星形成ゼミ
SFN #304 (19-28本目)
百瀬宗武（茨城大・理）ハイライトは以下の３本• #20: Unveiling the physical conditions of the youngest disks: A warm embedded disk in
L1527 (vanʼt Hoﬀ et al. A&A in press)• #23: ALMA OBSERVATIONS OF POLARIZATION FROM DUST SCATTERING IN THE IM
LUP PROTOPLANETARY DISK (Hull et al., ApJ in press)• #26: Distortion of Magnetic Fields in a Starless Core III: Polarization‒Extinction
Relationship in FeSt 1-457 (Kandori et al. ApJ in press)

20. Unveiling the physical conditions of the youngest disks:
A warm embedded disk in L1527【結果】

積分強度図とポンチ絵：ほぼedge onの円盤を見ている
CO系の赤と黒は redshifted, blueshifted の成分分布

van’t Hoﬀ et al. A&A in press
http://ads.nao.ac.jp/abs/2018arXiv180304515V【内容】原始星(Class 0天体)L1527に向けた13CO, C18O(J=2-1),
N2D+(J=3-2)をアーカイブデータを解析して，円盤の温度構造を議論。速度情報を用いて円盤成分とエンベロープ成分を切り分け。結果としては，COが凍結するほど低温な領域は円盤中には存在しない。Class 2と違って，活発な降着が円盤温度をあげている？

【観測データのサマリー】
Cycle 2のアーカイブデータを利用，Band 6。
COの同位体：基本的なガストレーサー，ただし，COのスノーラインや
depletionが最近話題。
+のD化物
N2D+：COが凍結する領域でEnhanceするNH
2
A&A proofs: manuscript no. Paper_L1527_arxiv
Table 1. Overview of the molecular line observations toward L1527.

Molecule Transition Frequency
(GHz)
13

CO
C18 O
N2 D+

J=2
J=2
J=3

Aul a
(s 1 )

1 220.39868 6.038⇥10 7
1 219.56035 6.011⇥10 7
2 231.32183 7.138⇥10 4

gup b
5
5
7

Eup c
(K)

vd
rms
rms
1
1
(km s ) (mJy beam(K
1channel ) channel 1 )15.87 0.32⇥0.18 (PA = 39.0 )
0.08e
10.2f
4.4f
15.81 0.32⇥0.19 (PA = 38.7 )
0.04e
8.7f
3.6f
22.20 0.93⇥0.71 (PA = 10.7 )
0.08
10.4
0.4

Notes. (a) Einstein-A coefficient. (b) Upper level degeneracy. (c) Upper level energy.

Beam
(00 )

CO系の積分強度のピークは連続波の位置にないが，連続波が光学的に厚いため(r ≲ 15auの範囲で)

lso

serterase.
18
O
the
with
ane
is,
ons.
met-

no
ant
3
for
m 2
and
ecold
t.

ore
abing
file
figrivCOsoomptithetirebeisk,
hat
adii
the
roure
Asthe
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ion

チャンネルマップ（前ページ）
における強度比
中心面付近では 2→ 13COは光学的に厚い
下図はRADEXを使った
CO柱密度と13CO/C18O(2-1)の関係，C18Oも中心面付近はoptically thickFig. 2. 13 CO (J = 2 1)/C18 O (J = 2 1) line intensity ratio in six
representative velocity channels. The velocity relative to the systemic
velocity of vlsr = 5.9 km s 1 is listed in the top right corner of each panel（塗りつぶしがτ > 1）
and the channel width is 0.16 km s 1 . The line ratio is only calculated inpixels with at least a 3 detection for both 13 CO and C18 O. The dottedand solid contours mark the respective 3 levels. The plotted spatial
scale is such as to highlight the inner ⇠ 100 (140 AU), that is, the disk 臨界密度に比べガス密度は十分高い
Merel L.R.peak
van ’t position
Ho↵ et al.:isUnveilingthewithphysicaland inner envelope. The continuum
indicateda conditions of the youngest disks:cross and the synthesized beam is shown in the lower left corner of theので，13COとC18Oの輝度温度を
lower left panel.

温度分布の導出に使える

The 13 CO (J = 2 1) / C18 O (J = 2 1) line ratio is calculated
for each pixel with at least a 3 detection (see Table 1) for both13
CO and C18 O, and six representative channels are presented inFig. 2. The 13 CO/C18 O ratio is below 3.0 in all regions, and below 2.0 in the pixels with the highest signal-to-noise (along the
midplane). The 3 uncertainty on the 13 CO/C18 O ratio is calculated for each pixel by propagating the rms in the 0.16 km s 1channels. Within this uncertainty, the ratio remains below 3.5,
except in a few pixels close to the 3 contour of the C18 O emission. Theoretically, the ratio cannot be smaller than 1.0, unless 13 CO emission is significantly more resolved out relative
to C18 O. The lowest values reached are ⇠0.5, but the 3 uncertainty in these pixels is 0.5. The values below 1.0 are thus not
significant.
Since
the
onlinetheintensity
gas temperature,
Fig.
A.1. 13COoptical(J = 2 depth1)/C18 Odepends(J = 2 1)
ratio calculatedthewith RADEX for a range of temperatures (di↵erent colors) and CO
12
13
16
18
quantity
we set out
to ratios
determine,
also
column densities.
Isotopeof [ C]/[theC]line= 77 ratio
and [ isO]/[O]dependent= 560 are adopted. In the left panel (right panel), open circles indicate that the1318
CO temperature.
(C O) emission Therefore,
is optically thin
⌧ <to1),determine
whereas filledwhether
circles arethe
used when the 13 CO (C18 O) emission is optically thick (i.e., ⌧ > 1).
on thein(i.e.,order
Values below
theeitherdotted or
lineboth
are observed
toward the L1527
emissionfrom
CO isotopologues
is midplane.
optically thick,we calculate line ratios for a range of temperatures (20–100
18

diskmanuscript
radiusno.dPof the radial temperature profile by determining the maximumA&A proofs:
temperature in the highest velocity channels and extrapolating small cold ou
snow
radius of 74
this assuming a radial power law dependence,
diskA【ナイーブな温度分布の導出】
simi
the entire dis
to ⇠no
T = T obs (R/Robs ) q ,(1) that CO ising
frac
18
For C prevO
qをいろいろ変化させてフィッティ
where T obs and Robs are the temperature and radius derived from emission are
4.3.
the observations.
ング Since the highest velocities are reached closest temperature
As m
to the star, these channels are the least likely to be contaminated (Fig. 4). The
callyby envelope emission. Often a power law exponent of q = 0.5 is bly lower disk
beemi○が輝度温度の観測値，
Fi
used,
appropriate for a black body flared disk with scale height plane (seethic
spon
1.25
h(r)
=
h
(r/R
)
(Kenyon& Hartmann 1987), but q can range the midplanee↵e?
実線/破線は 0redshifted/blueshiftedvest
between 0.33 and 0.75 (see e.g., Adams & Shu 1986; Kenyon the midplane
weのみのfitting
tion
et al. 1993;
Chiang & Goldreich 1997). However, full radiative that the tempthetransfer modeling shows that the temperature can not be charac- fected by ble
rang
of s
terizedbyasingle
power
law,
and
that
the
midplanetemperaturetemperaturesdept
r=74 auはAso + (2017)による，
profile can be flatter when the disk is embedded in an envelope In Sect. 5.1,⌧=
r=125auはT13による円盤半径(e.g., Whitney et al. 2003). We therefore compared a power law perature, andwith q = 0.5 to the midplane temperature profile from the T13 profile, unde
whe

of th
ener
and

T13はTobin+ (2013)で求められていた温度分布
それよりは低いが，あまり分解能の高くない観測の輝度温度で見積
もると運動温度を低く見積もる。
それも考慮すると，円盤の大部分
でCO freeze out (< 20K) の条件を満たしていないように見える（モデリングの項参照）

Fig. 4. Radial midplane temperature profiles for the L1527 disk andinner envelope. The top panel shows a comparison between the temperature profile from the T13 model and power laws with di↵erent exponents (q). In the middle and bottom panels, power law profiles arederived from 13 CO (orange lines) and C18 O observations (blue lines)
assuming q = 0.5 and q = 0.35, respectively. Solid lines are derived
from a redshifted channel, dashed lines from a blueshifted channel. The
measured values are indicated by filled circles. The orange shaded areashows the 1 uncertainty on the solid orange line (13 CO from redshiftedchannel). The gray shaded area marks the temperature range for which
CO is expected to be frozen out (i.e., below 20 K). The vertical dashdotted line marks a disk radius of 74 AU as determined by A17, and thedotted line a disk radius of 125 AU derived by T13.

turefittin
obse
ima
has
Kep
tatin
Mor
for t
deri
a sli
valu
sion
will
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a ga
danc
We
whi

line
mid

sion (⇠30 AU in this case) is smaller than the radius at which
the emission originate (⇠125 AU). At lower velocities only envelope emission will be observed (see Table 2), until at velocities
close to the systemic velocity the disk can be seen at the source
Merel L.R. van ’t Ho↵ et al.: Unveiling the physical conditions of the youngest disks:
position with envelope emission on both sides (see Fig. 6, thirdcolumn). The situation is similar for C18 O (see Fig. 6, bottom
warm
intermediate
coldrow), except that the emission becomes optically thick in the far
side of the disk (see Fig. 6, column one, two and five), instead
envelope（左）考える状況：3つのパターン
of the near side. In addition, the emission remains optically thinin most of the envelope at redshifted velocities (for the adopted
density structure; see Fig. 6, fourth column). The disk is thus not
hidden from our view.
（左下）チャンネルマップの再現例
For clarity, we have neglected the dust opacity in Fig. 6.
Adopting the dust opacity from T13, that is,  = 3.1 cm2 g 1 at
25分観測（上パネル）だとw, iが同じくらい
220 GHz, the dust would only become optically thick in the inner ⇠10 AU. If the velocity profile in the envelope is free fall
あっていて区別つかない。5時間観測（下パ
天球面と直交する面内で(Fig. D.1), the highest
envelope velocities are found along the
disk
disk
disk
line of sight toward the protostellar position, instead of o↵set toMerel’t Ho↵
et al.: Unveiling the physical conditions of the youngest disks:
ネル）だと区別が微妙につくようになる。
the northどの場所がどの視線速度に効くか(for redshifted emission)or toL.R.
the vansouth(for blueshiftedCOemission), as in the infalling rotating case. This changes the lothe envelope is blocked from our view except at velocities where
cation of the ⌧ = 1 contour at velocities that occur in both the
midplaneCOfrozenthe envelope extends further south than thedisk.
In thisoutcase, atdisk and envelope (that is, columns two and four in Fig. 6 will（右下）輝度温度（観測）から求めた温度はlook di↵erent; see Fig. D.2 and Table 2). For blueshifted velocismall angular o↵sets from the source center the emission origiFig.
Threewhiledi↵erentmodelsatfor
the CO
distribution
in theL1527 disk.
ties, the disk will be visible south of the protostar, while envelopenates in
the7.disk,
emissionlargero↵setstraces theenve仮定した温度より常に低く出る
Leftpanel:
a
warm
disk
with
no
CO
freezeout.
Middlepanel:aslightly
emission will be present to the north. For redshifted velocities,
lope (seeFig. 6, second column). In the redshifted channels thecolder disk where CO is frozen out in the outer most region (R > 71 AU)
Merel L.R. van ’t Ho↵ et al.: Unveiling the physical con
emission surrounding the protostar originates from within theemission
optically
thickin theenvelope
betweenus andnearbecomesthe midplane.
Rightpanel:a colddisk where
gaseous
CO is only
envelope, while the emission furthest north is coming from thethe disk.
At thehighestoccuring
theandenvelope,
disk layers.
present
in theinnervelocities
most region(R < 23inAU)
in the surface
power lawdisk.
emissionbe present
veryphaseclose
protostarFig.
6,
CO willis present
in the gas
in tothetheinnerenvelope(seein allmodels.
warm disk
In summary, for M? = 0.45 M , channels with velocity o↵fourth column), while envelope
emissionispresent
atlargerandone for t
A&A proofs: manuscript no. Paper_L1527_arxiv
sets larger than ⇠2.5 km s 1 are expected to contain only thepanel) are
gular o↵sets. However, the angular o↵set of the envelope emisthe model
disk’s emission. In contrast, channels with smaller velocity o↵sion (⇠30
AU
in
this
case)
is smallerthantheradius at iswhich
T
<
20
K
CO
is
frozenout
andtheabundance
set
to
zero.
ature is sy
sets show emission from both the disk and the inner envelope.
the emission
originate
(⇠125
AU).
At
lower
velocities
only
enover, these
For the warm model we adopt the temperature structure of the
The most important consequence of this is that the angular o↵setvelopeT13
emissionwill
be observed(see Table 2),
velocities
observatio
model,
becausethe temperatureis until
aboveat⇠25
K in the enof optically thick emission with respect to the protostellar posithan the teclose tothedisk
systemic
velocity
the
disk
can
be seen at andthe source
13
tire
(see
Fig.
C.2).
For
the
intermediate
coldmodels,Fig.
8.
Integrated
intensity(zerothmoment)
maps
for
CO
(left
panels)emission,tion in these channels does not always directly correspond to thepositionwithenvelope
emission
on
both
sides
(see
Fig.
6,
third
18
the T13 temperature is scaled down
radius, bec
and C O (right panels) in the warm (W; top row), intermediate (I; midradial position of the emitting material.
18 by a factor to create regions
column).
Thesituationissimilar
for
C
O
(see
Fig.
6,
bottomhwith T < 20 K (schematically shown in Fig. 7). For the interme- dle row) and cold (C; bottom row) model. Emission is simulated forstantiallyaDetermining exactly where the emission originates is nonbe
becauserow), exceptthat thebecomes
optically
theforfarthe cold five hour integration using an ALMA configuration resulting in a simidiate model
theemission
temperature
is lowered
bythick
40% inandtrivial and requires knowledge of, amongst others, the disk rathat playsside ofmodelthe disk(see Fig.
6, column
one,snowline
two andatfive),instead
lar
beam
size
as
the
observations.
The
intensity(color)scale
isdi↵erentby60%,
resulting
in
a
CO
71
AU
and
23
AU,
resolution
dius, density structure and CO abundance. Nonetheless, a globalof the respectively
near side. In (seeaddition,theemission
remains
optically
thin
for
each
panel.
The
positionofthecontinuum
peakismarkedbya
blackpact figura
Fig. C.2). In the intermediate model, 25% ofview of the temperature in the disk and inner envelope can be ob00
in most of the envelope at redshifted velocities (for the adopted
lar
the mass within 125 AU (disk outer radius) has T < 20 K. In cross and the beam is shown in the lower left corner of the panels. (⇠0.1
tained through the brightness temperature in the di↵erent velocprovide mstructure; see Fig. 6, fourth column). The disk is thus not
13 disk and innerFig. 5. Face-on
viewinofSect.
the midplanein theenvelope densitythe
ity channels.
As shown
4.2, for both
CO and C18 O thecold model this is 70%, roughly consistent with the result
To exahidden
from
our
view.
for
the
fiducial
model,
showingthevelocity
component
along
the
line
temperature along the midplane is above ⇠25 K at all velocities,
duce the ofrom Kama et al. (2016) that in TW Hya ⇠60% of the dust massFor
clarity,wehaveneglected
thedust
opacityinFig.
6.
of sight a(indicatedby theon topTheof thepanel;velocthe observer
the T13 disuggesting
warm disk
andblack
innerarrowenvelope.
highesthas temperatures below 20 K. The cold model thus
a overproduced for 13 CO. This could be because a canonical CO
2 resembles
1
creasing th
13
the dust opacity from T13, that is,  = 3.1 cm g at
is on the thatleft).contain
The envelopematerial
has a infalling
rotating
velocity Adopting
ity channels
resolved
CO emission,
and wereusedClass II disk. Although for the intermediate and cold model the abundance of 10 4 is too high for the more tenuous envelope
plane temprofile (Ulrich
1976; Cassen& Moosmanthe diskmaterial 220 GHz, the dust would only become optically thick in the into construct
the radial
temperature
profile,1981),
are whilev = -2.54
andtemperature is lowered to create CO freeze-out regions, the fidu- where photodissociation may play an important role in destroyFig. 11, b
rotating
Keplerian
adopteddisk channels
outer radius of ner ⇠10
AU. If the velocity profile in the envelope is free fallv =is 2.59
km withs 1 . For
stellar velocity.
masses The
0.45
M , thesemodel fromcial temperature is used for the excitation calculation. This way ing CO, unlike in the denser disk midplane where UV radiation
125 AU is marked by the dashed line. The outer edge of the envelope (Fig. D.1),
the highest envelope velocities are found along theobservatio
are indeed expected to trace only the disk. The velocity o↵sets
emission from the surface layers is not reduced through an cannot penetrate (see e.g., Visser et al. 2009a). Another reason
is set at18the radius where the temperature in the best-fit model from line ofthe(Fig. 11).
sight toward the protostellar position, instead of o↵set to
usedT13
fordrops
C Obelow
are slightly
smaller( v = -2.22
and v =The
2.43
excitatione↵ect.
Iftheseregions
would
in
fact
also
be
colder,
the
could
be
becausethelarge-scale
cloud
emission
is
not
taken
into
20
K,
i.e.
the
CO
freeze-out
temperature.
adoptedwarmer th
the north (for redshifted emission) or to the south (for blueshifted
km stellars 1 ), andcould
becontaminatedbyenvelope
emissionifthemass is 0.45 M (A17).
di↵erencewiththe fiducial
warm
model
would
increase.
account in the model. The e↵ect of emission being resolved out
emission),as in13theinfallingrotatingcase.
Thischangesthelostellar mass is 0.45 M . However, unlike for 13 CO, the redshifted
Fig. 11.
law temperature
profiles as(q =it0.35)
shown
in Fig. 4, 5.2. CO ainPowerthe simulated
imagesis inasthe
observations,The⌧ =COand C18 at
O velocities
emission forthreecation of the
1 contour
thattheoccur
in models
both theis simu- is not as large
C18 O
emissionis
expectedtobecomeopticallythickmostlyinbut now derived from images simulated with the observed uv-coverage
Table 2. Origin of emission at di↵erent velocity o↵sets.
simulated
imagesThe
still
some
inwarmthe cenlatedwith the(that3D radiativetransfer
codefour
LIME(Brinch& Hoger- that is, the(25minutes integration).
topshow
panel shows
theemission
results for the
The under
disk andenvelope
is, columns
two andin Fig.
6 willthe disk and not already in the envelope. Moreover, the measured
model,
thatis,thefiducial
T13
temperature
structure.
Thebottom
panel
tral
most
channelswhile
theobservations
do
not.
We
will
there13
spect to t
heijde 2010).
Again,
the velocity
the disk (Rveloci< 125 AU) islook di↵erent;
see Fig.
D.2 and
Table 2).inside
For blueshifted
temperatures are similar as for CO.
shows
the
resultsforthehot
model,
that
is,
the
T13
temperature
intributed tothe comparison between observations and models
Fig. 9. Comparison
between andCOtheobservations(left column)
andKeplerianenvelope
infallingve-to 26.1fore restrict
Envelope Mstar
Disk only
Envelope onlya
Fig. 10. has
As Fig.an
9,whilebut
for C envelope
O. The 3rotatinglevel corresponds
creased by 30%. The black lines show the midplane temperature profile freeze-out
ties,
the
diskmodels
will(seethebeofmaterialthethe threedi↵erent
Fig. visible7): warm (W,southsecond column),
in- protostar,(7.5) mJy beam channel for an integration time of 25 minutes (fiveto velocities
1
v > 1.1 km
s 1used. Lowerstellarmassesgivelinelower
termediate(I, third column),
and coldThe(C, fourthcolumn). Emission isis setin the corresponding
modelsto generate
the images.
The graylocityprofile.
inclination
to
85
,
that
is,
5
o↵set
from
velocity(M )(km sArticle) number, page(km7sof1 21)
hours).
emissionwillbepresent
tothenorth.
For
redshiftedvelocities,
van Zadelh
simulated with the same visibilities as the observations (25 minutes inrepresents
themidplanetemperature
used
by
A17.
The
shadedorangetegration)and for a longer withintegrationthe
time ofdisk
5 hours inanda configuration
residuals
intheintermediate
velocity
channels,butcannotreproedge-on
envelope
west
of
the
protostartiltedbemissionsurrounding
the protostar
originates from within thearea marks the 1 uncertainty on the temperature profile derived from in particul
producing a similar
beam size. Four representative
velocity channels are
UCM
0.19 v  1.74; v 1.79 0.98  v  1.30
pleted in thshown. The
top panels showat ±2.5 km (sees , whichis expected
emission,
consistent
with theresults
OhashiWith the is
the intermediate disk
towardtheemission
observer
Fig.
C.1).
This
theobservations,sameorientation
asmodelduce the disk
the observations(13 CO,
redshifted channel),as shown
in Fig.from
4.
envelope,
while
furthest
is current
comingfromtheto originate in the
disk. The themiddle emission
panels show blueshifted
emission northUCMb0.45 v  2.54; v 2.59 1.47  v  1.79
with CO frozen out at midplane radii larger than 71 AU would
et al. 2013
from theused
disk and farof the envelopev = 2008
1.58 km sand), and T13,
the be indistinguishable
fromtothe thewarmresults
disk modelpreusing 13 COet al. (2014). Since our goal is not to model the CO isotopologue
bysideTobinet (al.
but opposite
disk.
bottom panels show redshifted emission from the disk and near side ofover, a ge
Free fall
0.19 v  1.74; v 1.79
v = 0.98【エンベロープと円盤の切り分けについて】
外側：回転インフォール（Ulich 1976; Cassen &Fig. 5. Face-on view of the midplane in the disk and inner envelope
for the fiducial model, showing the velocity component along the lineof sight (indicated by the black arrow on top of the panel; the observeris on the left). The envelope material has a infalling rotating velocity
profile (Ulrich 1976; Cassen & Moosman 1981), while the disk material
is rotating with Keplerian velocity. The adopted disk outer radius of
125 AU is marked by the dashed line. The outer edge of the envelope
is set at the radius where the temperature in the best-fit model fromT13 drops below 20 K, i.e. the CO freeze-out temperature. The adopted
stellar mass is 0.45 M (A17).

Moosman 1981）

内側：純粋な回転

Table 2. Origin of emission at di↵erent velocity o↵sets.

Envelope
velocity
UCMb
UCMb
Free fall
Free fall
Modified
free fallc

Mstar
(M )
0.19
0.45
0.19
0.45

Disk only(km s 1 )v  1.74; v 1.79v  2.54; v 2.59v  1.74; v 1.79v  2.54; v 2.59

0.45v  1.10; v

Envelope onlya(km s 1 )0.98  v  1.30
1.47  v  1.79
v = 0.98
1.47  v  1.63

1.14 0.50  v  0.98

Notes. The fiducial model is highlighted in boldface. A disk radius of
125 AU is adopted for all models, except in the modified free fall casewhere a radius of 74 AU is used (A17). For a 125 AU disk, the disk-onlychannels would then be v  0.94; v 0.98 km s 1 and the envelopeonly channels would be 0.50  v  0.66 km s 1 . (a) Only for 13 CO forwhich the redshifted emission becomes optically thick in the envelope.
(b)
Infalling rotating envelope (Ulrich 1976; Cassen & Moosman 1981).
(c)Free fall velocity modified by a factor 0.3 and a disk radius of 74 AU(Ohashi et al. 2014, A17).

仮定する速度分布に応じて
disk only/ envelope onlyの
速度オフセットがどの範囲になるか

For these parameters, channels with velocity o↵sets larger than
⇠1.1 km s 1 would contain emission originating only in the disk.
For a 0.45 M star, both the disk and envelope containmaterial that has line-of-sight velocity o↵sets smaller than
±2.5 km s 1 . What material is responsible for the emission thus
depends on where the emission becomes optically thick. The⌧ = 1 surface for 13 CO in five representative channels ( v =-2.70, -2.04, 0, 2.10 and 2.74 km s 1 ) is shown in Fig. 6 (seeFig. D.2 for the case of a free falling envelope). At blueshiftedvelocities the southern half of the disk is observed, as well as thatpart of the envelope that is behind the disk for our line of sight(see Fig. 5). The emission becomes optically thick in the disk, so【温度分布を変えた時のモデリング・疑似観測】

どの速度チャンネルで両者がどのように混じるかを解析

コンサバに見積もって，速度オフセット2.5km/s以上はDisk only

13

18

1

1

1

1

4

Hull et al.

IM Lup1%
200 AU

4
2
0

20

30

40

50

60

70

80

Polarization angle (deg)Figure 2. Histogram of the polarization angles across the map of
IM Lup seen in Figure 1, where each individual polarization angle
measurement was binned presuming it had a Gaussian distribution
with the mean and standard deviation listed in Table 1. The solid
vertical line is the 48 orientation of the minor axis of IM Lup’s
disk. The two dotted vertical lines are plotted at 48 ± 3.6 , where
the 3.6 uncertainty is the mean of the statistical uncertainties
in all of the polarization angles detected toward IM Lup. The
broad distribution of polarization angles indicates that many of thepolarization angles deviate significantly from the 48 orientation of
the minor axis of the disk.

ウイング状に広がっているのは同心円っぽくなっていることによる
4.2. Maximum grain-size estimate for IM Lup

150

0.8

100

0.6

50

0.4

0
50

0.2

100
150

0.14

200 150 100

0.12

50

0

50

0.08

的特徴が再現できる。

0.06
0.04
0.02
0.00

100 150 200

15

10

5

0

5

10

15

amax = 50 µm
amax = 61 µm
amax = 70 µm

0.014

amax = 50 µm
amax = 61 µm
amax = 70 µm

0.012

HL Tauとの Polarization fractionの比較：HL Tauより高い
ダストサイズの微妙な違い？
短軸

Polarization fraction

Polarization fraction

長軸Pol. Patternは散乱にPolarization angle (deg)0.010
To estimate the maximum size of the dust grains respon0.010
sible for the observed scattering polarization, we adopt
0.008
0.008
the IM Lup disk model based on Cleeves et al. (2016),
0.006
taking into account only millimeter-sized dust grains. The
0.006
model is one of a viscous disk (Lynden-Bell & Pringle
0.0041974) with a column density of dust ⌃c = 0.25 g cm 2
0.004
0.002
at a radius Rc = 100 au. The scale height of the mil200
150
100
50
0
50
100
150
200
200
150
100
50
0
50
100
150
200
limeter grains is Hc = 3 au at Rc , which is 0.25 ⇥ theobservations of HL Tau reported in Stephens et al. (2017).
Position along major axis (au)(Projected) position along minor axis (au)scale height of the gas and the micron-sized grains. The
Note, however, that while the slight azimuthal curvature
Figure 4. Best-fit model to the ALMA 870 µm observations of IM Lup. Upper-left: The contours represent the total intensity (Stokes
surface-density power-law index = 1.0. For the dustI), which are logarithmically scaled beginning at 3 I = 300 µJy beam 1 . The color scale represents the polarized intensity, and is also
can be explained by pure self-scattering, it could also be
grain model, we adopted the one used by Kataoka et al.
logarithmically scaled with a minimum value of 3 P = 66 µJy beam 1 . The line segments represent the polarization orientation. Upper-right:
Normalized histogram of polarization angles in the model. 0 corresponds to the minor axis orientation. The two dashed lines are ±3.6 ,6 due to the superposition of polarized emission both fromHull(2015)
et al.and Yang et al. (2016a), which assumes sphericalthe same as Figure 2. Lower-left and lower-right: The polarization fraction along the major and (projected) minor axes, respectively, in the
self-scattering and from dust grains aligned with the dust
models. The thick, solid line reflects the best-fit model with a maximum grain size of 61 µm; the dotted and dashed lines reflect models withgrains comprising a mixture of water ice (62%), organics
larger (70 µm) and smaller (50 µm) maximum grain sizes, which do not fit the data as well.
emission gradient (Tazaki et al. 2017). This superposition(30%), and silicates (8%). All fractional abundances arecan be seen most clearly in the 1.3 mm ALMA image
IM Lupby volumechange.
(2017) have shown that, in inclined disks such as IM Lup,
0.012 and are taken from Pollack et al. (1994). The
of0.012
HL Tau reported in Stephens et al. (2017); however,the near side of the disk would be significantly brighter ingrain size distribution follows the Mathis, Rumpl & NordHL Tau
those data show far more curvature in the polarization
polarized intensity than the far side if the disk were opti3.5
4.3.2. Constraining dust settling in IM Lup
sieck (MRN)-type power-law n(a) / a(Mathis et al.
cally thick (i.e., ⌧? > 1) and the grains responsible for theangles than the 870 µm IM Lup data we present here,
Based on their modeling of ALMA 1.3 mm and 870 µm
1977),
witha
fixed
minimum
grain
size
a
=
0.25
µm.
scattering-induced polarization had a significant vertical
min
0.010
0.010
continuum and spectral-line observations toward IM Lup,suggestingthat alignment of dust grains with the dust
extent. There is no obvious near/far-side asymmetry deWe calculate the absorption/scattering opacities, as
Cleeves et al. (2016) found that the gas and the microntected in our observations (Figure 1); thus, if IM Lup wereemission gradient is not contributing significantly to oursized grains in the IM Lup disk have a significant verwell as the full scattering phase matrix, with Mie theory
marginally optically thick (⌧? ⇡ 1) in the region where
870 µm polarization observations.
tical extent (with a scale height of 12 au at a radius ofpolarization is detected, it would be reasonable to con(Bohren&
Huﬀman1983).
Wethencalculatepolarization
0.008
0.008
100 au; this is consistent with Very Large Telescope (VLT)
clude that the ⇠ 60 µm-sized grains that eﬃciently scatterby varying the maximum grain size amax with the MonteSPHERE/IRDIS observations of IM Lup in optical scat4. DISCUSSION
the 870 µm photons have already settled into a relatively
tered light; Avenhaus et al. 2018). The scale height ofCarlo Radiative Transfer code RADMC-3D (Dullemondthin layer (this is similar to what was found in HL Tau by
4.1. Comparing the polarization fraction profiles ofthe millimeter-sized grains is ⇠ 4 ⇥ less; however, to reStephens et al. 2017). This settling may be inconsistentetal.
2012).
Wecanreproducetheobserved
polarization
0.006
0.006
produce the significant depletion of spectral-line emission
IM Lup and HL Tau
withthe lofting of large grains14 expected based on thein our 870 µm IM Lup observations with amax of 61 µm.
in the inner ⇠ 40 au of the disk (which they attributework by Cleeves et al. (2016); however, based on their
to absorption by optically thick dust), some amount of
These observations of IM Lup are the second 850 µm
This is very close to the 72 µm value estimated by Yang
turbulent lofting of millimeter grains is necessary.
14 Here we categorize both the ⇠ 60 µm grains in IM Lup and
ALMA
et al. (2016a),
who modeled 1.3 mm HL Tau observations
0.004 polarization observations of an inclined, full (e.g.,
0.004
Dust
settling
canbeconstrainedwithscatteringmodelsthemillimeter-sized grains in the Cleeves et al. (2016) models as
Lupafter the HL Tau observations pub“large.”
non-transition)IMdisk,
by Stephens et al. (2014).
(e.g., Kataoka et al. 2015; Yang et al. 2017). Yang et al.
HL Taulished in Stephenset al. (2017). Both the IM Lup andWe show the results of our best-fit model in Figure
the
HL Tau observations show characteristics that are
4. We0.002
find that the histogram of the polarization angles
0.002
200consistent150 10050
0 of dust
50 self-scattering
100 150 200 in the model
200 (Figure
150 4,10050
0 panel)
50 is in100good
150 200
broadly
with models
upper-rightfrom an inclined disk.
agreement with thethat of theobservations(Figureaxis2), (au)in
Position along major axis (au)(Projected)
positionalong minor
However, substantial diﬀerences between the twothat both histograms show a significant spread in angles.
sources are revealed in Figure 3, which shows 1In addition, the predicted distributions of the polarization
Figure 3. 1-dimensional cuts of the polarization fraction Pfrac across the major axis (left) and minor axis (right) of both IM Lup (thisdimensional cuts of the polarization fraction P al. acrossfraction along the major and minor axes agree well withwork) and the 870 µm HL Tau data from Stephens etfrac(2017). The minor axis has not been deprojected. The gray shaded regions havethe major(left)
and(seeminor
axis1(right)
both
IMdata);Lup in theobservations,peaking at ⇠regions,
1.1% when
thestatisticalmaximumuncertainties
grain
a vertical
widthaxis
of ±
Pfrac
Table
for theofIMLuphighest signal-to-noise
these
can
and
HL
Tau
(the
870
µm
HL
Tau
data
are
from
Stephens
size
is
set
to
be
61
µm
(see
the
thick,
solid
lines
in
the
be smaller than the 0.03% systematic uncertainty in ALMA polarization observations. The gray curves in the lower-right are cuts of the
et al. 2017).
the horizontal
major- andminor-axis
bottom panels
Figure4). resolutions
We testedofslightly
synthesized
beams; Boththe plot’s
axis
lies at theprofiles
beams’ half-maximum
levels.ofThespatialthe IM larger
Lup and HL Tau
of IM Lupat the center.
HL Tau’s(70 µm) and smaller (50 µm) maximum grain sizes, andobservations
are peakapproximately
72 au However,and 55 au, only
respectively.
minor-axis profile peaks at the center; its major-axis profound that they do not reproduce the observations (see

光学的により薄いから？

ダストサイズ（半径）が61μmで色々な観測

0.10

Distance (au)0.012

amax = 61 µm

1%

200

0.014

7

【議論１：モデル計算】

1.2
1.0

Normalized No. of pol. orientations

・PA分布のヒストグラム

6

IM Lup model

Polarization fraction

8200

Polarized flux density (mJy beam 1 )No. of pol. orientations

10

ALMA observations of polarization from dust scattering in the IM Lup protoplanetary disk

Distance (au)

IM Lupit is unlikely that the depression is because of plane-ofsky smearing of the polarization signal within the central
synthesized beam (as opposed to the 3 mm ALMA map
of HL Tau shown in Kataoka et al. 2017 and Stephenset al. 2017, where the central beam is smeared as a result of the azimuthal polarization pattern). It is more
likely that the depression along the major axis is caused
by HL Tau’s high optical depth at 870 µm – 1.3 mm
wavelengths (ALMA Partnership et al. 2015; CarrascoGonzález et al. 2016; Jin et al. 2016), which is predicted
to reduce the polarization fraction (Kataoka et al. 2015;
Yang et al. 2017).
Alternatively, the fact that the polarization fraction in
IM Lup is substantially higher than in HL Tau could be
because IM Lup is a more evolved source than HL Tau,and thus has had a longer time for grains to grow larger.
However, considering the similar estimates for the maximum grain size in the two sources (see § 4.2), it is notclear that this is the case.

Polarization fraction

12

寄与するダストが鉛直方向に光学的に厚く，かつ，鉛直方向に広がっている場合，
near-far asymmetryが見えるはず。ダス
トは結構沈殿している？

【議論２：磁場整列？】

ダストサイズ61μmでも，降着が非定常だとして期待される1-5mG @50-100auで磁場整列は厳しい？（super-paramagnetic inclusionがあってギリギリ？）。降着率大きい原始星なら可能性あり。
【議論３：βによるダストサイズ見積もりとのinconsistency】
βから示唆されるサイズよりは小さい（HD 142527と同様の議論）

26. Distortion of Magnetic Fields in a Starless Core III:
Polarization–Extinction Relationship in FeSt 1-457
R. Kandori et al. (ApJ, 857, 100)
http://ads.nao.ac.jp/abs/2018ApJ...857..100K【内容】孤立したStarless Core FeSt 1-457では，近赤外線でのExtinction に伴うPolarization mapから３次元的な磁場構造が明らかになっていた。これを踏まえ，視線方向のde-polarizationを補正すると，減光度が高いほどpolarization degreeも高くなるという相関が導かれた。

of the data is 2.43(±0.05)% mag−1. If we choose the data
points of H - Ks < 0.9 mag, the slope of the relationship gets
steeper as 2.92(±0.15)% mag−1.
The observed polarization is the superposition of thepolarization from the core itself and ambient polarization,
which is unrelated to the core. A relatively linear PH versus
H - Ks relationship is obtained after subtracting ambientpolarization components, as shown in Figure 4(b) (see also
Figure6 of Paper I). The method of the subtraction of ambientpolarization components is described in Section 2 (see also
Paper I). The shape of the P–A relationship changes, with thedisappearance of the kink and ﬂattened feature in the region of
H - Ks 2 0.9 mag. The gray plus symbols show ambient
vectors (R>144″ and P/δP 10) after the subtraction
analysis. The ambient vectors show no trend with suppressionin polarization degrees. The subtraction analysis of ambientpolarization thus seems successful. It is now clear that
subtracting the background ambient polarization is crucial fora robust P–A relationship. The slope of the relationship is4.75(±0.33)% mag−1, which is similar to the average interstellar polarization slope (Jones 1989). Comparing the slopes in
Figure 4(a) and (b) data, the P–A relationship gets steeper after

the ambient subtraction. The shallow slope in Figure 4(a) data
seems to be due to the contamination effect of ambient and core
vector components that have different polarization angles.
FeSt 1-457 is known to be accompanied by inclined distorted
magnetic ﬁelds (Paper II), which causes depolarization effects
inside the core. With a known 3D magnetic ﬁeld structure, thedepolarization correction factor can be calculated as shown in
Figure 5. In Figure 5, severe depolarization regions around theequatorial plane clearly exist, resembling the model coresimulation in Figure 3 (for the case of 45°). The depolarization
correction factor was calculated by dividing the polarization
degree distribution of the (best-ﬁt) inclined distorted ﬁeld
model by that based on the inclined uniform ﬁeld.
Figure 4(c) shows the depolarization corrected P–A relationship obtained by dividing the Figure 4(b) relation by the
depolarization correction factor shown in Figure 5. In
Figure 4(c), the P–A relationship becomes steeper than that in
Figure 4(b), reﬂecting the correction. In the relationship, therelative location of the plot points moves slightly, because
the depolarization correction factor is not spatially uniform.
The slope of the relationship is 7.74(±0.51)% mag−1. The
correlation coefﬁcient is 0.79, which is larger than those in【結果】(a)は生データ，(b)が背景成分の補正，(c)がdepolarizationの補正，(d)が磁場の視線方向射影の補正（i=45度を仮定）

19
A New 3D Maser Code Applied to Flaring Events
M.D. Gray1, L. Mason1,2 and S. Etoka1
1 Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, M13 9PL, UK; 2
Pendleton Sixth Form Centre, Salford City College, M6 7FR, UK
E-mail contact: Malcolm.Gray at manchester.ac.uk
We set out the theory and discretization scheme for a new finite-element computer code, written specifically for thesimulation of maser sources. The code was used to compute fractional inversions at each node of a 3-D domain for arange of optical thicknesses. Saturation behaviour of the nodes with regard to location and optical depth were broadly as
expected. We have demonstrated via formal solutions of the radiative transfer equation that the apparent size of themodel maser cloud decreases as expected with optical depth as viewed by a distant observer. Simulations of rotation of
the cloud allowed the construction of light-curves for a number of observable quantities. Rotation of the model cloud
may be a reasonable model for quasi-periodic variability, but cannot explain periodic flaring.
Accepted by MNRAS
https://arxiv.org/pdf/1803.01625

21
Dust concentration and chondrule formation
Alexander Hubbard1, Mordecai-Mark Mac Low1, Denton S. Ebel2
1 Dept. of Astrophysics, American Museum of Natural History, New York, NY, USA; 2 Dept. of Earth and Planetary
Sciences, American Museum of Natural History, New York, NY, USA
E-mail contact: ahubbard at amnh.org
Meteoritical and astrophysical models of planet formation make contradictory predictions for dust concentration factors
in chondrule forming regions of the solar nebula. Meteoritical and cosmochemical models strongly suggest thatchondrules, a key component of the meteoritical record, formed in regions with solids-to-gas mass ratios orders ofmagnitude above background. However, models of dust grain dynamics in protoplanetary disks struggle to surpass
factors of a few outside of very brief windows in the lifetime of the dust grains. Worse, those models do not predictsignificant concentration factors for dust grains the size of chondrule precursors. We briefly develop the difficulty inconcentrating dust particles in the context of nebular chondrule formation and show that the disagreement is sufficiently
stark that cosmochemists should explore ideas that might revise the concentration factor requirements downwards.
Accepted by Meteoritics & Planetary Sciences
http://arxiv.org/pdf/1803.10047

22 Searching for Hα emitting sources around MWC758: SPHERE/ZIMPOL high-contrast imaging
N. Húelamo1, G. Chauvin2,3, H.M. Schmid4, S.P. Quanz4, E. Whelan5, J. Lillo-Box6, D. Barrado1, B.
Montesinos1, J.M. Alcaĺa7, M. Benisty2,3, I. de Gregorio-Monsalvo6,8, I. Mendigut́ıa1, H. Bouy9, B. Meŕın10,
J. de Boer11, A. Garufi12, and E. Pantin13
MWC758 is a young star surrounded by a transitional disk. Recently, a protoplanet candidate has been detectedaround MWC758 through high-resolution Lʹ-band observations. The candidate is located inside the disk cavity at a
separation of ∼111 mas from the central star, and at an average position angle of ∼165.5. We have performed
simultaneous adaptive optics observations of MWC758 in the Hα line and the adjacent continuum using SPHERE/ZIMPOL at the Very Large Telescope (VLT). We aim at detecting accreting protoplanet candidates through spectral
angular differential imaging observations. The data analysis does not reveal any Hα signal around the target. The
derived contrast curve in the B Ha filter allows us to derive a 5σ upper limit of ∼7.6 mag at 111 mas, the separationof the previously detected planet candidate. This contrast translates into a Hα line luminosity of LHα <∼ 5 × 10−5
L⊙ at 111 mas, and an accretion luminosity of Lacc < 3.7 × 10−4 L⊙. For the predicted mass range of MWC758b,
0.5–5 MJup, this implies accretion rates of Ṁ <∼ 3.4 × (10−8–10−9) M⊙ yr−1, for an average planet radius of 1.1
RJup. Therefore, our estimates are consistent with the predictions of accreting circumplanetary accretion models for
Rin = 1 RJup. In any case, the non-detection of any Hα emitting source in the ZIMPOL images does not allow us tounveil the true nature of the Lʹ detected source.
Accepted by A&A Letters
http://arxiv.org/pdf/1803.09264

24

The Herschel-PACS Legacy of Low-mass Protostars: The Properties of Warm and Hot Gas Components and Their
Origin in Far-UV Illuminated Shocks
A. Karska1,2,3, M. Kaufman4, L.E. Kristensen5, E.F. van Dishoeck3,2, G.J. Herczeg6, J.C. Mottram7, L︎ ukasz
Tychoniec3, J.E. Lindberg8, N.J. Evans II9, J.D. Green9,10, Y.-L.. Yang9, A. Gusdorf11, D. Itrich12 and N. Siodmiak13
Recent observations from Herschel allow the identification of important mechanisms responsible for the heating of gas
surrounding low-mass protostars and its subsequent cooling in the far-infrared (FIR). Shocks are routinely invoked toreproduce some properties of the far-IR spectra, but standard models fail to reproduce the emission from key molecules,
e.g. H2O. Here, we present the Herschel-PACS far-IR spectroscopy of 90 embedded low-mass protostars (Class 0/I).
The Herschel-PACS spectral maps covering ∼ 55 − 210 μm with a field-of-view of ∼50” are used to quantify the gas
excitation conditions and spatial extent using rotational transitions of H2O, high-J CO, and OH, as well as [O I] and [CII]. We confirm that a warm (∼300 K) CO reservoir is ubiquitous and that a hotter component (760 ± 170 K) is
frequently detected around protostars. The line emission is extended beyond ∼1000 AU spatial scales in 40/90 objects,
typically in molecular tracers in Class 0 and atomic tracers in Class I objects. High-velocity emission (> 90 km s−1) is
detected in only 10 sources in the [O I] line, suggesting that the bulk of [O I] arises from gas that is moving slower thantypical jets. Line flux ratios show an excellent agreement with models of C-shocks illuminated by UV photons for preshock densities of ∼105 cm−3 and UV fields 0.1-10 times the interstellar value. The far-IR molecular and atomic lines
are a unique diagnostic of feedback from UV emission and shocks in envelopes of deeply embedded protostars.
Accepted by The Astrophysical Journal Supplement Series
http://iopscience.iop.org/article/10.3847/1538-4365/aaaec5/meta
http://arxiv.org/pdf/1802.03379

25
The Coordinated Radio and Infrared Survey for High-Mass Star Formation III. A catalogue of northern ultra-compact H
II regions
I. E. Kalcheva1, M. G. Hoare1, J. S. Urquhart2, S. Kurtz3, S. L. Lumsden1, C. R. Purcell4,5 and A. A. Zijlstra6
A catalogue of 239 ultra-compact H II regions (UCHIIs) found in the CORNISH survey at 5 GHz and 1.5ʹʹ resolution in
the region 10◦ < l < 65◦, |b| < 1◦ is presented. This is the largest complete and well-selected sample of UCHIIs to dateand provides the opportunity to explore the global and individual properties of this key state in massive star formation at
multiple wavelengths. The nature of the candidates was validated, based on observational properties and calculated
spectral indices, and the analysis is presented in this work. The physical sizes, luminosities and other physical properties
were computed by utilising literature distances or calculating the distances whenever a value was not available. The
near- and mid-infrared extended source fluxes were measured and the extinctions towards the UCHIIs were computed.
The new results were combined with available data at longer wavelengths and the spectral energy distributions (SEDs)
were reconstructed for 177 UCHIIs. The bolometric luminosities obtained from SED fitting are presented. By comparing
the radio flux densities to previous observational epochs, we find about 5% of the sources appear to be time variable.
This first high-resolution area survey of the Galactic plane shows that the total number of UCHIIs in the Galaxy is ∼ 750− a factor of 3-4 fewer than found in previous large area radio surveys. It will form the basis for future tests of models of
massive star formation.
Accepted by A&A
https://arxiv.org/pdf/1803.09334

27

A radial velocity survey of the Carina Nebula’s O-type stars
Megan M. Kiminki1 and Nathan Smith1
1 Steward Observatory, University of Arizona, 933 N. Cherry Avenue, Tucson, AZ 85721, USA
E-mail contact: mbagley at email.arizona.edu
We have obtained multi-epoch observations of 31 O-type stars in the Carina Nebula using the CHIRON spectrograph on
the CTIO/SMARTS 1.5-m telescope. We measure their radial velocities to 1–2 km s−1 precision and present new or
updated orbital solutions for the binary systems HD 92607, HD 93576, HDE 303312, and HDE 305536. We also
compile radial velocities from the literature for 32 additional O-type and evolved massive stars in the region. The
combined data set shows a mean heliocentric radial velocity of 0.6 km s−1. We calculate a velocity dispersion of ≤9.1
km s−1, consistent with an unbound, substructured OB association. The Tr 14 cluster shows a marginally significant 5
km s−1 radial velocity offset from its neighbor Tr 16, but there are otherwise no correlations between stellar position andvelocity. The O-type stars in Cr 228 and the South Pillars region have a lower velocity dispersion than the region as awhole, supporting a model of distributed massive-star formation rather than migration from the central clusters. We
compare our stellar velocities to the Carina Nebula’s molecular gas and find that Tr 14 shows a close kinematic
association with the Northern Cloud. In contrast, Tr 16 has accelerated the Southern Cloud by 10–15 km s−1, possibly
triggering further massive-star formation. The expansion of the surrounding H II region is not symmetric about the Otype stars in radial velocity space, indicating that the ionized gas is constrained by denser material on the far side.
Accepted by MNRAS
http://arxiv.org/pdf/1803.07057

28
The CARMA-NRO Orion Survey
Shuo Kong1, H́ector G. Arce1, Jesse R. Feddersen1, John M. Carpenter2, Fumitaka Nakamura3, Yoshito Shimajiri4, Andrea
Isella5, Volker Ossenkopf-Okada6, Anneila I. Sargent7, Álvaro Śanchez- Monge6, Sümeyye T. Suri6, Jens Kauffmann8,
Thushara Pillai9, Jaime E. Pineda10, Jin Koda11, John Bally12, Dariusz C. Lis13,14, Paolo Padoan15,16, Ralf Klessen17,18,
Steve Mairs19, Alyssa Goodman20, Paul Goldsmith21, Peregrine McGehee22, Peter Schilke6, Peter J. Teuben23, Maŕıa Jośe
Maureira1, Chihomi Hara24, Adam Ginsburg25, Blakesley Burkhart26, Rowan J. Smith27, Anika Schmiedeke10,6, Jorge L.
Pineda21, Shun Ishii2, Kazushige Sasaki28, Ryohei Kawabe3, Yumiko Urasawa28, Shuri Oyamada29 and Yoshihiro Tanabe30
We present the first results from a new, high resolution, 12CO(1-0), 13CO(1-0), and C18O(1-0) molecular line survey of theOrion A cloud, hereafter referred to as the CARMA-NRO Orion Survey. CARMA observations have been combined withsingle-dish data from the Nobeyama 45m telescope to provide extended images at about 0.01 pc resolution, with a dynamic
range of approximately 1200 in spatial scale. Here we describe the practical details of the data combination in uv space,including flux scale matching, the conversion of single dish data to visibilities, and joint deconvolution of single dish andinterferometric data. A ∆-variance analysis indicates that no artifacts are caused by combining data from the two instruments.
Initial analysis of the data cubes, including moment maps, average spectra, channel maps, position-velocity diagrams, excitation
temperature, column density, and line ratio maps provides evidence of complex and interesting structures such as filaments,
bipolar outflows, shells, bubbles, and photo-eroded pillars. The implications for star formation processes are profound andfollow-up scientific studies by the CARMA-NRO Orion team are now underway. We plan to make all the data products
described here generally accessible; some are already available at https://dataverse.harvard.edu/dataverse/CARMA-NRO-Orion.
Accepted by ApJS
https://arxiv.org/pdf/1803.11522