星形成ゼミ 2015/ 8/7 Newsletter #271 No.21‒25
担当；清兼和紘•21. Pre- and Post-burst Radio Observations of the Class 0 Protostar HOPS 383 in
Orion
Galván-Madrid, Roberto; Rodríguez, Luis F.; Liu, Hauyu Baobab; Costigan, Gráinne; Palau,
Aina; Zapata, Luis A.; Loinard, Laurent
http://adsabs.harvard.edu/abs/2015ApJ...806L..32G•

22. A 1.3 cm line survey toward IRC +10216
Gong, Y.; Henkel, C.; Spezzano, S.; Thorwirth, S.; Menten, K. M.; Wyrowski, F.; Mao, R. Q.;
Klein, B.
http://adsabs.harvard.edu/abs/2015A%26A...574A..56G•

23. HNC in Protoplanetary Disks
Graninger, Dawn; Öberg, Karin I.; Qi, Chunhua; Kastner, Joel
http://adsabs.harvard.edu/abs/2015ApJ...807L..15G•

24. The Giant Molecular Cloud Environments of Infrared Dark Clouds
Hernandez, Audra K.; Tan, Jonathan C.
http://adsabs.harvard.edu/abs/2015arXiv150200183H•

25. Massive envelopes and ﬁlaments in the NGC 3603 star forming region
Hummel, C. A.; Stanke, T.; Galvan-Madrid, R.; Koribalski, B. S.
http://adsabs.harvard.edu/abs/2015arXiv150608597H
1

PRE- AND POST-BURST RADIO OBSERVATIONS OF THE CLASS 0 PROTOSTAR HOPS 383 IN ORION
Roberto Galván-Madrid1, Luis F. Rodríguez1,2, Hauyu Baobab Liu3, Gráinne Costigan4,
1
1
1
Palau
, Luis A. Zapataof, and• 21.1 Pre- and Post-burstAina
Radio
Observations
theLaurent
Class 0Loinard
Protostar HOPS 383 in Orion
Centro de Radioastronomía y Astrofísica, UNAM, Apdo. Postal 3-72 (Xangari), 58089 Morelia, Michoacán, México; r.galvan@crya.unam.mx
2
Astronomy
Department,
Faculty of Science,King
Abdulaziz
University,
P.O. Box
80203, Jeddah
21589, Saudi
ArabiaGalván-Madrid,Roberto;Rodríguez,
Luis F.;
Liu,
Hauyu
Baobab;
Costigan,
Gráinne;Palau,
Aina; Zapata,3
Academia Sinica, Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan
4Luis A.; Loinard, Laurent
Leiden Observatory, University of Leiden, PB 9513, 2300 RA Leiden, The Netherlands
Received 2015 April 3; accepted 2015 May 30; published 2015 June 18

http://adsabs.harvard.edu/abs/2015ApJ...806L..32G (ApJ Letter)

ABSTRACT
There is increasing evidence that episodic accretion is a common phenomenon in Young Stellar Objects (YSOs).
Recently, the source HOPS 383 in Orion was reported to have a q35 mid-infrared—and bolometric—luminosity
increase between 2004 and 2008, constituting the ﬁrst clear example of a class 0 YSO (a protostar) with a large
accretion burst. The usual assumption that in YSOs accretion and ejection follow each other in time needs to betested. Radio jets at centimeter wavelengths are often the only way of tracing the jets from embedded protostars.
We searched the Very Large Array archive for the available observations of the radio counterpart of HOPS 383.
The data show that the radio ﬂux of HOPS 383 varies only mildly from 1998 January to 2014 December, staying at
the level of ∼200–300 μJy in the X band (∼9 GHz), with a typical uncertainty of 10–20 μJy in each measurement.
We interpret the absence of a radio burst as suggesting that accretion and ejection enhancements do not follow eachother in time, at least not within timescales shorter than a few years. Time monitoring of more objects and speciﬁc
predictions from simulations are needed to clarify the details of the connection betwen accretion and jets/winds
in YSOs.
Key words: ISM: jets and outﬂows – radio continuum: stars – stars: formation – stars: protostars

・Class 0 Protostar HOPS 383 (Orion)で中間赤外とBolometric Luminosityで2004-2008にsubmillimeter-to-bolometric luminosity and in the character1. INTRODUCTION
x３５の増光が確認された。
istic temperature of the spectral energy distribution, Safron
The ・episodic
physical mechanisms
ofgas
accretion
andejectionareaccretionとoutﬂow/jetの相関を調べたい。
et al. (2015) conclude that the object is a class 0 YSO, making
fundamental aspects in the paradigm of low-mass star
HOPS 383 the youngest protostar known with an accretion・そこで、VLAのアーカイブデータを調べ、センチ波で同期した増光が確認できるかを調べたが、formation.
In its most commonly assumed form, this paradigm
related burst. The outburst is most notably seen at a wavelength
requires・結果としてはセンチ波での増光は確認できなかった。
the quasi-stationary evolution of the forming starsof 24 N m , in which the brightness increased by a factor q35(e.g., Shu et al. 1994). This assumption is seriously challenged
from 2004 to 2008. From their extensive IR to submillimeterby the ・議論ではこの状況を説明する解釈を３つほど提示している。
observational evidence of episodic accretion. First, afollow up, Safron et al. (2015) ﬁnd no evidence for signiﬁcanthandful・YSOsでのaccretionとjets/windsの詳細な関係を明らかにする為には、シュミレーションによるof individual bursting objects such as FU Ori and EX
fading of HOPS 383 up through 2012.
Lup provide evidence that large accretion bursts in YoungRadio observations are a key tool to study YSOs: the thermal
Stellar 予測と、より多くの天体での時間的なモニター観測が必要である。
Objects (YSOs) do actually happen, with durations
free–free emitting radio jets are often the only way to probe thefrom a few months to ∼100 years (e.g., Hartmann &bipolar ejection of material close to the protostar in the
Kenyon 1996; Herbig 2008, Audard et al. 2014). Second,
optically obscured class 0 objects (e.g., Rodríguez 1997;
surveys show that YSOs are systematically underluminous with
Anglada et al. 1998). Radio observations are also useful torespect to the expectation of steady-accretion models (e.g.,
Evans et al. 2009; Kryukova et al. 2012). Episodes of 2 probe non-thermal (gyro)synchrotron radiation from active
YSO magnetospheres (e.g., Güdel 2002; Forbrich et al. 2007).
signiﬁcantly increased accretion would alleviate this discre-The Astrophysical Journal Letters, 806:L32 (5pp), 2015 June 20

Table 1
Data SummaryEpoch15.01.1998
14.03.2008+01.06.2008
14.03.2008+01.06.2008
13.08.2011
13.08.2011
27.10.2014
03.11.2014
06.12.2014
13.12.2014
15.12.2014
All 2014

Frequency(GHz)

Array, HPBW, PA(arcsec × arcsec; deg)

Phase Center(h:m:s; deg:arcmin:arcsec)

HOPS383 Offset(arcsec)rms Noisea(μJy beam−1)

Flux Densityb
(μJy)

8.46
4.86
8.46
4.87
7.42
10.0
10.0
10.0
10.0
10.0
9.81

D; 9.2 q 7.5, −17.8
DnC+C; 6.0 q 4.2 , 49.8
DnC+C; 4.8 q 3.1, 57.4
A; 0.64 q 0.42 , −4.6
A; 0.44 q 0.27, −3.6
C; 2.18 q 1.66 , −1.9
C; 2.04 q 1.81, −7.2
C; 2.29 q 1.70 , −28.3
C; 2.35 q 1.61, −11.1
C; 2.28 q 1.65, −29.3
C, 2.32 q 1.62 , −14.1

05:35:24.395; −05:01:07.27
05:35:24.400; −05:01:43.25
05:35:24.413; −05:00:07.25
05:35:23.420; −05:01:30.52
05:35:23.420; −05:01:30.52
05:35:23.480; −05:01:32.30
05:35:23.480; −05:01:32.30
05:35:23.480; −05:01:32.30
05:35:23.480; −05:01:32.30
05:35:23.480; −05:01:32.30
05:35:23.480; −05:01:32.30

111.1
138.2
82.2
137.8
137.8
138.5
138.5
138.5
138.5
138.5
138.5

42
56
47
10
11
13
15
11
13
10
6

230 ± 50
224
188
270 ± 20310 ± 20
222 ± 12267 ± 16293 ± 13289 ± 15
246 ± 10260 ± 6

Galván-Madrid et al.

Notes.
arms noise corrected for primary-beam attenuation around the position of HOPS 383: B  5h 35m 29.s 81, E  4n59a51.´ 1 (Safron et al. 2015).
Figure 1. High angular resolution JVLA image of the radio counterpart ofb
Flux density of HOPS 383 corrected for primary-beam attenuation. The errors are only statistical, resulting
fromobserved
the noise
theAugustmaps 13,andandthecenteredqualityatof7.42
the GHz
ﬁt. The
HOPS 383,
on in
2011
(4.0 cm).
absolute uncertainty of the VLA ﬂux scale is a few percent (Perley & Butler 2013). Upper limits in HPBW
the 2008are.  9 ◦. 0 . The target clearly has two components.
, PA 0.epoch´ 26 q 0.
´ 214T

Figure 2. Radio light curve of HOPS 383. The red dots and inverted triangles
are X-band (8–10 GHz) measurements and upper limits, respectively. The blue
symbols are in the C-band (4–5 GHz). Error bars are o1T . See Table 1 fordetails of the observations. The vertical gray stripe marks the period of time
Galván-Madrid
etthatal.Safron et al. (2015) estimated for the occurrence of the infrared burst.
The peak insensities of the SE and NW components are 75and55 μJy beam−1,
respectively. The rms noise is 11 μJy beam−1.

The Astrophysical Journal Letters, 806:L32 (5pp), 2015 June 20

free emission from radio jets and accretion signatures as seen in
the near- to mid-infrared. Bursting protostars offer a unique
opportunity to perform this test under different conditions in agiven object.
In this Letter, we report on pre- to post-burst centimeter
wavelength observations of the bursting class 0 YSOHOPS 383.
2. DATA SETS AND BASIC RESULTS

consider a detection. Our reanalysis of these data shows a faintRegardless of the detailed jet emission mechanism, in thesource ・中間赤外でのx35の増光(2004-2008)にセンチ波での増光with ﬂux density of 0.23 ± 0.05 mJy that coincides
models it is assumed that an increase of accretion is followed
as the ﬂux density from the concatenated data. There appear to
spatially withbeHOPS
383.
by an ejection enhancement. Therefore, if the mid-IR burst
marginal ﬂux variations within 2014. The maximal ﬂux

は確認できない（1998-2014の間の観測で増光は確認できな
truly traces a burst of accretion (Safron et al. 2015), then adensity is 293 ± 13 μJy on 2014 December 12, which is 71 ±2.2.minimal
2008 value of 222 ± 12 μJy on 2014
signiﬁcant increase in the radio ﬂux should follow.
18 μJy above theWe ﬁnd that in HOPS 383 the infrared burst does not have a
October 27. Then, without considering the few percent
かった）。[センチ波の観測は埋もれたClass0のような若い天
These observations
on 2008March
14 andJune 2013),
1
counterpart in the radio. Figure 2 shows the radio light curves.
uncertaintywere
in themade
absolute
ﬂux scale(Perley& Butlerin the C and the
DnC
conﬁgurations
atwithin
4.86 andGHz asatpart
The long-term (decadal) radio ﬂux densities show mild
largestﬂux difference
20148.46
is signiﬁcant
4T . We体でejection(jets/winds)を確認するのに(唯一の?)有力な手
variations, staying at a level of ∼200–300 μJy. The earliest
of project AT359.
data fromthe bytwo
epochswere
also testedThefor intra-day
variability
splittingeach ofthe ﬁve
1998 ﬂux density is the same within 1T with the latest 2014
epochs into
two timebut
chunks.
In four of theconcatenatedobserving
for increased
sensitivy,
a counterpart
to ﬁve段である]
measurements(Table 1). The average ﬂux density of theepochs,theﬂux
measured
in
each
halve
of
the
data
is
within1σ
HOPS 383 was not detected with a 4σ upper limit of 0.22 andTo analyze the time behavior of HOPS 383 at centimeter
detections in the X band (8–10 GHz) is SX  265 μJy, with aof theandﬂux 8.46
reportedin Table
1. In the 2014 November 3 epoch,
0.19 mJy at 4.86
GHz,
respectively.
wavelengths, we looked for the available observations in thedispersionof 28 μJy, or about the noise in single epochs. Taken
the ﬂuxes in the halved data are within 2σ of the epoch average,
Very Large Array5 archive. In none of these observations isatface
value,
it may even seem that during the IR-burst theand this is the noisiest of the 2014 observations. We discard
2.3.
2011
HOPS 383 at the phase center; however, it falls inside the ﬁeld
radio
ﬂux
has
a minimum (see Figure 2). We constrain thesigniﬁcant intra-day variability. These results are consistent
今回の現象を解釈として３つ可能性を列挙している。
signiﬁcanceof
a possible radio ﬂux decrease in 2008 bywith
therelatively
longvariation
timescale
expectedfortheof view, allowing the determination of its parameters after
These observations
were
made
on
2011
Augustto13
withtheclassconsideringthenon-detection in this epoch to have a ﬂux
radio
emission
of
class
0
YSOs,
compared
class
II
and
correction for the primary beam response. In the following
１.Jansky
一般にaccretionとejectionの変動は相関しない可能性(観The hypothetical ﬂux
upgraded
Very
Large
Array
(JVLA)
in oftheA
between 0 and the 4T upper limit.
III
(Liu
et
al.
2014).
Finally,thehighsensitivity
these
JVLA
subsections we describe the observation epochs. Table 1
conﬁgurationobservations
under project
11A-220.
tuningis centereddecrease between the 1998 detection and the 2008 nonallowed
the data toThebe split
into smallerfrequency
summarizes the relevant properties of the data.
detection has a signiﬁcance between 3.3T and 0.6T . The
chunks.
Wewithmadetotal
imagesfrom the concatenateddata in fourat 4.81 測事実的に)and 7.42
GHz,
bandwidthsof approximately
In all the epochs the position of the radio source matches
hypothetical ﬂux increase between 2008 and 2011 (the epoch
windows,
each
1
GHz
wide.
The
intra-band
spectralindex
of
1 GHz. This is the only epoch with subarcsecond angularwithin 0″. 3–1″. 7 with the infrared position reported bywith the maximum ﬂux) has a signiﬁcance between 6.0T and
HOPS
383
is
consistent
with0,
indicating
optically
thin
free–
２.The
Large
BurstでProtosteller
accretionとejectionの変動はresolution.
source associatedwith HOPS 383 is for the ﬁrst
2.3T . These estimates do not include the absolute ﬂux
free emission.
Safron et al. (2015): B (J2000)  5h 35m 29.s 81, E (J2000) time resolved
into two
components
NW)a
Figure 2. Radiolight curveof HOPS(SE383. andThe reddots withand inverted
triangles
uncertainty of a few percent. In the following subsections,
4n59a51.´ 1. This is a fraction of the synthesized beam in all the
相関しない。”jetquenching”
andjoined
by
a
ridge
of
faint
emission.
In
separation
of
x
0.
´
45
are
X-band(8–10GHz)
measurements
and
upper
limits,respectively.
The
blue
we discuss several possible explanations for our result. These
Figure1.
Highangularresolution
JVLA
image
of
the
radio
counterpartof
3. DISCUSSION
epochs in which the radio source appears unresolved (seeforsymbolsare(uin, theC-band (4–5GHz).
barswearemeasureo1T . See Table 1possibilities
close
to Errornatural,maps done
withv ) weighting
do not necessarily exclude each other.
HOPS
383,
observedon
2011
August13,
andcentered
at
7.42
GHz
(4.0
cm).
Table 1), and about the synthesized
beam size for the 2011３.
トレースしているものの解釈；radio
continuum
emission
detailsoftheobservations.
The
verticalgray
stripemarks
theperiod
of
time
3.1.
No
Radio
Counterpart
totheInfraredBurst◦
has two components. ﬂux densities that are consistent with rest of the detections at
HPBW  0.´ 26 q 0.´ 21, PA  9 . 0 . The target clearly
6
Since
the
next
epoch
with
subarcsecond
angularresolution.
−1
that
Safronetal.
(2015)
estimatedfortheoccurrenceoftheinfrared
burst.
± 0.02 at 4.81
GHz
andwith0.31± 0.02mJy at
The peak insensities of the SE and NW components are 75and 55 μJy beam , other epochs: 0.27
MHD-launched
jets thatload
them
a fractionof the3.2. Protostellar Accretion and Ejection Variations−1 and the Spitzer-MIPS
away,
closest radio
はjetを見ていて、mid-Infraredはaccretionを見ているという.
respectively.
Thedetection
rms noise isis x11120
μJyaabeam
7.36 GHz (seeTablematerial
1). Theindex
betweenthese
accreted
are aspectral
basic featureof starformationmodelsAre Not Correlated in Generalangular resolution is _6aa, we conclude that the radio detection(e.g.,
Pudritz
et al. 2007;
Shang et al.
2007). Furthermore,
frequencies is
0.3,
indicating
free–free
emissionwith lowthere
Regardlessof thedetailedthisjet
emissionmechanism,
here reported is the counterpart of the bursting protostar
Accretion rates and jet/wind ejection rates have been shownis thecaseGHzeven doneforthewithyoungest in the
optical のは正しそう。けれども、異なる(複数の)component(binary
depth.is increasing
Figure 1 evidenceshows athatmap
at
7.36
to
systems:
the
class
0
protostars(e.g.,Lietal.
2014).
In
these
modelsit is assumed
that
an increaseof accretionis followedbe correlated within samples of the relatively evolved class II
bydensity
Safron etal. (2015).
asreported
the ﬂux
fromthe concatenated data. There appear to weighting
betweennatural anduniform
to emission
maximizeangular
YSOs and FU emissionはOrionis objects (e.g., Hartigan et al. 1995;
deeplyembedded
objects,
free–free
from
radio
jets
is
system)を見ているかもしれない。ﬁg1のradiobyanTheejection enhancement.
Therefore,themid-IRburst
be marginal ﬂux variations within 2014. The maximal ﬂux resolution.
naturethe radio
source if
ismaterial
clearly
Calvet et al. 1998). However, the idea that in a given object
one doubleof the best
ways ofto trace
the outﬂowing
in the2.1.2014
1998December 12, which is 71 ± seen, probably
truly traces
a ahundredburstofaccretionetal.
2015),
adensity is 293 ± 13 μJy on
frombipolarjet or 1997;
a(Safronbinary.
Thismapaccretionand ejection follow each other in time has only started
inner
few
AUradio(Rodríguez
Angladaet al.
1998). thenbipolarjetともbinaryとも解釈できる(両者の距離は~190au程
signiﬁcant
increase
in
the
radio
ﬂux
shouldfollow.
withthehighestangularresolution
already
suffersfromtobetested.
Contrary to the simplest expectations, it may be
Theoriginoftheradio
emissionisoften
consideredtobe
18 μJyabove
theminimal
value
of
222
±
12
μJy
on
2014
These observations were made on 1998 January 13 as part of
the
case
that
they do not follow each other. Ellerbroek et al.
shock-ionizedgas(Anglada
1996),
butmodelsthat
include)ﬂux
ﬁlteringofextendedemission.
signiﬁcant(
50%that in HOPS 383 the infrared burst does not have a
We_ﬁndOctober
27.AR387.
Then, The
without
considering
the few andpercent
the project
analysishas been presented度である)。
(2014)
could
not establish a relation between outﬂow and
X-ray ionization
withinthe
X-windscenario
also
reproducethecounterpart
in the3
radio. Figure 2 shows the radio light curves.
uncertainty
the absolute
ﬂux scale(Perley& Butlerat2013),
discussed in in
Reipurthet al. (1999).
The data
were obtainedaaccretion variability in the Herbig Ae/Be star HD 163296, theobserved properties of radio jets (Shang et al. 2004).
2.4. 2014 radio ﬂux densities show mild
The long-term (decadal)
frequency of 8.46 GHz in the D conﬁguration. Reipurth et al.

!

1

•

Purple Mountain Observatory & Key Laboratory for Radio Astronomy, Chinese Academy of Sciences, 2 West Beijing Road,
210008 Nanjing, PR China
22. A e-mail:1.3 cmline survey toward IRC +10216
ygong@pmo.ac.cn
2
Max-Planck Institute für Radioastronomy, Auf dem Hügel 69, 53121 Bonn, Germany
Gong,3 University
Y.; Henkel,
C.;Academy
Spezzano,
S.;No.
Thorwirth,S.; 100049
Menten,of Chineseof Sciences,
19A Yuquan Road,
Beijing,K.
PR M.;
ChinaWyrowski, F.; Mao, R. Q.;
4
Astronomy Department, King Abdulaziz University, PO Box 80203, 21589 Jeddah, Saudi Arabia
Klein,5 B.
I. Physikalisches Institut, Universität zu Köln, 50937 Köln, Germany
6
University of Applied Sciences Bonn-Rhein-Sieg, Grantham-Allee 20, 53757 Sankt Augustin, Germany

http://adsabs.harvard.edu/abs/2015A%26A...574A..56G (A&A)
Received 15 August 2014 / Accepted 19 November 2014

ABSTRACT

Context. IRC +10216 is the prototypical carbon star exhibiting an extended molecular circumstellar envelope. Its spectral properties
are therefore the template for an entire class of objects.
Aims. The main goal is to systematically study the λ ∼ 1.3 cm spectral line characteristics of IRC +10216.
Methods. We carried out a spectral line survey with the Eﬀelsberg-100 m telescope toward IRC +10216. It covers the frequency range
between 17.8 GHz and 26.3 GHz (K-band).
Results. In the circumstellar shell of IRC +10216, we find 78 spectral lines, among which 12 remain unidentified. The identified lines
are assigned to 18 diﬀerent molecules and radicals. A total of 23 lines from species known to exist in this envelope are detected forthe first time outside the solar system and there are additional 20 lines first detected in IRC +10216. The potential orgin of “U” lines isalso discussed. Assuming local thermodynamic equilibrium (LTE), we then determine rotational temperatures and column densitiesof 17 detected molecules. Molecular abundances relative to H2 are also estimated. A non-LTE analysis of NH3 shows that the bulkof its emission arises from the inner envelope with a kinetic temperature of 70 ± 20 K. Evidence for NH3 emitting gas with higher
kinetic temperature is also obtained, and potential abundance diﬀerences between various 13 C-bearing isotopologues of HC5 N are
evaluated. Overall, the isotopic 12 C/13 C ratio is estimated to be 49 ± 9. Finally, a comparison of detected molecules in the λ ∼ 1.3 cm
range with the dark cloud TMC-1 indicates that silicate-bearing molecules are more predominant in IRC +10216.
Key words. astrochemistry – ISM: molecules – stars: individual: IRC +10216 – radio lines: stars

・IRC+10216[AGB carbon star]の星周シェル(circumstellar shell)での1.3cm帯 (17.8 - 26.3 GHz)のラインサーベイ1. Introduction
manner. In addition, complex molecules have small rotational
constants so that their lowest-energy transitions arise at cm・結果；78輝線(12未同定輝線；18{異なる分子あるいはラジカル}；23輝線はこのエンベロープで存在がMolecular line surveys
are a powerful tool for analyzing both wavelenths. In this way, emission or absorption features fromphysical and chemical parameters of astronomical objects. In long carbon-chain molecules can be observed (e.g. HC N,
3
知られている分子種で太陽系の外で初検出、さらに20輝線はIRC+10216で初検出。
fact, they are the only means of obtaining a complete, unbiased HC N, HC N, HC N). Since systematic and highly sensitive
5
7
9
view
of the molecular inventory of a given source. However, de- surveys at these (low) frequencies that cover diﬀerent sources
・未同定輝線に関しては可能性を議論。
spite the long history of powerful large radio telescopes, few un- have rarely been carried out, one can only speculate as to what
・17分子でrotationaltemperaturaとcolumn
biasedfrequency surveys have been
conducted in the cm-wave densityを計算。
else may be observed. The enormous potential of cm-wave asregime.
A
4.4
GHz
bandwidthsurvey
from
17.6
to
22.0
GHz
totronomy for detecting new molecules in space has been high・H2に対する分子の存在比を見積もった。
ward the massive star-forming region W51 has been reported by lighted in the recent past through GBT detections of complex or・NH3のnon-LTE
temperatureの内側のエンBellet al. (1993b) showinganalysisにより、その放射の大半が70+-20Kのkinetic
many previously unidentified lines. ganic molecules such as CH CHCHO(propenal), CH3 CH2 CHO
2The dark cloud TMC-1 has been observed in selected bands (propanal), and CNCHO (cyanoformaldehyde; Hollis et al.
ベロープに起因することを示した。HC5Nの様々な13C同位体の間に存在比の違いを評価した。全般的に、from 4−6 and 8−10 GHz (Kalenskii et al. 2004), as well as at
2004; Remijan et al. 2008).
9−50
GHz(Kaifuetal.
2004).
More
recently,somesystematic
12C/13C比は49+-9と推定された。
For this paper, we performed a 1.3 cm line survey toward
observations toward the galactic center have also been reported
by・最後に、1.3cm帯の検出された分子に着いてTMC-1との比較を行い、IRC+10216でSi系分子がよRemijan et al. (2008) in the course of their GBT PRIMOS IRC +10216, which is also known as CW Leonis. It is found tobe the brightest object outside the solar system at 5 µm (Becklin
project1 .
り多いこと示した。
4
branch (AGB)
Spectroscopic features that are expected at long wavelengths et al. 1969). It is a well-studied asymptotic giant −5−1Table 1; これまでのIRC+10216でのLine surveyの研究；今回の観測周波数帯域は十分に調べられてきていない。

carbon chain molecule (eg. HC5N)の低励起輝線が多く検出された。またこれらの13C同位体も検出された。23輝線は太陽系外で初検出で、20輝線(13C同位体の炭素鎖分子など)はこの天体で初検出。（これらの輝線の幾つかは他の周波数帯ですでに検出されている）

5

柱密度とTrotを求めた
・Trot; 5.5 - 47.2 K
・N; 5.2E12 - 2.4E15 cm^-2!
例えばc-C3H2は最も低いTrotを持
つがそれは外側の冷たいエンベロー
プに存在することを示唆する。

!
Nはより複雑な分子で低い存在比の傾向がある。

ー＞ﬁg 5; IRC+10216とTMC-1の比較；
IRC+10216はSi系分子が多く存在する
6

HNC IN PROTOPLANETARY DISKS
Dawn Graninger1, Karin I. Öberg1, Chunhua Qi1, and Joel Kastner2

1• 23. HNC
in Protoplanetary
Disks
Harvard-Smithsonian
Center for Astrophysics,
60 Garden Street, Cambridge, MA 02138, USA; dgraninger@cfa.harvard.edu
2Center for Imaging Science, School of Physics and Astronomy, and Laboratory for Multiwavelength Astrophysics, Rochester Institute of Technology,
Graninger, Dawn; Öberg, KarinI.; Qi,
Chunhua;
Kastner,Joel USA
54 Lomb
Memorial
Drive, Rochester,
NY 14623,
Received 2015 February 19; accepted 2015 May 28; published 2015 July 1

http://adsabs.harvard.edu/abs/2015ApJ...807L..15G (ApJ Letter)
ABSTRACTThe distributions and abundances of small organics in protoplanetary disks are potentially powerful probes of disk
physics and chemistry. HNC is a common probe of dense interstellar regions and the target of this study. We use
the Submillimeter Array (SMA) to observe HNC 3–2 toward the protoplanetary disks around the T Tauri star
TW Hya and the Herbig Ae star HD 163296. HNC is detected toward both disks, constituting the ﬁrst spatiallyresolved observations of HNC in disks. We also present SMA observations of HCN 3–2 and IRAM 30 m
observations of HCN and HNC 1–0 toward HD 163296. The disk-averaged HNC/HCN emission ratio is 0.1–0.2
toward both disks. Toward TW Hya, the HNC emission is conﬁned to a ring. The varying HNC abundance in the
TW Hya disk demonstrates that HNC chemistry is strongly linked to the disk physical structure. In particular, theinner rim of the HNC ring can be explained by efﬁcient destruction of HNC at elevated temperatures,
similar to what is observed in the ISM. However, to realize the full potential of HNC as a disk tracer requires a
combination of high SNR spatially resolved observations of HNC and HCN and disk-speciﬁc HNC chemical
modeling.
Key words: astrochemistry – ISM: molecules – protoplanetary disks
原始惑星系円盤でのHNC(HCN)の観測から円盤のproperties(構造とか温度とか)を調べたい。
対象；(a)TW Hya, (b)HD163296
1. INTRODUCTION
nitrogen reservoir. Cyanides and iso-cyanides are also
望遠鏡；SMA: HNC 3-2(a,b) , SMA: HCN 3-2 (b)important for some prebiotic chemistry scenarios (Powner
The disks around youngstars
containthe gas and
dust
IRAM
30m:
HCN/HNC
1-0(b)from
et al. 2009). Observing HNC in disks has proven challenging,which planets form. From theory (Aikawa et al. 1996; Bergin(a,b) however. The intrinsic disk scales are small and the molecularet al. 2003)結果；HNC/HCN
and observations emission(Rosenfeld ratio;
et al.0.1-0.2
2012; Qiline ﬂuxes are weak compared to interstellar clouds andet al. 2013a), disks(a)areTWcharacterized
by steep temperature andHyaでHNCはリング構造。
protostars. To date, HNC has only been detected in a single
radiation gradientsHNCのリングの内縁はHNCの高い温度での効率的な破壊によって説明できる
both radially and vertically. These disk
disk, DM Tau (Dutrey et al. 1997), and no spatially resolved
structures affect planet formation (Takeuchi et al. 1996; Kretke
observations exist.
議論の中で、リング構造を作る機構の可能性を三つ列挙（後述）& Lin 2012),
motivating the quest for new emission-line
In this Letter, we report the ﬁrst interferometric detection ofdiagnostics that canこれらを切り分ける為には空間的に分離し高いSNRの観測とHNC/HCNの円盤に特化したtrace physical structure gradients in disks.
HNC in two disks, around the T Tauri star TW Hya
This study focuses on HNC, a common high-density gas(D = 54 pc; Qi et al. 2004) and the Herbig Ae star HD
化学モデルが必要である（ALMAで分かるだろう）
tracer in the ISM. At low temperatures, HNC forms efﬁciently
163296 (D = 122 pc; Qi et al. 2011), using the Submillimeter
together with HCN through ion–molecule reactions. HNC andArray (SMA). We combine these observations with new and
HCN are observed in a variety of different sources, but theirprevious SMA observations of HCN toward TW Hya and HD
relative abundances can vary by more than an order of 7 163296 and new IRAM 30 m observations of HNC and HCN

•

観測結果

Graninger et al.

Figure 1. 1.1 mm continuum and HNC 3–2 emission (upper panels) from TW
Hya using the compact + extended conﬁgurations of the SMA. Continuum
contours are [10, 20, 40,K] mJy, and the contours for HNC are [4, 6, 8]σ with
1σ = 0.064 Jy km s−1 beam−1. The lower left panel displays the spectrum
obtained from the HNC 3–2 image. The lower right panel displays the
deprojected, azimuthally averaged ﬂux as a function of radius.

TW Hya disk. This agrees with the ∼1″ separation between thetwo ﬂux peaks in the HNC-integrated emission image. The
depth and size of the hole may both be affected by the image
asymmetry, however. In Section 3.2, we evaluate the evidenceof an HNC ring and quantify its location based on an analysis
in the (u , v ) plane.
The HNC spectrum in Figure 1 was extracted using a
2.´ 5 q 2.´ 5 box. The integrated line ﬂux is 1.2 Jy km s−1 beam−1(∼12σ). Using the integrated HNC ﬂux from this study and theHCN ﬂux from Qi et al. (2008), the disk-averaged HNC/HCN
emission ratio is 0.18 ± 0.06, where the error is derived using the1σ rms for the two observations.
HNC is also detected toward HD 163296, but with a lower
SNR. Figure 2 displays the 1.1 mm continuum, HNC, and
HCN emission toward HD 163296, as well as the spectra andspatial velocity structure of emission from these molecules. The
velocity-integrated emission maps yield a maximum HNC andHCN ﬂux of 0.71 and 2.71 Jy km s−1 beam−1, respectively. The
HNC emission toward the HD 163296 disk is suggestive of aring viewed at a relatively high inclination, but the low SNR
inhibits any further spatial analysis.
Spectra were extracted using an elliptical mask, which
increased the SNR compared to a square mask. The total
integrated ﬂux from these spectra is 0.66 Jy km s−1 beam−1(∼4σ) and 4.47 Jy km s−1 beam−1 (∼9.5σ) for HNC and HCN,
respectively. This yields an HNC/HCN ratio of 0.14 ± 0.05,
similar to TW Hya.
The IRAM 30 m observations of HNC and HCN 1–0 toward
HD 163296 should be more sensitive to the coldest disk
material (Figure 3). The integrated intensities of HNC and
HCN were found to be 0.033 K km s−1 (3σ upper limit) and
0.104 K km s−1, respectively. This yields an upper limit on the
HNC/HCN ratio of 0.3, consistent with the SMA data.

Figure 2. 1.1 mm continuum (top left; contours = [20, 40, 80,K] mJy), integrated HN
162936 together with extracted spectra (bottom left) and integrated emission in two
Graningeral.
ﬂux contoursetare[2, 3]σ (HNC) and [3, 5, 7,K]σ (HCN) with 1T  0.21 and 0.14
bins are [3, 4]σ (HNC) and [8, 10, 12,K]σ (HCN).

The Astrophysical Journal Letters, 807:L15 (5pp), 2015 July 1

Figure 3. IRAM 30 m spectra of HNC and HCN 1–0 in HD 163296. The
dashed line at 5.8 km s−1 indicates the systemic velocity of HD 163296.
Figure 2. 1.1 mm continuum (top left; contours = [20, 40, 80,K] mJy), integrated HNC 3–2 (top middle), and HCN 3–2 (top right) ﬂuxes toward the disk around HD
3.2. HNC Line Modeling toward TW Hya
162936 together with extracted spectra (bottom left) and integrated emission in two velocity bins (bottom middle and bottom right). The HNC and HCN integrated
ﬂux contours are [2, 3]σ (HNC) and [3, 5, 7,K]σ (HCN) with 1T  0.21 and 0.14 Jy km s−1 beam−1, respectively. The HNC and HCN contours in the two velocity
To constrain the spatial extent of the HNC emission in TW
bins are [3, 4]σ (HNC) and [8, 10, 12,K]σ (HCN).

Hya, we follow the procedure outlined in Qi et al. (2008) andet al.
2015) to extract molecular column density
midplane(Tm(2012,)
constant abundance, bound by surface (Ts) and Öberg
proﬁles.
First,
parametric
HNC abundance models are deﬁnedboundaries deﬁned with respect to the vertically integrated
with
set respect
the HNCto the TW Hya disk temperature and density
hydrogen column in units of 1.59 × 1021 cm−2. We
structures from Qi et al. (2013b). The line emission is modeledboundaries using the values from Qi et al. (2008) for HCN (logusing the two-dimensional non-LTE Monte Carlo radiative
Ts = −0.5, log Tm = 1.5). The radial distribution of HNC wastransfer code RATRAN (Hogerheijde & van der Tak 2000)
modeled using an abundance power law with anouterradius cross-sections from Dumouchel et al. (2010).
using
collisionalof −1.5.
(Ro), hole radius (Rh), and power-law index The
HNC Theabundance models follow the form in Qi et al.
andparameter Ro was varied between 35 and 60 AU
h was assumes a molecular layer with a vertically(2008),Rwhich
varied between 20 and 50 AU. The minimum D 2 value wascalculated from the modeled and observed HNC emission in3
the (u , v ) plane using the data from the compact and extended
conﬁguration.
Figure 4 shows examples of emission models and observations in the extended conﬁguration. The best-ﬁt model based on
the D 2 value is a 10 AU wide ring with an Ro of 45 AU and an
Rh of 35 AU (Figure 5), shown in the lower left panel of
Figure 4. This model results in no signiﬁcant residuals whensubtracted from the data in the (u , v ) plane and then imaged.
The observed HNC azimuthal asymmetry toward TW Hya isthus not signiﬁcant. The presence of an HNC ring in the TW
Hya disk is by contrast a statistically signiﬁcant result. Figure 5
shows the models that ﬁt the data within 2σ and 3σ. Within 2σ,
the hole size is 28–45 AU (3σ: 26–47 AU), and the outer edge
is at 37–55 AU (3σ: 57 AU).

・TW Hya (Fig1); SMA: HNC 3-2(Fig1右下); リング平均フラックス（inclination=6deg, PA=155deg; Qi+2004, Williams&Cieza2011）=>中心にholeがある(0.5”)
TW Hya; HNC/HCN=0.18+-0.06 (HCNはQi+2008) [note: disk-averaged line ratio]!

・HD163296 (Fig2);
SMA: HNC/HCN 3-2 ＝＞HNC/HCN=0.14+-0.05 [note: disk-averaged line ratio]
HNCで若干リング構造っぽく見えるがSNRが低いので空間構造の解析は行っていない。
Figure 3. IRAM 30 m spectra of HNC and HCN 1–0 in HD 163296. The
dashed line at 5.8 km s indicates the systemic velocity of HD 163296.
IRAM 30m: HNC/HCN 1-0 ＝＞HNC/HCN＜0.3 (upper limit)

!
!

−1

3.2. HNC Line Modeling toward TW Hya
To constrain the spatial extent of the HNC emission in TW
Hya, we follow the procedure outlined in Qi et al. (2008) andÖberg et al. (2012, 2015) to extract molecular column density
proﬁles. First, parametric HNC abundance models are deﬁned
with respect to the TW Hya disk temperature and density
structures from Qi et al. (2013b). The line emission is modeledusing the two-dimensional non-LTE Monte Carlo radiativetransfer code RATRAN (Hogerheijde & van der Tak 2000)
using collisional cross-sections from Dumouchel et al. (2010).
The HNC abundance models follow the form in Qi et al.
(2008), which assumes a molecular layer with a vertically8

4. DISCUSSION
In this study, we detect and spatially resolve HNC emission
toward TW Hya and HD 163296. From these observations, we
extracted disk-averaged HNC/HCN emission ratios and theinner and outer HNC emission radii in the TW Hya disk. Here,we will discuss the constraints that these observations provideon the HNC chemistry and the disk structure.

The Astrophysical Journal Letters, 807:L15 (5pp), 2015 July 1

the small-scale emission structures.

Graninger et al.

which is reasonable in cloud cores, but may
approximation in disks (Meijerink et al. 200(2014) have also shown that the HNC/HCN
the C/O ratio, which may be elevated in disks.
detailed disk chemistry models that incorpor
HNC/HCN chemistry, disk-averaged line rat
useful probes of disk properties.

First, HNC and HCN are not co-spatial—some HCN ispresent toward the center of the disk, where HNC is absent.
This may be partially mitigated by averaging the HNC and
HCN emission only across disk areas where HCN and HNCemission are respectively detected. Second, the relationshipbetween line emission and abundance ratios is complicated byhigh HCN line opacities and different HCN and HNC
4.2. The HNC Inner and Outer Radii in the
excitation properties (Sarrasin et al. 2010). We used the large
In the disk toward TW Hya, there is a lack ovelocity gradient radiative transfer code, RADEX (van der Tak
interior to 28–45 AU, corresponding to tem
et al. 2007), with the collisional rate constants from Dumouchel
21–13 K using the Qi et al. (2013b) temperat
et al. (2010) to evaluate the relationship between thedoes not have an emission hole (Qi et al. 200
HNC/HCN line emission and the abundance ratios using
absence of HNC close to the star istemperatures and densities characteristic of HCN disk emission
temperature-dependent
HNC
destructio(Walsh et al. 2010) and the HCN column density from Qi et al.
et al. 2014). It is difﬁcult to explain the hole w(2008). We ﬁnd that the line and abundance ratios can differ by
UV ﬂux present at smaller radii since HNC
up to an order of magnitude, demonstrating that more detailed
thought
to have similar photodissociation
modeling is required to extract HNC/HCN abundance ratios(Roberge et al. 1991; van Dishoeck et al. 200
from the observed line emission.
HNC also has an outer emission radius
Even if abundance ratios are successfully extracted, it may
55 AU, with the best ﬁt at 45 AU, substantiallybe difﬁcult to infer a characteristic emission region temperaobserved radii of other gases, including CO, H
ture. Other environmental characteristics than temperature may
Figure 5. D 2 contour plot for the best-ﬁt models. The contours are 2σ and 3σ.
Figure 4. Comparison of HNC 3–2 data (upper left panel) and models with an(Qi et al. 2008). This outer cutoff in HN
regulate
theisHNC/HCN
ratiotheinbest-ﬁt
disks.model.
The The
model
by
The
3σ contour
in red, and theabundance
red cross marks
blackouter (Ro) and hole (Rh) radius. The contours are [3, 4, 5]σ, based on theunexpected. It coincides with the onset of C2HGraninger
al. of(2014)
inﬂuenceof radiation,and theshaded region
indicates
crossesmark theetvalues
modeledneglects
Rh and Ro,theobserved rms noise. Images are generated using the EXT data only to display
is proposed to be a UV photochemistry pRoreasonable.
whereRh⩾iswhich
in cloud cores, but may not be a goodthe small-scale emission structures.
et al. 2015) and is close to the radius whereapproximation in disks (Meijerink et al. 2007). Loison et al.
sharply decreases (Andrews et al. 2012)(2014) have alsothat Emission
the HNC/HCN
4.1. shown
HNC/HCN
Ratiosratio depends oncoincidences suggest three potential explanati
the C/O ratio, which may be elevated in disks. In the absence ofThe disk-averaged
HNC/HCNemissionratio theis complete
0.1–0.2
cutoff of HNC emission: (1) UV photodisso
detailed
disk chemistrymodels thatincorporate
toward
both
TW
Hya
and
HD
163296.
Based
on
ISM
due to deeper penetration of stellar and interste
HNC/HCN chemistry, disk-averaged line ratios are thus notobservations(Schilkeal. 1992) and astrochemistry models
in the low column density outer disk, (2) auseful probesof disketproperties.
(Graninger et al. 2014), the HNC/HCN ratio can be used tophotodissociation of H2 and carbon-bearing
・TW Hya に関しては、HNCのmodelingを行った(Fig4)
erating the HNC+H and HNC+C destructio
constrain the characteristic temperature of the emission region
4.2. The
HNC Innerand Outer
Radiibyin hydrogen
the TW Hya
Diskatgiven temperature (Graninger et al. 2014; Loi
becauseof
efﬁcient
destruction
of
HNC
atoms
HNCのabundanceをQi+2013bのtemperatureとdensityを使って決めた
apply
straightforand (3) grain-gas thermal decoupling due totemperaturestoward25 K. However,In the disk
TW Hya, to
there
is athislackresultof HNCemission
wardlyrequiresthat AU,
HNC
and HCN are co-spatial,that thereisgrain densities, causing a temperature reversal
interior
to code28–45
corresponding
to temperatures
above
non-LTE Monte Carlo ratdiative transfer
RATRANで放射を計算
a 21–13
simple K
relationshipbetweenemissiontemperature
and abundance
ratiosHCN
ofand thus the increase in thermally controlledusing the Qi
et al. (2013b)
proﬁle.
変数；Ro: outer radius [35-60au],HNC
Rh:
holeradius[20-50au]を振ってu,v空間でχ自乗検定
and
HCN,
and
that
temperature
dominatesthechemistry,(i.e., a similar process that sets the inner
does not have an emission hole (Qi et al. 2008). The observed
noneof whichhold intothethe
casestarof thesedisks.
emission).
absenceof clearly
HNC close
is consistent
with＝＞結果として、χ自乗が一番いいところで幅10AU,
Ro=45AU,
Rh=35AUのリング分布となった。(Fig5)
temperature-dependent
HNC
destruction(Graninger
et al. 2014). It is difﬁcult to explain the hole with the increased 4[ hole size ; 2σ: 28-45AU (3σ:26-47AU) ]UV ﬂux present at smaller radii since HNC and HCN arethought ＜57AU)
to have similar
photodissociation cross-sections[ outer edge ; 2σ: 37-55AU (3σ:](Roberge et al. 1991; van Dishoeck et al. 2006).
HNC also has an outer emission radius between 35 and
55 AU, with the best ﬁt at 45 AU, substantially smaller than theobserved radii of other gases, including CO, HCO+, and HCN
2Figure 5. D contour plot for the best-ﬁt models. The contours are 2σ and 3σ.
(Qi et al. 2008). This outer cutoff in HNC emission was
The 3σ contour is in red, and the red cross marks the best-ﬁt model. The black
unexpected. It coincides
with the onset of C2H emission, whichcrosses mark the values of modeled Rh and Ro, and the shaded region indicates9
is proposed to be a UV photochemistry product (Kastnerwhere Rh⩾Ro.

•議論[memo;HNCについて]・prestellerやprotostellerの環境 (10K)では、HNC/HCN ~ 1
温度が上がるとHNC/HCNの比は小さくなる＝＞HNC/HCNの比は温度を反映する
・低い温度でもC原子(H原子?)によりHNCは破壊される・shieldがあまりされていない領域では放射により破壊される (UV)!
[inner radii]
・HCNはHNCより内側までいる。
HNCの内縁は28-45AUで、Qi+2013bによる温度分布では21-13Kであり、HCNが破壊される温度と一致する。[温度が>25Kより大きいところで、HNCは効率的に破壊される][photodissociationでは、HCNとHNCが異なる内縁半径を持つことを説明できない][outer radii]
・HNCの外縁半径は35-55AU (best ﬁt:45AU)で、COやHCO+, HCNで観測されている半径よりもかなり小さい・他の観測では、C2Hでも同様に見られ、UV photochemistryによるものと示唆されている
また、dustの密度の急激な減少している半径と近い
これらのことからHNCの外縁半径について以下の三つの可能性が考えられる。
(1). HNCのUV photodissociation (外縁部では柱密度が小さく、星や星間のUVが入ってくる)(2). H2とC系のUV photodissociation の増加で、HNC+H, HNC+Cの破壊反応が促進される(3). gasとdustが低密度のため、grain-gas thermal decoupling が生じていて、HNCが高い温度となり破壊反応が進んでいる(内縁と同じ機構)。
＝＞これらを切り分ける為には空間的に分離し高いSNRの観測(ALMA)とHNC/HCNの円盤に特化した化学モデルが必要である

!
!

10

:1502.00183v3 [astro-ph.GA] 18 Jun 2015

Audra K. Hernandez
Department of Astronomy, University of Wisconsin, 475 North Charter Street, Madison, WI 53706, USA;
hernande@astro.wisc.edu•

24. The Giant Molecular Cloud Environments of Infrared Dark Clouds
Jonathan C. Tan

Hernandez, Departments
Audra K.;
Tan, Jonathan
of Astronomy& Physics, C.
University of Florida, Gainesville, FL 32611, USA;jt@astro.ufl.edu

http://adsabs.harvard.edu/abs/2015arXiv150200183H(ApJ)
Draft version June 19, 2015
ABSTRACT
We study Giant Molecular Cloud (GMC) environments surrounding 10 Infrared Dark Clouds(IRDCs), using 13 CO(1 0) emission from the Galactic Ring Survey. We measure physical propertiesof these IRDCs/GMCs on a range of scales extending to radii, R, of 30 pc. By comparing di↵erent
methods for defining cloud boundaries and for deriving mass surface densities and velocity dispersions,
we settle on a preferred “CE,⌧ ,G” method of “Connected Extraction” in position-velocity space plus
Gaussian fitting to opacity-corrected line profiles for velocity dispersion and mass estimation. We
examine how cloud definition a↵ects measurements of the magnitude and direction of line-of-sightvelocity gradients and velocity dispersions, including associated dependencies on size scale. CE,⌧ ,Gdefined GMCs show velocity dispersion versus size relations / s1/2 , which are consistent with thelarge-scale gradients being caused by turbulence. However, IRDCs have velocity dispersions thatare moderately enhanced above those predicted by this scaling relation. We examine the dynamical
state of the clouds finding mean virial parameters ↵¯ vir ' 1.0 for GMCs and 1.6 for IRDCs, broadly
consistent with models of magnetized virialized pressure-confined polytropic clouds, but potentiallyindicating that IRDCs have more disturbed kinematics. CE,⌧ ,G-defined clouds exhibit a tight correlation of /R1/2 / ⌃n , with n ' 0.7 for GMCs and 1.3 for IRDCs (c.f., a value of 0.5 expected for a
population of virialized clouds). We conclude that while GMCs show evidence for virialization overa range of scales, IRDCs may be moderately super virial. Alternatively, IRDCs could be virialized
but have systematically di↵erent 13 CO gas phase abundances, i.e., due to freeze-out, a↵ecting massestimations.
Subject headings: ISM: clouds - ISM: kinematics and dynamics - Stars: formation

IRDCsからGMCの接続を調べたい。
1. INTRODUCTION＝＞Galactic ring surveyの13CO(1-0)のデータを用い10個のIRDCsを調べた。＜30pc内の構造。
Galactic star formation resides mostly within Giant Molecular Clouds (GMCs), conventionally defined to have masses! 104 M and observed to extend up to several ⇥106 M (e.g., Blitz 1993; Williams et al. 2000; McKee & Ostriker2
物理量を求めるのに”CE,tau,G”
connected
G; gaussian2007). With typical mass surface methodがいい。CE;densities of ⇠ 100 M pc
, GMCs extraction,
have mean radial
sizes of ⇠ﬁtting
6 100 pc, assuming
simple circular symmetry. However, GMCs are highly irregular and hierarchical structures. Their dense clumps canspawn stellar clusters and associations, creating the bulk of the Galactic field star population. The efficiency and rateof star formation from these clumps is relatively low, i.e.,11a few percent per local free-fall time (Zuckerman & Evans
1974; Krumholz & Tan 2007). This appears to be mostly because much of the GMC material is stable with respect to

80

0.00

60
28.8

28.5

28.8

l

95
90
85
80
75
70

1.5
1.0

K deg

100

vlsr (km s−1)

I [K km s−1]

E

b

0.25

120

0.5
0.0

28.5

28.8

l

28.5

0.0

l

0.2

28.8

b

28.5

0.0

l

0.2
b

Fig. 1.— The IRDC/GMC sample: each row presents a single cloud (A through E) showing data out to a radius of 30 pc. Left to right:
GLIMPSE 8 µm image with the IRDC boundary marked with a pink ellipse (see Figs. 3 and A1-A4 for intensity scales); 13 CO(1-0) GRS
integrated intensity map (over the range v0 ± 15 km s 1 ) in units of K km s 1 ; Full projections of the cloud in position-velocity space:
simple-extraction (SE) 13 CO(1-0) vlsr vs l map, SE 13 CO(1-0) vlsr vs b map; connected extraction (CE) 13 CO(1-0) vlsr vs l map; CE
13 CO(1-0) v
lsr vs b map. The IRDC is marked with a white box of width equal to the extent of the elliptical boundary and height equal
to the FWHM of the IRDC co-added spectra. The solid black line shows the mass-weighted linear velocity gradient (dv0 /ds) across eachcloud at R = 30 pc.

3. METHODS: DERIVATIONS OF MOLECULAR CLOUD PROPERTIES
3.1. Definition of Cloud Boundaries
物理量を導出するのにいくつかの領域を選定する方法を試した(Table2)
We
firstsimple
determinedthe cloudcenter-of-mass in position-velocity space. To do this, we co-added the 13 CO spectra
・SE;
extraction(領域内の13COemissionをすべて使う）
within an elliptical boundary defined by twice the rmaj and rmin values, using the full GRS velocity range. Then, a
領域サイズを5,10,20,30pcと変化。
velocity intervalof 30 km s 1 was considered, centered on the peak of the emission profile. The cloud’s center-of-mass
was・CE;evaluatedfor both extraction
optically thinand opacity corrected gas, (l0 , b0 , v0 ) and (l0,⌧ , b0,⌧ , v0,⌧ ), respectively. We thenconnected(PV上で繋がっている領域のみを使う）；
use this as the reference point around which to search for GMC-scale gas, out to a projected radius of 30 pc and over a
抽出された領域で平均質量半径,領域半径、半分の質量半径で計算。
new velocity
range of v0 ± 15 km s 1 . This range is wide enough to sample the velocities of GMC gas, even those thatare potentially undergoing collisional interactions at velocities set by Galactic shear at the tidal radii of the clouds(Tan 2000; Tasker & Tan 2009). The GLIMPSE 8µm and GRS 13 CO integrated intensity maps centered on theselocations are presented in Figures 1 and 2.
We examined 13 CO emission within the v0 ± 15 km/s velocity range and within apertures of varying radii from thecloud center-of-mass coordinate. To explore how boundary definitions a↵ect estimated physical properties, we used
two di↵erent methods to select 13 CO line emission associated with the cloud. First, “Simple Extraction (SE)” selected
TABLE 2
GMC Physical Parameters Estimates
IRDC
Case

lCM ,lCM,⌧

bCM ,bCM,⌧

vCM ,vCM,⌧
1)

(pc)

(

)

( )

( km s

(1)

(2)

(3)

(4)

5

18.856,18.868-0.282,-0.276

64.3,64.3

10

18.862,18.875-0.282,-0.282

63.9,64.1

20

18.868,18.881-0.301,-0.313

63.5,63.7

30

18.868,18.881-0.307,-0.319

RM

18.881,18.893

RA

18.881,18.893

R1/2

18.881,18.893

s

v

R
1)

(pc)

M ,M⌧ ,MG
⇥104

(pc)

( km s

(M

(5)

(6)

(7)

(8)

2.06,2.12

0.0,0.0

5.00

2.58,2.57-0.4,-0.2

10.0

3.46,4.02-0.9,-0.6

62.6,63.1

3.72,4.37-0.325,-0.337

63.3,63.5-0.325,-0.337

63.3,63.5-0.325,-0.337

63.3,63.5), ⌧, G( km s1)

dv0 dv0,⌧
,
ds
ds
⇥10 2

✓v ,✓v,⌧

I0 ,I0,⌧⇥103

Erot,0 ,Erot,0,⌧
⇥1046

(M

pc2 )

,
(10

⌧
2)

log(↵vir ),log(↵vir,⌧ ),log(↵vir,G )(km/s/pc)

(deg)

(erg)

(9)

(10)

(11)

(12)

(13)

(14)

(15)

1.10,1.87,1.62

3.92,3.48,1.66

36.0,32.3

141,1240.0814,0.131

1.37,4.01

7.62,3.39

0.912,0.575,-0.00704

4.83,7.20,5.73

5.12,4.53,1.94

15.2,16.1

124,123

1.36,1.97

13.3,29.6

2.34,1.72

0.802,0.521,-0.115

20.0

16.1,24.0,19.2

5.87,5.15,2.36

16.8,15.6

139,142

16.8,25.7

74.2,165

6.38,3.78

0.696,0.411,-0.170-1.7,-1.3

30.0

31.0,44.1,34.9

6.58,6.01,3.02

12.7,12.1

96.8,96.3

68.2,96.7

183,370

5.97,3.80

0.688,0.457,-0.0389

5.10,6.60-1.7,-1.3

17.4,17.0

9.63,16.7,15.7

3.38,2.88,2.00

17.5,14.0

151,153

7.63,12.9

30.5,94.2

7.67,2.69

0.380,-0.00984,-0.298

5.10,6.60-1.7,-1.3

26.6,26.6

16.6,28.7,26.9

3.88,3.36,2.33

16.5,14.0

133,134

32.3,54.3

59.3,177

14.8,5.97

0.447,0.0862,-0.204

5.10,6.60-1.7,-1.3

17.1,16.2

9.38,15.5,14.7

3.34,2.81,1.99

17.2,13.5

151,153

7.18,11.0

29.4,84.9

7.18,2.36

0.374,-0.0177,-0.291

A

B

12

0.4

IRDC
A
B
C
D −s
E
F −s
G −s
H −s
I
J

CE,τ,G

SE,thinCE,thin

1.0
0.5
0.0

IRDC
A
B
C
D −s
E
F −s
G −s
H −s
I
J

GMC
0.2

log(σ/R1/2 / km s−1 pc−1/2)

1.5

−0.5

log(σ / km s−1)

1.5

SE,τ

CE,τ

0.0
=8
α vir

−0.2

−0.4

1.0

=2
α vir

−0.6
.5

=0
α vir

1.4

0.0

1.6

1.8

2.0

2.2

2.4

2.6

log (Σ /MO· pc−2)

−0.5

1.0

=4
α vir

=1
α vir

0.5

1.5

GMC

SE,τ,GCE,τ,G

Fig. 10.— Dependence of /R1/2 with mass surface density, ⌃, for the all clouds defined by CE, opacity corrected column densities,
and gaussian fitted velocity dispersions. As in Figure 9, IRDCs are shown by the filled triangles and the GMCs by the filled circles (seelegend). The error bar in the upper-left corner represents the mean uncertainties of ⇠ 14% in /R1/2 and ⇠ 30% in ⌃. The results of
RD10, a total of 329 clouds with M > 104 M , are shown by the gray crosses. The mean results of Solomon et al. (1987) are also shown(open black square). The best-fit power-law relation ( /R1/2 / ⌃n ) is shown for both the IRDCs (n = 1.27 ± 0.52; dot-dashed line) and
the GMCs (n = 0.70 ± 0.12; dashed line). Collectively, the IRDCs and GMCs have a best-fit power-law relation of n = 0.81 ± 0.10 (solidline). The dotted lines represents the expected scaling relation for virialized clouds with di↵erent values of ↵vir = 0.5, 1, 2, 4, and 8.

uncertainty in /R1/2 of ⇠ 14% based on 10% uncertainties in both and R.
Thus we considered the above scaling relationship for the results of our ten IRDCs and GMCs (Fig. 9). For IRDCs,
0.5⌃ was estimated for each ellipsoidal cloud using the geometric mean observed radius R and the total 13 CO-derived
mass. For GMCs, ⌃ was estimated using the circular area defined by the extraction radius R and its corresponding
13
CO-derived mass. We estimate a random uncertainty in ⌃ of ⇠ 30% based on the uncertainty in Tex and potential
0.0
variations in our assumed abundances. However, systematic uncertainties could be present at the factor of ⇠ 2 level.
For the ten GMCs, we find that the clouds analyzed using gaussian profile fitting on the opacity corrected column
density spectra to estimate the cloud mass and velocity dispersion, yield results most consistent with the scaling−0.5
expected for gravitational virial equilibrium. Specifically, we find the best-fit power-law relation /R1/2 / ⌃n has−0.5 0.0 0.5 1.0 1.5 −0.5 0.0 0.5 1.0 1.5
n = 0.54 ± 0.05 and n = 0.70 ± 0.12, for the SE and CE extraction, respectively. We gauge the significance of thesefits by estimating the Pearson’s correlation coefficient, ⇢, and the one-sided p-value. For all the GMCs defined by CE
log (R/ pc)
and gaussian profile fitting on the opacity corrected column density spectra, we estimate ⇢ = 0.71 and a p-value of
6.0 ⇥ 10 6 , indicating a significant linearity in the data.
On the other hand, for the ten IRDCs analyzed using gaussian profile fitting on the opacity corrected column
Fig. 6.— Velocity dispersion versus radius R of the clouds over which the dispersion is evaluated. Each panel represents a di↵erent
densityspectra to estimate the cloud mass and velocity dispersion, yield steeper best-fit power-law relations, withcloud definition: (a) Top left: simple extraction, optically thin column density derivation; (b) Middle left: simple extraction,
opacityn =(d)1.57
corrected column densities; (c) Bottom left: simple extraction, opacity-corrected column densities fitted by a Gaussian profile;
Top± 0.80 and n = 1.27 ± 0.52, for the SE and CE extraction, respectively. These power-law fits are lessright: connected extraction, optically thin column density derivation; (e) Middle right: connected extraction, opacity-correctedcolumn than those seen for the GMCs, with ⇢ = 0.56 and a p-value of 0.047 for the “CE,⌧ ,G” method.
significant
densities. (f ) Bottom right: connected extraction, opacity-corrected column densities fitted by a Gaussian profile. In each panel,the 10 on the “CE,⌧ ,G” extraction method, these IRDC and GMC results are displayed together in Fig. 10.
Focusing
IRDCs are shown by the colored filled triangles as defined in the upper right-hand corner. For Clouds D, F, G, and H, these estimates
Overallare based on the single component emission profile. GMCs are shown with filled circles. The best-fit power-law relation from McKeethe& IRDCs and GMCs together yield a best-fit power-law relation of n = 0.81 ± 0.10, somewhat steeper thann =and1/2 expected for virialized clouds. If a power-law relation with n = 0.5 is fit, then the normalization implies
Ostriker (2007) ( = pc (s/1 pc)↵ , with values of ↵ ' 0.5 and pc ' 0.72 km s 1 ) is shown by the dotted lines. The global the(IRDCs
GMCs) best-fit power-law relation for each panel is shown by the solid lines. In panel (f ), the best-fit power-law relations are↵
shown
¯also
1.60 and 1.03 for IRDCs and GMCs, respectively, and ↵¯ vir = 1.15 for the whole sample (with a Pearson’svir =for the IRDCs (dot-dashed line) and the GMCs (dashed line).
correlation coefficient of ⇢ = 0.78 with a p-value of 1.7 ⇥ 10 9 ). We note that some of the dispersion in the data isdue to cloud-to-cloud variation, which might be expected if there is systematic variation in degree of surface pressure
confinement, magnetization, distance uncertainty (a↵ecting R) or 13 CO abundance.

13

1European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germany ???
2
Centro de Radioastronomía y Astrofísica, Universidad Nacional Autónoma de México, Morelia 58090, Mexico
25. Massive envelopes and ﬁlaments in the NGC 3603 star forming region
3•
CSIRO Astronomy and Space Science, Australia Telescope National Facility, P.O. Box 76, Epping, NSW 1710, Australia

Hummel, C. A.; Stanke, T.; Galvan-Madrid, R.; Koribalski, B. S.
http://adsabs.harvard.edu/abs/2015arXiv150608597H (A&Ap)

1506.08597v1 [astro-ph.SR] 29 Jun 2015

Received: <date>; accepted: <date>; LATEXed: June 30, 2015

ABSTRACTThe formation of massive stars and their arrival on the zero-age main-sequence occurs hidden behind dense clouds of gas and dust. In
the giant H ii region NGC 3603, the radiation of a young cluster of OB stars has dispersed dust and gas in its vicinity. At a projecteddistance of 2.5 pc from the cluster, a bright mid-infrared (mid-IR) source (IRS 9A) had been identified as a massive young stellar
object (MYSO), located on the side of a molecular clump (MM2) of gas facing the cluster. We investigated the physical conditions
in MM2, based on APEX sub-mm observations using the SABOCA and SHFI instruments, and archival ATCA 3 mm continuum andCS spectral line data. We resolved MM2 into several compact cores, one of them closely associated with IRS 9A. These are likely
infrared dark clouds as they do not show the typical hot-core emission lines and are mostly opaque against the mid-IR background.
The compact cores have masses of up to several hundred times the solar mass and gas temperatures of about 50 K, without evidence
of internal ionizing sources. We speculate that IRS 9A is younger than the cluster stars, but is in an evolutionary state after that of thecompact cores.
Key words. stars: circumstellar matter - stars: early-type - stars: formation - stars: pre-main sequence - stars: individual: NGC 3603
IRS 9A・giant HII region であるNGC3603の隣の molecular clump (MM2)に対して、
APEX sub-mm Obs (SABOCA and SHFI)ATCA 3mm continuum and CS spectral line data
1. Introduction
disk, being photo-evaporated by a newly formed O star if one
considers the emission lines seen in a MIR Spitzer spectrum ofを用いて調べた。
The formation of high-mass stars is rapid (Zinnecker & Yorke IRS 9A (Lebouteiller et al. 2008). Adopting the distance to the・MM2の周辺環境としては、OBstarsの若いクラスターの放射によってガスやダストが晴らされていて2007) and leaves the new-born star
still enshrouded in gas and cluster, 7.2 kpc (Melnick et al. 1989) for the distance to MM2
クラスターから2.5pc(projected)に明るいmid-IR
9A)；massive young stellarが存在し、dust. Unlike the formationof low-mass stars viasource
simple (IRS
accreand IRS 9A (see also the discussion in de Pree et al. 1999), 1"
tion
disks,
theoreticalstudiesshow
thatfilaments
andnonaxIRS 9AはMM2の隣に位置している。
corresponds to 7200 AU.
isymmetricdisks(Krumholz
et
al.
2009)
that
funneltheradia・結果として、MM2を幾つかのcompact coresに分解し、そのうちの一つはIRS9Aの近くに付随していtive flux into the polar directions (Kuiper et al. 2015) provide
る。
protection for the accreting gas against the uniquely intense stelFig. 14.
Overlay of SABOCA contours on a 5 GHz image of the・これらはinfrared dark cloudである、（これらは典型的なhot-coreemission
H ii region obtained
from ATCA datalineを持たない）
by Mücke et al. (2002). The blue
lar radiation pressure until the final mass has been reached. colored image regions
correspond to the largest 5 GHz fluxes, while theThe (virial)emission.
massTheof
the molecular clump MM2 was deteryellow color denotes the faintest 5 GHzsourceの兆候は見ら
contour levels
・そのcompact
coresは数100Msunの質量を持ち、温度は〜50K(内部にionizing
Looking forcandidate massive young stellar objects

mined by Nürnberger et al. (2002) from observations of CS lines

are at 87%, 73%, 60%, 46%, 33%, 19%, and 06% of the maximum.
The ATCA beam is 2", the SABOCA beam is 7.8" and is shown in thelower left corner.

れない)である。
(MYSO) suitable for observation due to low foreground extinc- to be around 1500 solar masses. Due to the close association
tion, Nürnberger (2003) identified a bright infraredcoresより進化段階が進んでいると推量される
source (IRS of the MYSO IRS 9A and MM2 we performed observations at
・IRS9Aはクラスターの星よりも若いが、compact
the SHFI spectra do not show the typical hot core lines such as

9A: Frogel et al. 1977), on the side of a molecular clump (MM2:those from Methylcyanide. Based on their mid-infrared opacity,
wavelengths
to penetrate the obscuring dust and toclassifysub-mm
these cores as radio
infrared darkclouds even though their
Nürnberger et al. 2002) facing a massive cluster of OB starswe
at
14determine
temperatureof about 50 K istotal
higher than
expectedfor IRDCs.
15.
ATCA CS(2–1) contour overlay
an HSTH-↵ imagedust
massesandFig.physicalconditions
in ontheemisThe mid-IR source IRS 9A is associated with a massive core, (Brandner et al. 2000).
the center of the H ii region NGC 3603. Nürnberger (2003) esti-

SABOCAand
350um
image (MM2)
Observations
data reduction
ATCA 3 mm interferometry

S8

H13 CO+
H2 CS
C2 H
...
H↵
HCN
HCO+

et al.: Sub-mm spectroscopy of NGC 3603 IRS 9A
4-3
346.998347 C.A..
41.6Hummel
3.37
10(1,9)-9(1,8)
348.534250
105.2–
Table 1. ATCA 2005 observing parameters4(7/2,4)-3(5/2,3) 349.399342 41.9
2.884(9/2,4)-3(7/2,3) 349.339067 41.9
2.5
H26↵
353.622747 Rec.
–
Configuration
H75
4-3
354.505469 42.5
1.80
Obs. dates
2005, 5–7 August
4-3
356.734250 42.8
1.62

ATCA
CS 2-10mom

used archival data obtained with the Australia Telescope
Antennas5
mpact Array (ATCA)
in
its
most
compact
configurationof
Sg
Baselines
31,31,43,46,46,55,
size (unresolved by APEX). Increasing the number density ofecular lines in selected regions of the giant Hiiregionthe species instead would have led to saturation of the lines.77,77,82, and 89 m
C 3603 that were carried out in August 2005. ThisTwo
configlines near 230.232
230.841 GHz remained
TotalGHz
obs.andtime
2 ⇥ 6 h (MM2)unidentified.
The closest matches of lines from SO17 O,
S9 on the northern spur as well asCHthe
ion has antennas
eastT
sys respectively, predicted more lines 200–350 K
3 OCH3 , and CH2 CHCN,t track with baselines ranging from 31 to 89 m. ofFor
detailsthese species to be seen
in setting 4, but were not detected. 2 ⇥ 16 MHz
Bandwidth
Sc transitions, CS(2–1)
The fulland
spectra are shown in Fig. 9 and 10, while spectra ofhe observations
see
Table
1.
Two
Sa
ofinchannels
identified lines areNo.
shown
Figs. 5 to 8. In all plots, the 2 ⇥ 256
S(2–1), had been observed simultaneously with athe
bandwith
units are antenna temperaturesin Kelvinversus km/s.
Channelwidth
0.19 km s 1
6 MHz each, split into 256 channels, giving a channel width
Velocity range–15 to +35 km s 1
Sf
1
ATCA
.19 km Se
s . The rest frequencies of the two CS
lines are
3. Results
Center freq. IF1
97969 MHz color
980968 GHz and 96.412982 GHz, respectively. We mosaiced
CS 2-1
Center freq. IF2
96401 MHz3.1. Cores and filaments
molecular clump MM2 (Nürnberger et al. 2002) in the southPrimary
the molecular
cloud beam
clump MM2 into numerous 29 arcsec
part of NGC 3603, using a 3 ⇥ 3 grid, covering We
anresolvedareaofsources(“compactcores”)
bothinthe CS(2–1) line 1momandsubField center↵(J2000)= 11:15:11.5
0continuum
There is a clear correspondence of
1 . To achieve Nyquist sampling the pointings hadmmbeen
sep- emission.
Field
center to their counterparts(J2000) = –61:16:55
the compact cores seenat mm-wavelength
ed by 1500 , about half the primary beam width. The
weatherseen in the sub-mm. The
emissionhas
thegeneralappearanceof 5.6” ⇥ 4.9”
Synthesized beam
Fig. 4. (during
SABOCAthe
350 µmimageof NGC
3603reasonable,
MM2. The sourceand a system
ditions
day)
had
been
clumpyfilaments.
labeled “9” (S9 ) is associated with IRS 9A. This source and the source
RMS perchannel
35 mJy beam 1
The CS(2–1)
in NGC
3603 MM2 ranges fromlabeled “c”of(Sc220) wereKthewastargetsmeasured
for SHFI spectroscopy.
The forsourceperatures(±10%)
antennas
2–5, line emission
1. The mean widthCS(2-1) velocity 0.5
field shows
km s 1 (after smoothing)labeled “g” (Sg ) is at the tip of the eastern pillar (see Figs. 1 of Brandl at least 10 to 20 km s Channel
1
a significantly
higher
value
of
320
K
for
antenna1.
The
anbasicallytwofilaments
a
few
km
s
apart,
which
was
seen
in
et al. (1999) and Mücke et al. (2002)), while the source labeled “8” (S8 )Flux&bandpass
CS
spectraalreadybyNürnbergeretal.
(2002).
One
extends
is
at
the
tip
of
the
westernpillar,both
pointing
towardstheOBcluster.
a efficiency is around 25% at these frequencies. The observThe wedge at the bottom shows the scale for the image flux densities in from source Sg over S9 , calibrator
Sa , and Sf at a systemic velocityof 151253–055 (14.8 Jy)
PKS
procedure
calibrator
Jy/Beam. had comprised a setup on a strong 3 mm
–16
km s 1 (see Fig. 1), while the other one extends from source
Phase calibrator
PKS 1045–62 (0.45 Jy)over Sa to Se at a systemic velocity of about 12 km s 1 . Asigh elevation such as PKS 0537–441, bandpass S9calibration
EtaCar (9.0 Jy)
the two filaments overlap at the position of source S9 , the line
PKS 1253–055, and phase calibration on PKS 1045–62
and
0

5

10

15

20

25

ATCA
C34S 2-1
0momcontour30

in setting 1, we determined an excitation temperature of 45 K as

profiles are double-peaked here, also indicating optically thinarina.
Thereproduced
phase calibrationwasdifficult
PKS
1045–62
emission.
this value
theS9
measured line
ratios.
All lines as
were
Sc
Fig. 1. ATCA CS(2–1) molecular line emission in NGC 3603 MM2. Fig. 2. ATCA C34 S(2–1) molecular line emission in NGC 3603 MM2.
fit
with
this
temperature,
and
a
line
width
of
5
km/s.
Two
of
We
observed
the opticallythin/thick
Table
2. Properties
ofpair
SABOCAS8 andTop: integrated
intensitymap;
the contour
levelsofare lines
0.25, sources.
0.5, 1, 2, 3, 4, The
very
weak and ⌘ Carina is extended. Pointing had been perTop:sources
integratedlabeledintensity map;the contour levels are 0.16, 0.24, 0.32,
1
1
the lines, 13 CO and C18 O, are extended by about a factor of six H13 CO+/HCO+(Figs.
7/8)
which
can
be
used
to tracein1
and 5 JyS9
beam
km s first. Bottom:mean arevelocity
field;
the synthesizedandTable
0.40 Jy beam
km s 1 . Bottom: C34 S(2–1) contours overlaid onto
inthecolumnthesame
aslistedin
1ofLebouteiller
medoveronthe⌘ Carina
every hour,systemat temperatureother line emitting
regions, andwhichthewe adopted
5" in fall (Myerswas
etal. (5.6”
1996;
Klaassen& Wilsonal.panel, and the CS(2–1) emission map. The synthesized beam (5.6”⇥4.9”) is displayed
beam
4.9”)
is displayed
at the 2007;
bottomChen
of et
each
et ⇥al.
(2008).
isleftassociatedwith IRSat 9A.
Theleftsourcesize and✓ inthe grey-scale wedge for data values
grey scale/color
wedges for Source
data valuesS9on theright of each panel.
the bottomof each panel,
ined
every
30
minuteswithapaddle
above
the
3
mm
horn.
Article number, page 4 of 11
the deconvolved
right of each panel.size. At
column 5 is the geometric mean of the axes ofonthe
As part of a flux monitoring program for 3 mm calibrators,
Table 3.theSHFIdistanceobservationalparameters.
The
main
beam efficiencies
of
IRS
9A,
10"
correspond
to
0.33
pc.
The data in the last(⌘mb = 0.75 for APEX-1 and 0.73 for APEX-2) and conversion facng sources such as PKS 1253–055 (3C 279) and PKS 0537–
are41describedin Section
tors (39 three
K/Jy forcolumns
APEX-1 andK/Jy for APEX-2)
are taken3.3.
from grees). Even with the largest beam of 32" of APEX Band-1,
http://www.apex-telescope.org/telescope/efficiency/
are observed regularly with the ATCA. PKS 1253–055 was
the distance between S9 and Sc of 28" is large enough to prevent
However, sources S9 and Sa have a disSa
Se
sured to have a flux density of 14.8 Jy and 13.5 Jy (at 94.5
RA
F mutual✓ ofcontamination.
M
MV at position S9 might be
F measured
Receiver Setting
Range
NoiseDec
PWV
tance
16" sovthat thefluxz) on the 8th and 17th of August, 2005, respectively. We used
GHz
mK
mm
by emission
Jy contaminated"km/s
M from Sa
M, as half of the flux of Sa willbepicked
up
by
the
28"
beam.
APEX-1
1
216.9
–
220.9
10
2.0
e values for the absolute flux calibration.
Sa 2 11:15:12.07
90 9.22.5
498 122229.2 – 233.2 -61:16:58.8
30
2.0
The spectra with a bandwidth of 4 GHz were reduced using
1
Data reduction was carried out in the software package
APEX-2S8 3 11:15:02.88
346.1 – 350.1 -61:15:51.5
30
1.5
55 the6.7
0
CLASS software . 307
In order to identify
the lines, the systemic
353.2 – 357.2
30
0.4
velocityof sourcec wasmeasuredfrom the optically thin isoad (Sault et al. 1995) using standard procedures.
Sc 45 11:15:08.41-61:16:57.2
67
8.2
3.2
369
180
345.8
50
0.3
topologues of CO (C18 O) to be +13 km s 1 . NIST recommended
We used the program plboot to determine gain calibration
S9 11:15:11.34 -61:16:45.2 59 rest7.6
2.5of the lines
3292 were102
frequencies
used to identify the lines listed
in 8.5
Table 4, together
withtheir properties
and the antenna temperors of 2.64 and 2.45 for 97969 & 96401 MHz, respectively,Se
11:15:13.73-61:17:24.3
46
2.5
258
114
12, a 25-pointing map was executed in the CO (3-2) line (setting ature peak ratios between sources Sc and S9 (T c /T 9 ).
m observations of Mars with the H214 array on the 19th
11:15:11.25
45 9.6249 extension
0 (Maret et al. 2011) of
5). DetailsSf
are givenin Table 3. For all -61:17:12.0
observations, a velocity
We used the WEEDS
1resolution
of
0.5
km
s
was
chosen.
August, 2005. Using this factor, we derived 13.5 Jy for
Sg 11:15:10.33 -61:16:16.4 22 CLASS
4.9 to confirm the line
122identifications
0 on source S9 by checkEach observation sequence (except for the map) included ing for each species the presence of lines in any of our waveS 1253–055, in excellent agreement with its 3 mm flux as detwo back-to-back pointings towards sources S9 and Sc (28" dis- length settings. Using the three Formaldehyde lines we detected
tance), sandwiched between two observations of the sky at RA
mined on the 17th of August as part of the calibrator monitorand DEC = -60:09:57.8 which was selected for 1 http://www.iram.fr/IRAMFR/GILDAS
project (see http://www.narrabri.atnf.csiro.au/calibrators/).=its 11:19:58.5
absolute
calibrationwasestablishedwith2 http://physics.nist.gov/cgi-bin/micro/table5/start.pl
observations of VY
low 100µm emission(IRAS) a shortdistanceaway (1.2 deUsing natural weighting and mosaicing we made line cubes CMa and B13134 from the APEX primary and secondary cal- Article number, page 3 of 11
of the absolute calibration is estimatedcontinuum maps for NGC 3603 MM2. We used a veloc- ibrator list. The accuracy15
1esolution of 0.5 km s for the cubes and derived the mean to be 10%. The map is shown in Fig. 4.