the giant planets which are a Saturn-like planet and a Jupiter-like planet in this case are represented in grey and blackrespectively The two big colored planets with black borders represent the two Neptune-like planets The sizes of the Saturn-like andJupiter-like planets are fixed and are not in the same scale for the rest of the planets The same happens with the planetesimal populationeach planetesimal adopts a size of half the size of the smallest embryo independently of its mass Then the size for the rest of the planetsis represented by a lineal relation between the mass and the size where the size for the smallest embryo is the double of the size of theplanetesimals and the size of the biggest planet is smaller that the size for a Saturn-like planet like in the previous scenario The colorscale represents the fraction of water of the planets with respect to their total masses This scenario form 4 PHPs within the HZ but 3 ofthose 4 are the Neptune-like planets In this particular case the planet that remain within the HZ is a Super-Earth of 255Moplus classifiedas a PHP class B

systems and not to disperse them towards the outer zonesof the disc

Mass remotion In this scenario and considering allthe simulations a sim 50 of the embryos collides with thecentral star a sim 35 is ejected from the system and a sim

23 is accreted by other embryos The planetesimal popu-lation is excited mostly by the gas giants but also by theNeptune-like planet and in only sim 20 Myr this populationis completely removed from the systembull As unlike the other scenarios the planetesimals are ini-tially located inside the giants an important percentage ofthem collide with the central star (sim 11) while the greatmajority is ejected from the system (89)

Only a 03 of the planetesimals are accreted during thewhole evolution of the systems Thus although the locationof these bodies is different from the other scenarios the ac-cretion of planetesimals is still a very inefficient process

In conclusion the planetary systems of this scenarioresulting from S1 in figure 1 present between 2 and 3 planetsin the inner zone of the disc one of which is a Neptune-likeplanet In 6 of 10 simulations we find a water-rich Super-Earth PHP class A or B Then we find a Neptune-like planetaround 3 au and 2 Jupiter analogues beyond sim 10 au Thisscenario results to be interesting since formed PHPs withhigh primordial water contents



Figure 7 Time evolution of SIM10 of the scenario with planetesimals of 100 m and without considering type I migration The shadedareas represent the HZ regions already described in Section 1 Planets are plotted as colored circles planetesimals are plotted with blackdots the two big colored circles with black borders are Neptune-like planets and the two black circles represent Jupiter analogues of149MJ and 5MJ The size of the Jupiter-like planets is fixed and is not in the same scale for the rest of the planets The same happenswith the planetesimal population each planetesimal adopts a size of half the size of the smallest embryo independently of its massThen the size for the rest of the planets is represented by a lineal relation between the mass and the size where the size for the smallestembryo is the double of the size of the planetesimals and the size of the biggest planet is smaller that the size for a Saturn-like planetlike in the previous scenario The color scale represents the fraction of water of the planets with respect to their total masses This is thefirst scenario that form water-rich PHPs

53 Comparison between planetesimal sizes

Here we highlight the most important differences and simi-larities between the reference scenarios formed by planetesi-mals of different sizes and without migration rates during thegaseous phase focusing mainly on the inner regions wherepotentially habitable planets form

bull Following panels S1 S4 S6 and S9 in figure 1 it can beseen that the mass of the main perturbers of the differentscenarios grows as the size of the planetesimals decreasesDuring the evolution of the N-body simulations those lessgas giant planets like the Saturn-like planet in scenariosof 100 km and the Saturn-like planet in scenarios of 1 kmsuffered a significant inward embryo-driven migration Thismigration particularly placed the Saturn-like planet in sce-

narios of 100 km at the outer edge of the HZ preventing theformation of PHPs

bull Regarding the global configuration of the different sce-narios at the end of the gaseous phase we can see that sce-narios formed with planetesimals of 100 m are quite differentfrom those formed with planetesimals of 1 km 10 km and100 km Just by looking at the initial conditions in Fig 1and comparing S1 with S4 S6 and S9 it can be seen that1) the planetesimal population is located inside the giantspositions instead of outside 2) the gas giants are much fur-ther away from the central star than they are in the otherscenarios and 3) a water-rich embryo population is directlynext to the dry one in the inner region of the disc These dif-ferences in the initial conditions conduct to different resultsin the final configurations of the N-body simulations


14 M P Ronco et al

Table 4 Principal properties of the planets that remain withinthe HZ in scenarios with planetesimals of 100 m and without typeI migration The row in bold corresponds to the results of Fig 7


SIM1 491 126 369 4636 ASIM2 682 171 714 5228 BSIM3 mdash mdash mdash mdash mdashSIM4 528 154 363 5270 BSIM5 mdash mdash mdash mdash mdashSIM6 mdash mdash mdash mdash mdashSIM7 mdash mdash mdash mdash mdashSIM8 426 162 439 5116 BSIM9 426 166 264 5290 BSIM10 457 138 334 4628 A

bull The mass remotion mechanisms this is embryo andplanetesimal remotion via accretion collisions with the cen-tral star and via ejections of the systems are quite similar inall the scenarios being the most efficient mechanism the ejec-tion both of embryos and planetesimals The remotion ofplanetesimals is completely efficient for the scenario formedby planetesimals of 100 m

bull Unless the outer regions of the scenario formed by plan-etesimals of 100 m which do not present planets or planetes-imals at all the other three scenarios present a mixed popu-lation of water-rich and some dry planets and planetesimalsuntill 200 Myr of evolution

bull The inner regions of all the scenarios are generally con-figured quickly within the first 1 Myr of evolution for sce-narios of 100 km and 10 km and within the first 10 Myrand 50 Myr for scenarios of 1 km and 100 m respectively

bull The efficiency in the formation of PHPs seems to behigher for scenarios formed from small planetesimals thanfrom bigger ones being completely null for planetesimalsof 100 km The final masses of these PHPs are in a widerange of values The masses of the PHPs of the scenarioformed from planetesimals of 100 m are in the range of theSuper-Earths while the masses of the PHPs in the otherscenarios are in the range of the Super and Mega-Earths Itis important to remark that in scenarios with planetesimalsof 1 km 3 of the 4 PHPs founded where Neptune-like planetsthat do not present any interest Regarding water contentsthe only PHPs of real interest that present water contentsare those formed in scenarios with planetesimals of 100 mFigure 8 shows these mentioned characteristics of the PHPsformed within all the four reference scenarios The left panelshows the total number of formed PHPs per scenario as afunction of the planetesimal size and the right panel showstheir masses per planetesimal size

54 Results sensitivity to scenarios with type I

migration

The previous sections described the general results of thedefault N-body simulations performed considering differentplanetesimal sizes and no type I migration Since one of theimportant results of PI was that the most favorable scenar-ios for the formation of SSAs were planetary systems withlownull type I migration rates we also performed N-body

0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

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100m 1km 10km 100km

Num

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f PH

P pe

r sce

nario

Planetesimal size

Total of PHPsWaterminusrich PHPs

Figure 8 Total number of formed PHPs and number of water-rich PHPs per scenario as a function of the planetesimal size (leftpanel) and masses of the formed PHPs per planetesimal size

simulations for those SSAs at the end of the gaseous phasethat presented type I migration rates reduced to a 1 andto a 10 The goal of this section is not to describe in de-tail as we did before the particular results of each of thesesimulations but to give a global view of what are the ef-fects that this phenomenon produces and how the previousresults could change because of it

We performed 10 different simulations for each scenarioS2 S3 S5 S7 S8 and S10 that represent planetary systemswith type I migration rates reduced to a 1 and to a 10Those simulations with type I migration reduced to a 1do not present significant differences from the ones withouttype I migration For example the initial conditions of Slowast

9

(100 km without type I migration) and Slowast10

(100 km withtype I migration reduced to 1) are very similar and thusthe final configurations of the N-body simulations are simi-lar too Although the mass of the Saturn-like planet of Slowast

10

is smaller than the Saturn-like planet of Slowast9in a factor sim 3

the main result is that none of these scenarios form planetsin the inner region of the disc Something similar occurs be-tween S6 (10 km without type I migration) and S7 (10 kmwith type I migration reduced to 1) although S6 presents2 gas giants as main perturbers and S7 presents only 1 Themain results show that the inner region of the disc presentsbetween 1 and 2 dry planets with masses in the range of theSuper and Mega-Earths and that some of them between 3and 5 are completely dry PHPs In the same way the resultsconcerning scenarios of 1 km with type I migration reducedto a 1 and without type I migration are also similar themain result shows that both scenarios form dry PHPs Re-garding scenarios of 100 m both kinds of simulations withmigration reduced to a 1 and without type I migrationform water-rich PHPs

The results of PI showed that only very few planetarysystems formed from planetesimals of 100 m and 10 km wereable to form SSAs at the end of the gaseous phase with typeI migration rates reduced to a 10 Although this is an un-common result we still performed simulations for this sce-narios whose initial conditions are represented by S3 (forscenarios with planetesimals of 100 m) and by S8 (for sce-narios with planetesimals of 10 km)

For scenarios with planetesimals of 10 km the type Imigration during the gaseous-phase favors the giant planetto form closer to the central star In this case this result isdetrimental to the formation of PHPs However for scenar-ios with planetesimals of 100 m it seems to be just the otherway around We will describe these situations in more detailbelow



Figure 9 shows initial and final configurations of plane-tary systems formed from planetesimals of 100 m (bottom)and 10 km (top) The left panel shows the initial condi-tions at the beginning of the N-body simulations and theright panel shows the N-body simulations results On theone hand comparing the two scenarios of 10 km with andwithout type I migration we can see that type I migra-tion during the gaseous stage leads the giant planet of thesystem to locate closer to the central star than the inner gi-ant is in the scenario without migration This result whichis a natural consequence of type I migration and the factthat the position of the giant could also change during thepost-gas evolution of the system due to an induce migra-tion via the gravitational interactions with the surroundingembryos and planetesimals is unfavorable for the formationof PHPs In fact we already described that the simulationswith 10 km planetesimals and no type I migration were ableto form 5 dry PHPs However the 10 simulations developedwith 10 km planetesimals and type I migration rates reducedto a 10 were not able to form PHPs at all On the otherhand comparing the two scenarios of 100 m with and with-out type I migration we can see that in spite of the fact thattype I migration manages to locate the giant further insideregarding the position of the innermost giant in the scenariowithout type I migration in this case this effect favors theformation of PHPs In fact as we have already comment insection the simulations with planetesimals of 100 m andno type I migration were able to form 6 water-rich PHPsHowever the 10 simulations developed with type I migrationrates reduced to a 10 were able to form 11 PHPs 9 areclassified as PHPs class A and 2 class B The final massesof these PHPs are between 093Moplus and 373Moplus with watercontents between 243 to 335 It would seem then thatthere is a preferential location for the inner giants in termsof the efficiency in the formation of PHPs We will return tothis point later

To conclude scenarios with type I migration rateshighly reduced do not present significant differences fromthose scenarios without type I migration However scenar-ios with type I migration rates that represent a 10 of theTanaka et al (2002) type I migration rates show differentresults for planetary systems formed from small and bigplanetesimals This phenomena seems to favor the forma-tion of PHPs in scenarios with small planetesimals but def-initely avoids their formation in scenarios formed from bigplanetesimals

55 Results sensitivity to scenarios with

gap-opening giants

Another important phenomenon that could modify the finalconfiguration of the planetary systems is the type II migra-tion The type II migration occurs when a planet is massiveenough to fulfill the Crida et al (2006) criteria therefore toopen a gap in the disc which depends not only on the massand the position of the planet but also on the Hill radius ofthe planet the viscosity and the height scale of the disc (seePI for more details) This phenomenon which was includedin the model of planet formation of PI produced an inwardmigration during the gaseous phase only for this gap open-ing giant and an increase in its mass due to the accretionof gas and other planets in its path It is important to re-

Table 5 This table shows the range of masses for the giantplanets that opened and did not opened a gap in the scenar-ios with planetesimals of different size The ranges are in Saturnand Jupiter masses due to the fact that we find both types ofplanets in scenarios with 100 m 1 km and 10 km As it can beseen in the table the masses of those giants that opened a gap inthe disc during the gaseous phase is greater than the other ones

Scenario Giants that Giants thatdid not opened a gap opened a gap

100 m 081MS minus 5MJ 563MJ minus 892MJ

1 km 048MS minus 4MJ 723MJ

10 km 133MS minus 08MJ 188MJ minus 210MJ

100 km 056MS minus 190MS 323MJ

mark that the location of the gas giant at the end of thegaseous phase is not a consequence due only to the type IImigration but because once type II migration is turned onthe planet which moves inward accretes other planets in itspath Recall from PI that the resulting semimajor-axis of afusion between two planets is given by the conservation ofangular momentum (see eq 33 in PI)

It has been proved by many models that an inward gi-ant planet migration shepherds some planets captured atinner mean motion resonances with the giant which arethen dispersed into outer orbits leading to the formationof planetary systems with an inner compact planet pop-ulation and an outer much more scattered planet pop-ulation (Fogg amp Nelson 2005 2006 Raymond et al 2006Fogg amp Nelson 2007 Mandell et al 2007) ParticularlyFogg amp Nelson (2009) have showed that including a viscousgas disc with photoevaporation in an N-body code and al-though the type II inward migration it is still possible toform planets within the HZ in low-eccentricity warm-Jupitersystems if the giant planet makes a limited incursion into theouter regions of the HZ Our model of planet formations doesnot yet include the scattering effects from a migrating giantplanet or the gravitational interactions between planets thatcan lead to resonance captures Therefore the results of ourN-body simulations could be affected by the initial condi-tions which could had overestimated the number of planetsin the inner zone of the systems at the end of the gaseousphase

The main difference between the opening gap giantplanets with those which did not is their final mass Gapopening planets are in general more massive than thosethat did not Table 5 shows the ranges of masses in sce-narios with different planetesimal sizes for those giants thatopened and not opened a gap It is important to note thatthis table was drawn up from the initial conditions of Fig1 and from 4 more initial conditions for the different sizesthen used to run more N-body simulations which present atleast one gas giant that managed to open a gap in the disc

In order to see if there were significant differences be-tween the default scenarios and scenarios without type Imigration but with at least one giant that managed to turn

on type II migration we run more simulations We choseone initial condition for each planetesimal size that presentsa gas giant that has opened a gap This is we chose 4 newinitial conditions and we run 10 simulations per each of these4 scenarios changing as we did before the seed Figure 10


16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)

1 10 100Semimayor-axis [au]

Initial Conditions


Final N-body simulations results

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of w

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Figure 9 Initial conditions (left panel) and final N-body simulations results (right panel) of scenarios formed with planetesimals of100 m and 10 km without type I migration ( fmigI = 0) and with a type I migration reduced to a 10 ( fmigI = 01)

100m - fmigI=0 - Type II Mig OFF

100m - fmigI=0 - Type II Mig ON

100km - fmigI=0 - Type II Mig OFF

100km - fmigI=0 - Type II Mig ON


Initial Conditions



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04

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Figure 10 Initial conditions (left panel) and final N-body simulations results (right panel) of scenarios formed with planetesimals of100 m and 100 km without type I migration (fmigI = 0) and with and without type II migration The black points represent Jupiter-likeplanets while the grey points represent Saturn-like planets The black points with a white cross are those giant planets that had managedto open a gap in the disc during the gaseous phase evolution and had migrated inward

shows in a similar way as figure 10 does a comparison be-tween simulations of the two extreme reference scenarios of100 km and 100 m without type I and type II migration andscenarios of 100 km and 100 m without type I migration butwith the type II migration activated The left panels showthe initial conditions at the begining of the N-body simu-lations this is the final results at the end of the gaseousphase and the right panels show the N-body simulationsresults after 200 Myr of evolution

A global result found in all the new simulations is thatdue to the high mass of the new opening gas giants thosescenarios which present a planet and planetesimal popula-tion beyond the gas giants do not present a mix of dry andwater-rich material in the outer zone of the disc This is theinner dry embryos are directly ejected from the systems andare not scattered towards the outer regions

Regarding the efficiency in forming PHPs in scenarioswith planetesimals of 100 km it still remains completelynull None of the simulations with planetesimals of 100 kmwithout type I migration and with a planet that opened agap managed to form PHPs The gap opening Jupiter-likeplanet in the scenario of 100 km of figure 10 (the big black

point with a white cross) have 322MJ and is located aroundsim 2 au while the Saturn-like planet in the scenario of 100 kmwhich did not opened a gap presents a mass of 05MJ andis located initially at 296 au and finally at sim 17 au Forscenarios of 10 km the production of PHPs in the 10 simu-lations performed with a planet that opened a gap is lowerthan in the default scenarios This scenario formed 2 dryPHPs instead of 4 with masses of sim 2Moplus and 11Moplus Follow-ing with scenarios of 1 km in this case the PHP productionwas similar Both scenarios with and without an opening gasgiant produced 5 PHPs with masses ranging from 09Moplus tosim 6Moplus Finally scenarios with planetesimals of 100 m and agiant planet that opened a gap formed 9 PHPs with massesbetween 528Moplus and 11Moplus 3 more than in the default caseIn these scenarios the gap opening Jupiter-like planet has amass of 563MJ against the 149MJ of the Jupiter-like planetthat did not opened a gap and the first one is located around638 au while the second one is located at 10 au This analysisof the comparison between the whole group of simulationswith and without a gap opening planet shows us that thetype II inward migration coupled with the increase of themass of the giant seems to act in a similar way as type I



migration do Although this is the outcome of our simula-tions more simulations should be carried out to confirm thistrend These effects seem to favor the formation of PHPs forscenarios with small planetesimals and to be detrimental forscenarios formed from big planetesimals

56 Possible range of mass and location of the

inner giant to form PHPs

With the aim of exploring the possible range of masses andsemimajor-axis of the inner giant planet to form PHPs weperformed histograms taking into account all the simula-tions in where we found PHPs without differentiating themaccording to the size of the planetesimals the type I mi-gration rates or the gap opening As it is shown in the topleft panel in Fig 11 more than the 70 of the PHPs wereformed in planetary systems with the innermost giant planetlocated between 3 au and 6 au Then following the top rightpanel of Fig 11 it seems that a possible range for the massof the inner giant planet to form PHPs could be between1MJ and 3MJ since we find more than the 60 of the PHPswithin these values These top panels do not distinguish fromplanetary systems with one two or three gas giants butthe middle and the bottom histograms do If we distinguishplanetary systems with only one gas giant from those withtwo gas giants we find some differences For those systemswith one giant a possible range for the location and massfor the giant planet to form PHPs is between 3 au and 4 auand between 4MJ and 5MJ respectively However for thoseplanetary systems with two gas giants the majority of thePHPs were formed in systems with the inner giant around6 au and between 1MJ and 3MJ

It is however important to remark that these resultsare obtained only considering the results of the developedN-body simulations which do not present a detailed sampleof the semimajor-axis distribution of the inner gas giantsbetween 15 au and 10 au Thus although the sample of N-body simulations carried out in this work helps us to explorethe ranges of mass and position for the inner giant planetto form PHPs this sample is not enough to determine anoptimal mass and an optimal position for the inner giantplanet to form PHPs even more if we want to discriminatebetween systems with one and two giants In order to de-termine a precise value of the mass and location of a giantplanet to efficiently form PHPs it is necessary to developa greater number of N-body simulations from a better andalso greater selection of initial conditions Since we countwith these initial conditions from the results of PI we willreturn to this analysis in future works

6 CONTRAST WITH THE OBSERVED

PLANET POPULATION

One of the goals of this work is to contrast the final config-urations of our SSAs with those planetary systems alreadyobserved that present at least some similarities with thispopulation It is however important to remember that thegoal of PI and this paper is not to reproduce the currentexoplanet population but to generate a great diversity ofplanetary systems particularly SSAs without following anyobservable distribution in the disc parameters Then once

0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11

All the Planetary Systems

Planetary Systems with only 1 gas giant

Planetary Systems with 2 gas giants

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tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

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of P

HP

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01 015

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of P

HP

Semimajor-axis [au]

0 005 01

015 02

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Frac

tion

of P

HP

Mass [MJ]

Figure 11 Histograms showing the fraction of PHPs formed forall the simulations performed as a function of the semimajor-axisof the innermost giant planet (top left panel) and as a functionof the mass of the innermost giant planet (top right panel) Themiddle and the bottom panels show the same histograms butconsidering those performed simulations with only one gas giantand with two giants respectively

the late-accretion stage has finished and the final configu-rations of SSAs have been obtained our goal is to contrastour results with the current planetary systems populationwhich present on the one hand gas giants like Saturn orJupiter beyond 15 au and on the other hand rocky planetsin the inner regions particularly in their HZs To do this weare only going to take into account those planetary systemsaround solar-type stars with masses in the range of sim 08M⊙

to 12M⊙ which present at least one giant planet beyond15 au and those planetary systems which present planets ofinterest near or in the HZ As we have already mentionedat the beginning planetary systems similar to our own havenot yet been discovered However it is possible that plane-tary systems with detected gas giants present rocky planetsin the inner zone of the disc and also it could be possiblethat planetary systems with detected rocky planets presentgas giants not yet discovered particularly if they presentlarge semimajor-axis and low eccentricities

This section is then divided by taking into account thecomparison with two different observed planet populationsthe PHPs population and the gas giant planet population

61 PHPs population

Since the CoRoT (Barge et al 2008) and Kepler(Borucki et al 2010) missions began to work our con-ception of the existence of other worlds has changed


18 M P Ronco et al

Table 6 Potentially habitable planets around solar-mass starsfrom httpphlupredu planet candidateunconfirmed lowast es-timated value for rocky composition when not available

Name Mass [Moplus] Radius [Roplus] Spectral Type

Kepler-452 b 47lowast 16 GKepler-22 b 204lowast 24 GKepler-69 c 60lowast 17 GKepler-439 b 195lowast 23 GtauCet-e gt 43 16lowast GHD-40307 g gt 71 18lowast KKOI 723501 14lowast 11 GKOI 642501 35lowast 15 GKOI 722301 37lowast 15 GKOI 717901 15lowast 12 GKOI 445001 99lowast 20 GKOI 405401 101lowast 20 GKOI 758701 149lowast 22 GKOI 673401 131lowast 21 GKOI 713601 174lowast 23 GKOI 667601 70lowast 18 FKOI 547501 50lowast 17 FKOI 755401 99lowast 20 FKOI 523601 99lowast 20 FKOI 410301 144lowast 22 KKOI 520201 73lowast 18 FKOI 634301 84lowast 19 FKOI 747001 84lowast 19 KKOI 527601 152lowast 22 KKOI 734501 211lowast 24 FKOI 445801 259lowast 25 FKOI 394601 207lowast 24 FKOI 704001 147lowast 22 F

considerably During their working years these missionswere able to find a very high number of rocky planets withsimilar sizes to that of the Earth as never before givingrise to the exploration of this particular type of exoplanetsAt present both missions have ceased to function butthere are others on the way that will help to continuethe study detection and characterization of rocky planetssuch as TESS (Transiting Exoplanet Survey Satellite) andCHEOPS (CHaracterising ExOPlanets Satellite) TESS(Ricker et al 2010) is expected to discover thousands ofexoplanets in orbit around the brightest stars in the skythrough the transit method and CHEOPS (Broeg et al2013) which will complement TESS will characterizealready confirmed exoplanets using photometry of very highprecision to determine the exact radius of planetary bodiesof known mass between 1Moplus and 20Moplus

Till now Kepler has discovered a total of 2244 candi-dates 2327 confirmed planets and 30 small less than twiceEarth-size HZ confirmed planets around different type ofstars while the K2 mission which is still working hasdiscovered 480 new candidates and 292 confirmed planets(httpswwwnasagovmission_pageskepler)

In order to contrast our simulations with observationswe consider those N-body simulations that after 200 Myr ofevolution present PHPs Then we contrast the simulatedPHPs with the observed ones only considering those PHPsfrom httpphlupredu that orbit stars with masses inthe range of sim 08M⊙ to 12M⊙ (see table 6) It is importantto remark that we are not making a precise statistical com-

parison between the observed and the simulated populationof PHPs but trying to determine if there are observed plan-ets in the range of masses and radii of the ones we form in oursimulations within the HZ Figure 12 shows a semimajor-axisvs mass plane of these two populations Most of the massesof the observed exoplanets are estimated considering a rockycomposition when not available (see httpphlupredu)The blue and light-blue shaded areas in Fig 12 represent theclass A and B and C HZ regions respectively as we havedefined in section 51 As we can see in Fig 12 there is anoverlap of both populations the real one and the simulatedone in both panels that represent the mass vs semimajor-axis and the radius vs semimajor-axis diagrams To plot theradius of our simulated planets we consider physical densi-ties of 5 g cmminus3 The black rectangle delimits these regions ofoverlap between 0849 au and 1356 au and between 13Moplus

and 204Moplus (panel a) and between 0849 au and 1356 auand between 11Roplus and 25Roplus (panel b) in where we canfind both simulated and observed PHPs Particularly and tillnow the only confirmed exoplanets that remain within thisregion are Kepler-452 b and Kepler-22 b the rest are stillcandidates awaiting for confirmation Kepler-452 b whichpresents a radius of 163

+023

minus020Roplus is until now the fist planet

that presents the longest orbital period for a small transit-ing exoplanet to date and orbits a star of 1037

+0054

minus0047M⊙ at

1046+0019

minus0015au (Jenkins et al 2015) Kepler-22 b is the only

planet orbiting a star of 0970+006

minus006M⊙ in the inner edge of

its HZ at 0849+0018

minus0018au with a radius of 238

+013

minus013Roplus and

a mass estimation of lt 36Moplus (Borucki et al 2012)It is important to clarify that this comparison we are

doing here between the simulated PHPs in some SSAs andthe selected observed PHPs that orbit solar-type stars isnot a direct comparison since none of these observed PHPsis still part of a planetary system similar to our own thisis none of them has a Saturn or Jupiter analogue associ-ated beyond 15 au Besides it is worth remarking that thecombination of different detection techniques that allow usto measure both the mass and the radius of a planet haveshown that many of the planets discovered with Kepler couldhave compositions very different from those of the terrestrialplanets of our Solar system (Wolfgang et al 2016) Partic-ularly planets with M gt 6Moplus could include a mixture ofgas-rich and rocky planets which in fact could be very dif-ferent from the PHPs we form within our simulations How-ever we can see that we are able to form PHPs in a regionof the semimajor-axis vs mass and semimajor-axis vs radiusdiagrams which are already explored The lack of PHPs inmore extended orbits this is until sim 2 au is possible due tothe fact that the HZ planets associated with G-type mainsequence stars have longer orbital periods which means thatmore data for the detection is required in comparison withcooler K and M-type stars (Batalha 2014) It is then ex-pected that the advances in future missions will help us todetect a greater number of PHPs in these regions

62 Gas giant planet population

The population of giant planets in particular hot-Jupiterswhich was at first the most abundant and easy to be detectedby both the Doppler and the transit techniques is today infront of the enormous amount of terrestrial-type exoplanets



0

5

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15

20

25

30

04 06 08 1 12 14 16 18 2

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s [M

J]

Class C ZoneClass A and B Zones

Simulated PHPsCandidate PHPs

Confirmed

0

02

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06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

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2

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3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

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Keplerminus452 b

Keplerminus22 b

Figure 12 Semimajor-axis vs mass plane for the simulated PHPspopulation (colored diamonds) and the observed PHPs aroundsolar-mass stars (black points) The colorscale represents the ec-centricity data that we have for the simulated PHPs but not forthe majority of the observed ones The shaded areas representour definition of the HZ for a 5Moplus planet and the black squaredelimits the region in where we find both kinds of planets the ob-served ones and the simulated ones Kepler-452 b and Kepler-22b are the only two confirmed exoplanets within these region

discovered by Kepler a rare population (Batalha 2014) It isthen believed that the rate of occurrence of giant planets isgreater beyond 2 au and particularly around the location ofthe snowline (see Fischer et al 2014 and references therein)However it is still not so easy to detect gas giant planetson wide and circular orbits since Barnes (2007) showed thatplanets with eccentric orbits are more likely to transit thanthose planets with circular orbits with the same semimajor-axis by a factor of (1 minus e2)minus1 where e is the eccentricity ofthe planet

To date sim 500 gas giant planets (more massive thanSaturn) orbiting solar-type stars with masses in the rangeof sim 08M⊙ to 12M⊙ were confirmed From this only sim 142

are orbiting their star beyond sim 15 au However none ofthese giant planets has yet an associated rocky planet inthe inner regions of the system much less in the HZ Thisdoes not strictly mean that these rocky companions in theinner zone of the disc do not exist It is possible that thegiants are camouflaging the signals of the rocky planets andthat this makes it more difficult to detect them Thus againwe need to be careful when comparing the giant planets inour simulated planetary systems with those confirmed giantsbeyond 15 au since they may still be part of planetarysystems very different to Solar system analogues

In order to contrast our simulated gas giant popula-

tion with the observations we are only going to take intoaccount those single or multiple planetary systems withone or more gas giants with masses greater than 03MJ

beyond 15 au We also take into account those multipleplanetary systems with more than two gas giants beyond1 au Figure 13 shows semimajor-axis vs eccentricity vsmass planes for planetary systems with only one gas gi-ant (top panels) and with two gas giants (bottom panels)(httpexoplaneteu Schneider et al 2011) Each dia-gram shows the confirmed gas giant planets with massesgreater than Saturn in colored squares (planetary systemswith only one gas giant) and points (planetary systems withtwo gas giants) along with the gas giant planet populationof all our simulations Those gas giants in planetary systemsthat present PHPs at the end of the N-body simulations areshown in black squares (planetary systems with only one gasgiant) and points (planetary systems with two gas giants)while the gas giants in planetary systems that did not formPHPs are represented in grey

In general it can be seen that in those planetary sys-tems with only one gas giant (top panels) the range ofsemimajor-axis and masses of the simulated gas giants ispretty well represented by the confirmed population of gi-ant planets However this does not seem to happen for theeccentricity range The least eccentric confirmed exoplanetof the selected sample has an eccentricity of 002 whereas weform planets with up to more than one order of magnitudeless eccentric This population of planetary systems withvery low eccentric gas giants seems to be not yet detectedand in particular it seems to be the one that is able to formPHPs It is also interesting to remark that those simulatedgas giants with eccentricities higher than 02 semimajor-axisgreater than sim 15 au and masses greater than sim 03MJ andwhich are overlapped with the observed and confirmed gasgiants are part of planetary systems that were not efficientin forming PHPs (grey squares)

Something similar seems to happen in planetary sys-tems with two gas giant planets (bottom panels) In this kindof planetary systems although the range in mass is well rep-resented by the confirmed population of giant planets theranges of semimajor-axis and eccentricities are not In thissense we are forming a population of planetary systems withvery low eccentric gas giants in wider orbits extended untilsim30 au that seems to be not yet detected And again thispopulation seems to be efficient in forming PHPs

We then ask ourselves if in fact this population of verylow eccentricities and wide orbits is a natural result of theevolution of the simulations and therefore is indeed a pop-ulation not yet detected or if it is a consequence of the initialconditions imposed by our model of planet formation whereat the beginning and at the end of the gaseous stage the plan-ets are in circular and co-planar orbits A detailed study ofthe evolution of the semimajor-axis and eccentricities of ourgiants shows us that the gravitational interactions betweenthese bodies and the rest of the planets of the system whichare consider during the evolution of the N-body simulationslead to significant changes in their eccentricities althoughdo not lead to planet-planet scattering events that involvegas giant planets For planetary systems with only one gi-ant planet the registered changes in the giants eccentricitiesduring the late-accretion stage reach values higher than anorder of magnitude of the initial eccentricities For planetary


20 M P Ronco et al

systems with two gas giants these changes are even greaterand the eccentricities can increase even more Then thispopulation of planetary systems with giant planets of loweccentricities is within the framework of having consideredinitial (this is at the beginning of the N-body simulations)quasi-circular and co-planar orbits for the giants a naturalresult of its own orbital evolution It is however importantto highlight that more realistic treatments for the gaseousphase evolution of these systems could lead to different gasgiant planet configurations Our model of planet formationfor this stage does not include the gravitational interactionsbetween planets These interactions could lead to encoun-ters between giants increasing their eccentricities as the gasdissipates and scattering other planets in their way On theother hand this effect could also cause planets to be trappedin mean motion resonances Both situations can in turn leadto different final configurations for the planetary systems atthe end of the late accretion stage

7 DISCUSSION

In the present work we carry out N-body simulations aimedat studying the late accretion in Solar system analogues(SSAs) and the formation of potentially habitable planetsand water delivery To interpret the results of our simula-tions it is very important to mention some points concerningthe evolution of our systems during the gaseous phase as wellas once the gas has dissipated

According to that mentioned in Section 3 we select 10scenarios of work in order to carry out our analysis con-cerning the late-accretion stage in SSAs The evolution ofsuch systems during the gaseous phase was obtained fromthe work developed by Ronco et al (2017) PI

It is worth noting that as it was mentioned in Section2 the model of planet formation of PI from which we ob-tain the initial conditions to develope N-body simulationspresents limitations Some of these limitations are the con-sideration of isothermal discs and the corresponding type Imigration rates for such discs and the lack of the calculationof the gravitational interactions between embryos The con-sideration of more complex physics during the gaseous phasecould alter the final results of embryo and planetesimal dis-tributions at the end of the gaseous phase and thus signif-icantly change our results on the final planet mass theirlocations and water contents at the end of the late-accretionstages

Particularly as it was mentioned before type I mi-gration rates for non-isothermal discs could be very dif-ferent for those in isothermal discs such as they couldstop or even change the migration direction of the planets(Paardekooper et al 2010 2011) The consideration of grav-itational interactions between planets during the gaseousphase could allowed the planets to be trapped in mean mo-tion resonances (MMR) and this could avoid a fast orbitaldecay into very inner zones of the disc (Masset amp Snellgrove2001) Moreover the consideration of this effect as it wasshown by Morbidelli amp Crida (2007) within the frameworkof our Solar system is that if Jupiter was able to open a gapin the disc and migrate inward through type II migrationand if at the same time Saturn was able to migrate faster to-wards Jupiter they could have been locked in a 23 (MMR)

and could have stopped or even reversed the migration ofboth planets together On the other hand the interactionsbetween giant planets in the inner regions of the disc whenthe gas is already removed due to photoevaporation can leadto the ejection of one or two planets exalting the eccentric-ities of the remaining ones Then all these considerationswhich were not taken into account in PI can change thefinal results of our N-body simulations

It is also important to remark that the 10 scenarios ofwork chosen to carry out our investigation show differencesconcerning the number and the physical and orbital prop-erties of the giant planets that host On the one hand 5 ofthose 10 scenarios start the post-gas phase with a single gi-ant planet whose mass ranges from 045MS to 4MJ On theother hand the other five systems of work host two gaseousgiants at the beginning of the post-gas stage most of whichshow masses of the order of the Jupiterrsquos mass

In those 5 scenarios of work that host 2 gaseous giantsthe initial orbital separation between them ranges between41 and 117 mutual Hill radii It is worth noting that noneof the N-body simulations developed in the present researchproduced close encounters between the massive giant planetsof the system which seems to be consistent with the worksdeveloped by Gladman (1993) Chambers et al (1996) andMarzari (2014) Obviously an analysis concerning the sta-bility of the system of 2 giant planets is more complex sincethey are immersed in a population of embryos and planetes-imals in our scenarios of study However this simple consid-eration allows us to understand the absence of close encoun-ters between the massive giant planets in all our N-bodysimulations

From this we must remark that the SSAs formed in thepresent research did not experiment resonance capture pro-cesses during the gaseous phase while they did not undergostrong dynamical instability events that involve the massivegiant planets Thus the proposed analogy between the sys-tems formed in our simulations and the Solar system shouldbe carefully interpreted if scenarios of formation such as theGrand Tack model (Walsh et al 2011) and the Nice model(Tsiganis et al 2005) are accepted Studies concerning theformation and evolution of systems that host giant planetsand undergo strong instability events between them will beaddressed in future works

On the other hand it is interesting to note that ourrocky planets grow fundamentally from the giant impactsthey undergo once the gas disk dissipates a process thatlasts about 100 Myr Late formation stages like the ldquolateveneerrdquo or the ldquoheavy bombardmentrdquo on our Earth are notpresent in our simulations If our Solar system experiencedthree stages of formation given by the stage of giant im-pacts the late-veneer and finally the heavy bombardment(see Raymond et al 2013 and references therein) the plan-ets of our simulated systems only suffer the first stage notbeing affected by later ones in which a population of plan-etesimal remnants reshape them

Our N-body simulations form after 200 Myr of evolu-tion planetary systems similar to our own with rocky plan-ets in the inner zone of the disc and at least one gas giantbeyond 15 au One of the goals of this work is to analyzethe formation of planets within the HZ The simulations per-formed were able to form two very different kinds of PHPsdry and water-rich (this is planets that come from beyond



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Planetary systemswith two gas giants

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1

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J]

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10

100



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J]

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1

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0001 001 01 1


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ass

[MJ]

Eccentricity

1

10

100

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imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity

Figure 13 Comparison between the observations and the developed simulations This plot shows eccentricity vs semimajor-axis massvs eccentricity and mass vs semimajor-axis planes for planetary systems with one gas giant (top panels) and two gas giants (bottompanels) The colored squares and points represent the locations of the gas giant planets in planetary systems with only one and two gasgiants respectively beyond 15 au and with masses greater than Saturn The black squares and points represent the locations of thegas giants in simulated planetary systems that present PHPs in planetary systems with only one and two gas giants respectively Thegrey squares and points represent the locations of the gas giants in simulated planetary systems that do not present PHPs in planetarysystems with only one and two gas giants respectively The colored areas represent the range in semimajor-axis eccentricity and massin where we find planetary systems with one or two gas giants that were able to form PHPs

the snowline or planets that have accreted at least an em-bryo from this region) Taking into account the 4 defaultscenarios described in this work the most common type ofPHPs seems to be the dry ones since we formed them in 3 ofthose 4 scenarios However considering the final number ofPHPs formed within all the developed simulations and con-sidering that the most common result from PI was to formSSAs in scenarios with planetesimals of 100 m it seems thatthe most common PHPs are the water-rich ones instead ofthe dry ones since a 65 of the whole PHPs population isformed by water-rich planets (almost all of them formed inscenarios with planetesimals of 100 m) and a 35 is formedby dry planets

One of the limitations of our model which is importantto remark is that the high water contents in those water-rich PHPs must be considered as upper limits due to the factthat we do not take into account water loss during impactsHowever the great majority of our water-rich PHPs do notsuffer from a great number of giant collisions Thus theirprimordial water amounts could survive during the wholeevolution Moreover Dvorak et al (2015) showed that thewater retained in fragments after a collision remains on thetarget for impact velocities v le 13vesc while more than 80of the water stays as well in strongly inclined hit-and-runscenarios Another interesting result is that due to the factthat all our water-rich PHPs come from beyond the snow-line their water contents are primordial This result is inconcordance with the results of Salvador Zain et al (2017)who also found that all water contents of the HZ planetsin systems that harbor a Jupiter-mass planet around thesnowline are also primordial

If although the water contents have to be consideredas upper limits they could have remained in their greattotality until the end of the evolution We wonder thenwhat is the effect of such water contents in the habitability

Some authors argue that too large amounts of water maybe detrimental for the potential habitability (Alibert 2014Kitzmann et al 2015) These authors claim that a too largewater layer can lead to the existence of a high pressure icelayer at the bottom of the global ocean and this could pre-vent the Carbonate-silicate cycle from occurring This couldbe a problem since the Carbonate-silicate cycle is funda-mental to maintain the temperature of the surface for longperiods of time (Alibert 2014) and could be particularlyproblematic for stars such as the sun which suffer varia-tions in their flux at relatively short timescales MoreoverKitzmann et al (2015) showed that an unstable CO2 cyclecould limit the extension of the HZ for ocean planets Otherstudies such as the work developed by Abbot et al (2012)found that the weathering behavior is not sensitive to theland fraction for those planets which are covered by oceansas long as the land fraction is greater than 001 This resultsuggests that those partially covered water planets may havea HZ similar to a planet which is not Moreover althoughthese authors indicate that water worlds could have narrowerHZ they could be capable of ldquoself-arrestingrdquo a moist green-house effect and could become planets with partial oceancoverage Recently Noack et al (2016) also concluded thatwater-rich planets may be habitable and could allow thedeveloping of live as we know it if these water-rich plan-ets present shallow oceans a low planet mass or a highsurface temperature They find that small-massive planetswith less than one Earth mass can be habitable even forsubstantial amounts of water However Super-Earths cov-ered with oceans can only be considered as habitable planetsif they contain low percentages of water by mass Althoughthis is a topic that is still under debate and that the onlyexistence of liquid water is definitely not enough to deter-mine habitability water-rich planets seem to be commonin the Universe (Raymond et al 2004 Simpson 2017) We


22 M P Ronco et al

particularly find them in planetary systems without gas gi-ant planets in high (de Elıa et al 2013) and low-mass discs(Ronco amp de Elıa 2014 Ronco et al 2015) in planetary sys-tems around low-mass stars (Dugaro et al 2016) and inplanetary systems with different gas giant planets as mainperturbers (Salvador Zain et al 2017) and because of thatwe still consider they present a particular interest

Another important topic that could affect the habit-ability of our PHPs due to the possible absence of platetectonics is their final masses since the great majority ofthem are Super-Earths (59) and to a lesser extent Mega-Earths (31) with masses very different from that of ourplanet However Valencia et al (2007) show that plate tec-tonics which are a necessary condition for the developmentof life can also be found in Super-Earths These authorsshow that as the planetary mass increases the subductionprocess and therefore the plate tectonics becomes easierTherefore massive Super-Earths are very likely to exhibitplate tectonics and therefore to keep life possible On theother hand Morard et al (2011) showed that for massesgreater than 2Moplus the Super-Earths cores could be entirelysolid this precludes the existence of a liquid metallic coredriven by magnetic fields The lack of magnetic fields is detri-mental to the potential habitability given that they are thosethat favor the existence and permanence of an atmospherefor long periods of time

Finally and following with our model limitations webelieve that a more realistic treatment of the collisions in-cluding fragmentation and hit-and-run collisions (Chambers2013) rather than assuming that all impacts lead to a perfectmerger of the colliding bodies should be included in futureworks Quintana et al (2016) presented results of N-bodysimulations of the late stage formation of Earth-like plan-ets around sun-like stars with Jupiter and Saturn analoguesperturbing the system and including a collision model thatfollows Leinhardt amp Stewart (2012) and Chambers (2013)They found that globally when fragmentation is consideredthe number of final planets and their masses are similar tothose formed in the standard model without fragmentationand hit-and-run collisions However the collision history ofthe final planets is different and the accretion timescales areusually doubled in the fragmentation simulations Thus weconsider that this improvements in the treatment of colli-sions will allow us to determine in more detail the orbitaland physical properties of the planets formed in our simu-lations particularly the final masses and amounts of waterby mass of the formed PHPs

8 CONCLUSIONS

In this work we carry out a great number of N-body simu-lations aimed at analyzing the global formation of the post-gas phase of SSAs particularly focusing in the formation ofrocky planets within the HZ and their final water contentsThe initial conditions to start the simulations were providedby the results of PI (Ronco et al 2017) In that first workwe analyzed the evolution during the gas-phase of differ-ent planetary systems particularly of those which were atthe end of the gaseous phase similar to our own consideringdifferent planetesimal sizes and different type I migrationrates

The main results of our simulations which are analyzedfor planetary systems that did not experienced type I migra-tion rates during the gaseous phase show that the efficiencyin the formation of potentially habitable planets (PHPs)seems to strongly depend on the size of the planetesimalsfrom which these systems were formed Planetary systemsformed from small planetesimals are the most efficient onesMoreover this kind of planetary systems formed from smallplanetesimals of 100 m is the only one capable of formingwater-rich PHPs Bigger planetesimals of 1 km and 10 kmalso form PHPs but completely dry and planetary systemsformed from planetesimals of 100 km do not form PHPs atall

We also test the sensitivity of the results to scenariosthat experienced type I migration rates during the gaseousphase and that were also able to form giant planets mas-sive enough to open a gap in the gas disc and to turn on

type II migration Our results show that planetary systemswhich experienced very low type I migration rates do notpresent significant differences to the results of those SSAsthat did not and that planetary systems which experiencedtype I migration rated reduced to a 10 seems to favor theformation of PHPs in scenarios formed from small planetesi-mals but is detrimental for the formation of PHPs in scenar-ios formed from big planetesimals Those planetary systemswhich were able to form giant planets that opened gaps inthe gas discs during the gaseous phase seem to act in thesame way as type I migration does Analyzing all the de-veloped simulations it also seems that there might be anoptimal mass and location for the inner giant planet of thesystem regarding the efficiency of forming PHPs Althoughwe estimate those values around sim 6 au and between 1MJ

and 2MJ taking into account all the developed simulationsin this work a most detailed set of numerical simulationsconsidering a the semimajor-axis distribution of the innergas giants between 15 au and 10 au should be developedin order to properly determine these values It is howeverimportant to remark that we count with all the informationto do this work in a near future

Finally we make a comparison between our simulatedplanetary systems and the current potentially habitableplanet population and gas giant population Although thiskind of comparisons must be analyzed carefully since todate no planetary systems similar to our own have beenyet discovered this is none of the observed (confirmed andcandidates) potentially habitable planets present a gas giantcompanion beyond 15 au we are able to form PHPs in a re-gion of the mass vs semimajor-axis diagram that is currentlybeing explored Moreover the simulated giant planet popu-lation that was capable of forming PHPs occupies a regionof the semimajor-axis vs eccentricity diagram of extendedsemimajor-axis and very low eccentricities and to date apopulation with these characteristics has not been yet dis-covered We believe that the discovery of exoplanets withthese properties is a challenge of future missions Particu-larly Hippke amp Angerhausen (2015) explored the potentialof future missions such as PLATO 20 in detecting Solar sys-tem analogues and showed that the discovery of Venus andEarth-like planets transiting stars like our sun is feasible athigh signal-to-noise ratio after collecting 6 yrs of data Theyalso showed that due to the high number of observed starsby PLATO the detection of single-transit events by cold gas



giants like Jupiter Saturn Uranus and Neptune analogueswill be possible These future detections particularly thoseof cold gas giants will allow us to associate this populationwith the potential existence of planets in the HZ

ACKNOWLEDGEMENTS

The authors thank the anonymous referee for her or his valu-able comments which helped to improve the manuscriptMPR thanks Octavio Miguel Guilera for useful discussionsThis work has been supported by grants from the NationalScientific and Technical Research Council through the PIP043613 and National University of La Plata Argentinathrough the PIDT 11G144

REFERENCES

Abbot D S Cowan N B Ciesla F J 2012 ApJ 756 178Abe Y Ohtani E Okuchi T Righter K Drake M 2000 Water

in the Early Earth pp 413ndash433

Alibert Y 2014 AampA 561 A41Alibert Y Mordasini C Benz W Winisdoerffer C 2005 AampA

434 343Andrews S M Wilner D J Hughes A M Qi C Dullemond

C P 2010 ApJ 723 1241

Barge P et al 2008 AampA 482 L17Barnes J W 2007 PASP 119 986Batalha N M 2014 Proceedings of the National Academy of Science

111 12647

Bitsch B Johansen A Lambrechts M Morbidelli A 2015AampA 575 A28

Bolmont E Libert A-S Leconte J Selsis F 2016 AampA591 A106

Borucki W J et al 2010 Science 327 977Borucki W J et al 2012 ApJ 745 120Broeg C et al 2013 in European Physical Journal

Web of Conferences p 03005 (arXiv13052270)doi101051epjconf20134703005

Chambers J E 1999 MNRAS 304 793Chambers J E 2001 Icarus 152 205Chambers J 2008 Icarus 198 256

Chambers J E 2013 Icarus 224 43Chambers J E Wetherill G W Boss A P 1996 Icarus

119 261Crida A Morbidelli A Masset F 2006 Icarus 181 587

DrsquoAngelo G Marzari F 2012 ApJ 757 50Dauphas N Pourmand A 2011 Nature 473 489

Dugaro A de Elıa G C Brunini A Guilera O M 2016 AampA596 A54

Dvorak R Maindl T I Burger C Schafer C Speith R 2015Nonlinear Phenomena in Complex Systems Vol18 No3pp 310-325 18 310

Fischer D A Howard A W Laughlin G P Macin-tosh B Mahadevan S Sahlmann J Yee J C 2014Protostars and Planets VI pp 715ndash737

Fogg M J Nelson R P 2005 AampA 441 791Fogg M J Nelson R P 2006

International Journal of Astrobiology 5 199Fogg M J Nelson R P 2007 AampA 461 1195

Fogg M J Nelson R P 2009 AampA 498 575Gladman B 1993 Icarus 106 247Guilera O M Sandor Z 2017 AampA 604 A10

Guilera O M Brunini A Benvenuto O G 2010 AampA521 A50

Guilera O M Fortier A Brunini A Benvenuto O G 2011

AampA 532 A142

Guilera O M de Elıa G C Brunini A Santamarıa P J 2014AampA 565 A96

Hayashi C 1981 Progress of Theoretical Physics Supplement70 35

Hellary P Nelson R P 2012 MNRAS 419 2737

Higuchi A Kokubo E Kinoshita H Mukai T 2007 AJ134 1693

Hippke M Angerhausen D 2015 ApJ 810 29

Ida S Lin D N C 2004 ApJ 604 388

Inaba S Tanaka H Nakazawa K Wetherill G W Kokubo E2001 Icarus 149 235

Inamdar N K Schlichting H E 2015 MNRAS 448 1751

Jacobson S A Morbidelli A Raymond S N OrsquoBrien D PWalsh K J Rubie D C 2014 Nature 508 84

Jenkins J M et al 2015 AJ 150 56

Kasting J F Whitmire D P Reynolds R T 1993 Icarus101 108

Kitzmann D et al 2015 ArXiv e-prints 150701727

Kopparapu R K et al 2013a ApJ 765 131

Kopparapu R K et al 2013b ApJ 770 82

Kopparapu R K Ramirez R M SchottelKotte J Kasting J FDomagal-Goldman S Eymet V 2014 ApJ 787 L29

Laakso T Rantala J Kaasalainen M 2006 AampA 456 373

Lecuyer C Gillet Pand Robert F 1998 Chem Geol145 249aAS261

Leinhardt Z M Stewart S T 2012 ApJ 745 79

Lodders K Palme H Gail H-P 2009 Landolt Bornstein

Mandell A M Raymond S N Sigurdsson S 2007 ApJ660 823

Marty B 2012 Earth and Planetary Science Letters 313 56

Marzari F 2014 MNRAS 442 1110

Masset F Snellgrove M 2001 MNRAS 320 L55

Miguel Y Guilera O M Brunini A 2011 MNRAS 417 314

Morard G Bouchet J Valencia D Mazevet S Guyot F 2011High Energy Density Physics 7 141

Morbidelli A Crida A 2007 Icarus 191 158

Mordasini C Alibert Y Benz W Naef D 2009 AampA501 1161

Noack L et al 2016 Icarus 277 215

OrsquoBrien D P Morbidelli A Levison H F 2006 Icarus 184 39

Ohtsuki K Stewart G R Ida S 2002 Icarus 155 436

Paardekooper S-J Baruteau C Crida A Kley W 2010MNRAS 401 1950

Paardekooper S-J Baruteau C Kley W 2011 MNRAS410 293

Pfalzner S Steinhausen M Menten K 2014 ApJ 793 L34

Pringle J E 1981 ARAampA 19 137

Quintana E V Barclay T Borucki W J Rowe J F ChambersJ E 2016 ApJ 821 126

Rafikov R R 2004 AJ 128 1348

Raymond S N Izidoro A 2017 Icarus 297 134

Raymond S N Quinn T Lunine J I 2004 Icarus 168 1


Raymond S N OrsquoBrien D P Morbidelli A Kaib N A 2009Icarus 203 644

Raymond S N Schlichting H E Hersant F Selsis F 2013Icarus 226 671

Ricker G R et al 2010 in American Astronomical SocietyMeeting Abstracts 215 p 459

Ronco M P de Elıa G C 2014 AampA 567 A54

Ronco M P de Elıa G C Guilera O M 2015 AampA 584 A47

Ronco M P Guilera O M de Elıa G C 2017 MNRAS471 2753

Salvador Zain P de Elıa G C Ronco M P Guilera O M2017 preprint (arXiv171004617)


24 M P Ronco et al

Schneider J Dedieu C Le Sidaner P Savalle R Zolotukhin I

2011 AampA 532 A79Selsis F Kasting J F Levrard B Paillet J Ribas I Delfosse

X 2007 AampA 476 1373Simpson F 2017 MNRAS 468 2803Suzuki T K Muto T Inutsuka S-i 2010 ApJ 718 1289Suzuki T K Ogihara M Morbidelli A Crida A Guillot T

2016 AampA 596 A74Tanaka H Takeuchi T Ward W R 2002 ApJ 565 1257Tanigawa T Ikoma M 2007 ApJ 667 557Touboul M Kleine T Bourdon B Palme H Wieler R 2007

Nature 450 1206Tsiganis K Gomes R Morbidelli A Levison H F 2005

Nature 435 459Valencia D OrsquoConnell R J Sasselov D D 2007 ApJ 670 L45Veras D Evans N W 2013 MNRAS 430 403Walsh K J Morbidelli A Raymond S N OrsquoBrien D P Man-

dell A M 2011 Nature 475 206Williams D M Pollard D 2002

International Journal of Astrobiology 1 61Wolfgang A Rogers L A Ford E B 2016 ApJ 825 19de Elıa G C Guilera O M Brunini A 2013 AampA 557 A42

This paper has been typeset from a TEXLATEX file prepared bythe author


1 Introduction

2 Summary of the gas-phase planet formation model

3 Initial conditions

4 Properties of the N-body simulations

5 Results




54 Results sensitivity to scenarios with type I migration

55 Results sensitivity to scenarios with gap-opening giants

56 Possible range of mass and location of the inner giant to form PHPs

6 Contrast with the observed planet population

61 PHPs population


7 Discussion

8 Conclusions

2 M P Ronco et al

like stars Later Mandell et al (2007) explored the forma-tion of terrestrial-type planets in planetary systems with onemigrating gas giant around solar-type stars They also con-sidered planetary systems with one inner migrating gas giantand one outer non-migrating giant planet Fogg amp Nelson(2009) examined the effect of giant planet migration onthe formation of inner terrestrial planet systems and foundthat this phenomena does not prevent terrestrial planet for-mation Salvador Zain et al (2017) focused on the forma-tion and water delivery within the habitable zone in dif-ferent dynamical environments and also around solar-typestars These authors considered planetary systems that har-bor Jupiter or Saturn analogues around the location of thesnowline and found that the formation of water-rich planetswithin the habitable zones of those systems seems to be acommon process More recently Raymond amp Izidoro (2017)showed that the inner Solar systemrsquos water could be a re-sult of the formation of the giant planets These authors findthat when the gas disc starts to dissipate a population ofscattered eccentric planetesimals cross the terrestrial planetregion delivering water to the growing Earth

During our first work (Ronco et al 2017 hereafter PI)we improved our model of planet formation which calculatesthe evolution of a planetary system during the gaseous phaseto find via a population synthesis analysis suitable scenar-ios and physical parameters of the disc to form Solar systemanalogues (SSAs from now on) We were specially interestedin the planet distributions and in the surface density ec-centricity and inclination profiles for the planetesimal popu-lation at the end of the gas-phase obtained considering dif-ferent formation scenarios with different planetesimal sizesand different type I migration rates

The aim of this second work is to analyze the formationof SSAs during the post-gas phase through the developmentof N-body simulations focusing on the formation of rockyplanets and their water contents in the inner regions of thedisc particularly on the habitable zone To do so and inorder to analyze the formation of SSAs in a self-consistentlyway we make use of the results obtained in PI as the initialconditions for our N-body simulations and form final plane-tary systems from different planetesimal sizes Many worksin the literature that analyze the post-gas phase formation ofplanetary systems in particular the formation of rocky plan-ets in the inner regions of the disc considered arbitrary ini-tial condition (Chambers 2001 Raymond et al 2004 2006OrsquoBrien et al 2006 Mandell et al 2007 Raymond et al2009 Walsh et al 2011 Ronco amp de Elıa 2014) Howeveras it was already shown by Ronco et al (2015) more re-alistic initial conditions obtained from a model of planetformation lead to different accretion histories of the planetsthat remain within the habitable zone than considering ad

hoc initial conditions We also analyze the differences in theresults of planetary systems that were affected by the type Imigration regime (Tanaka et al 2002) during the gas-phaseand by planetary systems with gas giant planets that wereable to open a gap in the disc and were affected by type IImigration

This work is then organized as follow In Section 2 wesummarize our model of planet formation explained in de-tailed in PI in Section 3 we describe the initial conditionsfor the development of our N-body simulations obtainedfrom the results of PI in Section 4 we describe the most

important properties of the N-body code and then show ourmain results in Section 5 Then we compare our results withthe observed potentially habitable planets and with the cur-rent gas giant population in Section 6 Finally discussionsand conclusions are presented in Section 7 and Section 8respectively

2 SUMMARY OF THE GAS-PHASE PLANET

FORMATION MODEL

In PI we improved a model of planet formation based on pre-vious works (Guilera et al 2010 2011 2014) which we nowcalled PLANETALP This code describes the formation andevolution in time of a planetary system during the gaseousphase taking into account many relevant phenomena forthe planetary system formation The model presents a pro-toplanetary disc which is characterized by a gaseous compo-nent evolving due to an α-viscosity driven accretion (Pringle1981) and photoevaporation (DrsquoAngelo amp Marzari 2012)processes and a solid component represented by a plan-etesimal and an embryo population The planetesimal pop-ulation is being subject to accretion ejection (Laakso et al2006) and scattering (Ida amp Lin 2004 Alibert et al 2005)by the embryos and radial drift due to gas drag in-cluding the Epstein Stokes and quadratic regimes Theembryo population grows by accretion of planetesimals(Inaba et al 2001) gas (Ida amp Lin 2004 Miguel et al 2011Mordasini et al 2009 Tanigawa amp Ikoma 2007) and due tothe fusion between them taking into account their atmo-spheres (Inamdar amp Schlichting 2015) while they migratevia type I (Tanaka et al 2002) or type II (Crida et al 2006)migration mechanisms We also consider that the evolu-tion of the eccentricities and inclinations of the planetesi-mals are due to two principal processes the embryo gravi-tational excitation (Ohtsuki et al 2002) and the dampingdue to the nebular gas drag (Rafikov 2004 Chambers 2008)The particularities and the detailed description of the im-plementation of each phenomena can be found in PI and inGuilera et al (2014)

In PI we also described the method initial parametersand initial formation scenarios considered to synthesize apopulation of different planetary systems around a solar-mass star with the particular aim of forming SSAs Follow-ing Andrews et al (2010) the gas and solid surface densityprofiles are modeled as

Σg(R) = Σ0g

(R

Rc

)minusγeminus(

R

Rc

)2minusγ

(1)

Σp(R) = Σ0pηice

(R

Rc

)minusγeminus(

R

Rc

)2minusγ

(2)

where Σ0p and z0Σ

0g are normalization constants with z0 =

00153 the primordial abundance of heavy elements in thesun (Lodders et al 2009) and ηice is a parameter that rep-resents the increase in the amount of solid material due tothe condensation of water beyond the snowline given by 1if R ge Rice and given by 1β if R lt Rice where Rice is theposition of the snowline at 27 au (Hayashi 1981)

The variable parameters which were taken randomlyfrom range values obtained from the same observationalworks (Andrews et al 2010) are the mass of the disc Md




















fH2O(R) =

0 if R le Rice

β minus 1

β if R gt Rice

(3)


4 M P Ronco et al







Slowast10


10minus5

10minus4

10minus3

10minus2

10minus1

100



S1

Mas

s [M

J]

0 01 02 03 04 05 06



S2

0 01 02 03 04 05 06

100m minus fmigI=01


S3

0 01 02 03 04 05 06

1km minus fmigI=0


S4

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1



S5

Σ p [g

cm

minus2]

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1

100

101

1 10

10km minus fmigI=0


S6

Mas

s [M

J]


0 01 02 03 04 05 06

1 10



S7


0 01 02 03 04 05 06

1 10

10km minus fmigI=01


S8


0 01 02 03 04 05 06

1 10

100km minus fmigI=0


S9


0 01 02 03 04 05 06

1 1010minus5

10minus4

10minus3

10minus2

10minus1


S10

Σ p [g

cm

minus2]


0 01 02 03 04 05 06









SIMULATIONS













and 360 for




6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]




Zoom




5 RESULTS










Flowasteff=

Feff

(1 minus e2)12 (4)











minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions



















fH2O(R) =

0 if R le Rice

β minus 1

β if R gt Rice

(3)


4 M P Ronco et al







Slowast10


10minus5

10minus4

10minus3

10minus2

10minus1

100



S1

Mas

s [M

J]

0 01 02 03 04 05 06



S2

0 01 02 03 04 05 06

100m minus fmigI=01


S3

0 01 02 03 04 05 06

1km minus fmigI=0


S4

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1



S5

Σ p [g

cm

minus2]

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1

100

101

1 10

10km minus fmigI=0


S6

Mas

s [M

J]


0 01 02 03 04 05 06

1 10



S7


0 01 02 03 04 05 06

1 10

10km minus fmigI=01


S8


0 01 02 03 04 05 06

1 10

100km minus fmigI=0


S9


0 01 02 03 04 05 06

1 1010minus5

10minus4

10minus3

10minus2

10minus1


S10

Σ p [g

cm

minus2]


0 01 02 03 04 05 06









SIMULATIONS













and 360 for




6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]




Zoom




5 RESULTS










Flowasteff=

Feff

(1 minus e2)12 (4)











minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

4 M P Ronco et al







Slowast10


10minus5

10minus4

10minus3

10minus2

10minus1

100



S1

Mas

s [M

J]

0 01 02 03 04 05 06



S2

0 01 02 03 04 05 06

100m minus fmigI=01


S3

0 01 02 03 04 05 06

1km minus fmigI=0


S4

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1



S5

Σ p [g

cm

minus2]

0 01 02 03 04 05 06

10minus5

10minus4

10minus3

10minus2

10minus1

100

101

1 10

10km minus fmigI=0


S6

Mas

s [M

J]


0 01 02 03 04 05 06

1 10



S7


0 01 02 03 04 05 06

1 10

10km minus fmigI=01


S8


0 01 02 03 04 05 06

1 10

100km minus fmigI=0


S9


0 01 02 03 04 05 06

1 1010minus5

10minus4

10minus3

10minus2

10minus1


S10

Σ p [g

cm

minus2]


0 01 02 03 04 05 06









SIMULATIONS













and 360 for




6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]




Zoom




5 RESULTS










Flowasteff=

Feff

(1 minus e2)12 (4)











minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions



SIMULATIONS













and 360 for




6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]




Zoom




5 RESULTS










Flowasteff=

Feff

(1 minus e2)12 (4)











minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

6 M P Ronco et al

1

1 10 100 01

1

10

S8

Ecce

ntric

ity

Σ p [g

cm

minus2]




Zoom




5 RESULTS










Flowasteff=

Feff

(1 minus e2)12 (4)











minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions










minus4Moplus











9in





8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

8 M P Ronco et al




















SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions














SIM1 185 130 1080 0 A










10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

10 M P Ronco et al
















lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions





lowast 3910 ASIM4 343 114 1732



lowast 3891 A














12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

12 M P Ronco et al
















14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions









14 M P Ronco et al









migration


0

5

10

15

20

100m 1km 10km 100km

Mas

s [M

oplus]

Planetesimal size

0

10

20

30

40

50

60

o

f Wat

er b

y m

ass

0

2

4

6

8

10

100m 1km 10km 100km

Num

ber o

f PH

P pe

r sce

nario

Planetesimal size





9



10










gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions





gap-opening giants














16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

16 M P Ronco et al

100m - fmigI=0 - S1 - (SIM10)

100m - fmigI=01 - S3 - (SIM6)

10km - fmigI=0 - S6 - (SIM1)

10km - fmigI=01 - S8 - (SIM5)


Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater







Initial Conditions



0

01

02

03

04

05

06

Frac

tion

of w

ater














PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions








PLANET POPULATION


0 005

01 015

02 025

03 035

04

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

005

01

015

02

025

03

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 01 02 03 04 05 06 07 08 09

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0

01

02

03

04

05

06

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]

0 005

01 015

02 025

03 035

04 045

05

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Semimajor-axis [au]

0 005 01

015 02

025 03

035

0 1 2 3 4 5 6 7 8 9 10 11




Frac

tion

of P

HP

Mass [MJ]





61 PHPs population



18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

18 M P Ronco et al









+023



+0054

minus0047M⊙ at

1046+0019




its HZ at 0849+0018


+013

minus013Roplus and







0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions


0

5

10

15

20

25

30

04 06 08 1 12 14 16 18 2

Mas

s [M

J]



Confirmed

0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b

05

1

15

2

25

3

04 06 08 1 12 14 16 18 2

Rad

ius

[Roplus

]


0

02

04

06

08

1

Ecce

ntric

ity

Keplerminus452 b

Keplerminus22 b












20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

20 M P Ronco et al


7 DISCUSSION













0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions


0001

001

01

1



Ecce

ntric

ity

0001

001

01

1

1 10 100



Ecce

ntric

ity


01

1

10

100

Mas

s [M

J]

01

1

10

100



Mas

s [M

J]

01

1

10

0001 001 01 1



ass

[MJ]

Eccentricity

1

10

100

Sem

imaj

orminusa

xis

[au]

01

1

10

100



Mas

s [M

J]

01

1

10

1 10 100



ass

[MJ]


001

01

1

Ecce

ntric

ity







22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

22 M P Ronco et al




8 CONCLUSIONS










ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions



ACKNOWLEDGEMENTS


REFERENCES





C P 2010 ApJ 723 1241


111 12647

















AampA 532 A142




































Rafikov R R 2004 AJ 128 1348












24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

24 M P Ronco et al











1 Introduction




5 Results








61 PHPs population


7 Discussion

8 Conclusions

formationofsolarsystemanaloguesii: discsvian-body simulations

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