Imaging the water snowline in a protostellar envelope with H13 CO+?
Merel L.R. van ’t Ho↵1 , Magnus V. Persson2 , Daniel Harsono1 , Vianney Taquet1, 3 , Jes K. Jørgensen4 , Ruud Visser5 ,Edwin A. Bergin6 , and Ewine F. van Dishoeck1, 7
1

Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands
e-mail: vthoff@strw.leidenuniv.nl
2
Department of Space, Earth and Environment, Chalmers University of Technology, Onsala Space Observatory, 439 92 Onsala,Sweden3
INAF, Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy
4
Centre for Star and Planet Formation, Niels Bohr Institute & Natural History Museum of Denmark, University of Copenhagen,Øster Voldgade 5–7, 1350 Copenhagen K., Denmark5
European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany
6
Department of Astronomy, University of Michigan, 1085 S. University Ave., Ann Arbor, MI 48109-1107, USA
7
Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse 1, 85748 Garching, Germany
13+

背景：スノーラインの位置を観測的に制限することは、惑星形成メカニズムの解明に不可欠
最も重要な水のスノーラインは、原始惑星円盤での直接観測することがとても難しい

arXiv:1801.02636v1 [astro-ph.SR] 8 Jan 2018

目的：HCO＋は水のスノーラインの良いトレーサーなのか。

低質量星NGC1333-IRAAS2Aのエンベロープに対し、光学的に薄いH CO でマップピング。
H218O観測結果との比較で、本当にこれがよいトレーサーかどうかわかる。
ABSTRACT

A&A proofs: manuscript no. H13CO+_paper_arxiv

Context. Snowlines are key ingredients for planet formation. Providing observational constraints on the locations of the major snow1. Overview of the molecular line observations toward IRAS2A.
tiallycompositions
resolve the snowline.
the only thermal
water Tablelines is therefore crucial for fully connecting planet
to their Furthermore,formation mechanism.
Unfortunately,
the most important
lines
that
can
be
observedfromtheground(except
fortheHO
2 to the central star.
snowline, that of water, is very difficult to observe directly in protoplanetary disks due to its close proximity
Transition
Frequency Eup /k
Beamva
at 183 GHz)
arealines
the lesstracer
abundant
isotopologue
Aims. Based on chemical considerations, HCO+ line
is18predicted
to begoodfromchemicalof thewater snowline, because it is par00
(GHz)
(K)
( )
(km s 1 )such, even ALMA will have great difficulty locating+ ticularly abundant in dense clouds when water isHfrozen
2 O. As
out. This work aims to map the optically thin isotopologue
H+13 CO+ toward
13
the water snowline in protoplanetary disks.
H starCO thanJ = 3in disks.
2 260.255
25 0.93⇥0.68
0.09
the envelope of the low-mass protostar NGC1333-IRAS2A,
where
the snowline
is at larger distance from the18
An alternative
approachHsnowline.
22,0
203.408
204 0.87⇥0.72 0.115
13+is to use chemical imaging, as has
2 O 31,3
Comparison with previous observations of H18
O
will
show
whetherHCOisindeedagood
tracer
of
the
water
+
2
been done for the CO snowline. Due to CO being highly volatile,
J=2 1
10
2.1⇥1.7
4.06
Methods. NGC1333-IRAS2A was observed using
NOrthern
Extended(NOEMA)at ⇠ 000. 9DCO
resolution,
targeting144.068
itsthesnowline
is locatedtens Millimeterof AU from Array
the central
star (around13+(a)the H CO J = 3 2 transition at 260.255 GHz.⇠20
TheK;
integrated
emission
profilewasanalyzedusing
1D
radiative
transfermodeling
see e.g. Burke & Brown 2010). Although this is far Notes. Velocity resolution.
+of a spherical envelope with a parametrized abundancefor H13 CO. Thisprofilewaslocating
validatedwith a full chemical model.
enoughprofile
to be spatiallyresolved
withALMA,
it directly
13
+
00
18
remainsofdifficult;since COpeak,line emission
optically
Results. The H CO emission peaks ⇠2 northeast
the continuum
whereasisHgenerally
2 O shows compact emission on source.
it +does
not revealbytheatcold
midplane.
Reactions
within the inner ⇠360 AU to
Quantitative modeling shows that a decrease in thick,
H13 CO
abundance
leastdiska factorof sixisneeded+higher accretion rates) and the lower vapor pressure, the snow+ N 2 H+
gaseous
CO predicts
are the main
destruction
route forinNHCO
2 H , so
reproduce the observed emission profile. Chemical modelingindeeda steep increase
just outside
water snowline;
linetheis locatedfurther away from the star than in protoplanetaryis
expected
tobeabundant
only+ when CO is frozen out. Taking
the 50% decrease in gaseous H2 O at the snowline is not enough to allow HCO to be abundant. This places the disks
water(10s–100
snowlineAU
at 225instead of a few AU, Harsono et al. 2015;
simple chemical considerations into account, an upper limit for
AU, further away from the star than expected based
on
the
1D
envelope
temperature
structurefor
NGC1333-IRAS2A.
In
contrast,
+
Cieza
et
al.
2016).
In addition, H18
the CO snowline location can therefore be derived from N2 H
2 O has already been observed+
DCO observations show that the CO snowline isemission
at the expected
location,makinganHo↵outburst
scenario
unlikely.
toward four of these objects (Jørgensen & van Dishoeck 2010;(Aikawa
et
al.
2015;
van
’t
et
al.
2017),
as
has
13
+
13
et al.
2013).
Conclusions. The spatial anticorrelation of the Hbeen
COdoneandforH18
O emission
providea proofconcept
CO+ canbe 2012,
used asa The only thing lacking are thus highdisks around
TW Hyaand ofHD
163296that(QiH Persson
2thespatialresolution images of H13 CO+ . Here, we present NOEMA
tracer of the water snowline.
et al. 2013, 2015), and several protostellar envelopes (Jørgensen
observations of H13 CO+ toward one source, IRAS2A, and com2004; Anderl et al. 2016).
18
Key words. ISM: individual objects: NGC1333-IRAS2A–
ISM:
molecules–
astrochemistry–stars:
protostars– submillimeter:
For young protostars, the emission of complex organic pare these to the H2 O emission presented by Persson et al.
planetary systems(2012).
IRAS2Aisadeeply embedded Class 0 protostar in themolecules (COMs) has been used as tracer of the inner region
warm enough to sublimate water ice, the so called hot core, be- NGC1333 region of the Perseus molecular cloud. The quadruple(Jørgensen et al. 2004a),
cause these species are expected to be trapped in water ice. How- outflows suggest that it is a close binary
00
ever, the exact binding energy, and thus spatial extent, will di↵er which has recently been confirmed (0 . 6 separation, Tobin et al.
1. Introduction
freeze-out
of waterhas(e.g.,
Öbergal. 2011;
2015).
SimilarMadhusudhan
to Persson et al. (2012) we adopt a distance offor di↵erent molecules.
Compactradius
COM emission
been
de- et250
pc(Enochetal. al.
2006).
etal.
2014;
Ali-Dib
et
al.
2014;
Ali-Dib
2017;
Walsh
et
2015;
tected toward NGC1333-IRAS2A (hereafter IRAS2A) extendWater is probably the molecule that appeals mosting
to between
our imagiTheBoothobservations
results are presented in Sects. 2 and 3,
Mordasinial. 2016;
Eistrupet al. 2016;
et al.and2017).
⇠ 000. 4 and⇠ 100. 0 (e.g.,etJørgensen
et al.
2005; Maretand compared with the H1800nation, as it is essential for life as we know it. In star-forming
re- etInal.addition,
et al. 2014; Maury
2014), similar
to
the
size
(⇠
0
.
8)
of
2 O observations from Persson et al.
planetesimal formation is thought to be significantly
13+
18(2012).
In
Sect.
4,
the
integratedtheHOemission(Perssonetal.
2012).
The
exact
relationshipgions, water plays an important role as coolant of the warm
gas, enhanced in this region (e.g., Stevenson & Lunine 1988; Schoo- H CO emission is analyzed
2
using
1D
radiative
transfermodelingofa spherical envelope
between COMs
water snowline remains unclear though.
as dominant carrier of oxygen, and as major constituentof icyand thenenberg& Ormel2017).
The exact
and time evolution
with a parametrized
H13 CO+ abundance profile. In addition,
For example, the extent
of methanol
emissionin IRAS2A
is half location
grain mantles deep inside a planet-forming disk (seee.g., Mel- Oofemission
the watersnowline
in protostellar
systems
is thusa crucialabundance
profileis validated against the outcome of athe size of the H18(Maretet al. 2014),
while it is this2
18 2009; Bergin & van Dishoeck 2012; van Dishoeck+
nick
et al.
2
twice the sizein NGC1333-IRAS4A(Anderl
et al. 2016). In full chemical network and low-resolution DCO observations.
ingredient in planet
formation.
2013, for reviews). The transition from water being
frozenout et al. (2012) conclude that C2 H5 CN is likely The main conclusions, including that we confirm the predicted
addition,
Persson
H13 CO+ and H18
to sublimationof water ice. Since
detectionof anticorrelation
onto dust grains to being predominantly present innot
therelated
gas phase
2 O, are summarized in
Unfortunately,
the thewater
snowline
is very hardbetween
to observe
Sect.
5.
COMs
is
very
difficultinmatureprotoplanetarydisks,
with
only
occurs at the water snowline: the midplane radius at which 50%

観測：NOrthern Extended Millimeter Array (NOEMA) @ ~0.9ʼʼ 分解能
H13CO J=3-2 @ 260.255 GHz

H13CO+のemission peak：
一方、Hcontinuum peakの北~2ʼʼ

Oはコンパクトなemission

directly in protoplanetary disks. Because of the large binding

.: Imaging the water snowline in a protostellar envelope with H COe↵ect of the threeauses the H13 CO+y, it could perhaps
entary structure at=> 定量的モデル：

oward the continted in Fig. 4. Alnt on source, it iswhere water emis⇡ 0.8 km s 1 ) of
that the emission
on, the H18
2 O lineshown in Persson
he H182 O emission
ct component. The
spectrum is due to2).
d the systemic vege-scale envelope
channel maps inound the systemic
MT single dish obs that only ⇠7% ofervations, demond on larger scales.
n shows that somen the central changht at scales larger

このずれは、360AU以内でのH13CO+のabundanceの減少で説明できる。

A&A proofs: manuscript no. H13CO+_paper_arxiv

CO+O+ emission at thephysical-chemical
tive transfer code
the physical strucofiles for IRAS2A
he 1D spherically
Y (Ivezić & Elitzur
ameters were fitand the spatial exed on the chemiabundant when its
lest H13 CO+ disabundance insidend a higher abuner abundances arentegrated intensity
ence of the southannot be modelede restrict the comt. For a first analyi↵erence between
not be reproduced
ong the radial cut
fted channels arencluding all chand integrated radiale required to reach
wever, deriving thehis work. Finally,u, v) coverage are

Fig. 6. Left panels: Temperature (gray) and density (blue) profiles for the IRAS2A envelope from Kristensen et al. (2012) (top), and the corresponding H2 O (light blue), CO (blue) and HCO+ (purple) abundances predicted by the GRAINOBLE model (bottom). The dashed green line
shows the empirically inferred abundance profile for H13 CO+ (scaled down by a factor 100). Right panels: As left panels, but with the temperature
increased by a factor 2.2. The vertical dotted lines mark the H2 O and CO snowlines.

13>90 AU can roughly reproduce the observed integrated
intensity +20,000 AU radial extent of the model. An increased temperature
水スノーラインの丁度外側でH
CO が増加：(Fig. 5, second row). However, independent of the inner abun- is thus required to set the inner radius of the H CO abundance
13

Fig. 5. Di↵erent H13 CO+ abundance profiles (left panels) and the resulting simulated integrated emission along the northeastern part of the radial cut shown in Fig. 2 (right panels). The observed emission is shownin purple. The location of the abundance jumps (in Kelvin and AU) is
indicated in the top right corners. Bottom panel: Temperature (gray) anddensity (blue) profiles for the IRAS2A envelope from Kristensen et al.
(2012) used in the top two models. The location of the H2 O and CO
snowlines (at 100 K and 20 K, resp.) are marked by the dotted red andblue lines, respectively. The temperature profile increased by a factorof 2.2, used in the bottom two models, is shown in black. The resulting
H2 O snowline is indicated by the dashed red line while the CO snowline
now falls outside the adopted radial range. For IRAS2A, 100 corresponds
to ⇠250 AU.

produced from the Ratran images in GILDAS by creating a (u, v)
table with the function uv_fmodel.
Di↵erent H13 CO+ abundance profiles with their corresponding integrated emission profiles are presented in Fig. 5. For the
13physical+structure, the18
adopted
2 temperature drops below 100 K
around 90 AU. An H13 CO+ abundance of 1 ⇥ 10 10 at radii

H CO とH

dance, the emission always peaks on source, as in the case of
a constant H13 CO+ abundance at all radii (Fig. 5, top row). Although a large fraction (⇠90%) of the emission is already being resolved out with the (u, v) coverage of the observations, thisthus does not result in a dip in the H13 CO+ emission toward thesource position.
To13
reproduce +the location of the emission peak, the abundance jump has to be shifted outward to a radius of ⇠360 AU.
To move the 100 K radius, representing the water snowline, thisfar out, the temperature profile has to be multiplied by a factor of 2.2 (Fig. 5, third row). However, the observed integrated
intensity cannot be reached with this abundance profile, as the
H13 CO+ emission becomes optically thick for outer abundanceshigher than 3 ⇥ 10 10 . Varying the abundance does not a↵ect theposition of the emission peak.
The best match to the observed H13 CO+ radial profile canbe reached by introducing an additional drop in the abundancearound ⇠4000 AU (Fig. 5, fourth row). An increase in abundance of at least a factor of six is required around 360 AU to
get the emission peak at the observed location, but the exact radius of the outer drop does not influence the peak position. Sucha decrease in the H13 CO+ abundance can be expected around the
CO snowline because formation of HCO+ will be limited in thisregion due to a reduction in the parent gaseous CO molecules.

+at the water snowline, while the original temperature places theouter radius around the CO snowline. This discrepancy will befurther discussed in Sects. 4.2 and 4.3. The main result of thesemodels is that the observed position of the H13 CO+ emission
peak can only be reproduced by a reduction in the abundance ofat least a factor 6 in the inner ⇠360 AU.

化学モデルの予測と一致（水スノーラインが225AU付近で、一方H CO のabundanceピークは360AUからさらに135AU外側）
4.2. Comparison with a full chemical model

To inspect the legitimacy of the simple abundance profile, we
水スノーラインは225AUに置かれる。
compare it to the outcome of the astrochemical code GRAINO-

BLE (Taquet et al. 2012a,b, 2013, 2014). In short, GRAINOBLE
uses the rate equations approach (Hasegawa et al. 1992) considering three types of reactions (a so-called three-phase model
based on Hasegawa & Herbst 1993): reactions for gas-phase
species, reactions for species on the ice surface, and reactions forbulk ice species. The gas-phase chemical network is taken fromthe KIDA dabatase (Wakelam et al. 2015) supplemented by ionneutral reactions described in Taquet et al. (2016). A cosmic ray
ionization rate of 5 ⇥ 10 17 s 1 is used. In addition to thermal
desorption, UV-photodesorption and chemical desorption upon
formation of an exothermic surface reaction have been included
as desorption processes. The surface chemical network follows
the results of recent laboratory experiments in order to form the

1Dエンベロープの温度構造から予測される距離よりも、かなり遠い

Oの空間的不一致：

Article number, page 5 of 10「H13CO+は水スノーラインのトレーサーとして使える」ことを示す。

The Near-Infrared Outflow and Cavity of the proto-Brown Dwarf
Candidate ISO-Oph 200?
Whelan, E.T.1 , Riaz, B.2 , and Rouzé, B.1
1

Maynooth University Department of Experimental Physics, National University of Ireland Maynooth, Maynooth Co. Kildare,
proto-brown dwarf候補天体のNIR観測論文。
Ireland2

Max-Planck-Institut fr Extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany

天体：ISO-Oph 200（0,06+-0.01 Msun）。SEDからClass Iと考えられている（Evans+09）
ABSTRACT

arXiv:1801.08749v1 [astro-ph.SR] 26 Jan 2018

In this lettera near-infrared integral fieldstudy of a proto-brown dwarf candidate is presented. A ⇠ 0.00 5 blue-shifted outflow is detected
観測：integral ﬁeld
spectrograph
SINFONIを使ってVLTで観測。
1
1
1
Whelan, E.T. et al.: The ISO-Oph 200 NIR Outflow
in both H2 and [Fe II] lines at V sys = (-35 ± 2) km s and V sys = (-51 ± 5) km s respectively. In addition, slower (⇠ ± 10 km s ) H2
emission is detected out to < 5.00 4, in the Outflow
direction of
the blue
V sysbothO↵set
PA and red-shifted outflow lobes but along a di↵erent position angle
00
to the more compact faster emission. It isargued(kmthatemission is a jet and the extended H2 emission is tracing
Features 1 )the (more)(compact)
H2 Jet0.00 7± 1230
a cavity. The source extinction is estimatedat Av -35
= ±2
18 magmag and the outflow extinction at Av = 9 mag ± 0.4 mag. The H200
Fe II Jet
-51 ± 50. 5
215
outflow temperature is calculated to be 1422
density of the [Fe II] outflow is measured at ⇠ 10000 cm 3 .
F1 K ± 255
4 ± 3K and
4.00 0the electron
67
10
00
F2
3
2.=6 3.8 80
Furthermore, the mass outflow rate is estimated
at10Ṁ±out[H2]⇥ 10 M yr 1 and Ṁout[FeII] = 1 ⇥ 10 8 M yr 1 . Ṁout[FeII] takes-14 ±is5 investigated
5.00 4248 using the ratio of the [FeII] 1.257 µm and [PII] 1.188 µm lines.
a Fe depletion of ⇠ 88% into account. TheF3depletion
Outflowfeaturescorresponding
Using the Pa and Br lines Table
and a1.range
in emission
stellar mass
andandradiusṀacc is propercalculated to be (3 - 10) ⇥ 10 8 M yr 1 . Comparing theseties. O↵sets are with respect to the source position and along therates puts the jet efficiency inlinePA.
with
ofofmagneto-centrifugalmodelsof jet launching in low mass protostars. This is agiven
PAs predictions
are measured E
N. It is argued that F1, F2
andfurther case of a brown dwarfF3outflow
analogous
form partexhibitingof a cavity aroundthe jet. properties to protostellar jets.

コンパクトなblueシフトoutﬂow：

全てのラインで見えている

Key words. stars: brown dwarfs - stars:jets - stars:formation

solved H2 emission is seen at distance of 2.00 6 to 5.00 4 and along adi↵erent PA to the faster emission. The main emission features
here are marked F1, F2 and F3 in Figure 2. Theinvelocity,mor1. Introduction
protostellar
HH jets (Riaz et al., 2017). Here we present thephology and PA of this emission would suggest that it is tracing
resultsofanear-infrared
study of an outflow from a proto-BDa cavityaroundthe formation
collimated jet. processThis emission is not detected
Outflows play an important rolein thestarin [Fe II] and is discussed further in Section 4. candidate, which also shows characteristics similar to protosteland are driven by young stellar objects(YSOs)
in theClass
0 lar
A H2 excitation
diagramcan beconstructedby plotting
the driving source of the outflow is a YSO named ISOjets. The
loget
of al.,
the column
the statisticalto Class II evolutionary stages natural(Frank2014).densities
Thus,(divided
they byOph
200(↵ = 16:31:43.8, = - 24:55:24.5). The near- to midweight),
of a200
numberof H2 lines against the upper energy levelal.: Therange
ISO-OphinNIR Outflow
can be observed Whelan,acrossE.T.
a etlarge
wavelength(Whelan, infraredµm) spectral slope of the SED for ISO-Oph 200of the line. This can be used to derive the temperatureand(2-24
to2014a). Also, as di↵erent wavelengthcanconstrainregimes
the extinctionof o↵erthe linecompleregion
et al., with a Class I classification (Evans et al., 2009).
is(Davisconsistent
Outflow
Vemitting
O↵set
PA
sys
1 temperature2011).toIn launchinga thermalised
gas at a(kmsinglementary observational constraints
models,Featuresmulti)(00 ) We( )the points
Fig. 2. Continuum subtracted spectro-images of the ISO-Oph will lie along a straightH2line,00
the -35
slopeof which
is 230the havetempera-conducted physical and chemical modelling of subJet±2
0.
7
studieshave
proventureimportant(Nisiniet al.,
2016).
200 outflow in the H2 wavelength
1-0 S(1) and [FeII] 1.257µm lines.
(Caratti o Garatti
et al.,
is215
constrained
continuum and molecular line data for thisFe
II Jet2006).
-51The± 5 extinction
0.00 5 millimeter/millimeter
Velocities have been corrected
for the systemic
velocityof thefield
Near-infrared(NIR)
integralobservations
areby investigating
the scatterthe 4points
straight
line. find the results to be consistent with an early StageF1 aofparticularly± 3 about
4.00 0this67
object,and00source. Top panels: Herethe emission
in a studies
mainly blue-shifted
10 ±which
3 should
2. 6 reduce
80 the scatCorrecting
line fluxes
for extinction
usefultool forof jet launchingintheClass
0/IF2sources
0+I
stage (Riaz et al. in prep). The total (dust +55.00 4 varying
248 evolutionary
velocity range is shown to highlight the blue-shifted jet emission ter of the points about F3
a linear -14
fit. ±Therefore,
Av and
are
still
very
much
embedded
in
their
natal
clouds(Davisetal.,only. Black contours begin at 3 and increase in multiples of 1.5. analysing
the1.goodness
ofemission
the fit leads
to an andestimateof Amassv (Toddgas)
for
ISO-Oph
200 as derived from the sub-millimeter
Table
Outflowfeaturescorresponding
proper2
1
1 = 4.5 ⇥ 10 15 erg cm2011).
s 1 µm The. TheJ,H,
blue contours& RamsayHowat,
2006).
H2 excitation
for theand K show
NIRthebands
cover
Fe IIareforbidden
andHdiagramties. O↵sets
withArespect
tothe source
andISOalongthe
2 position
850
µm
flux
is
0.06±0.01
M with a bolometric luminosity ofcontinuum emission and they correspond to 50%, 75% and 95% Oph 200
jetPA.
wasPAsconstructed
Figureusing theare measured
Eof N. A.1)
It is argued
thatfluxes
F1, F2 andro-vibrational emission lines that aregiven
strongtracersof9 (see
jets.
These
⇠
0.09±0.02
L
.
The
mass
of the central object in this Class Iof the peak flux of the continuum. The black stars marks the po- given F3
in Table
1.
For
A
=
mag,
T
=
1442
±
255
K
is
meav
form part of a cavity
around the jet.
linesandcan
usedPA
toisinfer
jet properties
suchisas
temperatureandsition of the emission peaks
the beemissiondelineatedsured. This resultin agreementwith thetemperatures
reportedsystemcan be constrained to within the sub-stellar mass regime
by the dashed line). Bottom
panels:emission in
the extinction
by Garcia Lopezet al. (2013)
the very low mass protostar
densityandHeretothe
measurethetowards
the jetfor (Garcia
theofmeasured bolometric luminosity and numerical simvelocity ranges -80 km s 1 to 128 km s 1 (H2 ) and -96 km s 1 to IRS54. By varying the extinction it is found that using
the analysis
Lopez
et
al.,
2013).
Integralfield
spectroscopy
providesspectro1
114 km s ([Fe II]) is shown. This range was chosen to show si- the excitation diagrams is in a agreement with theAestimates
stellarevolution, as discussed in Riaz et al. (2016).
00ulations
00v
solved H2 emission is seen at distance of 2. 6 to 5. 4 andofalongamultaneously the extended
blue andred-shifted
low velocity
H2 presented
imageswhich
allow
the properties
of
the
jets
mappedasThea mainin Section
3.1.
di↵erent
PA totothebe
fasteremission.
emission features
Consideringthevery
low mass reservoir in the (envelope+disk)
emission which it is argued
traces the
The chosen
range
Using
the
lines
and
the models
ofvelocity,
Takami morfunction
ofcavity.
velocityand distance(Davisetratiosal.,of2011).
Such
herethearemarked
F1,[FeF2II]
andF3 studin Figure
2.
The
for
this
system,andthepresence of an outflow that will further
also shows that there is no red-shifted counterpart to the blue- et al. (2006), the electron density (ne ) can be estimated.
Studies
phologyandbe
PA tracingof this emission
would suggest that it is tracing
found
theis[FeII] andH2 foundlinesthattodi↵erent
shifted jet emission andies
that havethe cavity
emission
not seenin havethe 1.600µm / 1.644µm,This1.664µm/ 1.644dissipate
the envelope material, ISO Oph-200 will likely have aa cavityaroundthe collimatedjet.
emission
is notµm,
detected[FeII]. Contours start at components
3 and increase inof 1.5.
1 theandforbidden
1.677µmII]/ and1.644
ratiosassociatedaremoreinaccurate
tracers of theofmultiples
the jet,withemissionin [Feis µm
discussed
further
Sectionfinal
4.
massto below the sub-stellar limit, and can be considered as= 5 ⇥ 10 15 erg cm 2 s 1 µm 1 . The continuum emission is plot- density than the 1.534 µm / 1.644 µm ratio. However,withfaster,hot
densematerial
andtheAmolecular
H2 excitationemission
diagram cantracbe constructed
bydueplotting
theastrongcandidate
proto-BD.
ted in the same way as for the top panels.
the non-detection
of
the
1.600
µm,
1.664
µm
and
1.677
µm
naturallog (Carattiof the column
densitieset(divided by thelines
statisticaling lower excitation, slower moleculargasoGarattihere we are limited to using the 1.534 µm / 1.644 µm ratio. From

速度の遅い（~ +-10km/s）red&blueシフトの広がったemission：
H2のみで見える

方向は、速いコンパクトなemissionとは異なるPAを持っている。

コンパクトemission => jet

広がったH2 emission => cavityをトレースSourceの減光量：

Av = 18 mag +- 1 mag

Outﬂowの減光量：

Av = 9 mag +- 0.4 mag

H2 outﬂowの温度：

1422 K +- 255 K[Fe II]の電子密度：with color-color diagram (Davis+11)

Whelan, E.T. et al.: The ISO-Oph~10000 cm-3

mass outﬂow rate：
For H2,

M = 3.8 x 10-10 Msun/yr

For FeII,M = 1.0 x 10-8 Msun/yr

PaβとBrγ線のLuminosity、星の質量（0.06 - 0.1 Msun）、
半径（0.9 - 0.5 Rsun ）を用いて：
Macc ~ (3-10) x 10-8 Msun/yr

proto-BDのNIR観測は、BD質量天体のjetを調べるには良いツール。
広い角度のwindのようなoutﬂow成分の証拠を探すことができ、jetの減光が調べられ、jetの効率が測定できる。
近い将来、JWST/NIRSPECを使ったより高分解能でのjetのcollimationやstructureの観測が可能になる。

Gravity drives the evolution of infrared dark hubs:
JVLA observations of SDC13
G. M. Williams1 , N. Peretto1 , A. Avison2,3 , A. Duarte-Cabral1 , and G. A. Fuller2
1

School of Physics and Astronomy, Cardi↵ University, Queens Buildings, The Parade, Cardi↵ CF24 3AA, UK
email: williamsgwen12@gmail.com
2
Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Manchester, M13 9PL, UK
背景：フィラメントが集中している箇所（hub）は、星団・大質量星形成と深い関わり。
3
UK ALMA Regional Centre Node

ハブの進化とそこでの星形成の関係の理解は重要。

iv:1801.07253v1 [astro-ph.GA] 22 Jan 2018

Received Month Day, Year; accepted Month Day, Year

ABSTRACT
目的：SDC13の運動・密度構造を調べ、進化やfragmentationの駆動源を決める
Context. Converging networks of interstellar filaments, i.e. hubs, have been recently linked to the formation of stellar clusters andmassivethe relationshipbetween the evolutionof these systems
and theGreen
formation ofBankcores/stars
inside them is (0.07 pc)
観測：NH3(1,1)& stars.
NHUnderstanding
with Jansky
Very Large
Array,
Telescope3(2,2) 輝線

at the heart of current star formation research.
Aims. The goal is to study the kinematic and density structure of the SDC13 prototypical hub at high angular resolution to determine
G. M. Williams et al.: Gravity drives the evolution of infrared dark hubs: JVLA observations of SDC13
what drives its evolution and fragmentation.
Methods. We have mapped SDC13, a ⇠1000 M infrared dark hub, in NH3 (1,1) and NH3 (2,2) emission lines, with both the Jansky
Very Large Array and Green Bank Telescope. The high angular resolution achieved in the combined dataset allowed us to probe scales
down to 0.07pc. After fitting the ammonia lines, we computed the integrated intensities,
velocities and line widths, along
全てのcentroid
hubフィラメントはsupercriticalwith gas temperatures and H2 column densities.
Results. The mass-per-unit-lengths of all four hub filaments are thermally super-critical,
consistent with the presence of tens of（transonicの非熱的な速度分散を持つ）
gravitationally bound cores identified along them. These cores exhibit a regular separation of ⇠ 0.37 ± 0.16 pc suggesting gravitational
instabilities running along these super-critical filaments are responsible for their fragmentation. The observed local increase of thedense gas velocity dispersion towards starless cores is believed to be a consequence of such fragmentation process. Using energy
コアの間隔：
フィラメントに沿って~0.37pc
conservation arguments, we estimate that the gravitational to kinetic energy conversion
efficiency in the SDC13
cores is ⇠ 35%. We
see velocity gradient peaks towards ⇠ 63% of the cores as expected during the early stages of filament fragmentation. Another clear
observational signature is the presence of the most massive cores at the filaments’ junction, where the velocity dispersion is the largest.
We interpret this as the result of the hub morphology generating the largest acceleration gradients near the hub centre.
Conclusions. We propose a scenario for the evolution of the SDC13 hub in which filaments first form as post-shock structures in afew
Myer（非平衡
supersonic turbulent flow. As a result of the turbulent energy dissipation in theフィラメントの年齢：ashock, the dense gas within the filaments
is initiallymostly sub-sonic. Then gravity takes over and starts shaping the evolution of the hub, both fragmenting filaments and pulling the gas
towards the centre of the gravitational well. By doing so, gravitational energy is フィラメントの半解析的モデル（Clarke+16,
converted into kinetic energy in both local (cores) andglobal (hub centre) potential well minima. Furthermore, the generation of larger gravitational acceleration gradients at the filament
junctions promotes the formation of more massive cores.
17）を使用）Key words. stars:formation – stars:massive – ISM:clouds – ISM:kinematics and dynamics – ISM:structure

1. Introduction
The star formation process can be perceived as the mechanism
leading to the concentration of di↵use interstellar clouds into
compact, nuclear burning, balls of gas. While the importance of<= Peretto+14の力学的タイムスケールとconsistent

While a lot of e↵ort has focused on the properties of individual filaments in nearby star-forming regions (Arzoumanian et al.
2011; Peretto et al. 2012; Palmeirim et al. 2013; Panopoulou
et al. 2014; Salji et al. 2015), here we focus on hub filament
systems (HFS), i.e. a small network of spatially converging interstellar filaments (Myers 2009). The converging nature of such

コアでの重力エネル
ギーから運動エネル
ギーへ変換効率を求めた（ε~ 35 %）<= 理論値よりも大きい

G. M. Williams et al.: Gravity drives the evolution of infrared dark hu

フィラメントの結合部に最もmassiveなコアが付随（速度分散は最も大きい）<= ハブ形状によって生まれる加速度勾配に起因すると考えられる
SDC13の進化シナリオ：
超音速乱流でのpost shock構造として、フィラメントが形成
乱流エネルギーの散逸の結果、高密度コアは亜音速に重力が優位になり、hubの進化を促す。
それによって、重力エネルギーは運動エネルギーへと変換されるフィラメント連結部でのより大きな重力加速度勾配の発生が、大質量コア形成を促進