[N II] 1.46 THz 輝線による
銀河面サーベイ梅本智文 (国立天文台野辺山)
南極天文ワークショップ（極地研）

2019年12月12日

l=198l=236

FUGIN Project(Umemoto et al. 2017 PASJ, 69,78)• Using mulL-beam receiver FOREST,OTF mapping of the GalacLc plane
12
13
18
in CO, CO, C O(J=1-0),
simultaneously

Mapping area:
inner disk: l = 10°〜50° |b|≦1°

l=50

Spiral arms, interarm, bar/barend,outer disk: l = 198°〜236° |b|≦1°Comparison with inner disk

l=10
NASA/R. Hurt

# 01 Molecular Cloud Identification
Identified Strunctures by Dendrogram (Fujita et al. in prep.)銀河面における高密度ガスの割合

•

度
密
高
る
！
き
％
で
３
生
か
誕
ず
が
わ
星
合
割
の
ス
ガ

5kpcにわたる領域での高密度ガスの割合 ~ 3%•

K.Torii et al. 2019, PASJ, 71,

渦状腕では ~ 5%に対しbar 領域や渦状腕の間では ~ 0.1-0.4%– これらの領域での星形成率の違いか?

•距離 ~ 5.5–8kpc⇄~1pcの分解能•

FUGIN データは分子雲の総質量を求め高密度ガスが集中する星形成領域を検出できる4

warm, neutral, and atomic hydrogen clouds (WNM). The WNM clouds are converted to cool
HI clouds (CNM) through a thermal instability and phase transition. These clouds are mostoften observed through the collisionally excited, hyperfine, spin-flip transition of hydrogenat 1420 MHz (λ ~ 21 cm). The clouds also radiate away energy by emission in the ionizedcarbon fine structure line ([CII]). At densities >105 cm−3, gas thermal energy can also be

Lifecycle of the ISM(N II)(C II)(C II)

Warm neutral
and ionized gas

Formation ofcold HI clouds

Disruption ofmolecular clouds(C II)(C II)

(C I)

(C I)

CO

CO

Formation ofmolecular clouds

Stellarevolution
Star formation

HD
CO

FIGURE 1.7(C II)O

Walker (2016)

Lifecycle of the ISM. Starting at the top, overlapping supernovae shock waves sweep up

Atomic and molecular lines at THz

Mauna Kea 1.5mm
Chanantor 0.60mm
Dome Fuji 0.14mm
Dome Fuji 0.10mm

Kuno 資料より

The hotter a star burns, the shorter will be its lifetime on the main sequence, τMS,

[N II] as Probe
M ⎞ SFR and LIR⎛ of(1.22)
≈ 10
τ
−2.5

MS

10

⎜⎝ M ⎟⎠⊙The hotter a star burns, the shorter will be its lifetime on the main sequence, τMS,

• [N II] emission line

−2.5

where
10 ⎛ M ⎞
≈
τ
10
MS
⎜⎝ M
⎟⎠ (SFR) andAn
excellent
probe
of
star
formationrate⊙τ = stellar lifetime on the main sequence (years).
MS•(1.22)

infrared dust luminosity (LIR) (Zhao et al. 2013) because
where
Massive stars
have
greatest
luminosities, arefrom
the largest
contributors
to the galactic
the [N theII]τ islesscontaminatedtheemissionofolder
MS = stellar lifetime on the main sequence (years).
UV flux, and haveshortest
Theirpresencehigher
is indicative
recent star-forming
starthedue
to anlifetimes.
ionization
potential
thanofhydrogen.
activity and can be tracedby theandatomicfine structure
emissionthey
genMassive
starsPDRs
have thegreatest
luminosities,
are the line
largestcontributors
to thegalactic
Therefore,Land
may
be awith
more
indicator
ofofSFRUV flux,
havethe shortest
lifetimes.
Their presence
is indicative
recent star-forming
[N205
II]
erate. In particular,
the
[NII]
μm
line,
an accurate
ionization
potentialjusthigherthanactivity and cantobebe
tracedby the PDRs
and atomic
structureline emission
they genhydrogen, has thebeenthedemonstrated
an excellent
probe
ofestimates.
thefinestar
formation
rate (SFR)
more conventional
LIR -derived
erate. In particular, the [NII] 205 μm line, with an ionization potential just higher thanand infrared dust luminosity(LIRbeen) (Zhao
et al., 2013),
hydrogen, hasdemonstrated to be an excellent probe of the star formation rate (SFR)•

and infrared dust luminosity (LIR) (Zhao et al., 2013),

log SFR = ( −5.31 ± 0.32) + (0.95 ± 0.05) log L[NII](1.23)

log LIR = (4.51 ± 0.32) + (0.95 ± 0.05) log L[NII](1.24)log SFR = ( −5.31 ± 0.32) + (0.95 ± 0.05) log L[NII]
log LIR = (4.51 ± 0.32) + (0.95 ± 0.05) log L[NII]

where
where
= star
SFR = star formationSFR
rate(M⊙formationyr −1) rate (M⊙ yr −1)[NII] = luminosity of [NII] line (L⊙)
L[NII] = luminosity ofL[NII]
line (L⊙)
LIR = luminosity of dust in IR (L⊙)
LIR = luminosity of dust in IR (L⊙)(1.23)(1.24)

mitting clouds in velocity space. A sample [CII] spectrum taken along one line ofgh the Milky Way with the heterodyne instrument for the far-infrared (HIFI) (seeaboard the 3.5 m Herschel space observatory is shown in Figure 1.12. As part of the

COBE/FIRAS maps of [C II] & [N II]
158µm (CII) emission• All sky survey with a special
resolution of 7 & a velocity
resolution of 1000km/s
3 25 70–1120• [C II] line is the dominantcooling line of the ISM at~0.3% of infrared continuum
205µm (NII) emission• [N II] & [C I] are less intense
by a factor of 10 & 100–0.2

0.3 2.5 7

12• [N II] line will appear atstrongly ionized regions -->by comparing [C II] &[N II],
determine if [C II] is arising
Fixsen, Bennett, & Mather (1999)
COBE FIRAS images of Milky Way [CII] 158 μm emission (top) and [NII] 205 μm emisfrom ionized or neutral gasm). Images were made at 7° angular and ~1000 km/s velocity resolution. Black swaths

Herschel [C II] Galactic plane survey• Galactic Observations of Terahertz C+

and N+ (right) line intensities. The low velocity resolutionsr(GOT C+)

h
t
!
m
y
s
n
e
+
s
r
e
r
d Figure 2:GOTC+ observations along line of sights• [C II] survey by Herschel(3.5m)/
HIFI with 12” angular resolution& 0.1 km/s velocity resolution• 452 LOSs volume-weighted
sample of the Galactic plane• Every 0.87 (|l|<60 ), 1.3 (30 <|l| <60 ), 4.5 (60 < |l| <90 ), and
4.5 to 13.5 (|l| >90 )• b=0 , 0.5 , and 1.0 , 2.0 (|l|>90 )resulted in an under-sampled survey (from Pineda etal 2010)Langer et al. (2010), Pineda et al. (2013), Langer et al. (2014)

2

8

2.0
!"#$%"&'("

G013.90+0.0

G000.0 +0.0

7

[C II] and [N II] by Herschel/HIFI
1.5

6

5

)*%&)"&'("

3

1.0• [N II] emission to the GOT C+
T(K)

T(K)

4

0.5

W. D. Langer et al.: [C ii] and [N ii] from d• Can separate the contributions from highly ionized and2

Table 2. [C ii] emission and absorption components.

1

weakly ionized gas to the [C II] emission
0.0

0

•

[C ii] sourcecomponent
Itot (⌧)(K km s 1 )
103.5
33.6
88.5
42.3

LOS
# T
I /I
I /I
T
V
[C II] emission arising from strong sources
of
[N
II]
emissionlabel(K)
observed corrected (K) (kmG031.3 3 6.3
0.93
0.56
2.9
1
3.7
0.55
0.34
1.5
5
V (km/s)
V (km/s)
is frequently
absorbed by foregroundgasG049.1--> 2underestimate
G316.6 1 5.1
0.94
0.58
1.4
–
G337.0 1 4.3
2.37
1.37
1.0
–1
Langer et al.: [C ii] and [N ii] from dense highly
ionized regions ionized gas by only [C II] observations
-1
-200

-150

s

-100

-50

0

50

100

150

-0.5
-40 -30 -20 -10 0

200

10 20 30 40 50 60 70 80 90 100
lsr

lsr

1.5

6 2.0

5
1.5

[C II]

5
1.0

4

1.0

3

T(K)

T(K)
T(K)

3

4

T(K)

)"&'("

0.5

2

ion tot2
0.5

3.1.3. F[N II]

1

1

0.0
0.0

0

50

100

150

200

o

[C ii] em[N ii] em
G316.6+0.0
region b
Gaussian ﬁt of [CtheII]fract
four soufrom [N II] emissioncles, an aabsorptio
have thereducedhas a fra
nations,[C ii] lin

6
G049.10+0.0

G031.3+0.0
G013.90+0.0

G000.0 +0.0

ion obs0-1-0.5
0 -40 -30
20 -20 -10
40 0

-0.5
60 20 30
80 40 50
100 60 70
120 80 90
14010010

V V(km/s)(km/s)

m/s)

lsr lsr

0

10

20

30

40

50

60

V (km/s)
lsr

70

80

90

100

-1
-100

-75

-50

-25

0

25

V (km/s)lsr

Assumin[C ii] ingroundItot ([C iifunctionof the reT s is theand T o iglect [Cis likely
re-emissT s and T
the abso
The colu

Langer, Goldsmith & Pineda (2016)

Fig. 1. Main beam temperature versus velocity for [C ii] (blue solid spectra) and [N ii] (red solid spectra) taken with HIFI – see Table Fig.
1 for6.some
A Gaussian fit (black dash) to the line wings of the Vlsr ⇠of the line parameters, and 13 CO(1 0) (black dashed line) from Pineda et al. (2013). Each LOS is labeled with the GOT C+ LOS label.
Some
47 km
s 1 [C ii] component (blue) in G316.6+0.0 using the Vlsr andfeatures that are indicative of [C ii] absorption are highlighted with arrows.
line width of the Gaussian fit (not shown) to the [N ii] emission (red).
1.5

The Astrophysical Journal, 814:133 (21pp), 2015 December 1

Goldsmith et al.
Table 3
(Continued)

Herschel [N II] Galactic plane survey
LOS LabelLongitudeOBSID[N II] 122 μm

The Astrophysical Journal, 814:133 (21pp), 2015RDecember
1(a )Intensity±(Err)[×10−8](W m−2 sr−1)[N II] 205 μm

σ[×10−8]R (a )

Intensity±(Err)[×10−8](W m−2 sr−1)n(e)±(Err)

N(N+)±(Err)(cm−3)

Goldsmith[×1016] et al.
(cm−2)

G344.8+0.0
344.7830
1342267185(2)
1.37±0.020
0.42(3)0.80±0.056
36.4±4.3
1.3±0.02
G345.7+0.0
345.6520
1342267869(2)
12.60±0.044
7.55(2)
5.44±0.081
59.6±1.57.7±0.03
G346.5+0.0
346.5220
1342267178(2)2.90±0.020
0.77(2)
1.71±0.059
36.0±2.1
2.8±0.02
G347.4+0.0
347.3910
1342264238(2)
1.92±0.013
0.54(3)
1.32±0.049
27.5±1.8
2.4±0.02
G348.3+0.0
348.2610
1342264237(2)
3.02±0.020
0.90(2)2.07±0.064
27.7±1.5
3.7±0.02
G349.1+0.0
349.1300
1342267870(2)
21.31±0.096
4.41(2)
8.67±0.057
65.7±0.812.1±0.05
G350.0+0.0
350.0000
1342265935(3)
0.36±0.008
0.23(1)<0.22±...
...±...
...±...
G350.9+0.0
350.8700
1342265684(2)
1.42±0.013
0.41(2)
1.22±0.069
18.1±2.1
2.8±0.03
G351.7+0.0
351.7390
1342265683(2)
1.61±0.025
0.48(3)
1.68±0.055
12.1±1.0
4.9±0.07
G352.6+0.0
352.6090
1342265682(2)
1.47±0.021
0.36(3)
1.22±0.055
19.6±1.8
2.6±0.04
G353.5+0.0
353.4780
1342265681(2)
3.68±0.027
0.62(2)
2.95±0.066
20.9±0.9
6.1±0.04
G354.3+0.0
354.3480
1342265680(2)2.33±0.020
0.65(2)
1.53±0.048
29.9±1.7
2.7±0.02
G355.2+0.0
355.2170
1342267874(2)
5.22±0.017
1.34(2)
1.87±0.060
80.5±4.2
2.6±0.01
G356.1+0.0
356.0870
1342265934(2)
2.59±0.020
0.31(3)
1.63±0.038
32.3±1.4
2.8±0.02
G357.0+0.0
356.9570
1342265933(3)
0.35±0.010
0.21(1)<0.18±...
...±...
...±...
G357.8+0.0
357.8260
1342265932(2)
0.82±0.020
0.31(1)<0.19±...
...±...
...±...
G358.7+0.0
358.6960
1342267875(0)...±...
...
(0)...±...
...±...
...±...
Figure 19. Observed
intensity1342265931
ratio of the two (2)[N II] ﬁne 17.30±0.102
structure lines and derived
are plotted as a functionof Galactic longitude.
G359.5+0.0
359.5000
1.69 electron(2)densities 9.16±0.095
42.9±0.814.1±0.08• [N II] at 122 & 205 µm with PACS(5x5)• 149 LOSs selected from GOT C+, 10”(122um),15”(205µm)• Bothe lines are detected in the range -60 ≤ l ≤60• [N II] emission highly correlated with that of [C II]• High electron density--> extended envelopes of H II regions,and low-ﬁlling factor high-density ﬂuctuations of WIM

Figure 5. HIFI [N II] spectrum before baseline removal (with positive
continuum offset; shown in red), with continuum offset removed (lower
spectrum; shown in blue), and with third order baseline removed (lower
spectrum; shown in green). There is only a minimal difference between the
spectra with zeroth and third order baselines removed in this case.
Figure 4. Illustration of effect of polynomial baseline removal from PACS
spectra toward G024.3+0.0. A DC offset had previously been removed to yield
Figure20. spectra
Comparisonof theasintensitiesof thein[Nthis
II] of 122 and 205 μm lines.
the 25displayed
the red lines
122 μm observation. The
Thehigher-order
blue squarespolynomial
are the LOS
in
the
range
−30°<l< 30°,and theﬁtting
triangles(spectra shown in black after
baselineandare removal)
the LOS modestly
positions reducesin the region–60°<l<−30°
and30°<l<60°.
any residual offset in the vicinity of the line as The
well
black
equal
intensities(1:1 ofratio),
andII]theblue line
emission.
as dashed
reducingline
thecorresponds
amplitude oftothe
featureto the rightthe [Nspaxels. It is thus reasonable to consider the [N II] emission to

Figure 21. Intensities of [N II] 122 and 205 μm transitions compared with thosebe spatially extended and fairly uniform in each pointingof [C II] 158 μm. The red and blue solid lines are power law ﬁts of thedirection (;1[Narcmin
in size). Given that the PACS footprint
II] lines vs. [C II]. The black dashed line is a slope of unity,
corresponding
size is 47″
corresponds
pc at a with
distancetransitions
and alsoto[N1.5
II] measured
HIFI asofwellindicating
thatandboth that[N II]1′

Goldsmith et al. (2015)

[N II] 1.46 THz Galactic plane survey• [N II] emission line

• [N II] line will appear at strongly ionized regions --> bycomparing [C II] &[N II] maps, it is possible to determine
whether [C II] is arising from ionized or neutral gas• [C II] emission is frequently absorbed by foreground gas -->•underestimate highly ionized gas by only [C II] observations
An excellent probe of star formation rate (SFR) andinfrared dust luminosity (LIR) (Zhao et al. 2013) because the[N II] is less contaminated from the emission of older star
due to an ionization potential higher than hydrogen• [N II] 1.46 THz Galactic plane survey S.P.
@

• Wide area mapping of the Milky Way --> evolutionary

process of interstellar medium not only neutral gas but
also ionized gas

Detection of 205µm [N II] from Ground• AST/RO at South Pole (1.7m)• The ﬁrst detection of 205µm [N II] line from ground-base• [N II] emission reveals the fraction of [C II] emission arises
from the ionized gas and the neutral ISM
L128

OBERST ET AL.

• 27% of [C II] arises from the low-density ionized gas but[N ii] IN CARINANEBULA
73% from the neutral
ISM!

No. 2, 2006

L127
4.1. The [N ii] Line Ratio[N II] emission from Carina nebular

Fig. 1.—[N ii] 205 mm spectrum obtained using SPIFI on AST/RO. These
data were obtained at the peak of the radio continuum contours of the Carina IIwith velocity
km/s
H ii region (Retallack1983). Velocityresolution
resolution is 60 kmof
s!1.60The dashedlineshows the intrinsic instrumental profile. The data were smoothed with a Hann

The [N ii] lines provide a distinct probe of the ionization
structure and gas density in H ii regions. Formation of O""
requires energetic 35.1 eV photons, so that the [O iii] lines trace
H ii regions formed by very early-type (O7) stars. Lower energy
photons (14.53 eV) can form N" so that these lines arise from
H ii regions formed by the softer UV radiation of late O-type
or early B-type stars, or from the lower ionization “outskirts” of
H ii regions formed by earlier type stars. Since O"" and N""
have similar ionization potentials (35.1 and 29.6 eV, respectively), an O"" zone contains mainly N"" rather than N". The
twin [N ii] lines have significantly lower critical densities(ncrit ∼ 44 and 293 cm!3 (8000 K) for the 205 and 122 mm lines,
respectively) for thermalization than the [O iii] lines (ncrit ∼510 and 3600 cm!3 (8000 K) for the 88 and 52 mm lines, respectively), so that they make better probes of lower density gas.
Unfortunately, the undersampled ISO LWS mapping of Mizutani et al. (2002) does not directly sample the position where
we obtained our spectrum displayed in Figure 1. The nearest pos
sition is at l p 287!. 550, b p !0!. 583 [R.A. p 10 h 44 m 54.1,′′′
decl. p !59!37 17 (J2000.0)], or about 1".27 (more than a fullbeam) offset from our peak line flux position. However, this position is covered by other pixels in the array, for which we obtainan integrated line flux of 21.4 K km s!1, or (3.6 " 0.7) #10!15 W m!2 in our 54! beam. Using the ISO data archive, we
obtain a 122 mm [N ii] line flux of 1.66 # 10!14 W m!2, or making
the extended source correction, a flux of 1.12 # 10!14 W m!2 in
Fig. detector
2.—[N ii]
122(Gry
mm–to–205mm lineintensity ratio as a function of ionized
the 78!.2 LW2 ISO LWS
beam
et al. 2003).
Assuming
gasdensity,n(solidline).
The
open
circleour observed line ratio ofa uniform intensity source, ande correcting for the differencemarks
in
1.5
!
0.64
.
The
ratio
of
the
[C
!15
!2ii] 158 mm to [N ii] 205 mm lines as a function
beam size, the ISO flux is 5.33 # 10 W m referred to ourof ne (dashed line). The filled circle marks the observed line ratio of 9.2 ! 3.9.

Oberst et al. (2006)

and dips to be
the upper J p
ulated at dens
From the I
10!13 W m!2.
responds to 5.
LW4 beam. A
our 54! beam
30% calibrati
205 mm line r
gas density of
we would exp
an upper bou
observed [C iremaining 73%
results supporof the observe
regions, the G
in warm densThis work
OPP-0338149
NNG05GK70
P. A. R. Ade(C. A. Allen,
for their exce
local oscillato
grateful to K.
the TREND l2.2 Specifications of 10‐m Telescop
Small Telescope Designed as Survey• Diameter:
Diameter：>
10m
3-6 m
•
•
Surfaceaccuracy：<
20μSurfaceaccuracy:<
20µm•
the detection of 205µm [N II]•• ForFrequency：200GHz‐1.3
line from ground-base, < 10µm?
•
Field
of
view：1°Wideﬁeldheterodyne•
camera• receiver
Pointing
accuracy：2″
For wide area mapping, 100-250•• multi-beam
Tracking
accuracy：0.5″
receiverAngular resolution
GHz

Angularresolution（10m）
3m
4m
6m

460

55”

850

29”

200GHzNANTEN II (4m)

37″ 17”

1500

41”

27”

22”

15”

800GHz

1.3T

9.3″8.5”

5.8

13”

l=221l=236（注 2: Confusion limit は Bla

l=198

Visible Sky（注 3: ホーンを使用予定）(2) スペクトル線観測（冬季 50%レベル＠新ドームふじ）
周波数帯

周波数範囲

感度(5σ

τ=60seMapping area:(220 GHz)
210-275
0.096 K
Inner disk: l = 221°〜25°|b|≦1° 0.081
350
275-373450

385-5000.15

850787-950

0.34

1000

1010-1060

0.86

1500

1440-1540

1.05

650
600-710G.C.0.21
Spiral arms, interarm,bar,Partially overlap
area
1300 FUGIN
1260-1380
0.92
232 < l < 14 (EL >30 ), 221 < l < 25 (EL>20 )＊ON 点（注 1: 感度は 1 τ ∆V に比例（注 2: 多素子受信機を搭載）

l=50（注 3: 分光電波カメラ[R~10〇観測可能天域（＠新ドームふじ）

l=25
l=10

仰角(EL)赤緯(Decl.)

> 5°

<＋8°

>10°

<＋3°

>20°

<－7°

NASA/R. Hurt

Survey with Small Telescope• [N II] 1.46THz Galactic plane survey• [N II] emission reveals the fraction of [C II] emission arises
from the ionized gas and the neutral ISM• [C I] 809GHz observation when the weather is not good
• Mapping area: l = 221 - 25 (EL>20 ), |b| ≤1 (|b| ≤2 )• 20” grid, ΔT(5σ)=0.76K, 10x10 beams ̶> 20,000h (OTF)>> more low noise receiver and more beams (or smaller D)
Freq. band
(GHz)

Freq. range
(GHz)

Lines

Beam

460

385-540

CO (J=4-3), [C I] 3P1-3P0

250

0.054 K

41.3”

650

575-735

HCl, D2H+

250

0.092 K

28.5”

850

775-965

CO (J=7-6), [C I] 3P2-3P1

250

0.14 K

21.7”

1000

1000-1060

CO (J=8-7), NH+

100

0.35 K

18.7’

1300

1250-1380

CO (J=11-10), H2D+

100

0.47 K

14.2”

1500

1450-1550[N II]

100

0.76 K*

12.7”

Sensi>vity (5σ)
Angular Resolu>on(τ=10min, Δv=1km/s)(D=4m)(*:Tsys=6000K) From Kunoʼs document

Survey with Small Telescope• [N II] 1.46THz Galactic plane survey• [N II] emission reveals the fraction of [C II] emission arises
from the ionized gas and the neutral ISM• [C I] 809GHz observation when the weather is not good
• Mapping area: l = 221 - 25 (EL>20 ) some strip scans• 20” grid, ΔT(5σ)=0.76K, 10x10 beams ̶> 55h (OTF)/strip>> some strip scans at latitude b
Freq. band
(GHz)

Freq. range
(GHz)

Lines

Beam

460

385-540

CO (J=4-3), [C I] 3P1-3P0

250

0.054 K

41.3”

650

575-735

HCl, D2H+

250

0.092 K

28.5”

850

775-965

CO (J=7-6), [C I] 3P2-3P1

250

0.14 K

21.7”

1000

1000-1060

CO (J=8-7), NH+

100

0.35 K

18.7’

1300

1250-1380

CO (J=11-10), H2D+

100

0.47 K

14.2”

1500

1450-1550[N II]

100

0.76 K*

12.7”

Sensi>vity (5σ)
Angular Resolu>on(τ=10min, Δv=1km/s)(D=4m)(*:Tsys=6000K) From Kunoʼs document

Survey with 30cm Telescope• [N II] 1.46THz Galactic plane survey
Main reflector• Mapping area: l = 221 - 25 (EL>20 ), |b| ≤1 (|b| ≤2 )• Angular resolution of 2.8ʼ at 1.46THz• 3ʼ grid, ΔT(5σ)=0.76K at 10min, 10 beams ̶> 3400h (6800h)

THz Survey Telescopes• FIRSPEX (~1m)

On Oct. 15, 2009 STO had its test flight from Ft. Sumner, NM. During its ~12 hours at float altitude (~126,000ft.) key components of STO were tested to help ensure the system would meet the objectives of the upcoming
science flight. STO consists of 3 major components; a gondola, an 80 cm telescope, and a THz heterodyne
receiver system. The gondola and telescope have been refurbished from the successful Flare Genesis
Experiment. The gondola was upgraded by APL to use 3 gyroscopes for inertial guidance and an optical tracker
for absolute pointing. The telescope was light-weighted and its primary and secondary mirrors re-aluminized. A
room temperature receiver system and a cryostat were constructed and flown on the test flight. A computercontrolled, sliding weight was added to dynamically compensate for cryogen evaporation during flight.

(Rigopoulou et al. 2015)• Small satellite from LEO by ESA and CAS(China)• [C II] 1.9THz, [NII] 1.46THz, [C I] 809GHz, CO(6-5) 690

2.0 TEST FLIGHT INSTRUMENTGHz• STO (80cm)

The STO flight from New Mexico was designed to provide a faithful test of the Antarctic gondola, telescope
and key elements of the science instrumentation.

(Walker et al. 2016)• The Stratospheric TeraHertz Observatory by Balloon• [C II] 1.9THz and [NII] 1.46THz at 1 arcmin. angularresolution• SOFIA(2.5m)/GREAT

Figure 1: STO test-flight telescope and instrument package configuration, as conceptualized(left) and being integrated at Fort Sumner in early October, 2009 (right).

(Young et al. 2012)

Antarctic ready versions of the telescope/gondola systems, instrument electronics, computing, control, and dat
storage systems were flown. To faithfully characterize the performance of the telescope drives and pointin
system, the test flight instrument was designed with the same CG and weight as the full Antarctica fligh
instrument. The STO test flight carried a test cryostat and an un-cooled Schottky receiver. The receiver wa
used to point on the 13CO J=3-2 line. The dewar was used to test the ability of a sliding weight to compensat
for the loss of cryogens during flight.

• GREAT: 60‒240µm (Heyminck et al. 2012)• 1.25~5 THz ([N II] 1.46THz, [C II] 1.90THz, [O I] 2.06,
4.74THz)

Figure 1 shows the STO telescope and instrument package as it was configured for the test flight. Th
instrument flight system as deployed contained the following subsystems:

Hα, [C II}], CO(3-2)

Ionized Gas in the Scutum Arm• [C II]

158µm & [N II]205µm obs. by SOFIA• Highly ionized gas of the WIM at the inner edge of the
Scutum arm tangency sampled along18 LOS (l=30.0-31.75 )• Strong [N II] emission throughout the Scutum tangency anddecreases
withlatitudea[Cscaleheight~55pc
W. D.exponentially
Langer et al.: Ionized gas in the
Scutum spiralarm as tracedwith
in [N ii] andii]
Primary Ref

Secondary Ref

Tertiary Ref

Fig. 13
4. The seven [C ii] spectra observed by upGREAT (black) and theFig. 5. Left: emission spectra of [C ii] (blue and black), [N ii] (red), central
and CO(green)for thetenobservedlines of sightthat aretowards
within b(l,=b)±0[N ii] spectrum(red)by GREAT=.25(30are.0, plotted
as T mb versus velocity. The central [C ii] pixel is plotted in blue and0theof thesevenshapes[C ii] pixels
black.theTheScutum
single and.0).average
While theoverallof the is[Cplottedii] linesinacrossspi-average[C ii] emission profiles are similar but not exactly identical owing tovariations
inemissionacross1 the footprint of the upGREAT array. The blackral arm (Vlsr = 60 to 130 km s ) within each pixel are similar, the small
vertical dashed line marks the tangent velocity of the Scutum spiraldi↵erencesarm. Right:amongsame asin theleft panel foreightis linesof sight
b > 0 .25.
them
demonstrate
thatthethere
structure
duewith
to inThe downward black arrows mark a few of the clearly evident [C iidividual] absorption
features.
cloud components. The [N ii] emission in the central pixel isFig. 5. Left: emission spectra of [C ii] (blue and black), [N ii] (red),detected
and 13 COacross(green)the tenlines ofsightofthatb = ±0arm.
.25 are plotted
theforentirevelocityrange[C iiare] inwithin
the Scutum
as T mb versus velocity. The central [C ii] pixel is plotted in blue and the average of the seven [C ii] pixels is plotted in black. The single and average
additionprofiles
to theare
SOFIA
and[N iiidentical] observations
4000 , inrespectively.
Thethespectralin thearray.
HI VGPSis[C iiIn] emission
similar[Cbutii]notexactlyowing towevariations
emission across
footprintresolutionof the upGREATThe black13
1
13
verticaldashedHlinemarksCO(1–0)
the tangent
velocityof the Scutum
spiralarm.
the leftpanel
for ofthe reference
eightandlinesinof
sight COwithb > 0 .25.
use
auxiliary
i and
data
from publicarchives,[Nii] Right:
1.56
km
s as inwith
rms
noise
2K
the
GRS
the
within0 .8) same
primary/secondary/tertiary
position.
The
1
The
downward
black
arrows
mark
a
few
of
the
clearlyevident[Cii]
absorption
features.
from
Herschel
PACS
and
HIFI
at
(30
.0,
0
.0)
and
(31
.2766,
0.0)
spectralresolution
is
0.21
km
s
with
rms
of
2
K.
Fig. 3. The 18 lines of sight observed in [C ii] and [N ii] are indicated by calibrated emission spectra at the reference positions were then
(Goldsmith
et al. 2015;
Langeret al. 2016),[CScutumii] with HIFI
white crosses
superimposedon a 21cm continuum
map andof theFig.incrementally
5. Left: emissionof [C ii] manner(blue andtoblack),[N ii] (red), and 13 CO (green) for th
addedin spectraa hierarchicalthe respective

Langer et al. (2017)

Ionized Gas in the Scutum Arm• There is highly ionized gas within the arm with 1-20 times
electron density of the interarm WIM• [N II] emission arises from shock compression layers of the A&A 607, A59
WIM, accelerated A&A
by607,thegravitational potential of the arm
A59 (2017)

W. D. Langer et al.: Ionizedgas inthe Scutum spiral armtangent
spectraScutum spiral arm tangency: l = 30.0! to 31.75!

to the molecular spiral arm are also consistent with 13 CO in the(l–VLSR ) map shown in Fig. 2.

GaussFit components
In Molecular Gas[CII] & [NII] Residual:
In WIM Gas

[CII]
[NII]
13CO

Fig. 6. Top: spectra for the atomic, molecular, and ionic gas tracers, H i

In Fig.
a function o
4. Discussion(left panel)
Our SOFIA observations provide a rich data set of velocity rebols indica
solved [C ii] and [N ii] emissions in the Scutum tangency. Our
VLSR Ë
includes th
results show widespread clear detections of [C ii] and [N ii]
1Ádistance from Galactic center
along all LOSs, especially, at velocities VLSR = 60 to 125 km s
function ofand at |b| < 0 .25. Although weaker at higher latitudes there is
GRS data
detectable emission at least up to b = 1 .3. The unique combination of [C ii] and [N ii] provides a powerful tool to analyze thenear b = 0 .
ionized gas component in spiral arms and the ability to isolate
strong 13 CO[C ii] emission in the neutral and ionized gases.
itude value
In general the WIM is a highly ionized, low density, n(e) ⇠3WIM: [C II] & [NII]
0.02 to 0.1 cm , high temperature, T k = 6000 to 10 000 K
Molecular gas layer:
It is nearly
Scutumgas that fills a 2–3
kpc layerinneraroundedge
the Galactic midplane (cf.
12CO, [C II] & [NII]
scale heighFerrière 2001; Cox 2005; Ha↵ner et al. 2009). It is estimated tocontain about 90% of the ionized gas in the Galaxy and has a fillvelocity ran
ing factor of order 0.2 to 0.4, which may depend on height above
molecular
the plane. At the midplane the filling factor may be as small
consistentas 10%. It has been studied with pulsar dispersion measurements (Cordes & Lazio 2003), H↵ emission (Ha↵ner et al. 2009,
scale heigh2016), and nitrogen emission lines in the visible (Reynolds et al.
the uncerta2001). More recently the WIM has been detected in absorption
with [N ii] (Persson et al. 2014). The [C ii] and [N ii] emission
have simila
lines can also be used to map the WIM, however, the low densithe tangenc
ties of the WIM make it difficult to detect their emission throughof the gas dout most of+
the Galaxy.
However,the spiral arm tangencies o↵er
Galactic
centera unique geometry that enables the detection of the WIM in [N ii]
sion fromFigt
and [CFig.
ii] because
they correspond
to ofregionswitha very spirallong arm tangency showing the2
1.
Schematicview
theScutum
9.inAa schematicradialcrossdispersion.
sectionviewof the spiral arm perpen- n(e) , andpathFig.
lengthrelatively narrowvelocity
The
WIM
l ⇠
Langeretal.
(2017)
structure
of
the
gas
layers(adapted
from
Fig.
6
in
Velusamyetal.
2015).
has dicularbeen detected
in
[C
ii
]
emissionbyVelusamyetal.
(2012,and[Nii]sto the LOS to Scutum spiral arm tangency showing the inner toand2015) The
alongemission
the tangencies
ofthePerseus,Norma,
Scutum,andof key gas tracers along the tangency, [C ii], [N ii], and CO

toward sun

[CII]
[NII]
13CO
HI/50

Summary• To understand the evolutionary process of interstellar
medium not only neutral gas but also ionized gas,
the observations of atomic lines at THz are criucial• [N II] is an excellent probe of star formation rate(SFR) and infrared dust luminosity (LIR), and [N II]
emission reveals the fraction of [C II] emission arises
from the ionized gas and the neutral ISM.

• So I would like to propose the [N II] 1.46THz Galactic
plane survey with the heterodyne receiver camera atSouth Pole. There is diﬃculty to survey entire theMilky Way, but it is good to observe the limited area
(strips) or to survey by 30cm telescope.