星形成ゼミ (SFN #301 36-42) 2018/2/23 辰馬
36. Episodic accretion in binary protostars emerging from self-gravitating solar mass cores,
Riaz et al., http://arxiv.org/pdf/1712.09646
37. On fragmentation of turbulent self-gravitating discs in the long cooling time regime, Rice
& Nayakchin, http://arxiv.org/pdf/1712.05610
38. Phosphorus-bearing molecules in the Galactic Center, Rivilla et al., https://arxiv.org/pdf/
1712.07006
39. A new compact young moving group around V1062 Sco, Riser et al., http://arxiv.org/
pdf/1712.10143
40. Intensity-Corrected Herschel Observations of Nearby Isolated Low-Mass Clouds,
Sadavoy et al., http://arxiv.org/pdf/1712.00017
41. Accretion signatures in the X-shooter spectrum of the substellar companion to SR12,
Santamaría-Miranda et al., http://arxiv.org/pdf/1712.09297
42. The HR 4796A Debris System: Discovery of Extensive Exo-Ring Dust Material,
Schneider et al., http://arxiv.org/pdf/1712.08599
1

On fragmentation of turbulent self-gravitating discs in the long cooling time regime

Ken OnRice fragmentation, Sergei Nayakshin
37.
of turbulent self-gravitating discs in
Scottish Universities Physics Alliance (SUPA), Institute for Astronomy, University of Edinburgh, Blackford Hill,
thefor longtime
regimeEdinburgh, EH9 3HJ; Centre
Exoplanetcooling
Science, Universityof Edinburgh,
Edinburgh, UK; Department of
1,2

3

1

2

3

Physics and Astronomy, University of Leicester, Leicester, LE17RH, UK
E-mail contact: wkmr at roe.ac.uk

Ken Rice, Sergei Nayakshin

It has recently been suggested that in the presence of driven turbulence discs may be much less stable against gravitational collapse than their non turbulent analogs, due to stochastic density fluctuations in turbulent flows. Thismode of fragmentation would be especially important for gas giant planet formation. Here we argue, however, thatstochastic density fluctuations due to turbulence do not enhance gravitational instability and disc fragmentation in
the long cooling time limit appropriate for planet forming discs. These fluctuations evolve adiabatically and dissipate
away by decompression faster than they could collapse. We investigate these issues numerically in 2D via shearing box
simulations with driven turbulence and also in 3D with a model of instantaneously applied turbulent velocity kicks. In
the former setting turbulent driving leads to additional disc heating that tends to make discs more, rather than less,stable to gravitational instability. In the latter setting, the formation of high density regions due to convergent velocity
kicks is found to be quickly followed by decompression, as expected. We therefore conclude that driven turbulence
does not promote disc fragmentation in protoplanetary discs and instead tends to make the discs more stable. We
also argue that sustaining supersonic turbulence is very diﬃcult in discs that cool slowly.
Accepted by MNRAS
http://arxiv.org/pdf/1712.05610• Hopkins & Christiansen (2013)で提唱された、原始惑星系円盤内の超音速乱流によるガスの密度ゆらぎによって、高密度ガス領域が形成され、その領域が重力不安定になり収縮するという理論をシミュレーションで確かめた。
24• 原始惑星系円盤は冷却時間が局所的な力学時間(~1/Ω)に対して長いため、乱流により円盤が加熱されると冷却が不十分となり重力不安定にならないことがわかった。

• また、冷却時間の長い円盤内で超音速乱流を維持することが難しいこともわかった。
2

•2Dシミュレーション（乱流なし）

Fragmentation of turbulent discs

3

Fragm

Table 2. Baseline simulation results.
imposed, with initially a large ⌧c which then decays to theシミュレーション結果
prescribed value of ⌧c for a given simulation over the next 50
Disc fragments?
time units. In these simulations the instability often undergoes an initial transient burst phase before settling towards a4
Yesquasi-steady state, and this relaxation procedure avoids the5?
system artificially fragmenting during the initial transient
7
No
phase. The simulation is then run for a further 100 time
8
No
units before imposing the driven turbulence (i.e., the turbu- 10
Nolent forcing is typically turned on at t = 200). This, again,
Power spectra for the two simulations shown in Figure
ensures that the system is well past the transient burst phaseected, the simulation in which the turbulence is forced
and is likely close to the final state to which it would settle
cales (k = 1.0) has more inpowerat thesesmall
scales turbulent forcing, before inthe absenceof anadditional
simulation in which the turbulence
is forced
at larger
troducing theadditional
turbulent forcing. We will specify= 0.1). A k 5/3 power law(dotted
line)
is
shown
any simulations below that for
deviate from this initialization
procedure.

- β=(冷却時間)/(力学時間)=(冷却時間)Ω
- β≤4で円盤が分裂- Gammie (2001)の条件β≤3とほぼ一致• 2Dシミュレーション（乱流あり）
Fragmentation of→
turbulentdiscs 3
マッハ数Mが計算できる- パラメータ: 乱流強度fo、波数k

2.3 Turbulent forcing
Turbulentge⌧c which simulation
then decaysparameters.
to theWe wish to use a simple mathematicalmodel to describe
波数k=0.1
波数k=1.0
乱流パラメータ
ven simulation over the next 50
turbulent disc flows. For the PENCIL CODE simulations we
ons the instability
Figure 5. Reynolds,
fo oftenk underMa turbulent forcing method first introduced by Haugen,usephase before settling towards ain terms of a viscous ↵Brandenburg & Dobler (2004). In this method, a forcing
relaxation procedure
avoidsthe
0.1
0.1 function,
0.14
bulence and in which
f , which depends on position and time, describes
ng during the 0.1
initial 1.0
transient
0.27
into a quasi-steady st
thelocal acceleration that gas experiences due to turbulent
en run for a further
100
time
at ↵ = 0.028 shows th
0.25
0.1 eddy
0.30 motions. In this approach, the forcing function has theven turbulence0.25(i.e., theturbuwhich the e↵ective vi
1.0 form
0.57
ed on at t = 200).
rate at which the syst
0.5 This,
0.1again,
0.57
ll past the transientnois that if the system
0.5 burst
1.0 phase
0.95
f (x , t) = Re N f k (t) exp [ik (t) · x + i (t)] ,(3)l state to which it would settle
instability will act to
1.0
0.1
1.09
nal turbulent forcing, before inFigure
3. Density
structure for two simulations, one with a cool1.0
1.0 where
1.5 x is the position vector. The wave
vector,k (t), andbulent forcing. We will specify
ing timeof time= step.
4 (top panel) and the other with a cooling time
2.0
0.1 the2.0random phase, ⇡ < (t) < ⇡, change
at eachdeviate from this initializationof = is8 independent(bottom panel), but that do not include additional tur2.0
1.0 So2.32
that the time integrated forcing function
3.2 Turbulence
Figure4. Toomre Q profile for the simulation, without additionalbulentforcing.
The
top
panel
shows
a
numberofhigh-density
5.0
0.1 of3.90
the time step t, the normalisation factor,
N , has into whichbeturbulence,the cooling time is= 8. The cooling is
1/2
We carried out a seclumps,indicating
that
this
simulationisundergoing
fragmentaproportional to t . As in Haugen, Brandenburg&
Doblerturned on at t = 50 and allowed to decay to the prescribed
rate
1/2
ing times as in thebottomshows
spiral density
waves,
noconstant.
high density(2004), we take it to be N = fo cs (|k |ction.
t)The
.the
To
vary 50thes / overnext
timesteps,
after which
it isbutkept
Bytion 3.1) but in whclumps,indicatingthat
simulation
has settledstate
into ina which
quasistrength of the forcing, we vary the coefficient
fo .the systemt ⇠ 200
hasthissettledinto a quasi-steady
ULTS
athematical model to describe
described in §2.3, w
steadyin which
In each simulation we specify the magnitudeof theforc- the instability is primarily acting to transQ isstateroughly
constant.
PENCIL CODE simulations weFigure1.
Density
structurefor
twosimulations,one
forcedwithing wavenumber, |k |, but at each timestep
we randomly
seport angularmomentum.
andwithforcing wa
selinesimulations
od first introduced by Haugen,aforcing
wavenumber
of
k
=
0.1
(top
panel)
and
the
other
with
a
3 syslect the direction of this wavevector. We then force theillustrated
in Table04). In this method, a forcing
forcing wavenumber of k = 1.0 (bottom panel). Both simulations
our baseline simulationstem
to with
be those
initialised
as waves
nonhelical,
transverse

scales (k = 0.1). A kreference.

5/3 power law (dotted line) is shown for• 2Dシミュレーション（乱流あり）
Ken Rice & Sergei Nayakshin
Table 1.-Turbulent
parameters.
β=4:simulation
乱流なしで唯一分裂
乱流なし
fo

k

6

Ken Rice & Sergei Nayakshin

乱流ありで分裂しなくなる(波数1.0)

M

分
裂

0.1
0.1
0.14
0.1
1.0
0.27
0.25
0.1
0.30
0.25
1.0for
0.57two
wer spectra
Figurefor the2. twoPowersimulations
spectrashown
thein Figuresimulations shown in Figure
乱流
0.5the0.1
0.57, the simulation
1. As expected,
in whichthesimulation
turbulence
in which
is forced
the turbulence is forced
0.5= 1.0)(k = 1.0)
at smallhas morescales power
(k
at1.0
these
has0.95
more
small power
scales at these small scales
1.0 in 0.1
1.09
lation in
than
which
the thesimulation
turbulence
which
is forced
the atturbulencelargeris forced at larger
5/3 power
5/3
1.51). A k scales(k = law0.1).1.0(dottedA k 1.0
line)
power
is shownlaw (dottedforline) is shown forFigure 3. Density structure for two simulations, one with a cooling time of = 4 (top panel) and the other with a cooling time
2.0
0.1
2.0
reference.
of = 8 (bottom panel), but that do not include additional tur2.0
1.0
2.32
bulent forcing. The top panel shows a number of high-density
5.0
0.1
3.90
clumps, indicating that this simulationis undergoingoffragmentaFragmentation
turbulent discs 7
ulent
simulation
Table
1.the
parameters.
Turbulentsimulation
parameters.
tion. The bottom shows spiraldensitybut
no highe
2. Power
spectra fortwo simulations
shown in Figure
Figure6. Figurewaves,
showing thesurfacedensitydensity
structure in a simFigu
expected, the simulation in which the turbulence is forced
ulation
with
=
4
and
with
an
imposedturbulence(initiatedgraviclumps, indicating that this simulation has settled into a quasi3
RESULTS
all scales (k = 1.0) has more power at these small scales
at t = 200) with k = 1.0, and fo = 1.0. In the absence of thishoriz
steady
state
in
which
the
instability
isprimarily
actingtotransthe simulationfo
in whichkthe turbulence
isforcedatlargerimposedturbulence,
thissimulation
fragments
(top
panel
of
Figwe’dMfokM
5/3 power law (dotted line) is shown for
(k
=
0.1).
A
k
ure
3).
Turbulence
forcedatk=
1.0
appears,
however,
toinhibit=4
port angular momentum.
3.1 Baseline simulations
nce.
fragmentation.
expec
0.1
0.1
0.14
0.1
0.1
0.14
lence
We take 0.1
our baseline
simulations
to
those initialised as
1.0
0.27
0.1
1.0be 0.27
domi
disap
described
in Section
2.2 withan energy
equation and cool0.25
0.1parameters.
0.30
0.25
0.1
0.30
1. Turbulent
simulation
However, for k = 1.0 there appears to be a forcing amriumThe= 4 simulation shows
clear
signs
of fragmentation,ing, but0.25
without
additional
forcing. In other
1.0 any0.57
0.25 turbulent
1.0
0.57
plitude above
which
the turbulent
forcing inhibits fragmenof theand w
tation.
Figure6showsasimulation
with
=
4,
f
=
1.0,
f
k
M
o
o
乱流
0.5 only0.1
0.57
0.1
0.57
while the = 7, 8 and 10
simulations all settle into a quasiwords, theturbulence
is0.5
the gravito-turbulence
drivenand k = 1.0. There are no indications of any high density
0.5 0.1 1.0
0.95
0.5itself.
1.0Figure0.95 3 shows a base0.1 0.14
steady state. The = 5 simulation
does
havethesome
regions
ofvalues forby the gravitational
instability
fragments. Figure
7 showse↵ectiveviscous ↵
0.1
1.0
0.27
1.0
0.1
1.09
1.0
0.1
1.09
simulation. For fragment.
= 4, we would
expectthe total ↵ tohigh density, but doesn’tthis
unequivocally
This
could
line simulation=80.25 with
0.1 0.30 = 4 (top panel) and one with
Figure3.
Density
structure
Figurefor
3.
twoDensitysimulations,structure
one
for
with
two
a
simulations,coolonewitha oncoolsettle
somewhere
close
to
↵
=
0.056.
Prior
toturning
the
1.0
1.0
1.5
1.0
1.0
1.5
0.25 1.0
be an indication of stochastic
fragmentation(Paardekooper(bottom panel).
The0.57 = 4 simulation has a number of dense
5 sim
turbulent
forcing
at tthe= 200,
the total↵ isa heading
towards
0.1 0.57
ing time of = 4 (toping
panel)
time andof the= 4other(topwithpanel)
a cooling
and
time
other
with
coolingtime
2.0 0.5 0.1
2.0
2.0
0.1
2.0
there
2012; Young & Clarke 2016),
or thesehighregionsforcing isclumps, indicating
the expectedvalue, but
dropsdensity
once the turbulent
1.0 that
0.95 it is undergoing fragmentation. Theof=
8(bottom
panel),
of
but
=
8
that
(bottomdonotpanel),
includebut
additional
that
do
not
turinclude
additional
tur2.0 0.5
1.0
2.32
2.0
1.0
2.32
the
1.0
0.1does
1.09 not, and instead shows the presence
initiated.
The Reynolds
contribution
continues
to provide
may
simply
shear
out
and
the
systemmaintainanapproxi= 8 simulationby thbulentforcing.
The topbulentpanelshows Thea about
numbertophalfpanel
high-density
shows
numberhigh-density
Figure 3. Density
structure
for two simulations,one
withforcing.
a cool5.0 1.0 0.1
3.90
5.0
0.1
3.90
1.0
1.5
ofoftheexpectedastress,
but theofgravitational
conhavematelyquasi-steadyof spiral density
waves,
intimea quasi-steadyof =clumps,
4 (top panel)and theotherwitha coolingtime state.
2.0
0.1
2.0 indicating that it ising
virtually
disappears.
indicatingthatclumps,thissimulation
indicating
istribution
that
undergoing
thissimulation
fragmentais undergoing fragmentaagainof=8(bottom
panel),
but
that
do
not
includeadditional
tur2.01.0 2.32
What
appears
todensity
bewaves,happeningisthat
the density
imposedstate in whichtheinstability
isactingtosteadilytransportForthesimulations
that
do
not
fragment,
we
can
also
tion.
The
bottomshowstion.
spiralThedensitybottomwaves,
showsbutspiralnodensity
high
but
no
high
bulent forcing. The top panel shows a number of high-densityof k
5.0
0.1 3.90
turbulence is non-helical and, hence, does not necessarily
clumps,indicatingthat thisindicating
simulationis undergoing
fragmentaanglar
momentum(Lodato&Rice2004).
The
fragmentaconsider
otherproperties.
Figure4shows
the
Q-profilefor
0.1 aFigure9.
Qvalues
from
a
set
of
simulations
withcoolingtimes
clumps,thatclumps,thissimulation
indicatinghas
thatsettledthissimulationintoaquasihassettledintoaquasiFigure
showing
the baseline
simulation
for a coolingmanifest
itself 8.
as Figurea stressand does
not necessarilytransTS
3 RESULTS
tion. The bottom shows spiral density waves, but no high densityof the
=
10
and
with
turbulent
forcing
amplitudes,
but
allsamethere
time
ofmomentum.
= to7various(top
panel)
thesimulation
withthesteady
state
inthe
which
steadythe8 into
instability
state
which
primarily
instability
actingistransprimarily
actingtothetranstion boundary in these simulations ( ⇠ 4)clumps,
is similar
to thatthat=
simulation.
This
shows
that
we
run
the
port
angularItinitially
does, andhowever,heat
disc,
indicating
thissimulation
has settled
a quasi-in isforcedcoolingat a wavenumber
of
k
=
0.1.
They
all
settletoquasi-steady
RESULTStime bybutconsideringthat also includes
turbulent
forcing,
at twhich can be seen
Q, whichincreasesfrom forcedthe
statein portwhich
theinstability
is primarily
actingto transangularmomentum.
portangularmomentum.
nesimulations
3.1Gammie
Baseline
simulations
foundby(2001)(  3), but we steady
don’t
claim
that
simulation,
until
t
=
50,
withoutany
cooling.
We
then
turn
stateswithconstant
Q,
but
those
with
largerforcing
amplitudes
4
a wavenumber of k = 0.1 and with a forcing amplitude of f = H

- β=5

β=7

乱流なし

乱流ありで分裂する(波数0.1)

分裂- β=8, 10: 乱流ありでも分裂しない

Baseline simulations

port angular momentum.

ours is a precise representation of the fragmentation bound-just over 2 to about 5. The system therefore settles into aotodecay
states withhigherQbaseline
values andare,
1.0(bottom
panel).
simulationshows nomore
signsplitu
ofon the cooling and allow settle
it tototheTheprescribed
ratetherefore,• 波数の小さい（空間的には大きな）乱流で、冷却時間と比べて絶妙な強度（小さすぎるとあまり効かず、大きすぎると壊してしまう）であれば、乱流は重力不安定を促
進させる。

10

10
Ken Rice& Sergei
Nayakshin
Ken Rice& Sergei
Nayakshin• Hopkins & Christiansen (2013)で言われていた、乱流があるとβ≫5の円盤でさえ分裂する、というのはシミュレーションにより否定された。

• 3Dシミュレーション（乱流あり）- 乱流なしだとβ≤8で分裂、β>9は分裂しない- 初期条件はβ=10とする（理論通りだと乱流があれば分裂するはず）
乱流はvelocity kick fieldとして計算に含める（z方向は無視）- & Sergei
Nayakshinn Ken
Rice Rice& Sergei
Nayakshin
0 years

80 years

320 years

5

800 years（分裂しない）（100 AUでの公転周期は160 years）

• なぜHopkins & Christiansen (2013)と異なる結果が得られたのか？
- 原始惑星系円盤は冷却時間が長い(β≫1)ため、等温ではない。
‣ Toomre parameter Q=csΩ/πGΣと乱流強度が独立という仮定が成立しない。
‣ 乱流が超音速になるとQとcsが両方増加する。
‣ csが増加するのは、乱流による加熱とガスの冷却を釣り合わせるため。
‣ csが増加すると、最初は超音速であった乱流が亜音速になる。
→ 超音速乱流を維持することは難しい‣ そもそも超音速乱流ではα>1となるが、観測による円盤のαはもっと小さい（HL Tauでα~10-4 (Pinte et al. 2016)）

6

Amelia Bayo1,6 , Steven G. Parsons4 , Mariusz Gromadzki5 , Aurora Belén Aguayo Villegas1,6

41. Accretion signatures in the X-shooter spectrum of thesubstellar companion to SR12
1

Instituto de Fı́sica y Astronomı́a, Universidad de Valparaı́so, Av. Gran Bretaña 1111, 5030 Casilla, Valparaı́so, Chile;2
European Southern Observatory, Ave. Alonso de Cordova 3107, Casilla 19, 19001, Santiago, Chile; 3 Departamento
de Ciencias Fisicas, Facultad de Ciencias exactas, Universidad Andres Bello. Av. Fernandez Concha 700, Las Condes,
Santiago de Chile; 4 Department of Physics and Astronomy, University of Sheﬃeld, Sheﬃeld, 53 7RH, UK; 5 Warsaw
University Astronomical Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; 6 Núcleo Milenio FormaciónPlanetaria - Santamaría-Miranda,
NPF, Universidad de Valparaı́so,
Av. Gran
Bretaña Matthias
1111, Valparaı́so,Chile; 7 Millennium
Alejandro
Claudio
Cáceres,R. Schreiber,
Adam Nucleus
Hardy,“Protoplanetary Disks in ALMA Early Science”, Universidad de Valparaı́so, Chile

Amelia Bayo, Steven G. Parsons, Mariusz Gromadzki, Aurora Belén Aguayo Villegas
E-mail contact: asghaloth at gmail.com

About a dozen substellar companions orbiting young stellar objects or pre-main sequence stars at several hundred
au have been identified in the last decade. These objects are interesting both due to the uncertainties surrounding
their formation, and because their large separation from the host star oﬀers the potential to study the atmospheresof young giant planets and brown dwarfs. Here, we present X-shooter spectroscopy of SR 12 C, a ∼2 Myrs young
brown dwarf orbiting SR 12 at an orbital separation of 1083 au. We determine the spectral type, gravity, and eﬀective
temperature via comparison with models and observational templates of young brown dwarfs. In addition, we detectand characterize accretion using several accretion tracers. We find SR 12 C to be a brown dwarf of spectral typeL0±1, log g = 4 ± 0.5, an eﬀective temperature of 2600±100 K. Our spectra provide clear evidence for accretion at
a rate of ∼10−10 M⊙ yr−1 . This makes SR 12 one of the few sub-stellar companions with a reliable estimate for its
accretion rate. A comparison of the ages and accretion rates of sub-stellar companions with young isolated browndwarfs does not reveal any significant diﬀerences. If further26accretion rate measurements of a large number of substellar
companions can confirm this trend, this would hint towards a similar formation mechanism for substellar companionsat large separations and isolated brown dwarfs.

Accepted by MNRAS

au以上も離れている惑星/褐色矮星の形成メカニズムは謎に包まれている。降着• 100
http://arxiv.org/pdf/1712.09297
率を測ろうにも中心星に近すぎると観測できない。
The HR 4796A Debris System: Discovery of Extensive Exo-Ring Dust Material

Glenn, John0.013±0.007
H. Debes , Carol
A. ,Grady, AndrasGáspár, Thomas Henning , Dean C.
12 Schneider
C (~ 2 Myrs,
Msun
1083±217
au)のX-shooter分光観測をした。
• SR
Hines , Marc J. Kuchner , Marshall Perrin , John P. Wisniewski
Steward Observatory
and the Department of Astronomy, The University of Arizona 933 North Cherry Avenue,companionや孤立した若い褐色矮星の質量降着率と明確な違いはない• 他のsubstellar
Tucson, AZ 85721 USA; Space Telescope Science Institute 3700 San Martin Drive, Baltimore, MD 21218 USA;
1

2

2

5

3

2

1

4

6

1

2

3

Eureka Scientific 2452 Delmer, Suite 100, Oakland, CA 96002 USA; 4 Max-Planck-Institut für Astronomie Königstuhl
ため、これらの形成メカニズムは似ているかもしれない。
5

17, 69117, Heidelberg, Germany; NASA/Goddard Space Flight Center, Exoplanets & Stellar Astrophysics Laboratory
Code 667, Greenbelt, MD 20771 USA; 6 H. L. Dodge Departmentof Physics and Astronomy, University of Oklahoma
7
440 West Brooks Street, Norman, OK 73019 USA

Hβ

Hγ

8

Hα

OI
CaII triplet

A. Santamarı́a-Miranda et al.

Table 2. Accretion rates measurements for each emission line
1)

Line

Wavelength (Å)

Ṁacc (M yr

H↵
H↵
H
H
H11
OI
CaII
Pa

6562.8a6562.8b
4861
4341
3771
8446
8544.212800(4.90 ± 5.64) ⇥ 10
(1.16 ± 0.88) ⇥ 10
(1.04 ± 0.74) ⇥ 10
(7.01 ± 4.98) ⇥ 10
(2.93 ± 2.06) ⇥ 10
(1.02 ± 0.75) ⇥ 10
4.83 ⇥ 10 10(3.95 ± 2.94) ⇥ 10

11
11
10
11
11
10
13

Notes
a Measured as explained in Section 4.1
b Measured as explained in Section 4.5

ure 6. Bottom pannel: SR 12 C uvb spectra, with the Balmer series labeled. Top left pannel: H
emission line.

mass and radius of SR 12 C as above. The given uncernty does not take into account the standard deviation of
linear fit from Herczeg & Hillenbrand (2008) as theseues are not available to us. However, the included untainties related to the mass and radius of SR 12 C likely
minate the error estimate.

Accretion rates derived from other Hydrogenlines

e entire Balmer and Paschen emission line series hasn used to measure accretion (Ferguson 1997). We used
(Lacc /L ) = b+a⇥log(Lline /L ) to relate the accretion
minosity and the luminosities of the lines (see Fang et al.
9; Herczeg & Hillenbrand 2008 and Table 6 and 8 fromliaco et al. 2012). We applied this method to the H↵,• SEDフィッティングより: L0±1, log g = 4±0.5,

emission line. Top rightpannel:Figure5. The X-shooter spectrum (covering the visible and J, H, K bands) of SR 12 C shows clear evidence for ongoing accr(bottom pannel). In the top panels we highlight the strong H↵ emission line (left) and CaII triplet and OI line (right).

Figure 9. Based on Fig. 6 from Bowler et al. (2011). Mass ac-

2600±100 K

H , H and H11 lines identified in the spectrum4.4of SR 12 C.
rate basedon theOI line
SR 12 C. First, we calculate the line luminosity in a sim
cretion rates vs Accretion
stellar mass.
Triangles
shows
upper limits. We
Figure 6 shows the UV spectrum (bottom panel) of SR 12 C.
fashion as in the previous subsection. We then use Fi
2
TAU b, GSC 06214-00210 b, CT Cha B and SR 12 fromC Herczeg & Hillenbrand (2008) to estimate the aEquivalent width, flux line (W/minclude), log(LPFW
a / L ) and acAdditional emission lines that can be related to accretion
cretion luminosity in solar unitsas
areblack
given squares.
in Table
1.DataThe
blue etarerates given intion luminosity of SR 12 C (note that the coefficients fo
are those
fromshown
OI (e.g. in
Joergens
al. accretion2012). Our X-shooter
resulting accretion rates can be found
in
Table
2.
shown linear correlation are not given) and finally obta
Natta et al. spectrum(2006) whichmeasuredusing
P a linewhile
those incontainsare
a strong
OI 8446
emissionthat weaccretion rate of Ṁ = (2.85 ± 1.98) ⇥ 10 10 M yr 1 u
here use to add another measurement of the accretion rate in• それぞれの輝線の幅から質量降着率を求めた- Paβを除くとほぼ一致- ~ 10-10Msun/yr

red are measured with Br . Purple symbols represent measurements from Herczeg et al. (2009) and green pentagons have been
5 DISCUSSION
obtained from Muzerolle et al. (2005). We selected only class II
objectsNatta
et al. (2006). Keep in mind that this plot comWe have measured accretion rates
for thefrom
substellar
companion (SSC) to the young binarybines
SR 12 usingeight di↵erent
accretion
rate measurements with di↵erent methodologies
methods. In general, the measured accretion rates are veryof objects with di↵erent ages. However, it seem that SSCs roughly
similar and agree with each other. However, the accretionthe correlation
rate measurement based on Pa follow
is significantly
lower than between mass and accretion rate of young
the other values we obtain, as illustrated
Fig. 7. As8all valstellar inobjects.
ues agree very well except the one obtained from Pa , the

MNRAS 000, 1–?? (2CT Cha b Wu et al. (2015b) estimate a relatively high accre1
tion
rate
of
log
Ṁ
(M
yr
)=-9.22 from interpreting excess
Table 2. Accretion rates measurements for each emission line
emission in the r’ -band as caused by H↵ emission. This issimilar to the value we obtained for SR 12 C using all techLine
Wavelength (Å)
Ṁacc (M yr 1 )niques (i.e. log Ṁ (M yr 1 ) ⇠ 10) except the one basedH↵6562.8a(4.90 ± 5.64) ⇥ 10 11
on Pa . For GSC 6214-210 B and FW Tau B published acb
11H↵
6562.8(1.16 ± 0.88) ⇥ 10
cretion rate measurements are based on Pa emission andH
4861(1.04 ± 0.74) ⇥ 10 10
the obtained values are relatively low, i.e. 10.7 ± 1.3 andH
4341(7.01 ± 4.98) ⇥ 10 11
11.0 ± 1.3 respectively. Thus one might get the impresH113771(2.93 ± 2.06) ⇥ 10 11
sion that indeed be that Pa accretion rates of SSCs are
OI
8446(1.02 ± 0.75) ⇥ 10 10
systematically lower. However, Joergens et al. (2013) invesCaII
8544.2
4.83 ⇥ 10 10
tigated accretion on the isolated brown dwarf OTS 44 andPa
12800(3.95 ± 2.94) ⇥ 10 13
find that the rate derived from Pa is significantly larger
Notes
than the one obtained from H↵. Thus, further simultanea Measured as explained in Section 4.1
ous measurements of accretion rates based on di↵erent lines
b Measured as explained in Section 4.5
are needed before any firm conclusions can be drawn. Nevertheless, we conclude that the higher values we obtain for
SR 12 C with 8 di↵erent methods probably represent a morereliable estimate of the accretion rate in SR 12 C. The cur（左）> 100 au, < 15 MJupのsubstellar companionの年齢、質量、降着の有無
rentlyavailable accretion rate measurements for SSCs as
Figure 8. Age in Myrs versus mass of the SSCs listed in 2. Red
Figure 9. Based on Fig. 6 from Bowler et al. (2011). Mass aclisted
in rates
Tablevs3.stellar
In addition
the threeobjects
indicates that the objects are accreting, green represents non accretion
mass. toTriangles
shows
uppermentioned
limits. We<
10Myrsで降着が見られる
above,
accretion
rateChaestimates
basedcreting objects and in blue we show SSCs without information
includetwo
FWmoreTAUobjects
b, GSC have
06214-00210
b, CT
B and SR
12 C
onaccretion
excesscontinuum
emissionon their accretion status. It seems that objects younger than 10
as theblacksquares. luminosity
Data shownfrom
in blue
are accretion
rates
given inMyrs
are usually accreting.
Natta etet al. (2006)(Zhou
2014). which are measured using P a while those in（右）substellar
companion、孤立した低質量星、褐色矮星の降着率vs質量
red InareFig.
measuredwith Br. Purplesymbols
represent systems
measure8 we show
accretingandnon accreting

•

-

•

ments from Herczeg et al. (2009) and green pentagons have beenobtained from Muzerolle et al. (2005).
We selected
only (2017)class II
MNRAS
000, 1–??
objects from Natta et al. (2006). Keep in mind that this plot combines accretion rate measurements with di↵erent methodologies
of objects with di↵erent ages. However, it seem that SSCs roughly
follow the correlation between mass and accretion rate of youngstellar objects.

- 4つの黒四角がsubstellar companion- Natta et al. (2006)はclass II天体のみプロット- 若い星とsubstellar companionは同じ相関を持つ

empirical relation used for Pa seems to provide accretion
降着率が似ていることから、substellar
companionの形成シナリオは、コア集積+散• Figure
7. Accretion rates obtained from di↵erent emission lines.
rates that are perhaps less reliable. This hypothesis is someApparently, the obtained values agree very with each other. The
only exception is the accretion derived from Pa which is significantly lower.

what supported by looking at the accretion rates obtained
乱ではなく、円盤不安定性かprotostellar cloudsの収縮だと思われる。
9for other SSCs. Accretion rate measurements exist for threeother SSCs(CT Cha b,GSC 6214-210 B and FW Tau B). For

Glenn Schneider1 , John H. Debes2 , Carol A. Grady3 , Andras Gáspár1 , Thomas Henning4 , Dean C.
Hines2 , Marc J. Kuchner5 , Marshall Perrin2 , John P. Wisniewski6

42.
The
HR 4796A
DebrisSystem:Discovery
Extensive
Steward
Observatory
and the Departmentof Astronomy,
The Universityof Arizona 933 ofNorth
Cherry Avenue,
Tucson, AZ 85721 USA; Space Telescope Science Institute 3700 San Martin Drive, Baltimore, MD 21218 USA;
Exo-Ring
Eureka Scientific 2452 Delmer, Suite
100, Oakland, CA Dust
96002 USA;Material
Max-Planck-Institut für Astronomie Königstuhl
1

2

3

4

17, 69117, Heidelberg, Germany; 5 NASA/Goddard Space Flight Center, Exoplanets & Stellar Astrophysics Laboratory
Code 667, Greenbelt, MD 20771 USA; 6 H. L. Dodge Department of Physics and Astronomy, University of Oklahoma
Glenn Schneider, John H. Debes, Carol A. Grady, Andras Gáspár, Thomas Henning,
440 West Brooks Street, Norman, OK 73019 USA

Dean C. Hines, Marc J. Kuchner, Marshall Perrin, John P. Wisniewski

E-mail contact: gschneider at as.arizona.edu

The optically and IR bright, and starlight-scattering, HR 4796A ring-like debris disk is one of the most (and best)studied exoplanetary debris systems. The presence of a yet-undetected planet has been inferred (or suggested) fromthe narrow width and inner/outer truncation radii of its r = 1.′′ 05 (77 au) debris ring. We present new, highly sensitive,
Hubble Space Telescope (HST) visible-light images of the HR 4796A circumstellar debris system and its environment
over a very wide range of stellocentric angles from 0.′′ 32 (23 au) to ∼15′′ (1100 au). These very high contrast images wereobtained with the Space Telescope Imaging Spectrograph (STIS) using 6-roll PSF-template subtracted coronagraphy
suppressing the primary light of HR 4796A and using three image plane occulters and simultaneously subtracting thebackground light from its close angular proximity M2.5V companion. The resulting images unambiguously reveal thedebris ring embedded within a much larger, morphologically complex, and bi-axially asymmetric exoring scattering
structure. These images at visible wavelengths are sensitive to, and map, the spatial distribution, brightness, and radialsurface density of micron size particles over 5 dex in surface brightness. These particles in the exo-ring environment
may be unbound from the system and interacting with the local ISM. Herein we present a new morphological andphotometric view of the larger than prior seen HR 4796A exoplanetary debris system with sensitivity to small particles
at stellocentric distances an order of magnitude greater than has previously been observed.
Accepted by AJ

4796Aはリング状(77 au)のデブリ円盤を持ち、惑星があると考えられている。
• HR
http://arxiv.org/pdf/1712.08599
Telescope Imaging Spectrograph (STIS) 6-roll PSF template subtracted• HSTのSpace
The Physics of Protoplanetesimal Dust Agglomerates. X. High-velocity Collisions between Small andLarge Dust Agglomerates as Growth Barrier
coronagraphy(6R/PSFTSC)による観測と、近くのM型伴星による光を除去して、HR
Rainer Schräpler1 , Jürgen Blum1 , Sebastiaan Krijt2 and Jan-Hendrik Raabe1

14796Aのデブリリングを詳細に調べた。
Institut für Geophysik und extraterrestrische Physik, Technische Universität Braunschweig, Mendelssohnstr. 3, D-

38106 Braunschweig, Germany; 2 Department of the Geophysical Sciences, The University of Chicago, 5734 South
Ellis Avenue, Chicago, IL 60637, USA
E-mail contact: r.schraepler at tu-bs.de10Dist. c Age dDiskInitial HST Disk ImagingeReference
steepness from: (a) the ansal peak at r = 1.04" to r = 1.6"(pc)steeply (Myr)declining asLr-7.8(goodnessofInstrument
fit:
IR/L
star
-3.6
-2.7
RHR
= 0.978),< r < 4.5" withSB ~ r A0V(R = 0.977)72.8, and (c)84.5"
with SB ~ r NICMOS(R
4796 (b)A 1.6" 5.774+0.012± 2< r < 12"
0.0042
Schneider et al. 1999
TargetVmag bB-V bTable 1 – HR 4796 System Components(A) Primary Star and Debris Disk
Spec a= 0.967). (c) is beyond the angular extent of the exo-ring material on the SW major axis. The
photometric uncertainties (internal measurement errors)
the individual
points are closely(B) ofM-StarCompanion
approximated by the local dispersion
aroundthese sectionalradial power-law ffits: a few percent atabbCompanion
Spec
Vmag
Separation/(2015)
ReferencesΔMag
r < 4.5" on both sides of the star, and increasing to ~ 15% at 12" on theP.A.
NE sideof the star at the
HR 4796
B periphery
M2.5of the debris
13.3system 7.5
7.92"
0.02" / 29.81°
Jura et al. 1995, Lagrange et al. 2012outer
detectableas enumerated
in §±6.1.
afrom Houk et al. 1982; VizieR on-line catalog III/80; b from Hog et al. 2000; VizieR on-line catalog I/259
c
N catalog I/311
Hipparcos parallactic distance from Van Leeuwen 2007; VizieR on-linedAge estimation from Stauffer et al. 1995. e IR (4µm – 1 mm) excess from Pascual et al. 2016.
f
As determined from this work with: internal measurement uncertainty in PA
E ± 0.04° and absolute uncertainty of
the HST guide-star frame orientation ≈ ± 0.1°.

The ≈ 76° inclined ring-like debris disk of HR 4796A was 長軸方向の表面輝度
initially imaged in 1998 by
Schneider et al. (1999) in near-IR scattered-light at 1.1 and 1.6 µm with HST/NICMOS. These• 緑線が長軸方向
scattered-light discovery images were followed-up with higher fidelity(but relatively shallow
depth) observations in 2001 with HST/STIS 2-roll coronagraphy in visible light (Schneider et al.
右上の白ドットはHR
2009) that better revealed and detailed its steeply “edged” ring-like •nature. Subsequent NICMOS
4796 Bの位置
observations (Debes et al. 2008), with multi-band diagnostic filters, suggested
the possibility ofradiationally evolved complex organic materials in the ring. As a bright, geometrically favorable,• 左右で軌道長半径依存性
debris disk target for episodically improving ground-based near-IR
and narrow-field angle
adaptive optics (AO) technologies, the archetypical HR 4796A system が異なる
was extensively observedwith capability-driven focus on the debris ring itself. E.g.: With Subaru/HiCIAO (Thalmann et al.
2011), VLT/NaCO (Lagrange et al. 2012), Gemini/NICI (Wahhaj et al. 2014), Gemini/GPI(Perrinal. 2015),(Rodigas,
Perrinet al. ibid suggested that the GPIFig.
10. Plot ofeta 1-pixel(0.05077")Magellan/MagAO
wide radial SB profile of the HR
4796A debris 2015).
structure withcorresponding
image
and powerpolarized
law fits. Alongintensity
the extension signatureof the debris-ring
major
axis
(green ringline on may
image), implicate
starlightdetectedofthedebrisit actually being opticallyscattering debris is robustly detected to r ≈ 12" on the NE side of the star, but on the SW side to only ≈ 4.5" with an
11 disk
abrupttruncationor outer edge self-shadowed
boundary as the peripheryonof thedebrissidestructureto the south.
thick
or partiallyoneof“bends”
themajor axis. Ground-based observations ofHR 4796AはLocal
Interstellar
Cloud (LIC)とHyades Cloudの中を進んでいるため、
bow shockが生じている。
もしくはHR
4796Bの輻射圧が影響している。
12

発表しなかった論文のアブストラクト

with temperature lower than 10000 K. The true level of the veiling continuum can be derived by measuring the weakest
photospheric lines with equivalent widths down to ≈10 mÅ. A limited spectral resolution and/or low signal-to-noise
ratio results in overestimation of the veiling continuum. In the three very active stars, the veiling continuum is aminor contributor to the observed veiling, while the major contribution comes from the line-dependent veiling.
Accepted by A&A
http://arxiv.org/pdf/1712.03784

Episodic accretion in binary protostars emerging from self-gravitating solar mass cores

R. Riaz, S. Vanaverbeke, and D.R.G. Schleicher
36.
Episodicaccretionin binary protostars emerging from
Departamento de Astronomı́a, Facultad Ciencias Fsicas y Matemáticas, Universidad de Concepción, Av. Esteban
self-gravitating
solar
cores Plasma-Astrophysics,
Iturra s/n Barrio Universitario,
Casilla 160-C, Concepción,Chile; mass
Centre for mathematical
1

2

1

1

2

Department of Mathematics, KU Leuven, Celestijnenlaan 200B, 3001 Heverlee, Belgium

R. Riaz, S. Vanaverbeke, and D.R.G. Schleicher
E-mail contact: rriaz at astro-udec.cl
Observations show a large spread in the luminosities of young protostars, which are frequently explained in the context
of episodic accretion. We here test this scenario using numerical simulations following the collapse of a solar mass
molecular cloud using the GRADSPH code, varying the strength of the initial perturbations and the temperature of
the cores. A specific emphasis of this paper is to investigate the role of binaries and multiple systems in the context of
episodic accretion, and to compare their evolution to the evolution in isolated fragments. Our models form a varietyof low mass protostellar objects including single, binary and triple systems with binaries more active in exhibiting
episodic accretion than isolated protostars. We also find a general decreasing trend for the average mass accretion rateover time, suggesting that the majority of the protostellar mass is accreted within the first 105 years. This result canpotentially help to explain the surprisingly low average luminosities in the majority of the protostellar population.
Accepted by A&A
http://arxiv.org/pdf/1712.09646
13

de Córdova 3107, Vitacura 763 0355, Santiago, Chile; European Southern Observatory, Alonso de Córdova 3107,
molecular clouds located in the Galactic Center, which are characterized by diﬀerent types of chemistry. PN is
Vitacura Casilla 763 0355, Santiago, Chile; 5 Centro de Astrobiologı́a (INTA-CSIC), Ctra. de Ajalvir Km. 4, Torrejóndetected in five out of seven sources, whose chemistry is thought to be shock-dominated. The two sources with PN
de Ardoz, 28850 Madrid, Spain; 6 Observatorio Astronómico de Quito, Escuela Politécnica Nacional, Av. Grannon-detections correspond to clouds exposed to intense UV/X-rays/cosmic-ray radiation. PO is detected only towards
Colombia S/N y Av. Diez de Agosto, Quito 170403, Ecuador; 7 Department of Physics and Astronomy, UCL, Gower
the cloud G+0.693−0.03, with a PO/PN abundance ratio of ∼1.5. We conclude that P-bearing molecules likely form
St., London, WC1E 6BT, UK; 8 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany;in shocked gas as a result of dust grain sputtering, while are destroyed by intense UV/X-ray/cosmic ray radiation.
9
Space
Telescope
Science Institute, 3700
MartinBaltimore,
MD21218, USA
V. M.
Rivilla,I. Jiménez-Serra,
S. SanZeng,
S. Drive,
Martín,J. Martín-Pintado,
J. Armijos-Abendaño,
Accepted by MNRAS Letters
E-mailat arcetri.astro.it
S. Viti, R. contact:
Aladro,rivilla
D. Riquelme,
M. Requena-Torres, D. Quénard, F. Fontani, and M. T. Beltrán
https://arxiv.org/pdf/1712.07006
Phosphorus (P) is one of the essential elements for life due to its central role in biochemical processes. Recent searches
have shown that P-bearing molecules (in particular PN and PO) are present in star-forming regions, although theirformationroutes remain
poorlymoving
understood.
In thisLetter,V1062
we report
observations of PN and PO towards seven
A new compact
young
group
aroundScomolecular clouds 1,2
located in the Galactic Center, which are characterized by diﬀerent types of chemistry. PN is
Siegfried Röser , Elena Schilbach1,2 , Bertrand Goldman3,4 , Thomas Henning3 , Attila Moor5 , and Alizdetected 5,6
in five out of seven sources, whose chemistry is thought to be shock-dominated. The two sources with PN
Derekasnon-detections correspond to clouds exposed to intense UV/X-rays/cosmic-ray radiation. PO is detected only towards
1
Zentrumfür Astronomiewith
der Universität
Heidelberg,ratio
Landessternwarte,Königstuhl
69117 Heidelberg,
Germany;thecloud G+0.693−0.03,
a PO/PN abundanceof ∼1.5. We conclude
that 12,P-bearingmolecules likely
form
2
3
für asAstronomie
Universität
Heidelberg,
69120
Heidelberg, Germany;
Maxin Zentrum
shocked gas
a result ofder
dust
grain sputtering,
whileMönchhofstraße
are destroyed by12-14,
intenseUV/X-ray/cosmic
ray radiation.
4
Planck-Institut für Astronomie, Königstuhl 17, 69117 Heidelberg, Germany; Observatoire astronomique de StrasAcceptedby MNRAS
Lettersbourg, Universitéde Strasbourg- CNRS UMR 7550, 11 rue de l’Université, 67000, Strasbourg, France; 5 Konkoly
https://arxiv.org/pdf/1712.07006
Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, H-1121 Budapest,
Konkoly Thege Miklós út 15-17, Hungary; 6 ELTE Eötvös Lornd University, Gothard Astrophysical Observatory,
Szombathely,
Hungary Elena Schilbach, Bertrand Goldman, Thomas Henning, Attila Moor,
Siegfried Röser,38. Phosphorus-bearing molecules in the Galactic Center

39. A new compact young moving group around V1062 Sco
A new compact young moving group around V1062 Sco

E-mail contact: roeser at ari.uni-heidelberg.de and Aliz Derekas
Siegfried Röser1,2 , Elena Schilbach1,2 , Bertrand Goldman3,4 , Thomas Henning3 , Attila Moor5 , and AlizAims. We5,6are searching for new open clusters or moving groups in the Solar neighbourhood.
DerekasMethods. We used the Gaia-TGAS catalogue, cut it into narrow proper motion and parallax slices and searched for
1
Zentrum spatialfür Astronomie
der Universität
Heidelberg,
Landessternwarte,Königstuhl
12, over-densities
69117 Heidelberg,
Germany;
significant
over-densities
of
stars
in
each
slice.
We
then
examinedstars
forminginopticaland2
3
Zentrum fürcolour-magnitude
Astronomie der diagrams
Universität
12-14,
69120
Heidelberg,
Germany; Maxnear-infrared
to Heidelberg,
determine ifMönchhofstraße
they are compatiblewithisochronesof a cluster.
4
Planck-Institutfür Astronomie,
17, 69117
Heidelberg,
astronomiquede StrasResults. We detected
a hitherto Königstuhl
unknown moving
group
or cluster Germany;in the UpperObservatoire
Centaurus Lupus(UCL) 5sectionofbourg,
Universitéde Strasbourg
CNRS UMR
11 ruel’Université,
67000,ofStrasbourg,
France;Konkoly
the Sco-Cen
OB-association
at a- distanceof 1757550,
pc fromthedeSun.
It is a group
63 co-moving
stars withages
Observatory,
Research
Centrefor
Astronomy
andEarth
Sciences,
HungarianAcademyof
Sciences,
H-1121
Budapest,of less than 10 to about 25 Myr. For the brightest
stars, which are present in the Gaia-TGAS catalogue the mean
6
Konkoly
Thege
Miklósút
15-17,
Hungary;
ELTE
Eötvös
Lornd
University,
Gothard
Observatory,
diﬀerence between kinematic and trigonometric distancemoduliis about
0:01 magwith aAstrophysicalstandard deviationof 0.11
Szombathely,
Hungarymag. Fainter cluster
candidates are found in the HSOY catalog, where no trigonometric parallaxes are available. For
a subsetcontact:of our candidatestars, we obtained radial velocity measurements at the MPG/ESO 2.2-metre telescope in
E-mailroeser at ari.uni-heidelberg.de
La Silla. Altogether we found twelve members with confirmed radial velocities and parallaxes, 31 with parallaxes or
Aims.
We are searching for new open clusters or moving groups in the Solar neighbourhood.
radial velocities, and 20 candidates from the convergent point method. The isochrone masses of our 63 members range
Methods. We used the Gaia-TGAS catalogue, cut it into narrow proper motion and parallax slices and searched forfrom 2.6 M⊙ to 0.7 M⊙ .
significant spatial
over-densities of stars in each slice. We then examined stars forming over-densities in optical andAccepted by A&A
near-infrared
colour-magnitude diagrams to determine if they
14 are compatible with isochrones of a cluster.

1

Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA, 02138, USA; 2 Square Kilometre
Array Organisation, Jodrell Bank Observatory, Lower Withington, Cheshire SK11 9DL, UK; 3 Department of Physics,
State University of New York at Fredonia, 280 Central Ave, Fredonia, NY 14063, USA; 4 National Research Council
Canada, 5071 West Saanich Road, Victoria BC Canada, V9E 2E7; 5 Department of Physics and Astronomy, Universityof Victoria, PO Box 355, STN CSC, Victoria BC Canada, V8W 3P6; 6 Departmento de Astronomı́a, Facultad Ciencias
Fı́sicas y Matemáticas, Universidad de Concepción, Concepción, Chile 0000-0003-2300-8200; 7 Max-Planck-Institut
8für
Astronomie(MPIA),
Königstuhl
17,
D-69117
Heidelberg,
Germany;
Homer
Dodge Departmentof Physics
andSarah I. Sadavoy, Eric Keto, Tyler L. Bourke, Michael
M. L.
Dunham,PhilipC.
Myers,
Astronomy, University of Oklahoma, 440 W. Brooks Street, Norman, OK 73019, USA; 9 Leiden Observatory, Leiden
Ian W. Stephens,
Francesco,
Kristi Webb, Amelia Stutz, Ralf Launhardt,
University,
P.O. Box 9513,James
2300-RA Di
Leiden,The Netherlands

40. Intensity-Corrected Herschel Observations of Nearby
Isolated Low-Mass Clouds

E-mail contact: sarah.sadavoy at cfa.harvard.edu

John Tobin

We present intensity-corrected Herschel maps at 100 µm, 160 µm, 250 µm, 350 µm, and 500 µm for 56 isolated
low-mass clouds. We determine the zero-point corrections for Herschel PACS and SPIRE maps from the Herschel
Science Archive (HSA) using Planck data. Since these HSA maps are small, we cannot correct them using typicalmethods. Here, we introduce a technique to measure the zero-point corrections for small Herschel maps. We useradial profiles to identify oﬀsets between the observed HSA intensities and the expected intensities from Planck. Most
clouds have reliable oﬀset measurements with this technique. In addition, we find that roughly half of the clouds
have underestimated HSA-SPIRE intensities in their outer envelopes relative to Planck, even though the HSA-SPIRE
maps were previously zero-point corrected. Using our technique, we produce corrected Herschel intensity maps forall 56 clouds and determine their line-of-sight average dust temperatures and optical depths from modified black
body fits. The clouds have typical temperatures of ∼14–20 K and optical depths of ∼10−5–10−3 . Across the wholesample, we find an anti-correlation between temperature and optical depth. We also find lower temperatures thanwhat was measured in previous Herschel studies, which subtracted out a background level from their intensity maps tocircumvent the zero-point correction. Accurate Herschel observations of clouds are key to obtain accurate density andtemperature profiles. To make such future analyses possible, intensity-corrected maps for all 56 clouds are publicly
available in the electronic version.
Accepted by ApJ
http://arxiv.org/pdf/1712.00017

Accretion signatures in the X-shooter spectrum of the substellar companion to SR12
Alejandro Santamarı́a-Miranda1,2,6,7 , Claudio Cáceres3,6 , Matthias R. Schreiber1,6,7 , Adam Hardy1,7 ,Amelia Bayo1,6 , Steven G. Parsons4 , Mariusz Gromadzki5 , Aurora Belén Aguayo Villegas1,61

Instituto de Fı́sica y Astronomı́a, Universidad de Valparaı́so, Av. Gran Bretaña 1111, 5030 Casilla, Valparaı́so, Chile;2
3
European Southern Observatory, Ave. Alonso de Cordova
15 3107, Casilla 19, 19001, Santiago, Chile; Departamento