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beam size 01

kpc

0; Mpc10% 10$ 10# 10(10# 10 m (CI ?$ − (?%((CI?
−?$#10$0; Mpc
10% 10$ 10# 10(30 m

10%10&$10&#(ACIIArH D E = 2 − 1( I( I?(⁄# − ?$⁄#

10&' 10&( 10&# 10&$ 10%10&' 10&( 10&# 10&$ 10%
redshiftredshift

NH2+
OH NH
H2O 202–111
NHOH+
CO (9–8)
HCN (12–11)
H2O 312–303
H2O+ H OH+2O 111–000
2HCN (13–12)
CO (10–9)H2O 312–221
H2O 321–312

sub-mm –THz 20100

CO (13–12)
HCN (17–16)

CO (12–11)
H2O 523–514
HCN (16–15)
NH2
NII NH2NH2
OH+CO (8–7)

CO (7–6)
CI (2–1)
CH+

CO (6–5)
H2O+
H2O+
H2O 211–202

(a)

NH3
HF

Flux Jy

CO (4–3)
CI (1–0)
HCN (6–5)
HCO+/HOC+ 6–5(?)
CO (5–4)• Arp 220 (ULIRG)
Herschel, 17## − 40##
~ 6 − 15 kpc

HCN (15–14)

Rangwala et al.

HH22OO 422–413
220–211
HCN (14–13)
CO (11–10)

Astrophysical Journal, 743:94 (19pp), 2011 December 10(b)

(c)

(b)

500 600 700 800 900 1000
1200
1400
3 Arp 220. TheGHz
Figure 1. Herschel SPIRE-FTS spectrum of4567
spectrum shows the FIR continuum between 190 and 670µm in (a). Red solid points i(Rangwala+measurements from the SPIRE photometer and the dotted curves show the photometerbandpasses11)with arbitrary normalization. Line identifi
several molecular and atomic species in (b) and (c) with like colors for like species.
(A color version of this figure is available in the online journal.) • CO! • [CI] • [NII] • CH+•OH • H2O+
 • ArH+# • HF• H2O+

 '" &  
SFR%
AGN%
AGN%

AGN%

$  &  
CO ( )Kamenetzky et al.

68Upper J

10107

Sgr B2 (M)

10&SgrB2(M)NGC 6240

SgrB2(N)Mrk 231
Sgr A*UGC 05101

NGC 1266

108

Arp 299-A
NGC 253

Arp 220
NGC 1068

NGC 253
M82
M82

IRAS F17207-0014

M83

NGC 4038
NGC 1365-SW

NGC 1222

M83

Cen A
Cen A

GalacticCenter

SgrB2
Extended10&InnerGalaxy

Sgr B2 extended

2 4 6 8 10 12Upper @

12cm+,

Mrk 231

Mrk 273

2468
Upper J

1012S. All J = 1 → 0 luminosities are scaled to match that of Mrk 231 (3.7 ×105 L⊙ ). SLEDs are colored to indicate increasing LFIR
nt in the left or right panel is for clarity only. On the left panel, the SLEDs of the Galactic center (|l| < 2.◦ 5) and the Inner Galaxyed, are shown in green for comparison (Fixsen et al. 1999). On the right panel, we show the SLEDs of two star-forming cores andEtxaluze et al. 2013) and that of Sgr A* (Goicoechea et al. 2013). Because these SLEDs begin at J = 4 → 3, we scale J = 4 → 3 to
he shapes. None of the SLEDs are corrected for dust extinction.
vailable in the online journal.)' H%

Luminosity (I⨀ scaled to
Mrk 231 CO @ = 1 − 0)

95:174 (39pp), 2014 November 1010$10%coolwarm

10#10#

10%

10$-./0 K

10&(Kamenetzky+ 14; Kamenetzky+ 14)

Table 6
Photometry Flux Densities in Jy: SPIRE Photometer
250 µm Fν

σtot

350 µm Fνσtot

500 µm Fν

σtot

5.40

0.878

1.92

0.354

0.567

0.238

5070
60

[CI] T [K][CI] 3P1–3P0 & 3P2–3P140
5040

303020• à 13CO J=1–0!  (e.g., Ikeda+ 02)•  10#
$Dust Temperature [K]

80

20

1.0
1.5
2.0
2.5log CO Cold Component T [K]CO #!#[CI]7log [CI] LTE Mass [M ], 25 678 ,⨀Figure 22. Carbon, CO, and dust temperatures. Note that the y-axis and color
bar are linear. There is no apparent correlation between temperatures.
(A color version of this figure is available in the online journal.)

10 6
10%510&4
10'37 # 8 ( 9 ) 10*+ 11**log Molecular
Mass from10
CO [M ] 10
10 10
10,-./.,23 ,⨀Figure 23. C i and CO mass. The horizontal line represents the average(X[C i]/H2 ) in Section 4.5.

relationship (Figure 23). They adopted a relative abundance(Kamenetzky+ 14 ")
X[C i]/H2 = 3 × 10−5 . The weighted average abundance of
X[C i]/H2 = MC /(Mgas /1.4) × mH2 /mC ; using the gas massderived from the cold component CO fitting (and thereforedependent upon XCO ), we found log(X[C i]/H2 ) = −4.0 ± 0.5.
−5
In a linear ratio, this is 10+20−7 × 10 , a not particularly wellconstrained value but consistent with the values presented inOuCO,
exam
from
differgalax
dustusingbut stconclfor gmay
just r“thegas pnot recautio
speciis nohigh[C i]formaTh[C ii]
a var
descrof H2et al.
emissdark(∼4%
CO-dhaveour gin sta
the esline r
thanwithrange
2006of th(propdark

[CI] 3P1–3P0 & 3P2–3P1CO (M83)[CI] Publications of the Astronomical Society of Japan, (2014), Vol. 00, No. 010#

103
10$
10%

Michiyama et al.

&, '6 789 + 78: K km s +$ pc 3

Σ>?@ A⨀ pc +3

The Astrophysical Journal Letters, 897:L19 (7pp), 2020 July 1

11104

10%

$103$3;105

10 10 10 10< '6<'( )*$+%10#10$%,+$3+$of Σ as a function of I +$ , where the dataFig. 8. (a) Plotswith I<K4 σ km
are represented
by gray points, a&spc'( )*$+%
Krepresentedkm sin the same mannerKaskm
sin Figure 6. The best-fit relation with a fixed slope of unity to the data with I>4Figure 3. (a) L [C I](1 - 0) L CO(4 - 3) vs. L [C I](1 - 0) LFIR for PDRT calculations and observations. The red solid and blue dashed lines indicate the theoretical tracks for−3points indicate the nearby U/LIRG sample
by SPIRE/FTS (Kamenetzky et al. 2016). The yellow
constant nH (cmmol) and Uuv [G0], respectively. The[ C Iblack](1−0)[ CdetectedI](1−0)
star shows the [C I](1–0) upper limit of NGC6052 presented in this project. L [C I](1 - 0) L TIR) is calculated without any aperture correction, indicating that this ﬁgureWe
does not constrain the Uu,v . (b) [C I](1–0) line luminosity vs. CO(1–0) line luminosity for U/LIRG galaxies. The symbols are the same comparison sample as[ C(a).
I](1−0)note that CO(1–0) luminosity is obtained
the upper limit of [C I](1–0) luminosity is obtained by array observation (the missing ﬂux is not−2 by single dish,−1but−1
correspondingcorrected).
to α[ C I] ∼ 3.8 [M⊙ pc(K km s ) ]. The dashed line is the mean value of the local galaxies (7.3 [M⊙ pc−(Miyamoto+ 21 Michiyama+ 20)by Crocker et al. (2019). In panels (b) and (c), the radial variation of α[ C I] is plotted, but in (c), the data are separated into tw
Tdust = 23 K.characterize
The open theboxes
in panelvalues binned
a widthpredominantlyof r = 10′′from. The
upwarddownward tri
PDR.
The (b)ratiorepresent
suggests theUuv=4–5.
likelywith
to originate
extended,lowandexcitation′′(Díaz-Santosal. 2017).
gas.
On theother
hand, higher-J d)
COPlots
emission
is likely
the values binned
with a etwidthof r = 10 as well, but for Tdust > 23K and≤ 23K, respectively.
of Σmolas a functio
dominated by compact, high excitation
gasandthecontribution−2−1 −1
relation with a fixed slope of unity is shown by the solid line, corresponding
to α diffuse∼ 0.5[M⊙ pc(K km s ) ]. The dashe
from more CO
gas. Thus,
the [C I](1−0)/CO(4–3) and4.2. Molecular−2
Gas Mass −1 −1
may representof very
[C I](1–0)/CO(1–0)
the averaged value α CO ∼ 3.1 [M⊙ pc(K km s ) ] in nearby galaxies
measuredratios
by Sandstrom
et properties
al. (2013).
Panels (e) a[CI] 3P1–3P0 & 3P2–3P1•  #
 NGC 1808 (  + 21)
 M83 (Miyamoto+ 21)
 NGC 6052 (Michiyama+ 20)
à CO J=1–0 ( HI)
! "(Michiyama+ 20)  (1)A&A 587, A45 (20*,-.&&Σ)*+/⨀ yr $% kpc $&• [NII] (1461 GHz)
• 
"
$SFR (• SF#   (Hughes+ 16)
à SFR % '!(?)

we drepre
relati10#10

log10$%

whic
andpreviresid
A sim
NGC
4.38
cavea
ple siestim

10$&
10$'

10$6 10$7 10$8
In th$& on the$%Fig. 3. The ⌃9 – Fcorrelation
superimposedpredictions emisWmsr:;; &#<from the Starburst99 model. For varying values of ", the dashed lines of-sig

5. CSFR[NII]205

correspond to a stellar population with a constant SFR and td = 2 Myr,and solid lines correspond to td 20 Myr. The dash-dotted line is thebest-fitting linear relation from the M 51 and NGC 891 data.

(Hughes+ 16&)Simply combining all the data to define an average relation isprevented by the aforementioned o↵set in the individual galaxyand Sand M4039
alsoappethe ntion i• ArH+ (618 GHz, 1235 GHz)

H' fraction abundance to H  (2)10#$/

• 
10#%&
 10#%'• H2O (620 GHz)10&
PKS 1830-211

1.00.90.8

H2O+
38ArH+H2O+

O3 atom
36ArH+−2000200< = 0.88582  km s #%Ar+ArH+10#'
10#(10#) &10 10#' 10#( 10#)78 from surface mag(Müller+15  Schilke+ 14)  (3)0123 K• HF (1232 GHz)
• HF/H2 ~3.6×10()•  The Astrophysical Journal, 785:22 (5pp), 2014 April 10

Flux densitymJy⁄beam
6 7 ⁄ 68

Flux⁄Continuum

No. 1, 2005

1.2
1.0emission
0.8 absorption
0.601.0HI
0.8−100
0.6 absorption
−200
0.4
−300
−200 0 200 400
Velocity km s (/Figure 2. Upper panel: spectrum of the ground-state transition of HF J = 1–0(black) normalized by the corresponding continuum toward NGC 253 and theGaussian fit (green) of the emission component. Lower panel: HF normalized
absorption profile (black) obtained after subtracting the Gaussian line. H iabsorption spectrum (blue, in LSR velocities) toward the central continuum
source from Koribalski et al. (1995) is also shown for comparison, the conversion
between heliocentric velocities and LSR velocities is ν LSR = ν HSR – 2.95 km s−1 .
(A color version of this figure is available in the online journal.)

3.1. Column Densities of the Absorbing HF Gas

8060
40201010

10INTERSTELLAR FLUORIN(.

Monje et al.

C+(/-H2C0F0CO
HF

249: mag

CF+6Fig. 2.—Gas temperature and abundances in a one-sided PDR with nH ¼102 cm#3 and "UV ¼ 1.

Figure 3. Normalized spectrum of the ground-state transition of HF J = 1–0(black) with respect to the continuum
towardNGCtwo Gaussianfits F/O abundances were inferredsightline
to 4945
HD and208440.
Similar(dashed green lines) center at 543 km s−1 and 637 km s−1 with δV(FWHM) =fromFUSE observationsof F i in HD 209339A and from the
47 km s−1 and δV(FWHM) = 40 km
s−1 , respectively.
H i absorption spectrum(blue, in LSR velocities) toward thecentral continuum
sourceOtt et al. ! Sco (by Copernicus; Snow &original
discovery
offromF i toward(2001) is shown for comparison. York 1981). The assumed gas-phase fluorine abundance cor(A color version of this figure is available in the online journal.)(Monje+ 14  Neufeld+ 05)

responds to roughly 0.6 times the solar abundance of fluorine.
possible
fluorine
continuum originating fromThethe nucleus
andeffects
movingofawayfrom depletion in dense regions areconsideredinx
3.4
below.
it and toward us along the line of sight. The redshifted emission"UcurandfacexH2noalw  (4)• CH (532/536 GHz, 1471/1477 GHz)• ' •  ! CH⁄CO ~2×10+,
à AGN  #&• OH+ (909 GHz, 971 GHz, 1033 GHz)
• H2O+ (1115 GHz, 1139 GHz)
• H+%)" !$   %
• AGN *(

The Astrophysical Journal, 779:75 (12pp), 2013 December 10

Scoville & Murchikova



@A
implies an H ii region emission measuren n vol)of2220#$% 30(EM = 24
EM = 4.60×10 cm (using α = 2.6×10 cm s at T =7 ergs cm#& s #(79:75 (12pp), 2013 DecemberThe10 Astrophysical Journal, 779:75 (12pp), 2013 December 101062−3−13

B

Scoville & Murchikova

e p3 −1HI>

10 K
He i has an ionization#&(potential of 24.6 eV, 7500
and its photoionK
2.2. He Line Emission

10

Scoville & Murchikova

e104 K). Using the specific emissivity of 3 × 10−31 erg cm−3 s−1
for H i-26α from Figure 2, one finds the recombination line
luminosity will be LH26α = 1.38 × 1032 erg s−1 . For a source
distance of 1 Mpc and a line width of 30 km s−1 , this corresponds
to a peak line flux density of ∼3.3 mJy. This flux density isreadily detectable at a signal-to-noise ratio of 10σ within ∼1 hrwith ALMA Cycle 1 sensitivity.

10#&'@A50 41 37 34 32

HeII
10> K7500 K

7C,CEF ⁄7C
10> K, 10> cm#&•  
• He/H 5000 K

17500 K
10000 K
15000 K
20000 Kization requirements are not very different than those of H i.
Thus, the He i recombination lines probe the ionizing UV radiation field in much the same way as H i (see Section 3). Since theHe i submm lines will be weaker than those of H i because of
the lower He abundance, we do not examine the He i emission
extensively here and instead focus on He ii.
Theforionization
potential of He ii is 54.4 eV, corresponding to
α recombination lines are shown
T =er unit ne np , and the x-axisphotons
has the lower
with λ = 228
Å for conversion of He+ to He++ . Since
Figure 2. Emissivities of the H i submm recombination linesFigure
are shownfor
3. Emissivitiesof the He ii submm recombination lines are shown foritions. The emissivities werethe
calculatedfor
4 K.havemost massivestarinan andSB10will
surface temperatures
T = 7500
The emissivity
is given per unit ne np .TThetop border= 7500and 104 K. For He ii, it is per unit ne nHe++ . The top border has theand the curves for all densities are coincident.
Figure4.
Emissivityof He ii to H ifrequencies
are shown
as a function of principal quantum number n for ne = 102 –108 and Te = 7500, 10,000, and 15,000 K. The∼50,000 K, the ionizing
fromsuchofathe
population
will
has the EUVquantumnumberα transitions,andhave
their rest frequencies
are on of the αratios
quantum numbertransitions, and their rest
are on the bottom
4 K and density
4 cm−3 , so one can see the dependence on ion, temperature, and density. The line ratios are very nearly
emissivities
arenormalized
toHi
at
10
10−3
thebottomaxis.
Theemissivities
werecalculatedfor
n=10,1000,andforaxis. The emissivities were calculated for ne = 10, 1000, and 104 cm, andonly a very small fractionof the photons with energy sufficient to e−4/3
cm−3, and the curves forall ofdensities
104the+++the curves fordensities are
Note that the emissivities
for Heperociated nonlinear amplification(seeproduce He. Thus,
recombinationlines
He ii are(Hecoincident.
) that independent
ofalldensity
forcoincident.
all the temperatures,
but theydoii dependon temperature as Te
for quantum number n ∼ 10–50. The latter is clearly shown in the++unitnn
are
4–5
times
greaterthanthoseforHiat
similarfrequency—this(A color version ofofthis
online journal.)e He4−3
bmm H i lines (and theare free–freeproduced by recombinatione figure+ Heis++available
can bein athestrong
lower-rightpanel where
theloweremissivity
ratiorelative
is shown
fivetemperatures at ne = 10 cm .
partially compensates
for theHe abundance
to H iffor
theirionized
ly thin, their line fluxesdiscriminant
are a linear for the existence of an AGN with a relatively
volumesare similar
be thefor a very(A colorversionof (asthiswould
figureis caseavailable
in hardthe ionizing
online continuum).
journal.)mission measure (EM =ne nEUV–X-ray
hardcontinuum.
In SBs, there
can besome andHe iiStorey (A&color
are from Hummer& Storey(1987)
Hummerp vol).
version of this figure is available in the online journal.)cellent probe of the EUV
luminosity
emissionassociated(1995a,with Wolf–Rayetthe EMpopulation transfer1995b). Thestars.
latterHowever,work includes+
ssuming the EUV photons
notregion relativeof theareHe++
to thatandof theregion will
be hasmuchemissivities for H iby electron
ionHcollisions
andopticaldepths,
and that
hence
maser
amplification
galaxies. We have therefore rescaled the optical depths toclear
fromthese figuresincreasingTe leads
to a decrease inof the linest). Lastly, we note thatlessin thanvirtually
for an AGN.
and He ii up to principal quantum number n = 50.
Theyalso
theofHeii/Hratio. In Figure(lower right), the ratio
emission.
Ini emissivity
such instances,the 44 recombination
lineintensity∆vFWHM = 100 km s−1 . We have also scaled the optical depth
ction of the recombination
atThelines
emissivitiesof the Heii recombinationlines arecalculate
opticaldepth parameters
for atakenwide range
electron−3isshownfora
singledensityn=10cm,
butT
=5000
be insignificant.
from Storey & Hummer(1995a,For thedensities
interested(ne ).
temperatures(Te1995b).
) and electron
Wenot
makeusewill
accuratelyreflect thee ionized gas EMe and the associatedto a specific optical depth τ per unit ne nion Lpc , where Lpc is theto20,000 theK and the temperature dependence is clear and themm wave recombinationlines awere
reader,simple modelfor thescaling resultsof rate incoefficientsof thesenumerical
this paper;be-Figure
1showsLyman forcontinuum
emissionof the stellar
population. For
path length in parsecs and the volume densities in cm−3 .
all n transitions.
Thus,rates
the temperatureand density
H ii—in these regions thecontinuum
tweenhydrogen andhydrogenic
ions isdeveloped
analytically linesame
Storey& Hummer(1995b)
H i recombination
emissivities
ii toiiHlines,
i line ratios
at fixednα canthe
be possibility
thedependence
submm Hofi theandHeHe
we cananalyze
Figures 5 and 6 show the specific optical depths for the H i410#&$200 500 800 200 500 800
frequency GHz

0 10 20 30 40 50@A(Scoville & Murchikova 13)

20 K!" W m%& Hz %) rad%&.0 0 5==0• )#!+/ K001K• 4&, 0001%)/3*%10K
à 2010%)015.
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
.
%&&
1
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• 1 -
10%&2 )
1010& 103 102• )$0."
frequency GHz• ,"7 %3(-[CI] + $!9)/ + CO J=1–0 + HI 15• #10. …
$!9)/' (4)36ArH+, HF0+
[CI]&'8  26 *&