星形成ゼミ
2018/02/16
SFN #300 #51-#60
長谷川哲夫（国立天文台）

1

Dust Concentration and Emission in Protoplanetary Disks Vortices
Anibal Sierra1 , Susana Lizano1 and Pierre Barge2

300
Dust Concentra:on and Emission in Protoplanetary Disks Vor:ces
Instituto de Radioastronomı́a y Astrofı́sica, UNAM, Apartado Postal 3-72, 58089 Morelia Michoacán, México
Anibal Sierra1,
SusanaLizano1and Pierred’Barge2
2#51
Aix Marseille
Université,
CNRS,
Laboratoire
Astrophysique de Marseille, UMR 7326, 13388, Marseille, France

1

ApJ 850,115 (2017)
E-mail contact: a.sierra at irya.unam.mx

We study the dust concentration and emission in protoplanetary disks vortices. We extend the Lyra-Lin solution forthe dust concentration of a single grain size to a power-law distribution of grain sizes n(a) ∝ a−p . Assuming dust
conservation in the disk, we find an analytic dust surface density as a function of the grain radius. We calculate
the increase of the dust to gas mass ratio ϵ and the slope p of the dust size distribution due to grain segregation
within the vortex. We apply this model to a numerical simulation of a disk containing a persistent vortex. Due tothe accumulation of large grains towards the vortex center, ϵ increases by a factor of 10 from the background disk
value, and p decreases from 3.5 to 3.0. We find the disk emission at millimeter wavelengths corresponding to synthetic
observations with ALMA and VLA. The simulated maps at 7 mm and 1 cm show a strong azimuthal asymmetry. This
happens because, at these wavelengths, the disk becomes optically thin while the vortex remains optically thick. The
large vortex opacity is mainly due to an increase in the dust to gas mass ratio. In addition, the change in the slope ofthe dust size distribution increases the opacity by a factor of 2. We also show that the inclusion of the dust scattering
opacity substantially changes the disks images.
3
39
Accepted by ApJ
https://arxiv.org/pdf/1710.03324

ALMA Reveals Transition of Polarization Pattern with Wavelength in HL Tau’s Disk
Ian W. Stephens1 , Haifeng Yang2 , Zhi-Yun Li2 , Leslie W. Looney3 , Akimasa Kataoka4,5 , Woojin
Kwon6,7 , Manuel Fernández-López8 , Charles L. H. Hull5,9,10, Meredith Hughes11 , Dominique SeguraCox3,12 , Lee Mundy13 , Richard Crutcher3 and Ramprasad Rao14
1

Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA, USA
Astronomy Department, University of Virginia, Charlottesville, VA 22904, USA
3
Department of Astronomy, University of Illinois, 1002 West Green Street, Urbana, IL 61801, USA
4
National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan
5
NAOJ FellowFigure 1. Gas surface density (left) and gas temperature (right) of a disk containing a large scale vortex (after 100 vortex
rotations) .
6
Korea Astronomy
and Space Science Institute, 776 Daedeok-daero, Yuseong-gu, Daejeon 34055, Korea
7
Korea Universityof Science
Gajang-ro,
Yuseong-gu,
Daejeon 34113,
Integrating
the eqs. and(6, 7)Technology,
in all the disk, 217
the massofdust
particles, we approximate
the dustKorea
mass of the8
Instituto Argentino
Radioastronomı́a,(CCT-La
Plata,grain
CONICET;CICPBA),
C.C. No.
5,i.e.,1894,Villa Elisa,
gas and thede
massof dust with grain radiusa arepopulation withsize a by the cumulativemass,c
Md (a) ⇠ Md (a). Thus, equation (12) can be rewritten
Argentina
2

5

6

Figure 2. Dust particle accumulation within the vortex. Left panel:
Slope p of the power law fit to the dust particle size
8

distribution n(a)da / a p da (see text). Middle panel: fractional standard deviation
p/p. Right panel: Map of the local dust to gas mass ratio (normalized to ✏).

and 3 mm for the configurations C4-6 (Band 7) and C48 (Band 4). The precipitable water vapor (PWV) was
set to 1.3 and 1.5 mm in the Band 7 and Band 4 configuration, respectively. These antenna arrays provide asimilar angular resolution ✓B ⇠ 0.00 10, which corresponds
to 12 AU at the assumed distance of 120 pc.
We also simulate VLA observations at 7 mm and 1
cm with CASA in the A configuration (the highest resolution). The FWHM beam of these configurations are✓B ⇠ 0.00 043 and 0.00 19, respectively. In both simulations
the PWV was set to 1.5 mm.
Figure 3 shows the ALMA simulated maps in the upper panels (left: 1 mm, right: 3 mm). Lower panels
show the VLA maps (left: 7 mm, right: 1 cm). The
vortex emission starts to show at 3 mm and becomesvery evident at larger wavelengths, in the VLA images.
The left panel of Figure 4 shows the SED of the disk(yellow solid line). We also include the SED of an axisymmetric disk (black dashed line) and the contributionof a 2M star (blue dotted line). Note that the SED is
not modified by the emission associated to the vortex,
except at mm wavelengths where the disk becomes optically thin. The right panel shows the ratio between theflux from the vortex disk and the axisymmetric disk.
The ratio has a maximum value of 1.25 at ⇠ 7 mm.
The reason for this behaviour is that the numerical simulation of B17 uses an adiabatic equation of state, thus,
Figure 3. Upper panels: Simulated ALMA images at 1 mm (left) and 3 mm (right). Lower panels: Simulated VLA images atthe disk temperature increases with the surface density;
7 mm (left) and 1 cm (right). The beam is shown in the left bottom corner of each image.
therefore, the flux associated to vortex is larger than theingredient to produce strong azimuthal
asymmetries
in same1 mm,withamax =in⌫ > ⌫ for mm wavelengths;
flux
from theregionthe axisymmetricdisk. thus,
the dust thermal emission at 7 mm and 1 cm. Note,however, that jets and photoevaporated disk winds will

scattering increases the total opacity and the albedo islarge !⌫ > 0.8 (see Figure A2).

centration of the dust particles around the vortex center,as shown in the Figure 3. The strong disk asymmetry
comes from the enhancement of the dust to gas mass ratio due to size segregation inside the vortex, the opacity
is increased by a factor of 10 at the vortex center. The
smaller slope of the size distribution p within the vortex
also increases the opacity at mm wavelengths (see Appendix A). For example, for p = 3 and = 7 mm, theopacity increases by a factor of ⇠ 1.8 compared with thecase p = 3.5. Both e↵ects (higher dust to gas mass ratio
and lower slope of the size distribution) are responsible
for the increase of the opacity by a factor of 18 at thevortex center.
In order to explore the e↵ect of dust concentration in
the vortex, we made maps of the disk model at 1, 3,
7 mm and 1 cm assuming a constant dust to gas mass
ratio of ✏ = 0.0131 and a fixed slope of the dust particle
size distribution p = 3.5 throughout the disk. These
maps are shown in Figure 5 with the same observationalparameters described in the subsection §3.2.
The ALMA maps at 1 and 3 mm do not change muchfrom those in Figure 3 where segregation is taken into
account. However, there is a dramatic change for theVLA observations at 7 mm and 1 cm. If the dust segregation is not included, the high contrast between thevortex and the disk vanishes.
In the model with dust segregation, the vortex domiFigure 5. Upper panels: Simulated ALMA images at 1 mm (left) and 3 mm (right). Lower panels: Simulated VLA images atnates
the emission at long wavelengths. For those wave7 mm (left) and 1 cm (right). These models do not include dust segregation inside the vortex (see text). The beam is shown in
the lengths
left bottomthe
cornerof eachimage.
disk
becomes
optically thin while the vortex
remains
optically
thickdue
to the increasethat
ofatthedust
7 mm and 1 cm the vortex emission will dominate
compared to the previous model (amax = 1 mm). Howopacity
in thisFigure6 shows
logarithmof but will have a small spatial extent,
disk emission,ever,
the maximum
dustregion.
to gas mass
ratio within
the the the9National Astronomical Observatory of Japan, NAOJ Chile Observatory, Alonso de Córdova 3788, Oﬃce 61B, 7630422,
Vitacura, Santiago, Chile
10
Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura, Santiago, Chile
300
ALMA
Reveals Transi:onof Polariza:on
PaUernwithMiddletown,
Wavelengthin HL Tau’s
Disk
11
Van Vleck
Observatory,
Astronomy Department,
Wesleyan
University,
Connecticut
06459,USA
Ian
W.
Stephens1,
HaifengYang2,
Zhi-Yun
Li2,
LeslieW.
Looney3,
Akimasa
Kataoka4,5,
WoojinKwon6,7,
12#52
Max-Planck-Institute for Extraterrestrial Physics, Giessenbachstrasse 1, 85748 Garching, Germany
13
ManuelofFerńandez-Lo University
́pez8, Charles
L. H. Hull5,9,10,
Meredith
Hughes11,
Segura- Cox3,12, Lee
Department
Astronomy,of Maryland,CollegePark, MD
20742,DominiqueUSA
14
Mundy13,
Richard andCrutcher3
and Ramprasad
Rao14
Instituteof Astronomy
Astrophysics,
Academia
Sinica, 645 N. Aohoku Place, Hilo, HI 96720, USA
ApJ 851,55 (2017)
E-mail contact: ian.stephens at cfa.harvard.edu
The mechanism for producing polarized emission from protostellar disks at (sub)millimeter wave- lengths is currently
uncertain. Classically, polarization is expected from non-spherical grains aligned with the magnetic field. Recently,
two alternatives have been suggested. One polarization mechanism is caused by self-scattering from dust grains ofsizes comparable to the wavelength while the other mechanism is due to grains aligned with their short axes along
the direction of radiation anisotropy. The latter has recently been shown as a likely mechanism for causing the dust
polarization detected in HL Tau at 3.1 mm. In this paper, we present ALMA polarization observations of HL Tau for
two more wavelengths: 870 µm and 1.3 mm. The morphology at 870 µm matches the expectation for self-scattering,
while that at 1.3 mm shows a mix between self-scattering and grains aligned with the radiation anisotropy. The
observations cast doubt on the ability of (sub)millimeter continuum polarization to probe disk magnetic fields for atleast HL Tau. By showing two distinct polarization morphologies at 870 µm and 3.1 mm and a transition between the
two at 1.3 mm, this paper pro- vides definitive evidence that the dominant (sub)millimeter polarization mechanism
transitions with wavelength. In addition, if the polarization at 870 µm is due to scattering, the lack of polarization
asymmetry along the minor axis of the inclined disk implies that the large grains responsible for the scattering havealready settled into a geometrically thin layer, and the presence of asymmetry along the major axis indicates that the
HL Tau disk is not completely axisymmetric.
Accepted by ApJ
http://adsabs.harvard.edu/pdf/2017arXiv171004670S

The spatial extent of Polycyclic Aromatic Hydrocarbons emission in the Herbig star
HD 179218
Anas S. Taha1 , Lucas Labadie1 , Eric Pantin2 , Alexis Matter3 , Carlos Alvarez4 , Pilar Esquej5 , Rebekka
Grellmann1 , Rafael Rebolo6 , Charles Telesco7 and Sebastian Wolf81

I. Physikalisches Institut, Universität zu Köln, Zülpicher Strasse 77, 50937 Cologne, Germany

The Astrophysical Journal, 851:55 (6pp), 2017 December 10
The Astrophysical Journal, 851:55 (6pp), 2017 December 10

Stephens et al.
Stephens
The Astrophysical Journal, 851:55 (6pp), 2017 December
10 et al.

Stephens et al.

The transition of polarization morphologies from one
wavelength to another is also apparent in the StokesQ and U
images, as seen in Figure 2. Speciﬁcally, at 3.1 mm the
StokesQ image shows negative regions from north to southand positive regions from east to west. At 1.3 mm, the negative
regions are much brighter than the positive regions, and at
870 μm, the positive regions almost completely disappear.
StokesU shows two positive and two negative regions at
3.1 mm. The negative blobs disappear at 1.3 mm, while thepositive regions connect, and at 870 μm, the positive regions
turn into a single, large region.
For the Stokes V images (Figure 2), 3.1 mm appears to be
purely noise. For both 1.3 mm and 870 μm, StokesV has anegative ﬂux blob offset to the top-right of the Stokes I
intensity, and this blob is signiﬁcantly brighter in 870 μm.
Although circular polarization can in principle be produced byscattering of linearly polarized light off non-spherical grains(e.g., Tazaki et al. 2017), StokesV has not been wellcharacterized for ALMA, so its detection could be spurious
and will not be discussed in detail in this paper.
The central polarization vectors for 1.3 mm and 870 μm are
55°. 4 ±0°. 96 and 55°. 3 ±0°. 41, respectively. The major axisof HL Tau’s disk is 138°. 02 ± 0°. 07 (ALMA Partnership et al.
2015), indicating the polarization vectors are ∼7° from being
perpendicular to the major axis. We would expect thepolarization to be perfectly perpendicular to the major axis.
The reason for this small 7° discrepancy is uncertain, but could
indicate asymmetry in the disk.
The right panels of Figure 2 show P for each wavelength.
For 1.3 mm and 870 μm, there is a slight asymmetry for P
along the major axis of the disk, with the southeast morepolarized than the northwest. For 1.3 mm, the absolutedifference between the two P peaks along the major axis is0.0013 ± 0.0014, which is not statistically signiﬁcant. The
Figure2. ALMAof1.3
theStokesparameters
IQUV
andthe linear
polarization
fraction P for 3.1 mm (top), 1.3 mm (middle), and 870 μm (bottom). Stokes I
Figure 2. ALMA observations of the Stokes parameters IQUV and the linear polarization
fraction
P for
3.1
mm
mm
(middle),and 870(bottom).
StokesI
asymmetryfor
870observationsμm
is(top),more
apparent,
although
theμm
southeastis given in Figure
1. 851:55 (6pp), 2017 December 10
is shown
in eachpanel for at[4,the25, 100, 500] ´ sI , wheresIa isparticular´contours
sIof, wheregiven notinwavelength
Figurecontours for a particular wavelength is shown in each panel for [4, 25, 100, 500]partThesIAstrophysicalJournal,thefordisk
does
peak 1.(i.e.,
it is still
increasingedge of the detected polarized emission). The P peak towardthat theislarge
grainsresponsible
for scattering
870 μmthat that the large grains responsible for scattering at 870 μm thethat
northwest
0.0077
± 0.0010.
Drawinga line fromatthe
are alreadyintothea thinlayerof(Yanget al.at2017).
Evidence
peaksettled
through
centerthe disk,
the sameare already settled into a thin layer (Yang et al. 2017). Evidence northwest
for dust
settling
in HL
was also±
foundin Kwonon thesoutheast
side,TauP=0.0094
0.0010.
P rises et al.
for dust settling in HL Tau was also found in Kwon et al. radius∼0.01152015)
at the edgeof theTherefore,
based
on disk.
modelingof thetheP asymmetry
millimeter dust
(2011, 2015) based on modeling of the millimeter dust to (2011,
the major
for 870is ∼0.002–0.004.
Thesecontinuum,
andaxisin Pinteet μm
al. (2016)
based on the
shape ofcontinuum, and in Pinte et al. (2016) based on the shape of along
asymmetries
arecontinuum
relatively images
moderate, but the fact that theygaps in dust
gaps in dust continuum images
appear both at 1.3 mm and 870 μm suggests they couldPT =be a real 4.feature.
we consider total
polarization
1.3 mmIf MorphologicalModelof HLTau
4. 1.3 mm Morphological Model of HLTau
Q 2 + U 2 + V 2 I , the asymmetry is similar; for 1.3 mm, it isFor
illustrative
purposes,
weμm
have
developed two simple
±
0.0014, while
for 870
it isalso∼0.001–0.004.
For illustrative purposes, we have also developed two simple 0.0012
modelset for1.3 showed
mm data:onefor the polarization
orientation
Yang
al. the(2017)
thatasymmetryalong the major
models for the 1.3 mm data: one for the polarization orientation
andis thefor thedistribution ofpolarization
the polarized
nototherexpected
forspatial
scattering-induced
in intensity.
anand the other for the spatial distribution of the polarized intensity. axis
The polarization
model
disk, soorientation
axisymmetry
mustis abesimple
broken50–50
if themix of
The polarization orientation model is a simple 50–50 mix of axisymmetrica purely azimuthal
pattern(mimicking
patternatis really due
to scattering
eventhethoughtheobserved
total
a purely azimuthal pattern (mimicking the pattern observed at polarization
3.1 mm)and aI )uni-directional
pattern along
theotherminor
axis (an(Stokesis highly axisymmetric.
On thehand,
3.1 mm) and a uni-directional pattern along the minor axis (an intensity
theidealization
polarizationof along
the minor
axis
is expected
be thethe patternobserved
at 870μm). We tocreatedFigure 1. ALMA
polarimetricat 8703.1 mm(top,We
Kataoka
et al.theidealizationof thepattern observations
observed at
μm).
createdasymmetric
the dust the
diskStokes
is opticallythickand twohas patternsamodel by ifaveraging
Q and Uof the2017),
1.3
mm(middle),
and
870
μm
(bottom),
where
the
red
vectors
show
the
model
by averaging the Stokes Q and U of the two patterns signiﬁcant angular width (for a cartoon illustration, see Figure 6>3σ polarization morphology (i.e., vectors have not been rotated). Vector
separately. Figure 3 shows the model compared with the
Figure 3. Morphological 1.3 mm model for a 50–50 mix of azimuthal andseparately.
Figure
3
shows
the
model
comparedwiththelengths are linearly proportional to P. The color scale shows the polarizedin observations.
Yang2017).
HL matchTaummdisk
is known
to bemix
optically
Figureet3.al.
Morphological
1.3
model
for a 50–50
of azimuthal
The tworemarkably
well,
which
strength-and uniformFigure
4. isDistribution
polarized
intensity
withoverlaid
polarization
patterns
shown withof
yellow
vectors,
which areon topvectorsof theintensity, which The
is masked
only show
3σ detections.
Stokes
I contours
inobservations.
two tomatch
remarkably
well,
which
strengthuniform
patterns
is
shown
with
yellow
vectors,which
areoverlaid
on
top
of
the
thick
at
870
μm
(Carrasco-Gonzálezetal.
2016),
butnosuperposed
for anmorphology
illustrative disk
whereThethe contours
polarizationareis StokesI,
producedens the notion that the 1.3 mm emission is a transition caseobserved
polarization(redmodel
vectors).
eachthepanel
are shown[3, 10,25,mm
50, 100,
200, 325, is
500,
750,
1000] × scaseI,ens
notionthatfor the
1.3
emissionatransition
observed
polarization
morphology(redvectors).
The
contours
areStokesI,byacombinationof
directemissionandscattering
byazimuthallyalignedasymmetryis detected
the minor axis. The lack ofwith the levels given in Figure 1.
between 3.1
mm and along
870 μm.
where sI is 44, 154, and 460 μJy bm−1 for 3.1 mm, 1.3 mm, and 870 μm,
grains. Color scale has been normalized to have a peak of 1.
with the levelsin theFigure1.
between
3.1 mm and 870 μm.
asymmetryfor P given
along
optically
minor onaxisa implies
respectively.
The polarized
intensitymodelthickis based
disk model

The polarized intensity model is based on a disk model
similar to ModelA of Yang et al. (2017), but with a reduced
scale height for the dust layer. The grains are assumed to be

similar to ModelA of Yang et al. (2017), but with a reduced
emission and scattering by such grains are computed under the
scale height for the dust layer. The grains are assumed to beelectrostatic approximation (see Yang et al. 2016b for details).
oblate spheroids with their shortest axes aligned radially, as

emission and scattering by such grains are computed under theelectrostatic approximation (see Yang et al. 2016b for details).
does not match all details of the data; a detailed modeling will
The resulting polarization pattern is shown in Figure 4 andprovides an independen
from the shape of the ga
By observing high-res
different wavelengths, w
polarization mechanism
lengths. Without observ
and 870 μm, the polari
have been difﬁcult to imodeling polarization in
multi-wavelength, high-r

H.F.Y. is supported byrted in part by NASA NN
1716259. A.K. acknowle
number JP15K17606. W
Research Program throug
Korea (NRF-2016R1C1B
following ALMA data: A
ADS/JAO.ALMA#2016
ESO (representing its m(Japan), together with
(Taiwan); and KASI (Rethe Republic of Chile.
operated by ESO, AUI/N
Astronomy Observatory
Foundation operated unde
Universities, Inc.

Laboratoire AIM, CEA/DRFCNRSUniversite Paris Diderot, IRFU/DAS, F-91191 Gif-sur-Yvette, France
Laboratoire Lagrange, Observatoire de la Cote dAzur, BP 4229, 06304 Nice Cedex 4, France
4
W. M. Keck Observatory, 65-1120 Mamalahoa Hwy, Kamuela, HI 96743, USA
5European Space Astronomy Centre, P.O. Box 78, 28690 Villanueva de la Canada, Madrid, Spain
The spa:al extent of Polycyclic Aroma:c Hydrocarbons emission in the Herbig star HD
1300
I. Physikalisches Institut, Universität zu Köln, Zülpicher Strasse 77, 50937 Cologne, Germany
6#53
179218Institutode Astrofisica de Canarias, C/ Via Lactea s/n, La Laguna, 38200 Tenerife , Spain
7
Anas S.ofTaha1,
Lucas Labadie1,
MaUer3,
Alvarez4,
Department
Astronomy,
UniversityEricof Pan:n2,
Florida, Alexis
Gainesville,
FLCarlos
32611,USA Pilar Esquej5, Rebekka Grellmann1,
8
RafaelRebolo6, Charles
Telesco7
and Sebas:anUniversity
Wolf8
Institutefor Theoretical
Physics
and Astrophysics,of Kiel, Leibnizstr. 15, 24118 Kiel, Germany
A&A
E-mail contact: labadie at ph1.uni-koeln.de
3

We investigate, in the mid-infrared, the spatial properties of the Polycyclic Aromatic Hydrocarbons (PAHs) emission
in the disk of HD 179218, an intermediate-mass Herbig star at ∼300 pc. We obtained mid-infrared images in the PAH1, PAH-2 and Si-6 filters centered at 8.6, 11.3, and 12.5 µm, and N-band low-resolution spectra using CanariCam on
the 10-m Gran Telescopio Canarias (GTC). We compared the point spread function (PSF) profiles measured in the
PAH filters to the profile derived in the Si-6 filter, where the thermal continuum emission dominates. We performed
radiative transfer modeling of the spectral energy distribution (SED) and produced synthetic images in the threefilters to investigate diﬀerent spatial scenarios. Our data show that the disk emission is spatially resolved in the
PAH-1 and PAH-2 filters, while unresolved in the Si-6 filter. Thanks to very good observing conditions, an average full
width at half maximum (FHWM) of 0.232′′ , 0.280′′ and 0.293′′ is measured in the three filters, respectively. Gaussian
disk fitting and quadratic subtraction of the science and calibrator PSFs suggests a lower-limit characteristic angular
diameter of the emission of ∼100 mas, or ∼30 au. The photometric and spectroscopic results are compatible withprevious findings. Our radiative transfer (RT) modeling of the continuum suggests that the resolved emission should
result from PAH molecules on the disk atmosphere being UV-excited by the central star. Simple geometrical modelsof the PAH component compared to the underlying continuum point to a PAH emission uniformly extended out tothe physical limits of the disk model. Furthermore, our RT best model of the continuum requires a negative exponentof the surface density power-law, in contrast with earlier modeling pointing to a positive exponent. We have spatiallyresolved - for the first time to our knowledge - the PAHs emission in the disk of HD 179218 and set constraints on its
spatial extent. Based on spatial and spectroscopic considerations as well as on qualitative comparison with IRS 48 and
HD 97048, we favor a scenario in which PAHs extend out to large radii across the flared disk surface and are at thesame time predominantly in an ionized charge state due to the strong UV radiation field of the 180 L⊙ central star.
Accepted by A&A
http://arxiv.org/pdf/1711.05202

ATLASGAL-selected massive clumps in the inner Galaxy: VI. Kinetic temperature andspatial density measured with formaldehyde
X. D. Tang1,2,3 , C. Henkel1,4 , F. Wyrowski1 , A. Giannetti1,5 , K. M. Menten1 , T. Csengeri1 , S. Leurini1,6 ,1Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany
Xinjiang Astronomical Observatory, Chinese Academy of Sciences, 830011 Urumqi, PR China
3
Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, 830011 Urumqi, PR China
4
Astronomy
Department, King Abdulaziz
80203,
21589 Jeddah,
Sauditemperature
Arabia300
ATLASGAL-selectedmassive University,
clumps inPO
theBoxinner
Galaxy:VI. Kine:cand
5
INAF–Istituto di Radioastronomia & Italian ALMA Regional Centre, Via P. Gobetti 101, I–40129 Bologna, Italy
spa:al density
measuredwith formaldehyde
6#54
INAF–Osservatorio
Astronomico
di Cagliari,
Via della Scienza 5, I–09047, Selargius (CA), Italy
7
Tang1,2,3,
C. Henkel1,4,
F. Wyrowski1,
A. Giannej1,5,
M. Menten1,Kent
T. Csengeri1,
S. Leurini1,6,
J. S.
School X.
ofD.
PhysicalSciences,
Universityof Kent, Ingram
Building,K.Canterbury,CT2 7NH,
UK
8
C.and
Koenig1,
Guesten1,
Y. X. Lin1,
X. W. Zheng8,
Esimbek2,3
J. J. Zhou2,3
School Urquhart1,7,of Astronomy
SpaceR.Science,
Nanjing
University,
210093J. Nanjing,
PRand
China
2

A&A

E-mail contact: xdtang at mpifr-bonn.mpg.de
Formaldehyde (H2 CO) is a reliable tracer to accurately measure the physical parameters of dense gas in star forming
regions. We aim to directly determine the kinetic temperature and spatial density with formaldehyde for the ∼100
brightest ATLASGAL-selected clumps (the TOP100 sample) at 870 µm representing various evolutionary stages ofhigh-mass star formation. Ten transitions (J = 3–2 and 4–3) of ortho- and para-H2 CO near 211, 218, 225, and
291 GHz were observed with the Atacama Pathfinder EXperiment (APEX) 12 m telescope. Using non-LTE models with
RADEX, we derive the gas kinetic temperature and spatial density using the measured para-H2CO 321 –220 /303 –202 ,422 –321 /404 –303 , and 404 –303 /303 –202 ratios. The gas kinetic temperatures derived from the para-H2 CO 321 –220 /303 –41
202 and 422 –321 /404 –303 line ratios are high, ranging from 43 to >300 K with an unweighted average of 91 ± 4 K.
Deduced Tkin values from the J = 3–2 and 4–3 transitions are similar. Spatial densities of the gas derived fromthe para-H2 CO 404 –303 /303 –202 line ratios yield 0.6–8.3 × 106 cm−3 with an unweighted average of 1.5 (±0.1) × 106
cm−3 . A comparison of kinetic temperatures derived from para-H2 CO, NH3 , and the dust emission indicates thatpara-H2 CO traces a distinctly higher temperature than the NH3 (2,2)/(1,1) transitions and the dust, tracing heatedgas more directly associated with the star formation process. The H2 CO linewidths are found to be correlated withbolometric luminosity and increase with the evolutionary stage of the clumps, which suggests that higher luminosities
tend to be associated with a more turbulent molecular medium. It seems that the spatial densities measured with
H2 CO do not vary significantly with the evolutionary stage of the clumps. However, averaged gas kinetic temperaturesderived from H2 CO increase with time through the evolution of the clumps. The high temperature of the gas tracedby H2 CO may be mainly caused by radiation from embedded young massive stars and the interaction of outflows withthe ambient medium. For Lbol /Mclump >∼ 10 L⊙ /M⊙ , we find a rough correlation between gas kinetic temperature andthis ratio, which is indicative of the evolutionary stage of the individual clumps. The strong relationship between
H2 CO line luminosities and clump masses is apparently linear during the late evolutionary stages of the clumps,
indicating that LH2 CO does reliably trace the mass of warm dense molecular gas. In our massive clumps H2 CO line
luminosities are approximately linearly correlated with bolometric luminosities over about four orders of magnitudein Lbol , which suggests that the mass of dense molecular gas traced by the H2 CO line luminosity is well correlatedwith star formation.

Accepted by A&A
https://arxiv.org/pdf/1711.10012

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Emission from magnetized accretion disks around young stars

1
C.
Tapia
and S. Lizano
300
Emissionfrom magne:zed accre:on disks around young stars
1
C. Tapia1
and S. Lizano1 y Astrofı́sica, UNAM, Apartado Postal 3-72, 58089 Morelia, Michoacán, México
Instituto
de Radioastronomı́a#55
ApJ 849, 136 (2017)
E-mail contact: c.tapia at irya.unam.mx

We calculate the emission of protoplanetary disks threaded by a poloidal magnetic field and irradiated by the central
star. The radial structure of these disks was studied by Shu and collaborators and the vertical structure was studied
by Lizano and collaborators. We consider disks around low mass protostars, T Tauri stars, and FU Ori stars with
diﬀerent mass-to-flux ratios λ . We calculate the spectral energy distribution and the antenna temperature profiles
at 1 mm and 7 mm convolved with the ALMA and VLA beams. We find that disks with weaker magnetization (highvalues of λsys ) emit more than disks with stronger magnetization (low values of λ ). This happens because theformer are denser, hotter and have larger aspect ratios, receiving more irradiation from the central star. The levelof magnetization also aﬀects the optical depth at millimeter wavelengths, being larger for disks with high λsys . In
general, disks around low mass protostars and T Tauri stars are optically thin at 7 mm while disks around FU Ori are
optically thick. A qualitative comparison of the emission of these magnetized disks, including heating by an external
envelope, with the observed millimeter antenna temperature profiles of HL Tau indicates that large cm grains are
required to increase the optical depth and reproduce the observed 7 mm emission at large radii.
Accepted by The Astrophysical Journal (849, A136, 2017)

Table 1. YSOs parameters

https://arxiv.org/pdf/1711.04803v1
YSO

Ṁd
(M yr

1

)

Md

R⇤

Lc

(M )

(R )

(L )

0.03

2

0.93

0.02

7

230

Transient events in bright debrisLMP
discs:
Collisional
avalanches
revisited2 ⇥ 10
0.20
3
7.1
6

Philippe Thebault1 and Quentin Kral2T Tauri
1

FU Ori1 ⇥ 10 8
2 ⇥ 10 4

LESIA, Paris Observatory, France
2
Institute of Astronomy, University of Cambridge, UK

E-mail contact: philippe.thebault at obspm.fr
A collisional avalanche is set oﬀ by the breakup of a large planetesimal, releasing vast amounts of small unboundgrains that enter a debris disc located further away from the star, triggering there a collisional chain reaction that
42

10

11
sys = 4sys = 12sys = 24sys = 4

102

<T >
Tc
Tz90

<T >
Tc
Tz90

<T >
Tc
Tz90

Tirr

Tirr

Tirr

z
zirr
z90

sys = 24

z
zirr
z90

z
zirr
z90

101

z [AU]T [K]103

sys = 12

102

100

10

1

10

2

101

100

101

102

100

101

$ [AU]

100

101

102

100

101

102

100

101

$ [AU]$ [AU]= 1AU= 10AU= 50AU= 1AU= 10AU= 50AU= 1AU= 10AU= 50AUF [erg cm 2 s 1 ]

103

T [K]

102

$ [AU]

102

10

8

10

9

10

10

10

11

10

12

10

13

10

14

102

100

101

$ [AU]

102$ [AU]

=0
= 60
= 80

=0
= 60
= 80

=0
= 60
= 80

101

10

2

10

1

100

101

10

2

1

10

100

z [AU]

101

10

2

10

1

z [AU]

100

0.02

0.03

0.04

✓[00 ]
0.05

0.06

0.07 0.01

0.02

0.03

0.04

102

103

0.06

0.07 0.01

0.02

0.03

0.04

103

100

0.05

0.06

sys = 12

101

102

103

[µm]

21

✓[00 ]
0.1

0.07

0.2

0.3

✓[00 ]

0.4

0.5

0.6

sys = 240.7
0.87mm
1.30mm
2.90mm
7.00mmamax = 1 mm
102

0.1

0.2

0.3

0.4

0.5

0.6

0.7
0.87mm
1.30mm
2.90mm
7.00mmamax = 1 cm

101

TB [K]

102⌧1mm

TB1mm [K]102

Figure 2. LMP disk models with di↵erent mass-to-flux ratios, sys = 4, 12, 24, that label each column. The upper panels
show the the di↵erent surfaces: the black solid lines show the surface of the disk z1 ; the green dot lines show the irradiation
surface zirr ; the blue dot-dashed lines show the disk mass surface z90 . The lower panels show the spectral energy distribution(SED) of the star plus disk system at di↵erent inclination angles ✓ between the disk rotation axis and the l.o.s: ✓ = 0 , 60 ,and 80 (red, black and blue lines, respectively). The dotted lines show the disk contribution.

103
sys = 4101[µm]✓[00 ]
0.05

100

[µm]

LMP disk models with di↵erent mass-to-flux ratios, sys = 4, 12, 24, that label each column. The upper panels
show the radial temperature profiles of the disks: the solid black lines correspond to the mass weighted temperature < T >;the red dashed lines show the mid-plane temperature Tc ; the blue dot-dashed lines show the temperature of mass surface z90 ;the green dot lines show the temperature at the irradiation surface zirr . The lower panels show vertical temperature structure
at the radii indicated in the upper right boxes. The star symbol corresponds to the location of the irradiation surface zirr andthe
20 diamond symbol corresponds to the location of the mass surface z90 .

0.01

101

z [AU]

Figure 1.

✓[00 ]

100

101

101

101

100

100
2

3

4

5

6

7

8

9

10 1

2

3

4

$ [AU]

5✓[00 ]
0.01

0.02

0.03

0.04

6

7

8

9

10 1

2

3

4

$ [AU]

5✓[00 ]
0.05

0.06

0.07 0.01

0.02

0.03

0.04

6

7

8

9

10

$ [AU]
101✓[00 ]
0.05

0.06

0.07 0.01

0.02

0.03

0.04

0.05

0.06

0.07

Optical depth

1103
sys = 4sys = 12sys = 24

100

⌧7mm

TB7mm [K]

101
102

10

101

1

10

20

30

40

50

$ [AU]
100

100
1

2

3

4

5

6

$ [AU]

7

8

9

10 1

2

3

4

5

6

$ [AU]

7

8

9

10 1

2

3

4

5

6

7

8

9

10

$ [AU]

Figure 11. Antenna temperature TB and optical depth ⌧ at 1 mm and 7 mm of FU Ori disks with di↵erent mass-to-flux ratios,
sys = 4, 12, 24, as a function of distance to the central star. The description of the panels and the lines is the same as in
Figure 9. The upper middle panel and the lower right hand panel do not show the optical depth profile because, for numericalconvenience, the integration of the optical depth ends at the value 25.

60

70

80

90

10

20

30

40

50

60

70

80

90

$ [AU]

Figure 12. Antenna temperature TB and optical depth ⌧ at 0.87, 1.3, 2.9, and 7 mm of LMP disks as a function of distanceto the central star. The radial profiles have been convolved with the ALMA beams: ✓0.87 mm = 0.03400 , ✓1.3 mm = 0.02900 ,✓2.9 mm = 0.06600 , and the VLA beam at 7 mm, ✓7 mm = 0.04300 . The color code is shown in the boxes in the upper right corners.
This color code is similar to the one used in Figure 3 of Carrasco-González et al. (2016) that plots the observed ALMA andVLA antenna temperature profiles of HL Tau. The disks have a mass-to-flux ratio sys = 24 . The disk parameters are shownin Table 4. The upper panels show TB and the lower panels show ⌧ for each wavelength. The left panels correspond to models
with a dust grain distribution with amax = 1 mm: Model I (dotted lines) without envelope heating (Te = 0 K); Model II (solidlines) with Te = 50 K; Model III (dashed lines) with Te = 100 K. The right panels correspond to models with amax = 1 cm:
Model IV (dotted lines) without envelope heating (Te = 0 K); Model V (solid lines) with Te = 50 K; Model VI (dashed lines)
with Te = 100 K.

Transient events in bright debris discs: Collisional avalanches revisited
Philippe Thebault1 and Quentin Kral2

300
Transient events in bright debris discs: Collisional avalanches revisited
LESIA, Paris Observatory, France
Philippe
Thebault1University
and Quen:n
Kral2
2#56
Instituteof Astronomy,of Cambridge,
UK

1

A&A
E-mail contact: philippe.thebault at obspm.fr

A collisional avalanche is set oﬀ by the breakup of a large planetesimal, releasing vast amounts of small unboundgrains that enter a debris disc located further away from the star, triggering there a collisional chain reaction that
could potentially create detectable transient structures. We investigate this mechanism, using for the first time a fully
self-consistent code coupling dynamical and collisional evolutions. We also quantify for the first time the photometric42
evolution of the system and investigate if avalanches could explain the short-term luminosity variations recently
observed in some extremely bright debris discs. We use the state-of-the-art LIDT-DD code (Kral et al., 2013, 2015).
We consider an avalanche-favouring A6V star, and two set-ups: a ”cold disc” case, with a dust release at 10 au andan outer disc extending from 50 to 120 au, and a ”warm disc” case with the release at 1 au and a 5-12 au outer disc.
We explore, in addition, two key parameters, which are the density (parameterized by its optical depth τ ) of the mainouter disc and the amount of dust released by the initial breakup. We find that avalanches could leave detectablestructures on resolved images, for both ”cold” and ”warm” disc cases, in discs with τ of a few 10−3 , provided that20
22
large dust masses (>∼10 –5×10 g) are initially released. The integrated photometric excess due to an avalanche isrelatively limited, less than 10% for these released dust masses, peaking in the λ∼10–20µm domain and becoming
insignificant beyond ∼40–50µm. Contrary to earlier studies, we do not obtain stronger avalanches when increasing τ
to higher values. Likewise, we do not observe a significant luminosity deficit, as compared to the pre-avalanche level,
after the passage of the avalanche. These two results concur to make avalanches an unlikely explanation for the sharp
luminosity drops observed in some extremely bright debris discs. The ideal configuration for observing an avalanchewould be a two-belt structure, with an inner belt (at ∼1 or ∼10au for the ”warm” and ”cold” disc cases, respectively)of fractional luminosity f ≥10−4 where breakups of massive planetesimals occur, and a more massive outer belt, with
τ of a few 10−3 , into which the avalanche chain reaction develops and propagates.

Accepted by Astronomy and Astrophysics
https://arxiv.org/pdf/1711.06482

Non-linear Development of Secular Gravitational Instability in Protoplanetary Disks
Ryosuke T. Tominaga1 , Shu-ichiro Inutsuka1 , Sanemichi Z. Takahashi2,3,4
1

Department of Physics, Nagoya University, Nagoya, Aichi 464-8692, Japan
Astronomical Institute, Tohoku University, Sendai 980-8578, Japan3
Department of Applied Physics, Kogakuin University, Hachioji, Tokyo, 192-0015, Japan2

2

Thebault & Kral: Collisional avalanches revisited

2013; Wyatt & Jackson, 2016) 2 . The observed dust could betransient, possibly produced in the aftermath of the breakup ofa massive body (Meng et al., 2014; Wyatt & Jackson, 2016), or
be sustained by comet or asteroid scattering in closely packed
multiple planet systems (Bonsor et al., 2013).
The discovery of this new class of extreme discs has spawned
Thebault & Kral: Collisional avalanches revisited5
a renewed interest in the collisional avalanche
scenario, which
could indeed be greatly facilitated by these system’s high level of
dustiness. This interest has been compounded by the fact that a
handful of these discs exhibit brightness variations on very shorttimescales, of the order of a year, for which avalanches havebeen suggested as a possible explanation (Meng et al., 2015).
It is worth noting that avalanches have been mostly invoked as
being able to explain the rapid luminosity decrease in some ofthese systems, in particular the ∼50% decay over ∼370 days following the initial brightening of ID8 (Meng et al., 2014, 2015),
and the abrupt luminosity drop (by a factor 30) observed forFig. 1. Schematic description of the avalanche process. First, a the TYC 8241 2652 1 system over the course of less than twolarge body is shattered in the inner region of the planetary sys- years (Melis et al., 2012). The argument is that the passage oftem, releasing very small unbound grains that are blown out by the avalanche would leave the main disc depleted in " 100µm
radiation pressure. Then these small grains penetrate an outer particles, as they have been "sandblasted" by the high-velocity
disc located further away from the star. There they will collide, at impacts with the outbound avalanche grains, making the postvery high velocities, with "native" target grains, producing a new avalanche system (after the transient luminosity excess due togeneration of small unbound grains, which will also be blown unbound grains) less luminous than it was prior to the avalanche.
out and collide with other target grains in the disc, creating new Interestingly, this luminosity drop consequence of an avalanche
fragments, etc.. If the main disc is dense enough, this chain re- had not been mentioned by GAT07, who insisted only on the
Thebault& exceeding
Kral: Collisional
avalanchesincrease
revisitedthat it might trigger. It is true that the tem-7
luminosity
action can produce an excess of small grains
largelyporal evolution of the amplification factor F showed a progresthe amount initially produced by the breakup.
grainsof themain
disc and
thatincrease,
we are here
theparameter
Wien domain
sive decay
after
a strong
initialbut inthis
onlyofthe Planck
which a temperature
increase
results
measuresthe function,
avalancheforgeometricalcross section
with respectinexponential
increaseof the grains,
Bλ (λ, T )but
parameter.
to an
that
of the initiallyreleased
not that of the toa less massive object to reach observability than for an isolated tal avalanche+disc
In addition, it issystem.
importantnote ofthat,
reference
As a tomatterfact,forthethis
photometric−3
planetesimal breakup.
τevolution
case,
thesystem
breakupitself createsinaGAT07,
luminosity
total
wasevent
not investigatedthe⊥ ∼2×10of theThis scenario has first been quantitatively investigated by excess
whichessentiallyexceeds, atonallspatial
wavelengths,
thatof localized
the subsequentfocus being
signatures
andlumiGrigorieva et al. (2007) (hereafter GAT07), using a then unique avalanche(see inset in Fig. 7). In other words, from a photometnosity contrasts.
numerical code coupling the dynamical and collisional evolu- ric point of view, the avalanche is not able to amplify the signal
tion of a disc. This study has shown that collisional avalanches produced in the immediate aftermath of the breakup. This seems
1.3.
the avalanche
mechanism
can indeed strongly amplify the geometrical cross section of the to
berevisiting
in disagreementwith GAT07,
who obtained an "amplificainitially breakup-released dust, and produce spiral shaped pat- tion" factor F∼200 for a similar set-up. This diﬀerence is mostly
In
the
light
of
this
new
context,
itappears
important to reinvesterns in the disc that can last for a few orbital periods. GAT07 due to the fact that F measures geometricalcross
sections
tigate the collisional avalanche mechanism andreassess
its andrelquantified an avalanche’s amplitude by a cross section "ampli- notactualluminosities.
As
such
it
does
not
account
for the fact
evancewithrespecttoobserved
systems.
Suchareinvestigation
fication factor" F and found that this parameter increases expo- that the initial breakup occurs closer to the star, where both thewas in any case necessary and long overdue, firstly because
thenentially with the outer disc’s density (quantified by its radial scattered(becauseof thedilutionin 1/r2(pho) andpioneeringlightstudy
of GAT07
leftstellar
manyflux
issues
unexplored
optical depth τr ). However, for the nominal case they consid- thermal(becausedustcloser
to
the
star
is
warmer
and
thus
more
tometric evolution, wavelength dependence, etc.), which maybe
ered, i.e., a breakup releasing 1020 g of dust and a β Pictoris-like luminous,
least in
24-70µm
range) emission
dust
opened theat door
to thesome
misconceptions
about ofthea given
avalanche
disc, GAT07 found that collisional avalanches induce luminos- grainarehigherthanfurther
out
in
the
disc.
When
considering
phenomenon, but also because new debris disc modelling codes
ity increases much too small to be observed. They extrapolated thegeometricalcross
section,retrieve
havedisc-integrated
been developedtotal
in thepast decade,whichallowwea moreacFig.
Cold disccase:5system-integrated(using
GRaTerthat 8.
a systemat least
times denser thanSEDβ Pic,
a discthealreadyat GAT07’s
results,
whichis thatthe avalanche
does advanced
indeed havecurate
study
of
this
complexmechanism.
The
most
of
packageforasystematadistanceof30pc),
at
the
peak
of
the
the upper end of bright debris discs, should be required in order an amplification eﬀect when compared to the initially released−3
this new generation of codes is probably the LIDT-DD model
avalanche
luminosity,
the referenceτ⊥debris∼2×10disc.
disc
for an avalanche
to beforobservable
in aGAT07
resolved
At dustpopulation(fullin Fig. which
7). is to date the only that al−2
by Kralet lineal. (2013),
andfor a only
denser∼10
The
lines givethe lu- developed
the time,oneτ⊥suchverydisc.
brightdiscdashed
was known,
BD+20307
Fig. 3. Normalized
synthetic
images,in scattered
light,
showing
evolution of an avalanche in a cold debris disc (extending fromWhat
theavalanchedoes,
however,
is toextendthetheduration
lows
for
a
coupledstudyofcollisions
and−3 dynamics taking intominosity
excessduetotheavalanche(computed
bysubtract50 luminosity
to 120 au) of average
opticaldepthτ⊥ ∼2×10, forthean initial
breakup, exat t=0, releasing 5×1022 g of dust at 10 au from the star.
Song et al. (2005), an unresolved system, thus not allowing a of theexcessperiod.
Indeed,while
luminosity
account
thecrucialsize-dependenteﬀectofradiationpressure.
ing
the
SED
of
a
reference
avalanche-freecase).
The
cutoﬀ
ofstraightforward application of the avalanche hypothesis.
cesscreatedby
thebreakupitselfquickly
drops
to
a
much
lower
LIDT-DD allows to relax some of the simplifications done in
the dashed-blue line beyond 50µm reflects the fact that, for this level,
be themeasuredat some locationin scattered light excess
images, we
findfor a much
Fig. 3 shows synthetic head-on images of the system insubsequentavalanche-induced
GAT07,
the
fact
that
"avalanche"(i.e.,haslasts
theinitially
redense disc case, the avalanche-system is slightly less luminous, longer
that,time,tomainlymeet
thiscriteria,thereleased
amountroughlyof dustto be
scattered
of
the
order
of
∼100
years,
a
third
of
the
dy-light, for an unfolding avalanche in this massive
22
leasedall new
debrisis they
will tocollisionally
1.2.
avalanches
in "extreme”
debrisdiscs?
verydebris
large, mandg, which
equivalent
the mass of a generate)
minit ∼5×1022g breakup case. It is qualitatively similar to what
init ∼5×10
attheselong wavelengths,
than
an avalanche-freecounterpart. namicaltimescaleat theininnermost
main disc.
This
du- by GAT07, with a tail of high-β grains escaping on
was obtained
km (i.e.,
object.grainsand ∼150"field"the outer partdisc ofnottheaﬀectedby collisions
This is because, beyond 50µm, the flux is mostly dominated by rationis, however,shorterthan whatwasfoundby GAT07(see trajectory from the breakup location, and then pena spiral-like
However, in the past decade, several other very luminous with avalanchegrains)particles
weretreatedseparately,but
also
large grains, which are slightly eroded by the passageof thediscs (with IR fractional luminosities f ! 10−2 ) have been Fig.3 of that paper), a diﬀerence that is once again due to the fact
avalanche (see discussion in Sec.4.4.1)
discovered (e.g. Meng et al., 2015; Genda et al., 2015; Melis, that
2 GAT07 were considering the system’s cross section and not
Note that a debris disc does not necessarily have to be "extremely"
photometry.
Indeed, the cross section amplification happen2016). For these "extreme" debris discs, the amount of dust re- itsluminous in order to exceed the maximum possible luminosity Lmax exing
in from
the outerparts ofcascade.
the disc,
stronglycould
contributes
toquired to explain their current luminosities, especially at wave- pected
for exama collisional
L ≥which
Lmax situations

Ryosuke T. Tominaga1 , Shu-ichiro Inutsuka1 , Sanemichi Z. Takahashi2,3,4
1

Department of Physics, Nagoya University, Nagoya, Aichi 464-8692, Japan
Astronomical Institute, Tohoku University, Sendai 980-8578, Japan
3300
Non-linear
Developmentof Secular
Gravita:onal
Instability
Protoplanetary Disks
Department
of Applied
Physics, Kogakuin
University,
Hachioji, Tokyo,
192-0015,inJapan
4
RyosukeT. Tominaga1,
Shu-ichiro
Inutsuka1,
Sanemichi
Takahashi2,3,4
National
Astronomical
Observatoryof Japan,
Osawa,
Mitaka,Z. Tokyo
181-8588, Japan#57
PASJ
E-mail contact: tominaga.ryosuke at a.mbox.nagoya-u.ac.jp
2

We perform non-linear simulation of secular gravitational instability (GI) in protoplanetary disks that has been
proposed as a mechanism of the planetesimal formation and the multiple ring formation. Since the timescale of thegrowth of the secular GI is much longer than the Keplerian rotation period, we develop a new numerical scheme fora long term calculation utilizing the concept of symplectic integrator. With our new scheme, we first investigate thenon-linear development of the secular GI in a disk without a pressure gradient in the initial state. We find that thesurface density of dust increases by more than a factor of one hundred while that of gas does not increase even by afactor of two, which results in the formation of dust-dominated rings. A line mass of the dust ring tends to be very
close to the critical line mass of a self-gravitating isothermal filament. Our results indicate that the non-linear growthof the secular GI provides a powerful mechanism to concentrate the dust. We also find that the dust ring formed via
the non-linear growth of the secular GI migrates inward with a low velocity, which is driven by the self-gravity of thering. We give a semi-analytical expression for the inward migration speed of the dusty ring.
Accepted by PASJ

102

101

Σ/Σ0, Σd/Σ0

http://arxiv.org/pdf/1711.05948

0

10Implications for planetary systemformation from interstellar object 1I/2017 U1 (‘Oumuamua)
10-1

1David E. Trilling , Tyler Robinson1 , Alissa Roegge1 , Colin Orion Chandler1 , Nathan Smith1 , Mark1,2
10-2
Loeﬄer1 , Chad Trujillo1 , Samuel Navarro-Mezaand Lori M. Glaspie1
1
2

Department of Physics and Astronomy Northern
University
Box 6010 Flagstaﬀ, AZ 86011
90
100 Arizona
110
120
130 PO 140
Distance from the Central Star [au]Instituto de Astronomı́a Universidad Nacional Autónoma de México Ensendada B.C, 22860 México
Fig. 8. Snapshot of the distribution of the surface densities of dust and gas. The thick blue line shows the surface density of dust, the thin red line shows

E-mail contact: david.trilling
atarenau.edu
that of gas. Both of themnormalized by the unperturbed surface density of gas. The solid, dashed and dotted-dashed lines show the surface density
distribution at t = 26148 yr, 36479 yr, 54689 yr respectively. The most unstable mode of the secular GI at r = 110 and 130 au grows. While the gas surface
density does not almost increase, the dust surface density becomes a hundred times larger than the initial surface density.

43

0.2

Line

1
0.8
0.6
2.4

10000

20000

2.2

30000
Time [yr]

40000

50000

Line Mass [ 10-3MJ au-1 ]Fig. 10. Time evolution of the line mass of the2dust ring. The solid line shows the line mass of the dust ring, and the dashed line shows the critical line massof filament in which a sound velocity is equal1.8
to cd we used in this work. We find that the line mass of the dust ring formed by the secular GI is comparable tothe critical line mass.

1.6
1.4

1.2
between the self-gravity of thering and that of filament, whose mass that hold equilibrium by the1pressure gradient force is determinedby the critical line mass. Gap structure also forms around the
0.8
ring. The dust-to-gas ratio in0.6these gaps decreases by a factor of about ten. We find the dust-to-gas
10000

20000

30000
Time [yr]

40000

50000

Fig. 10. Time evolution of the line mass of the dust ring. The solid line shows the line mass of the dust ring, and the dashed line shows the critical line massof filament in which a sound velocity is equal to cd we used in this work. We find that the line mass of the dust ring formed by the secular GI is comparable to140
Distance from the Central Star [au]the critical line mass.

between the self-gravity of the
135 ring and that of filament, whose mass that hold equilibrium by thepressure gradient force is determined by the critical line mass. Gap structure also forms around thering. The dust-to-gas ratio in130these gaps decreases by a factor of about ten. We find the dust-to-gas
125

from
Specific Angular MomentumDistance[ jkep(1au)] the Central Star [au]

140
120

0

10000

20000
30000
Time [yr]

40000

50000

Fig. 11. Time evolution of positions of dust cells.
135 We reduce the number of cells in the plot. Because we solve the evolution with the Lagrangian fluid equation,
the more cells there are, the larger the surface density is. The radius where the cells concentrate represents the radius of the ring. We can see that the ringmoves inward after it forms.

130
11.55

ratio reaches the minimum value
at the vicinity of the dense ring.
11.5
125

We also find that the dust ring migrates inward. Because this disk initially has the flat density
11.45

distribution, the radial drift 120
of dust grains does not occur in the unperturbed state. Nevertheless, we
0
11.4

10000

20000

30000

40000

50000

Time [yr]find that non-linear growth of the secular GI results
in a radial drift of the dust ring. Figure 11 shows
11.35
Fig. 11. Time evolution of positions of dust cells.
We reduce the number of cells in the plot. Because we solve the evolution with the Lagrangian fluid equation,16

the more cells there are, the larger the surface density is. The radius where the cells concentrate represents the radius of the ring. We can see that the ring
11.3
moves inward after it forms.

11.25

32000

36000

40000

44000

ratio reaches the minimum value at the vicinity ofTime
the[yr]dense ring.

48000

Fig. 12. Time evolution of specific angular momentums of dust and gas. The thin gray line represents that of dust, the thick red line represent that of gas. The

We also find that the dust ring migrates inward. Because this disk initially has the flat density

vertical axis shows the specific angular momentum normalized by the Keplerian specific angular momentum jkep (r) at r =1au. We reduce the number of

distribution,
the radial drift of dust grains does not occur in the unperturbed state. Nevertheless, wecells we plot here with a constant interval. The specific angular momentum of dust decreases with the dust ring migrating, while that of gas increases when
find
that
of the ratio
secularGIresults
resultsin ratea radialof thedustring. ofFigure
11 ofshows
the dust
ring non-linear
passes through. growth
The high dust-to-gas
in the ringin largerof changedriftof specificangularmomentumgas than thatdust.
16
the time evolution of positions of dust cells. From t ≃ 30000yr the secular GI grows non-linearly, and300 Implica:ons for planetary system forma:on from interstellar object 1I/2017 U1#58 (‘Oumua- mua)David E. Trilling1, Tyler Robinson1, Alissa Roegge1, Colin Orion Chandler1, Nathan Smith1, Mark Loeﬄer1, Chad
Trujillo1, Samuel Navarro-Meza1,2 and Lori M. Glaspie1
ApJL 850, L38 (2017)

The recently discovered minor body 1I/2017 U1 (‘Oumuamua) is the first known object in our Solar System thatis not bound by the Sun’s gravity. Its hyperbolic orbit (eccentricity greater than unity) strongly suggests that it
originated outside our Solar System; its red color is consistent with substantial space weathering experienced overa long interstellar journey. We carry out an simple calculation of the probability of detecting such an object. We
find that the observed detection rate of 1I-like objects can be satisfied if the average mass of ejected material fromnearby stars during the process of planetary formation is ∼20 Earth masses, similar to the expected value for our
Solar System. The current detection rate of such interstellar interlopers is estimated to be 0.2/year, and the expected
number of detections over the past few years is almost exactly one. When the Large Synoptic Survey Telescope begins
its wide, fast, deep all-sky survey the detection rate will increase to 1/year. Those expected detections will provide
further constraints on nearby planetary system formation through a better estimate of the number and properties ofinterstellar objects.
Accepted by ApJL
http://arxiv.org/pdf/1711.01344

The VLA Nascent Disk And Multiplicity Survey of Perseus Protostars (VANDAM).
III. Extended Radio Emission from Protostars in PerseusL! ukasz Tychoniec1,2 , John J. Tobin1,3 , Agata Karska4 , Claire Chandler5 , Michael M. Dunham6,7 , ZhiYun Li8 , Leslie W. Looney9 , Dominique Segura-Cox9 , Robert J. Harris9 , Carl Melis10 and Sarah I.
Sadavoy11
1

Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300RA Leiden, The Netherlands
Astronomical Observatory, Faculty of Physics, Adam Mickiewicz University, Sloneczna 36, PL-60268 Poznan, Poland
3
Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, 440 W. Brooks Street, Norman,
Oklahoma 73019, USA
4
Centre for Astronomy, Nicolaus Copernicus University in Torun, Faculty of Physics, Astronomy and Informatics,
Grudziadzka 5, PL-87100 Torun, Poland
2

5

Grudziadzka 5, PL-87100 Torun, Poland
National Radio Astronomy Observatory, P.O. Box O, 1003 Lopezville Road, Socorro, NM 87801-0387, USA
6
Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA, USA
7
Departmentof Physics,
StateDisk
Universityof New York Fredonia,
Fredonia,
New
York 14063,(VANDAM).
USA300
The VLA
Nascent
And Mul:plicity
Survey ofPerseusProtostars
III.
8
Department of Astronomy, University of Virginia, Charlottesville, VA 22903, USA
ExtendedRadio Emission
Protostars
in61801,
Perseus
9#59
Departmentof Astronomy,
University fromof Illinois,
Urbana, IL
USA
10
Tychoniec1,2,
John
J. Tobin1,3,
Karska4,
Chandler5,
MichaelM.92093,
Dunham6,7,CenterLukasz
for Astrophysics
and
Space
Sciences,Agata
Universityof Claire
California,
San Diego,
CA
USA Zhi- Yun Li8,
11
Leslie W. Looney9,
Segura-Cox9,
J. Harris9,
Carl Melis10
and Sarah I. Sadavoy11
Max-Planck-Institut
fur DominiqueAstronomie,
Konigstuhl Robert
17, D-69117
Heidelberg,
Germany
ApJ 852, 18 (2018)
E-mail contact: tychoniec at strw.leidenuniv.nl
5

Centimeter continuum emission from protostars oﬀers insight into the innermost part of the outflows, as shock-ionizedgas produces free-free emission. We observed a complete population of Class 0 and I protostars in the Perseus molecular
cloud at 4.1 cm and 6.4 cm with resolution and sensitivity superior to previous surveys. From a total of 71 detections,
8 sources exhibit resolved emission at 4.1 cm and/or 6.4 cm. In this paper we focus on this sub-sample, analyzing
their spectral indices along the jet, and their alignment with respect to the large-scale molecular outflow. Spectral
indices for fluxes integrated toward the position of the protostar are consistent with free-free thermal emission. The
value of the spectral index along a radio jet decreases with distance from the protostar. For six sources, emission iswell aligned with the outflow central axis, showing that we observe the ionized base of the jet. This is not the case fortwo sources, where we note misalignment of the emission with respect to the large-scale outflow. This might indicatethat the emission does not originate in the radio jet, but rather in an ionized outflow cavity wall or disk surface. For
five of the sources, the spectral indices along the jet decrease well below the thermal free-free limit of -0.1 with > 2σsignificance. This is indicative of synchrotron emission, meaning that high energy electrons are being produced in theoutflows close to the 16
disk. This result can have far-reaching implications for the chemical composition of the embedded
disks.
Table 1. Properties of powering sources of resolved jets
Accepted by The AstrophysicalJournal
SourceRegionOther namesa

Classbhttps://arxiv.org/pdf/1711.03035.pdf

Lbol b

Tbol b

LK

PAc

Per-emb-8

IC 348

PER22, IC 348a, IRAS 03415+3152, YSO 480

2.6±0.543.0±6.0

135 (5)Per-emb-18

NGC 1333

NGC 1333 IRAS7, YSO 240

2.8±1.759.0±12.0

159 (2)Per-emb-20

L1455

L1455-IRS40

1.4±0.2

65.0±3.0

115 (2)Per-emb-30

B1

PER19, B1 SMM11, YSO 40

0/I

1.1±0.093.0±6.0

109 (2)Per-emb-33

L1448

PER02, L1448 N(A), L1448 IRS3B, YSO 20

8.3±0.857.0±3.0

105 (3)Per-emb-36

NGC 1333

PER06, NGC 1333 IRAS2B, YSO 16I

6.9±1.085.0±12.0

24 (1)

44

L1448 IRS3A

L1448

-

I

9.2±1.3

47.0±2.0

38 (4)

SVS 13C

NGC 1333

VLA20

1.5±0.2

21.0±1.08 (1)a Names: YSOXX - Jørgensen et al. (2007), PERXX - Karska et al. (2014), VLAXX - Rodrı́guez et al. (1997)b References: Enoch et al. (2009), Sadavoy et al. (2014)c References: (1) Plunkett et al. (2013), (2) Davis et al. (2008), (3) Kwon et al. (2006), (4) Lee et al. (2016), (5) Tobin etal. in prep.

the molecular outflow could indicate that the free-free emission is tracing a portion of an ionized outflow cavity. Thismay be similar to what was observed toward a more luminous protostar in Serpens (Hull et al. 2016). The 9 mm
observations
SVS 13C is a Class 0 object, as classified by Sadavoy et al. (2014), located in the NGC1333 region in Perseus.
This of L1448 IRS3A are also extended in the same direction and the spectral index was found to be relatively(Tobin et al. 2016b). Thus, the free-free emission is likely contributing also at shorter wavelengths.
source was also detected at millimeter wavelengths by Looney et al. (2000). It produces the most prominentflat
extended5
radio jet in our sample. The position angle for its centimeter emission is 9 ± 1 . It extends ⇠400 (900 AU)8 in boththe north and south direction. The extended centimeter emission toward this source had been previously reported byRodrı́guez et al. (1997) (their source VLA2) and later by Reipurth et al. (2002) and Carrasco-González et al.
(2008).
valueof 0.25 ± 0.16 along the extension, and the spectral index remains consistent with free-free emission despite
However, our observations are the highest fidelity maps to date taken toward this source. We find irregular,
clumpybeing
marginally steeper than optically thin free-free emission. This source did not have previously published CO
structure along the jet, with the spectral index significantly decreasing outwards. At 200 (460 AU) distancethewhich would indicate outflow direction and the Spitzer data are not well-enough resolved to be
or Hin2 observations,6
northern
direction, we can observe a source of emission detached from the main component, suggesting that definitive.
there is a However, recent ALMA observations resolved a disk around Per-emb-8 (Tobin et al. in prep.) and observed
clump of denser matter in the area. Plunkett et al. (2013) tentatively identified an outflow with 8 position theangle
and in 12 CO. The observed disk has a position angle of 45 and the CO outflow is orthogonal to this. The
outflowa strongly
negative
spectralindexposition( 1.10 angle± 0.47)
theofjetdirection.
The directionof extended
extended radio
radioextended
emission, bothwhich
are consistent
with anjet.
The 4.1 cm and 6.4 cm emission has a position angle of about ⇠45 di↵erent from the CO outflow direction.
Leeet al. (2016)
found
a similarof 0along
Figure6. Images of Per-emb-8 with contours as in Figure 1 ( 6.4 cm = 4.90 µJy and 4.1 cm = 3.70 µJy). Synthesized beamfeatureoriginating
fromthethisobservation
source (Davisal. 2008),
havinga position
is consistent
with the toH2beoutflow
also appearedin the(HH790)
plane of thesky in bothfrometPlunkett
et al.
(2013)and Thus,
Lee angle
etlikeal. L1448 IRS3A, the extended radio emission might
trace the edge of an ionized outflow cavity or the ionizedis shown
in the left bottom corner (Sp. Index and 6.4 cm: 0.00 36⇥0.00 34, 4.1 cm: 0.00 26⇥0.00 25). The star marks the position of theof ⇠109This
which
is consistentof 117motions± 4 . Stormet al.
also byfinds
evidence
for an HCO+
outflow(2016).
is consistent
withwiththe ourlargefitproperof ⇠100km (2014)
s 1 foundRaga
et al. (2013).
surfaceofthe disk.
protostar basedon Ka-band observations (Tobin et al. 2016b). Red and blue arrows are orthogonal to the major axis of the+
toward Per-emb-30 with a similar position angle. Interestingly, the H2, radio, and HCO outflow appear mono-polar;
disk based on ALMA observations (Tobin et al. in prep.).
upcoming CO observations from the MASSES survey (e.g., Lee et al. 2014) will be more definitive. The mono-polar
appearance
may reflect
the spatialSpectral
distribution
dense
around [-2,the-1.6,
protostar;south-easternof 1.6,theoutflow
Figure1. Imagesof Per-emb-36.
index ofmap
withgascontours:-1.2, -0.8,-0.4, 0.0, 0.4,portion
0.8, 1.2,
2.0].
6.4
Figure5. Images of L1448 IRS3A with contours as in Figure 1 ( 6.4 cm = 5.26 µJy and 4.1 cm = 4.20 µJy). Synthesized
is interacting
mediumthe12,north-western
portionis not.
Per-emb-18
cm
and 4.1 cmwith
mapsambientwith contours[-3,while
3, 6, 9,15, 20] x where
beam is shown in the left bottom corner (Sp. Index and3.7.
6.4 cm:0.00 45 x 0.00 36, 4.1 cm: 0.00 29 x 0.00 23). The star marks the position
6.4 cm = 4.83 µJy and 4.1 cm = 3.90 µJy. Synthesized

3.2. SVS 13C

beam is shown in the left bottom corner (Sp. Index and 6.4 cm: 0.00 41⇥0.00 35, 4.1 cm: 0.00 26⇥0.00 22). The stars mark theof position
the protostar based on Ka-band observations (Tobin et al. 2016b). Red and blue arrows indicate outflow direction from Lee
Per-emb-18
is a Class 0 system in the NGC1333 IRAS7 region consisting of two protostars separated by only 20 AUof the protostars based on Ka-band observations (Tobin et al. 2016b) and the red and blue arrows indicate outflowetdirection
al. (2015) with a position angle of 38 .
(Tobinet al. 2016b). Thus, C-band observations were unable to resolve the system. The radio emission is asymmetricfrom Plunkett et al. (2013) with a position angle 24 . 6.4 cm and 4.1 cm maps are not corrected for a primary beam response
and the spectral index decreases steeply from 0.37 ± 0.14 at the protostellar position to 1.13 ± 0.42 in the southern(spectral index derived from PB corrected map).

with the value of the H2 outflow (159 ) observed by Davis
outflow. Measured position angle of 164 ± 16 is consistent
3.6. Per-emb-8
et Per-emb-8
al. (2008).isThe
steep
negative
spectral
index
in
this
sourcepartlyto thelackof emission
detected(⇠2200AU)
from
Per-emb-55
which atisa Class 0 protostar in the IC 348 region, and it isis alsolocated9.00 6 due
3.2. SVS 13C
cm multiple
where extended
6.4 cmetemission
is present.
a4.1close
itself (Tobin
al. 2016b).
Per-emb-8 shows 4.1 cm and 6.4 cm emission extended north and slightly
SVS 13C is a Class 0 object, as classified by Sadavoy et al. (2014), located in the NGC1333 region in Perseus.
east, This
the position angle is 13 ± 2 . The spectral index smoothly decreases further away from the protostar, with asource was also detected at millimeter wavelengths by Looney et al. (2000). It produces the most prominent extended
00
radio jet2.in Naturally-weighted
our sample. The position
angle
forwithits centimeteremission± 1cm. =It4.83
extends⇠4 4.1(900
in µJy).
both
Figuremap of SVS
13C
contours as in
Figure is
1 (9 6.4µJy and= Figure
3.90
6. Images of Per-emb-8 with contours as in Figure 1 ( 6.4 cm = 4.90 µJy and 4.1 cm = 3.70 µJy). Synthesized beam
cm AU)
00
00
the north andsouth
direction.
extended
centimeteremissionsourcepreviously
reportedbySynthesized
beam
is shown
in the The
left bottom
corner(Sp. Index
and 6.4toward
cm:: 0.00this41⇥0.
35, 4.1had
cm:been0.00 26⇥0.
22). Thestar
marks
is shown
in the left bottom corner (Sp. Index and 6.4 cm: 0.00 36⇥0.00 34, 4.1 cm: 0.00 26⇥0.00 25). The star marks the position of the9thepositionetof al.
the(1997)
protostarbased
on Ka-band
et al.et2016b)
and theand blue arrows indicate
outflowbased on Ka-band observations (Tobin et al. 2016b). Red and blue arrows are orthogonal to the major axis of the
protostar
Rodrı́guez(theirsourceVLA2) observations
and later by(Tobin
Reipurthal. (2002)
andredCarrasco-González
et al.
(2008).
direction
et al.are(2013)
Lee etal. (2016).
disk
based on ALMA
observations(Tobinet al.
in as
prep.).
Per-emb-20
is a Class
0 protostar
known
also
L1455 IRS4. An extension of the 4.1 cm and 6.4 cm is observed with
However,from
our Plunkett
observations
the andhighestfidelitymaps to date taken toward this source. We find irregular,
clumpya position
structure along the jet, with the spectral index significantly decreasing outwards. At 200 (460 AU) distancein the angle of 137 ± 8 . This position angle is reasonably consistent with the 115 position angle found by Davisal. (2008)
northern direction, we can observe a source of emission
detached from the main component, suggesting thatet
thereis a from H2 observations. The extension is small, but is present at both wavelengths, and other sources in
3.3. Per-emb-30
3.7. Per-emb-18
mapclump of denser matter in the area. Plunkett et al. (2013) tentatively identified an outflow with 8 position theangle
and00do not show a similar extension. The spectral index is flat in the central part
5 (0.00±0.18) and is decreasing
Per-emb-30
is afoundClass a0 similarsource position
located in
the ofBarnard1 region.
radio emission
6.4 cm extending⇠1Per-emb-18
is a Classalong
0 systemin the NGC1333
IRAS7 region
consistingalongof twothe
protostars
separatedby onlybut
20 AU
the northwest
the extension.
The spectral
index measured
extensionis -1.38±0.51,
theof whichWe
arefindconsistent
with anatextended
radiotoward
jet.
The
Lee
et al. (2016)
angle
0 , both
(230
AU)
from
the source
morecompact
at 4.1
The morecompact
emission
at 4.1 andcm extension
results
in is(Tobinet2016b).
Thus, C-band
marginallyresolvedobservations
at 6.4 cm. were unable to resolve the system. The radio emission is asymmetricoutflow
also
appeared
to beandin theplaneof the emission
sky in boththe cm.
observation
from
Plunkett
et al. (2013)
Lee et al.
7al.only
and the spectral index decreases steeply from 0.37 ± 0.14 at the protostellar position to 1.13 ± 0.42 in the southern(2016). This is consistent with the large proper motions of ⇠100 km s 1 found by Raga et al. (2013).
outflow. Measured position angle of 164 ± 16 is consistent with the value of the H2 outflow (159 ) observed by Davis
Figure7. Images of Per-emb-18 with contours as in Figure 1 ( 6.4 cm = 6.77 µJy and 4.1 cm = 5.54 µJy). Synthesized beam
Figure 3. Images of Per-emb-30 with contours as in Figure 1 ( 6.4 cm = 4.87 µJy and 4.1 cm = 4.01 µJy). Synthesized
etal.beam(2008).
The steep negative spectral index in this source is also partly due to the lack of emission detected at
is shownofin the left bottom corner (Sp. Index and 6.4 cm: 0.00 56⇥0.00 38, 4.1 cm: 0.00 37⇥0.00 25). The stars mark the position ofis shown in the left bottom corner (Sp. Index and 6.4 cm: 0.00 34 x 0.00 37, 4.1 cm: 0.00 24 x 0.00 22). The star marks the position
4.1
cm
where
extended 6.4 cm emission is present.
the protostars based on Ka-band observations (Tobin et al. 2016b). Red and blue arrows indicate outflow direction from Davisthe protostar based on Ka-band observations (Tobin et al. 2016b) and the red and blue arrows indicate the outflow direction
et al. (2008) with a position angle of 159 .
from Davis et al. (2008) with a position angle 109 .

3.8. Per-emb-20

8. Images of Per-emb-20 with contours as in Figure 1 ( 6.4 cm = 5.30 µJy and 4.1 cm = 4.05 µJy). Synthesized beam
Figure 4. Naturally-weighted images of Per-emb-33 with contours as in Figure 1. ( 6.4 cm = 4.87 µJy and 4.1 cm =Figure
4.01 µJy).
is mark
shownthein the left bottom corner (Sp. Index and 6.4 cm: 0.00 38⇥0.00 35, 4.1 cm: 0.00 24⇥0.00 22). The star marks the position of theSynthesized beam is shown in the left bottom corner (Sp. Index and 6.4 cm: 0.00 34 x 0.00 37, 4.1 cm: 0.00 24 x 0.00 22). Stars
1. 3.4.
Imagesmap
contours:[-2, -1.6, -1.2, -0.8, -0.4, 0.0, 0.4, 0.8, 1.2, 1.6, 2.0]. 6.4
Figure 2. Naturally-weighted map ofFigure
SVS 13C with
contours
asofin Per-emb-36.
Figure 1 ( 6.4 cmSpectral= 4.83 µJyindex
and 4.1=with
3.90 µJy).
Per-emb-33
based on Ka-band observations (Tobin et al. 2016b). Red and blue arrows indicate outflow direction from Davis et al.
positions of the protostars based on Ka-band observations(Tobinet al. 2016b).00 The
sourceon the east
is 0033-C,cm
whileprotostar
A (south)00
00
Synthesized beam is shown in the left bottom
corner(Sp.
Index
and with
6.4 cm::contours
0. 41⇥0. 35,[-3,4.1 3,
cm:6,0. 9,
26⇥0.
22).
star
marks
cm
and
4.1
cm
maps
15,The20]
where
6.4 angle
cm =of4.83(2008)
witha position
115 . µJy and 4.1 cm = 3.90 µJy. SynthesizedandB (north) forma tightsystembinary. ofThe
red andblue
arrowsindicate
outflow
directionfrom(Looney
Kwon
et et
al.12,(2006)
witha xposition
Per-emb-33
isatripleClass
0protostars,
also
known
as
L1448
IRS3B
al.
2000)
and
L1448
the position of the protostar based on Ka-band observations (Tobin et al. 2016b) and the red and blue arrows indicate outflow 00
00
angle
of
105
.
beam
is
shown
in
the
left
bottomcorner
(Sp.
Index
and
6.4
cm:
0.
41⇥0.
35,
4.1
cm:
0.00 26⇥0.00 22). The stars mark the position
N(B) (Curiel
et al. 1990).
directionfrom Plunkett
et al. Per-emb-33-A(2013) and Lee iset dominating
al. (2016). the emission in the C-band and Per-emb-33-C is marginally
Figure7.
ImagesofPer-emb-18
with
as in Figure
1 ( 6.4 cm= 6.77 µJy
and 4.1 cm = 5.54 µJy). Synthesized beamof the protostars
based
et al.The
2016b) and the red andcontours
blue arrows
indicate
outflow
direction
detected, in contrast to the 9 mm observations,where source
C hasontheKa-band
highest fluxobservations
density (Tobin(Tobin
et al. 2016b).
4. cm:DISCUSSION
is shown in the left bottom corner (Sp. Index and 6.4
0.00 56⇥0.00 38, 4.1 cm: 0.00 37⇥0.00 25). The stars mark the position of
from
Plunkettetal.
(2013)
withaposition
angle
24
.
6.4
cm
and
4.1
cm
maps
are
not
corrected
foraprimary
beam
response
resolution was not sufficient to resolve Per-emb-33-B
A. The radio emission is also extended the
along
the
3.5. from
L1448source
IRS3A
protostars
based
on Ka-bandwavelengths
observationshave(Tobinet al. 2016b).
Red and blue
arrowsindicate
outflow(e.g.,
direction
Davis
Resolvedjets
at centimeter
typically
been associatedwithmassive
protostars
HH from
80-81;
see
3.3.making
Per-emb-30
jet outflow direction measured by Lee(spectral
et al. (2015),
it difficult
determinewhetheroris with a position angle of 159 .
index
derivedfromtoPB
corrected
map).
et al. (2008)
00 not Per-emb-33-C

Marti
et 00al. 1995) and with a few remarkable low-mass protostars that have exceptionally powerful jets (e.g., L1527,
LooneyL1448 IRS3A is a Class I protostar and a wide companion of Per-emb-33, separated by 7. 3 (⇠1700 AU;
actually
detected
if thisis just locatedthe extended
emission.
Along We
the findradioradio
jet/outflow
direction
to thewest, there
Per-emb-30
is a orClass
0 source
in thejet
Barnard1 region.
emissionat 6.4 cm
extending
⇠1
HHcm
1-2
VLA 1, SVS 13C; see Rodrı́guez et al. 1997, 2000; Reipurth et al. 2004). However, analysis of the frequency
et al. 2000; Tobin et al. 2016b). It has extended centimeter emission in00 the western direction at both 4.1
andAU), which
alsoatappears
tobe dueis a clumpof emission
separatedmore
from the source
position
distance
2. 4 (540(230
AU)Thefromthe source
emission
at by
4.1acm.
The±moreemission
4.1 emission
cm ofresultsin radio jets toward an unbiased sample of low-mass protostars has only recently become possible with theextended
6.4 cm.
spectralindexandnortheast compactof the protostar
positionis 0.61
0.24.compact
This indicates
that
the
could
3.8. Per-emb-20
to the radio jet. The spectral index along the outflow is consistent with optically thin free-free emission, varying
from
3.2.
SVS
13C
adventof theupgraded Jansky VLA. While some of the previous studies looked at a reasonably large number of sources
be produced by synchrotron emission, but the fact that emission arises on the edge of the source casts a doubton
this
0.14 ± 0.14 in the A source position to 0.30 ± 0.74 in the north-western clump and 0.53 ± 0.50 in the position(e.g.,
et al. 1998), the observations were often conducted at resolutions of ⇠100 or greater and the observations
detection. The portion more extended to the west has a spectral index of 0.13 ± 0.12 and is consistent with
free-free
SVSjet13C
0 object,
asCO
classified
by Sadavoy
etAngladaal. (2014),
located in the NGC1333 region in Perseus. Thisof source C. Position angle of the extended
107 is± 5a Class
is consistent
with theoutflow position
angle obtained

4

Astronomy Data Center, National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588,
Japan
5
Université de Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France
6300
High spectral
resolu:onobserva:onsof Alfonso
HNC3 andin the L1544
Observatorio
Astronómico
Nacional (OAN,
IGN), Calle
XII, HCCNC
3, 28014 Madrid,
Spain prestellar core7
C. Vastel1,Institut,
K. Kawaguchi2,
D. Quzúenard3,
Ohishi4,Str.
B. Leﬂoch5,
R. Bachiller6,
H.S.P. Mu ̈ller7
I. Physikalisches
Universität
Köln,M.
Zülpicher
77, 50937
Köln, Germany#60
MNRAS 474, L76 (2018)
E-mail contact: cvastel at irap.omp.eu
HCCNC and HNC3 are less commonly found isomers of cyanoacetylene, HC3 N, a molecule that is widely found indiverse astronomical sources. We want to know if HNC3 is present in sources other than the dark cloud TMC-1 andhow its abundance is relative to that of related molecules. We used the ASAI unbiased spectral survey at IRAM 30m
towards the prototypical prestellar core L1544 to search for HNC3 and HCCNC which are by-product of the HC3 NH+
recombination, previously detected in this source. We performed a combined analysis of published HNC3 microwave
rest frequencies with thus far unpublished millimeter data because of issues with available rest frequency predictions.
We determined new spectroscopic parameters for HNC3 , produced new predictions and detected it towards L1544. We
used a gas-grain chemical modelling to predict the abundances of N-species and compare with the observations. The
modelled abundances are consistent with the observations, considering a late stage of the evolution of the prestellar
core. However the calculated abundance of HNC3 was found 5–10 times higher than the observed one. The HC3 N,
HNC3 and HCCNC versus HC3 NH+ ratios are compared in the TMC-1 dark cloud and the L1544 prestellar core.

5

Accepted by MNRAS Letters• IRAM 30m, Astrochemical Studies At IRAM
http://arxiv.org/pdf/1711.10255(ASAI) unbiased spectral survey• 実験室でのマイクロ波分光とミリ波天体観測の合わせ技で分子定数を決めてラインを同定
Probing
the massive star forming environment - a multiwavelength investigation of the
し、prestellarIRDC
core L1544でHNC
filamentary
G333.73+0.37
3を初めて検出• Nを含む分子について、gas-grain chemical
V. S. Veena1 , S. Vig1 , B. Mookerjea2 , Á. Sánchez-Monge3 , A. Tej1 and C. H. Ishwara-Chandra4
modellingと観測を比較
1
Indian• InstituteScience and Technology, Figure
Thiruvananthapuram,695 547, India
HC3NH+,ofHCSpace
3. Column densities for HC N, HNC and HCCNC as a function of the age of the phase 2 for diﬀerent initial gas densities of the
3N, HNC3, HCCNCはコアの外
2ambient cloud phase: 10 cm (full line), 3 × 10 cm (dashed line), and 2 × 10 cm (dash-dotted line). The area of confidence of theTata Instituteof Fundamental Research, Mumbai, observed
600 column
005,densities
India
for these species is also shown. Grey areas shows the timespan area of confidence of each model based on the
縁部（他の分子でnon-thermal
observed HCNH . These areas are labelled solutions 1, 2 and 3 (Quénard et al. 2017). Left panel: Model EA1 of initial atomic abundances
3
Physikalisches
Institut, Universität zu Köln, ZülpicherStr.
77, Right
50937
Köln,(”low-metalabundances”).
panel: Model
EA2 ofGermany
initial atomic abundances (”high-metal abundances”). All the informations on thedesorp:onが見られる）から出ている
modelling can be found in Quénard et al. (2017).
4
National
Centre for Radio Astrophysics(NCRA-TIFR), Pune, 411 007, India• L1544がprestellar
core後期の段階にあACKNOWLEDGEMENTS
Kawaguchi K., Ohishi M., Ishikawa S.-I., Kaifu N., 1992a, Theることを考慮すると、モデルの予想は観E-mail contact:veenavs.13 at iist.ac.in
Astrophysical Journal, 386, L51
Based on observations carried out with the IRAM 30 m Telescope.
Kawaguchi K., et al., 1992b, ApJ, 396, L49
IRAM is supported by INSU/CNRS (France), MPG (Germany),
測された強度とコンシステント（例外：Kawaguchi K., Kasai Y.,
Ishikawa
S.-I., OhishiM., Kaifu N.,
We present a multiwavelength study of the filamentary
dark
cloud
G333.73+0.37.
The
regioncontains
and infrared
IGN (Spain). K.K.
and M.O.
are very (IRDC)
grateful to S. YaAmano T., 1994, The Astrophysical Journal Letters, 420, L95
mamoto
and
S.
Saito
for
the
opportunity
to
carry
out
laboratoryE., dust
Rawlingsemission
J.,
Casellifrom
P.,
2014,
3）
two distinctHNC
mid-infrared
sources S1 and S2 connectedby investigations
dark lanes
dust.KetoColdtheat thegas
Nagoyaand
University.
spectroscopicon HNC of
3

2

3

−3

3

−3

4

−3

+

3

H.S.P.M. thanks Y. Endo for comments on the FTMW data of
HNC3 . M.O. is supported by the JSPS Kakenhi grant number
JP15H03646. C.V. would like to thank J.-C. Loison for a fruitful discussion on the reaction rates and branching fraction of the
DR of the HC3 NH+ ion. D.Q. acknowledges the financial support
received from the STFC through an Ernest Rutherford Grant

Monthly Notices of the Royal Astronomical Society,440,

2616 in the region. The physical
IRDC is detected at seven wavelength bands and we have identified 10 high density clumps
Krüger M., Dreizler H., Preugschat D., Lentz D., 1991, Angew.
Chem. Int. Ed.,30, 1644
properties of the clumps such as temperature: 14.3 − 22.3 K and mass: 87 − 1530 MKrüger
determinedby fitting a⊙ areM., Stahl W., Dreizler H., 1993, J. Mol. Spectrosc., 158,
298
modified blackbody to the spectral energy distribution of each clump between 160 µmOsamuraandY.,1.2
mm. The total mass
Fukuzawa K., Terzieva R., Herbst E., 1999,