Star Formation News Letter No. 296 #11-16
楠根貴成（国立天文台）#11 Search for grain growth towards the center of
L1544#12 Hydrodynamics of embedded planetsʼ ﬁrst

atmospheres - III. The role of radiation transport forsuper-Earth planets#13 Molecular clumps photoevaporation in ionizedregions#14 How fast in mass-segregation happening inhierarchical formed embedded star clusters?

#15 The evolution of Giant Molecular Filaments#11 Search for grain growth towards the center of L1544 ?
A. Chacón-Tanarro1 , P. Caselli1 , L. Bizzocchi1 , J. E. Pineda1 , J. Harju1, 2 , M. Spaans3 , and F.-X. Désert4
Department of Physics,
P.O. Box 64, 00014
University of Helsinki,
天体L1544(d~140pc)Finland
Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands

1
2

Max-Planck-Institüt für extraterrestrische Physik, Giessenbachstrasse 1, 85748 Garching, Germany

A. Chacón-Tanarro et al.: Search for grain growth towards the center of L15443

観測1.2 & 2 mm (IRAMに搭載されたNIKAを使用)
10 7

use (MJy/sr), Eq. (1)4 for each pixel transforms

10

Univ. Grenoble Alpes, CNRS, IPAG, Grenoble F-37000, France
9

Received - / Accepted (6)Main beam width8

T d (K)12.5ʼʼ at 1.2mm、18.5ʼʼ at 2mm(1800au)(2600au)chemical structure of dense cloud cores, including the dynamically important ionization
fraction, and the future evolution of solids
in protoplanetary disks. To study this further, we focus on L1544, one of the most centrally concentrated pre-stellar cores on theth along the line of sight, which has a direct relaius r, and i and j represent the pixel coordinates

arXiv:1707.00005v1 [astro-ph.GA] 30 Jun 201710 6

(cm-3)slos

B⌫ [T d (ri j )]nH2 (ri j )ds,

ABSTRACT2ZnHµH2 mH7

10 5
previous section, the ratio between the surface
Inandensecold.2 and 2 mm gives us
estimate andof , andcon- molecular clouds dust grains are surrounded by thick icy mantles. It is however not clear if dust growth and6
lue of ⌫ , although coagulationin this case we willtake
deriveplaceitbefore the switch-on of a protostar. Thisis an important issue, as the presence of large grains may a↵ect theRS ⌫ (ri j )

.

(7)

10 4

5
10000

A. Chacón-Tanarro et al.: Search for grain growth towards the center of L1544

Grainのopacity変化と成長を調べる

B⌫ [T d (ri j )]nH2 (ri j ) ds

0

2000

4000

6000

8000

verge of star formation, and with a well-known physical structure. We observed L1544 at 1.2 and 2 mm using NIKA, a new receiverFig. 3. Temperature (red) and density (blue) profiles of L1544 fromat the IRAM 30 m telescope,
and we used data from the Herschel Space Observatory archive. We find no evidence of grain growth
Keto et al. (2015).
h is more accurate thanthe method
in Sec-of L1544 at the available angular resolution. Therefore, we conclude that single dish observations do not allow us totowards
theusedcenterse it does not assume a constant value for thegrain
growth toward the pre-stellar core L1544 and high sensitivity interferometer observations are needed. We predict that
rence frequency, butinvestigate
it also has the clear
wrong
the spectral index andthe opacity
not change
dust
grainsdocan
grow to 200 µm in size toward the central ⇠300 au of L1544. This will imply a dust opacity change by a factor of ⇠2.5
f sight. This may not be the case, as if changes
atcore1.2aremm,
which
can be detected using the Atacama Large Millimeter and submillimeter Array (ALMA) at di↵erent wavelengths andts to the center of thepresent,
the same
pply along the linewithof sight
the center.
antoward
angularresolution of 200 .
slos

R (au)

lso assumes that the spectral index does not deperature, but as weKey
have seenin Section
1, no clouds - ISM: individual objects: L1544 - ISM: dust, extinction - opacity - stars: formation
words.
ISM:
erature dependence is expected within the temlevant for L1544 (Agladze et al. 1996; Mennella
det et al. 2005; Coupeaud et al. 2011; Demyk.
1.2mm maps are shown in Fig. 4. The behaviour
ndex seems not to di↵er much from the one pren 3.2.1, where a constant temperature was asr, here the spatial variations are not significant,igible within the errors(⇠25%). This is due toRratio of the term s B⌫ [T d (ri j )]nH2 (r)ds at thelosquencies does not di↵er significantly from thessuming a constant temperature of 10 K. On theopacity value changes significantly, ⇠60% with
20%. Thus, the overall 1.2mm value dispersion isthan before, and surprisingly, the opacity is now
asing trend towards the central regions, contraryn of grain growth.
explanation of this apparently contradictory reodel assumes spherical symmetry, while L1544
Fig. 9. Variation of the opacity with time and projected radius deduced Fig. 10. Variation of the opacity with time and projected radius deduced
ngated core, with an aspect ratio of about 2. Thismodel cannotHbe used for 5the outer3parts of thewith ALMA (on the left) and NIKA (on the right) simulated observa- with ALMA (on the left) and NIKA (on the right) simulated observa2
lar along the major axis. Moreover, the deriva- Fig. 4. Maps of the (a) spectral index and (b) the opacity at 1.2 mm
tions for the case of a static cloud. The opacity for the NIKA simulated tions for the case of a dynamical cloud. The change on the opacity betral index using the method of the fluxes ratio derived with the method described in Section 3.2.2. The angular resoluchangeda factor of ⇠1.5, while for ALMA the change tween the center and the outer edge for the case of NIKA is less thanavelengths in the Rayleigh-Jeans limit leads to tion is 1900 (HPBW shown by observations
the white circle in thebottom rightby
corner).
s due to the uncertainties on the fluxes (Shetty The crosses mark the peak ofcorresponds
the emission corresponding
to the final
1.2
to a factorof ⇠14.
15%, while for ALMA it is a factor of ⇠2.5 after 1.06 Myrs of evolution.

1. Introduction

been questioned by Jones et al. (2013), who suggested that amorphous hydrocarbon material could actually produce the observed
coreshine without the need of large grains. Moreover, Andersen
et al. (2014) found a similar threshold for the coreshine and water ice, with the scattering efficiency at 3.6 µm increasing with
the increase of the water-ice abundance, suggesting that water
ice mantle growth in dense clouds may be at least partially responsible for the coreshine e↵ect.

Pre-stellar cores are self-gravitating starless dense cores withclear signs of contraction motions and chemical evolution(Crapsi et al. 2005). They are formed within molecular clouds,
due to the influence of gravity, magnetic fields and turbulence. They are considered to be on the verge of star formation, and therefore represent the initial conditions in the process of star formation (Bergin & Tafalla 2007; Caselli & CecThe study of the dust emission at long wavelengths, sensicarelli 2012). These systems are characterized by high densities(n > 10 cm ) and low temperatures (T < 10 K) towards their tive to the larger grains, is therefore important to gain a better
understanding on the grain size distribution on the early phases
central regions (Crapsi et al. 2007).
In dense and cold molecular clouds dust grains are sur- of star formation. This emission depends on dust grains properrounded by thick icy mantles(e.g. Boogert
et al. 2015, and ref- ties such as structure, size and composition. Here we will focus
mm map convolved
to 19 .
erences therein). It is however not clear if dust growth and coag- on the dust opacity and the spectral index, two parameters thatdex and opacity of the dust using NIKA andof the dust. is
Theborn.
adoptedvaluesforthe variation
opacityand thedependresolution
directly on
the observablegrain size distribution.
fluxes. This means that no signatures of grain growth are present
ulation take place before index
a protostar
animportant(Fig.This10),istheat NIKA’s
is notspectral index are derived from the best least squares fit towiththespectral
sensitivity
reached, as it is below 15% and the RMS of towards L1544, within the limits of our data.
issue, dependence
as dust coagulation
may
a↵ect
the formation
andevolutionplex and unknownof the spectral thesurfacebrightnesses
inthe five
windows
available
The opacity, ⌫ , is a measurement ofFinally,
the dusta absorption
cross model of grain growth applied
temperature,
uncertainties in the disks,surface (3from Herschel/SPIRE
and 2datafrom corresponds
NIKA)
toward (Zhaothetopeak
simple analyticalour⇠15% of the peak emission.
This result isof theprotoplanetary
formingfrom molecular
cloudsd the spherically symmetric model, here we emission of L1544 at each band, assuming that the opacity andsections
weighted
by
the
mass
of
the
gas
and
dust.
As
explained
consistent
with ourobservations,which show that no change on to L1544 shows that NIKA observations cannot detect opacity
erent approach.
modeled Foster
densities andthe (2013)
spectral index
are aindependent,
possibledegeneracies
et al.The2016).
et al.
findstrong ascorrelation
be- by Kruegel&NIKA
Siebenmorgen(1994),
when dust
areandsmall
files are now used to predict the emission at between both parametersthe
are opacity
beyond the is
scopeof thispaperchanges
acrossgrains
the corethat only interferometers, in particseenorneededtoreproducedata. However,
tween
the
visualandthe slope
thetheextinction
fixing the valuesof theopacityand theextinction
spectral and not
well-known.
Beforeof
fitting,
five maps havelaw
been
comparedtothewavelength,theso-called
Rayleigh
limit,
theular ALMA, can provide information
about grain growth toward
Fig.
10
shows
that
the
dust
opacitychangeshouldbedetected
toward the Perseus molecular cloud. This can beinterpreted as mass absorption coefficient does not depend on the grain size,
Article number, page 5 of 12
the central regions of L1544.
withhigh-resolution
telescopes,
like
ALMA,
as
an
opacityingrain growth, but it is not clear if grain coagulation is needed

Grain成長の証拠は見つからず

中心から300au以内 200μmに成長していると予測
opacityは2.5倍ほどになり、
00

この変化の検出にはALMAが必要

#12 Hydrodynamics of embedded planets’ first atmospheres -III. The role of radiation transport for super-Earth planets
Nicolas P. Cimerman,1? Rolf Kuiper,1 Chris W. Ormel2

1 Institut für Astronomie & Astrophysik, Universität Tübingen, Auf der Morgenstelle 10, D-72076 Tübingen, Germany
2 Anton Pannekoek Institute, University of Amsterdam, PO Box 94249, NL-1090 GE Amsterdam, the Netherlands

低質量惑星@~0.1au

Hydrodynamics of embedded planets’ first atmospheres - III

4

N. P. Cimerman et al.

ABSTRACTyTable 1. List of simulations. The first column gives the simulation name. The second column shows the core mass Mc . The third colum
gives the injection time tinj for the planet’s gravity and the fourth column shows the total simulation time tend . The spatial resolutio
for the domain [rc, 5H] ⇥ [0, ⇡/2] ⇥ [0, ⇡] is given in the fifth column. Comments and references to figures are given in the sixth column
Standard resolution simulations were conducted for three di↵erent opacities.

The population of close-in super-Earths, with gas mass fractions of up to 10% represents a challenge for planet formation theory: how did they avoid runawayName
gas accretion
Mc [M ]
tinj [⌦ 1 ]tend [⌦ 1 ]Nr ⇥ N ✓ ⇥ N
Comments
circularorbit
and collapsing
to hot Jupiters despite their core masses being in the critical range of
RT-M1-loRs
14
1000
64 ⇥ 16 ⇥ 32
resolution study, Appendix A
Mc ' 10 M ? Previous three-dimensional (3D) hydrodynamical simulations
indicate
RT-M0.1
0.1
4500
128 ⇥ 32 ⇥ 64
lowest mass casethat atmospheres of low-mass planets cannotRH/B
be considered isolated RT-M1
from the pro-1
4
1000
128 ⇥ 32 ⇥ 64
Figs. 2, 4, 6, 7, 8, 9
starplanettoplanetary
disc, contrary to what is assumed in 1D-evolutionary calculations.
This2
RT-M2
5
1000
128 ⇥ 32 ⇥ 64
Figs. 3, 4
x
RT-M5
5
1000
128 ⇥ 32 ⇥ 64
highest mass casefinding is referred to as the recycling hypothesis. In this Paper
we investigate
the5
4
290
256 ⇥ 64 ⇥ 128
resolution study, Appendix A
recycling hypothesis for super-Earth planets, accounting for realistic RT-M1-hiRs
3D radiation-12
RT-M2-hiRs
5300
256 ⇥ 64 ⇥ 128
resolution studyahydrodynamics. Also, we conduct a direct comparison in terms of the evolution of theentropy between 1D and 3D geometries. We clearly see that 3D atmospheres mainTable
2. List ofparametersradiativefor di↵erent core masses Mc in units of Earth masses. We also give the core radius, Bondi and Hill radius itain higher entropy: although gas in the atmosphere loses
entropythrough
fifth column gives
the dimensionless thermal mass m = RB /H of the planet, which depends on orbital location and th
cooling, the advection of high entropy gas from the discEarthintoradii.
theTheBondi/Hill
spherelocal disc properties. In the last two columns, we give the radius at which we find the transition between transient and close to circula
potentiallyarrestinggrowthglobal frameet al. slows down Kelvin-Helmholtz contraction,b) local frame
around planet
streamlinesenvelopein the mid-plane
and atthe aenvelope mass fraction env ⌘ Menv /Mc , both after t = 1000 ⌦ 1 . Values of env range from resul
8 N. P.a)Cimerman
sub-critical gas mass fraction. Recycling, therefore, operates
vigorously,
withfor standard
opacity toin
VLOlineopacity.
Figure 1. Schematic drawing of local shearing
sheetbyapproximation.
a) The planet
is fixed onwea circularorbit with an
the “innerseperationcore”
aresults
previous studies.
However,also identify– in size ⇡ 25% offrom the star. b) We align our co-moving frame that is centred on the planet such that the x-direction is pointing radially away from
Mlowerrc [R ]
RB [R ]
RH [R ]
m
rcirc /RBc [M ]
env
the direction
Bondiofradius– where streamlines are more circular and3 entropies are much
1.0 the y-direction follows the orbital
the star andthe planet. Arrows indicate the gas velocity of the shearing flow. 10
Without athanthe
Futurestudies
0.1
0.48
1.77
1.09
0.04
0,018 – 0.024 %“headwind” arising from a pressure gradientin thein
disc,thisouterflow is atmosphere.
symmetric with respect
to theplanet. at higher resolutions are needed to assess
1.03
17.7
23.4
0.38
0.25
1.5 – 1.6 %
0.5
whether this region can become hydrodynamically-isolated on long time-scales. 1
2

1.30

35.4

29.5

⇢ [g/cm3 ]y [RHill ]

5
1.77
88.5
40.0
due to the Coriolis force, a cor = 2 Key⇥ v 3 andthe cenhas already reached
95% ofit’s satellites:
limit at unity.atmospheres
We note that - planets and satelwords:
hydrodynamics- planets
and0.0
trifugal force a cent =⇥ ( ⇥ r ). In a linear approximathe chosen injection times are usually a few ⌦ 1 and arelites: formation - protoplanetary
discs.
tion (for x, y, z ⌧ a), the latter can be put together withthus much shorter than the actual time for the coagulation
0.5 of stellar gravity to give the tidal acceleracontributionsof the core. While for massive envelopes self-gravityof theapproximation
it is given by (see also Fig. 1b)
2
2tion a tid = 3x ⌦ e x z ⌦ e z . The vertical component of a tid
gas is the important e↵ect potentially leading to runaway
4 mass 3
gives rise1.0
to a density stratification of the disc in z and wasgas accretion, we can neglect it while the atmosphere’s
v10x ⌦ ey.
(13)
11
1
1 (x) =
t = 300.0
t = 500.0
2
is small comparedto the core mass, Menv ⌧ Mc .
not consideredt =in50.0
the low-mass limit (Paper I; Paper
II), be1causeINTRODUCTION
source of the luminosity (e.g. Ikoma et al. 2001; Lee & Chithere RBondi /H ⌧ 1. Super-Earths, however, are moreIn the frame where the planet is at rest, the co-rotation
ang 2015).
massive planets,
so that this additional term cannot be ne1.0
orbits at x = 0 are so as well. Gas that orbits closer to theGiantglected.planets form in gas rich discs either by gravitational
2.3 Numerical methods
star isis typical
overtaking
due to faster angular motion,
The ⇠10 M for the CCM
forthetheplanet
formation
instability(Kuiper
1951;
Cameron
1978)
or,
as
is
assumedWe only
0.5 consider a Keplerian disc, so we do not takewhile
gas further
lags behind.
Jupiter
assuming
ISM-like
graintheopacities(inoutregionscool
We solve the ofabove
describedhydrodynamic
equationsusinto accountthe globalpressure
force, which
results inamore
commonly,
by coreaccretion(Safronov1969;
PollackWe
follow Leea etsteadyal. (2014)
in adopting the minimumenoughforsilicate
grainstobepresent),
influxing
amodifiedversionof
the
Godunov-typehydrodynamisub-Keplerianmotion
of
the
gas,
seen
as
a
“headwind”
inet al. 1994).
In the latter scenario the final rapid gas accre0.0 frame,
masswith
extrasolar
nebula (MMEN, Chiang & Laughlin 2013)
cal code PLUTOof(Mignone
et al. 2007),
complementedathe planet’s
as described in Paper I; Paper II. Lastly,planetsimals,
anda
H/He-dominatedenvelope(allsolidstion
proceeds only after a massive enough ‘core’ has module
been to solve for radiation transfer in the FLD approximato provide
parameters for the unperturbed state of the
we have to take care of the gravitational interaction of the5
10et
end up in the core) (Pollack
al. 1994;
et yields
al. values for the middisc.
Taken
at a Hubickyj= 0.1 au, thisassembled.
numberfor
criticaltion
core
0.5 A
(Kuiper et al. 2010). We use a Roe solver (Roe1981)
disc gas withthetypicallyplanetary quoted
core. As in
Paper II,
wethisuse an2005; Mordasini
al. equation2009).
However,in ⇢thelast 6years 3 and the temperatureplane
gasis density
to solve the Riemannproblem. TheetFLD(10)
0 = 6 · 10 g cm
exponential
smoothingin space
to modify
the Newtomass(CCM)
is ⇠10term
Earth
mass
(Stevenson
1982; Pollack
it implementation
has been realizedthat some
these
model
T1Minimal=of1000
K (Leeet al.parameters
2014, Eqns. (12) & (13)). The subsolved using anof the Generalized
nianthat results
a force-freeboundary
1.0in a– way
etal.force
1994)
a numberthatinfollows
fromsolvingat1D stellar1
t = 700.0 1
t = 900.0 1
t = 1000.0
1standsfor unperturbed
mid-plane quantities. The
Residual (GMRES)
method (Saad& Schultz
1986)script
includingthe core’s surface:arechosenratherconservatively.
In
envelopes,
grainswill
structure like hequations.
Assuming radial symmetry, restarts
these for faster convergence. For more technicalobtained
details ondensityetis al.
a factorof Mora few larger compared to theGMc 1.0 ⇣0.5
r0 ⌘ p i0.0
0.5
1.0
1.0
0.5
0.0
0.5
1.0
1.0
0.5
0.0
0.5
1.0
quickly
coagulateandsettle(Movshovitz
2010;ã 2b = compute
exp Ahow thee rtemperature,,(12)modelspressure,
anddensityminimum-mass solar nebula (MMSN, Weidenschilling 1977;
x [RHill] the code, please refer to Kuiper
x [RHill et] al. (2010).
r x [RHill ]r2

y [RHill ]arXiv:1707.08079v1 [astro-ph.EP] 25 Jul 2017

Accepted 2017 July 22. Received 2017 July 18; in original form 2017 June 7

dasini 2014; Ormel 2014). A zero grain opacity may be more
profilesevolveas functionof time.
Asepointed
out by Piso &Hayashi 1981). The exact value of the background density
where r is
the distanceto the planet
center,r is the radial
realistic, especially in the casewithout
solid accretion. It isshouldYoudin(2014),
thesemodelscomeindi↵erent
flavours:
staticunit
vector= 10, p =of8 theare density
parametersthat
ensurea in the mid-plane for simulation RT-M2-VLO. Dashed circles
Figure3. andTimeAevolution
and theflow
structuremarknotthea↵ect the qualitative outcome of our simulations,also
restrictive
to assume
Bondi(blue)
andHill(green)radius.
Theandarrowsshow
gasavelocity
andsubthe grey lines
are rather
corresponding
streamlines.
By symmetry
sincecontinuous
we thecover a planetesimalrange of opticalacdepths by using di↵erent
quicktransition
betweenthe smoothed
Newtonianmodelsconsideraccretionof solids
topureprovidesteadyverticalbetween
velocity rcomponent
vanishes
here.
3
PHYSICALINITIAL
CONDITIONS
cretion.
Without
anexternal
sourceforluminosity
andintroduce
grain an additional scaling
regime
⇠
r
2
r
.
In
this
work,
the
smoothingopacities.
However,it might
00stantial luminosity to support the envelope (e.g. Stevenson
length r0 is given by the physical radius of the planetaryfactorfor the density profiles that we find.
opacity& Ikoma(2010)
calculateIn this section,
we first Hori
describe
the physicalmodelof the that cores as low as
1982;2006),assuming
while ina quasi-static
modelshydrostatic
core r Rafikov. It is calculated
mean density of
the core

Recyclingが作用する<=先行研究結果と一致

0.75
1.9

0.2
0.07

5.2 – 5.4%
5.1 – 5.4 %We confirmed that without including the planet’s grav
itational potential, the code accurately maintains the initia
equilibrium state for the full simulation time.
3.2

Opacities

As mentioned above, a dust-free environment is a viable sce
nario for super-Earth planets that we are focusing on. Due t
the close-in location, the atmospheres easily stay hot enoug(T & 1500 K) to e↵ectively evaporate dust grains, rendeing their contribution to the opacity negligible (see e.g. Le& Chiang 2015, their Fig. 6). Another reason to conside
a dust-free scenario is that the planetary core has alread
formed and incorporated most of the dust in its surroundings. Where the background disc is below the sublimatio
temperature of dust, coagulation of dust grains and settlin
is another e↵ective mechanism for grain removal (Orm2014; Mordasini 2014). These mechanisms result in opacitiethat are reduced by a few orders of magnitude compared ttypical ISM opacities. More generally, the opacity of ho
pre-planetary atmospheres will be determined by the ga
and exhibit a complex dependence on temperature, pressur

1.0

0.8
0.6

Hydrodynamics of embedded planets’ first atmospheres - III

0.2
0.0

0.0
0.2

kB
baryon0.4

0.2
kB
baryon

0.4

0.5

0.5

0.0
x [RHill ]0.5

1.0 1.0

0.5

0.0x [RHill ]

0.5

1.0

y [RHill ]

0.5
0.0
0.5
1

t = 1000.0

1

t = 2000.0

RT-M1-VLO

t = 100
t = 500
t = 1000
t = 2000

1

0.8

1.0

1
1
1

0.30

0.100
0.075
0.050
0.025
0.000
0.025
0.050
0.075
0.100

1.0 t = 100.0

RT-M1-LO

0.6

1.0

0.25
kB
S ) baryon

1.0 1.0

3D,1D (

0.5

RT-M1

0.4

0.8

kB
baryon0.0
x [RHill ]0.8

1

t = 2000.0

S

1.0

1

t = 1000.0S

0.6
1.0 t = 100.0 1
1.0
0.5

13

0.0

S

y [RHill ]

0.5

0.20
0.15
0.10
0.05
0.00
0.2

0.4

1

0.6
r [RBondi ]

0.8

1.0

0.2

0.4

0.6
r [RBondi ]

0.8

1.0

0.2

0.4

0.6
r [RBondi ]Figure 9. Time evolution of the radial profile of the spherically averaged specific entropy for a 1 M core. Colours mark di↵erent times
while the three columns display results for decreasing opacity from left to right. The top row shows both 1D (solid) and 3D (dashed)results while the bottom row shows their di↵erence. As expected, the radiative cooling is fastest for the lowest optical depth. Quickly,
the entropy in the 1D atmosphere deviates from the 3D result, which e↵ectively remains at the initial entropy for r & 0.4 RB . This
Figure 7. Mid-plane flow structure for RT-M1-VHO (top) and RT-M1-VLO (bottom), similar to Fig. 3. Dashed circles mark the Bondi (blue)
discrepancy grows as the cooling stalls in the 3D case, while the 1D atmosphere continues to lose entropy. To illustrate the convergenceand Hill (red) radius. Coded with colors is di↵rence in specific entropy to the unperturbed mid-plane value. The two bottom rows show
to a steady state for the 3D model RT-M1-VLO, the cyan lines in the right panel show the state at twice our usual simulation time t = 2tend .
the same data, but with di↵erent scaling of the colormap. The bottom row emphasizes the mixing of gas and resulting lower entropy
Note that the gradient inversion close to r = 0 is an artefact of the smoothed gravity.

1.0

0.5

0.0
x [RHill ]0.5

1.0 1.0

0.5

0.0x [RHill ]0.5

1.0 1.0

0.5

0.0x [RHill ]

0.5

1.0

3D大気：高いentropyを維持

outflow in the mid-plane.

disc Bondi/Hill球への高entropyガスの移流：

time however, the inner region is also a↵ected by shearing
and horseshoe orbits, which inject high entropy gas whileremoving low entropy gas from the very inner regions. This
process is emphasized in the bottom panel of Fig 7. Comparing the results for di↵erent opacities, we see that more
efficient cooling (and thus lower entropy) promotes the transition towards more circular – and arguably more bound –motion around the core. We find a similar structure for simulation RT-M2-VLO.

Paper II. Apart from the more realistic thermodynamics,

1 RB, 0.5 RB and 0.35 RB . All quantities are averaged overitashould also be realized that in this work we focused on
the long-term behaviour, while sacrificing resolution (whichtime of one ⌦ 1 . Due to the symmetry of our model, we deaccording to Paper II to be important). In adscribe the longitude of only one feature, the other will matteredbedition, the adopted gravitational smoothing in our model
found at 0 = +⇡. The longitude of = 3⇡/2 corresponds torenders the study of the innermost regions somewhat problooking at the planet from the y direction, following its orlematic; a pure Newtonian potential is far preferable (butbital direction. On the surface of the Bondi sphere the massruns into problems with PLUTO, see Paper II). Runs at higher
flux through the atmosphere is dominant in the mid-plane,
resolution (Fig. A1) and di↵erent values for rsmooth (Fig. A3)
where mainly the shearing and horseshoe flow is passing perdo show a slightly di↵erent flow pattern for the inner corependicularly through it. The intensity of the mass flux region,ismainly a↵ecting the rotational feature, but also thatthose parameters do not a↵ect the main results. This is dishighest there and reduces towards higher latitudes, where
further in the Appendix A.
the outward flow dominates. Comparing the flow pattern cussedtothe entropy plot below, we recognize the low-S outflow from A relevant question is how this inner low-S region will
evolve on evolutionary time-scales (recall again that the ⇠103
Fig. 8 in the mid-plane. For this feature at ' 3⇡/4, theit is clear that the efficiency of recycling strongly depends
on the fraction of atmospheric gas that is replenished withgas from the disc. Given the above described topology of theflow, the distribution of mass as a function of distance fromthe core plays an important role. From Fig. 4, we see thatin our simulations, where we fixed = 7/5 for simplicity, themass of the atmosphere is dominated by the outer parts thatshow significant recycling. However, Piso & Youdin (2014)
and Lee et al. (2014) argue that close to the planet, where
the temperature is sufficiently high (T 2500 K), accounting
for dissocation of H2 will result in more realistic values of 4/3. Depending on the ratio of the core radius to theouter radius of the atmosphere rc /RB,H , this value of canrepresent a transition to an atmospheric structure, wheremost of the mass accumlates close to the core. If this is thecase, recycling time-scales deep in the atmosphere could dramatically increase. However, by integration of the exact hydrostatic density profile (not the power-law approximation)
we find that for values of rc /RB,H (see Table 2) corresponding
to our planets, this transition occurs at an adiabatic index
of only ' 1.2. Still, non-ideal e↵ects as H-dissociation (resulting in non-uniform ) could yet render the atmospherecentrally condensed, but this we leave as an exercise to be investigated by future models. For a more realistic treatment,
future studies should also consider tabulated, non-constantopacities. Altogether, given the short physical times covered,
it is hard to predict the long-term thermal and dynamical
evolution of this critical region based on the simulations of
this work.

Kelvin-Helmholtz収縮を鈍くし、envelopの成長を止める（放射冷却により大気中ガスはentropyを失うが）We emphasize however that 3D e↵ects are key. Fig. 8
displays the mass flux Fm (in units of Earth masses per
inverse orbital frequency per surface of the Bondi sphere)
and the specific entropy S on a spherical surface at r =inner coreを同定

orbits only corresponds to 5 yr). From Fig. 7 one observes
that the low-S region expands with time. Possibly, this may
imply that the core becomes more isolated, although Fig. 4
MNRAS 000, 1–16 (2017)illustrates that radial recycling rates remain robust.6 Likewise, the low-opacity run in Fig. 9 also shows that the point
where the 3D entropy profile detaches from the 1D curve
moves inwards with time: r = 0.19 at t = 500 ⌦ 1 to r = 0.14
at t = 103 ⌦ 1 . These are clear manifestations of continued recycling. Finally, physical e↵ects, like dissociation ofH2 and evaporation of silicate grains or pebbles, may decide
the evolution of the inner core (possibly already during theassembly of the super-Earth planet). From the above results

・サイズ Bondi半径の25%・流線 円・entropy その外側大気よりもかなり低い
6

Unfortunately, because of the low resolution, we were unable toconduct a topology analysis, that is, it is possible that some ofthe streamlines in the inner region are closed.
MNRAS 000, 1–16 (2017)

molecules, determining the formation of a PDR (Tielens &Hollenbach 1985; Kaufman et al. 1999; Röllig et al. 2007;
Bisbas et al. 2012) at the surface of the clump itself. Furthermore, radiation above the Lyman limit (h⌫ > 13.6 eV)
ionizes neutral atoms, and it is completely absorbed within
1
1,2 a shallow layer.
3,1,4,5
1
6
The goal of this paper is to understand the evolution of
1 Scuola Normale Superiore, Piazza dei Cavalieri 7,
I-56126 Pisa, Italy
and the density profile of molecular clumps exposed
2 Kavli IPMU, The University of Tokyo, 5-1-5radius
Kashiwanoha, Kashiwa 277-8583, Japanto a UV
radiation field including hydrogen-ionizing photons
Figure1. Density
3 Centro Fermi, Museo Storico della Fisica e Centro
Studi e Ricerche “Enrico Fermi”, Piazza del Viminale 1, Roma,
00184,
Italy profile of molecular spherical clumps at temproduced
byMadingleymassive stars
or Cambridge
quasars. The
key0HA,
pointUKis thatperature Th2 = 10 K, confined by a medium with Ticm = 100 K
4 Kavli Institute for Cosmology, University of Cambridge,
Road,
CB3
and nicm = 103 cm 3 , computed with the BE model The outeran ionized
and anAve.,atomicshell form
at theedge
5 Cavendish Laboratory, University of Cambridge,
19 J. J.shellThomsonCambridge
CB30HE,
UKof thedensity depends only on Th2 and the confining ICM pressure,The dynamics
of thislayered structure23,
is determined
6 Nordita, KTH Royal Institute of Technology clump.
and Stockholm
University,
Roslagstullsbacken
SE-10691 Stockholm, Sweden
while the mass sets the density at the centre of the clump.
by the fact that each layer is at a di↵erent temperature and53pressure. For a clump with initial density n0 ' 10 cm ,typical temperatures deep into the clump are T ' 10 100
28 July 2017
nicm [ cm 3 ]Ticm [K]Th2 [K]K, while an atomic (ionized) region can be heated up to
3
4around T ' 10 K (T ' 10 K). We denote this type of
Stellar case10
100 1000
10
ionization/photodissociation regions as iPDR.
Quasar case60
2.2 ⇥ 107
100
ABSTRACT In particular, we apply our model to two scenarios.

#13 Molecular clumps photoevaporation in ionized regions

D. Decataldo , A. Ferrara , A. Pallottini, S. Gallerani , L. Vallini

arXiv:1707.08574v1 [astro-ph.GA] 26 Jul 2017

UV放射下における分子クランプの進化（半径・Life time）を調べた
クランプ

isothermalのBonnor-Ebert球We study the photoevaporation of molecular clumps exposed to a UV radiation field
Stellarpropagatingcase: a molecularclump
in pressure
equilib・OB型星の場合
Figure 7. Schematic representationof •wavesin a clumpby radiation,
where
including hydrogen-ionizing
photons(h⌫ >issuddenly
13.6eV)heated
producedby massive
starswavy
or lines are shock waves
rium within a GMC in the proximity of an OB star, and it

Table 1. Clump and ICM properties at the onset of the photoevaporation process.

and double dashed lines arequasars.
rarefaction
Highpropagationpressure shells
a shock
into adjacent
inner clumps
shells, andWe waves.
follow theanddrive
collisionof shockwaves inside
andas a result a rarefactionissuddenly
engulfedbythe
expanding
HII
region.
take into accountself-shieldinge↵ects,the evolutionof theclump
size andspeed. The result is anwave propagates back. Discontinuity
interfaces
move in thesame determining
direction of shockwaves, atpost-shock
Quasarcase:
a clump forms
asionization-photodissociationa result of thermal inpressure
with•time.
The
theforregionequilibrium(iPDR) with atomic gas, whose temperature
expansion of the two shellsdensity
and a contraction
of
thestructurecore. See oftext
the a detailed description.

距離 0.3pc、0.1Msun、初期半径~0.02pc
stabilities within the outflow, finds itself embedded in thedepends on the distance and luminosity from the stellar
is obtained for di↵erent initial clump masses (M = 0.01 104 M ) and impinging
ionized wind and exposed to the quasar radiation.
source. We takenat ' 103 cm 3 as a typical density for thisfluxes (G0 = 102 105 in units of the Habing flux). The cases of molecular
clumpssurrounding
gas,areand we compute the corresponding temworks inliterature
concentrated
mostlywithinonengulfed inPrevious
the HII regionof thean OBstar andclumps carried
quasaroutflows
peratureaccordingthe e↵ectsWeof non-ionizing
photons
onundergoes
photoevaporationof cases an initial shock-to the FUV flux, using the same method
treated separately.
find that theclump
in bothused to computethe temperature of the atomic phase in theHollenbachexpansion
2002; Adams
et al. which
2004; Chamcontractionclouds
phase(Gorti
and a& following
phase,
lets the radiation
penetrate
clump (see Sec. 2.3). Such temperature ranges between 102
pion
et
al.
2017).
The
evolution
in
their
case
is
simplified
byin until the clump is completely evaporated. Typical evaporation time-scales are' 0.01
K and 103 K. The clump is exposed to the radiation of thefact thatclump
is exposed
to radiation
below the clump
Myr in thethe
stellarcasetheand
0.1 Myrin theonly
quasarcase, where
mass is 0.1
massive star when it is engulfed in the growing HII region,
3 Lyman limit. As a result, the clump develops a single
M and 10theM respectively. We find that clump lifetimes in quasar
outflows are
for which the density is taken to be nicm ' 10 cm 3 .
structure.
compatibleshellwiththeir observed extension, suggesting that photoevaporationis theIn the caseof quasars, molecular clumps likely form
Theregulating
paper is organized
In Sec.outflows.
2 we describe
main mechanism
the sizeasoffollows.
molecularfrom thermal instabilities within the outflow. Clumps detachthe model adopted for the structure of molecular clumps,

from thehot phase at the discontinuity between the fast
togetherthe physics
of the photodissociation
dynamical and thermal
pro- - quasars:
Key words:
ISM:withclouds,evolution,
regiongeneral1

INTRODUCTION

wind and the ISM (Zubovas & King 2014), starting fromcesses involved. In Sec. 3 we present the results for the evoluthe distance at which the outflow has become energy-driven.
tion of the clump radius, and compute the evaporation time
This critical radius has been estimated by Zubovas & Kingof a clump for di↵erent parameters of the system. The re(2012):
sults are presented separately for clumps located near stellar
Werf4 weet al.
1993; Luhman
et al. 1998). Detections of fineor quasar sources. In Sec.
summarize
our results.
1/2
Rc = 520 200 M8 v0.1 pc ,(1)
structure lines of [CI] and [CII], high-J CO rotational lines,

The di↵use interstellar medium (ISM) is characterized by
and J = 3 2 lines of HCN and HCO, show
that200
PDRs
arevelocity dispersion in the host galaxy
whereis thea turbulent multi-phase structure, showing a broad range
2 4
made of a low-density, more di↵use component
10 km s 1 , M8 is the mass of the SMBH inin units(nof '200
2 MODEL In
of densities, temperatures and chemical compositions.
3
6 7
83cm ), and high-density structures (n ' units
10 ofcm), such10 M
and v0.1 is the wind velocity in units of 0.1
some regions, gravitational forces and pressure
the environment
2.1compress
Gaseoustheandspeedof light. In the outflow, the gas is heatedas in M17SW (Hobson 1992; Howe et al.times
2000),
in thegas to sufficiently high densities, so that the formation oftoTicm 'must
2.2 ⇥have
107 K and has a typical density nicm ' 60
Orion
barclumps(Lis &are
Schilke2003).
These
clumpsThe
gaseousenvironment
where
locatedplays
a3
molecules such as H2 and CO is allowed.
& King 2014; Costa et al. 2014). The outflow
sizes properties
as small assuch
one astenth
or a hundredthcm
of pc,(Zubovas
often showcrucial role in determining
temperature,
CO
maps
have
revealedthat
Giant
Molecular
CloudsFigure 8. Left: radius, R, of the molecular
core
of
a
clump
as
a
function
of
time,
when
the
clump
is
exposedstellar radiation
fragments(King
2010;toNayakshin& Zubovasof2012) because
elongatedshapes.
Thethepresence
density and confining ing
pressure.
We now
describe
inter- of such clumps a↵ects(GMCs)
containa very rich
structurefilofthermalinstabilties,
so
that
one
componentdi↵erent
luminosity
L. internal
The clumphasclump
afeaturing
mass
of
0.1
M
,
initialradius
r
'
0.02
pc,
and
is
located
0.3
pc
away
from
the
source.
Right: cools to acemission
the clumps
spectrummedium (ICM) significantly
which surrounds
in thestellar of stellar PDRs.
amentsand clumps
et al. (located1987; Bertoldi&pcMcKeelowtemperature.
The
existenceofanequilibrium
betweenLifetime(tc ) of (Bally
the clumps
at
0.3
from
the
source)asafunction
of
their
mass
and
stellarsourceluminosity.
Clumpsin theand quasar case. The ICMAccording
propertiestoaresummarizedinrecentobservations(Cicone et al. 2014),

励起星の光度が大きいほど、shock-contraction phaseが早くなる・Quasarのアウトフローの場合
距離10

1 kpc、104Msun、

Decataldo et al.

Contraction
phaseの時間はほぼ一定
Figure 10. Left panel:radius (R) of the molecular core of a clump as a function of time, when the clump is exposed to quasar radiation.
The clump has a mass of 10 M and di↵erent luminosities are considered. Right panel: Lifetime (t ) of a clump exposed to the radiationof a quasar, for a range of values of the initial clump mass and di↵erent source luminosities. The distance of the clump4fromthe source
全ての場合において、クランプ外層を完全に電離し、10Kに温めるため
is fixed at 1 kpc.
3

c

L
F
G
クランプがcontact
discontinuity(CD)で出来ると仮定
ion

103 L
104 L
105 L

?

0.006
0.5
14

0

80
1.2 ⇥ 103
9.5 ⇥ 103・quasarでの分子ガス構造の進化において、photoevaporationが重10 erg s
3.4
6 ⇥ 10Quasar
要なメカニズムかも。
10 erg s
34
6 ⇥ 10(1 kpc)OB star(0.3 pc)

45

1

2

46

1

3

1047 erg s 1

340

6 ⇥ 104

Table 2. FUV and ionizing fluxes at a distance D = 0.3 pc in・クランプがアウトフロー中で連続的に生まれる必要がある？
stellar and D = 1 kpc in the quasar case. Fion is in units oferg s 1 cm 2 , G0 is in units of the Habing flux.

#14 How fast is mass-segregation happening in hierarchicalarXiv:1707.07710v1 [astro-ph.GA] 24 Jul 2017

formed embedded star clusters?
R. Domı́nguez1⋆ , M. Fellhauer1, M. Blaña2, J.P. Farias3, J. Dabringhausen4

1 Departamento de Astronomı́a, Universidad de Concepción, Casilla 160-C, Concepción, Chile

2 Max-Planck-Institut für extraterrestrische Physik, Gieβenbachstraβe 1, 85748 Garching bei München, Germany
3 Department of Astronomy, University of Florida, Gainesville, FL 32611, USA
4 Astronomical Institute, Charles University of Prague, V Holesovickach 2, 180 00 Praha 8, Czech Republicmass segregation (MS) の進化を調べる（5Myrまで）
Accepted —–. Received —–; in original form —–２つの星団を考える
We investigate the evolution of mass segregation in initially sub-structured young
embedded star clusters with two diﬀerent background potentials mimicking the gas.
Our clusters are initially in virial or sub-virial global states and have diﬀerent initialvirial or sub-virial、最大質量の星の分布が異なる（以下）distributions for the most massive stars: randomly placed, initially mass segregated or
ABSTRACT

3つの初期分布even inverse segregation. By means of N-body simulation we follow their evolution for
5 Myr. We measure the mass segregation using the minimum spanning tree method
ΛMSR and an equivalent restricted method. Despite this variety of diﬀerent initialconditions, we find that our stellar distributions almost always settle very fast intoa mass segregated and more spherical configuration, suggesting that once we see aspherical or nearly spherical embedded star cluster, we can be sure it is mass segregatedno matter what the real initial conditions were. We, furthermore, report under which
circumstances this process can be more rapid or delayed, respectively.

・NO-SEG ランダムに質量を設定
Key words: methods: numerical — stars: formation — galaxies: star formation —galaxies: star cluster: general
・SEG-IN 全ての大質量星>4Msunを0.5pc以内に配置
・SEG-OUT 大質量星を1.0pcよりも大きな半径に配置
Mass segregation in hierarchical clusters

3

namical (McMillan et al. 2007; Allison et al. 2009; Yu et al.
1.52011) and fast for
sub-structured
clusters
Massivecool
particles and
Massive particles
Low mass particles
Low mass particles
Mass segregation (MS) has been detected in embed(Allison et al. 2010; Parker et al. 2016) reaching some level
ded star clusters1 (Lada et al. 1996; Hillenbrand 1997; 1
of MS in ∼ 1 Myr. Primordial MS is 1also possible, as
Hillenbrand & Hartmann 1998; Bonatto & Bica 2006;a result of the star formation process (Zinnecker 1982;
Chen et al. 2007; Er et al. 2013), i.e., that the most massive
Murray & Lin 1996; Elmegreen & Krakowski
2001; Klessen
0.5
0.5
0.5
stars are more centrally concentrated. Advances in obser2001a; Bonnell et al. 2001; Bonnell & Bate 2006) where
vations have shown that star clusters form sub-structured
the massive stars tend to form in thecentre of the00(Larson 1995; Elmegreen
2000; Testi et al. 2000; Williams 0
star-forming regions because of higher accretion rates in
2000; Cartwright & Whitworth 2004; Gutermuth et al.
the central parts or as a result of competitive
accretion-0.5
2005; Schmeja &-0.5Klessen 2006; Carpenter & Hodapp 2008; -0.5
(Larson 1982; Murray & Lin 1996; Bonnell et al. 1997). The
Schmeja et al. 2008) and this has also been supported
later has been used to explain very high levels of MS-1-1
by simulations (Klessen& Burkert 2000, 2001; Bate et al. -1
in very young objects, when dynamical processes are not2003; Bonnell et al. 2003; Bate 2009; Oﬀner et al. 2009;
fast enough to explain the observed MS (Bonnell & Davies
Girichidis et al. -1.5
2012; Dale et al. 2012). Observations also -1.5
Raboud&0 Mermilliod
1998).
On-1.5-1.5
the otherhand 0
-1.5
-1
-0.5
0
0.5
1
1.5
-1.5 1998;
-1
-0.5
0.5
1
1.5
-1
-0.5
0.5
1
1.5
show that stars form dynamicallyx [pc]cool (sub-virial state)x[pc],, primordial mass seg- x [pc]Parker et al. (2015) pointed out that
(Peretto et al. 2006, 2007; Proszkow & Adams 2009) andregation rarely occurs if massive stars form by competitive
this property
is supported
by hydrodynamical
simulations
Figure
2. Examplesof the three diﬀerent
initial distributions of massive stars assumed in this work (shown in 2D). Red filled circles
INTRODUCTIONy [pc]

1.5

Massive particles
Low mass particles

y [pc]

1.5

y [pc]

1the minimum spanning tree methodと8 R. Domı́nguez
et al.
equivalent
restricted
methodを使用
1.5

Massive particles
Low mass particles

1.5

Massive particles
Low mass particles

1

1

0.5

0.5

0.5

0

y [pc]

1y [pc]

y [pc]

1.5

球状の分布になる

0

-0.5

-0.5

-1

-1

Massive particles
Low mass particles

0

-0.5

-1

Mass segregation in hierarchical clusters
-1.5
-1.5

-1

-0.5

0

0.5

1

-1.5
-1.5

1.5

-1

-0.5

0

0.5

1

1.5

-1.5
-1.5

-1

-0.5

9
0

0.5

1

1.5

Table 1. Summary of results. The first, second and third column are giving information of the initial conditions: type of BG potential,x [pc]
x [pc], and Λrest respectively. Thex [pc]mass distribution and virial ratio respectively., The fourth and fifth column show the final values for ΛMSR
values are calculated by taking an average value of all measurements between 4 and 5 Myr. The sixth column gives tseg i.e. the information
Figure 7. Examples
thestays
final
state
the
three diﬀerent
distributionof massive
shown
Fig.
when (tsegof) Λrestlargerthanfor2, i.e.,a significant
level of MSinitialis detected.
The last two columns
are the stars
final values
for r8inand
Rh 2. Left panel: Final state
respectively.
we NO-SEG.
calculate themCentral
as averagespanel:from all Finalmeasurements
4 and 5 Myr.
of a fractal whichstartsAgain
withstatebetweenof a fractal
which starts with SEG-IN. Right panel: Final state of afractal which starts with Potential
SEG-OUT.
Mass placement
Uniform
Uniform
PlummerPlummer

NO-SEG
NO-SEG
NO-SEG rest
NO-SEG

Qinit

ΛMSR,final

Λfinaltseg [Myr]r8 [pc]Rh [pc]

0.5
0.2
0.5
0.2

2.35 ± 0.94
2.32 ± 1.38
2.65 ± 1.53
2.78 ± 1.13

2.68 ± 0.86
2.76 ± 0.82
2.51 ± 0.78
2.74 ± 0.74

0.4
0.4
0.8
0.6

0.82 ± 0.20
0.69 ± 0.19
0.90 ± 0.24
0.65 ± 0.12

0.79 ± 0.14
0.65 ± 0.11
0.86 ± 0.17
0.64 ± 0.10

this initial contraction of the low-mass stars as the main
stant plateau. The maximum value of Λ
= 3.03 is reachedreason for the fastly declining MS values. In all four panels
after 2.16 Myr. The final value Λfinal = 2.74. We measure a
Uniform
SEG-IN
0.5
2.58 ± 1.24
2.69 ± 0.69
± 0.170.73 ± 0.13
wealwayssee an0.77
initialdeclineof r8 as well, i.e. also the alreadytseg = 0.59 Myr.
Uniform
SEG-IN
0.2
2.03 ± 1.39
2.56 ± 0.66
always0.71 ± 0.21
0.64 ± 0.11
± 2.282.89 ± 0.82 centrally
always
0.79 ± 0.16
0.77 ± high-mass
0.14
concentrated
stars are moving towards
In the lower half of Plummer
this panel,SEG-IN
we see that0.5the3.49
cold
initialPlummer
SEG-IN
0.2
2.20 ± 1.55
2.79 ± 0.68
always0.64 ± 0.18
0.61 ± 0.12
the centre as well. This should increase the level of MS;
conditions produce a similar collapse or contraction in all
Uniform
SEG-OUT
0.5
1.43 ± 0.58
2.12 ± 0.57
3.5
± 0.320.87 ± 0.14
instead
we1.02see
a rapiddecrease in both methods. While insimulations resulting inUniform
small error-bars
for R. The steeper
SEG-OUT
0.2h 1.72 ± 0.91
2.70 ± 0.72
1.3
0.68 ± 0.16
0.63 ± 0.07
Plummer
SEG-OUT
0.5
1.18
±
0.50
1.61
±
0.39
never
1.22
±
0.31
0.97
±
0.16 this could be due to ejected highthe non-restricted method
potential of the Plummer BG helps to damp the oscillations
Plummer
SEG-OUT
0.2
1.43 ± 0.69
2.14 ± 0.60
3.8
0.86 ± 0.38
0.69 ± 0.12
mass stars, our restricted method is designed to avoid suchrather quickly. It settles at about the same value as in the
results.
uniform, cold case.
to be behaviour
primordial (e.g.
Bonnell
Davies to the NO-SEG cases
fastcases
evolution
MS but a slow increase
with thethe uniformof R
similar
In all four
we inapproximately
obtainsame val-where MS hasTheh is &
1998; Hillenbrand & Hartmann 1998; Raboud & Mermilliod
BG (Λfinal = 2.12 ± 0.57 and tseg = 3.5 Myr) or almost no
above. Parker et al. (2015) foundues of MS at evolution
the endofwiththethesimulations.
methods to1998) is explained
not strictly necessary.
at allPlummer BG (ΛBoth
final = 1.61 ± 0.39
MS istop-leftless common
whenagainin the processtseg to= never).
shows that
small
contraction ofwiththat primordialThepanelshowsofthe uniform BG–Qinit =measure MS andtend
agreeThiswithin
thethe1σ
deviations
star formation the influence of feed-back is taken into acour stellar distribution is not enough to ensure suﬃcient in0.5 case.
The simulations reach a final value of Λfinal =the restrictedteractions
methodshowing slightly higher values. Thecount. However,
we cannot exclude primordial MS, which
between the high-mass stars and the lower mass
2.69by±0.69 with simulations
tseg = (e.g.
always(in fact all panels show
agreement seems
todrivebe thembetterwhentowards
a Plummerstars toeﬃciently
the centre. backgroundHow realis supportedhydrodynamical
Klessenisticsimulations
such initial conditions
is debatable.
al. 2001; Bonnell
Bate 2006) oreven ashows the results for thetseg =et always).
The&top-rightpanel
was used. The
with are,uniformbackground have2001b; Bonnell

ほぼ全ての場合において、短いタイムスケール<1 MyrでMS

#15 The evolution of Giant Molecular Filaments
4

1,2? , and C. L. Dobbs2
Ana
Duarte-Cabral
1

Duarte-Cabral et al.

School of Physics & Astronomy, Cardi↵ University, Queen’s building, The parade, Cardi↵, CF24 3AA, U.K.

15.2 Myrs
2 Schoolof Physics, University of Exeter, Stocker Road, Exeter, EX4 4QL, U.K.

y (kpc)Spiral galaxyのシミュレーションで、GMFsの時間進化を調べる(11Myr)（modelの詳細はDobbs et al. 2008, 2011, 2013, 2014, 2015, 2016）
GMF1

ABSTRACTThe evolution of Giant Molecular Filaments

3

In recent years there has been a growing interest in studying giant molecularfilaments (GMFs), which are extremely
GMF2elongated (> 100 pc in length) giant molecular
clouds (GMCs). They are often seen as inter-arm features in external spiral galaxies,
but have been tentatively associated with spiral arms when viewed in the Milky Way.
In this paper, we study the time evolution of GMFs in a high-resolution section ofa spiral galaxy simulation, and their link with spiral arm GMCs and star formation,over a period of 11 Myrs. The GMFsGMF1
generally survive the inter-arm passage, although
they are subject to a number of processes (e.g. star formation, stellar feedback anddi↵erential rotation) which can break the giant filamentary structure into smaller
sections. The GMFs are not gravitationally bound clouds as a whole, but are, to someextent, confined by external pressure. Once they reach the spiral arms, the GMFs tend
to evolve into more substructured spiral arm GMCs, suggesting that GMFs may be
precursors to arm GMCs. Here, they become incorporated into the more complex andalmost continuum molecular medium that makes up the gaseous spiral arm. Insteadof retaining a clear filamentary shape, their shapes are distorted both by their climb
up the spiral potential and their interaction with the gas within the spiral arm. The
GMFs do tend to become aligned with the spiral arms just before they enter them(when they reach the minimum of the spiral potential), which could account for theobservations of GMFs in the Milky Way.

The evoy (kpc)y (kpc)

9.0 Myrs
17.2 Myrs

11.1 Myrs
19.0 Myrsy (kpc)y (kpc)arXiv:1706.05421v1 [astro-ph.GA] 16 Jun 2017

Accepted 2017 June 15. Received 2017 June 15; in original form 2017 March 1

Key words: ISM: clouds – galaxies: ISM, star formation, spiral – Methods: numerical13.2 Myrs

1

x (kpc)INTRODUCTION

to the larger-scale Galactic environment. In particular, after

Figure 2. ofFig. 1,
continued. where the background grey scale shows the top-down view of the total column density, the cyan
Figure 3. Time evolution
GMF1,
the discovery of the striking ⇠ 80 pc long “Nessie” filament
Observationsofnearbygrand-design
spiral while
galaxies
havehighlight
thebluecontoursof the smaller
clouds ofthat were or will be once part of the(Jacksonet &al. follow2010) insome
the Galactic
plane, a number
minimum1 . contours
Some of these GMFs
can in factthespan main> 500 pc filament,
In accordance withthe dark
results from
Duarte-Cabralbeen
in revealing
thedistribution
ofthatdark
2 are
in length in H2long, but asis clearfascinatingfrom the left panelsof Figs. 1
Dobbs(2016), we findgiant molecularfilaments
otherobservational
studies(e.g. clusters)
Beuther et al.
2011;
Li etcolour-coded
al.
main
GMF.
Circlesshow
thepopulationofstar
particles(representing
stellarsizeandby their age (large/blueand 2, the COlanes(seen asof
white-pink
only intense
present on regions of star-formation,
dense areas)
gas isand2013;
Goodman
et
al.
2014;
Ragan
et
al.
2014;
Zuckeretal.
the denser ridgesof theseclouds, splittingthese
GMFs into being 5 Myrs or older). The dashed lines delineate the bottom of the potential well, as in Fig. 1. The only
being
young,
andandsmall/red
acrossspiralarms
inter-arm
regions (e.g. Elmegreen
smaller filaments of the order of ⇠100 pc in length.
2015;
Wang
et al. 2015a, 2016; Abreu-Vicente et al. 2016)
2 GMFs were identified by selecting the sample
of clouds
from our
1980, 1989; star
Scoville
et al. 2001;event
La Vigneal. 2006;
Corderat t = 12
significant
formation
in et
GMF1
occursMyrs,
whichresults in a visible shell on the following timesteps. At the end of
have
sinceaxistried
to find other Nessie-like filaments in the
5 pc-grid cubes with an aspect ratio larger than6, a majorofet al.
2008;
Schinnerer
et al.close
2013,to2017).
atleastThough30
pc, and a limited
minor
axis
ofin
lesspotential
than 15 pc. The aspectratio and has yet to enter the gaseous part of the spiral arm (where there is
the
run,
GMF1
is
still
the
bottom
of
the
well,
Milky
Way, other ‘giant molecular filaments’ (GMFs), and
1 Note that the spiral potential minimum is slightly shifted withestimated using
the major andfrom the ellipsoidal
their ability
to probe anddefinedresolve
theisfitdetailed
structuresofminor axes
respect to thethe gaseous spiral arm (as ofbyof the 3D
density structure.
However,long
filaments canthem
often with the Galactic large-scale spiral structure.
associatea position
largeof concentrationbothstellarclusters
andgas).

inter arm領域のみで見られ、spiral armの軸に沿う
y (kpc)

GMFs：

GMFsは重力的に束縛されておらず、外圧によって制限されている

theof interstellar
medium,thosewith of filaments, in which case thethe higher densitygas), due to the specificresponseof the studies
gasbehavecurvedprovided
and be part ofus
networks

0
8

10

12

14
16
Time [Myrs]

18

20

22

0
8

10

12

14
16
Time [Myrs]

18

20

22

Figure 5. Time evolution of the total mass of GMF1 (left) and GMF2 (right). As the giant filaments evolve, they do not survive asa single entity but are are made up of smaller clouds as the main filament breaks up. The black solid line shows the evolution for theentire cloud (i.e. the total mass within the several fragments), which remains relatively constant over time. The coloured lines show theevolution of di↵erent cloud fragments, following the same colour scheme as Figs. 3 and 4: solid turquoise line for the main GMF1, solid
turquoise and solid magenta for the two main filaments of GMF2; and blue dashed lines for all the remaining smaller individual clouds
The evolution of Giant Molecular
Filamentsthat wereonce part of the7 GMFs.

molecular gas mass fraction

total mass
5

4

4

3

Feedback

2

GMF1

3

2

Feedback
1
0
8

Star
Formation
10

12

1

Bottom ofpotential well
14
16Time [Myrs]

18

20

0
22 8

Feedback

0.5

Entry ofgaseousspiral arm

Bottom ofpotential well
10

12

14
16
Time [Myrs]

0.4

GMF2

Bottom ofpotential well

0.5
Molecular mass fraction

5

0.6

0.6

GMF2

Molecular mass fraction

6

GMF1

log (Mtotal [Msun])log (Mtotal [Msun])

6

Star
Formation

0.3

0.2

0.4

0.3

0.2

Feedback
0.1 Bottom of

0.1

potential well
18

20

22

0.0
8

10

12

14
16
Time [Myrs]

18

20

0.0
8
22

10

12

Entry ofgaseousspiral arm
14
16Time [Myrs]

18

20

22

Figureevolution
gure 5. Time evolution of the total mass of GMF1 (left) and GMF2 (right). As the giant filaments evolve,
they 6.
doTimenot survive
as of the molecular gas mass fraction of GMF1 (left) and GMF2 (right). Solid and dashed lines are as in Fig.5.
Althoughthe total massof the clouds remains relatively constant over time, the molecular gas fraction increases with time, especiallyingle entity but are are made up of smaller clouds as the main filament breaks up. The black solid line showsthe evolution
for thewhilstthe clouds
are being
ire cloud (i.e. the total mass within the several fragments), which remains relatively constant over time. The
colouredlines show
the formed in the inter-arm (which is the case of GMF1). The molecular mass fraction of GMF2 also increases,
GMF2
does
not reachlution of di↵erent cloud fragments, following the same colour scheme as Figs. 3 and 4: solid turquoise line but
for themain
GMF1,
solidmore than 60% molecular due to the input of stellar feedback in the denser regions as soon as SF takes place.

・GFsは時間とともに壊れ、より小さなフィラメントへ
9

quoise and solid magenta for the two main filaments of GMF2; and blue dashed lines for all the remaining smaller individual clouds
The evolution of Giant Molecular Filamentst were once part of the GMFs.

GMF1 at t = 18 Myrs and GMF2 at t = 9 Myrs). In effect, as the gas is accelerated towards the bottom of the
DISCUSSION
0.6
0.6
spiral potential, the density contrast of the ridge of the filaGMF1
GMF2
5.1 ments
No ‘bones’
ofasthegalaxy? of the gas feeling a “frontal
Bottom ofincreases,
a consequence
Feedback
potential well
0.5
0.5
wind”,
orrampressure.
This havee↵ectively
works thatas a compresPrevious observational studies
suggested
the
sion
force
acting formon thesome
smalleraxis
of these orfilaments,
GMFs
in
the
Galaxysortofskeleton,“bonesofwhilst
0.4
0.4
the galactic
shearthefrom
the arms
di↵erential
rotation,continues
Star
the Galaxy”,
tracingspiral(Goodmanet al.
2014; toFormation
stretchcloudsin their(seeFig. 3). someThis shear
Zuckeret al.the
2015;
Wang
et al.longer
2016).axis
Thiscontrasts
0.3
0.3
suchobservations
that the starofformation
activity, where
althoughit does
whatiswith
external galaxies,
there
are ocis not particularly
strongin the moments
approachclearcur,filamentary
structures in
the inter-arm
regions,ofbut
in
0.2
0.2
of the arms,
bottomof gasthe tends
spiral topotential.
the spiraltheexhibit a more complex
Feedback Entry ofmorphology.
resultsthe
are bottom
clearly of
more
AfterOurreaching
theconsistentpotential withwell (i.e.
gaseous0.1 Bottom of
0.1
the observations
externalfila- theGMF1 at theof endof thegalaxies.
run, or Giant
GMF2molecularat the start),spiral arm
potential wellments
exist start
everywherein back
the inter-arm
but theycloudsto climb
up. The regions,accelerationthat the0.0
0.0
are simply
shearedcloudsofasgas,and doonnotconstituteavelocclouds
feel
is
now
actinga“break”
theircircular
8
10
12
14
16
18
20
22
8
10
12
14
16
18
20
22
physically
relevant
structure.
Furthermore
they
being
ities, andthis causesmaterial
to crowd
on cease
the downstream
Time [Myrs]
Time [Myrs]
giant
filaments
once
become or
joinFig.
the4,giant
sideofthearespiral(see
for moleculart > 14 Myrs).
gure 6. Time evolution of the molecular gas mass fraction of GMF1 (left) and GMF2 (right). Solid anddashedlines
astheyin potential
Fig.5.
that
make
upespeciallythe gaseous (orspiralCloudsstart
toincorporate
arearms.
being incorporated by)hough the total mass of the clouds remains relatively constant over time, the molecular gas fraction clouds
increases
with
time,
Althoughstriking
in theirandtheform,
giantpotential,
molec- efgas
cloudsalsoshape
exitingspiralilst the clouds are being formed in the inter-arm (which is the case of GMF1). The molecular mass fractionof from
GMF2otheralso increases,
8. Countdue
of numberof starevents inoverular fectively
filaments
aretakesnotanplace.
gravitationallyboundclouds asgas,
a i.e.
almost continuumof molecular
GMF2 does not reach more thanFigure
60% molecular
to the inputof formation
stellar feedback
thetime,denser regions
assoon ascreatingSF
for GMF1 (solid red line) and GMF2 (dashed blue line). This waswhole.
are spiralstructuresnaturally
shapedcombined
andtheThesegaseousarm. that
Thisaremergingof clouds,
done by tracking all the gas particles that make up each GMF
confined by their global galactic environment (shaped byover time, directly on the SPH data. We can see that GMF2
the rotational
shear,
andby the external pressure),
MNRAS
000,
1–13
(2017)
MF1 at t = 18 Myrs and GMF2
att=9
Myrs).
In
efwith
the
fact
that
cloudsoften
leave
the potential
wellatconfinedanis very actively forming stars between t = 9 11 Myrs (when it
and thus along
they areunsurprisingly
coherent in velocity. Wet, as the gas is accelerated towards
thebottom
of
the
angle
(i.e.
feelingdi↵erent
accelerations
theirlength),is at the bottom of the potential well), and then becomes morecaution,however,that
thisshouldnot be interpreted asral potential, the density contrastof theoverridgeof thefila-is much resultsin overall,the clouds
becoming more distorted, no longer
constant
time.
GMF1
less active
but does
proof thatare contrast
physically relevant structures. For innts increases, as a consequence of
thetogasfeelinga “frontal
imprintsof theirpreviousstatetheseas highseem
have
an increase
number of holding
SF eventsthe
towards
the endof
stance, we find that GMFs often break before they reachthe simulation,it also reachesgiant
the bottomof thefilaments.
potentialnd”, or ram pressure. This e↵ectively
works aswhen
a compresmolecular
the gaseous spiral arm, and the resulting smaller filamenat t =filaments,
17 19 Myrs.
The late SF events of GMF1, however,n force acting on the smaller axiswell,of thesewhilstIn terms
ofinchemicalas the GMFs
approachtary sections
maintain
their relative position and motion,
are not associated with the main GMF that
we track
Sect. 4.1, evolution,
galactic shear from the di↵erential rotation, continues to
but instead take place within some ofthespiral
smallerfragmentsthat mass particularly
thearm,
the totalremains relatively
up until constantthey cross the bottom of the potential
Molecular mass fraction

Molecular mass fraction

5

・GMFs内のガス

with the fact that clouds often leave the potential well at anangle (i.e. feeling di↵erent accelerations along their length),results in the clouds becoming more distorted, no longer
holding the imprints of their previous state as high contrast
giant molecular filaments.
In terms of chemical evolution, as the GMFs approach
the spiral arm, the total mass remains relatively constantover time (Fig. 5), while the increase in density on the ridgesof the filaments naturally increases the amount of both
molecular gas and CO (Fig. 6). Once in the gaseous spiral arm, the molecular mass fraction of the clouds still increases (Fig. 6), but the molecular fraction does not reachmore than ⇠60%. However, this limit is likely due to ourprescription of the feedback, which prevents gas from reaching very high densities, which is where we would expect it
to become mostly molecular.
Throughout their evolution, both GMF1 and GMF2
fragment into smaller clouds (e.g. see dashed lines on Figs. 5
and 6), which also alters their overall appearance. The reasons for clouds to break are a combination of stellar feed-

ガス状spiral armに入る前に星形成が起こっている

spiral armに入っても、星形成イベントは増えていない・GMFsはspiral armに沿う

銀河系でのGMFsの観測結果と一致