THE STAR FORMATIONNEWSLETTER
No. 265 (#16-20)
塚越崇(茨城大)13 Feb. 2015

Herschel Far-infrared Spectral-mapping of Orion BN/KL Outflows:
Spatial Distribution of Excited CO, H2O, OH, O, and C+ in Shocked Gas
J.R.Goicoechea et al.
Spatial distribution of excited CO, H2 O, OH, O and C+ in Orion BN/KL outflows• Herschelによる複数輝線イメー

3

H2 v=1−0 S(1)

c)

a)

ジング観測

b)

• Orion BN/KL周辺• hot core
• low- velocity outflow• high- velocity outflow• 二つのH2 Peak
Peak1 and Peak2• Herschel HEXOS GT Key
Programによる観測d)

H2O(423-312)

[CII]

[OI]

e)

f)

OH

79µm cont.

g)

h)

i)

CO(35-34)

CO(24-23)

CO(15-14)• 高分解能によってPeakを分解• 多くの輝線を検出•

Peak1とPeak2で強度が異なる

Figure 1. Line surface brightness maps obtained with Herschel /PACS toward Orion BN/KL. The FoV is 20 ⇥20 . The areas used tocompute the line luminosities in the three representative regions discussed in the text (Peak 1, Peak 2 and ambient cloud) are shown withblack and pink squares respectively. The black crosses and the pink star show the exact positions where the spectra shown in Figure 2
are extracted. The position of the BN source and that of the Orion Hot core are also shown in each panel as colored stars. Panel a) H⇤2
emission distribution in the region (Bally et al. 2011). The 4 blue contours show the SCUBA 450 µm continuum emission at 3 and 20, 50
and 90% of the emission peak respectively Johnstone & Bally (1999). Panel b) shows the [C ii] 158µm line surface brightness. The positionof the Trapezium stars are marked with blue stars. The white contours in panels d) and e) represent H2 O and OH line absorption contours

8ライン強度比Table 1Model components and resulting source-averaged column densities toward Peak 1 (assuming a line-width of 30 km s
†Assuming no contribution in the hot component (see10 text).

a)d)rotation diagram

Goicoechea et al.

Componentb) shock-excited
Hot
Warm shock-excited
Cool plateau outflow(s)

Tk
(K)
⇠2500
⇠500
⇠200

n(H2 )(cm 3 )c)7⇠2⇥10
⇠2⇥107⇠2⇥106Size(arcsec)
⇠10
⇠10
⇠30

Trot (CO)'1000 K for J>40 are very high, in the range
Tk '3500 to 1000 K respectively.
The warm and hot components dominate the CO line
emission for J > 20. Assuming the same H2 densities inferred from the CO v=1-0 band analysis, we
run models varying Tk and N (CO) around the values derived by González-Alfonso et al. (2002) and
Sempere et al. (2000). A good fit to the data ise) '500 K and Tk, hot '2500 K. Thef) infound for Tk, warm
ferred source-averaged CO column densities are similar to those inferred from the CO v=1-0 vibrational
band: N (CO)warm ' 2.0 ⇥ 1019 cm 2 and N (CO)hot '1.5 ⇥ 1016 cm 2 .
A third cooler and more extended component is needed
to reproduce the 12 CO and 13 CO emission with J<20.
We associate this component with the extended plateau(outflowing and swept-up material). The gas density inthe plateau is estimated to be n(H2 )⇡(1-2)⇥106 cm 3(e.g., Blake et al. 1987; Melnick et al. 2010; Tercero

Figure 4. Line surface brightness ratios for selected lines observed with Herschel /PACS toward Orion BN/KL. The areas used to computethe line luminosities in the three representative regions discussed in the text (Peak 1, Peak 2 and ambient cloud) are shown with black andpink squares respectively. The black crosses and the pink star show the exact positions where the spectra shown in Figure 2 are extracted.
The position of the main sources discussed in the text: the Trapezium cluster stars (blue stars), the BN source and the position of the Hot
Core (colored stars) are also shown.

non-LTE model of Peak1

in the region. The exact values will depend on the actual
column density of material toward each line-of-sight, thelocal excitation conditions and radiative transfer e↵ects.
Finally, the line ratio maps shown in Figures 4a, 4d and5 suggest that there are not major di↵erences between
the physical conditions prevailing in the Peak 1 region
and those of Peak 2 (see also sec. 4.2). In the following,
we focus our analysis in the interpretation of the spectra
toward Peak 1.
4.2. CO, H2 O and OH rotational diagrams
Figure 6 shows all CO, H2 O and OH detected lines toward Peak 1 in a rotational diagram. In these plots weassumed extended emission over a ⇠3000 ⇥3000 aperture.
Owing to the high densities previously inferred toward
Peak 1, &106 cm 3 from H2 line observations (Rosenthal et al. 2000) and &107 cm 3 from CO and H2 O rovibrational lines (González-Alfonso et al. 2002), the rotational temperatures (Trot ) derived from the CO rotational diagrams are a good lower limit to the gas temperature (Tk &Trot ).
The 12 CO rotational diagram shows a notable contribution of very high-J lines (J>40) that cannot be fittedwith a single rotational temperature in LTE (a straightline in the plot). A smilar population diagram of the mid-1 ).

curved H2 level distribution is not compatible with a lowdensity isothermal component (see also Sect. 6.2). As
H2 , the 12 CO rotational diagram suggests the presenceof at least two temperatures components, a warm component with Trot ⇡400 K and a hotter component with
Trot ⇡1000 K. For comparison, the CO emission in theambient cloud is dominated by a cooler component with
Trot ⇡180 K (blue triangles in Fig. 6a). This is the warmand extended face-on PDR (the interface between theforeground H ii region and the molecular cloud), heatedby FUV photons from the Trapezium stars (see Peng etal. 2012b, for large-scale mapping of several mid-J CO
narrow lines).
A rotational diagram of the 13 CO lines detectedwith PACS toward Peak 1 provides Trot ⇡200 K, thus
lower than the Trot inferred from 12 CO. The measured
12
CO/13 CO intensity ratio for the J=16-15 and 15-14
lines is 55±5, lower than the 12 C/13 C isoptopic ratio of⇠79±7 in Orion(Langer & Penzias 1990). For J>16,
however, the 12 CO/13 CO intensity ratio becomes equalor larger than the isotopic value, showing that while thesubmm 12 CO low- and mid-J lines are optically thick,the high-J far-IR lines (J>16) are unambiguously optically thin.
H O and OH have a more complex rotational spec-

N (CO)(cm 2 )
⇠1.5⇥1016⇠2.0⇥1019⇠2.0⇥1019

N (H2 O)(cm 2 )
⇠2.0⇥1016(⇠2.0⇥1017 )†⇠6.7⇥1017a)

N (OH)(cm 2 )
⇠1.5⇥1016(⇠5.0⇥1016 )†⇠1.5⇥1017

et al. 2010). This is, at least a factor 10 lower thanin the shock-excited gas traced by the very high-J CO
rotational lines and by the CO v=1-0 band (hot andwarm components). Modelling all components together(adding the emission from each component) we obtaina satisfactory fit to the CO ladder with Tk, plat '200 K,n(H2 )'2⇥106 cm 3 and N (CO)plat '2⇥1019 cm 2 forthe cool component (Fig. 8). The complete model agrees
with observations within a ⇠20 % and is also consistent
with the 13 CO lines assuming a 12 C/13 C isotopic ratio of
80.

b)4.3.2. H O model
2

The extended plateau in Orion BN/KL shows enhanced H2 O column densities compared to more quiescent interstellar environments. Still, the exact water
vapor abundances are controversial (cf. Cernicharo et al.
1994, 1999, 2006; Harwit et al. 1998; Wright et al. 2000;
Melnick et al. 2000, 2010). Most of these observations
were carried out toward the core of Orion BN/KL wherea very strong mid- and far-IR radiation field from the Hot
Core region exists. The angular resolution of the above
observations range from the ⇠30 and ⇠8000 of SWAS andthe ISO/LWS, to the 4400 -1200 beam of Herschel /HIFIin the submm. The ⇠1200 resolution of Herschel/PACS
allows us to resolve the intrinsic H2 O emission toward
Peak 1. In addition, the continuum emission is significantly weaker and lines are not contaminated by absorption toward the Hot Core.
A single temperature component can not fit theobserved H2 O emission lines toward Peak 1. However, most of the H2 O lines with Ek /k<750 K
can be reasonably fitted with the extended plateau
component (Tk, plat '200 K, n(H2 )'2⇥106 cm 3 andN (H2 O)plat '6.7⇥1017 cm 2 ). This agrees with detailed
models of H2 O lines detected by Herschel /HIFI toward
the Hot Core position (Melnick et al. 2010).
Even taking into account the moderate far-IR radiation field toward Peak 1, the H2 O lines arising fromhigher energy levels require higher excitation conditions.
The detailed excitation analysis of the H2 O v=1-0 vibrational band at ⇠6.7 µm toward Peak 1 concluded
that hotter and denser gas contributes to the mid-IR
H2 O ro-vibrational emission. Since abundant water va-

c)

Goico

dances with accuracy (e.g., Chen et al. 2014).
Table 2
Peak1のgascoolingに寄与しているのは？
23
200
king N (H2 )⇡10 cm
in a 30 beam for theObserved luminosities and cooling budget (assuming d=414 pc)
plateau component (Blake et al. 1987; MelSpecies
Lobs (L )%aLobs (L )%b4
et al. 2010),
we
obtain(CO)⇡2⇥10,plat• 半分はH2、続いてCOin Peak 1c in Peak 1c
in mapc
in mapc
5
5(H2 O)⇡0.7⇥10
and plat (OH)⇡0.2⇥10
abunFIR 12 CO
18.7⇠(22±5)%
84
⇠31%
es.
Hence, in the less excited swept-up gas,
FIR H2 O
7.0⇠(8±2)%
32⇠12%
r vapor is •notFIRだけで見るとCOが突出
the dominant O-bearing species.
FIR OH3.5⇠(4±1)%
13⇠5%ver, if we only consider the hot shock-excited
FIR [O i]
3.5⇠(4±1)%
18⇠7%
onent, the derived H2 O/CO abundance ratio
FIR [C ii]
0.3⇠(0±0)%
7⇠3%
13
FIR CO0.3⇠(0±0)%
2⇠1%
uch higher (⇠1.3).
Rosenthal et al.
(2000)
12 CO v=1-0
d
d
⇠6.1
⇠7%
⇠14
⇠5%
red that the column density of hot H2 with
d
d
20
2
shock
H2 O v2 =1-0
⇠1.3
⇠2%
⇠3
⇠1%
>1200 K toward
Peakproperties
1 is N (H2 ).10 cm . Takee
MIR H2
⇠43
⇠51%
⇠102
⇠37%
he columns from this component alone, we de4
Total
⇠84
100%
274
100%
C-shockだとH2Oが重要な
ne hot (H2•O)&2⇥10, hot (CO)&1.5⇥10 4 and3
5
FIR dust⇠4⇥10⇠10OH)&0.8⇥10 5 .

••coolantになるので観測と合わない

5. Model uncertainties and Peak 1/2 di↵erencesa In a ⇠3000 ⇥3000 aperture (⇠12500 AU ⇥ 12500 AU).
b In the entire ⇠20 ⇥20 map (⇠0.25 pc ⇥ 0.25 pc).
c Absolute calibration accuracy up to ±30%.

• H2Oを減らすメカニズム
r multi-component
model satisfactorily reproduces
d Scaled from ISO/SWS observations (González-Alfonso et al.
fluxes of most observed lines. Despite this model
2002).
e Scaled from ISO/SWS observations (Rosenthal et al. 2000).
ion is certently •not初期のOを減らすかUV
unique, the agreement withobdissociationか -> J-type shock?
tions suggests that the model captures the averagecal conditions of the shocked-exited gas.
C- and
J-typeshocks?
e best fit to•themized
CO J=41
to 48
absolute
line fluxes
the cooling budget in shocked gas. This tabulation also
63
gives Tk '2300-3000 K and n(H2 )'(3-4)⇥10 cm
includes an extrapolation of the mid-IR H2 line fluxes
shock?
16• UV irradiatedN (CO)=(5.0-2.5)⇥10cm 2 C-type
for Orion
Peak 1.
and of the CO and H2 O ro-vibrational emission detected
e values suggest that the postshock4gas -3
where theISO/SWS(Rosenthal
et al. 2000; González-Alfonso et
4by
6 cm-3 & G <100
二つの解:
n=10
cm
&
G
=5x10
or
n=10
•
0
0
excitation CO, H2 O and OH lines arise might be
al. 2002). Their measured luminosities have been multio ⇠5 times less dense than our assumed value
plied by
7
3
(H2 )'2⇥10 • cmPeak1はextremeな環境である(but see Sempere et al. 2000;(3000 ⇥3000 )/(1400 ⇥2000 )'3 to approximately take into acález-Alfonso et al. 2002). In such a case, the H2 O
count the di↵erent aperture sizes. Roughly half of
mn densities in the hot component will be a factor
the total line luminosity is provided by H2 rotational
gher and the H2 O/CO abundance ratio will increase
and vibrational line cooling. CO is the next most im⇠1.3 to ⇠1.6. Finally, the presence of a moderate

A GALEX based search for the sparse young stellar populationin the Taurus-Aurigae star forming region
Ana I Comez de Castro et al.
– 33 –• UVによるTTS同定も有効なツールの一つである•観測例がほとんど無い•

UVデータを使った新しいcriteriaを提案– 29 –• GALEX all sky surveyのUVデータ– 30 –を使ってTaurus領域で確認••

50cm NUV&FUVの宇宙望遠鏡

Fig. 6.— Density of NUV GALEX sources in the TMC. The densities are color coded instars per 3 arcmin2 (see lateral bar). The molecular gas distribution, as unveiled by the
2MASS extinction map (Lombardi et al. 2010) is overlaid for reference.

可視や赤外データも使用• 63天体の新検出に成功

G0
G1 G5
K2
K1
K5
K4
M0

•

既知の天体もしっかり検出• 今後のプロシージャとなる

M2
M4e
M1
M2

A new approach to determine optically thick H2 cooling and its
effect on primordial star formation
Tilman Hartwig et al.

• 光学的に厚い部分のH2 cooling rateを見積もる新しい手法•

TeeCol algorithmがベース•

Sobolev approximationと比較• 収縮するコアにおいて10%程度のエラーでcooling rateを見積もれる•

原始星+円盤が出来ると対称性が崩れてエラーが大きくなる• 温度が低く出るのでframentationが起きやすい• 質量降着率も小さくなりPop III星の質量が小さいものが出来るPre-main sequence accretion in the low metallicity Galactic starforming region Sh2-284
V.M.Kalari and J.S.Vink
a Identification no. in Delgado et al. (2010).

考えられている
観測でも低Zでは早い散逸時間が指摘されている•

質量降着率はどうか？

• 低Z領域Sh2-284でPMSの質量降着率を観測的に調べる• Halpha emissionを観測•

dures. Wavelength calibration was performed using Ne
arc spectra using the IRAF identify task. Spectra werecorrected for sensitivity using the Sutherland extinction
curve. The variable pupil design of the SALT telescope
makes absolute flux calibration impossible. The final ex10m Southern African Large Telescope(SALT)による観測
tracted 1-dimensional spectra are shown in Fig 2. The
average signal to noise is ⇠ 20 - 40.

••

Fig. 1.— Sh 2-284 in H↵ taken from Parker et al. (2005). Northis up and east is to the left. Stars observed with SALT are marked
by circles, and stars with VIMOS spectra are marked by crosses.
The central cluster Dolidze 25 is also marked. At d = 4000 pc, 0.15
corresponds to ⇠ 10 pc.

IRデータを元に6天体を選定

VIMOS/VLTの観測データも使用 [Cusano+2011]•

サンプルは全部で24天体に

F
can
top
sta
the

2.1. Basic analysis of the SALT data

Equivalent widths (EW) of emission and absorption
lines were measured using the IRAF task splot. Spectral types were derived from the EW of the (i) Mg II
5711 Å(ii) Mn II 6015 Å(iii) Ca I 5711 Å lines following
the method devised by Hernández et al. (2004). E↵ective temperatures (Te↵ ) were estimated using the spectral
type-Te↵ calibration of Kenyon & Hartmann (1995). H↵
line profiles are displayed in Fig. 3. Five stars show H↵
in emission. Combined with the fact that these stars allshow near to mid-infrared excess resembling PMS stars,

Normalized Intensity•

Normalized Intensity• metallicity(Z)↓で質量降着率↑の関係がNormalized Relative flux

b Identification no. in Puga et al. (2009).

TABLE 4Median Ṁacc of PMS stars in the 1 -2 M

• 物理量の導出

Region

N

Nexp

Age
(Myr)

log Ṁacc(M yr 1 )Reference

L1641

23

-

1

7.1

Orion nebula cluster⇢-Ophiuchus
NGC 2264
Taurus-Auriga
IC 348
Lupus

41
5
7
3
2
3

75±10
5±3
15±4
2±1
6±2
-

1.5
0.5-2a3
3
1.5-4b
3

7.65
7.95
7.8
7.4
7.8
7.6

Cha II
Trumpler 37

7
42

(b)

3-4
4

7.9
7.85

Fang et al. (2009);
Caratti o Garatti et al. (2012)
Manara et al.(2012)
Natta et al.(2006)
Barentsen et al.(2013)
Calvet et al.(2004)
Dahm (2008)
Alcalá et al.(2014);
Herczeg & Hillenbrand (2008)
Antoniucci et al.(2011)
Sicilia-Aguilar et al.(2010);
Barentsen etal.(2011)

Z = 0.4 Z
LMC Field3 Pop.2
LMC Field3 Pop.1
LMC Field2
LMC Field1
SN 1987 A

146
96
325
490
104

6
8
9
13
14

7.2
7.2
7.4
7.6
7.6

Spezzi et al.(2012)
Spezzi et al.(2012)
Spezzi et al.(2012)
Spezzi et al.(2012)
De Marchi et al.(2010)

Z = 0.2 Z
Sh 2-284

20

3.5

7.5

This work

NGC 346dd
NGC
346
V. M. Kalari & J. S. Vink -

1
20

7.1
7.7

De Marchi et al.(2011)
De Marchi et al.(2011)•

質量と年齢; HR diagram•

降着ルミノシティ; log Lacc ∝ log LHa•

質量降着率; Macc∝Lacc•

円盤のIR class; SED

Z =Z• 質量降着率とZの関係•

Z = 0.1 Z

ほとんど違いは見られない

Note. — (a) ⇢-Ophiuchus is oft-quoted with an age of 0.5 Myr (Natta et al. 2006), however theages of the accreting PMS stars are thought to be as old as 2 Myr (Wiliking et al. 2008; Rigiliaco etal. 2010). (b) The total cluster mass of Trumpler 37 is debated, as studies of stellar content reveal adearth of near6.0
solar mass members, and a deviant stellar mass function slope (Errman et al. 2012).
(c) The age of IC 348 is debated, with initial studies suggesting an age as young as 1.5 Myr (Herbig1998), but later revised to around 4 Myr (Mayne & Naylor 2007). (d) Taken from Figure 8 of Spezzi
et al. (2012). 6.5

6.0

6.5

log Ṁacc (M yr 1 )10

Maccの中間値で比較

log Ṁacc (M yr 1 )•at different Z

7.0

7.5

8.0

7.0
7.5
8.0
8.5

8.5
9.0
9.0
0.05

1M <M <2M

0.00

0.05

0.10

0.15

log M (M )

0.20

0.25

0.30

9.5

5.5

6.0

6.5

log t (yr)

7.0

7.5

The HIFI spectral survey of AFGL2591(CHESS). III.
Chemical structure of the protostellar envelope
Maja Kazmierczak-Barthel et al.

M. Kaźmierczak-Barthel et al.: Chemical• アバンダンス分布をモデル化して、観測で得られているスペクトルを再
cal radiative transfer model (ratran) for the envelope of AFGLI
Table 1. Abundances (relative to H2 ) estimated with the spheri2591, based on HIFI and JCMT data.

現する• well known isolated massive protostar AFGL2591• disk-like structure, outflow, envelope…

• 物理モデルを与えて輝線モデリング• RATRAN• chi-squareで合わせる• アバンダンス分布• constant• jump-like distribution•

100K (water関連; van der Tak+2006)•

230K (N-beargin species Rodgers & Charnley 2001)• アバンダンスは５つに分類された

Molecule
CO
CN
CS
HCO+
H2 CO
N2 H+
CCH
NO∗
OCS∗
OH∗
H2 CS∗
C∗
C+∗
CH∗
O∗∗
HCN
HNC
NH3
SO
SO2
H2 S
H2 O
HCl
CH3 OH

Inner abundance Outer abundance
2 × 10−4
1 × 10−9
4 × 10−8
3 × 10−11
3 × 10−8
1 × 10−11
1 × 10−8
8 × 10−13
8 × 10−10
8 × 10−12
8 × 10−8
5 × 10−8
4 × 10−8
2 × 10−8
4 × 10−9
4 × 10−6
7 × 10−6
1 × 10−8
2 × 10−5
5 × 10−6
5 × 10−7
1 × 10−7
1 × 10−8
4 × 10−6
4 × 10−10
5 × 10−7
5 × 10−9
5 × 10−6
5 × 10−9
8 × 10−8
8 × 10−10
1 × 10−4
1 × 10−9
4 × 10−6
4 × 10−10
8 × 10−7
8 × 10−8

Jump at
const
const
const
100 K
100 K
100 K
100 K
const
const
const
const
const
const
const
const
230 K
230 K
100 K
100 K
100 K
100 K
100 K
100 K
100 K

2

C
a
g
a
e

(
c
d
(
o
l

W
a
a
l
i
s
H
V
i

a
t
l
i
u

s

• 時間発展化学計算••

化学モデル: 13元素からなる1268分子種、38812個の反応•

cosmic-ray ionization rate of 5E-17 1/s•

photoreaction by interstellar UV•

grains are uniform amorphous olivine particles 3 g/cm^3 and 0.1µm•

gas-grain interaction•

surface recombination

初期アバンダンス導出の為0D分子雲の化学計算を実行•

三種類の初期値を使用 - low metallicityモデルがベストフィット(C/O=0.44)• 20個の計算を走らせ観測(RATRAN)と比較•

基本的には良く合うが合わせにくいケースもある•

HCN/HNC, complex organics, S-bearing species• 初期アバンダンスモデルを変えても合う。複雑な分子は再現するが、シンプルな分子の再現がうまくいかない•

初期温度が高くなる為

1e−02

radiative transfer model
chemical model

1e−03
1e−04
1e−05

[CO]

1e−06

1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12
1e−13
1e−14

radiative transfer model
chemial model

[CN]
1e+13

1e+13

1e+14

1e+15

1e+14

1e+15

1e−04
1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12
1e−13

radiative transfer model
chemical model

[CS]
1e+13

Radius [m]constant
1e+14

1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12
1e−13

radiative transfer model
chemical model

radiative transfer model
chemical model

1e−06
1e−08
1e−10

[HCO+]
1e+13

1e−04

1e+14

1e+15

radiative transfer model
chemical model

1e−06
1e−08
1e−10
1e−12

[CCH]

1e−14
1e+13

1e+14

1e−12

[H2CO]

1e−14

radiative transfer model
chemical model

1e+14

radiative transfer model
chemical model

1e−04
1e−06
1e−08
1e−10
1e−12

[NO]

1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12
1e−13
1e−14

1e+14

1e−14

[OCS]

1e−16

1e+15

1e+13

1e+15

1e−04
1e−06
1e−08
1e−10
1e−12

[H2CS]

[CI]

1e−14
1e+13

1e+13

1e+14

1e+15

1e+14

radiative transfer model
chemical model

[OH]
1e+13

radiative transfer model
chemical model

1e−02

radiative transfer model
chemical model

1e+14

1e−04
1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11

1e+15

1e+14

1e+00
1e−02
1e−04
1e−06
1e−08
1e−10
1e−12
1e−14
1e−16

1e+15

radiative transfer model
chemical model

[CII]
1e+13

1e+14

Radius [m]1e+15

Radius [m]+[N2H ]
1e+13

1e+13

radiative transfer model
chemical model

1e+13

1e+15

Fig. 2. The radial distribution of abundances with error-bars of CO, CN, and CS at the time of 104 years. The ratran-derived values(dashed lines) are compared with the results of the best-fit 1D chemical model (solid lines); see Table 4 for the specifications.
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12
1e−13
1e−14
1e−15

1e−03
1e−04
1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11

1e+15

100K jump
外側でX大

1e+14

1e+15

Radius [m]1e−04

radiative transfer model
chemical model

1e−06
1e−08

1e−04

1e−10

1e−06

constant

radiative transfer model
chemical model

1e−02

1e−08

1e−12

[CH]

1e−14

1e−10

[OI]

1e−12

1e−16
1e+13

1e+14

(観測少ない)

1e+15
1e+13

Radius [m]

1e+14

1e+15

Fig. 4. The same as Fig. 2. Radiative transfer models with constant abundance are based on only a few observed transitions.

1e+15

Radius [m]Fig. 3. The same as Fig. 2. Radiative transfer models are with a jump at 100 K (lower inner abundance) .
M. Kaźmierczak-Barthel et al.: Chemical structure of the AFGL 2591 envelope

1e−03
1e−04
1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11

radiative transfer model
chemical model

1e−05
1e−06
1e−07
1e−08
1e−09
1e−10
1e−11
1e−12

radiative transfer model
chemical model

230K jump
内側でX大lack of major high-temperature production channels may also
play a role.
1e−02
[HNC]
[HCN]
radiative transfer model+
1e−02 UV intensity
radiative transfer model
radiative transfer
Where
the modeled abundance
profiles for CCH,
We modelincreased the external
bychemicalfactors
1e−03
model of 50 and
1e−04C, C , CH,
chemical model
chemical model
1e−04
CS, 1e−04
SO2 , and HCl coincide with the respective ratran profiles 500
compared to the best-fit
model, but found that is has no eﬀect
1e+13
1e+14
1e+15
1e+13
1e+14
1e+15
1e−05
1e−06
1e−06
at someradii, the HCN and CH3 OH abundancesin the best-fit on the calculated abundances
because the enhanced interstellar
Radius [m]Radius [m]
1e−06
1e−08
model
are
lower
than
the
ratran-derived
values.
The
first
reaUV field cannot penetrate deeply into the massive AFGL 2591
1e−08
1e−07
1e−10
1e−08
son forthis disagreement is a fundamental restriction of the uti- envelope and the resulting
UV-irradiated volume is negligible in Fig. 5. The same as Fig. 2. Radiative transfer models are with a jump at 230 K (higher inner abundance).
1e−12
1e−10
1e−09
[SO2] withlized 1D approach with which one cannot properly capture the the 1D spherically-symmetriccase, unlike in the 2D model
1e−10
1e−14
[NH3]
physicalprocesses that occur on thewalls of UV- the outflow[SO](e.g., as described in Bruderer et al. 2010b).
1e−12
1e−11 and chemicalaction converting HNC to HCN by collisions with atomic H, 4.3. Complex organics
1e+13
1e+14
1e+15
irradiated outflow cavity. The presence of the outflow cavity andThus,
the 1D approach that is often employed to model and quantum-chemically calculated to be ∼1200 K, should in fact
1e+13
1e+14
1e+15
1e+13
1e+14
1e+15
Radius [m]even a disk-like rotating structure in the AFGL 2591 is well es- understand the chemical composition of lowand high-mass pro- be lower, ∼200 K. Alternatively, a high-temperature destruction
The problem with predicting CH3 OH abundances is that thistablished (see, e.g., Tamura & Yamashita (1992), Bruderer et al. tostellar objects appears inappropriate for PDR-like ions andchannel of HNC may be missing in the models, and one opradiative transfer model
1e−02 and Wang etradiative
1e−02
model
radiative transfer model
species is produced as ice on dust grain surfaces during the cold,(2010b),al. transfer(2012)).
Clearly, in thismolecules.
1e−02case such achemical model To explain their elevated abundances, a 2D or 3D tion is the reaction of HNC with O. The KIDA database gives a
1e−04
chemical model
chemical model
1e−04
complexa 2D or even model with a detailed descriptionof the outflow physics (central rate constant of ∼10−12 cm3 s−1 and a barrier of >1000 K, which pre-AFGL 2591 phase and is later injected into the gas phase
1e−06 environment is better characterized by
1e−04
1e−06
3D model.
1e−08
X-ray or FUV source irradiating
outflow cavity walls, outflow- would suggest that this reaction only plays a minor role in the as a result of heating. Because our 1D model does not include
1e−06
1e−10
a warm-up phase, but rather a simple two-stage cold-to-warm
1e−08etc.) must be utilized.
Asshown by Bruderer et al. (2010b), a multidimensional ap- induced shocks at the walls,
AFGL 2591 envelope. However, the reaction data are uncertain
1e−12
jump-like evolution, it probably does not allow methanol pre1e−08
1e−10
and
we
recommendanaccuratemeasurementof
thispotentiallyproachisrequiredtoexplaintheobserved
abundances
of
light
1e−14
cursor ices to be suﬃciently accumulated in grain mantles (see
+
[H
O]
1e−10
1e−12
1e−16
2
importantreaction.
Third,
thedynamical
andthermalevolution
hydrides and hydride ions, such as. As we do not include 4.2. HCN and HNC[HOH
Garrod et al. 2008). In addition, we model the pre-AFGL phase
[HCl]
2S]
1e−18
1e−14
of the environment may play a significant role in the HCN/HNC
a central
UV source in our model (since its radiation1e+13
in the 1D1e+14
1e+15
assuming 0D physical conditions, which may not be suﬃcient
1e−20
approach1e+13
is quickly1e+14
absorbed 1e+15
by the massive envelope inner re- RadiusOur[mmodelsare able to fit either
HNC abun- isomer-specific formation and destruction chemistry; our model for the 1D model used after that. It might also be that our ap1e+13 the HCN
1e+14 or the1e+15]
gions), our model tends to under-predict the abundances of sim- dance, but not both. This discrepancy reflects the limitations of does not consider this eﬀect as it is based on a static physical proach to model CO hydrogenation on the dust surfaces does
1e+00
ple photostableradicals
and
PDR-like
species. This may explain modern astrochemical models that are not capable of explaining structure of the AFGL 2591 envelope.
radiative
transfermodel
not include all the necessary details such as a multi-phase ice
chemical model
why1e−02
the abundances of species
like C, C+ , CH, CN, CCH, and why the observed HNC/HCN abundance ratio decreases with in1e−04
structure (chemically inert, bulk ice with chemically active ice
14
HCl1e−06
at small radii (<∼ 10 m) are lower in the model than the creasing temperature in regions of low- and high-mass star forlayers), porosity, and the presence of surface binding sites of
1e−08
ratran-derived
observed values. For HCl, which has a very lim- mation (see recent discussion in Graninger et al. (2014)). One
diﬀerent energies (see Taquet et al. (2012), Vasyunin & Herbst
ited 1e−10chemistry
network
in
our
model
of
only
58
channels,theexplanation
is
that
the
energybarrieroftheisomerization
re1e−124.1. Photostable species located in the outflow cavity• パラメータ依存性

10

1e−14
1e−16
1e−18

100K jump
内側でX大[CH3OH]
1e+13

1e+14

1e+15

Radius [m]Fig. 6. The same as Fig. 2. Radiative transfer models are with a jump at 100 K (higher inner abundance).

•

grain sizeには依らない - iceは溶けている為•

UVに依らない - denseな為•

ionization rateに依存 - 特にイオン分子•

initial C/O比に依存 - ほぼ全ての分子

11•chemical ageに依存 - 特に重いS,N系分子