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

•（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。

•

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

•

（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？

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

•（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
けは、ちょっと傾きがきつい→
小さな中心星ほど円盤が早くな
くなる？

surveys. We assume Tdust=20K for all disks and adopt
distances of 140pc for Taurus (Kenyon et al. 2008), 150pc or
200pc for Lupus (Comerón 2008), 160pc for ChamaeleonI(Luhman 2008), and 145pc for Upper Sco (de Zeeuw
et al. 1999). For Upper Sco, we only include the “full,”“evolved,” and “transitional” disks from the sample ofBarenfeld et al. (2016), as these represent the “primordial”
disks that do not yet show signs of disk clearing. Only sources
with M* 0.1Me were considered in order to exclude browndwarfs, while also maintaining a common stellar mass limit
among the surveys. Stellar masses taken from the literature
were derived using the Siess et al. (2000) evolutionary tracks
for all regions except ChamaeleonI, which used the Baraffe
et al. (2015) models; the stellar masses derived using these two
grids are generally consistent, thus any effects should be
negligible.
The ﬁtted linear regression parameters for each region are
given in Table 4 and plotted in Figure 7. Our results further•（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。

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

•

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？

•#1は分解。#2は分解できてない•#1はアウトフローと直行した方向(東西)に広がっている•

CH3CN(37k − 36k) というラインに着目。励起温度は588-1300 K。

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

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

•

#1に着目

•

1st moment map (左上)はvelocity gradientはoutflowと直行している
ことを示す

•

PV図はKepler回転+何か(infallかoutflow)を示す•

Toomre のQ値のマップ（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
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perpenents shcompo
as posTh
with l
we id
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Fi
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An
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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.

•

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

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