An ALMA Survey of Protoplanetary Disks in the σ Orionis Cluster
M. Ansdell1, J. P. Williams1, C. F. Manara2, A. Miotello3, S. Facchini4, N. van der Marel1, L. Testi5,6, and E. F. van Dishoeck3,4
The Astronomical Journal, 153:240 (15pp), 2017 May

The σ Orionis cluster is important for studying protoplanetary disk evolution, as its
intermediate age (∼3–5 Myr) is comparable to the median disk lifetime. We use ALMA to
conduct a high-sensitivity survey of dust and gas in 92 protoplanetary disks around σ Orionis
members with M*
0.1 Me. Our observations cover the 1.33 mm continuum and several CO
J = 2–1 lines: out of 92 sources, we detect 37 in the millimeter continuum and 6 in 12CO, 3 in
13CO, and none in C18O. Using the continuum emission to estimate dust mass, we find only
11 disks with Mdust
10 M⊕, indicating that after only a few Myr of evolution most disks
lack sufficient dust to form giant planet cores. Stacking the individually undetected continuum
sources limits their average dust mass to 5× lower than that of the faintest detected disk,
supporting theoretical models that indicate rapid dissipation once disk clearing begins.
Comparing the protoplanetary disk population in σ Orionis to those of other star-forming
regions supports the steady decline in average dust mass and the steepening of the Mdust–M*
relation with age; studying these evolutionary trends can inform the relative importance ofdifferent disk processes during key eras of planet formation. External photoevaporation fromthe central O9 star is influencing disk evolution throughout the region: dust masses clearly
decline with decreasing separation from the photoionizing source, and the handful of CO
detections exist at projected separations of >1.5 pc. Collectively, our findings indicate that
giant planet formation is inherently rare and/or well underway by a few Myr of age.

•

Figure 2. Continuum images at 1.33mm of the 37 detected disks in our σOrionis ALMA sample, ordered by decreasing ﬂux density (as reported in Table 1). The last
two images show the stacked non-detections described in Section 5.3. Images are 2″×2″ and the typical beam size of 0 31×0 25 (Section 3) is shown in the ﬁrst
panel.

11mJybeam−1 within 1kms−1 velocity channels for the12
CO, 13CO, and C18O lines, respectively.

•

1.33 mmの連続波と12CO, 13CO, C

Oを使った。

4. Alma Results

4.1. 1.33mm Continuum
Emission
millimeter
continuum
and 6 in 12CO, 3 inNearly all sources in our sample are unresolved, thus wein C18O.
measure 13CO,
continuum ﬂux anddensities none
by ﬁtting point-source•

感度は連続波が2地球質量くらい、COが3木星質
量くらい•

他のALMAサーベイ(Upper Sco, Lupus,

Chameleon)とSMAサーベイ(Taurus)と相補的

and centered on the continuum emission (for detections) or the
expected stellar position (for non-detections). We measure theimage rms using a 4″ to 9″ radius annulus centered on the ﬁtted
or expected source position. Candidate detections are identiﬁed
as those with emission exceeding 3× the rms in multiple
nearby channels within 0–25kms−1 (LSRK frame), which
covers the range of RVs found for σ Orionis members (e.g.,
Jeffries et al. 2006).
For each candidate detection, we create zero-moment maps
by integrating across the velocity range where the emission
exceeds the noise. The integrated ﬂux (F12CO) is then measured
using circular aperture photometry, where the aperture radius
for each source is determined by a curve-of-growth method inwhich successively larger apertures are applied until the ﬂux
stabilizes to within errors. Uncertainties (E12CO) are estimatedby taking the standard deviation of the ﬂuxes measured within
the same sized aperture placed randomly within the ﬁeld ofview but away from the source. We consider sources asdetections when F12CO>4×E12CO. We adopt this highdetection threshold because, as pointed out in Barenfeld et al.
(2016), this procedure selects both the velocity range andaperture size that maximize the signal, thus can produce false
detections at lower signiﬁcance levels. We detect only six
sources in 12CO using this procedure. For these sources, we
also search for 13CO and C18O emission using the same

結果：out of 92 sources, we detect 37 in themodels to the visibility data using the uvmodelﬁt routine in
CASA. The point-source model has three free parameters:
integrated ﬂux density (Fλ), right ascension offset from thephase center (Δα), and declination offset from the phase center
18(Δδ). For the ﬁve resolved sources (1036, 1152, 1153, 1274,
and540), we ﬁt an elliptical Gaussian model instead, which
has three additional free parameters: FWHM along the major
axis (a), aspect ratio of the axes (r), and position angle (PA).
We scale the uncertainties on the ﬁtted parameters by thesquare root of the reduced χ2 value of the ﬁt.
Table 1 presents the 1.33mm continuum ﬂux densities andassociated uncertainties (F1.33 mm), where the uncertainties are
statistical errors and do not include the 10% absolute ﬂux
calibration error (Section 3). We detect only 37 out of the 92
observed sources at >3σ signiﬁcance (Figure 2). For detections, the source locations in Table 1 are the ﬁtted source
Akimasa Kataoka (NAOJ fellow, Nationalcenters
Astronomicalof Japan)
output byObservatory
uvmodelﬁt, while
for non-detections they are

σ-Ori (3-5Myr)の原始惑星系円盤のALMAサーベイ。

Ansdell et al.

•

参考：1太陽≈1000木星、1木星≈300地球

The Astronomical Journal, 153:240 (15pp), 2017 May

Ansdell et al.

•（左上）連続波が受かった37天体を暗い順に並べたもの。一番左の点は受からなかった55天体のスタックしても受からなかったことから出した上限値。このことから、急速な円盤の消失がわかる。

•
Figure 4. Dust masses for the 37 continuum-detected sources in our σOrionis ALMA sample. Dust masses are from Table 1 and error bars include the 10% absolute
ﬂux calibration uncertainty (Section 3). The downward-facing triangle is the typical 3σ upper limit for individual continuum non-detections, while the star shows theconstraint on their average dust mass from the stacked continuum non-detections (“Stack #1” in Section 5.3). Squares indicate TDs identiﬁed by their SEDs
(Section 5.4). Sources outlined in blue are also detected in 12CO (Section 5.2) and sources outlined in red are  0.5 pc from the central OB system (Section 6.1);
sources
closest to theJournal,
central OB
systemhavetheMaylowest dust masses, while sources with the highest dust masses tend to also be detected
The Astronomical
153:240(15pp),
2017
AnsdellinetCO.
al.

10地球質量を超えていたのは11天体→他の天体ではもうcore accretionでガス惑星を作れない

We use this simpliﬁed approach, which assumes a single
for particularly cold or low-mass disks). The CO isotopologue
grain opacity and an isothermal disk temperature, to easelines, speciﬁcally 13CO and C18O, are especially useful forcomparisons with other disk surveys (Sections 6.2 and 6.3).
constraining Mgas as their moderate-to-low optical depths mean•
Moreover, although Andrews et al. (2013) derived athat they trace the bulk gas content rather than the temperature
Tdust = 25 K ´ (L* L :)0.25 relation from two-dimensional
proﬁle of the disk. Ansdell et al. (2016) and Miotello et al. (2017)
continuum radiative transfer models, we use an isothermal
have used these CO isotopologuelines to estimate Mgas fordisk temperature becausemore detailed modeling of resolved
protoplanetary disks in the Lupus clouds that have been surveyed
disks suggests that Tdust is independent of stellar parameters. In
by ALMA.
particular, Tazzari et al. (2017) ﬁt continuum observations of
Unfortunately, because we ﬁnd no C18O detections in
36 resolved disks in Lupus directly in the uv-plane to two-layerσOrionis, we cannot use the same combination of CO
disk models by solving energy balance equations at each disk
isotopologue lines to estimate Mgas in this region. However,radius, ﬁnding no dependence of Tdust as a function of stellar
we can still place rough constraints on Mgas by comparing ourparameters.
measured 12CO and 13CO line luminosities or upper limits toTable 1 presents our Mdust estimates, derived using Equation (1)the WB14 model grid. The uncertainties on Mgas are larger for•with our F1.33 mm measurements (Section 4.1). Figure 4 shows
this line combination because 12CO is optically thick andthe continuum-detected disks in order of increasing Mdust as welltherefore more sensitive to other disk parameters, such as theas the typical 3σ upper limit of ∼2.0M⊕ (F1.33 mm ~ 0.45 mJy).
temperature proﬁle. Nevertheless, Mgas can still be estimated
Only fourdisks have Mdust>30M⊕ (F1.33 mm > 6.7 mJy), thus
using this method because the combination of integrated line
Figure all
6. Left:
map of σ Orionisdiskswith ourcontinuum
detections
circled in wellorange and gas detections circled in blue; the sizes of the circles scale with the ratio
nearlyprotoplanetary
inALMAσOrionis
have
dust masses
ﬂuxes still primarily depends on bulk gas mass rather thanof disk dust mass to stellar mass (Mdust M*) and squares indicate TDs. Non-detections are shown by gray triangles. The central OB system, σ Ori, is marked by the
below
theandminimumof solids
neededto form
planetsother gas
disk
parameterstheexploration
white crossthe dashed mass
white circles
show radial
distancesof 1 andthe2pc.
Notably,inthe sourcesthesewith detectableemission
are among the(seefurthest
fromparameterσOri. Top
oursystem1977).
Moreover,only
11orange points
separationfrom σOri,
where
are continuum
detectionsas
andwellgray triangles
3σ upper limits.
right:solar
disk dust
mass (M(Weidenschillingdust) as a function of projected
described
in WB14as the are
separationof gas masses inatsmaller
separations
fromthecentralOBsystem,andmassive
disks
(M
3
M
)
are
missingwithin∼0.5pc
ofσOri.
Bottom
There
is
a
clear
decline
in
M
dust
dust
e
disks have Mdust>10M⊕ (F1.33 mm > 2.2mJy), thus by
Figure5). holds even after correcting for the Mdust–M* relation
trend stillright: Mdust/M* as a function of projected separation from σOri, illustrating that the declining∼3–5Myr
most
protoplanetarydisks
appearsufﬁcient(Section 6.2). Our
ALMAcontinuum detection
fraction,showntobylack
the thickorange line, also staysFigure
relatively5constantuntil ∼2pc,
after whichmodel
it doubles.
comparesthe WB14
grid to our measured
12
13
dust to form giant planet cores. Note that signiﬁcant amounts of
CO Akimasa
and COKataoka
line luminosities
for sources
detectedin at least Observatory of Japan)(NAOJ fellow,
National
Astronomicalsolids may still exist in objects greater than a few centimeters in（左下）OB型星からの距離との相関。オ
レンジが連続波(ダスト)青がCO(ガス)検出。OB型星近くのガスは晴れている。
2pc以内では円盤検出率30%、以遠では70%。Spitzerではこの傾向はなかったのは、光蒸発は外側だけに効くから？

The Astronomical Journal, 153:240 (15pp), 2017 May

Ansdell et al.

SMA

遠い(385pc)

The Astronomical Journal, 153:240 (15pp), 2017 May•（左上）中心星の質量と円盤のダスト質量との相関。他のサーベイ結果のどの領域を見ても、大きい星ほど円盤は重い。

•

相関は、年老いたUpper Scoだ

古い

Ansdell et al.

けは、ちょっと傾きがきつい→

steeper, this would only further distinguish σOrionis from theyounger regions. Second, the source with the highest stellar
7. Disk dust ismass(Mdust) as a functionof stellar
mass (M*) which
for disk populations in ﬁve star-forming regions with ages spanning the disk dispersal timescale
massFigurein σOrionis
notdetectedin thecontinuum,Myr). Colored circles are (sub-)millimeter continuum detections and gray triangles are 3σ upper limits. For σOrionis, the black triangles indicate 3σ upper
could(∼1–10
potentiallyskew the ﬁt. We could not ﬁnd a valid reason
limits from stacks of the non-detections in three stellar mass bins. For Lupus, the 20 sources with unknown stellar masses that were included in the analysis via an MCfor discarding
this source(e.g.,
no representative
evidence forbinarity),
method (see Ansdell
et al. 2016)
are givenvaluesand identiﬁed by thick gray outlines. For each region, the solid lines show the Bayesian linear
regression
ﬁts to the data,take into accountupperspectral
limits, intrinsic
scatter, and measurement errors on both axes (Kelly 2007). The lower right panel comparesthoughits spectral
typewhichis uncertain(±2.5
types;
the ﬁts in all
the ∼1dex whendifferencein Mdust between
Hernández
et ﬁveal. regions,2014).illustrating
Nevertheless,re-ﬁttingthe the youngest and oldest regions at low stellar masses, and the convergence in Mdust at
high stellar masses.
distribution without this source, we recover thesameAnsdell
et al.
parameters. Third, we have shown in Section 6.1 that externalsurveys. We assume Tdust=20K for all disks and adopt
Table 4photoevaporationis reducing disk dust masses in σOrionis.
Mdust–M* Bayesian Fit Parameters
distances of 140pc for Taurus (Kenyon et al. 2008), 150pc or•This may serve to steepen the Mdust–M* relation, as external
200pc for Lupus (Comerón 2008), 160pc for ChamaeleonIβbδRegionAge (Myr)αbphotoevaporation should be more effective at removing gas and(Luhman 2008), and 145pc for Upper Sco (de Zeeuw
smallet dust
grains around lower-mass stars whose orbiting
Taurus1–2
1.2±0.1
1.7±0.2
0.7±0.1
al. 1999). For Upper Sco, we only include the “full,”a
1–3
1.2±0.2
1.8±0.4
0.9±0.1
Lupus
material
is lessand
gravitationally
this of“evolved,”“transitional”bound.
disks Unfortunately,
from the sample
ChaI
2–3
1.0±0.1
1.8±0.3
0.8±0.1
possibleeffectet al.
has (2016),
not yetbeen represent
tested withexternal
Barenfeld
as thesethe “primordial”
σOrionis
3–5
1.0±0.2
2.0±0.4
0.6±0.1
photoevaporation
models.
Given
thesepotential
issues,amoredisks that do not yet show signs of disk clearing. Only sources
Upper Sco
5–11
0.8±0.2
2.4±0.4
0.7±0.1
sensitive(sub-)millimeter
continuum
surveyofσOrionis,aswith M* 0.1Me were considered in order to exclude brownwell dwarfs,
as additional
and adetailed
modeling,
will beFigurewhile observations
also maintaining
commonstellar masslimit
Notes.8. Dust-mass cumulative distributions for Taurus, Lupus, Chamaeneeded
to conﬁrm
the Stellar
linear regression
ﬁt frompresented
here.
a 0.1Me.
leonI,
UpperetSco
host
stars
M*methodology
among
the surveys.
masses taken
the literature
Fit σOrionis,
taken fromand
Ansdellal. disks(2016),aroundas theyusedthewithsame
Thedescribedaverage indust
mass forgiven
for reference.
Thefordistributions
Furthermore,
we using
note thethatSiessseveralregression
Section
6.2,eachbut region
also anisMC
analysisto account
20 Lupus
were derived
et al. other(2000)linear
evolutionary
tracksandtheir
1σ
conﬁdence
intervals
werecalculated
using
the
Kaplan–Meiermethods
left-censoreddata sets areavailable
thesources with unknown stellar masses.
for allfor
regions
except ChamaeleonI,
whichused theinBaraffeestimator
in the ASURV package (Lavalley et al. 1992) to properly account forbstatistical
literatureand theshouldtestedfor using
this speciﬁc
We use the convention of Kelly (2007), where β and α represent the slope
et al. (2015)
models;stellar bemassesderivedthese two
upper limits using well-established techniques for left-censored data sets.
and intercept, respectively. This differs from that of Pascucci et al. (2016), who
astronomical
grids are problem.
generally consistent, thus any effects should beswitched these symbols.
negligible.
Alternatively, to account for stellar mass differences between
The ﬁtted 6.3.
linearregressionDistributions
parameters for each region are
Dust-mass
distributionsObservatory of Japan)
regions, one
couldKataoka
instead (NAOJ
comparethe M
Akimasafellow,
National
dust/M*Astronomical
given in Table 4 and plotted in Figure 7. Our results further

小さな中心星ほど円盤が早くなくなる？
（左下）円盤ダスト質量の累積関数。平均値ではTaurusで15地球質量、Upper Scoで5地球質量。

Fragmentation and disk formation in high-mass star formation:
The ALMA view of G351.77-0.54 at 0.06ʹʹ resolution
H. Beuther1, A.J. Walsh2, K.G. Johnston3, Th. Henning1, R.
Kuiper4,1, S.N. Longmore5, and C.M. Walmsley6
Context. The fragmentation of high-mass gas clumps and the formation of the accompanying accretion disks lie at the heart ofhigh- mass star formation research.
Aims. We resolve the small-scale structure around the high-mass hot core region G351.77-0.54 to investigate its disk andfragmentation properties.
Methods. Using the Atacama Large Millimeter Array at 690 GHz with baselines exceeding 1.5 km, we study the dense gas, dust
and outflow emission at an unprecedented spatial resolution of 0.06ʹʹ (130 AU @2.2 kpc).
Results. Within the inner few 1000 AU, G351.77 fragments into at least four cores (brightness temperatures between 58 and 197
K). The central structure around the main submm source #1 with a diameter of ∼0.5ʹʹ does not show additional fragmentation.
While the CO(6-5) line wing emission shows an outflow lobe in the north-western direction emanating from source #1, the dense
gas tracer CH3CN shows a velocity gradient perpendicular to the outflow that is indicative of rotational motions. Absorption
profile measurements against the submm source #2 indicate infall rates on the order of 10−4 to 10−3 M⊙ yr−1 which can beconsidered as an upper limit of the mean accretion rates. The position-velocity diagrams are consistent with a central rotating
disk-like structure embedded in an infalling envelope, but they may also be influenced by the outflow. Using the CH3CN(37k −36k) k−ladder with excitation temperatures up to 1300 K, we derive a gas temperature map of source #1 exhibiting temperaturesoften in excess of 1000 K. Brightness temperatures of the submm continuum never exceed 200 K. This discrepancy between gas
temperatures and submm dust brightness temperatures (in the optically thick limit) indicates that the dust may trace the disk midplane whereas the gas could be tracing a hotter gaseous disk surface layer. In addition, we conduct a pixel-by-pixel Toomre
gravitational stability analysis of the central rotating structure. The derived high Q values throughout the structure confirm that
this central region appears stable against gravitational instability.
Conclusions. Resolving for the first time a high-mass hot core at 0.06ʹʹ resolution at submm wavelengths in the dense gas anddust emission allowed us to trace the fragmenting core and the gravitationally stable inner rotating disk-like structure. A
temperature analysis reveals hot gas and comparably colder dust that may be attributed to different disk locations traced by dust
emission and gas lines. The kinematics of the central structure #1 reveal contributions from a rotating disk, an infalling envelope
and potentially also an outflow, whereas the spectral profile toward source #2 can be attributed to infall.

•

ALMAで大質量星形成領域G351.77を観測した•

分解能が0.06秒角(=130AU@2.2kpc, 基線
長1.5km, Band9:690GHz)と大変高いのがウリ•少なくとも４つのコアに分解されているのを確認した。その詳細を説明する（次ページ）

Akimasa Kataoka (NAOJ fellow, National Astronomical Observatory of Japan)•

天体までの距離は議論があり、ここでは2.2kpc(または1kpc)を採用。

Fig. 3. 438 µm continuum image of G351.77-0.54. The left panel shows a larger area encompassing the two central sources as well as two additionalsubmm sources in the south-east and north-west (all marked). The white crosses mark the cm continuum sources by Zapata et al. (2008). The rightpanel then shows a zoom into the center with the triangles presenting additionally the class II CH3 OH masers from Walsh et al. (1998). The−1′′ of 0.06′′ and a scale-bar for an
contour levels start at 4σ H.
andBeuthercontinue
steps
with thecore
1σ levelof 7 mJy beam
. The
beam
et in
al.:2σThe
high-mass
G351.77-0.54
resolvedby synthesized
ALMA at 0.06
resolution
assumed distance of 2.2 kpc are presented as well.

3. Results•

分解能は0.06秒角=123(60)AU•

4つのsubmmソースを確認(#1-#4)•#1と#2はbinaryか？

is clearly extended in the east-west direction, approximately
perpendicular
to the main outflows reported in Leurini et al.
•
3.1. Submm continuum emission at 438 µm(2011a). As will be shown below, #1 is extremely strong in line
Figure 3 presents the 438 µm continuum image of the G351.77- emission whereas #2 is weaker. Without kinematic evidence we
0.54 region. The first impression one gets from this image is cannot provide firm proof of this, but it seems likely that #1 andthat on scales of several 1000 AU we find only a small de- #2 could form a high-mass binary system. It seems likely thatgree of fragmentation. In the remainder of the paper, we will the central part (essentially the central beam of 132(60) AU dialways state the quantitative parameters for both distance esti- ameter) is optically thick in the 438 µm dust continuum emismates of 2.2 and 1.0 kpc (see Introduction), where the 1.0 kpc sion. This is partially based on the high Planck brightness temparameter will be stated in brackets. While within approximately perature of the central beam towards source #1 of approxithe inner 500(250) AU the region fragments into two cores (#1 mately 201 K (see4 section 3.4). For a protostar of luminosity⊙ and assuming a typical dust temperature
and #2), outside of that we have several 1000 AU scales with- of ∼ 2 − 10 × 10 L
−0.4
T ∝ r , temperatures on the order of 200 K atout any further submm core, and then we find the two addi- distribution•
tional cores #3 and #4. Assuming optically thin dust emission 100 AU distance are expected, agreeing reasonably well with theat 200 K (see more details below), our 3σ detection limit of high brightness temperature and moderate optical depth. At such
21 mJy beam−1 corresponds to mass and column density sensi- comparably high optical depth of the submm continuum emis3
tivities of ∼0.004(∼0.0009) M⊙ and ∼ 5 × 1023 cm−2 , respec- sion, we cannot unambiguously distinguish whether source #1
tively. While we can obviously miss some low-column density is a single source or whether multiplicity is hidden by opticalcores, assuming a distribution of column densities it seems un- depth. As we discuss in section 3.4, the optical depth should
likely that we are missing a whole population of lower-mass decrease rapidly with radius. It is worth noting also that the
cores
in µm
this continuum
region. Figure
also CO(6–5)
shows thefive cm continFig.
4. 438
image3 with
red-shifted
emission CH3 CN(37k − 36k ) emission (see section 3.3) appears to follow
sources reported
by Zapatacontours
et al. (2008).
in uumred contours.
The continuum
are in We
4σ alsostepspresent(1σ = an arc or ring centered on the continuum maximum which also
−1 class II maser positions from Walsh et al. (1998)
the beam
CH3 OH
7 mJy. The blue-dotted contours show the corresponding neg- suggests dust obscuration towards the peak of source #1. Ourin Figure
andCO
they
seem toassociatedwith#2.s−1 mass estimates given below based on the optically thin assumpative
features.3,The
emissionwasbeintegratedfrom
25 source
to 42 kmandis shown thein contours
15 positions
to 95% (stepof the peakemis- tion are thus underestimates though likely not by a large factor.
However,
absolutefrom
maser
of 10%)
this dataset
are only
Akimasa Kataoka(NAOJ
fellow,
Nationalsion.
The centerof associate
the phase them
centerunambigupresented incertainwithincoordinates
1′′ , hence, are
we thosecannotThe submm continuum
emission canalso beemployed
to es-Astronomical Observatory of Japan)

上図は連続波で☓は cm波•#1は分解。#2は分解できてない

源、△はメタノールメーザーの位置を表す•#1はアウトフローと直行した方向(東西)に広がっている

下図はCO(6-5)の図で、red側
だけoutflowが見えている。

最小基線長40mより大きな構
造は見えていないことに注意•

CH CN(37k − 36k) というラインに着目。励起温度は588-1300 K。

H. Beuther et al.: The high-mass core G351.77-0.54 re

•

#1に着目

•

1st moment map (左上)

T=Trotはvelocity gradientはoutflowと直行していることを示す
Fig. 8. CH3 CN(37k − 36k ) 1st moment maps of the the k-levels 2 to 10 toward the central sources #1 and #2. The top-left panel shows thecorresponding
integrated
0th
mapmapsof theofCHtransition(integrated
from−13 to#1+10
s−1The). Thewhitepanel
contour
outlines
Fig. 8. CH3 CN(37
1st moment
thethe k-levels
10 toward
the central
sourcesandkm#2.
top-leftshows
the3 CN(37
2 − 3622) tok − 36k ) moment−1−1
the corresponding
3σ level of 0.51
Jy km s 0th. Theothermap
panelsall showintensity-weighted
velocity
maps, andk =contour
2 panels.
While
integrated
momentof theCH3 CN(37(integrated
fromwedges−13 toare+10shown
km s above). The thewhiteoutlines
2 − 362 ) transition−1level isofdone0.51 Jy
km s any. Theclipping,other panels
all showintensity-weighted
velocity
maps,
andThewedgesshown above= 2 µmpanels.
While
the thek =3σ
0 map
without
the first
momentmaps are clipped
at a 4σlevel.
blackarecontourspresentthe
thek 438
continuum−1
the k =in0 4σ
mapstepsis doneanybeamclipping,
the first
moment beams
maps area 4σ
Theandblack
contours (@2.2
presentkpc)
the 438µm continuum
emission(1σwithout= 7 mJy). Thesynthesized
areclippedshownatfor
all level.
panels,a scale-bar
is presented
in the−1emission in
4σ steps(1σ = 7 line
mJy inbeam). The synthesized
are shownall panels,
andfora scale-bar(@2.2 kpc) isdiagrams
presented
in theinbottom-rightpanel.
The diagonal
the top-middle
panel of beams
CH3 CN(37the axis
the position-velocity
shown
3 − 36for3 ) marks
panel. The diagonal line in the top-middle panel of CH3 CN(373 − 363 ) marks the axis for the position-velocity diagrams shown in
Fig.bottom-right9.
Fig. 9.

•

T=200K

PV図はKepler回転+何か(infallかoutflow)を示す•

Toomre のQ値のマップ

T∝r-0.4（右）は、強く温度の過程に依存する。が、Q値は非常に高いので、
spiralがたったりする程9. Position-velocity
diagrams
nineCHCH
36
transitionstoward
toward source
source #1#1 along
along the axis marked
Fig.Fig.
9. Position-velocity
diagrams
forforthethenine
marked inin thethetop-middle
top-middlepanel
panelofof
3 CN(37
3 CN(37k k−−36k )k )transitions
8. Theverticalshows
approximatevelocity
velocityofofrest
restandandthethetwotwofull-line
full-linecurves
curves present
present a Keplerian rotation
Fig.Fig.
8. Theverticallinelineshowsthetheapproximate
rotation curve
curveforforaa10(4.5)
10(4.5)MM⊙⊙star
star
at distancesof 2.2(1)kpc,
respectively.
Thedasheddashedcurves
curvescorrespondcorrespondtotomasses
massesofof5(2.3)5(2.3) andand 20(9)20(9) M⊙⊙ (inner andat distancesof 2.2(1)kpc,
respectively.
Theand outerouter curves,
curves,respectively)
respectively)atat
2.2(1)
distance.
2.2(1)
kpckpcdistance.

度はあるかもしれないが、重力的にかなり安定してそう。

Akimasa Kataoka (NAOJ fellow, National Astronomical Observatory of Japan)

Fig. 12. The color-scales show in all three panels the Toomre Q map(see text for details) of the disk seen toward source #1. While the top
panel uses the rotational temperature derived from CH3 CN (Fig. 10) at
each pixel to determine the sound speed and column density, the middle
panel uses a uniform lower temperature of 200 K for these estimates.
The lower panel uses a dust temperature distribution of T ∝ r−0.4 for the
calculations (see section 4.3). The black contours show in each panel
the 438 µm continuum emission in 4σ steps of 28 mJy beam−1 . Toomre
Q values outside the CH3 CN(372 − 362 ) 3σ level (Fig. 10) are blanked.
The beam is shown at the bottom-right of each panel.

core inand teWview w
104 asourceumn dlower
stellarTrape
Th
lines,
that c1300 K
Furthetify hiuum s
matic
perpenents shcompo
as posTh
with l
we id
the cothe maby outbluefeatur
10−3 Mthe acin #2.
Fi
isotoptempe
tures,
ently
so farable cAssumemissiture abrighta diﬀethe digaseou
cases
An
minin
structusolvab

On the chemistry of the young massive protostellar core NGC 2264 CMM3
Zainab Awad1 · Osama M. Shalabeia1
Abstract We present the first gas-grain astrochemical model of the NGC
2264 CMM3 protostellar core. The chemical evolution of the core is affectedby changing its physical parameters such as the total density and the amountof gas-depletion onto grain surfaces as well as the cosmic ray ionisation rate,ζ. We estimated ζCMM3 = 1.6 × 10−17 s−1 . This value is 1.3 times higher
than the stan- dard CR ionisation rate, ζISM = 1.3 × 10−17 s−1. Species
response differently to changes into the core physical con- ditions, but they
are more sensitive to changes in the deple- tion percentage and CR ionisation
rate than to variations in the core density. Gas-phase models highlighted theimpor- tance of surface reactions as factories of large molecules and showedthat for sulphur bearing species depletion is impor- tant to reproduce
observations.
Comparing the results of the reference model with the most recent millimeter
observations of the NGC 2264 CMM3 core showed that our model is capableof reproduc- ing the observed abundances of most of the species during earlystages (≤ 3 × 104 yrs) of their chemical evolution. Models with variations in
the core density between 1–20 × 106 cm−3 are also in good agreement withobservations dur- ing the early time interval 1×104 < t (yr) < 5×104. In addition, models with higher CR ionisation rates (5–10) × ζISM are often
overestimating the fractional abundances of the species. However, models
with ζCMM3 = 5ζISM may best fit observations at times ∼ 2 × 104 yrs. Our
results suggest that CMM3 is (1–5) × 104 yrs old. Therefore, the core ischem- ically young and it may host a Class 0 object as suggested by previous
studies.

•

大質量クラスター形成領域NGC2264にある原始星
コアCMM3を想定した化学進化計算を実施した.

•比較する観測値は、野辺山45mとASTEによる
4,3,0.8 mm帯のスペクトルサーベイ(Watanabe etal. 2015)の値を用いている

83

Page 4 of 12Z. AwadTable 2 The grid of models we performed, given the main different
physical parameters compared to the reference model (RM). Note thatnone of the other parameters stated in Table 1 have been changed
Model

Density (cm−3 )† Depletion

ζRM

5.4 × 106

full

ζCMM3

2.7 × 107

full

ζCMM3

1.0 × 106

full

ζCMM3

5.4 × 106

partial

ζCMM3

5.4 × 106

none

ζCMM3

5.4 × 106

full

5ζISM

5.4 × 106full

10ζISM

M1
M2
M3
M4
M5
M6

† The depletion is partial means that it is 15% less than the RM follow-

ing Awad et al. (2010). Pure gas-phase chemistry is mimicked in modelM4 with depletion of 0% (none)

3 Results and discussion

RMはreference
We calculated
the fractional abundances (x(Xmodel•of a selected set of species includes H2 CO, l(a• species is considered large when it conta
atoms as defined by Herbst and van Dishoe
sulphur bearing species to represent part ofmolecules
in CMM3 by Watanabe et al. (201•
are illusterated in Figs. 1, 2, and 4 that show
time evolution of the molecular abundances
der different physical conditions; core density
depletion,
and CR ionisation rate, respectively•
estimation of the fractional abundance of CS
the core age are also shown and discussed in

M1,2は密度依存性

M3,4はダストへの吸着依存性
M5,6は電離率依存性を調べている

3.1 Chemical trends
In addition
to theObservatory
RM we rana grid of six other modAkimasa Kataoka (NAOJ fellow, National
Astronomicalof Japan)On the chemistry of the young massive protostellar core NGC 2264 CMM3

Page 5 of 12
On the chemistry of the young massive protostellar core NGC 2264 CMM3

•

83

-3

Page 7 of 12 83(5-10)e6 cm くらいの密度のコアのほうが、密selected molecules are affected by changing the depletion
efficiency. Sulphur bearing species show an increase in theirabundances in both models M3 and M4 while abundancesof large molecules are enhanced in model M3 and become
invisible in model M4. The latter result emphasises the role
• of surface reactions in the formation of large molecules (e.g.
Tielens and Hagen 1982; Hasegawa et al. 1992; Garrod et al.
2008).
The chemical analysis unveiled that the abundance ofgaseous species is enriched through their non-thermal desorption via CR induced photodesorption from grain mantles,
and when the species are frozen back onto grain surfaces
their gaseous abundance are reduced. In model M3 and during early stages of the evolution, t < 104 yrs, the species areproduced via few gas-phase channels in addition to their inFig. 3 The chemical network for CH3 OH and CH3 OCH molecules asrevealed by the chemical analysis of the RM. The network shows the
duced photodesorption from grain surfaces. This leads to anmain formation and destruction pathways of the species. The prefix ‘m’increase of their abundances. Moreover, the species are deindicates mantle species. Different line styles for the arrows represents •stroyed through a series of gas-phase reactions that have a
different situations: short dashed arrows are the pathways that become
very small net loss rates < 10−22 , and hence do not strongly
weak with time while thick solid lines are the dominant routes until latetimes of evolution > 105 years. Dash-dotted arrows represent extra
affect the molecular abundances. We, also, have noticed thatformation routes of CH3 OH that are minor in the RM, during early
the formation and destruction channels as well as their rates
times < 2 × 104 yrs, but dominate the chemistry in less dense cores• in models M3 and the RM become closer at later stages of(model M2) at the same time interval; see the text in Sect. 3.2.1
the evolution (t > 5 × 105 yrs) which causes the conversionof the results.
ution ofline, and low-density core model (M2, nH = 1.0 × 106 cm−3 ) is theof gaseous CH3 OH in less-dense cores for longer times beIn order to clarify the above discussion on depletion, wet differdash-dot line. Observed values, as taken from Watanabe et al. (2015),fore it declines to converge to the values obtained in the RM
discuss the case of CH3 OCH. The production of CH3 OCH
are represented by solid gray line
) is thearound 105 yrs.
ISMgrainin the RM is via its induced photodesorption fromdashedCH3 OH is mainly destroyed via Si+ which starts to desurfaces and the gas-phase reaction ‘CH + CH3 OH’. The
crease in abundance at times ∼ 2 × 105 yrs. This decreaseabundance of the molecule is decreased by its depletion
allows CH3 OH to stay longer in the gas causing its abunonto grains to either form icy CH3 OCH or hydrogenate into
5 yrs.
dancetore-increase
until
itsaturates
at
times
>
5
×
10
es and(1) the core density, (2) depletion onto grains, and (3) the• CH3 OCH2 , see Fig. 3. In model M3, from the chemical analAfter that time, the chemical analysis showed that the main
rate of CR ionisation.
given
ysis we found that the formation of CH3 OCH is governed bypathways as well as the net rates of formation and destruc4 yrs).
the gas-phase reaction ‘CH + CH3 OH’ because, at any time,tion of CH3 OH in all the three models become close. This
the abundance of CH3 OH is higher in model M3 than in the
3.2.1 The core density
ompaleads to the observed conversion of the abundances at late
RM by 2 orders of magnitudes. The reason for this is thattimes.
The
discussed
scenarioformethanolisapplicable,inhe obthe formation of methanol in model M3 occurs by the reacgeneral terms, for the other species in question.
By
increasing
thecore
densityanincreaseintheabundances+t timesFig. 1 The effect of varying the core density on the time evolution ofline, and low-density core model (M2, nH = 1.0 × 106 cm−3 ) istions
the ‘CH3 OCH + H+
’
and
‘NH
+
CH
OH
3
3
3
2 ’, in which theofthespeciesis
expected,
but
this
is
not
what
we
obtained,the calculated molecular abundances. Different curves represent differdash-dot
line.depletion
Observedefficiency
values, as taken from Watanabe et al. (2015),
latter does not weaken with time as the case in the RM, andstage,
3.2.2
The
6
−3
see Fig. 1. We found that changes in the core
are represented by solid gray linecm )density
is the haveits destruction by Si+ ions is minor. This situation keeps theegin ent
to models: the reference model (RM, nH = 5.47 × 10
−3 ) is the dashed= 2.7 (model× 10 cmM1,
solid line,
denseeffects
core model(M1, nHcoresminor
in denser
dash line) by re-The role of grain surfaces as factories of various interstellar abundance of methanol higher in model M3 than in the RM
ructiveducing all molecular abundances while in less-dense coresmolecules is now well-established (e.g. Tielens 2013; Lin- and hence increases the yield of CH3 OCH.
hed atWhen depletion
et al. 2015,
and references
therein).
Therefore,
we ex- Observatory
Akimasa
Kataoka(NAOJ fellow,
National
Astronomicalof Japan)become zero, in model M4, large(model M2, dash-dot line) the abundances are enhanced.nartz

度の小さいコアより観測をよく説明する
depletion(ダストへの吸着)があるものや高密度モデルは大きな分子とOCSの観測値を説明する。depletionが少ない/ないモデルはSO, SO2,
CSを説明する→これらは気相反応期限っぽい
気相反応のみのモデルでは観測値を説明しない
宇宙線による電離率はISMの1.3倍をリファレン
スモデルにしたが、更にその5倍でも観測を説明する→電離率は(1.3-5)ζ

か。

リファレンスモデルは1-2e4年でよく観測と一致。また、高密度や高電離度モデルでは2e4-5e4
年だとよく一致。これはだいたいClass0天体の年齢と一致。なので、たぶんCMM3は年齢が(15)e4年くらいのClass 0天体だろう。
−13

arXiv:1703.08609v1 [astro-ph.

range, with the detection of 94 new rotational transitions. Additionally, 46 line positions taken from the literature have been accurately
remeasured.
Results. The newly-measured lines have significantly enlarged the available data sets for HOCO+ and DOCO+ , thus enabling thedetermination of highly accurate rotational and centrifugal distortion parameters. Our analysis showed that all HOCO+ lines with
Ka ≥ 3 are perturbed by a ro-vibrational interaction that couples the ground state with the v5 = 1 vibrationally-excited state. This
resonance has been explicitly treated in the analysis in order to obtain molecular constants with clear physical meaning.
Conclusions. The improved sets of spectroscopic parameters provide enhanced lists of very accurate, sub-millimetre rest-frequencies
of HOCO+ and DOCO+ for astrophysical applications. These new data challenges a recent tentative identification of DOCO+ toward
a pre-stellar core.

Accurate sub-millimetre rest-frequencies for HOCO+ and DOCO+ ions
L.Bizzocchi1,V.Lattanzi1,J.Laas1,S.Spezzano1,B.M.Giuli
a n o 1 , D . P r u d e n z a n o 1 , C . E n d r e s 1 , O . S i p i l ä 1 , a n d P. Caselli1
Key words. Molecular data – Methods: laboratory: molecular – Techniques:
spectroscopic – Radio lines: ISM•

Context. HOCO+ is a polar molecule that represents a useful proxy for its parent molecule
CO2, which is not directly observable in the cold interstellar medium. This cation has been
detected towards several lines of sight, including massive star forming regions, protostars,
1. Introduction
and cold cores.
Despite the obvious astrochemical relevance, protonated CO2 and its
deuterated variant, DOCO+, still lack an accurate spectroscopic characterisation.

were
as proxies
for important
Aims. The aim Protonatedof this work is ionsto extend
thefirst
studysuggestedof the ground-statepure rotational+ well intomolecules
bytheHerbst
et al. region.
(1977), shortly after the
spectra of interstellar
HOCO+ and DOCO
sub-millimetre
Methods. Ground-statetransitions
have been polyatomic
recorded in the species
laboratoryin
using
a frequencyfirst detectionof charged
space(HCO+ ,
modulationBuhl
absorption
spectrometerequippeda free-spaceglow-discharge
cell.
The
& Snyder1970;
N2 H+with, Turner1974).
These firstpioneerions were produced in a low-density, magnetically-confined plasma generated in a suitable
ing studies demonstrated that ion–molecule reactions must occur
gas mixture. The ground-state spectra of HOCO+ and DOCO+ have been investigated inin the interstellar medium (ISM), and are capable of generating
the 213–967 GHz frequency range, with the detection of 94 new rotational transitions.
ionic
forms
of non-polar
as N2 , Cremeasured.
2 , CO2 , andAdditionally,
46 linepositionstaken from themolecules,
literature havesuch
been accurately
HCCH. These species are likely to be present to a large extent inResults. The newly-measured lines have significantly enlarged the available data sets forthe dense gas, but they escape radio-telescope detection owing
HOCO+ and DOCO+, thus enabling the determination of highly accurate rotational andthe lackof rotational
spectra.
centrifugalto
distortion
parameters.
Our analysis
showed that all HOCO+ lines with Ka ≥ 3
are perturbed byCarbon
a ro-vibrational
interaction
the groundwith the1
dioxide(CO2 )thatiscouples
widespread
instatespace.
It vis5 =abunvibrationally-excited
This resonancehas been explicitly
treatedthe analysis ininterdant in state.
planetary
atmospheres,comets,andin especiallyorder to obtain molecular constants with clear physical meaning. Conclusions. The
stellar
ices,
where
it
has
been
extensively
detectedinthese
improved sets of spectroscopic parameters provide enhanced lists of very accurate, sublatterenvironments
by
ISO
and
Spitzertelescopes
towards++
millimetre rest-frequencies of HOCO and DOCO for astrophysical applications. These
several lines of sight (e.g., Whittet et al. 1998; Bergin et al.
new data challenges a recent tentative identification of DOCO+ toward a pre-stellar core.

CO2を量を推定するのに有用なHOCO+についてよく調べたい(Ioppolo et al. 2009; Mennella et al. 2006; Pontoppidan et al.
2003; Watanabe & Kouchi 2002). Nonetheless, speculation on•theHOCO+とDOCO+の基底状態をよく調べた
possible contribution of a gas-phase formation route remains(e.g., Sakai et al. 2008), and the diﬃculties involved in the direct
observation of CO2 in the infrared hinder the clarification of this•matter.
たとえばL1544では以下のネットワークから、+Protonatedcarbondioxide(HOCO)
providesa
useful,in[HOCO+]/[HCO+]を観測することで、[CO2]/[CO]を推
direct way to trace gaseous CO2 in the ISM. Vastel et al. (2016)
constrained
the CO2 abundance in the L1544 pre-stellar coreus- 2016)
測するような試みがなされている(Vastel
et al.
ing an extensive chemical model that considered the following
HOCO+ main formation channels:
CO2 + H+3 → HOCO+ + H2 ,CO + H+3 → HCO+ + H2 ,followed by

HCO+ + OH → HOCO+ + H .

(1a)
(1b)
(2)

2005). In the solid phase, the CO2 :H2 O ratio has been ob- At the steady state, they derived an indirect estimate of theserved to vary in the range of 0.15–0.5 in molecular clouds [CO2 ]/[CO] ratio from [HOCO+ ]/[HCO+ ]. The assumptions inand protostars (e.g., Boogert et al. 2015;
Whittet et al. 2009; volved in this approach hold for the external layers of dense
Akimasa Kataoka (NAOJ fellow, National Astronomical Observatory of Japan)

+

L. Bizzocchi et al.: HOCO+ and DOCO+ rest frequencies+

L. Bizzocchi et al.: HOCO and DOCO rest frequencies
HOCO+ , JKa ,Kc = 300,30 − 290,29

DOCO+ , JKa ,Kc = 161,15 − 151,14

parameterunit
5. Chemical
model
ground state
A
B
C
DJ
DJK
DK
d1
d2
HJ
HJK
HKJ
′′
HK

detection signal / a.u.

detection signal / a.u.

B
C

Table 2. Spectroscopic parameters determined for HOCO+ .a
fit I

ab initiob

fit IIDMHz
789 947.786(23)789 944.610(21)784 759.5++
Predictions for MHz
the HOCO
andDOCO
abundancesin
10 773.73262(26)
10 773.68964(34)10 787.1
MHz
10 609.43083(28)10 610.3413(23)10 623.7
L1544 can be derived
fromthechemical
model
developedbykHz
3.49883(20)3.54176(24)3.433
0.93613(12) a static
0.6454(38)0.852Sipilä et al. (2015), MHz
whichconsiders
physicalstructureMHz
1 123.57
1 123.57
728.341kHz−0.05153(12)(as described−0.04555(20)−0.04148
and an evolving gas-grainchemistryin Sipilä
et al.
kHz−0.01635(10)0.01744(15)−0.00914
2016). Table 4 shows
column densitiesofHz the model−0.002computed−0.002−0.0020.402(83)14.5(11)0.885+
HOCO+ and DOCOHz
fordiﬀerent
cloud
evolutionary
ages
avkHz
3.597(28)19.756(52)−0.089415
+
MHz
3.418
3.418
3.418
eraged over a 30 beam.
Vastelet al. (2016)
found a HOCO
v =1
column densityof about
2 × 1011 cm−2767. Comparison
with the
A
MHz
920.6
762 735.5
BMHz
10 789.9
10 803.2
predicted column
densities,
shown
in
Table
4
suggeststhat theCMHz
10 641.6
10 655.0
D
kHz
3.433
3.433
best agreement
is
found
at
late
evolutionary
times.
However,
D
MHz
0.852
0.852
5
MHz
728.341
since abovedD3 × 10 kHz
yr, the model predicts1 300
too
much
CO
freeze−0.04148−0.04148d−0.00914−0.00914 weout comparedto thekHzobservations of Caselli
et al. (1999),
H
Hz−0.002−0.0025consider 3 ×
as the best agreement time.
At this0.885
stage,
H 10 yr Hz
0.885
kHz
−0.089415 +the DOCO+HH/HOCO+MHz
is about 10% and the−0.089415
predicted
DOCO
,
3.418
3.418
J = 50,5 −∆E40,4 mainis 6 mK
cm beam brightness temperature
535.6
η
MHz
3 816.6(47)(assuming an
excitation
temperatureof
8.5
K, as deduced byηMHz
0.180(11)+of lines for their observed
74
93 line).
Vastel et al.no.(2016)
HOCO
c

D
F

c

d

d

d

d

5

F
G
G
H
I

e

e

e

640995640997640999641001641003641005323119323121Frequency / MHz

323123323125Frequency / MHz

d

J

d

JK

f

K

+

Fig. 1. (left panel) Recording of JKa ,Kc = 300,30 − 290,29 transition of HOCO . Total integration time is 254 s with 3 ms time constant.
(right panel) Recording of JKa ,Kc = 161,15 − 151,14 transition of DOCO+ . Total integration time is 218 s with 3 ms time constant.
(both panels) The red traces plot the computed spectral profile obtained with the proFFit code using a modulated Voigt profile (seetext).

d

1

d

2

d

J

d

K

について、以前よりよくもとまった

The search for new rotational transitions of HOCO+ was
guided by the spectroscopic parameters previously reported byBogey et al. (1988), thus their assignment was accomplished in• way. However, at frequencies above 500 GHz,a straightforward
increasingly larger discrepancies (∼ 500 kHz) were found between observed and predicted line positions. Forty-three newrotational lines were recorded, reaching a maximum quantum
number J values of 29 and a frequency as high as 967 GHz.
These data included 12 b-type transitions belonging to theb
P+1,−1 , b Q+1,−1 , and b R+1,+1 branches. In addition, 18 lines previously reported by Bogey et al. (1986, 1988) were remeasured

−1

d

N

O
O
P

The analysis clearly showed that the ground-state spectrumof HOCO+ is perturbed. An anomalous slow convergence of the
σw
0.82
1.03
rotational Hamiltonian had already been noted by Bogey et al.
(1986), such that high-order terms were used to achieve a sat- a See text for details on fit I and fit II. Values in parentheses represent 1σ uncertainties, expressed in units of the last quoted diP
isfactory fit of the measured frequencies. Similar perturbations b Theoretical values computed by Fortenberry et al. (2012).
+
+
P
Table
4.
Chemicalmodel
predictions
for
HOCO
and
DOCO
c
have been observed in many quasi-linear molecules, for which d Fixed at the value determined by Amano & Tanaka (1985b).
Fixed
the ab initio value.
in atL1544.
HNCO (isoelectronic with HOCO+ ) serves as a case study e Derived
from the ground state value plus ab initio α (Fortenberry et al. 2012).
Pf
Adjusted by step-by-step procedure.
(Niedenhoﬀ et al. 1995b). These anomalies reflect the breakR
down of the Watson-type asymmetric rotor Hamiltonian because−2−2time / yr NHOCO+ / cm
NDOCO+ / cm
NDOCO+ /NHOCO+
of the accidental ∆Ka = ±1 degeneracy occurring betweenSµb = 1.8 D. Th
ues determined for the
4 H JK and HK J sextic 9constants should be formed
8 using µa = 2.7 D andground-state rotational levels and those of a low-lying, totally- considered
+
5only
× 10
9.2 × 10 since they include
3.9 × contains
10 57 lines for 0.042
HOCO and 72 lines for S
as eﬀective approximations,symmetric excited state.
5contributions not explicitly
10 treated by this are 9selected in the frequency range 60 GHz ≤spurious resonance
1
×
10
2.4
×
10
2.1 × 10
0.084
S
+
and
applying
the
cut
oﬀs
E
/k
<
100
K.
Th
analysis.
Finally,
it
is
to
be
noted
that
fit
I
and
fit
II,
when
limu
We carried out the analysis of the HOCO spectrum follow3 2×lines,
105yield spectral
6.3
× 1010that are coincident
6.7 × these
109 rest-frequencies0.106
is at least 4 × 10−8 , whicT
ited to Ka ≤predictions
ing two diﬀerent approaches. The first, simpler analysis was per- within the 1σ computed
5 uncertainties.
10
9 km s−1 (or better) in unit of radial velociT
to
0.011
5
×
10
7.7
×
10
6.6
×
10
0.086
formed by applying a cut-oﬀ at Ka = 2. This excluded from the
New HOCO+ and DOCO+ line catalogues, based on the
10
9T
1 ×constants
106 of Table8.8
0.067
2 (fit×I,10
only) and Table 3,5.9
have× 10
least-squares fit all the lines aﬀected by the resonance, and al- spectroscopicU
computed and are provided as supplementary data availlowed us to consider the HOCO+ ground state as isolated. The beenable at the CDS. These data listings include: the 1σ uncertaintiesv
Our analysis suggests that the line observed
Ka = 0, 1, 2 lines were fitted using a single state Hamiltonian, (calculated taken into account the correlations between spectroand these results are reported in the first column of Table 2 scopic constants), the upper-state energies, the line strength fac- Vastel et al. (2016) and tentatively assigned to DOV
50,5 − 40,4 , cannot actually be attributed to thisW2(referred to as fit I hereafter). This analysis provides a com- tors S i j µg (g = a, b), and the Einstein A coeﬃcients for sponta- served
feature, once red-shifted by the 7.2 km s−1
neousemission,
Akimasa
Kataokaparameters,(NAOJ fellow,
Nationalfour
Astronomical
Observatoryof Japan)pact setof rotational
includingquartic andtwo
has a rest frequency of 100 359.81 MHz, whereaW

実験室実験から、HOCO+とDOCO+のスペクトル3.1. HOCO+

d

KJ

ab
12
abJ
12

N
N

d

JK

v

estimated to be in the range of 20–50 kHz, depending on the line
width, the achieved signal-to-noise ratio (S/N), and the baseline.
With a magnetic field applied during the plasma discharge, theions are produced primarily in the negative column, which is anearly field-free region. Therefore we assume the Doppler shift
caused by the drift velocity of the absorbing species is negligible(see e.g., Tinti•et al. 2007).

K
M

L1544のモデル計算もした（1段落）COSMIC-RAY INDUCED DESTRUCTION OF CO IN STAR-FORMING GALAXIESTHOMAS G. BISBAS1,2 , EWINE F. VAN DISHOECK1,3 , PADELIS P. PAPADOPOULOS4,5,6,7 , LÁ SZLÓ SZŰ CS1 , SHMUELBIALY8 AND ZHI-YU ZHANG7,9•(省略します)•モチベーション:H2ガスの量を観測するのに、COが使われ

ABSTRACT
We explore the effects of the expected higher cosmic ray (CR) ionization rates ζCR
on the abundances of carbon monoxide (CO), atomic carbon (C), and ionized carbon(C+) in the H2 clouds of star-forming galaxies. The study of Bisbas et al. (2015) is
expanded by: a) using realistic inhomogeneous Giant Molecular Cloud (GMC)
structures, b) a detailed chemical analysis behind the CR-induced destruction of CO,
and c) exploring the ther- mal state of CR-irradiated molecular gas. CRs permeating
the interstellar medium with ζCR 10×(Galactic) are found to significantly reduce the[CO]/[H2] abundance ratios throughout the mass of a GMC. CO rotational line
imaging will then show much clumpier structures than the actual ones. For ζCR
100×(Galactic) this bias becomes severe, limiting the utility of CO lines forrecovering structural and dynamical characteristics of H2-rich galaxies throughout
the Universe, including many of the so-called Main Sequence (MS) galaxies where
the bulk of cosmic star formation occurs. Both C+ and C abundances increase withrising ζCR, with C remaining the most abundant of the two throughout H2 clouds,
when ζCR ∼ (1 − 100)×(Galactic). C+ starts to dominate for ζCR 103×(Galactic).
The thermal state of the gas in the inner and denser regions of GMCs is invariant withTgas ∼ 10 K for ζCR ∼ (1 − 10)×(Galactic). For ζCR ∼ 103 ×(Galactic) this is nolonger the case and Tgas ∼ 30 − 50 K are reached. Finally we identify OH as the key
species whose Tgas −sensitive abundance could mitigate the destruction of CO athigh temperatures.

ている。
•

ALMAでは更に高赤方偏移の天体まで見るだろう。

•

その時、宇宙線の量も多くなるはずである。→COが壊れる。

•

その効果を定量的に見積もるため、分子雲コアの3次元化
学計算を行った。特に、CRの電離率を1e-17 - 1e-14 per
sec で降って結果を見た。

•

CO/H2 比は1e-14 (=3×Galactic)でもガクッと減った。むしろ、Atomic Carbonがガス質量のよいトレーサーである

Akimasa Kataoka (NAOJ fellow, National Astronomical Observatory of Japan)