Luminous Infrared Galaxies

FIRSEDApril 3, 2013

FIRSEDApril 3, 2013

• IRAS-detected galaxies with f60μm > 5.24 Jy & |b| > 5o

• 201 galaxies with LIR ≥ 1011.0 L⊙

• HST, Chandra ~ 88 LIRGs; Spitzer, Galex, Herschel ~ 201 LIRGs

• http://goals.ipac.caltech.edu

The Great Observatories All-sky LIRG Survey (GOALS)

Luminous & Ultraluminous Infrared Galaxies

• LIRGs: LIR[8-1000 μm] ≥ 1011-11.99 L⊙; ULIRGs: LIR ≥ 1012 L⨀

• Spiral galaxies which show an increasing tendency to be involved in interactions or mergers with increasing luminosity

• Rich in optically-visible star clusters

NGC 5257/8

• Spiral galaxies which show an increasing tendency to be involved in interactions or mergers with increasing luminosity

• Rich in optically-visible star clusters

• Rich in dust, and both neutral and star-forming molecular gas

Contours = HI

CO(1→0)

CO(1→0)

Luminous & Ultraluminous Infrared Galaxies

Contours = HI

• Spiral galaxies which show an increasing tendency to be involved in interactions or mergers with increasing luminosity

• Rich in optically-visible star clusters

• Rich in dust, and both neutral and star-forming molecular gas

• Evidence of Active Galactic Nuclei (AGN) activity

AGN

Mrk 231

Luminous & Ultraluminous Infrared Galaxies

March 14, 1998 9:20 Annual Reviews SANDER2 AR12-18

756 SANDERS & MIRABEL

⊗⊗

–2

–4

–6

–8

–10

log Φ

(M

pc–3

Mbo

l–1)

9 10 11 12 13log Lbol (L ).

⊗

IRAS BGSIRAS 1Jy ULIGsMrk SeyfertsMrk StarburstcDsPG QSOsNormal galaxies

Figure 1 The luminosity function for infrared galaxies compared with other extragalactic objects.References: IRAS RBGS (Sanders et al 1996a), IRAS 1-Jy Survey of ULIGs (Kim 1995), Palomar-Green QSOs (Schmidt & Green 1983), Markarian starbursts and Seyfert galaxies (Huchra 1977),and normal galaxies (Schechter 1976). Determination of the bolometric luminosity for the opticallyselected samples was as described in Soifer et al (1986), except for the adoption of a more accuratebolometric correction for QSOs of 11.8⇥ ⌫L⌫(0.43 µm) (Elvis et al 1994).

comes from ULIGs at flux levels f60 = 1–2 Jy corresponding to sources at z >⇠

0.13. No evidence for evolution is found for the subsample of 2 Jy ULIGs, i.e.n = 3.8± 3 (Kim & Sanders 1996). These results appear to be consistent withprevious debates in the literature which find n ⇠ 5.6–7 for redshift surveys withflux limits f60 ⇠ 0.5 Jy (Saunders et al 1990, Oliver et al 1995) but only n ⇠ 3–4for surveys with flux limits f60 >

⇠1.5 Jy (Fisher et al 1992), and with analyses

of IRAS extragalactic source counts (Hacking et al 1987, Lonsdale & Hacking1989, Lonsdale et al 1990, Gregorich et al 1995) that show evidence for strongevolution only at relatively low flux levels ( f60 <

⇠1 Jy). More definitive tests of

U/LIRGs and luminosity function

•The local space density of LIRG is low• and overtakes that of normal galaxies beyond a Lbol ~ 1011.4 L⦿

U/LIRGs and the Main Sequence

• The high end luminosity of the LIRG population lie above the main sequence with SFR ~ 100s...

• ... if the IR is tracing only star formation

Main Sequence

Starbursts

♁ LIR > 1011.4 L⊙

+ LIR ≤ 1011.4 L⊙

Ultraluminous IR Galaxies & the evolution of QSOs?

Gas-rich Disk Progenitors with seed SMBH

(Sanders et al. 1988)

(Sanders et al. 1988)

Enhanced Star Formation;SMBH building

Ultraluminous IR Galaxies & the evolution of QSOs?

Optically luminousQSO phase

(Sanders et al. 1988)

Ultraluminous IR Galaxies & the evolution of QSOs?

Elliptical/SO Galaxy

(e.g., Barnes & Hernquist 1992)

Evolutionary byproduct

Gas-richProgenitors

EllipticalByproduct

Starburst/AGN;Dust-enshrouded

phase; cool IR colors

Dust-clearing;AGN-dominant

phase;warm IR colors

SED evolution: Cool → Warm

(f25μm/f60μm < 0.20) (f25μm/f60μm ≥ 0.20)

Dust-enshroudedphase

Dust-clearing phase

(Sanders & Mirabel 1996): see Vivian U’s talk for an update

The importance of IR for LIRG studies

• νIRLIR / νoptLopt ↑ as Lbol ↑

– 29 –

Fig. 2.— Left: Histogram showing the ratio of SFR(UV) to SFR(UV+IR). The solid line is the full GOALSGALEX sample. Colored lines show the GOALS GALEX sample divided into luminosity bins as in Fig. 1:1011 < LIR < 1011.4L! (red dotted line), 1011.4 < LIR < 1011.8L! (green dashed line), and LIR > 1011.8L!

(blue dashed line). The FUV contribution to SFR is small for (U)LIRGs and decreases with increasing LIR.Right: LIR plotted against the ratio of SFR(UV) to SFR(UV+IR). Median ratios of the star formationrates are shown for each luminosity bin (red: 1011 < LIR < 1011.4L!, green: 1011.4 < LIR < 1011.8L!, blue:LIR > 1011.8L!) along with 1! standard deviations of the mean. Although anticorrelated (Spearman rankcorrelation coe!cient of -0.47) the correlation is not linear.

The importance of IR for LIRG studies

(Howell et al. 2010)

•The SFR(UV) vs. SFR(IR) discrepancy grows with increasing LIR

HST Far-UV

(Inami et al. 2010)

5 kpc

• Unobscured massive star formation

An imaging census of light from a LIRG

HST Optical• Unobscured star

formation

(Inami et al. 2010)

5 kpc

HST Optical• Unobscured star

formation

(Kim et al. 2013)

5 kpc

HST NIR• embedded NRG

sources + old stellar population

(Inami et al. 2010)

5 kpc

1.6μm

70 INAMI ET AL. Vol. 140

5 10 15 20 25 30Rest wavelength (microns)

0.01

0.10

1.00

10.00

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL + LL

6 8 10 12 14Rest wavelength (microns)

0.0

0.1

0.2

0.3

0.4

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL only

PAH

[SIV]

[NeII]

Figure 6. Spitzer IRS spectra of source D (see Figure 3). The full low-resolution spectrum is at the top left, while a close-up of the SL spectrum is shown at the bottomleft. The SH and LH spectra are shown at the top right and the bottom right, respectively. Some residual order tilting is evident in the two reddest orders of LH. Allspectra are shown in the rest frame.

originates from source C+D. When this is done, we derive aratio of FPAHtot/FIR ! 0.02. Since the continuum emission atthe overlap wavelengths of SL and LL has a difference less than10%, dust emission in our low-res spectrum is dominated bysource C+D.

3.5. AKARI Near-infrared Spectroscopy

The spectrum of sources C and D taken with AKARI iswell separated from that of the other sources (A and B) inII Zw 096, although they are all in the Np aperture, becausethe direction of spectral dispersion is nearly perpendicular tothe vector separating C, D, and A. The AKARI near-infraredspectrum of source C+D is shown in Figure 7. Strong 3.3 µmPAH emission, 3.4 µm emission from aliphatic hydrocarbongrains, and Br! emission lines are detected. The line fluxes are(2.37± 0.14) " 10#16 W m#2 and (0.50± 0.05) " 10#16 W m#2

for the PAH and the Br!, respectively, after fitting both witha Gaussian function (see Tables 3 and 4). The EQW of the3.3 µm PAH is 0.14 µm. Because the AKARI spectral resolutionbetween 2.5 and 5 µm is only R ! 120, the Br! recombinationline is not resolved.

The spectrum of source A (not shown) also exhibits 3.3 µmPAH and Br! emission. The 3.3 µm PAH has a flux of(1.77 ± 0.12) " 10#16 W m#2 and an EQW of 0.08 µm. TheBr! line flux is (0.33 ± 0.06) " 10#16 W m#2.

3.6. Bright Star Clusters in the Remnant

The high-resolution HST imaging data show that a largenumber of bright star clusters surround the II Zw 096 nucleiand are present throughout the merging disks (see Figure 2).From the B- and I-band images, we identify 128 clusters withapparent magnitudes of 18.3 mag ! mB ! 26.7 mag (Vega)and 17.9 mag ! mI ! 25.7 mag (Vega) with 80% completenesslimits of 26.0 mag and 25.5 mag, respectively. Of these, 97 and88 are detected in the FUV and H-band images, respectively.

Figure 7. AKARI near-infrared spectrum of sources C and D. The wavelengthscale is given in the rest frame.

For all clusters, we use a 0.$$1 radius circular aperture for theHST photometric measurements. We estimate that the clusterscontribute % 14% of the total FUV luminosity (& 8 " 1036 W)in II Zw 096 (uncorrected for reddening).

In Figure 8, we present an FUV–optical (F140LP#F435Wversus F435W#F814W) color–color diagram of the clusters.We additionally show the colors of a number of diffuse emissionregions throughout the merger system. For the diffuse emissionregions, we use 0.$$5–0.$$8 radius circular apertures, depending onthe location in the FUV image to avoid crowded regions.

Two potential explanations for the observed color distributionare considered. First, if all the clusters are coeval, then theirdifferent colors result from variable extinction. In other words,those clusters or diffuse emission regions which have colors ofF435W#F814W " 1.0 mag or F140LP#F435W " 0.0 magare heavily reddened (AV " 3 mag). When we de-redden an

Spitzer mid-IR• Embedded Star

Formation & AGN • Dust emission

(Inami et al. 2010)

5 kpc

70 INAMI ET AL. Vol. 140

5 10 15 20 25 30Rest wavelength (microns)

0.01

0.10

1.00

10.00

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL + LL

6 8 10 12 14Rest wavelength (microns)

0.0

0.1

0.2

0.3

0.4

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL only

PAH

[SIV]

[NeII]

Figure 6. Spitzer IRS spectra of source D (see Figure 3). The full low-resolution spectrum is at the top left, while a close-up of the SL spectrum is shown at the bottomleft. The SH and LH spectra are shown at the top right and the bottom right, respectively. Some residual order tilting is evident in the two reddest orders of LH. Allspectra are shown in the rest frame.

originates from source C+D. When this is done, we derive aratio of FPAHtot/FIR ! 0.02. Since the continuum emission atthe overlap wavelengths of SL and LL has a difference less than10%, dust emission in our low-res spectrum is dominated bysource C+D.

3.5. AKARI Near-infrared Spectroscopy

The spectrum of sources C and D taken with AKARI iswell separated from that of the other sources (A and B) inII Zw 096, although they are all in the Np aperture, becausethe direction of spectral dispersion is nearly perpendicular tothe vector separating C, D, and A. The AKARI near-infraredspectrum of source C+D is shown in Figure 7. Strong 3.3 µmPAH emission, 3.4 µm emission from aliphatic hydrocarbongrains, and Br! emission lines are detected. The line fluxes are(2.37± 0.14) " 10#16 W m#2 and (0.50± 0.05) " 10#16 W m#2

for the PAH and the Br!, respectively, after fitting both witha Gaussian function (see Tables 3 and 4). The EQW of the3.3 µm PAH is 0.14 µm. Because the AKARI spectral resolutionbetween 2.5 and 5 µm is only R ! 120, the Br! recombinationline is not resolved.

The spectrum of source A (not shown) also exhibits 3.3 µmPAH and Br! emission. The 3.3 µm PAH has a flux of(1.77 ± 0.12) " 10#16 W m#2 and an EQW of 0.08 µm. TheBr! line flux is (0.33 ± 0.06) " 10#16 W m#2.

3.6. Bright Star Clusters in the Remnant

The high-resolution HST imaging data show that a largenumber of bright star clusters surround the II Zw 096 nucleiand are present throughout the merging disks (see Figure 2).From the B- and I-band images, we identify 128 clusters withapparent magnitudes of 18.3 mag ! mB ! 26.7 mag (Vega)and 17.9 mag ! mI ! 25.7 mag (Vega) with 80% completenesslimits of 26.0 mag and 25.5 mag, respectively. Of these, 97 and88 are detected in the FUV and H-band images, respectively.

Figure 7. AKARI near-infrared spectrum of sources C and D. The wavelengthscale is given in the rest frame.

For all clusters, we use a 0.$$1 radius circular aperture for theHST photometric measurements. We estimate that the clusterscontribute % 14% of the total FUV luminosity (& 8 " 1036 W)in II Zw 096 (uncorrected for reddening).

In Figure 8, we present an FUV–optical (F140LP#F435Wversus F435W#F814W) color–color diagram of the clusters.We additionally show the colors of a number of diffuse emissionregions throughout the merger system. For the diffuse emissionregions, we use 0.$$5–0.$$8 radius circular apertures, depending onthe location in the FUV image to avoid crowded regions.

Two potential explanations for the observed color distributionare considered. First, if all the clusters are coeval, then theirdifferent colors result from variable extinction. In other words,those clusters or diffuse emission regions which have colors ofF435W#F814W " 1.0 mag or F140LP#F435W " 0.0 magare heavily reddened (AV " 3 mag). When we de-redden an

70 INAMI ET AL. Vol. 140

5 10 15 20 25 30Rest wavelength (microns)

0.01

0.10

1.00

10.00

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL + LL

6 8 10 12 14Rest wavelength (microns)

0.0

0.1

0.2

0.3

0.4

Flux

den

sity

(Jy)

IIZw096 source D - IRS SL only

PAH

[SIV]

[NeII]

Figure 6. Spitzer IRS spectra of source D (see Figure 3). The full low-resolution spectrum is at the top left, while a close-up of the SL spectrum is shown at the bottomleft. The SH and LH spectra are shown at the top right and the bottom right, respectively. Some residual order tilting is evident in the two reddest orders of LH. Allspectra are shown in the rest frame.

originates from source C+D. When this is done, we derive aratio of FPAHtot/FIR ! 0.02. Since the continuum emission atthe overlap wavelengths of SL and LL has a difference less than10%, dust emission in our low-res spectrum is dominated bysource C+D.

3.5. AKARI Near-infrared Spectroscopy

The spectrum of sources C and D taken with AKARI iswell separated from that of the other sources (A and B) inII Zw 096, although they are all in the Np aperture, becausethe direction of spectral dispersion is nearly perpendicular tothe vector separating C, D, and A. The AKARI near-infraredspectrum of source C+D is shown in Figure 7. Strong 3.3 µmPAH emission, 3.4 µm emission from aliphatic hydrocarbongrains, and Br! emission lines are detected. The line fluxes are(2.37± 0.14) " 10#16 W m#2 and (0.50± 0.05) " 10#16 W m#2

for the PAH and the Br!, respectively, after fitting both witha Gaussian function (see Tables 3 and 4). The EQW of the3.3 µm PAH is 0.14 µm. Because the AKARI spectral resolutionbetween 2.5 and 5 µm is only R ! 120, the Br! recombinationline is not resolved.

The spectrum of source A (not shown) also exhibits 3.3 µmPAH and Br! emission. The 3.3 µm PAH has a flux of(1.77 ± 0.12) " 10#16 W m#2 and an EQW of 0.08 µm. TheBr! line flux is (0.33 ± 0.06) " 10#16 W m#2.

3.6. Bright Star Clusters in the Remnant

The high-resolution HST imaging data show that a largenumber of bright star clusters surround the II Zw 096 nucleiand are present throughout the merging disks (see Figure 2).From the B- and I-band images, we identify 128 clusters withapparent magnitudes of 18.3 mag ! mB ! 26.7 mag (Vega)and 17.9 mag ! mI ! 25.7 mag (Vega) with 80% completenesslimits of 26.0 mag and 25.5 mag, respectively. Of these, 97 and88 are detected in the FUV and H-band images, respectively.

Figure 7. AKARI near-infrared spectrum of sources C and D. The wavelengthscale is given in the rest frame.

For all clusters, we use a 0.$$1 radius circular aperture for theHST photometric measurements. We estimate that the clusterscontribute % 14% of the total FUV luminosity (& 8 " 1036 W)in II Zw 096 (uncorrected for reddening).

In Figure 8, we present an FUV–optical (F140LP#F435Wversus F435W#F814W) color–color diagram of the clusters.We additionally show the colors of a number of diffuse emissionregions throughout the merger system. For the diffuse emissionregions, we use 0.$$5–0.$$8 radius circular apertures, depending onthe location in the FUV image to avoid crowded regions.

Two potential explanations for the observed color distributionare considered. First, if all the clusters are coeval, then theirdifferent colors result from variable extinction. In other words,those clusters or diffuse emission regions which have colors ofF435W#F814W " 1.0 mag or F140LP#F435W " 0.0 magare heavily reddened (AV " 3 mag). When we de-redden an

Spitzer mid-IR• Embedded Star

Formation & AGN• Dust emission

(Inami et al. 2010)

5 kpc

Chandra X-ray• Embedded Star

Formation & AGN

K. Iwasawa et al.: C-GOALS survey

10.5 2 5

110

FE /

10!1

4 erg

keV

!1cm

!2s!1

Energy (keV)

UGC8387

10.5 2 5

110

FE

Energy (keV)

UGC8696

10.5 2 5

0.1

1FE

Energy (keV)

F14348!1447

10.5 2 5

0.1

1

FE

Energy (keV)

F14378!3651

10.5 2 5

110

FE

Energy (keV)

VV340 N

10.5 2 5

110

FE

Energy (keV)

VV705 N

10.5 2 5

0.1

1

FE

Energy (keV)

F15250+3608

10.5 2 5

110

FE

Energy (keV)

UGC9913

10.5 2 5

110

FE

Energy (keV)

ESO69!IG6 N

10.5 2 5

1010

0

FE

Energy (keV)

NGC6240

10.5 2 5

110

FE

Energy (keV)

F17132+5313

10.5 2 5

110

FE

Energy (keV)

F17207!0014

10.5 2 5

110

FE

Energy (keV)

F18293!3413

10.5 2 5

0.1

110

FE

Energy (keV)

ESO593!IG8

10.5 2 5

0.1

1

FE

Energy (keV)

F19297!0406

10.5 2 5

110

0.5

25

20

FE

Energy (keV)

19542+1110

10.5 2 5

110

FE

Energy (keV)

CGCG448!020

10.5 2 5

110

FE

Energy (keV)

ESO286!IG19

10.5 2 5

0.1

1

FE

Energy (keV)

21101+5810

10.5 2 5

110

FE

Energy (keV)

ESO239!IG2

10.5 2 5

110

FE

Energy (keV)

F22491!1808

10.5 2 5

110

FE

Energy (keV)

ESO148!IG2

10.5 2 5

0.1

1

FE

Energy (keV)

ESO148!IG2 N

10.5 2 5

0.1

1

FE

Energy (keV)

ESO148!IG2 S

Fig. 4. continued.

A106, page 11 of 50

(Inami et al. 2010)

5 kpc

JVLA• Embedded Star

Formation & AGN

(Barcos et al. 2013)

5 kpc

ALMA• Dust continuum

distribution

• star-forming molecular gas tracers

(Stierwalt et al. 2013)

Cycle

0

Stay T

uned

CO(6→5)1.6μm

Isolated System?

Mid-Stage - progenitors intact

NGC 5257/8

Mid-Stage - overlapping disks, double AGN + outflow

(Mazzarella et al. 2012)

Late Stage - single nucleus, weak AGNNGC 2623

(Evans et al. 2008)

The utility of IR/submm/radio emission

• assessment of the ionization mechanisms/source(s) in obscured regions and the optical depth towards these regions

• assessment of the physical state in molecular clouds

• measurement of the physical sizes of active regions and thus the IR surface brightness

• measurement of the kinematics of the gas, and thus a assessment of rotation (e.g. for dynamical mass estimated) and importance of infall/outflow in obscured regions

...at wavelengths that are energetically important

Arp 220

Mrk 231

1010

1011

ALMA

Spitzer

JWST

JVLA

•JWST, ALMA and JVLA provide sub-” resolution

Herschel

SED as a function of telescope λ coverage

The utility of IR/submm/radio emission

1. Spitzer: AGN/Starburst diagnostics (see Veilleux’s talk)

2. Herschel: the [C II]/FIR decrement

3. Herschel: CO SLED

4. ALMA: CO(6→5) emission

(1) AGN-Starburst diagnostics: ISO data

(Genzel et al. 1998)

high

/low

ioni

zatio

n →

←destruction by AGN

AGN-Starburst diagnostics

ISO

FIG. 2.ÈISOPHOT-S spectra (j/*j D 90) of 15 ULIRGs. Observed wavelengths are shown at the top and rest wavelengths at the bottom of each spectrum. Flux densities (not including the correctionfor chopped mode of a factor of D1.4) are in janskys. Error bars as provided by PIA are shown for UGC 5101.(° 2.1)

(Genzel et al. 1998)

AGN-Starburst diagnostics

(Genzel et al. 1998; Armus et al. 2007)

ISO

Fig. 1.—IRS Short-Low and Long-Low spectra of the 10 BGSULIRGs in order of decreasing 25 !mflux density. Prominent emission features and absorption bands (thelatter indicated by horizontal bars) aremarked on representative spectra. Not all features aremarked on all spectra; see Tables 3 and 4 formeasured features. Expanded views ofthe 5Y14 !m regions (Short-Low) of each spectrum are shown at the end.

FIG. 2.ÈISOPHOT-S spectra (j/*j D 90) of 15 ULIRGs. Observed wavelengths are shown at the top and rest wavelengths at the bottom of each spectrum. Flux densities (not including the correctionfor chopped mode of a factor of D1.4) are in janskys. Error bars as provided by PIA are shown for UGC 5101.(° 2.1)

Spitzer

AGN-Starburst diagnostics

(Armus et al. 2007)

Fig. 1.—IRS Short-Low and Long-Low spectra of the 10 BGSULIRGs in order of decreasing 25 !mflux density. Prominent emission features and absorption bands (thelatter indicated by horizontal bars) aremarked on representative spectra. Not all features aremarked on all spectra; see Tables 3 and 4 formeasured features. Expanded views ofthe 5Y14 !m regions (Short-Low) of each spectrum are shown at the end.

Spitzer

PAH

absorption by ionizing source

•E.g., Strength of PAH anti-correlated with AGN strength

AGN-Starburst diagnostics

(Armus et al. 2007)

Figure 2, normalized at 60 !m (Fig. 2a) and 25 !m (Fig. 2b). Inboth normalizations, IRAS 08572+3915 has the flattest (bluest)andArp 220has the steepest (reddest) spectrumof theBGSULIRGsample. There is nearly a factor of 50 (25) spread among theULIRGs at 5.5 !m when normalized at 60 !m (25 !m).

We also detect absorption features from gas-phase interstellarC2H2 and HCN at 13.7 and 14.0 !m, respectively, in 4 of the 10BGS ULIRGs, namely Mrk 231, Arp 220, IRAS 08572+3915,and IRAS 15250+3609. Both absorption features are detectedin all galaxies, although the HCN line in Mrk 231 is uncertaindue to a continuum ‘‘jump’’ just redward of the absorption. Thesefeatures are labeled in Figure 3. Weak C2H2 absorption mayalso be present in IRAS 05189!2524 andMrk 273, but in thesegalaxies the HCN line is not detected. Fluxes and equivalentwidths (EQWs) for the C2H2 absorption-line range from 5Y10 ;10!21 W cm!2 s!1 and 0.002Y0.012!m, respectively, with IRAS15250+3609 having the deepest C2H2 absorption and Mrk 231having the shallowest absorption among the four ULIRGs. Fluxesand EQWs for the HCN line are smaller, in the range 2Y6 ;10!21 W cm!2 s!1 and 0.002Y0.007 !m, respectively, again withIRAS 15250+3609 showing the deepest absorption. The strengthsof the C2H2 and HCN features, as well as the measured "9:7 arelisted for the BGS ULIRGs in Table 5, below. The strengths ofthese features, when fit to models that vary the excitation tem-perature and column density of the gas, can be used to constrainthe temperature of the gas and, in some cases, its chemical evo-lutionary state (e.g., Lahuis & vanDishoeck 2000; Boonman et al.2003). A detailed analysis of these lines in ULIRGs will be pre-sented in a future paper (Lahuis et al. 2007). However, the posi-tions of the line centroids at 13.7 and 14.0 !m already suggestthat a warm gas component of "200Y400 K dominates the ab-sorption in Mrk 231, Arp 220, IRAS 08572+3915, and IRAS15250+3609.

3.2. Emission Lines

The high-resolution (SH and LH) ULIRG spectra are domi-nated by unresolved atomic fine-structure lines of Ne, O, Si, andS covering a large range in ionization potential (see Figs. 3 and4 andTables 3 and 4). Ratios of the fine-structure lines can be used

Fig. 2.—Normalized low-resolution IRS spectra. Spectra are normalized at(a) rest frame 60 !m, and (b) rest frame 25 !m.

Fig. 3.—IRS SH spectra of the 10 BGS ULIRGs. Prominent spectral features are marked.

ARMUS ET AL.152 Vol. 656

Figure 2, normalized at 60 !m (Fig. 2a) and 25 !m (Fig. 2b). Inboth normalizations, IRAS 08572+3915 has the flattest (bluest)andArp 220has the steepest (reddest) spectrumof theBGSULIRGsample. There is nearly a factor of 50 (25) spread among theULIRGs at 5.5 !m when normalized at 60 !m (25 !m).

We also detect absorption features from gas-phase interstellarC2H2 and HCN at 13.7 and 14.0 !m, respectively, in 4 of the 10BGS ULIRGs, namely Mrk 231, Arp 220, IRAS 08572+3915,and IRAS 15250+3609. Both absorption features are detectedin all galaxies, although the HCN line in Mrk 231 is uncertaindue to a continuum ‘‘jump’’ just redward of the absorption. Thesefeatures are labeled in Figure 3. Weak C2H2 absorption mayalso be present in IRAS 05189!2524 andMrk 273, but in thesegalaxies the HCN line is not detected. Fluxes and equivalentwidths (EQWs) for the C2H2 absorption-line range from 5Y10 ;10!21 W cm!2 s!1 and 0.002Y0.012!m, respectively, with IRAS15250+3609 having the deepest C2H2 absorption and Mrk 231having the shallowest absorption among the four ULIRGs. Fluxesand EQWs for the HCN line are smaller, in the range 2Y6 ;10!21 W cm!2 s!1 and 0.002Y0.007 !m, respectively, again withIRAS 15250+3609 showing the deepest absorption. The strengthsof the C2H2 and HCN features, as well as the measured "9:7 arelisted for the BGS ULIRGs in Table 5, below. The strengths ofthese features, when fit to models that vary the excitation tem-perature and column density of the gas, can be used to constrainthe temperature of the gas and, in some cases, its chemical evo-lutionary state (e.g., Lahuis & vanDishoeck 2000; Boonman et al.2003). A detailed analysis of these lines in ULIRGs will be pre-sented in a future paper (Lahuis et al. 2007). However, the posi-tions of the line centroids at 13.7 and 14.0 !m already suggestthat a warm gas component of "200Y400 K dominates the ab-sorption in Mrk 231, Arp 220, IRAS 08572+3915, and IRAS15250+3609.

3.2. Emission Lines

The high-resolution (SH and LH) ULIRG spectra are domi-nated by unresolved atomic fine-structure lines of Ne, O, Si, andS covering a large range in ionization potential (see Figs. 3 and4 andTables 3 and 4). Ratios of the fine-structure lines can be used

Fig. 2.—Normalized low-resolution IRS spectra. Spectra are normalized at(a) rest frame 60 !m, and (b) rest frame 25 !m.

Fig. 3.—IRS SH spectra of the 10 BGS ULIRGs. Prominent spectral features are marked.

ARMUS ET AL.152 Vol. 656

high-ionization [Ne V]

low-ionization [Ne II]

•E.g., [Ne V] / [Ne II] correlated with AGN strength

Spitzer IRS: AGN vs. starburst diagnostics

•Mid-IR diagnostics used: [NeV]/[NeII], [OIV]/[NeII], 6.2 μm PAH EQW & MIR continuum ratios (e.g., Veilleux et al. 2009)

•Punchline 1: Based on the MIR: 10% of local LIRGs have dominant AGN

•Punchline 2: AGN are responsible for 12% of LIR of LIRGs in the local Universe

•Is the mid-IR a sufficiently long wavelength to diagnose AGN/starburst fractions?

(Petric et al. 2011; see Veilleux’s talk for more in-depth discussion)

The Astrophysical Journal, 730:28 (11pp), 2011 March 20 Petric et al.

Figure 1. Mid-infrared [Ne v]/[Ne ii] vs. 6.2 µm PAH EQW excitation diagram.The red triangles are [Ne v] detections while the black arrows are upper limits.In all cases 1! error bars are shown. The black circles indicate ULIRGs wherethe filled symbols show detections and empty symbols show upper limits. Theblue empty stars mark upper limits for SB galaxies from Bernard-Salas et al.(2009). The solid black lines indicate the fractional AGNs and SB contributionto the MIR luminosity from the [Ne v]/[Ne ii] (vertical) and 6.2 µm PAH EQW(horizontal) assuming a simple linear mixing model. In each case, the 100%,50%, 25%, and 10% levels are marked. The 100% level is set by the averagedetected values for the [Ne v]/Ne ii] and 6.2 µm EQW among AGNs and SBs,respectively, as discussed in Armus et al. (2007). The blue line traces where thesummed SB and AGN contribution equals 100%. For most LIRGs the [Ne v]/[Ne ii] ratio suggests that the AGN contribution to the nuclear MIR luminosityis below 10%.(A color version of this figure is available in the online journal.)

empirically it has been shown that emission line ratios of [Ne v]/[Ne ii] ! 0.75 and [O iv]/[Ne ii] ! 1.75 indicate that more than50% of the nuclear MIR emission is produced by an AGN (e.g.,Armus et al. 2007 and references therein). Three LIRG nucleihave [Ne v]/[Ne ii] ! 0.75, implying that only 1% of LIRGshave more than 50% of their MIR emission powered by anAGN (see Figure 1).

For comparison with the LIRGs, nine SB galaxies are in-cluded in the figures using data from Bernard-Salas et al.(2009): NGC 660, NGC 1222, IC 342, NGC 1614, NGC 2146,NGC 3256, NGC 3310, NGC 4088, NGC 4676, NGC 4818,NGC 7252, and NGC 7714. Three of these galaxies are LIRGswhile the rest have lower IR luminosities. These particular ob-jects were used by Armus et al. (2007) to determine the zeropoints in the ionized gas diagnostics. Data for seven well-studiedULIRGs from the RBGS sample analyzed in Armus et al. (2004,2007) are also presented for comparison: Mrk 231, Arp 220,IRAS 05189!2524, Mrk 273, IRAS 08572+3915, UGC 5101,and Mrk 1014.

The [O iv] 25.890 µm line is detected in 120 nuclei (53% ofthe sources). Four LIRGs have ratios of [O iv]/[Ne ii] ! 1.75implying that only 1% of LIRGs have more than 50% of theirMIR emission powered by an AGN. The ratios of [O iv]/[Ne ii] range between 0.002 and 5.5 and have a median of0.03 and a mean of 0.24 with a dispersion of 0.74 (seeFigure 2).

For several sources, as observed in similar investigations (e.g.,Farrah et al. 2007), the estimated contributions to the MIRluminosity from star formation and an AGN do not add upto 100%. In Figures 1 and 2, these sources are located eitherbelow or above the mixing line between the diagnostics, that isas the PAH EQW drops the [Ne v]/[Ne ii] should go up, andgalaxies should follow this line. However, many sources fallwell below this mixing line. We suggest that this is partially dueto uncertainties in the AGN zero points, that is the values of the

Figure 2. Mid-infrared [O iv]/[Ne ii] vs. 6.2 µm PAH EQW excitation diagram.Symbols have the same definition as in Figure 1 except that red triangles indicate[O iv] detections.(A color version of this figure is available in the online journal.)

Figure 3. Mid-infrared diagnostic diagram, first used by Laurent et al. (2000) andlater modified for Spitzer IRS by Armus et al. (2007), comparing the integratedcontinuum flux from 14 to 15 µm, the integrated continuum flux from 5.3 to 5.5,and the 6.2 µm PAH flux. The three vertices, labeled as AGN, H ii, and PDR,represent the positions of 3C273 from Weedman et al. (2005), and M17 andNGC 7023 from Peeters et al. (2004). These vertices were chosen to facilitatecomparison with ULIRGs as presented in Armus et al. (2007) and Brandl et al.(2006). The red lines from left to right indicate a 90%, 75%, and 50% fractionalAGN contribution to the nuclear MIR luminosity, respectively. GOALS sourcesare shown as open circles.(A color version of this figure is available in the online journal.)

[Ne v]/[Ne ii] and [O iv]/[Ne ii] flux ratios for a source in whichthe AGN contributes 100% of the MIR luminosity. In additiondifferential extinction of the [Ne v] and [O iv] lines originatingfrom the obscured region around the nucleus versus the [Ne ii]line originating from the SB (likely more extended) may lead toan underestimation of the contribution of the AGN to the MIRemission.

3.2. Dust Diagnostics

For the sample of LIRGs in GOALS the 6.2 µm PAH EQWsrange between 0.01 and 0.94 µm with a mean of 0.47 and amedian of 0.53. About 16% of LIRG nuclei have 6.2 µm PAHEQW below 0.27 µm, less than one-half of the value seen inlocal SB systems (Brandl et al. 2006). This suggests that 16%of LIRGs have an AGN which dominates the MIR emission.

Figure 3 adopted from Laurent et al. (2000) and Armus et al.(2007) shows the 15–5.5 µm continuum flux ratios versus the6.2–5.5 µm continuum flux ratios for the LIRGs. The continuum

5

(2) The [C II] 158μm/FIR decrement in U\LIRGssize e!ects in the nearby starburst nuclei, further measure-ments are warranted.

Why might high-U [=!(H)/nec] conditions exist inULIRGs and what possible configurations might weenvisage? Given the evidence for relatively higher densitiesin ULIRGs, the high-U conditions likely arise from high!(H) due to high Lyman continuum luminosities impingingon material at small radii, rather than from low electrondensities. These conditions could arise in highly compactH ii regions, or they could arise from di!erences in theglobal conditions inULIRGs compared to luminous IR gal-axies, such as might be present in very compact luminousstarburst regions where the ionized regions of large numbersof individual H ii regions merge and are possibly sur-rounded by a global PDR. Here, we might expect high U tobe accompanied by highG0.

For the most extreme high-U cases, such as might bepresent in buried or partially obscured AGNs, the dustemission at the dust sublimation radius must peak at near-IR wavelengths (Dudley & Wynn-Williams 1997). In thiscase, radiative transfer e!ects could produce excess FIRemission arising from a central source that is not associatedwith the normal PDR tracers in these galaxies and may suf-fer higher extinction than the PDR components. Thisexplanation may be responsible for the observed FIRmolecular absorption seen toward ULIRGs (Fischer et al.1999). Such entities could be present together with a compo-nent of more normal and possibly less obscured PDR andstarburst regions. These PDR and starburst regions coulddominate the line emission and thus mask the presence of anobscured central source. In this paper, we do not try toascertain quantitatively which of these scenarios isdominant, although pioneering work along these lines byLaurent et al. (2000) has attempted to quantify separatePDR, H ii region, and AGN contributions to the mid-IRspectral signatures of galaxies. Rather, one of the aims ofthis paper is to quantify the observational evidence thathigh-U e!ects are important and to emphasize the need forfurther observational and modeling e!orts.

3.6. Implications forMore Distant Sources

As discussed in Paper I, an important implication of the[C ii] deficit in ULIRGs is that the usefulness of the [C ii] linein cosmological work could be diminished if cosmologicalsources resemble this local sample. As a result of the presentwork, we can begin to provide an indication of what thismeans quantitatively. In the present sample, only 3 of 15ULIRGs resemble normal and starburst galaxies in theratio of [C ii]-to-dust continuum, so that an experiment thatcould just detect a cosmological ultraluminous IR source,such as those detected in continuum using SCUBA withnormal [C ii] properties, would need to observe five sourcesfor one [C ii] detection. Specifically, for the high-redshiftsources considered by Stark (1997), the per season detectionrate estimate would be reduced from 200 to 40. Such a detec-tion rate should still provide useful information for under-standing star formation as a function of redshift (e.g.,Madau, Pozzetti, & Dickinson 1998), since lower limits tothe massive star formation rate can be calculated directlyfrom the measured [C ii] intensity. Given the presence ofupper limits to the [C ii] flux in our sample, it is not possibleto say what would be required to detect the entire sample atcosmological distances, but it is clear that should the [C ii]

deficit hold, it may become a driver for more ambitiousfuture surveys.

In Figure 9, the ratio of [C ii] to FIR flux is plotted against8–1000 lm IR luminosity for both the sample consideredhere and the Malhotra et al. (2001) sample. For galaxies notlisted in Table 2, we take the IR luminosities from the IRASBright Galaxy Sample (Soifer et al. 1989) or its extension(Sanders et al. 1995). As a precaution, we have plotted asopen symbols sources that are listed as resolved at 25 or 60lm or both, although we have made use of the LWS data toestimate the FIR flux where appropriate in the present sam-ple. It is of interest that for the strongest upper limit in thisplot (NGC 4418), the power source of the FIR is consideredto be a buried AGN based on its mid-IR spectral properties(Roche et al. 1986; Dudley & Wynn-Williams 1997; Spoonet al. 2001). The data in Figure 9 suggest a break occurringnear 1012 L!, although additional data, say with theHerschel and Astro-F space telescopes, are required to ruleout a continuing decline at higher luminosities and to betterdefine the sharpness of the break.

Since the [C ii] 158 lm line is one of the brightest lines innormal and starburst galaxies (and ULIRGs) in the wave-length range where dust extinction plays a lesser role, itsusefulness as an astrophysical diagnostic remains undimin-ished notwithstanding the cautionary implications of thepresent work. In the context of the high-U hypothesis dis-cussed above, sources with low dust abundance relative tohydrogen would have to achieve an even higher U to pro-duce a similar diminution of the [C ii] line intensity relativeto the continuum. For such scenarios, we might expect thebreak suggested in Figure 9 to migrate to higher luminosi-ties ifU and L are generically correlated and enrichment is agradual function of cosmic time in individual sources. Ifsuch is the case, then mid- and far-infrared fine-structurelines, which are preferred ISM abundance diagnostics (e.g.,Simpson et al. 1998, 1995), as well as the [C ii] 158 lm lineitself, will be easily observed in moderate-luminosity cosmo-logical ultraluminous IR sources and can provide a checkon this scenario. These could be the majority of such sources

Fig. 9.—Ratio of [C ii] intensity to FIR flux plotted as a function of IRluminosity for the present sample (circles) and for those galaxies in theMalhotra et al. (2001) sample (diamonds) that appear in the IRAS BrightGalaxy Sample. Sources plotted as open symbols may have both con-tinuum and line emission that extend beyond the LWS beam (see text). Thesample medians calculated (ignoring upper limits) for the filled diamondsymbol galaxies and for the filled circle galaxies with LIR greater than orequal to 1012L! are plotted by solid lines.

772 LUHMAN ET AL. Vol. 594

(e.g., Luhman et al. 2003)

•L[C II] is ~ 1% of LFIR in normal, star-forming galaxies•Major coolant in the neutral ISM• ionization potential = 11.26 eV

[C II]/FIR vs. the star formation efficiency

(Gracia-Carpio et al. 2011)

• [CII]/FIR ↓ as the star formation efficiency ↑•Similar decrements are seem for other forbidden submm lines

The Astrophysical Journal Letters, 728:L7 (5pp), 2011 February 10 Gracia-Carpio et al.

10!5

10!4

10!3

10!2

10!1[C II] 158µm

FIR

[O I] 145µm

FIR

[N II] 122µm

FIR

[O III] 88µm

FIR

100 101 102 103

10!5

10!4

10!3

10!2

10!1

100 101 102 103

NGC 4418

Arp 220

IRAS 17208!0014

[O I] 63µm

FIR

100 101 102 103

[N III] 57µm

FIR

100 101 102 103

[O III] 52µm

FIR

LFIR / MH2 [ LSun / MSun ]

Line

flux

/ FI

R

HII galaxyLINERSeyfert & QSOBlue Compact DwarfUnclassifiedHigh!z galaxy

Figure 2. Global line to FIR continuum ratios in galaxies with different optical activity classifications as a function of LFIR/MH2 . PACS observations are representedwith squares. Open symbols indicate 3! upper limits to the line flux. All the line to FIR ratios start to decrease at LFIR/MH2 ! 80 L" M"

#1 (vertical lines). Thehorizontal lines indicate the approximate minimum values for the different line to FIR ratios observed in galaxies with LFIR/MH2 < 80 L" M"

#1.(A color version of this figure is available in the online journal.)

Our main result is that we now find line deficits in both the PDRand the H ii lines. All the line to FIR ratios start to decreaseat LFIR/MH2 ! 80 L" M"

#1. This result also does not seemto depend on the optical activity classification of the galaxies.We note, however, that the integrated emission of the [O i] and[O iii] lines in some Seyfert galaxies might have a significantcontribution from X-ray dominated regions close to the activegalactic nucleus (AGN; e.g., Maloney et al. 1996).

The [O i] 63 µm line was detected in absorption at somevelocities in NGC 4418, Arp 220, and IRAS 17208-0014. The[O i] 63 µm/FIR ratios used in Figure 2 correspond to theemission components only and should be regarded as lowerlimits. This will reduce the deficit observed for this line to someextent, but will not eliminate it, as other galaxies (e.g., Mrk 231;Fischer et al. 2010) with no clear absorption features in theirline profiles do still have lower [O i] 63 µm/FIR ratios thangalaxies with LFIR/MH2 < 80 L" M"

#1. The [O i] 145 µm lineis less likely to be affected by this kind of absorption and alsoshows a deficit at high LFIR/MH2 . There is no sign of absorptionby foreground material in the other fine structure lines at thecurrent PACS velocity resolution (!100–250 km s#1).

In summary, our observations show that the line to FIRcontinuum ratio drops by a factor of 3 to 10 above LFIR/MH2 !80 L" M"

#1 in all the far-infrared fine structure lines wehave observed. This drop seems to be a universal feature ofgalaxies at different redshifts and with different optical activityclassifications. Neutral and ionized gas tracers are affected,as are lines with different rest-frame wavelengths, intrinsicoptical depths, and critical densities. This eliminates mostof the proposed explanations for the [C ii] deficit, leavingthe dependence on the LFIR/MH2 ratio as the most plausibledominant factor driving the deficits. We have assumed in thisLetter that LFIR is mostly due to star formation. This might

not be true for some galaxies like Mrk 231, but even an AGNcontribution of 50% to LFIR will fall short of explaining thefactor of 10 decrease in some of the line to FIR ratios.

In the next section, we show with quantitative modeling thathigh ionization parameters can indeed account for most of theobserved line deficits. The key factor is that in dusty star-formingregions the fraction of luminosity absorbed by dust particles inionized as well as neutral gas increases with U.

4. COMBINED H ii AND PDR LINE MODELS

Abel et al. (2009) used the spectral synthesis code Cloudy(Ferland et al. 1998) to model the [C ii]/FIR ratio in cloudsexposed to different radiation fields. They found that they couldreproduce the [C ii] line deficits observed in some local LIRGsand ULIRGs, as well as their warmer IRAS S60µm/S100µmcolors, by increasing the value of the ionization parameter at thesurface of the clouds. As U increases, the H ii region is extendedto higher AV into the cloud. As a result, a larger fraction of theUV photons are absorbed by the dust in the ionized region andreemitted in the form of infrared emission. The net effect is thatthe fraction of UV photons available to ionize and excite thegas is reduced at high U, decreasing the relative intensity of thefine structure lines compared to the FIR continuum (see alsoVoit 1992). Since the number of UV photons per dust particleincreases with U, the dust temperature (and the S60µm/S100µmcolor) increases.

We have extended the analysis in Abel et al. (2009) to includeall the fine structure lines observed with PACS. We used Cloudyto simultaneously calculate the emission from the molecular,neutral, and ionized gas assuming pressure equilibrium in theclouds. The ionizing source was represented by a non-LTECoStar stellar atmosphere (Schaerer & de Koter 1997) with T$ =36,000 K, and the calculations were stopped at AV = 100 mag

3

[CII]

/FIR

LFIR/M(H2) = star formation efficiency →1 10 100 1000

low [C II]/FIR galaxies have high ionization parameters (U)

(Zhao et al. 2013; Diaz-Santos et al. 2013; see also Abel et al. 2009; Gracia-Carpio 2011)

Does [C ii] Emission Trace Current SF in LIRGs? First GOALS Results From Herschel/PACS 5

sible to disentangle.As mentioned in §2.3, the angular size of a PACS spaxel

is roughly similar to that of the aperture used to ex-tract the Spizer/IRS spectra of our galaxies. Because thePACS beam is under-sampled at 160µm (FWHM! 12!!

compared with the 9.6!! size of the PACS spaxels), andmost of the sources in the sample are unresolved at 24µmin our MIPS images (which have a similar angular res-olution as PACS at ! 80µm), an aperture correctionhas to be performed to the spectra extracted from theemission-peak spaxel of each galaxy to obtain their totalnuclear fluxes. This was the same procedure employed toobtain the mid-IR IRS spectra of our LIRGs. The nom-inal, wavelength-dependent aperture correction functionprovided by HIPE v8.0 works optimally when the sourceis exactly positioned at the center of a given spaxel. How-ever, in some occasions the pointing of Herschel is notaccurate enough to achieve this and the target can beslightly misplaced ! 3 !! (up to 1/3 of a spaxel) from thecenter. In these cases, the flux of the line might be under-estimated. We explored whether this e!ect could be cor-rected by measuring the position of the source within thespaxel. However, some LIRGs in our sample show lowsurface-brightness extended emission, either because oftheir proximity and/or merger nature, or simply becausethe gas and dust emission are spatially decoupled. This,combined with the spatial sub-sampling of the PACS/IFSdetector and the poor S/N of some sources prevented usfrom obtaining an accurate measurement of the spatialposition and width of the [C ii] emission and thereforefrom obtaining a more refined aperture correction. Thus,we performed only the nominal aperture correction pro-vided by HIPE.The individual far-IR luminosities of galaxies belong-

ing to a LIRG system formed by of two or more com-ponents were derived in a similar manner to the methodused to calculate their LIR (see §2.1). In this case, wescaled the integrated IRAS far-IR luminosity of the sys-tem (L42.5"122.5µm

FIR , as defined in Helou et al. 1985) withthe ratio between the continuum flux density of each indi-vidual galaxy evaluated at 63µm in the PACS spectrum(measured in the same spaxel as the [C ii] line), and thetotal IRAS 60µm flux density.

4. RESULTS AND DISCUSION

4.1. The [C ii]/LFIR Ratio: Dust Heating and Cooling

The far-IR fine structure line emission in normal star-forming galaxies as well as in the extreme environmentshosted by ULIRGs has been extensively studied for thepast two decades. A number of works based on ISO dataalready suggested that the relative contribution of the[C ii]157.7µm line to the cooling of the ISM in PDRs ascompared to that of large dust grains, as gauged by thefar-IR emission, diminishes as galaxies are more IR lumi-nous (Malhotra et al. 1997; Luhman et al. 1998; Brauheret al. 2008).

4.1.1. The Average Dust Temperature of LIRGs

Figure 1 (upper panel) shows the L[C II]157.7µm/LFIRratio for the GOALS sample as a function of the far-IRPACS S! 63µm/S! 158µm continuum flux density ratio.We chose to use this PACS-based far-IR color in the x-axis instead of the more common IRAS 60/100µm color

Fig. 1.— The ratio of [C ii]157.7µm to far-IR luminosity (upperpanel) and [C ii]157.7µm to monochromatic luminosity at 158µm(bottom panel) as a function of the S! 63µm/S! 158µm continuumflux density ratio for individual galaxies in the GOALS sample.Circles of di!erent colors indicate the LFIR of galaxies (see colorbar). Red diamonds mark galaxies with LIR ! 1012 L!, ULIRGs.These two plots show that the decrease of the L[C II]/LFIR ratiowith warmer far-IR colors seen in our LIRGs is primarily causedby a significant decrease of gas heating e"ciency and an increaseof warm dust emission.

mainly because of two main reasons: (1) this way weare able plot data from individual galaxies instead of be-ing constrained by the spatial resolution of IRAS, whichwould force us to show only blended sources; (2) by usingthe 63/158µm ratio we are probing a larger range of dusttemperatures within the starburst (T ! 50 " 20K);with the colder component probably arising from regionslocated far from the ionized gas-phase, and closer to thePDRs where the [C ii] emission originates. For reference,we show the relation between the PACS 63/158µm andIRAS 60/100µm colors in the Appendix. The ULIRGsin the GOALS sample (red diamonds) have a medianL[C II]/LFIR =6.5 # 10"4, a mean of 7.0 (± 0.9)# 10"4,and a standard deviation of the distribution of 4 # 10"4.LIRGs span two orders of magnitude in L[C II]/LFIR,from ! 10"2 to ! 10"4, with a mean of 3.5 # 10"3

and a median of 2.8 # 10"3. The L[C II] ranges from

The Astrophysical Journal Letters, 765:L13 (6pp), 2013 March 1 Zhao et al.

Figure 3. Ratio of [N ii] 205 µm to IR luminosity is plotted against ((a) and (d)) IR luminosity, (b) IRAS 60 µm/100 µm color, and (c) [O iii] 88 µm/[N ii] 205 µmflux ratio. The symbols are the same as in Figure 2 except that here larger circles represent ULIRGs. The symbols enclosed by an additional red circle in panels (a)and (b) are the same sources as shown in panels (c) and (d). Equation (1) is shown by the dashed line in panel (a). The ordinate of panel (d) is a nominal value afterthe correction for the hardness effect (using the equation shown in this panel). Mrk 231 is outside the plotted range of panel (d).(A color version of this figure is available in the online journal.)

only !5 eV higher than that needed for the formation of N++,and thus the [O iii]/[N ii] line ratios are even stronger indicatorof the hardness. Furthermore, this ratio is insensitive to the gasdensity (Rubin 1984). Except for M82 (Duffy et al. 1987) andMrk 231 (Fischer et al. 2010), all of the [O iii] 88 µm fluxes areadopted from Brauher et al. (2008).

Figure 3(c) shows the relation between L[O iii]/L[N ii] andL[N ii]/LIR. Clearly, there exists a good correlation such as L[N ii]/LIR decreases as L[O iii]/L[N ii] increases. There is only one out-lier, the ULIRG Mrk 231, which shows both small L[N ii]/LIRand L[O iii]/L[N ii]. Excluding Mrk 231, and using an ordinaryleast-squares linear fit to the data, we obtain log(L[N ii]/LIR) =("3.63 ± 0.06) + ("0.61 ± 0.06) log(L[O iii]/L[N ii]). The Spear-man rank correlation coefficient ! of this correlation is "0.93at a >5" level of significance.

Galaxies plotted in Figure 3(c) (hereafter the [O iii] sample)are also marked with red circles in Figures 3(a) and (b), whichindicate that the hardness variation of the heating radiation fieldamong galaxies can largely account for the scatter shown inthe LIR–L[N ii] relation, at least for LIR < 1012 L# galaxies.Actually, this is also supported by the fact that the scatter ofless luminous galaxies (0.2 dex for Brauher et al. sample; thusprobably less active in star formation and lower U) is smallerthan that of our luminous sample (0.28 dex). Keeping in mindthat the conversion from [N ii] 122 µm flux to [N ii] 205 µm fluxmay contribute largely to the scatter of Brauher et al. sample,

the real scatter in the L[N ii]–LIR relation might probably be evensmaller for these galaxies.

To obtain a direct view of how much the scatter can be re-duced, we also display the hardness-corrected L[N ii]/LIR, whichwas calculated according to the fitted L[O iii]/L[N ii]–L[N ii]/LIRrelation, for the [O iii] sample in Figure 3(d) (arbitrarily scaled).Comparing to Figure 3(a), one can see that the scatter is reducedsignificantly (from 0.26 dex to 0.11 dex by excluding Mrk 231,which is still an outlier due to that the hardness correction haslittle effect on it, and is outside the plotted range).

All of the seven ULIRGs in our sample show some deficitin L[N ii] relative to LIR. Four of the seven, including Arp 220and Mrk 231, were observed by ISO LWS, but no [O iii] 88 µmemission was detected at the 3" limit (e.g., Brauher et a. 2008),despite all showing some AGN characteristics. These galaxiesalso show very warm FIR colors (see Figure 3(b)). Therefore,very dusty H ii regions with high U might be needed to explainthe deficit of [N ii] emission in these galaxies (e.g., Fischer et al.1999; Abel et al. 2009; Draine 2011).

4. SUMMARY

In this Letter we have presented our initial results of the [N ii]205 µm emission for a sample of 70 (U)LIRGs observed bythe Herschel SPIRE FTS. Combining with the ISO and IRASdata, we have studied the correlation between L[N ii] and LIR,

5

• ... as inferred from the [OIII] / [NII]• ...and dust temperature

high ionization / low ionization → warmer dust →

low [C II]/FIR galaxies are optically thick

(Diaz-Santos et al. 2013; spectra - Brandl et al. 2006)

• [CII]/FIR ↓ as optical depth ↑

6 Dıaz-Santos et al.

! 107 to 2" 109 L!.Our results are consistent with ISO observations of

a sample of normal and moderate IR-luminous galaxiespresented in Malhotra et al. (1997) and further analyzedin Helou et al. (2001). The GOALS sample, though,populates a warmer far-IR color regime and extends thetrend inferred from the ISO data more than one order ofmagnitude to higher LIR. Despite the increase in disper-sion at 63/158µm! 1.25 or L[C II]157.7µm/LFIR " 10"3

(basically in the ULIRG domain), the fact that wefind the same tight trend independently of the rangeof IR luminosities covered by the two samples suggeststhat the main observable linked to the variation of theL[C II]/LFIR ratio is the average temperature of the dust(Tdust) in galaxies.This interpretation agrees with the last physical sce-

nario described in the Introduction, in which an increaseof the ionization parameter, <U>, would cause the far-UV radiation from the youngest stars to be less e!cientin heating the gas in those galaxies. At the same time,dust grains would be on average at higher temperaturesdue to the larger number of ionizing photons per dustparticle available in the outer layers of the H ii regions,close to the PDRs. Indeed, the presence of dust withinH ii regions has been recently observed in several star-forming regions in our Galaxy (Paladini et al. 2012).Both e"ects combined can explain the wide range ofL[C II]157.7µm/LFIR ratios and far-IR colors we observe inthe most warm LIRGs. Variations in nH, though, couldbe responsible for the dispersion in L[C II]/LFIR that wesee at a given 63/158µm ratio.To further support these findings, the bottom panel of

Figure 1 shows that the drop in the ratio of [C ii]157.7µmluminosity to the monochromatic continuum at ! 160µmis smaller by an order of magnitude than the correspond-ing drop in L[C II]/LFIR (albeit with a larger dispersion).This implies that the decrease of the L[C II]/LFIR ratiowith warmer far-IR colors seen in our LIRGs is primar-ily caused by a significant increase in warm dust emission(peaking at ! " 60µm) that is not followed by a propor-tional enhancement of [C ii] emission, contrarily to whatis normally seen in star-forming galaxies with lower IRluminosities. This suggest that the [C ii]157.7µm lineis not sensitive to the youngest stars responsible for thewarm dust emission. We will elaborate on this point in§4.2.

4.1.2. A Link Between Mid-IR Dust Obscuration andFar-IR Re-emission

The strength of the 9.7µm silicate feature is defined asS9.7µm # ln(fobs

!P/f cont

!P); with f cont

!Pand fobs

!Pbeing the

un-obscured and observed continuum flux density mea-sured in the mid-IR IRS spectra of our LIRGs and eval-uated at the peak of the feature, !P , normally at 9.7µm(see Stierwalt et al. 2013a for details on how it was cal-culated in our sample). Negative values indicate absorp-tion, while positive ones indicate emission. By definition,S9.7µm measures the apparent optical depth towards thewarm, mid-IR emitting dust. Figure 2 shows that thereis a clear trend (r = 0.65, pr = 0;" = 0.47) for LIRGswith stronger (more negative) S9.7µm to display smallerL[C II]/LFIR ratios, implying that the dust responsiblefor the mid-IR absorption is also accountable for the far-

Fig. 2.— L[C II]157.7µm/LFIR ratio for individual galaxies inthe GOALS sample as a function of the strength of the 9.7µmsilicate absorption feature, S9.7µm, measured with Spitzer/IRS.Galaxies are color-coded as a function of the S! 63µm/S! 158µmratio, with values ranging from 0.10 to 4.15 (see Figure 1). Thesolid line represents an outlier-resistant fit to the bulk of the star-forming LIRG population with S9.7µm ! "2. The dotted lines arethe ± 1! uncertainty.

IR emission. The formal fit (solid line) can be expressedas:

log(L[C II]/LFIR) = $2.30(±0.02) + 0.82(±0.05)S9.7µm

(1)

with a dispersion in the y-axis of 0.30 dex.Within the context described in the previous section,

the contrast between the inner layer of dust that is beingheated by the ionizing radiation to T ! 50K and thatof the cold dust at T " 20K emitting at ! ! 150µmwould create both: (1) the silicate absorption seen at9.7µm due to the larger temperature gradient betweenthe two dust components and (2) the increasingly higher63/158µm ratios seen in Figure 1 due to the progres-sively larger amount of dust mass that is being heated tohigher temperatures. This scenario is consistent with thephysical properties of the ISM found in the extreme envi-ronments of ULIRGs, in which the fraction of total dustluminosity contributed by the di"use ISM decreases sig-nificantly, and the emission from dust at T ! 50$ 60Karising from optically-thick ”birth clouds” (with ages" 107"8 Myr) accounts for ! 80% of their IR energy out-put (da Cunha et al. 2010). Furthermore, our findingsare also in agreement with recent results showing that theincrease of the silicate optical depth in LIRGs is relatedwith the flattening of their radio spectral index (1.4 to8.44GHz) due to an increase of free-free absorption, sug-gesting that the dust obscuration must largely be origi-nated in the vicinity of the starburst region (Murphy etal. 2013).There are a few galaxies that do not follow the cor-

relation fitted in Figure 2, showing very large silicatestrengths (S9.7µm < $2) and small L[C II]/LFIR ratios(< 10"3) typical of ULIRGs or, in general, warm galax-ies (see color coding or Figure 1). We would like to note

=ln(fobs/fcont)9.7μm

log space, i.e., the IRS estimated infrared luminosities agreewithin 23% with LIR. This is much more accurate than the esti-mates by Chary & Elbaz (2001) and Takeuchi et al. (2005). Theexcellent correlation suggests that, at least for a homogeneoussample of starburst galaxies, F15 !m and F30 !m can be used toaccurately derive LIR.

4.3. A Large Variety in Silicate Absorption

Figure 10 shows, from top to bottom, a series of starburst spec-tra, sorted by increasing absorption of the 9.8 !m silicate reso-nance. Since the peak of the resonance coincides with a minimumbetween the 7Y9 and 11Y13 !m PAH emission complexes, the ef-fect of silicate absorption only becomes apparent toward the lowerhalf of the plot. Nevertheless, the figure strikingly illustrates thestrong effect of amorphous silicates on the overall 5Y38 !m spec-tral shape and the large variations evenwithin one class of objects.Figure 10 shows that the 10 !m trough can in fact become a dom-inating feature of the spectral shape of starbursts.

The top two spectra also show a broad emission structure be-ginning at about 16 !m and extending to 21Y25 !m. The shapeof this feature is consistent with that of an 18 !m silicate emis-sion feature, which would make these the first detections of thisfeature in starburst galaxies. However, the emission structurecould also be due to theC!C!C in-plane and out-of-plane bendingmodes of PAHs (Van Kerckhoven et al. 2000). The identificationof this feature will be discussed in a future paper. The remaining

starburst galaxies in Figure 10 do not show evidence for a similaremission structure. The 13Y35 !m spectra, from NGC1222 (top)to NGC 4945 at the bottom, show an increasingly pronounceddepression, peaking at 18.5!m, signaling increasingly strong sili-cate absorption. The latter result is in full agreement with the trendfound for the 9.8 !m silicate feature.Among all the galaxies classified as starbursts, NGC 4945

(bottom spectrum in Fig. 10) is a ‘‘special case,’’ exhibiting byfar the strongest dust obscuration to its nuclear region (Spoonet al. 2000). Based on the IRS spectrum, the apparent opticaldepth in the 9.8 !m silicate feature is at least 4 and may be higher,depending on the choice of the local continuum. Apart fromstrong amorphous silicate absorption, the line of sight also revealsthe presence of a 23!mabsorption feature. Following the analysisof deeply obscured lines of sight toward ULIRG nuclei (Spoonet al. 2006), we attribute the 23 !m feature to crystalline silicates(forsterite). Their detection in NGC 4945 suggests that crystal-line silicates are perhaps a more common component of the ISMand not just limited to ULIRG nuclei.Figure 11 compares the average starburst template from Fig-

ure 6 to NGC 4945, the most extincted source within our sample.For comparison we also show the heavily embedded ULIRGIRAS 08572+3915 from Spoon et al. (2006). The usually rathershallow 18 !m silicate band reduces the continuum by a largefactor compared to the average starburst spectrum. While NGC4945 has a similarly strong silicate absorption feature as theULIRG IRAS 08572+3915, the starburst also shows strong PAHcomplexes near 16.4 and 17.1!m,which are absent in theULIRGspectrum.Crystalline silicates, causing the features at 16 and 19!min IRAS 08572+3915, appear to be absent in the starburst tem-plate, although they are more difficult to discern at our spectralresolution given the presence of PAH features and the [S iii] linein that range.While NGC 4945 shows by far the most extreme silicate ab-

sorption within our sample, it is not an exotic object but definesthe endpoint of a sequence of increasing optical depths. Recently,Dale et al. (2006) reported on the Spitzer SINGS survey of 75nearby galaxies. The typical target in their sample shows onlymodest dust obscuration, consistent with an AV " 1 foregroundscreen and a lack of dense clumps of highly obscured gas and dust.Unfortunately, we cannot distinguish between the local and global

Fig. 10.—IRS spectra of Mrk 52, NGC 7714, NGC 1222, NGC 7252, NGC2146, NGC 520, NGC 3628, and NGC 4945 arranged to illustrate the gradualeffect of increasing silicate absorption from top to bottom. The flux densitiesFk have been arbitrarily scaled for better comparison.

Fig. 11.—Comparison of the average starburst template (dotted line) fromFig. 6 to the most extincted source within our sample, NGC 4945 (dashed line),and to the ULIRG IRAS 08572+3915 (solid line), chosen as an extreme case ofextinction from the sample of Spoon et al. (2006). The spectra have been nor-malized at 14.5 !m.

BRANDL ET AL.1140 Vol. 653

←in

crea

sing

opt

ical

dep

th

low [C II]/FIR galaxies have compact IR-emitting regions

(Diaz-Santos et al. 2013)

• [CII]/FIR ↓ as the compactness of the MIR emission ↑

8 Dıaz-Santos et al.

Fig. 3.— L[C II]157.7µm/LFIR ratio as a function of the lumi-nosity surface density at 15µm, !15µm (top), and the fraction ofextended emission at 13.2µm, FEE13.2µm (bottom), for individualgalaxies in the GOALS sample. Galaxies are color-coded as a func-tion of the 63/158µm ratio (top) and the strength of the silicatefeature, S9.7µm (bottom). A linear fit to the datapoints withoutlimits is shown in the top panel as a solid line. The dotted linesare the ± 1! uncertainty.We note that the 15µ luminosities aremeasured within the Spitzer/IRS LL slit while the 13.2µm sizeswere obtained from the SL module.

able to destroy a significant fraction of the smallest PAHmolecules (Voit 1992; Siebenmorgen et al. 2004). In par-ticular, a galaxy is regarded as mid-IR AGN-dominatedwhen its 6.2µmPAH EW ! 0.3, and it is classified asa pure starburst when 6.2µmPAH EW " 0.5 (althoughthese limits are not strict). Sources with intermediatevalues are considered composite galaxies, in which bothstarburst and AGN may contribute significantly to themid-IR emission.

4.2.1. [C ii]157.7µm Deficit in Pure Star-Forming LIRGs

Figure 4 shows the same data as Figure 3 (top) butcolor-coded as a function of the 6.2µmPAH EW. As wecan see, when only pure star-forming galaxies are con-sidered (yellow and red circles, where 6.2µmPAH EWs! 0.5µm), the L[C II]/LFIR ratio drops by an order ofmagnitude, from 10!2 to" 10!3. This indicates that thedecrease in L[C II]/LFIR among the majority of LIRGs is

Fig. 4.— L[C II]157.7µm/LFIR ratio as a function of the lumi-nosity surface density at 15µm, !15µm for individual galaxies inthe GOALS sample. This Figure is the same as Figure 3 butcolor coded as a function of the 6.2µmPAH EW to highlightthose sources whose mid-IR emission might indicate the presenceof an AGN (see text for details). The solid line is a fit to purestar-forming LIRGs only, defined as to have 6.2µmPAH EWs! 0.5µm.

not caused by a rise of AGN activity but instead is afundamental property of the starburst itself. It is only inthe most extreme cases, when L[C II]/LFIR < 10!3, thatthe AGN could play a significant role. In fact, power-ful AGN do not always reduce the L[C II]/LFIR ratio, asshown also in Sargsyan et al. (2012). Stacey et al. (2010)also find that the AGN-powered sources in their high-redshift galaxy sample display small L[C II]/LFIR ratios.However, they speculate that except for two blazars, thedeficit seen in these sources could be caused compact, nu-clear starbursts (with sizes less than 1 # 3 kpc) perhapstriggered by the AGN. We will return to this discussionin §4.2.2.The result obtained above also implies that the

[C ii]157.7µm line alone is not a good tracer of the SFRin most LIRGs since it does not account for the increaseof warm dust emission (Figure 1) seen in the most com-pact galaxies that is usually associated with the mostrecent starburst. In Figure 4 we fit those sources with6.2µmPAH EWs ! 0.5µm to provide a relation be-tween the L[C II]/LFIR ratio and the !15µm for pure star-forming LIRGs. The analytic expression of the fit is:

log(L[C II]

LFIR) = 1.87(± 0.39) # 0.47(±0.05) $ log(!15µm)

(3)with a dispersion in the y-axis of 0.20 dex. The slopeand intercept of this trend are indistinguishable (withinthe uncertainties) from those obtained in Eq. (2), fur-ther supporting the idea that the influence of AGNactivity is negligible among IR-selected galaxies with10!3 <L[C II]/LFIR < 10!2.

4.2.2. The Influence of AGN in the [C ii] Deficit

Figure 5 shows the L[C II]157.7µm/LFIR ratio as a func-tion of the 6.2µmPAH EW for the LIRGs in GOALS.

Fit to “SB-dominated” U/LIRGs

low PAH EQW -AGN-dominant

U/LIRGs

H+ H+

H+

H+

H+

H+

H+

H+

H+H+

H+

Ho

PAH

PAH

PAH

PAH

PAH

Ho

Ho

Ho

Ho

Ho[CII]Ho

Ho

[CII]

[CII]

>>[CII]<<

[CII]

>>[CII]<<

e-

H+

H+

H+

HII Region

PDR

•As U ↑, the gas opacity in the HII regions ↓•chance of UV photons → dust → FIR photons high, especially if the starburst region is optically thick

Molecular Cloud

UV

FIR

U/LIRGs with high SSFRs

(Diaz-Santos et al. 2013)

... form their stars in warm, compact, optically thick regions

Does [C ii] Emission Trace Current SF in LIRGs? First GOALS Results From Herschel/PACS 11

Fig. 6.— L[C II]157.7µm/LFIR ratio versus LIR/M! normalized

to the LIR/M! of MS galaxies at z! 0 (SSFRMS " 0.09Gyr!1),for individual galaxies in the GOALS sample. Galaxies are color-coded as a function of their 6.2µmPAH EW. Colored circles indi-cate sources for which an EW is available. Squares indicate lowerlimits. Small black circles are sources for which there is no infor-mation. The solid line is a fit to pure star-forming LIRGs only,excluding those sources that may harbor an AGN that could dom-inate their mid-IR emission (6.2µmPAH EWs < 0.5µm; see Fig-ure 5). The dotted lines are the ± 1! uncertainty. The dashed lineindicates the SSFR of MS galaxies. The dotted-dashed line repre-sents 3#SSFRMS, the limit above which galaxies are consideredto be starbursting.

SSFR; see, e.g., Gracia-Carpio et al. 2011) then thisequation can be applied to predict the [C ii] luminos-ity of star-forming galaxies at any redshift for which ameasurement of the (far-)IR luminosity and stellar massare available as long as their SSFR is normalized to theSSFRMS at that z to account for the increase in gas massfraction (Mgas/M!) and therefore SSFR of MS galaxiesat higher redshifts (Daddi et al. 2010a; Magdis et al.2012). Conversely, in future large IR surveys like thoseprojected with the Cornell-Caltech Atacama Telescope(CCAT), this relation could be used for estimating theSSFR or stellar mass of detected galaxies when their [C ii]and far-IR luminosities are known.As an example, in Figure 6 we show two interme-

diate redshift z ! 1.2 galaxies from Stacey et al.(2010) (SMMJ123633 and 3C368) and three high redshiftz ! 4.5 sub-millimeter galaxies (SMGs) (LESS 61.1 andLESS 65.1: Swinbank et al. 2012; LESS 033229.3: DeBreuck et al. 2011; Wardlow et al. 2011) marked as blacksquares and diamonds, respectively. To calculate theSFR of the SMGs we assumed that their LIR " 2#LFIR.We estimated the SSFRMS at the redshift of each galaxyin the same manner as described above. However, sinceEq. 13 from Elbaz et al. (2011) is calibrated only up toz ! 3, we normalized the SSFR of the SMGs by theSSFRMS at this redshift (" 4.6Gyr!1). As we can see,three galaxies follow the correlation suggesting that theyare starbursting sources with L[C II]/LFIR ratios consis-tent with their normalized SSFRs. On the other hand,the two remaining high-z SMGs show significantly lowerL[C II]/LFIR ratios than the average of LIRGs for thesame normalized SSFR. In particular, one of them dis-play a L[C II]/LFIR more than an order of magnitude

lower than the value predicted by the fit to our localgalaxy sample. Interestingly, both SMGs lie in the pa-rameter space where most of the mid-IR identified AGNare located (red stars). This may suggest that these twogalaxies could harbor AGN or unusually week [C ii] emis-sion for their normalized SSFR.Eq. 3 can also be used to predict the mid-IR luminosity

surface density of star-forming galaxies at high redshiftsfor which the [C ii] and far-IR luminosity are known,such as those which may be found in future spectroscopicsurveys with CCAT. Moreover, if a measurement of theirrest-frame mid-IR luminosity is also available, this corre-lation can be further used to estimate the physical size ofthe star-forming region in the mid-IR. This is particularlyuseful for Herschel detections of z> 3 sources in deepfields that have been already observed. In this cases, thepredicted mid-IR size of galaxies could be compared withdirect measurements of the size of their far-IR emittingregion as observed with ALMA on physical scales simi-lar to those we are probing in our GOALS LIRGs withPACS.Finally, because GOALS is a complete flux-limited

sample of local LIRGs, we are able to predict the con-tamination of sources hosting AGNs in future large-scalesurveys with both [C ii] and far-IR measurements. InTable 1 we provide the percentages of galaxies with mid-IR detected AGNs classified in di!erent L[C II]/LFIR andS" 63µm/S" 158µm bins. The values provided in the Ta-ble were computed using two conditions for the detectionof the AGN that serve as upper and lower limits for theestimated fractions (columns (2) and (3)). The first wasbased on the 6.2µmPAH EW only and the second re-quired of an additional mid-IR diagnostic to classify thegalaxy as harboring an AGN (see above). For exam-ple, we predict an AGN contamination of up to 65% forL[C II]/LFIR < 5 # 10!4, which implies that at least 1/3of IR-selected sources with extremely low L[C II]/LFIRratios will be powered by starbursts. Moreover, at thelevels of L[C II]/LFIR $ 5 # 10!4 or 63/158µm< 2, theAGN fraction is expected to be % 20%.

6. CONCLUSIONS

We obtained new Herschel/PACS [C ii]157.7µm spec-troscopy for 200 LIRG systems in GOALS, a 60µm flux-limited sample of all LIRGs detected in the nearby Uni-verse, which are the most likely local counterparts of thedominating IR-selected population of star-forming galax-ies at high redshift. A total of 240 individual galaxieswhere observed in [C ii]. We combined this informationtogether with Spitzer/IRS spectroscopic data to providethe context in which the observed [C ii] luminosities and[C ii]/LFIR ratios are best explained. We have found thefollowing results:

• The LIRGs in GOALS span two orders of mag-nitude in L[C II]/LFIR, from ! 10!2 to 10!4,with a median of 2.8 # 10!3. ULIRGs have amedian of 6.5 # 10!4. The L[C II] range from! 107 to 2# 109 L" for the whole sample. TheL[C II]/LFIR ratio is correlated with the far-IRS" 63µm/S" 158µm continuum color. We find thatall galaxies follow the same trend independentlyof their LIR, suggesting that the main observable

Main Sequence Starbursts

= SFR/M*

U\LIRGs with high SSFRs

(Elbaz et al. 2011)

A&A 533, A119 (2011)

Fig. 16. Left: SFR – M! correlation at z " 0. Galaxies classified as compact are marked with large filled red dots. Solid line: fit to the main sequenceSFR–M! relation: SFR#M!/[4$ 109 M%]. Dotted lines: 16th and 84th percentiles of the distribution around the sliding median (0.26 dex). Right:relation of the sSFR and IR surface brightness of galaxies for which a radio size was estimated. The vertical dashed line illustrates the thresholdabove which galaxies have been classified as compact. The solid and dashed red lines are a fit to the sliding median of the relation (Eq. (11)) andits 68% dispersion.

RSB, which measures the excess in sSFR of a star-forming galaxy(which we label its “starburstiness”), as defined in Eq. (10):

RSB = sSFR/sSFRMS = "MS/" [> 2 for starbursts], (10)

where the subscript MS indicates the typical value for main se-quence galaxies at the redshift of the galaxy in question. A star-burst is defined to be a galaxy with RSB & 2. 75% of the galax-ies with compact star formation (!IR & 3$ 1010 L% kpc'2) haveRSB > 2, hence are also in a starburst mode, and 93% of themhave RSB>1. Conversely, 79% of the starburst galaxies are “com-pact”. Globally, starburst galaxies with sSFR > 2 $ (sSFR) havea median !IR " 1.6$ 1011 L% kpc'2, hence more than 5 timeshigher than the critical IR surface brightness above which galax-ies are compact. The size of their star-forming regions is typi-cally 2.3 times smaller than that of galaxies with sSFRMS.

The sSFR and !IR are correlated with a 0.2 dex dispersion(Fig. 16-right) following Eq. (11), where sSFR is in Gyr'1 and!IR in L% kpc'2:

sSFR = 1.81 ['0.66,+1.05]$ 10'4 $ !0.33IR . (11)

Both parameters measure specific quantities related to the SFR:the sSFR is measured per unit stellar mass, while !IR is relatedto the SFR (derived from Ltot

IR) per unit area. Note, however, thatit was not obvious a priori that these quantities should be corre-lated, since the stellar mass of most galaxies is dominated by theold stellar population, whereas the IR (or radio) size used to de-rive !IR measures the spatial distribution of young and massivestars.

Because the starburstiness and the IR8 ratio are both en-hanced in compact star-forming galaxies, they are also corre-lated as shown by Fig. 17. The fit to this correlation is given inEq. (12):

RSB = (IR8/4)1.2. (12)

The dispersion in this relation is 0.3 dex. Hence we find that itis mostly compact starbursting galaxies that present atypicallystrong IR8 bolometric correction factors, although there is nota sharp separation of both regimes, but instead a continuum ofvalues.

Fig. 17. RSB = sSFR/sSFRMS versus IR8 (=LtotIR/L8) for z " 0 galaxies

(AKARI and GOALS samples). The red line is the best fit (plain) andits 0.3 dex dispersion (dashed).

6. IR8 as a tracer of star formation compactnessand “starburstiness” in distant galaxies

In the previous section, we have defined two modes of starformation:

– a normal mode that we called the infrared main sequence, inwhich galaxies present a universal IR8 bolometric correctionfactor and a moderate star formation compactness, !IR, and

– a starburst mode, identified by an excess SFR per unit stellarmass, hence sSFR, as compared to the typical sSFR of mostlocal galaxies.

Galaxies with an enhanced IR8 ratio were systematically foundto be forming their stars in the starburst mode and to show astrong star formation compactness (!IR > 3$ 1010 L% kpc'2).In order to separate these two modes of star formation in dis-tant galaxies as well, we first need to define the typical sSFR

A119, page 16 of 26

radio calibratedsizes

... form their stars in warm, compact, optically thick regions

U\LIRGs with high SSFRs

(Murphy et al. 2013)... form their stars in warm, compact, optically thick regions

4 MURPHY ET AL.

Figure 1. Top: The linear 8.44GHz size of each source plottedagainst the radio spectral index (generally) measured between 1.49and 8.44GHz (see Table 1). Bottom: The same as the top panel,except that the 8.44GHz brightness temperature is now plottedagainst the radio spectral spectral index. As already shown byCondon et al. (1991), there is a general trend for sources havingflat spectral indices to be small in size and have large 8.44GHzbrightness temperatures. In both panels, sources with 8.44GHzbrightness temperatures larger than 104.5 K, indicating the pres-ence of an AGN, are identified by filled circles. We additionallyidentify those sources categorized as being AGN dominated bytheir low 6.2µm PAH EQW (circles with crosses); note, each ofthe sources with an 8.44GHz brightness temperature larger than104.5 K also has a 6.2µm PAH EQW < 0.27µm .

3.2. Identifying Potential AGN-Dominated Sources

In Figure 3 we look at the 6.2µm PAH EQW, whichmeasures the relative strength of the PAH flux with re-spect to the amount of hot dust emission at 6.2µm, ver-sus radio spectral index. The 6.2µm PAH EQW hasbeen used to indicate the presence of an AGN since theytypically have very small PAH EQWs (e.g., Genzel et al.1998; Armus et al. 2007). Specifically, starburst domi-nated systems appear to have 6.2µm PAH EQWs thatare ! 0.54µm (Brandl et al. 2006), while AGN have6.2µm PAH EQWs are " 0.27µm. We find that thegalaxies having steep spectral indices (i.e., !8.44GHz

1.49GHz !0.6) never have extremely low 6.2µm PAH EQWs (i.e.,"0.1µm). Additionally, all 3 sources identified as AGNvia their excessively large 8.44GHz brightness tempera-ture sit in the PAH-defined AGN region.Using the 6.2µm PAH EQW to split the sources up

into starburst and AGN dominated systems, we find thatthe mean radio spectral index for starburst dominatedsystems is 0.74, with a standard deviation of 0.15. For

Figure 2. The optical depth of the 9.7µm silicate feature plot-ted against the radio spectral index (generally) measured between1.49 and 8.44GHz (see Table 1). Sources with 8.44GHz bright-ness temperatures larger than 104.5 K, indicating the presence ofan AGN, are identified by filled circles. We additionally identifythose sources categorized as being AGN dominated by their low6.2µm PAH EQW (circles with crosses); note, each of the sourceswith an 8.44GHz brightness temperature larger than 104.5 K alsohas a 6.2µm PAH EQW < 0.27 µm . The dotted line is an ordi-nary least squares fit to all galaxies while the dashed line excludessources identified as possible AGN. Both fits illustrate a generaltrend for sources to have decreasing silicate optical depths as theirradio spectra steepen. The source to the lower left, havinga negative radio spectral index, is UGC08058 (Mrk 231)and is not included in the ordinary least squares given thatwe only have a lower limit for its silicate optical depth.

Figure 3. The 6.2µm PAH EQW plotted against the radio spec-tral index (generally) measured between 1.49 and 8.44GHz (seeTable 1). The horizontal lines indicate regions for which sourcesare thought to be AGN or star formation driven based on their mea-sured 6.2µm PAH EQW. Sources having spectral indices steeperthan !8.44GHz

1.49GHz ! 0.6 never appear to have extremely low 6.2µmPAH EQWs (i.e., "0.1µm). Sources with 8.44GHz brightness tem-peratures larger than 104.5 K, indicating the presence of an AGN,are identified by filled circles. The source to the lower left, havinga negative radio spectral index, is UGC08058 (Mrk231).

the AGN dominated systems, we find a significantly flat-ter mean spectral index of 0.45, albeit with a much largerscatter of 0.27.

4. DISCUSSION

Taking archival high resolution 8.44GHz data from theliterature, we have compared the radio and mid-infraredspectral properties for a sample of local (U)LIRGs. Wenow use these results to investigate how the dust is dis-

α is correlated with radio extent

RADIO AND MID-INFRARED PROPERTIES OF COMPACT STARBURSTS 5

Figure 4. The infrared-derived star formation rate of each galaxyplotted against its stellar mass. The size of each plotting symbolis a function of the galaxies radio spectral index, with smaller sizesindicating flatter spectral indices. The solid and dashed lines in-dicate the main-sequence and 1! (0.26 dex) dispersion for z ! 09µm-detected AKARI galaxies as measured by Elbaz et al. (2011,SFR/M! ! 0.25Gyr"1). Each of the infrared-luminous galax-ies in this sample is located well above the main-sequence rela-tion. Sources having 8.44GHz brightness temperatures larger than104.5 K, indicating the presence of an AGN, are identified by filledcircles.

optical depth most likely indicates that the dust obscu-ration in these sources occurs in the compact starburstitself, and not by extended dust located in the foregroundgalaxy disks.

4.2. Radio Spectral Indices and the Main Sequence ofStar-Forming Galaxies

In Figure 4 we plot the z ! 0 main sequence rela-tion as given in Elbaz et al. (2011), which is definedusing a local sample of galaxies detected at 9µm withthe AKARI/Infrared Camera (IRC; Onaka et al. 2007).Most galaxies lie along a linear relation between starformation rate and stellar mass, such that the specificstar formation rate (sSFR = SFR/M!) is constant, beingsSFR ! 0.25Gyr"1 with a dispersion that is a factor of1.82. It is worth noting that the z ! 0 main-sequence re-lation reported by Elbaz et al. (2007) has a slope that isslightly flatter than unity, being 0.77. If we instead adoptthis relation for the z ! 0 main sequence, our conclu-sions are not strongly a!ected. Like Elbaz et al. (2011),we decide to chose the main-sequence relation having aslope of unity as this is consistent with findings at z ! 1(Elbaz et al. 2007), z ! 2 (Pannella et al. 2009), z ! 3(Magdis et al. 2010), and z ! 4 (Daddi et al. 2009).We overplot our sample galaxies and find that each

source lies more than 1! away from the main sequence,as expected for starbursting galaxies. The size of eachplotting symbol is a function of the radio spectral index ofthe source, with smaller symbols indicating flatter spec-tral indices. Sources having the flattest spectral indicestend to be furthest from the main sequence.In Figure 5 we investigate the relation between the

distance of these galaxies from the main sequence versusthe compactness of the starburst as measured by theirradio spectral index. Since we have adopted a constantspecific star formation rate to define the z ! 0 main se-quence, the distance from the main sequence is given bythe specific star formation rate itself. We find a gen-eral trend where the specific star formation rate among

Figure 5. The specific star formation rate of each galaxy plot-ted as a function of radio spectral index (generally) measured be-tween 1.49 and 8.44GHz (see Table 1). Sources having 8.44GHzbrightness temperatures larger than 104.5 K, indicating the pres-ence of an AGN, are identified by filled circles. We additionalidentify those sources categorized as being AGN dominated bytheir low 6.2µm PAH EQW (circles with crosses); note, each ofthe sources with an 8.44GHz brightness temperature larger than104.5 K also has a 6.2µm PAH EQW < 0.27µm . The solid lineis a fit to the non-AGN sources, as identified by either their high8.44GHz brightness temperature or low 6.2µm PAH EQW, of theform sSFR = "(1"e!"!MS )+sSFRMS e!"!MS (see §4.2), which isforced to fit through the typical radio spectral index (#MS ! 0.8)and specific star formation rate (sSFRMS ! 0.25Gyr"1) of z ! 0star-forming galaxies (diamond). The corresponding error bars in-dicate the 1! dispersions of these values.

the sample increases as the radio spectral index flattens.While a general trend does exist, it is worth noting thatthe entire sample lies well above the z ! 0 main se-quence. The minimum and mean specific star formationrates among the sample are a factor of !2.5 and 15 timeslarger than the z ! 0 main sequence value, respectively.However, many sources have radio spectral indices con-sistent with normal star-forming galaxies (i.e., " ! 0.8;Condon 1992). Thus, it seems that a flat radio spec-tral index does indicate that a source is above the mainsequence, but galaxies having a “normal” star-formingradio spectra (i.e., " ! 0.8) are not necessarily on themain sequence.Since we know that typical z ! 0 star-forming galaxies

lying on the main sequence have radio spectral indicesnear " ! 0.8, it may be more appropriate to fit the datawith something other than an ordinary least squares fitthat passes through this normalization point. The radiobrightness temperature, which is a measure of the surfacebrightness or “compactness” of a source, is given by theexpression Tb = Te(1 " e"!ff ) + Tb(0)e"!ff , where Te isthe thermal electron temperature and #! is the free-freeoptical depth. The specific star formation rate has beenshown to be related to the compactness (brightness tem-perature) of a source (e.g., Elbaz et al. 2011), and in Fig-ure 2 we have found a relation between the radio spectralindex and the optical depth of the 9.7µm silicate feature.Thus, let us assume that the specific star formation rateis related to the radio spectral index with a function hav-ing the form of sSFR = $(1"e"""MS)+sSFRMS e"""MS ,where "MS and sSFRMS are the typical radio spectral in-dices and specific star formation rates for galaxies on thez ! 0 main sequence, respectively, and $ is a normaliza-tion constant related to the maximum specific star for-mation rate, which is determined by the fit to the data.

log space, i.e., the IRS estimated infrared luminosities agreewithin 23% with LIR. This is much more accurate than the esti-mates by Chary & Elbaz (2001) and Takeuchi et al. (2005). Theexcellent correlation suggests that, at least for a homogeneoussample of starburst galaxies, F15 !m and F30 !m can be used toaccurately derive LIR.

4.3. A Large Variety in Silicate Absorption

Figure 10 shows, from top to bottom, a series of starburst spec-tra, sorted by increasing absorption of the 9.8 !m silicate reso-nance. Since the peak of the resonance coincides with a minimumbetween the 7Y9 and 11Y13 !m PAH emission complexes, the ef-fect of silicate absorption only becomes apparent toward the lowerhalf of the plot. Nevertheless, the figure strikingly illustrates thestrong effect of amorphous silicates on the overall 5Y38 !m spec-tral shape and the large variations evenwithin one class of objects.Figure 10 shows that the 10 !m trough can in fact become a dom-inating feature of the spectral shape of starbursts.

The top two spectra also show a broad emission structure be-ginning at about 16 !m and extending to 21Y25 !m. The shapeof this feature is consistent with that of an 18 !m silicate emis-sion feature, which would make these the first detections of thisfeature in starburst galaxies. However, the emission structurecould also be due to theC!C!C in-plane and out-of-plane bendingmodes of PAHs (Van Kerckhoven et al. 2000). The identificationof this feature will be discussed in a future paper. The remaining

starburst galaxies in Figure 10 do not show evidence for a similaremission structure. The 13Y35 !m spectra, from NGC1222 (top)to NGC 4945 at the bottom, show an increasingly pronounceddepression, peaking at 18.5!m, signaling increasingly strong sili-cate absorption. The latter result is in full agreement with the trendfound for the 9.8 !m silicate feature.Among all the galaxies classified as starbursts, NGC 4945

(bottom spectrum in Fig. 10) is a ‘‘special case,’’ exhibiting byfar the strongest dust obscuration to its nuclear region (Spoonet al. 2000). Based on the IRS spectrum, the apparent opticaldepth in the 9.8 !m silicate feature is at least 4 and may be higher,depending on the choice of the local continuum. Apart fromstrong amorphous silicate absorption, the line of sight also revealsthe presence of a 23!mabsorption feature. Following the analysisof deeply obscured lines of sight toward ULIRG nuclei (Spoonet al. 2006), we attribute the 23 !m feature to crystalline silicates(forsterite). Their detection in NGC 4945 suggests that crystal-line silicates are perhaps a more common component of the ISMand not just limited to ULIRG nuclei.Figure 11 compares the average starburst template from Fig-

ure 6 to NGC 4945, the most extincted source within our sample.For comparison we also show the heavily embedded ULIRGIRAS 08572+3915 from Spoon et al. (2006). The usually rathershallow 18 !m silicate band reduces the continuum by a largefactor compared to the average starburst spectrum. While NGC4945 has a similarly strong silicate absorption feature as theULIRG IRAS 08572+3915, the starburst also shows strong PAHcomplexes near 16.4 and 17.1!m,which are absent in theULIRGspectrum.Crystalline silicates, causing the features at 16 and 19!min IRAS 08572+3915, appear to be absent in the starburst tem-plate, although they are more difficult to discern at our spectralresolution given the presence of PAH features and the [S iii] linein that range.While NGC 4945 shows by far the most extreme silicate ab-

sorption within our sample, it is not an exotic object but definesthe endpoint of a sequence of increasing optical depths. Recently,Dale et al. (2006) reported on the Spitzer SINGS survey of 75nearby galaxies. The typical target in their sample shows onlymodest dust obscuration, consistent with an AV " 1 foregroundscreen and a lack of dense clumps of highly obscured gas and dust.Unfortunately, we cannot distinguish between the local and global

Fig. 10.—IRS spectra of Mrk 52, NGC 7714, NGC 1222, NGC 7252, NGC2146, NGC 520, NGC 3628, and NGC 4945 arranged to illustrate the gradualeffect of increasing silicate absorption from top to bottom. The flux densitiesFk have been arbitrarily scaled for better comparison.

Fig. 11.—Comparison of the average starburst template (dotted line) fromFig. 6 to the most extincted source within our sample, NGC 4945 (dashed line),and to the ULIRG IRAS 08572+3915 (solid line), chosen as an extreme case ofextinction from the sample of Spoon et al. (2006). The spectra have been nor-malized at 14.5 !m.

BRANDL ET AL.1140 Vol. 653

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(3) The CO Spectral Line Energy Distribution - Mrk 231

(van der Werf et al. 2010)

•Transitions up to J = 14→13 observed•Model: dense gas clump in a 560 pc radius star-forming disk exposed to UV + X-rays from an AGN heating the inner 160 pc.

A&A 518, L42 (2010)

Fig. 2. Luminosities of CO lines from Mrk 231. Filled symbols repre-sent measurements from the SPIRE FTS spectrum, while ground-basedmeasurements are denoted with open symbols. Coloured lines indicatetwo model PDR components (red and green lines) and an XDR compo-nent (blue line). The sum of these three components is indicated by theblack line and fits the CO measurements. In the legend, n denotes thenumber density of hydrogen nuclei (n = nH+2nH2 ) in cm!3, G0 denotesthe incident UV flux in units of 1.6 " 10!3 erg s!1 cm!2 for the PDRs,and FX the incident X-ray flux for the XDR. The legend also indicatesthe relative emitting areas of the three components.

the fraction of gas within 0.3 pc from an O5 star in Mrk 231to much less than the #50% required to produce the highestCO lines with a high excitation PDR. This problem does notexist for the XDR model, which predicts most of the dust to becooler.

3.2. XDR chemistry

The extraordinarily luminous emission from the molecular ionsH2O+ and OH+ reveals the chemical signature of an XDR.Assuming that the emission arises from a disk with 160 pc ra-dius (as derived above), we can derive column densities in theupper levels of the relevant transitions, which results in valuesNup # 3.0!3.5 " 1013 cm!2 for both species. Modeling the nu-clear molecular gas disk in Mrk 231 with a radius of 520 kpc andH2 mass of 2.2 " 109 M$ then results in lower limits to the to-tal OH+ and H2O+ abundances relative to H2 of #2"10!10 in thecentral 160 pc. Given the short radiative lifetimes (<60 s) of theupper levels involved, most H2O+ and OH+ molecules will be inthe ground state, and total abundances will exceed these lowerlimits by large factors. Such abundances require an e!cient andpenetrative source of ionization in the molecular gas, since theproduction of OH+ is mainly driven by H+ + O % O+ + H fol-lowed by O+ +H2 % OH+ +H and H+ +OH% OH+ +H. In themodels considered in Sect. 3.1, the key species H+ is 2!3 ordersof magnitude more abundant in the XDR model than in the highexcitation PDR, and the OH+ abundance is larger by a compara-ble amount. A similar argument can be made for H2O+, which isformed by OH+ + H2 % H2O+ + H.

The extraordinary luminosity of the OH+ and H2O+ (andto a lesser extent CH+) lines in Mrk 231 is underlined by a

comparison with the SPIRE spectrum of the Orion bar PDR(Habart et al. 2010), which shows no trace of OH+ or H2O+,and only weak CH+ emission, while in Mrk 231 the lines areonly a factor 2!3 fainter than the CO lines. While enhanced cos-mic ray fluxes in a starburst environment will increase the de-gree of ionization and hence the production of OH+ and H2O+

in a PDR, they do not elevate the temperatures to the level re-quired to produce the highest CO lines (Meijerink et al. 2006). Itis thus the combination of strong high-J CO lines and high OH+

and H2O+ abundances that reveals X-ray driven excitation andchemistry in Mrk 231.

4. Outlook

We have shown that the SPIRE spectrum of Mrk 231 revealsboth PDR and XDR emission lines, and made a separation ofthese components. A key goal of the HerCULES project will beusing this decomposition for a quantitative separation betweenstar formation and black hole accretion as power sources for theinfrared luminosities of dusty galaxies. In the case of Mrk 231,this issue will be addressed in a forthcoming paper (Meijerinket al., in prep.). Data will be obtained as part of the HerCULESprogram for an additional 28 objects, which will enable us to putthe results presented here on a statistically significant footing.

Acknowledgements. We thank Ewine van Dishoeck, Xander Tielens, andThomas Nikola for useful discussions. We especially thank Ed Polehampton,Peter Imhof-Davies and Bruce Swinyard for their help with the FTS data pro-cessing. JF thanks MPE for its hospitality. The Dark Cosmology Centre isfunded by the DNRF. The following institutes have provided hardware and soft-ware elements to the SPIRE project: University of Lethbridge, Canada; NAOC,Beijing, China; CEA Saclay, CEA Grenoble and LAM in France; IFSI, Rome,and University of Padua, Italy; IAC, Tenerife, Spain; Stockholm Observatory,Sweden; Cardi" University, Imperial College London, UCL-MSSL, STFC-RAL, UK ATC Edinburgh, and the University of Sussex in the UK. Fundingfor SPIRE has been provided by the national agencies of the participating coun-tries and by internal institute funding: CSA in Canada; NAOC in China; CNES,CNRS, and CEA in France; ASI in Italy; MEC in Spain; Stockholm Observatoryin Sweden; STFC in the UK; and NASA in the USA. Additional funding supportfor some instrument activities has been provided by ESA.

ReferencesAalto, S., Spaans, M., Wiedner, M. C., & Hüttemeister, S. 2007, A&A, 464, 193Boksenberg, A., Carswell, R. F., Allen, D. A., et al. 1977, MNRAS, 178, 451Braito, V., Della Ceca, R., Piconcelli, E., et al. 2004, A&A, 420, 79Davies, R. I., Tacconi, L. J., & Genzel, R. 2004, ApJ, 613, 781Downes, D., & Solomon, P. M. 1998, ApJ, 507, 615Fixsen, D. J., Bennett, C. L., & Mather, J. C. 1999, ApJ, 526, 207González-Alfonso, E., Smith, H. A., Ashby, M. L. N., et al. 2008, ApJ, 675, 303González-Alfonso, E., et al. 2010, A&A, 518, L43Gri!n, M. J., et al. 2010, A&A, 518, L3Habart, E., et al. 2010, A&A, 518, L116Meijerink, R., & Spaans, M. 2005, A&A, 436, 397Meijerink, R., Spaans, M., & Israel, F. P. 2006, ApJ, 650, L103Meijerink, R., Spaans, M., & Israel, F. P. 2007, A&A, 461, 793Papadopoulos, P. P., Isaak, K. G., & van der Werf, P. P. 2007, ApJ, 668, 815Pilbratt, G. L., et al. 2010, A&A, 518, L1Poglitsch, A., et al. 2010, A&A, 518, L2Sanders, D. B., Mazzarella, J. M., Kim, D.-C., Surace, J., & Soifer, B. T. 2003,

ApJS, 126, 1607Spaans, M., & Meijerink, R. 2008, ApJ, 678, L5Swinyard, B. M., et al. 2010, A&A, 518, L4Taylor, G. B., Silver, C. S., Ulvestad, J. S., & Carilli, C. L. 1999, ApJ, 519, 185Veilleux, S., Rupke, D. S. N., Kim, D., et al. 2009, ApJS, 182, 628

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•>95% of the emission centered on a single, resolved source•If the molecular gas is distributed in a rotating disk, Mdyn = 3.3 x 109 Msun

(Inami et al. 2010; Stierwalt et al. 2013)

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Figure 2. Molecular-line spectra at the two nuclei of Arp 220. Continuum has been subtracted. The right ordinate is the fractional absorption depth fa with respectto the continuum intensity of each nucleus, and can be converted to optical depth via !a = ! log(!fa). The 1" noise is 32, 12, and 28 mJy beam!1 for HCO+(4–3),(3–2), and CO(3–2), respectively. The CO(3–2) spectra overplotted in black are from 0.""5 resolution data in Sakamoto et al. (2008). The systemic velocity that weestimated for each nucleus is shown by a dashed line.

of 3.7 (or >1.4 when allowing for 1" error) toward the easternnucleus. The total line absorption

!!a dV integrated over the

velocities where !a ! 0.1 is, in the order of HCO+(4–3),HCO+(3–2), and CO(3–2), 75, 150, and 50 km s!1 towardthe eastern nucleus and 85, 82, and 15 km s!1 toward thewestern. The absorption profiles vary among the lines, morenotably toward the western nucleus. There are two caveatsto keep in mind. First, our observations provide lower limitsto total absorption because emission from the absorbing gasitself and from ambient (i.e., nonabsorbing) gas in our beammasks the absorption. Second, our synthesized beam is larger forHCO+(3–2) than for HCO+(4–3) and CO(3–2).

The absorbing column density NH2 is on the order of1022 cm!2 for the CO absorption and 1023 cm!2 for the HCO+

if the absorbing gas is in 50 K LTE and abundances are[CO/H2] # 10!4 and [HCO+/H2] # 10!8. The high abun-dance adopted for HCO+ is consistent with recent models ofthe photon-dominated region (PDR) and the X-ray-dominatedregion (XDR) (Meijerink & Spaans 2005), as well as models ofHCO+ enhancement in protostellar outflow shocks (Rawlingset al. 2004). The absorbing gas will have a depth on the orderof 1–10 pc if its density is at the CO(3–2) critical density of104 cm!3. The actual absorbing material can be distributedover a longer length if it is inhomogeneous as in numericalhydrodynamical simulations of multiphase gas obscuring activenuclei (e.g., Wada & Norman 2002). The column densities ofthe absorbing gas calculated here have large uncertainties be-cause the absorption data provide lower limits, gas excitationconditions are poorly known, and molecular abundances cansignificantly vary. Still, it is possible for the absorbing columndensity to take these values because they are smaller than thetotal column density of the nuclei ($1025 cm!2; S08). The re-markable detection of vibrationally excited HCN in absorptionin Arp 220 (Salter et al. 2008) suggests that the excitation ofHCO+ may also be more unusual than described by our sim-ple analysis. More realistic models of the excitation and ra-

diative transfer of both molecules, including radiative couplingto the intense infrared continuum, will be explored in a futurepublication.

3.4. Gas Motion—Rotation and Outflow

Rotation of gas around each nucleus is evident in the CO(3–2)position–velocity diagrams (Figure 3) as the negative to positiveshift of V ! Vsys across the nucleus. The CO emission rangesfrom 5000 to 5900 km s!1 around the western nucleus. Thehigh-velocity gas around 5800–5900 km s!1 is detected at >3"for the first time. The fall-off of the rotation velocity away fromthe nucleus agrees with the CO(2–1) observations by Downes &Eckart (2007) and suggests truncation of the mass distributionof the nucleus beyond a "100 pc core.

Radial motion of gas, in addition to the rotation, is suggestedby the blueshifted absorption mentioned above. The absorbinggas must be in front of the continuum core of each nucleusand moving away from it toward us. Although absorptiononly tells us the gas motion along our lines of sight to thecontinuum nuclei, the detection of similar blueshifted absorptionon both nuclei whose nuclear disks are misaligned suggeststhat the gas motion is more likely a radial outflow in mostdirections than noncircular motion only along our sight lines.The redshifted emission in the P Cygni profiles must thencontain emission from gas on the far side of the continuum nucleiand moving away from the nuclei and us. The other emissionin the spectra must be from the gas rotating in the nucleardisks.

The typical outflow velocity of the absorbing gas is probablyabout 100 km s!1 because the deepest absorption in the P Cygnispectra is between (Vsys ! 45) and (Vsys ! 135) km s!1. High-velocity components up to $500 km s!1 are suggested towardthe western nucleus by the HCO+(3–2) absorption with largevelocity offsets. The absorbing gas is turbulent or has a velocitygradient across each continuum nucleus because absorption isalso seen at slightly positive velocities with respect to systemic.

No. 2, 2009 P CYGNI PROFILES TOWARD ARP 220 NUCLEI L105

Table 1Observed Parameters of Arp 220

Parameter East West All ftota

SCO(3!2) (Jy km s!1) 250 301 (7.0 ± 1.1) " 102 0.23 ± 0.06SHCO+(4!3) (Jy km s!1) 19 61 (1.1 ± 0.2) " 102 1.0 ± 0.3SHCO+(3!2) (Jy km s!1) 35 63 (1.0 ± 0.1) " 102 0.64 ± 0.13S0.87 mm (Jy) 0.18 0.49 0.68 ± 0.10 1.0 ± 0.2S1.1 mm (Jy) 0.07 0.20 0.27 ± 0.03 0.82 ± 0.15CO(3–2) max Tb (K) 58 (66) 40 (47) · · · · · ·HCO+(4–3) max Tb (K) 16 (24) 26 (34) · · · · · ·HCO+(3–2) max Tb (K) 12 (17) 15 (21) · · · · · ·0.87 mm max Tb (K) 16 (23) 42 (50) · · · · · ·1.1 mm max Tb (K) 7 (12) 20 (25) · · · · · ·Vsys(radio, LSR) (km s!1) 5415 ± 15 5355 ± 15 · · · · · ·

Notes. We conservatively estimate flux-calibration 1! errors to be 10% and 15% for 0.86 and 1.1 mm bands, respectively. Line fluxes areintegrated over 4800–5900 km s!1 (5200–5900 km s!1 for HCO+(3–2)) including both emission and absorption. Peak brightness temperaturesare measured at or around each nucleus without correcting for beam dilution or missing flux. The first one is Rayleigh–Jeans Tb and the secondin parentheses is Planck Tb.a Fraction of the flux (or flux density) recovered in our 0.##3 resolution data with respect to the total flux of Arp 220. The total fluxes are 106 ± 23Jy km s!1 for HCO+(4–3) (Greve et al. 2009), (3.0 ± 0.6) " 103 Jy km s!1 for CO(3–2) in the central 14##(Wiedner et al. 2002), 152 ± 26 Jykm s!1 for HCO+(3–2) (this work; SMA subcompact array observations with the minimum projected baseline of 9 m), 0.7 ± 0.1 Jy for 0.87mm continuum (S08 and references therein), and 0.33 ± 0.05 Jy for 1.1 mm continuum (this work; SMA subcompact configuration).

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Figure 1. Integrated intensity maps of HCO+(4–3), (3–2), and CO(3–2) linesand 0.87 mm continuum. The line moment-0 maps were made using positivesignals only. The lowest contour is at 2! and the nth contour is at 9.5n, 4.3n,and 10.3n Jy beam!1 km s!1 in (a), (b), and (c), respectively, and at 14.4n1.5

mJy beam!1 in (d). Crosses are at continuum peaks and filled ellipses showsynthesized beams.

there in conditions favorable for the transitions of high criticaldensity.

3.2. P Cygni Line Profiles at Both Nuclei

The line spectra at the two nuclei, in Figure 2, show bothemission and absorption. In particular, all the observed linesexcept CO(3–2) toward the western nucleus show blueshiftedabsorption and redshifted emission, i.e., P Cygni profiles. Thisfeature is absent in our calibrator spectra.

The western nucleus has a narrow, deep absorption feature at5340 km s!1 and emission at higher and lower velocities in our

CO data. This profile is consistent with the previous CO(2–1)profile in Downes & Eckart (2007) except that the absorption isdeeper in our CO(3–2) data. The deep absorption is almost at thesystemic velocity of the nucleus, marked with a dashed line inFigure 2, as expected if the CO-absorbing gas is in circularrotation in the western nuclear disk. The systemic velocity,Vsys(W) = 5355 ± 15 km s!1, was estimated from the centralvelocity of the CO profile in the low-resolution (0.##5) data of S08,measured at half-maximum. The HCO+(4–3) and (3–2) linesare also in absorption at the systemic velocity. In addition, bothHCO+ lines have absorption at blueshifted velocities while mostof their emission appears at redshifted velocities. The HCO+(3–2) profile shows multiple velocity components in absorption;the minima at 4890, 5220, and 5340 km s!1 are at 3.6! 4.6! ,and 4.6! , respectively. Absorption deeper than 3! is seen downto Vsys(W) ! 495 km s!1.

The eastern nucleus has very asymmetric line profiles withrespect to its systemic velocity, which we estimate to beVsys(E) = 5415 ± 15 km s!1 in the same way as for the westernnucleus. In all the observed lines, absorption is predominantlyseen at or blueward of the systemic velocity while most emissionis in redshifted velocities, making the profiles P Cygni type. Thedeepest absorption is at 45–75 km s!1 below Vsys(E) and is 2.8! ,4.5! , and 3.0! in HCO+(4–3), (3–2), and CO(3–2), respectively.

3.3. Absorption Depth

Absorption (i.e., negative intensity) in our continuum-subtracted data means absorption of the continuum, becauseline self-absorption alone cannot produce negative intensity. Inorder to absorb continuum the excitation temperature of theforeground gas must be below the brightness temperature ofthe background continuum. The (deconvolved) western nucleushas Tb= 90–160 K and a size of 50–80 pc in 860 µm contin-uum (S08; for the spectral energy distribution (SED) of the twonuclei see Matsushita et al. 2009). The high Tb, due to highdust opacity in submillimeter, makes line absorption againstdust continuum more likely to occur (and deeper) than at longerwavelengths.

The maximum apparent optical depth of each absorption lineis in the range of 0.4–1.0 except the HCO+(3–2) optical depth

•gas disk has an inner region with significant outflow motion and outer region dominated by rotation

•dense absorbing gas deep with the gas disk•source of acceleration - mechanical energy from supernovae or radiation pressure from an AGN

(Sakamoto et al. 2009)

Summary

• U\LIRGs have complex substructures which vary as a function of wavelength

• Mid-IR diagnostics has provided insight to the relative importance of starburst and AGN to the IR output of U/LIRGs

• At the high luminosity end, the star formation in U/LIRGs is compact and highly dust-enshrouded

• ALMA will enable the study of the star-forming molecular gas and dust, the importance of starburst and AGN heating to the chemistry observed in molecular clouds, the extent of the star-forming regions, the feedback in the form of molecular outflows