星形成ゼミ
2018/06/08
SFN #304 #49-#58
長谷川哲夫（国立天文台）

1

304 Mul$-epoch monitoring of the AA Tau like star V 354 Mon#49 Indica$ons for a low gas-to-dust ra$o in the inner disk warp

P. C. Schneider1; 2, C. F Manara2; 3, S. Facchini4, H. M. Günther5, G. J. Herczeg6, D. Fedele7, and P. S. Teixeira8; 9; 10
A&Ap in press観測： VLT X-Shooter（λ300 – 2500 nm分光測光）、連続3夜
結果：1) 輝線（Hα、HeI、Li）弱い→accreQonの変化ではなくoccultaQonの変化
2) SEDは大きく変化→warped inner diskのdustによる減光
3) 吸収線（NaとK）→大部分photosphericだがinner diskによる変化もある
Gas/dust比はinterstellarに比べて1桁以上低い
仮説1: inner diskはdebris diskに進化中
P. C. Schneider et al.: Multi-epoch monitoring of the AA Tau like star V 354 Mon
仮説2: outer diskから落ちてきた粒径の大きな個体の破砕

Fig. 3. Top: Observed spectral energy distributions. Bottom two panels: Absolute di↵erence between model and observed spectra. Models for thedimmer spectra (18th and 19th Jan.) assume that the di↵erence between the brightest spectrum (20th Jan.) and the dimmer spectra is caused (1)only by additional extinction allowing E(B V) and RV to vary (blue), and (2) our AA Tau model in red (see text for a detailed description). To
di↵erentiate between the two 18th Jan. spectra, the shorter one is displayed as a dashed line. Spectral ranges strongly a↵ected by telluric absorption
are removed from the fit and shown in lighter colors. The parameters corresponding to the blue and red curves in the bottom two panels are given
in Table 2.

3

Patrick D. Sheehan

and Josh A. Eisner

Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721, USA

304 MulQple
Gaps
inJulythe11;Diskof2018
theFebruary
Class I7;Protostar
GYFebruary
91 7; published 2018 April 9
Received
2017
revisedaccepted 2018
D. Sheehan and Josh A. Eisner#50 Patrick
ApJ 857, 18
Abstract

We present the highest spatial resolution Atacama Large Millimeter/submillimeter Array (ALMA) observations todate of the Class I protostar GY 91 in the ρ Ophiuchus L1688 molecular cloud complex. Our 870 μm and 3 mmdust continuum maps show that the GY 91 disk has a radius of ∼80 au, and an inclination of ∼40°, but mostinterestingly that the disk has three dark lanes located at 10, 40, and 70 au. We model these features assuming they
The Astrophysical Journal, 857:18 (7pp), 2018 April 10& Eisneraregaps in the disk surface density proﬁle and ﬁnd that their widths are 7, 30, Sheehan
and 10au. These gaps bear a striking
resemblance to the gaps seen in the HL Tau disk, suggesting that there may be Saturn-mass planets hiding in the
disk. To constrain the relative ages of GY 91 and HL Tau, we also model the disk and envelope of HL Tau and ﬁndthat they are of similar ages, although GY 91may be younger. Although snow lines and magnetic dead zones canalso produce dark lanes, if planets are indeed carving these gaps then Saturn-mass planets must form within theﬁrst ∼0.5 Myr of the lifetime of protoplanetary disks.
Key words: planets and satellites: formation – protoplanetary disks – stars: formation – stars: individual (GY 91)(n5–12=−0.25). However, their measured value is athe border between disks with foreground extinction andnets form in protoplanetary disks. When observed at highwith envelopes (of n5–12=−0.2), and the extinction-cor
ion with Atacama Large Millimeter/submillimeter Array
spectral index (a ¢IR =0.31) and bolometric temp
A),Figurea number
these
disks
show
interesting
patterns
1. Our ALMAof
345 GHz(left) and
100 GHz(right) mapsof the GY 91 protoplanetary
disk. inTwo dark lanes are readily apparent in the 345 GHz map, while a¢ =470(TGHzK) still qualify the source as a Class I pr
third dark lane is also apparent in inner regions of the disk at 100 GHz because of the higher resolution of our 100
maps.
bol
millimeter emission proﬁles, in some cases including(Dunham et al. 2015). van Kempen et al. (2009) alsoof bright and dark rings (ALMA PartnershipTable
et 1al. 2015;
HCO+ emission toward GY 91 that was bright enoughLog of ALMA Observations
et al. 2015; Andrews et al. 2016; Isella et al. 2016; Fedele
cutoff for an embedded source, but that em
Observation Date
ALMA Band
Baselines
Integration Time above theCalibrators
2017a,
2017b; Loomis
et al. 2017; Dipierro
et al. Total2018).
(UT)(m)(s)(Flux,
Bandpass,
Gain)
seems to be associatedwith a patch of cloud that peaof these
disks
may
even
be
quite
young
(1
Myr;
ALMA
2015 Oct. 31
3
84–16, 200
169
1517–2422, 1625–2527
away fromthe source.
2015 Nov. 26
3
68–14, 300
169
1517–2422, 1625–2527
rship
et
al.
2015;
Dong
et
al.
2015;
Fedeleetal.
2017b;
2016 Apr. 17
3
15–60058
1733–1304,
1427–4206,new1625–2527
HereJ1517–2422,
we present
ALMA data, which, when com
May 19730
J1625–2527
ro 2016et al.
2018). There are
a number15–640
of explanations forwith theJ1517–2422,
observed
SED, show that GY 91 is indeed a C
2016 Sep. 11
7
15–314060
J1625–2527
eatures, including chemical processes that alter dust
protostar with a circumstellar disk embedded in4 an env
es wereanddevoted
sticking/fracturing
processes
near snow
et al. lines
2002; Natta et al.
2006;
Stankeandet al.870
2006;μm
Alves
deto spectral line observations,although nothing
Our
3 mmimagesalso reveal the presence oOliveira & Casali 2008; Jørgensen et al. 2008; Padgett
was detected. We list details of the observations in Table 1.
1. Introduction

Table 1
Log of ALMA Observations
Observation Date
(UT)2015 Oct. 31
2015 Nov. 26
2016 Apr. 17
2016 May 19
2016 Sep. 11

ALMA Band

Baselines(m)

Total Integration Time
(s)

3
3
3
7
7

84–16, 200
68–14, 300
15–600
15–640
15–3140

169
169
58
30
60

were devoted to spectral line observations, although nothing
was detected. We list details of the observations in Table 1.
The data were reduced in the standard way with the CASA
pipeline and the calibrators listed in Table 1. After calibrating, we
imaged the data by Fourier transforming the visibilities with the
CLEAN routine. After our initial imaging, we found that we could
improve the sensitivity of the 345 GHz image by self-calibrating.
We ran four iterations of phase-only self-calibration on thecompact conﬁguration track and a single iteration of phase-only
self-calibration on the extended conﬁguration track. This
improved the rms in an image produced with natural weighting(i.e., a robust parameter of 2) from 0.36 to 0.27 mJy beam−1. We
were unable to improve the 100 GHz image by self-calibration.
Our ﬁnal images were produced using Briggs weighting with arobust parameter of 0.5, which provides a good balance between
sensitivity and resolution, to weight the visibilities for both data
sets. The 3 mm image has a beam of size 0 06 by 0 05 with a P.
A. of 81°. 9 and an rms of 36 μJy beam−1. The 870 μm image has

Calibrators(Flux, Bandpass, Gain)
1517–2422, 1625–2527
1517–2422, 1625–2527
1733–1304, 1427–4206, 1625–2527
J1517–2422, J1625–2527
J1517–2422, J1625–2527

et al. 2002; Natta et al. 2006; Stanke et al. 2006; Alves deOliveira & Casali 2008; Jørgensen et al. 2008; Padgett
et al. 2008; Wilking et al. 2008; Evans et al. 2009; Gutermuth
et al. 2009; Barsony et al. 2012). When modeling the SED, as
we discuss below, we assume a constant 10% uncertainty on
any photometry from the literature to account for any ﬂux
calibration uncertainties between the measurements.
In addition to the broadband photometry, we also download
the Spitzer IRS spectrum of GY 91 from the CASSIS database(Lebouteiller et al. 2011, 2015). Rather than consider the entire
SED, which can be computationally prohibitive for the radiativetransfer calculations described below, we sample the IRS
spectrum at 25 points ranging from 5 to 35 μm. We also assume
a 10% uncertainty on these ﬂuxes, like we do for the broadband
photometry.
3. Results

5

Johansen 2013; Banzatti et al. 2015). The change in the opticalproperties of dust grains across the snow line could causefeatures like those seen in HL Tau or GY 91 (e.g., Zhang
et al. 2015).
Alternatively, the “sintering” of dust grains just below thesublimation temperature produces brittle grains that fragment
more readily and therefore grow to smaller sizes just outside of
the snow line. Because the sintered grains have smaller sizes,
they undergo slower radial drift, causing pile-ups near snowlines. This process could also produce features similar to thoseseen in HL Tau or GY 91 (Okuzumi et al. 2016).
We compare the midplane disk temperature inferred for GY
91 with the temperatures of snow lines of common volatiles(e.g., Zhang et al. 2015). The outermost gap does roughlymatch the freeze-out region of N2. However, no obvious
counterparts are seen for the inner two gaps.
Zonal ﬂows produced by magneto-rotational instability
driven turbulence (Johansen et al. 2009) have also been shown
to produce axisymmetric pressure bumps that can trap large
dust grains and may produce gap-like features in millimeter
images (Pinilla et al. 2012; Dittrich et al. 2013; Simon &Armitage 2014). In this case, the pressure bumps are created bylarge-scale variations in the magnetically driven turbulence that
produce variations in the mass accretion rate that in turn causesmaterial to pile up. This effect can also be seen at the outeredge of magnetic dead zones, where there is strong radial
variation in the mass accretion rate. These ﬂows can produce
gap-like features in disks (Pinilla et al. 2012; Flock et al. 2015).

6

ESPOL Polytechnic University, Escuela Superior Politécnica del Litoral (ESPOL), Physics Department, FCNM, Campus Gustavo Galindo km. 30.5
Vía Perimetral, P.O. Box 09-01-5863, Guayaquil, Ecuador
Received2018 March 1; accepted2018 March 8; published
AprilBinary
11
304 From Large-scale to Protostellar
Disk FragmentaQon
into2018
Close
Stars
Di G. Sigaloi1, Fidel Cruz1 , Ruslan Gabbasov2, Jaime Klapp3,4 , and José Ramírez-Velasquez5,6#51 Leonardo
ApJ 857, 40
Abstract

Recent observations of young stellar systems with the Atacama Large Millimeter/submillimeter Array (ALMA)
and the Karl G. Jansky Very Large Array are helping to cement the idea that close companion stars form via
fragmentation of a gravitationally unstable disk around a protostar early in the star formation process. As the disk
grows in mass, it eventually becomes gravitationally unstable and fragments, forming one or more new protostars
in orbit with the ﬁrst at mean separations of 100 au or even less. Here, we report direct numerical calculations down
to scales as small as ∼0.1 au, using a consistent Smoothed Particle Hydrodynamics code, that show the large-scale
fragmentation of a cloud core into two protostars accompanied by small-scale fragmentation of their circumstellar
disks. Our results demonstrate the two dominant mechanisms of star formation, where the disk forming around a
protostar (which in turn results from the large-scale fragmentation of the cloud core) undergoes eccentric (m=1)
fragmentation to produce a close binary. We generate two-dimensional emission maps and simulated ALMA
1.3 mm continuum images of the structure and fragmentation of the disks that can help explain the dynamical
processes occurring within collapsing cloud cores.
Key words: circumstellar matter – methods: numerical – planets and satellites: formation – protoplanetary disks –stars: formation – stars: pre-main sequence
1. Introduction

lar systems, consisting of two or more stars, are formedo different mechanisms—large-scale fragmentation ofcores of gas and dust during their early isothermal
se phase (Larson 2001; Fisher 2004), leading to widely
ted companion protostars (∼1000 au or larger), andscale fragmentation of their circumstellar disks early inollapse due to gravitational instabilities (Rodríguez
2005; Connelley et al. 2008; Pech et al. 2010; Murillo
2013; Tobin et al. 2013), leading to bound systems
ted by tens of au. This scenario of star formation seems

lack physical reality because protostellar disks form andin mass within a collapsing environment as a result of theof material from the outer cloud envelope (Eisner 2012;
et al. 2012). On the other hand, if more protostars form w
the same environment, the gravitational interaction athem may limit the disk size and affect the way anmomentum is transferred. However, the disk fragmen
model has succeeded in reproducing the constraints impos
the observed statistical properties of low-mass objects, su
low-mass binary systems, brown dwarfs, and even planmass objects, which in turn have not yet been7 explain
means of other formation mechanisms (Stamatellos & Whi

The Astrophysical Journal, 857:40 (11pp), 2018 April 10

Sigalotti et al.

Figure 1. Equatorial density evolution in a 266 au region around the cloud center for model U2 after large-scale fragmentation into a binary protostar: (a) 44.14 kyr,(b) 44.39 kyr, (c) 44.43 kyr, and (d) 44.68 kyr. The logarithm of the density ρ is plotted. For (a)–(d), the maximum density is: log10ρmax=−10.86, −10.63, −10.62,
and −10.39 (ρmax in g cm−3), respectively. The insets show a magniﬁcation of the binary cores and their whirling disks. The disks develop a two-armed spiral patternassociated with a linear growth of a gravitational instability. In (b) and (c), the gravitational instability in disk B enters a nonlinear growth phase and so one of its spiralarms is seen to fragment into a second protostar.

8

9

we calculate the temperature for all particles by combining thebarotropic pressure–density relation with the ideal gas law to give⎡⎛ r ⎞g- 1⎤
T = T0 ⎢1 + g ⎜ ⎟ ⎥ ,⎢⎣⎝ rc ⎠ ⎥⎦(4 )where T0=10 K, rc = 5.0 ´ 10-12 g cm−3, and γ=5/3. At
temperatures below 100 K, a value of γ=5/3 is appropriate
because the rotational and vibrational degrees of freedom ofmolecular hydrogen are frozen out, and so only translational
degrees of freedom need be considered (Boss et al. 2000).
However, precise knowledge of the dependence of temperature
on density at the transition from isothermal to nonisothermal
collapse will require solving the radiative transfer problem
coupled to a fully self-consistent energy equation. Thistransition was studied for the collapse of an initially centrally
condensed core obeying a spherically symmetric Gaussian
density proﬁle, using nonisothermal thermodynamics andsolving the mean intensity equation in the Eddington approximation with detailed equations of state (Boss et al. 2000). The

collapse was seen to remain strictly isothermal up toρ∼10−16 g cm−3. At higher densities, the collapse was nearisothermal, with the temperature rising very slowly from 10 K
to ∼11.2 K when ρ∼10−14 g cm−3. At densities higher thanthis, the temperature was seen to rise sharply, with the density
and the collapse becoming nonisothermal. Therefore, a value ofthe critical density rc = 1.0 ´ 10-14 g cm−3 would be more
representative of the near isothermal phase. However, a valueof rc = 5.0 ´ 10-12 g cm−3, while prolonging the isothermal
phase of collapse to higher densities, makes convergence
more difﬁcult to attain and resolution requirements signiﬁcantly
more demanding. Larger and more massive disks than those
calculated here can therefore be obtained by anticipating thenear isothermal phase using lower transition densities inbarotropic collapse calculations. Figure 5 shows the projection
on the disk midplane of the disk particle densities (normalized
to the critical density rc = 5 ´ 10-12 g cm−3) and temperatures(calculated using Equation (4)) as functions of radius for disks
A and B at ≈44.39 kyr. The prominent peaks in density andtemperature at R=0 represent the central protostars, while

10ALMA 1.3mm観測
シミュレーション
The Astrophysical Journal, 857:40 (11pp), 2018 April 10

Sigalotti et al.

Figure 7. ALMA 1.3 mm continuum images of (a) disk A at 44.97 kyr, (b) disk B at 44.47 kyr, both at a viewing angle of 0◦, and (c) disk A at 44.97 kyr, (d) disk B at
44.47 kyr, both at a viewing angle of 45◦. The side of the disks nearest the observer is oriented toward the bottom of each frame. The source distance is assumed to be
at 100 pc from the observer. The angular resolution of each image is shown with the blue circle drawn in the lower left corner, corresponding to a synthetic beam sizeof 0″.05 (9 au at this distance). The simulation images were produced with an integration time of 1 hour. The color bar and numbers on the right indicate the ﬂux
intensity (in mJy beam−1).

at a distance of ≈15 au from the primary protostar at about
44.38 kyr, reaching a maximum separation of ≈19 au by
44.43 kyr. During this stage, the fragment accretes mass fromthe disk at a high rate as shown by the steep increase of MGI,
losing orbital angular momentum during this stage andmigrating to a minimum separation of ≈10 au by ≈44.6 kyrwhen it has already completed half of an orbital period.
Thereafter, as the disk dissipates, the accretion rate declines andthe secondary moves away in its orbit until it reaches amaximum separation of ≈18.2 au at ≈44.76 kyr by the time it
has completed a full revolution about the primary. The
secondary continues in its elliptic orbit and completes a second
period by ≈45 kyr when the calculation is terminated. At thistime, the mass of the secondary has become comparable to theprimary. During this second orbital period, the minimum andmaximum distances from the primary are ≈13.4 and ≈18.2 au,
respectively. This gives an orbital period of ≈225 years, whichcorresponds to a much shorter timescale compared to the Class0 lifetime of ∼30 kyr (André et al. 2000). The eccentricity was
found to decrease from ≈0.36 for the ﬁrst period to ≈0.151 forthe second period, suggesting that the secondary is approachinga more circular orbit as it comes close to the primary. However,sense, it is useful to simulate ALMA images to help interpretthe observations. Figure 7 shows simulated ALMA images ofdisks A and B at viewing angles of 0◦ and 45◦ and at two
different times during the evolution. These correspond to full
resolution images obtained by transforming the 2D emission
maps of Figure 6 with the aid of the simobserve andsimanalyze tools under Common Astronomy Software
Application for a wavelength of 1.3 mm (230 GHz; ALMA
band 6, conﬁguration C43-8 consisting of a full array of 50
antennas of size 12 m). A reference distance of 100 pc fromthe source was assumed. These images bear a remarkable
resemblance to the ALMA images of the Class 0 triple
protostar L1448 IRS3 (consisting of a tight binary and a more
widely separated tertiary) as observed when the tertiary, which
has the largest peak intensity in the system, is removed (Tobin
et al. 2016b). However, compared to this protostellar system,
the length scale in our results for disk B is smaller by a factorof ∼4.
3.4. Orbital Evolution of Second Protostar in Disk B
The orbital evolution of the second protostar around theprimary in disk B is shown in Figure 8. The secondary emerges
9

Figure 6. (a) Gas surface density (top row) and midplane temperature (bottom row) of disk A viewed face-on at four different times. (b) The same as (a) but for disk B.
By ≈44.97 kyr, when the calculation was terminated, disk A shows signs of fragmentation into a second protostar. The logarithm of the surface density Σ (in g cm−2)and the temperature T (in K) are plotted. Each row of images uses the color bars on their right sides. In all images, the central protostar is drawn with a black (surface
density) and dark blue (temperature) circle to increase the surface density and temperature contrasts.

11

304 Dynamical evolu$on of stars and gas of young embedded stellar sub-clusters
Sills,1‹ Steven Rieder,2 Jennifer Scora,1 Jessica McCloskey1 and Sarah Jaﬀa3#52 Alison
MNRAS 477, 1903

12

Figure 1. Initial conditions for DR21. X- and Y-axes correspond to RA and FigureFigure
3. Initialconditionsfor SeeNGC
6357.
Fig. 1 for aThedescription.
2. Initialconditionsfor M17.
Fig.
1 forSeea description.
lowest The
Dec., respectively. Stars are shown as blueish-white objects; gas is included contour
lowestlevel corresponds
a surfacedensity103
M⊙eachpc−2 , andlevelcontour
corresponds
to a surface todensityof 167
M⊙of
pc−2
, and
0.5
0.5
as diffuse emission, reflecting light from nearby bright stars. Gas is also subsequent
each subsequent
lineofis 10
a factorof 10 higher.
line is a factor
higher.
indicated with contour lines. The lowest contour level corresponds to a−2
surface density of 150 M⊙ pc , and each subsequent line is a factor of plane of the sky). The stars are represented as objects with intensity
100.5 higher. Figure created using FRESCO.1
and colour related to their mass and radius. Gas is rendered as

diffuse emission, and its smoothed surface density is also shownwith contour lines.
In Figs 4 and 5, we show the positions of stars in our model ofDR21 in the x–y plane (equivalent to the RA–Dec. plane of the sky).
Observed stars are shown as red circles, and the inferred low-massstars are shown as small black points.

Figure 4. Initial stellar positions for the DR21 simulations. Observed MYStIX stars are shown as red circles, and the inferred population of low-mass

13

describes the range of scales over which the fractal sub-structure
exists, with a low value indicating that the smallest sub-structures
are not much smaller than the overall cloud and a high value indicating that the smallest structures are very small compared to thecluster size. The density scaling exponent describes where most of
the mass is distributed in the hierarchy, with a high value indicatingthat most of the mass is confined to the smallest sub-structures anda low value indicating that the sub-clumps are not much more dense
than the larger clumps. Clusters that are not sub-structured but centrally condensed can be characterized by a radial density exponent
using the method described in Cartwright & Whitworth (2004) thatis incorporated into the Q+ analysis code.
What we found was that the fractal dimension of the DR21 fiducial simulation was consistently 1.58 (quite clumpy) until the subclusters merged at around 0.5 Myr. However, during that time, thenumber of levels decreases steadily and the density scaling exponent
increases, suggesting that the stars from the smallest sub-clump(s)are spreading out to fill the larger hierarchy. After the clumps merge
at ∼0.5 Myr, the region completely loses its fractal structure and isconsistent with a centrally concentrated system. The central stellar
density increases with time, as does the radial density gradient.
If we limit our analysis to only the high-mass stars (M >0.83 M⊙ ) that were taken from the MYStIX observations and exclude the low-mass stars added for completeness, we can compare
the structure that would be theoretically seen by an observer to thefull structure we have in the simulations. These results are shown in
Fig. 8. For the initial conditions of DR21, this mass limited analysis
gave a lower fractal dimension and density scaling exponent thananalysis of the full cluster but a similar number of levels. This suggests that the observable higher mass stars populate a more clumpy
distribution, but are not confined to the smallest sub-clumps. Theirevolution showed the same trend as the full cluster. The high-massstars in M17 and NGC 6357 exhibit similar behaviour; the initialconditions have a low fractal dimension and high number of levels,
but the sub-structure is quickly erased (within 1–2 Myr) and thecluster becomes centrally concentrated. It is interesting to note thatwhen the full clusters are analysed, M17 would not be classified as
fractally sub-structured and NGC 6357 only just so. It is only when
we analyse the high-mass component of these two clusters that the
sub-structure is detected. However, this difference in structure isambiguous and may be caused by limitations in the Q+ analysis
rather than an actual difference in structure between the high-mass
and low-mass components.

←10万年毎のスナップショット（5pc角の領域を示す）
By comparing our results to observaQons from the MYStIX
regions, we conclude that;
1. Our models match the observaQons well, and so theiniQal condiQons and choices of a priori unknown
parameters we used are realisQc and reasonable.
2. The subsequent evoluQon of these regions is quite
insensiQve to reasonable values of the parametersthat we varied.

Figure 7. Snapshots of the first nine snapshots of the evolution of DR21
in our fiducial simulation, from top left to lower right. The snapshots aretaken 100 000 yr apart, with the first snapshot 100 000 yr after the simulation
begins. Each snapshot is 5 pc on a side, otherwise as in Fig. 1. The subclusters have almost entirely merged
withinthe first million years ofclusterYoung
sub-clusters
1909
evolution, and the gas density has become quite spherical.

α vir = 2T/|W|. Our simulations have an additional gravitational
potential from the gas mass, which effectively decreases the virial
parameter and shortens the time for collapse.
3.1 Sub-structure
To characterize the sub-structure seen in our simulations, we used
the Q+ algorithm (Jaffa, Whitworth & Lomax 2017), which is not
related to the virial parameter but rather characterizes the fractal distribution of the stars. The Q+ algorithm describes a sub-structured
star cluster in terms of three parameters: the fractal dimension, thenumber of levels, and the density scaling exponent. The fractal dimension describes how clumpy or smooth the distribution is, with
1 being very clumpy and 3 being smooth. The number of levels
MNRAS 477, 1903–1912 (2018)

ed from https://academic.oup.com/mnras/article-abstract/477/2/1903/4939295
nal AstronomicalObservatory, Japan usery 2018

We found that our iniQally very sub-structured systems
became spherical, monolithic, and smooth very quickly,
in both their spaQal and velocity distribuQons.
EvoluQon of embedded clusters at young ages appears to
be dominated by gravitaQonal interacQons between thestars, with only a small inﬂuence from the surrounding gas.

14018 March 7; in original form 2017 November 21

304 The diverse lives of massive protoplanets in self-gravita$ng discs
Stamatellos1‹ and Shu-ichiro Inutsuka2#53 Dimitris
MNRAS 477, 3110

ABSTRACTGas giant planets may form early on during the evolution of protostellar discs, while these are
relatively massive. We study how Jupiter-mass planet-seeds (termed protoplanets) evolve in
massive, but gravitationally stable (Q ! 1.5), discs using radiative hydrodynamic simulations.
We find that the protoplanet initially migrates inwards rapidly, until it opens up a gap in thedisc. Thereafter, it either continues to migrate inwards on a much longer time-scale or starts
migrating outwards. Outward migration occurs when the protoplanet resides within a gap withgravitationally unstable edges, as a high fraction of the accreted gas is high angular momentum
gas from outside the protoplanet’s orbit. The effect of radiative heating from the protoplanetis critical in determining the direction of the migration and the eccentricity of the protoplanet.
Gap opening is facilitated by efficient cooling that may not be captured by the commonly used
β-cooling approximation. The protoplanet initially accretes at a high rate (∼10−3 MJ yr−1 ),and its accretion luminosity could be a few tenths of the host star’s luminosity, making theprotoplanet easily observable (albeit only for a short time). Due to the high gas accretionrate, the protoplanet generally grows above the deuterium-burning mass-limit. Protoplanet
radiative feedback reduces its mass growth so that its final mass is near the brown dwarfplanet boundary. The fate of a young planet-seed is diverse and could vary from a gas giant
planet on a circular orbit at a few au from the central star to a brown dwarf on an eccentric,
wide orbit.
Key words: accretion, accretion discs – hydrodynamics – protoplanetary discs – browndwarfs – planetary systems.

last 20 yr have made possible thebe long when compared to the observed lifetimes of circumstellar discs (of the order of a few Myr; Haisch, Lada & Lada 2001;
Hernández et al. 2008; Muzerolle et al. 2010). Moreover, this model

15

Table 1. Simulation parameters.
Run idOpacity

a
αmin

b
αp,i
(au)

c
Mp,f
( MJ )

d
αp,f
(au)

efe

Run 1
Run 2
Run 3
Run 4
Run 5

BL94f
BL94f
BL94f
BL94f
SEM03g

0.1
0.1
0.2
0.1
0.1

40
40
40
40
40

24
15
29
25
31

40
17
37
36
53

0.035
0.0012
0.13
0.017
0.17

Notes. a Minimum SPH artificial viscosity. b Initial semimajor axis of the protoplanet. c Final protoplanet mass. d Final semimajor axis of the protoplanet.
e Final protoplanet eccentricity. f Bell & Lin (1994) opacity. g Semenov et al.
(2003)
3116opacity.D. Stamatellos and S.-i. Inutsuka

studies (Nayaksh& Montesinos 20
The way thatthe dynamical sta
difference betwee
protoplanet’s Hil
and the angular moutside the protop
mass of gas withthat moves towar

1 The same result is

the gas within dista

MNRAS 477, 3110–3135 (2018)

Downloaded from https://academic.oup.com/mnras/article-abstract/477/3/3110/4961145
by National Astronomical Observatory of Japan user
on 25 May 2018

Figure 2. The evolution of the disc surface density in the five simulations listed in Table 1. The star and the protoplanet as depicted as wide dots. Each row
corresponds to different snapshots of each run at times as stated on the graph. All runs show similar features, apart from the run in which the radiative feedbackof the protoplanet is taken into account (Run 2; second row).

16

Figure 18. The evolution of a protoplanet in a 0.1-M⊙ protostellar disc. The protoplanet is placed at different initial orbital radii within the disc: 5, 10, 20,
40, and 80 au (top to bottom row; sRun1, . . . , sRun5, respectively). There is a difference in the migration pattern depending on the position of the protoplanetin the disc.

17

3048; inTheforma$onof rings
18 March
original
form 2017 December
17 and gaps in magne$cally coupled disc-wind systems:#54 ambipolar diﬀusion and reconnec$on

Scoq S. Suriano,1‹ Zhi-Yun Li,1 Ruben Krasnopolsky2 and Hsien Shang2
MNRAS 477, 1239

ABSTRACT

Radial substructures in circumstellar discs are now routinely observed by Atacama Large
Millimeter/submillimeter Array. There is also growing evidence that disc winds drive accretionin such discs. We show through 2D (axisymmetric) simulations that rings and gaps develop
naturally in magnetically coupled disc-wind systems on the scale of tens of au, where ambipolar
diffusion (AD) is the dominant non-ideal magnetohydrodynamic effect. In simulations where
the magnetic field and matter are moderately coupled, the disc remains relatively laminar withthe radial electric current steepened by AD into a thin layer near the mid-plane. The toroidal
magnetic field sharply reverses polarity in this layer, generating a large magnetic torque thatdrives fast accretion, which drags the poloidal field into a highly pinched radial configuration.
The reconnection of this pinched field creates magnetic loops where the net poloidal magnetic
flux (and thus the accretion rate) is reduced, yielding dense rings. Neighbouring regions withstronger poloidal magnetic fields accrete faster, forming gaps. In better magnetically coupledsimulations, the so-called avalanche accretion streams develop continuously near the discsurface, rendering the disc-wind system more chaotic. Nevertheless, prominent rings and gaps
are still produced, at least in part, by reconnection, which again enables the segregation of
the poloidal field and the disc material similar to the more diffusive discs. However, thereconnection is now driven by the non-linear growth of magnetorotational instability channel
flows. The formation of rings and gaps in rapidly accreting yet laminar discs has interesting
implications for dust settling and trapping, grain growth, and planet formation.
Key words: accretion, accretion discs – MHD – protoplanetary discs – ISM: jets and outflows.

18oh0.26
oh2.6
oh26
beta1e2
beta1e4

0.922
0.922
0.922
0.0922
9.22

4.408
4.408
4.408
4.408
4.408

0.25
0.25
0.25
0.25
0.25

9.71
9.71
9.71
97.1
0.971

2.5
25
250
–
–

Figure 1. A representative (‘reference’) axisymmetric simulation. Shown are the mass volume density (logarithmically spaced colour contours in units ofg cm−3 ), the poloidal magnetic field lines (magenta), and the poloidal velocity unit vectors (black). Panels (a)–(d) corresponding to simulation times of 0, 200,
1000, and 2500 inner orbital periods, respectively. (See the supplementary material in the online journal for an animated version of this figure.)

3 REFERENCE MODEL AND THE
F O R M AT I O N O F R I N G S A N D G A P S T H RO U G HRECONNECTION(left to right), where t0 = 1 yr is the orbital period at the inner edgeof the simulation domain (r0 = 1 au). The disc-wind launching
19 from
proceeds in an inside-out fashion, i.e. the wind is launchedlarger disc radii as the simulation progresses. By 1000t0 (one orbit

20

1248

S. S. Suriano etRing
al. and gap formation

1247

compared to its poloidal component. The now weaker poloidal fieldstrength can be seen clearly in the last panel from the lack of fieldlines in the post-reconnection region.
3.3.3 Reconnection-driven ring and gap formation
The general picture for the reconnection-driven ring and gap formation is as follows. A poloidal magnetic field line that initiallythreads the disc rather smoothly is dragged by preferential accretion near the mid-plane into a highly pinched radial configuration(see the upper-right panel of Fig. 7 for an illustration). Reconnection of the highly pinched field line produces a poloidal magnetic loop next to a poloidal field line that still threads the discmore smoothly (see the lower-left panel of Fig. 7). After reconnection, the material trapped on the magnetic loop is detached fromthe original (pre-reconnection) poloidal field line, which resultsin the separation of matter and (poloidal) magnetic flux. SpecifFigure 7. Poloidal magnetic field lines at four different times are shown in grey, with a reconnecting field line highlighted in black. The colours show the ratio
ically, there is no net poloidal magnetic flux passing through theof the toroidal to poloidal magnetic field, Bφ /Bp .
matter on the loop, and there is less matter remaining on the original poloidal field line (since some of the matter on the original field
diffusivity, such as resistivity), prominent avalanche streams would
3.3.2 Reconnection of pinched radial magnetic field
is now on the detached magnetic loop). The net effect is a redishave developed near the disc surface, which would driveline
the entire
The secondary component of the mid-plane current sheet develops
disc and its envelope into an unsteady state and make tribution
the forma- of the poloidal magnetic flux away from the loop-forming
as a result of mass accretion in the disc, which is concentrated intion of a laminar mid-plane current sheet impossible (seePaperintoI
regionanadjacent
regionwhere
the poloidal
flux accumulates.
near the mid-plane(see Fig. 6).
theprimary radialcurrent(Jr ) sheet
and Section 4.1 below).1 Secondly, as first stressed by BrandenThe mass accretionisdrivenby the removalof angularmomentumSince
mass
accretion
tends
to
be
fasterinregionswithstrongerburg & Zweibel (1994), AD tends to steepen the magnetic gradientwhich
due to the Lorentz force (∝Jr Bθ ) in the azimuthal
poloidalthat thedirection,
reconnectionnear a magnetic null, i.e. where the magnetic field changes
polarity. magnetic fields, it is not surprisingpeaks in the radial current sheet where the toroidal magnetic field
The reason is that the Lorentz force associated with theinduced
magnetic variation of the poloidal field strength with radius wouldchanges sign. Pictorially, as the φ component of the magnetic fieldpressure gradient drives the ions (relative to the neutrals) towards
lead to a spatial
variation
massbelow
accretion
would inchangesfrom thein+φthedirection
the discrate
to thethat−φ direction
the null from both sides. Since the field lines (of opposite polarabove thedisc in a thin
mid-plane
the magnetic
fieldi.e.
linestheturnlead to spatial
variationin themass layer,(surface)
density,
ity across the null) are tied to the ions, they are dragged
towards
become severely kinked
in the azimuthal direction. The sharp field3
the null as well, leading to a steepening of the magnetic
gradiformation of ringsand gaps.
kink generates a magnetic tension force in the −φ direction thatent, which in turn yields a stronger magnetic force that drives the
The preferential
poloidallines
insideexerts a concentration
braking torque andof
drivesthe discfieldaccretionin thecurrenttheions and the field lines even closer to the null. Since the ambipolar
sheet. of the disc is illustrated pictorially in Fig. 8(a),low-density
gaps
magnetic diffusivity, ηA , is proportional to the field strength (seeThe accretion through the mid-plane current sheet now drags theequation 5) and thus vanishes at the null, this steepening
would
which
is a snapshotof the inner part of the reference simulation
poloidal magnetic field lines inwards with the flow. This results inresult in an infinitely thin, singular, current sheet in principle.
In(up to 15 au)a pronounced
at a representative
time
when magnetic
the simulationhasradial pinch of
the poloidal
field, which
practice, the thickness of the current sheet is limited by the gridtransports thestate,poloidalflux The
inwards
and concentration
yields anotherfieldreached a quasi-steadyt/tmagnetic
0 = 2500.
resolution.
in the azimuthal direction, Jφ . Eventually, theis quantified
incurrent
panelcomponent(b), whichplots the vertical component of the
Sharp magnetic field reversals that give rise to thincurrent
radial pinch becomes so severe that the magnetic field reconnects,
sheets are prone to reconnection. However, this is not thecaseforred) atthe mid-plane
andaresurface
densitymagnetic fieldforming(Bz , inpoloidal
magnetic
field loops that
reminiscentof the ofthe mid-plane current sheet shown in Fig. 6, because it is protearing bothmode relative
instabilityto(Furthal. 1963;
see also&
the disc (in black),
theiretinitialvaluesat tZweibel= 0. Two
duced by the reversal of the toroidal field component Bφ , andYamada 2009). An example of the reconnection process is shown
featuresit is impossible to reconnect oppositely directing toroidal
fields stand out: (1) there is a strong anti-correlation between
in Fig. 7, which plots lines of constant poloidal magnetic flux, i.e.
curunder the adopted axisymmetry.2 Nevertheless, this primary
the surface
densityverticalfieldstrength,
as expected
in tothepoloidal andmagnetic
field lines,
alongwith the ratioof the toroidal
rent sheet leads to another current component that does allow
reconnection-induced
scenario
ring
and
gap
formation,
and(2)
contours),
at a radius centred
poloidal magnetic
field, ofBφ /Bp (colour
for reconnection.
230. At
the first frame
shown,
on r = 6.5 aufieldfrom time
t/t0 = 200 to
the poloidal magnetic
is typicallymuchweakerin thehighthe poloidal magnetic field highlighted in bold already has a kinkclose
to zero) than in the lowdensity ring regions
z ispinch
across the (where
mid-plane.BThegrows with time until the field lineisincreasedoveritsapproximately
initial value
densitygap
regions(whereBisstretchedalmostparalleltothemid-planeoverfourbyz
1 In this reference run, AD does not appear to suppress the development ofradial grid cells at t/t = 210. By the next frame shown at t/t0 = 220,
factorflux
the MRI completely. Channel-like features are evident at largearadii
whereof 2–4). This drastic 0segregation of poloidal magnetic
the field has reconnected, forming a loop near r = 6.4 au. In thethe disc is better coupled to the magnetic field as measured by relative
the Elsasserto matter
appears
to
behas
quite
stable withtimethe(persistingforlastframe,
theloop
disappeared;
however,
processof
its
number (see Fig. 1c).
of
innerorbits),despite
the
fact
that
there
is
continued
2 Reconnection of the highly pinched toroidal field is expectedthousands
formation
has
left
a
lastingmarkonthemagnetic
field
structure.
in 3D, andTheregion now
has athemuchlargerand
toroidalmagnetic field component
will be explored in a future investigation.
mass accretion
through
both
rings
gaps.

Figure 8. The density, magnetic field strength, and velocities plotted up tor = 15 au at t/t0 = 2500. The panels show (a) the logarithm of the density(g cm−3 ) in colour and the poloidal magnetic field lines (i.e. magnetic flux
contours) in black, (b) the surface density normalized to the initial radial
distribution, ! i ∝ r−1/2 , and the vertical magnetic field strength at the
mid-plane normalized to its initial distribution, Bz, i ∝ r−5/4 , (c) the radial
velocity (km s−1 ) of neutrals (black) and ions (red) at the mid-plane, and(d) the density-weighted vertical average of the radial velocity (km s−1 )of neutrals (black) and ions (red) over z = ±2h. (See the supplementarymaterial in the online journal for an animated version of this figure.)

presence of significant AD, this is not necessarily the case. Indeed,at the time shown in Fig. 8, the ions are moving outwards in severalof the gaps, especially the two near r = 5 and 6.5 au. The ionsare forced to expand against the infalling neutrals by an outward
Lorentz force due to either a temporary poloidal flux concentrationnear the inner edge of the gap [in a manner that is reminiscent
21ofthe forced ion–neutral separation in the AD shock in magnetizedReceived September ....; accepted ...

304 Unveiling the remarkable photodissocia$on region of M8
Tiwari1,⋆,, K. M. Menten1, F. Wyrowski1, J.P. Pérez-Beaupuits1, 2, H. Wiesemeyer1, R. Güsten1, B. Klein1, C. Henkel1, 3#55 M.
A&Ap in press
ABSTRACT

Aims. Messier 8 is one of the brightest HII regions in the sky. We collected an extensive dataset comprising multiple submillimeter spectral lines from neutral and ionized carbon and from CO. Based on it, we aim to understand M8’s morphology and that ofits associated photo dissociation region and to carry out a quantitative analysis of the regions’ physical conditions such as kinetic
temperatures and volume densities.
Methods. Using the Stratospheric Observatory For Infrared Astronomy (SOFIA), the Atacama Pathfinder Experiment (APEX) 12 m
and the Institut de Radioastronomie Millimétrique (IRAM) 30 m telescopes, we have performed a comprehensive imaging survey ofthe emission from the fine structure lines of [C II] and [C I] and multiple rotational transitions of carbon monoxide (CO) isotopologues
within 1.3 ⇥ 1.3 pc around the dominant Herschel 36 (Her 36) system composed of at least three massive stars. To further explore themorphology of the region we compare archival infrared, optical and radio images of the nebula with our newly obtained fine structureline and CO data, in particular with the velocity information they provide. We perform a quantitative analysis, using both LTE andnon-LTE methods to determine the abundances of some of the observed species as well as kinetic temperatures and volume densities.
Results. Bright CO, [C II] and [C I] emission has been found towards the HII region and the photodissociation region (PDR) in M8.
Our analysis places the bulk of the molecular material in the background of the nebulosity illuminated by the bright stellar systems
Her 36 and 9 Sagitarii. Since the emission from all observed atomic and molecular tracers peaks at or close to the position of Her
36, we conclude that the star is still physically close to its natal dense cloud core and heats it. A veil of warm gas moves away fromHer 36 toward the Sun and its associated dust contributes to the foreground extinction in the region. One of the most prominent star
forming regions in M8, the Hourglass Nebula, is particularly bright due to cracks in this veil close to Her 36. By using radiativetransfer modeling of di↵erent transitions of CO isotopologues, we obtain H2 densities ranging from ⇠ 104 – 106 cm-3 and kinetic
temperatures of 100 – 150 K in the bright PDR caused by Her 36.
Key words. galactic: ISM — galactic: individual: M8 — submillimeter: ISM — ISM: HII region — ISM: clouds

Introduction

e influence of bright stars on their surrounding interstellar
dium is immense. Their strong ultraviolet and far-ultraviolet
lds give rise to bright HII regions and photodissociation reons (PDRs). These are the best grounds to study the e↵ect of
V and far-UV photons on the heating and chemistry of the instellar medium (ISM). The fine structure lines of C+ and O,

ous molecular and atomic species allows us to derive the physi
conditions in PDRs.
Messier 8 (M8) is located in the Sagittarius-Carina arm, n
our line of sight towards the Galactic Center. It is located at a dtance ⇠1.25 kpc (10 corresponds to 0.36 pc) from the Sun (Damani et al. (2004); Arias et al. (2006)) with an error of ⇠ 0.1 k
22(Tothill et al. 2008) and is about 34 ⇥ 12 pc in diameter.
M
is associated with the open young stellar cluster NGC6530,

9 SgrHer 36

1.5 pcFig. 1. The left panel shows the Lagoon Nebula, Messier object 8 (M8) or NGC 6523, in the constellation of Sagittarius, as seen by the Kitt Peak
4-meter Mayall telescope in 1973. North is up, east to left. Credits to "National Optical Astronomy Observatory/Association of Universities for
Research in Astronomy/National Science Foundation" Copyright WIYN Consortium, Inc., all rights reserved. The right panel shows the WISE
3.4 µm image of the M8 region around Her 36 (marked with a star) investigated in this paper. The contour levels are 10% to 100% in steps of 10%of the peak emission ⇠ 4000 data number. The 1.5 pc indicated in the lower left correspond to 40 .

tral baselines of first order and subsequently produced data cubeswith beam/3 sampling. For the beam sizes, see Table 1.
2.2. The APEX data

Observations of low- and mid-J 12 CO and 13 CO transitions were
performed with the Atacama Pathfinder Experiment (APEX)
12m submillimeter telescope (Güsten et al. 2006) during 2015
June - August and 2016 July, September and October. As shownin Table 1, we used the following receivers: PI230 to map the
low-J 12 CO and 13 CO transitions, FLASH+ in the 345 and 460
GHz bands to map the mid-J CO transitions and CHAMP+ in
low and high frequency sub arrays to map the higher frequency
mid-J CO transitions.
We used the PI230 receiver to map C18 O J = 2 ! 1 at
219.560 GHz , 13 CO J = 2 ! 1 at 220.398 GHz and 12 CO J
= 2 ! 1 at 230.538 GHz. FLASH+ was used in the 345 GHzband to map the 12 CO J = 3 ! 2 transition at 345.795 GHz.
FLASH+ was also used in the 460 GHz band to map the 13 CO J
= 4 ! 3 transition at 440.765 GHz, C18 O J = 4 ! 3 at 439.088
GHz and [C I] 3 P1 ! 3 P0 fine structure line at 492.160 GHz. The
CHAMP+ receiver was used to map 12 CO J = 6 ! 5 at 691.473
GHz and 13 CO J = 6 ! 5 at 661.067 GHz in the low frequency
sub array complemented by 12 CO J = 7 ! 6 at 806.651 GHz
and 13 CO J = 8 ! 7 at 881.272 GHz in the high frequency sub
array.
All maps shown in Fig. 3 were observed in on-the-fly(OTF) total power mode centered on R.A. 18h 03m 40s .3; Dec.
24 22’4300 (J2000), which corresponds to the position of Her
36. The maps obtained from the observations done in July,
September and October 2016 have a size of 24000 x 24000 . Mapsthat were obtained from the observations done in 2015 are comparatively smaller in size. We integrated 0.7 s per dump for allmaps and enough coverages were performed to reach the rms
noise levels as mentioned in Table 1. The o↵set position relative to the centre at (300 , 300 ) was chosen as reference, similar

to the SOFIA observations. The pointing accuracy (< 300 ) was
maintained by pointing at bright sources such as RAFGL5254
and R Dor every 1 – 1.5 hrs. A forward efficiency Fe↵ = 0.95
was used for all receivers, and the beam coupling efficiencies
Be↵ = 0.62, 0.69, 0.63, 0.43 and 0.32 were used for the PI230,
FLASH+ 340, FLASH+ 460, CHAMP+ 660 and CHAMP+ 810 receivers, respectively.
2.3. The IRAM 30 m data

Observations of low-J 12 CO, 13 CO and hydrogen recombinationline observations were performed with the IRAM 30 m telescopein August 2016. We observed the whole 3 mm range using the
EMIR receivers (Carter et al. 2012). We simultaneously mappeda region of 24000 x 24000 , which is similar to the size of most othermaps previously observed with SOFIA and APEX, in the 12 CO
and 13 CO J = 1 ! 0 transitions (Table 1) and the hydrogen recombination lines H40↵ to H43↵ at 99.023 GHz, 92.034 GHz,
85.688 GHz and 79.912 GHz, respectively. Molecular high density tracers, also detected in our wide spectral band observation,
will be analyzed in a subsequent paper.
All maps shown in Fig. 4 were observed in OTF total powermode centered on Her 36. Each subscan lasted 25 s and the integration time on the o↵-source reference position was 5 s. Theo↵set position relative to the centre at (300 , 300 ) was similar tothat used for the SOFIA and APEX mapping and the pointing
accuracy (< 300 ) was maintained by pointing at the bright calibrator 1757-240 every 1 – 1.5 hrs. A forward efficiency Fe↵ =0.95 and a beam coupling efficiency Be↵ = 0.69 were adopted
for the EMIR receivers. These values were taken from the latest(2015) commissioning report4 .

4

www.iram.es/IRAMES/mainWIKI/IRAM30mEfficiencies
Article number, page 3 of 18

23

Fig. 2. Colour maps of the integrated intensity of the [C II] 158 µmand J = 11 ! 10, J = 13 ! 12 and J = 16 ! 15 transitions of 12 CO toward
Her 36 which is the central position ( ↵ = 0,
= 0) at R.A.(J2000) = 18h 03m 40.3s and Dec.(J2000) = 24 220 4300 , marked with an asterisk. The
contour levels are 10% (> 3 ⇥ rms, given in Table 1) to 100% in steps of 10% of the corresponding peak emission given in Table 1. All maps areplotted using original beam sizes shown in the bottom left of each map.

All data reduction was done using the CLASS and MIRA
programs that are a part of the GILDAS5 software package andall the observations are summarized in Table 1.

3. Results
3.1. Peak intensities of the molecular line emission

The maxima of the distributions of the velocity integrated intensities of the emission in the [C I] and [C II] lines and vari5

www.iram.fr/IRAMFR/GILDAS/

ous transitions of 12 CO, 13 CO and C18 O are presented in TaA&A
proofs:no. maps
mainof the 12 CO
ble 1. Fig.
2 shows
velocitymanuscript
integrated intensity
J = 11 ! 10, 13 ! 12 and 16 ! 15 transitions. The emissionin all the lines has a similar spatial distribution and peaks arefound at about the same o↵set position ( ↵ = 5.000 ,= 5.000 )north-west of Her 36.
Figs. 3 and 4 show velocity integrated intensity maps of low
and mid-J transitions of 12 CO, 13 CO and C18 O, i.e. the J = 1 !
0, 2 ! 1, 3 ! 2, 6 ! 5, 7 ! 6 transitions of 12 CO, the J = 1! 0, 2 ! 1, 4 ! 3 and 6 ! 5 transitions of 13 CO and the J = 1! 0 transition of C18 O. The intensities of the low-J transitions

Article number, page 4 of 18

24

018 March 22; in original form 2017 October 13

304 Tes$ng the Larson rela$ons in massive clumps
Traﬁcante,1‹ A. Duarte-Cabral,2 D. Elia,1 G. A. Fuller,3 M. Merello,1 S. Molinari,1 N. Pereqo,2 E. Schisano1 and A. Di Giorgio1#56 A.
MNRAS 477, 2220

ABSTRACT

We tested the validity of the three Larson relations in a sample of 213 massive clumps
selected from the Herschel infrared Galactic Plane (Hi-GAL) survey, also using data from
the Millimetre Astronomy Legacy Team 90 GHz (MALT90) survey of 3-mm emission lines.
The clumps are divided into five evolutionary stages so that we can also discuss the Larson
relations as a function of evolution. We show that this ensemble does not follow the three Larson
relations, regardless of the clump’s evolutionary phase. A consequence of this breakdown isthat the dependence of the virial parameter α vir on mass (and radius) is only a function of
the gravitational energy, independent of the kinetic energy of the system; thus, α vir is not agood descriptor of clump dynamics. Our results suggest that clumps with clear signatures of
infall motions are statistically indistinguishable from clumps with no such signatures. The
observed non-thermal motions are not necessarily ascribed to turbulence acting to sustain thegravity, but they might be a result of the gravitational collapse at the clump scales. This seems
to be particularly true for the most massive (M ≥ 1000 M⊙ ) clumps in the sample, where
exceptionally high magnetic fields might not be enough to stabilize the collapse.
Key words: surveys – stars: formation – stars: kinematics and dynamics – stars: massive –stars: statistics – infrared: stars.

Larson (1981)
0.5(ii) power-law
Clouds are in approximately
with a virial
1. A size–linewidth
relaQon, σvirial∝ !equilibrium,
2/Evirial
R/GM ≃ 1, where
M isathemass parameterofparameter α vir = Ekin
G = 5σ equilibrium,
2. Clouds are approximately
withvirialdominated by highly supersonic
the region
and G is the gravitational constant. This relation implies
2!/$&%≃
=
"
/"
%=
5σ
1 of the system Ek ∝ σ 2 M is of similar intensity
αvirclouds# $ that the kinetic energy
persions in giant molecular
as the gravitationalenergy of'%∝the system,on et al. 1987), massive
clumps density
3. A volume'-size relaQon,!−1.1 EG ∝ M2 /R.

. Urquhart et al. 2014; Traficantemassive cores (size ≃ 0.1 pc; Zin-(iii) A volume density n-size relation, n ∝ R−1.1 . This relation
implies that GMCs are universal structures, with a mostly uniform

25

Table 2. Classification scheme of our 213 clumps following the evolutionary scenario described in Merello et al. (in preparation). In column 2, we give thewavelengths at which each evolutionary stage becomes bright. In column 3, we list the surveys that have a visible counterpart in the data. Column 4 shows themedian value of L/M and column 5 gives the number of identified objects.
Evolutionary
phase
Starless
Protostellar MIR dark
Protostellar MIR bright
YSOs
H II regions (UCH II + ext. H II)

Brightwavelengths (µm)

Survey

>70
≥70
≥21
≥8
≥8 and radio emission

Hi-GAL
Hi-GAL
Hi-GAL, MIPSGAL, WISE
Hi-GAL, MIPSGAL, WISE, RMS
Hi-GAL, MIPSGAL, WISE, RMS

L/M(L⊙ /M⊙ )
1.7
8.4
6.4
32.1
36.3

Count
14
12
106
25
56 (42+14)

Figure 2. L/M ratio distribution of our 213 clumps divided into the various
evolutionary phases. The entire sample spans more than four orders ofmagnitude and there is an evident increase of the L/M ratio going fromstarless to H II regions.

(ii) Protostar MIR dark, which is when a clump becomes visible
at 70 µm but it is still dark in the MIR, or the emission is too faint tobe identified. These clumps are bright at all Hi-GAL wavelengths,with no counterparts in the MIPSGAL, WISE and RMS surveys.
(iii) Protostar MIR bright, which is when the clumps also become visible in the MIR and their bolometric luminosity increases
significantly. These clumps have at least one counterpart in one of
the MIPSGAL, WISE and MSX surveys but they do not pass theRMS criteria to be classified as YSOs (Lumsden et al. 2013).
(iv) YSOs, which is when the protostars have reached the zeroage main sequence and have also become bright in the NIR regime.
They are classified as YSOs in the RMS survey.
(v) H II regions, where the thermal bremsstrahlung emission of
the gas ionized in the envelope of the more massive stars can be26

Dipartimento di Fisica e Chimica, Università di Palermo, Via Archirafi 36, 90123 Palermo, Italy

304 Structure of X-ray emiJng jets close to the launching site: from embedded to disk#57 bearing sources

Received 1 December 2017; accepted 15 March 2018

S. Ustamujic1, S. Orlando2, R. Bonito2, 3, M. Miceli3, 2, and A. I. Gómez de Castro1
ABSTRACT
A&Ap in pressContext. Several observations of stellar jets show evidence of X-ray emitting shocks close to the launching site. In some cases, theshocked features appear to be stationary, also for Young Stellar Objects (YSOs) at di↵erent stages of evolution. We study the case of
HH 154, the jet originating from the embedded binary Class 0/I protostar IRS 5, and the case of the jet associated to DG Tau, a more
evolved Class II disk-bearing source or Classical T Tauri star (CTTS), both located in the Taurus star-forming region.
Aims. We aim at investigating the e↵ect of perturbations in X-ray emitting stationary shocks in stellar jets; the stability and detectability in X-rays of these shocks; and explore the di↵erences in jets from Class 0 to Class II sources.
Methods. We performed a set of 2.5-dimensional magnetohydrodynamic numerical simulations that modelled supersonic jets ramming into a magnetized medium. The jet is formed by two components: a continously driven component that forms a quasi-stationary
shock at the base of the jet; and a pulsed component constituted by blobs perturbing the shock. We explored di↵erent parameters forboth components. We studied two cases: a jet less dense than the ambient medium (light jet), representing the case of HH 154; anda jet denser than the ambient (heavy jet), associated with DG Tau. We synthesized the count rate from the simulations and comparedwith available Chandra observations.
Results. Our model explains the formation of X-ray emitting quasi-stationary shocks observed at the base of jets in a natural way,
being able to reproduce the observed jet properties at di↵erent evolutionary phases (in particular, for HH 154 and DG Tau). The jet iscollimated by the magnetic field forming a quasi-stationary shock at the base which emits in X-rays even when perturbations formed
by a train of blobs are present. We found similar collimation mechanisms dominating in both heavy and light jets.
Conclusions. We derived the pysical parameters that can give rise to X-ray emission consistent with observations of HH 154 and
DG Tau. We have also performed a wide exploration of the parameter space characterizing the model; this can represent a useful tool
to study and diagnose the physical properties of YSO jets over a broad range of physical conditions, from embedded to disk-bearing
sources. We have shown that luminosity did not change significantly in variable jet models for the range of parameters explored.
Finally, we provided an estimation of the maximum perturbations that can be present in HH 154 and DG Tau taking into account theavailable X-ray observations.
Key words. ISM: jets and outflows – magnetohydrodynamics (MHD) – X-rays: ISM – stars: pre-main sequence – ISM: individual
objects: HH 154 – stars: individual: DG Tau

Introductionung Stellar Objects (YSOs) are stars at their early stages ofdisk, and most of its complex phenomenology can be mode
as a star interacting with a circumstellar accretion disk
27 (Stone
Norman 1994; Stone et al. 1996; Romanova et al. 2011; Orlan

304 Star forma$on toward the H II region IRAS 10427-6032
Vaidya⁎,a,
A B S Kaushar
T R A
C T Souradeep Bhaqacharyaa,b, Vatsal Panwara,c, Manash R. Samald, Wen-Ping Chend, Devendra K. Ojhae#58
New Astronomy 63, 27

The formation and properties of star clusters formed at the edges of the H II regions are poorly known. In thispaper, we study stellar content, physical conditions, and star formation processes around a relatively unknown
young H II region IRAS 10427-6032, located in the southern outskirts of the Carina Nebula. We study this region
by making use of the near infrared (near-IR) data from VISTA, mid infrared (mid-IR) from Spitzer and WISE, farinfrared (far-IR) from Herschel, sub-mm from ATLASGAL, and 843 MHz radio-continuum data. Using multi-band
photometry, we identify a total of 5 Class I and 29 Class II young stellar object (YSO) candidates, most of which
newly identiﬁed, in the 5′ × 5′ region centered on the IRAS source position. Modelling of the spectral energy
distribution for selected YSO candidates using the radiative transfer models shows that most of these candidates
are intermediate mass YSOs in their early evolutionary stages. A majority of the YSO candidates are found to becoincident with the cold dense clump at the western rim of the H II region. Lyman continuum luminosity calculation using radio emission indicates the spectral type of the ionizing source to be earlier than B0.5–B1. We
identiﬁed a candidate massive star possibly responsible for the H II region with an estimated spectral type
B0–B0.5. The temperature and column density maps of the region constructed by performing pixel-wise modiﬁed
blackbody ﬁts to the thermal dust emission using the far-IR data from the Herschel show a high column density
shell-like morphology around the H II region, and low column density (0.6 × 1022 cm −2 ) and high temperature(∼ 21 K) matter within the H II region. Based on the morphology of the region in the ionized and the moleculargas, and the comparison between the estimated timescales of the H II region and the YSO candidates in theclump, we argue that the enhanced star-formation at the western rim of the H II region is likely due to compression by the ionized gas.

dense clumps in which star formation is likely to occur (Kendrew et al.,
28 the2016). An expanding H II region may trigger star formation via