A&A 379, 823-844 (2001)
DOI: 10.1051/0004-6361:20011377

Infrared to millimetre photometry of ultra-luminous IR galaxies[*],[*]: New evidence favouring a 3-stage dust model

U. Klaas1 - M. Haas1 - S. A. H. Müller1,2 - R. Chini2 - B. Schulz3 - I. Coulson4 -
H. Hippelein1 - K. Wilke1 - M. Albrecht2 - D. Lemke1

1 - Max-Planck-Institut für Astronomie (MPIA), Königstuhl 17, 69117 Heidelberg, Germany
2 - Astronomisches Institut, Ruhr-Universität Bochum, 44780 Bochum, Germany
3 - ISO Data Centre, Astrophysics Division, Space Science Dep. of ESA, Villafranca, PO Box 50727, 28080 Madrid, Spain
4 - Joint Astronomy Centre, 660 N. Aohuku Place, University Park, Hilo 96720, Hawaii, USA

Received 22 December 2000 / Accepted 24 September 2001

Infrared to millimetre spectral energy distributions (SEDs) have been obtained for 41 bright ultra-luminous infrared galaxies (ULIRGs). The observations were carried out with ISOPHOT between 10 and 200 $\mu $m and supplemented for 16 sources with JCMT/SCUBA at 450 and 850 $\mu $m and with SEST at 1.3 mm. In addition, seven sources were observed at 1.2 and 2.2 $\mu $m with the 2.2 m telescope on Calar Alto. These new SEDs represent the most complete set of infrared photometric templates obtained so far on ULIRGs in the local universe. The SEDs peak at 60-100 $\mu $m and show often a quite shallow Rayleigh-Jeans tail. Fits with one single modified blackbody yield a high FIR opacity and small dust emissivity exponent $\beta < 2$. However, this concept leads to conflicts with several other observational constraints, like the low PAH extinction or the extended filamentary optical morphology. A more consistent picture is obtained using several dust components with $\beta = 2$, low to moderate FIR opacity and cool (50 K > T > 30 K) to cold (30 K > T > 10 K) temperatures. This provides evidence for two dust stages, the cool starburst dominated one and the cold cirrus-like one. The third stage with several hundred Kelvin warm dust is identified in the AGN dominated ULIRGs, showing up as a NIR-MIR power-law flux increase. While AGNs and SBs appear indistinguishable at FIR and submm wavelengths, they differ in the NIR-MIR. This suggests that the cool FIR emitting dust is not related to the AGN, and that the AGN only powers the warm and hot dust. In comparison with optical and MIR spectroscopy, a criterion based on the SED shapes and the NIR colours is established to reveal AGNs among ULIRGs. Also the possibility of recognising evolutionary trends among the ULIRGs via the relative amounts of cold, cool and warm dust components is investigated.

Key words: infrared: galaxies - galaxies: active, evolution, fundamental parameters, photometry, starbursts

1 Introduction

Ultra-luminous IR galaxies (ULIRGs) emit the bulk of their energy in the mid- and far-infrared with $L_{\rm IR} > 10^{\rm 12}$ $L_{\odot}$. Optical and near-infrared images show a disturbed morphology and signatures of interaction and merging. Spectra show emission lines characteristic for strong starbursts, shocks (LINERs) and in some cases also Seyfert types (AGNs). Sanders et al. (1988a) suggested that the ULIRGs form an evolutionary link between starburst galaxies and quasars. The current literature and the research state are excellently reviewed by Sanders & Mirabel (1996), and most recently by Genzel & Cesarsky (2000).

Here we focus on the following topics:

While the luminosity of an ULIRG originates from warm T > 100 K and cool T > 30 K dust, the bulk of the dust in normal spiral galaxies is much colder at $T \approx 15$ K. Do ULIRGs also have such cold cirrus-like dust, or is all the dust already heated up to T > 30 K? This is observationally related to: how does the IR SED continue beyond 100 $\mu $m into the submm range? For bright but less luminous IRAS galaxies a three component model - with cirrus, starburst and AGN dust - has been proposed and successfully applied by Rowan-Robinson & Crawford (1989).
What can we learn about the evolution of ULIRGs from such a cold dust stage? How far can the merging state be inferred from the relative amounts of cool and cold dust? For example, with respect to the Sanders et al. (1988a) conjecture this means: do AGN-ULIRGs have less cold dust than starburst dominated SB-ULIRGs?
Can the dust mass and opacity provide clues in favour of or against hidden AGNs? How does this compare to the near- and mid-IR emission of ULIRGs whose optical and MIR spectra favor their power to be dominated by either AGNs or SBs? Observationally this relates to measuring the NIR-MIR SEDs, which are currently known only for a few ULIRGs.
These topics introduce new aspects for exploring the ULIRGs. In this paper, based on IR to mm photometry, we present a sample of observational templates, and confine the discussion to a few basic implications, which may further stimulate more detailed analyses.

Table 1: Measured flux densities in Jy as a function of wavelength in $\mu $m. The objects are listed by IRAS name, except when another name is more common, then the IRAS name is given in the footnote. The statistical errors from signal processing are typically about 5% at 100-150 $\mu $m, and reach in the other passbands 10-20% of the quoted fluxes depending mainly on the source/sky brightness contrast as seen by the apertures used. In case of uncertainties larger than 30% 3$\sigma $ upper limits are listed. The apertures used are 30 $\hbox {$^{\prime \prime }$ }\times 30\hbox {$^{\prime \prime }$ }$ at 1.2 and 2.2 $\mu $m, circular with 52 $\hbox {$^{\prime \prime }$ }$ diameter at 10-25 $\mu $m, 46 $\hbox {$^{\prime \prime }$ }\times 46\hbox {$^{\prime \prime }$ }$ at 60 and 90 $\mu $m, 184 $\hbox {$^{\prime \prime }$ }\times 184\hbox {$^{\prime \prime }$ }$ at 120-200 $\mu $m, and circular with 8 $\hbox {$^{\prime \prime }$ }$, 15 $\hbox {$^{\prime \prime }$ }$, 24 $\hbox {$^{\prime \prime }$ }$ diameters at 450, 850, 1300 $\mu $m, respectively.

1.2 2.2 10 12 15 25 60 90 120 150 180 200 450 850 1300

    0.057 <0.120 0.09 0.25     7.50 5.34 4.21 3.00      
00262+4251* 0.0054 0.0055 <0.087 0.043 0.10 0.36 3.36 2.60 2.90 2.40 2.40 <2.10      
00406-3127                 0.77 0.63 0.44 0.32      
03068-5346     <0.063 0.040 0.05 0.15 3.60 3.34 3.33 2.78 1.72 1.35      
03158+4227     <0.165 0.085 <0.21 0.25 5.25 3.88 3.19 1.75 <1.79 1.05      
03538-6432                 1.37 1.07 0.85 0.62      
04232+1436 0.0055 0.0106 <0.108 0.065 0.07 0.39 3.88 4.14 3.50 2.30 1.60 1.10      
05189-2524     0.600 0.800 1.00 3.08 13.02 10.70 10.00 6.80 4.36 3.84     <0.036
06035-7102     0.090 0.130 0.15 0.53 5.86 5.05 5.15 2.90 2.01 1.35      
06206-6315     <0.072 0.040 0.11 0.24 4.84 4.20 4.57 2.90 1.87 1.50      
12112+0305     0.065 0.100 0.20 0.44 9.94 8.95 8.31 5.87 3.93 2.93      
MK231a     1.425 2.400 2.90 8.66 31.68 27.34 24.32 14.74 9.75 6.88      
MK273b     0.100 0.250 0.50 2.07 27.45 23.78 19.99 13.10 8.69 7.40      
MK463c     0.425 0.580 0.80 1.79 2.33 1.94 1.60 0.99 <0.84 0.45      
14348-1447     <0.096 0.108 <0.27 0.42 8.21 7.23 6.66 5.61 3.76 2.69 0.21 0.024 <0.009
14378-3651     <0.276 <0.399 <0.36 0.41 8.53 7.49 5.90 3.60 2.80 2.45     <0.010
15245+1019     <0.084 0.051 <0.18 0.22 5.25 5.43 5.58 3.60 1.94 1.31      
15250+3609     <0.267 0.238 0.56 1.31 8.63 6.66 4.16 2.37 1.60 1.16      
Arp 220d     0.147 0.600 1.14 8.28 113.35 111.50 109.00 87.89 63.95 54.81      
15462-0450*     0.064 0.100 0.15 0.38 2.81     2.80 2.44 <2.5      
16090-0139     <0.252 0.100 <0.21 0.20 6.70 5.80 4.00 2.75 2.00 1.21 <0.13 <0.017  
NGC 6240e     0.259 0.750 1.00 3.31 23.60 26.70 25.90 18.91 12.73 9.00 1.00 0.150  
17208-0014     0.080 0.200 0.25 1.32 32.22 31.90 30.00 23.00 17.50 12.50 1.07 0.155  
17463+5806                 0.63 0.58 0.37 0.23      
18090+0130* 0.0200 0.0195 0.187 0.300 <0.38 0.59     16.63 14.50 10.48 9.74      
18470+3234 0.0027 0.0036 <0.081 0.135 0.15 0.44 4.15 3.55 3.20 2.45 1.70 1.25      
19254-7245f     0.123 0.200 0.40 1.32 5.57 5.26 4.27 3.06 2.30 1.61     0.012
19458+0944*     <0.396 <0.540 <0.77 0.26   5.82   7.61   7.05      
20046-0623 0.0039 0.0040 <0.084 0.125 <0.39 0.47 4.50 4.71 3.69 3.18 2.79 1.40 <0.16 <0.026  
20087-0308     <0.117 0.070 0.10 0.25 5.54 4.85 5.96 4.05 2.95 2.30 <0.10 <0.010  
20100-4156     <0.090 0.090 0.14 0.34 6.53 5.11 4.51 3.36 2.09 1.86     <0.014
20414-1651     <0.168 <0.117 <0.18 0.25 6.16 4.22 3.60 2.56 1.54 <1.35      
ESO 286-19g     <0.204 0.275 0.50 2.11 13.90 9.26 6.73 4.14 2.56 1.80      
21130-4446     <0.090 0.055 <0.45 0.15 4.03 3.60 4.50 3.94 3.14 2.90     <0.023
21504-0628 0.0055 0.0057 0.040 0.050 0.15 0.28 3.33 2.82 2.80 1.52 1.26 0.70 <0.11 <0.016  
22491-1808     <0.063 0.080 0.11 0.50 6.80 5.45 3.55 2.70 1.95 1.70 <0.20 0.019  
ESO 148-2h     0.126 0.250 0.36 1.55 12.50 11.36 10.00 6.40 4.70 3.50      
23230-6926     0.035 0.060 0.10 0.29 4.80 3.98 4.10 3.00 2.00 1.70     0.010
23365+3604 0.0075 0.0065 <0.120 0.080 0.15 0.65 7.78 6.76 6.75 4.81 3.35 3.01 0.17 0.020  
23389-6139     <0.075 0.031 <0.26 0.20 3.88 3.30 4.03 2.83 1.99 1.52     <0.005
23515-2917                 0.80 0.64 0.48 0.33      
* ISO data partly distorted and uncertain; a 12540+5708; b 13428+5608; c 13536+1836; d 15327+2340; e 16504+0228; f "Super-Antennae''; g 20551-4250; h 23128-5919.

2 The sample

From the Queen Mary & Westfield College IRAS galaxy catalogue (QIGC, 17711 entries, Rowan-Robinson et al. 1991) we chose the 48 brightest sources with (1) $F_{\rm 60\mu m} >
3$ Jy and (2) $L_{\rm FIR} = L_{\rm 40-120~\mu m} > 10^{\rm 12}$ $L_{\odot}$ (based on $H_{\rm0} = 50$ kms-1/Mpc). Throughout this paper, however, we use $H_{\rm0} = 75$ kms-1/Mpc, thus, some of the sources reach only about 10 $^{\rm 11.5}$ $L_{\odot}$. In the course of the ISO mission 37 out of these 48 sources could be observed (77%), all objects having z < 0.2. Thus, the sample is representative for the nearby bright ULIRGs. In coordination with the other ISO Guaranteed Time programmes the two slightly less luminous objects NGC 6240 and Mrk 463 were included. Later we observed four more objects with higher redshifts ( 0.4 > z > 0.2, $F_{\rm 60~\mu m} \approx 1$ Jy) which were visible to the satellite.

3 Observations and data reduction

3.1 ISOPHOT 10-200 $\mu $m data

The observations were carried out with ISOPHOT (Lemke et al. 1996), the photometer on board ISO (Kessler et al. 1996). A description of the ISOPHOT observation templates (AOTs) is given in Laureijs et al. (2000). The MIR 10 to 25 $\mu $m observations were done in triangular chopped mode (AOT P03) with a 52 $\hbox {$^{\prime \prime }$ }$ aperture and a chopper throw of 60 $\hbox {$^{\prime \prime }$ }$ which matched the compact size of the targets well. The FIR 60 to 200 $\mu $m observations were performed with the array cameras C100 and C200 in the "sparse map'' off-on mode (AOT P37-39) with background positions 3-4 $\hbox{$^\prime$ }$ north of the source. Beam sizes used for photometry are given in Table 1. The four faint sources at higher redshift were observed in mini-map mode (AOT P22) in a $4 \times 2$ raster and restricted to the filters at 120, 150, 180, and 200 $\mu $m.

The data were reduced using the PHT Interactive Analysis tool (PIA[*]) (V7.3.3e) in standard processing mode, together with the calibration data set V4.0 (ISOPHOT Data User Manual V4.0, Laureijs et al. 1998). This includes correction for non-linearity of the electronics, deglitching (removal of data disturbed by cosmic particle events), and correction for signal dependence on the reset interval time. To handle the signal transients in the time series of bright FCS illuminations, only the last half was taken, when the values approached the final signal level. The calibration of detector responsivity and its changes was performed using associated measurements of the thermal fine calibration source (FCS) on board.

For the chopped measurements with the P1 and P2 detectors at wavelengths 10-25 $\mu $m we inspected the sequence of chopper plateaux and removed outliers, with the criterion that the uncertainty was larger than the average three sigma of the measurements. By this procedure typically the first chopper plateau of a series and those with residual glitches were removed.

Finally the fluxes were corrected for aperture/beam size effects. For the $3 \times 3$ pixel C100 array (60 and 90 $\mu $m) the fluxes were derived with two methods using (1) only the central pixel (46 $\hbox {$^{\prime \prime }$ }\times 46\hbox {$^{\prime \prime }$ }$) and (2) the whole array (138 $\hbox{$^{\prime\prime}$ }\times 138\hbox{$^{\prime\prime}$ }$). Both methods yield essentially the same flux (with some higher noise for the whole array), providing evidence that the objects are point-like and not resolved in the FIR.

The error propagation in ISOPHOT data reduction is described in Laureijs & Klaas (1999). The statistical errors derived from signal processing are about 5-20%, depending on the wavelength range and object brightness (see caption to Table 1), but systematic errors due to absolute calibration accuracy are estimated to be 30% (Klaas et al. 1998b). To account for the overall uncertainty in the signal derivation as well as relative and absolute photometric calibration we have adopted a general photometric uncertainty of 30%.

3.2 SCUBA 450 and 850 $\mu $m data

The 450 and 850 $\mu $m observations were obtained on July 1st and 3rd, 1999, using the Submillimetre Common User Bolometer Array (SCUBA, Holland et al. 1999) at the James Clerk Maxwell Telescope (JCMT) on Mauna Kea, Hawaii. In general, photometry with 1 $\hbox{$^\prime$ }$ chopper throw, mini-jiggle and beam switching was performed. We used the SCUBA narrow band filters in order to minimize possible CO line contributions (see also Sect. 4.1). The atmospheric transmission was determined every hour from measurement series at various zenith distances (sky dips). The conditions were excellent and stable ( $\tau_{850~\mu{\rm m}} = 0.134 \pm 0.002$ on July 1st and $0.209 \pm 0.003$ on July 3rd, respectively). $\tau_{450~\mu{\rm m}}$ was extrapolated from the 850 $\mu $m skydips ( $\tau_{450~\mu{\rm m}} = 0.65 \pm 0.01$ on July 1st and $1.09 \pm 0.02$ on July 3rd, respectively). Mars and Uranus served as standard calibrators, and repeated photometry measurements gave a reproducibility of better than 5% and 14% at 850 $\mu $m and 450 $\mu $m, respectively. The observing time per source was 10-20 min (actually the dome carousel driver was out of order, so that the targets could only be observed during their oblique rising passage across the dome slit pointing towards south-east). The data were reduced using the SCUBA User Reduction Facility (SURF) with special emphasis on identifying noisy bolometer pixels and removing the sky noise determined from off-source pixels.

In order to check for possible extended submm flux, jiggle maps covering 2 $\hbox{$^\prime$ }$ were obtained for three sources. These sources were NGC 6240 and 17208-0014, both revealing a submm flux excess (as discussed below), and 20046-0623 providing the same good upper limits for different beams. None of the sources seems to be extended at the spatial resolution of about 8 $\hbox {$^{\prime \prime }$ }$ and 15 $\hbox {$^{\prime \prime }$ }$ FWHM, respectively. These test cases suggest that for the other sources the bulk of submm flux is also contained within our photometry beams and that some possible extended submm flux is small and lies within the total adopted uncertainties of about 20-30%, in particular since NGC6240 and 17208-0014 lie in the low redshift range of our sample.

3.3 SEST 1300 $\mu $m data

The $1300~\mu$m observations were obtained between July 4th and 9th, 1999, using the MPIfR single channel facility bolometer (Kreysa 1990) at the Swedish ESO Submillimetre Telescope (SEST) on La Silla, Chile. The system provides a beamsize of 24 $\hbox {$^{\prime \prime }$ }$ and was used in the ON-OFF mode with dual beam switching and a beam separation of 68 $\hbox {$^{\prime \prime }$ }$ in azimuth. Pointing and focus were checked at least every two hours and the pointing accuracy was always better than 4 $\hbox {$^{\prime \prime }$ }$. The atmospheric transmission as determined by sky dips was stable at about $\tau_{1300} \approx 0.1$. Uranus served as a standard calibrator. The observing time per source was about 1 hour (on-source). From repeated calibrator measurements we estimate the absolute accuracy to be about 20-30%.

\end{figure} Figure 1: Spectral energy distributions of ultra-luminous IR galaxies, ordered along the columns by RA. Upper limits are indicated by a downward arrow and in the case of ISO, SCUBA and SEST measurements in addition by open symbols. The redshifts and, if known, the optical spectral types are listed. The wavelength and frequency ranges are as observed and not corrected with regard to the rest frame of the objects. The dotted lines represent modified blackbody fits. Emissivity exponent $\beta $ and opacity $\tau_{\rm 100~ \mu m}$ are listed in the plots. The resulting temperatures T are corrected for redshift. For those sources with no submm/mm data available we used an average $\beta = \langle$1.6$\rangle$. The solid lines indicate the shapes of the SEDs, dash-dotted lines are used where they are uncertain. At NIR-MIR wavelengths the SED shape is obtained by connecting the data points, and at long FIR and submm wavelengths it is determined by the modified blackbody fit. Luminosities were derived by integrating the SEDs between 10 and 1000 $\mu $m. Dashed lines in the mm range indicate synchrotron extrapolations from cm observations. The literature data are from Sanders et al. (1988a,b), Crawford et al. (1996), Marx et al. (1994), Murphy et al. (1996), Rigopoulou et al. (1996), Duc et al. (1997), Rigopoulou et al. (1999) and Lisenfeld et al. (2000). For some sources the NIR fluxes of the central regions are plotted with a "$\times $'' (Diameter = 1 $\hbox {$^{\prime \prime }$ }$, from Scoville et al. 2000, except 5 $\hbox{$.\!\!^{\prime\prime}$ }$5 for 18090+0130 from Hill et al. 1988).

\end{figure} Figure 1: continued.

\end{figure} Figure 1: continued.

\end{figure} Figure 1: continued.

3.4 MAGIC 1.2 and 2.2 $\mu $m data

J and K-band images were obtained at the Calar Alto 2.2 m telescope on August 21 and 22, 2000, using the $256 \times 256$ pixel NIR camera MAGIC (Herbst et al. 1993) with a pixel scale of 0 $\hbox{$.\!\!^{\prime\prime}$ }$6416. The observations were performed in standard dithering mode with a total exposure time per source of about 10 and 30 min in J and K, respectively. The seeing was about 2 $\hbox {$^{\prime \prime }$ }$ preventing the identification of spatial details in the sources (see Appendix A). But the conditions were photometric, in particular during the second night. The UKIRT standard stars FS2, FS27 and FS35 were observed for flux calibration. The data reduction followed the standard procedures. Based on the cross calibration of the standard stars, we estimate the photometric accuracy to be about 15%. In the case of IRAS18090+0130 the uncertainty is somewhat higher due to superposition of two stars.

4 Results

4.1 Spectral energy distributions

The fluxes from our observations are listed in Table 1. The 60 and 90 $\mu $m values agree within 20% with the IRAS 60 and 100 $\mu $m ones. For many sources, where IRAS provided only upper limits at 12 or 25 $\mu $m, now the 10-25 $\mu $m fluxes could be measured. Also, for 16 sources out of 41, submm/mm fluxes and good upper limits are provided, as well as NIR 1.2 and 2.2 $\mu $m fluxes for seven sources.

The spectral energy distributions (SEDs, as measured, not corrected for redshift) are shown in Fig. 1, supplemented by literature data. The remarkable features of the SEDs are:

For each galaxy the maximum of the SED can now be clearly determined. It lies between 60 and 100 $\mu $m. The ISOPHOT long wavelength filters beyond 100 $\mu $m clearly outline the beginning of the Rayleigh-Jeans branch. In some cases (e.g. NGC 6240, 17208-0014) a somewhat plateau-like broad maximum is revealed, suggesting a high opacity even in the FIR and/or the presence of various cool to cold dust components (see also Fig. 3);
For 20 sources the shape of the Rayleigh-Jeans branch can now be determined by the mm and submm data points. As discussed in Sects. 3.1 and 3.2, within the error budget the effects of different instrumental beams and a possible contribution of extended flux appear to be negligible. In particular for those cases (e.g. NGC 6240, 17208-0014), where ISOPHOT already indicated a broad maximum, the submm measurements confirm the high flux prediction from the 100-200 $\mu $m measurements. Also, where ISOPHOT indicated a steep Rayleigh-Jeans branch (e.g. 16090-0139, 23365+3604, 23389-6139) the mm and submm fluxes and upper limits are low, demonstrating the consistency between the instruments;
Shortward of the maximum at around 60-100 $\mu $m the SEDs exhibit two basic shapes:
a flat NIR plateau followed by a jump-like flux increase at about 10 $\mu $m (e.g. Arp 220 and 12112+0305). In some cases also indications of PAH emission around 7.7 $\mu $m and silicate 9.7 $\mu $m absorption features (e.g. Arp 220 and NGC 6240) are recognized;
a power-law-like flux increase from the NIR to the MIR (e.g. Mrk 463) or FIR (e.g. Mrk 231). Spectral PAH emission or silicate absorption features may be present, e.g. for Mrk 463, compare also with Fig. 3 in Rigopoulou et al. (1999), but due to the high continuum level they are diluted in the broad band photometry.

4.2 CO line and synchrotron emission contamination of submm/mm fluxes

The determination of the Rayleigh Jeans branch allows a detailed analysis with respect to the dust emissivity $\lambda^{-\beta}$ and the opacity $\tau_{\rm 100~ \mu m}$, as carried out in the next section, under the condition that the emission is of thermal nature. Therefore, beforehand one has to check possible contamination of the submm and mm fluxes by CO lines and/or synchrotron emission:

For all our sources observed at 1300 $\mu $m with SEST the CO(2-1) line at 230 GHz (1304 $\mu $m) moves out of the filter band pass due to their redshifts of $z \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
$\displaystyle ...;
The CO(3-2) line at 345 GHz (869 $\mu $m) may contribute to the SCUBA 850 $\mu $m fluxes for those sources with $z \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
$\displaystyle .... Direct CO(3-2) observations are only available for Arp220 (Mauersberger et al. 2000). Therefore, we estimated CO(3-2) line strengths from the CO(1-0) fluxes (Solomon 1997; Downes & Solomon 1998; Gao & Solomon 1999), adopting an intensity ratio $R = I_{\rm CO\,(3-2)}$/ $I_{\rm CO\,(1-0)}
= 0.9$ as found for Arp 220. The strength of the redshifted CO(3-2) line, however, is reduced according to the transmission of the SCUBA 850 $\mu $m filter which lies between 55% and 30% for our sources with $z \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
$\displaystyle .... It turns out that only four sources have a CO(3-2) line contribution to the 850 $\mu $m flux which exceeds 5%: Arp 220 and Mrk273 (both 11%), Mrk231 (22%), and NGC6240 (35%). These contributions are smaller or of the order of our adopted photometric uncertainties. Therefore we decided not to correct for the CO(3-2) line contribution;
In 20 cases with available radio fluxes (Condon et al. 1990; Crawford et al. 1996) extrapolations with spectral indices between 0.5 and 1.0 from the cm range towards shorter wavelengths are far below (<10%) our measured submm and mm fluxes (many of the extrapolations lie even below the range plotted in Fig. 1). The only exception is Mrk463 (with two Seyfert nuclei) where the 1.3 mm flux is variable, thus dominated by synchrotron emission (Chini et al. 1989a; Marx et al. 1994).
Our conclusion is therefore, that, like the FIR emission, the bulk of the submm/mm flux is of thermal nature being emitted by dust.

4.3 Dust parameters

In order to characterise the dust emission, the SEDs are fitted with modified blackbodies. Such fits, however, are not unique. They rely largely on the mass absorption coefficient $\kappa$ and its wavelength dependence $\beta $, both still being a matter of debate. Values of $\beta $ between 1 and 2 are commonly used (e.g. Hildebrand 1983). In case of a flat Rayleigh-Jeans tail the SEDs can also be modelled by several dust components. Since the interpretation of the dust emission as well as the derivation of the dust mass depend on the blackbody models used, we investigate the two main cases. They represent simplified formalisms, each relying on implicit assumptions, and a realistic description probably lies between these two extremes. In the following two subsections the FIR-submm range is investigated, and the MIR part is addressed in the third subsection.

4.3.1 Single modified blackbody

We used the following model:

S_{\lambda} = {B_{\lambda}(T)} \cdot (1 - {\rm e}^{- \tau_{\lambda} }) $, with $\end{displaymath} (1)

\tau_{\lambda} = \tau_{ 100~\mu{\rm m}} \cdot (100~\mu{\rm m} / \lambda )^{\beta}.
\end{displaymath} (2)

The SEDs are fitted between 60 and 1300 $\mu $m with a single modified blackbody leaving T, $\beta $ and $\tau_{\rm 100~ \mu m}$ free (minimising $\chi$$^{\rm 2}$ in a grid search).
\end{figure} Figure 2: Distribution of $\tau_{ 100~\mu\rm m}$ versus $\beta $ for those sources with measured submm/mm fluxes. Different symbols correspond to optical spectral types as in Table 4: + Seyfert1, $\times $ Seyfert2, $\blacklozenge $ LINER, $\blacksquare $ HII/SB, and $\bullet $ for not classified; the quasar PG0050+124 denoted by $\ast $ is also included for comparison (from Haas et al. 2000a). Arrows indicate lower limit cases.

Although the emissivity exponent $\beta $ is a free parameter in Eq. (2), it is kept constant over the whole wavelength range. In order to keep the parameter space under control, we have decided not to introduce a $\lambda $ dependence of $\beta $.

The relation of the parameters $\beta $, $\tau_{ 100~\mu\rm m}$ and T with the SED shapes is:

$\bullet $
$\beta $ corresponds to the slope of the Rayleigh-Jeans tail, whether it is flat (small $\beta $) or steep (large $\beta $);
$\bullet $
$\tau_{ 100~\mu\rm m}$ determines, whether the peak plateau of the SED is narrow (small $\tau_{ 100~\mu\rm m}$) or broad (large $\tau_{ 100~\mu\rm m}$);
$\bullet $
Like for normal blackbodies (Wien's displacement law), T is related to the wavelength, at which the SED maximum is located.
Fits were performed only for the "mm-subsample'', i.e. for those 22 sources with submm/mm fluxes available. We also used the 8 cases of upper limits[*], treating them formally as detections, and after the fitting procedure taking into account that they provide lower limits for $\beta $. In some cases, where the ISO 60 or 90 $\mu $m fluxes show a large error or deviate strongly from the IRAS 60 and 100 $\mu $m fluxes, we also included the IRAS data (e.g. for 16090-0139).

The fitted parameters $\beta $, $\tau_{ 100~\mu\rm m}$ and T are listed in Table 2, together with $\chi^{2}$. A visual impression of the quality of the fits is given in Fig. 1. In summary, the results for the mm-subsample are:

$\bullet $
    1.2 < $\beta $ $\mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
$\displaystyle ... 2.0 (Fig. 2)
$\bullet $
    0.5 < $\tau_{ 100~\mu\rm m}$ < 5 (Fig. 2)
$\bullet $
    50 K < T < 70 K
with an uncertainty of about 0.2 for $\beta $, 25% for $\tau $ and 5-10 K for T. The quoted uncertainties have been estimated considering "bent'' SEDs with the FIR fluxes reduced and the submm fluxes increased by their measurement uncertainty and vice versa. It should also be noted that for sources with significant MIR flux (an extreme case is Mrk463) additional uncertainties in the determination of T and/or $\tau_{\rm 100~ \mu m}$ are introduced when using the full 60 $\mu $m flux, which is partially due to a dust component not considered here.

The quoted parameter values should be considered with some tolerance and their interdependence borne in mind:

$\bullet $
For two SEDs with maxima located at similar wavelengths, T decreases with $\beta $. Due to the steeper fall-off at long wavelengths the width of the peak plateau can only be maintained, if the maximum is shifted to longer wavelengths. A complementary pair are 19254-7245 ( $\beta = 1.2$) and 20046-0623 ( $\beta = 1.8$).
$\bullet $
For two SEDs with similar $\beta $, but largely different $\tau_{\rm 100~ \mu m}$, T increases with $\tau_{\rm 100~ \mu m}$. Due to the broader width, but similar shape on the long wavelength side, the maximum is shifted towards shorter wavelengths. A complementary pair is Mrk231 and 14348-1447.
On the other hand, there seems to be no correlation between $\tau_{\rm 100~ \mu m}$ and $\beta $ as can be seen from Fig. 2.

The dust parameters were determined via Eq. (1) only for the mm-subsample. For the remaining sources with wavelength coverage limited to 200 $\mu $m (IR-subsample) $\beta $ could not be fitted reliably (as we found from tests with the mm-subsample using only the 60-200 $\mu $m fluxes). For the IR-subsample we kept $\beta $ fixed using the average value $\beta = 1.6$ derived from the mm-subsample. Then $\tau_{ 100~\mu\rm m}$ and T could be determined reasonably well from the 60-200 $\mu $m fluxes alone. (Exceptions are 00262+4251, 15462-0450, 18090+0130 and 19458+0944 which have less complete spectral coverage due to bad quality measurements as flagged in Table 1. In these cases $\tau_{ 100~\mu\rm m}$ was fixed to 6.0). The resulting values lie in the same range as for the mm-subsample (Table 2 and Fig. 1). As a check, we fitted also $\tau_{ 100~\mu\rm m}$ and T of the mm-subsample with a fixed $\beta = 1.6$ using only the 60-200 $\mu $m fluxes. The results are basically consistent with those obtained from the longer wavelength coverage, except for the sources with extremely low or high true $\beta $. Hence, in the discussion below we can mostly use the full sample, and only where $\beta $ plays a role, we confine it to the mm-subsample.

4.3.2 Multiple modified blackbodies

As derived in the previous section, for the majority of the "mm-subsample'' sources (11 out of 14, not having lower limits for $\beta $, one of them having $\tau_{100~\mu{\rm m}} = 5$, Fig. 2) it is not possible to fit the FIR-submm SEDs properly with one single modified blackbody with an emissivity law of $\lambda $$^{\rm -2}$, rather the superposition of two or more modified blackbodies is required.

In the low opacity case Eq. (1) can be approximated by

S_{\lambda} = {B_{\lambda}(T)} \cdot \tau_{\lambda} \propto {B_{\lambda}(T)} \cdot \lambda ^{- \beta}.
\end{displaymath} (3)

Based on Eq. (3), implicitly assuming $\tau_{100~\mu{\rm m}} \ll 1$, we fitted several modified blackbodies with free T and fixed $\beta = 2$ to all observed SEDs, as illustrated for some examples in Fig. 3. They have temperatures in the range between 30 and 50 K (cool), and between 10 and 30 K (cold). However, the decomposition into the various components is not unique, because of the sparse data coverage in the sub-mm range.
\rotatebox{90} {\resizebox{6.5cm}{!} {\inclu...
...90} {\resizebox{6.5cm}{!} {\includegraphics {ms10599f3f.eps} } } }\end{figure} Figure 3: Examples for fits with several optically thin ( $\tau_{100~\mu\rm m} \ll$ 1) modified blackbodies (BBs) with an emissivity law of $\lambda $$^{\rm -2}$. Top row: superposition of two BBs; middle row: possible superposition of three BBs; bottom row: fits can also be done for those SEDs with no submm data: approximation by two BBs.

In the case of multiple blackbodies no direct conclusion about the opacity $\tau_{ 100~\mu\rm m}$ can be drawn (nevertheless, in Sect. 5.1.2 below, $\tau_{ 100~\mu\rm m}$ will be constrained using CO data). The most realistic case might be that of several blackbodies with $\beta \approx 2$, and a range of opacities from low to partly high.

4.3.3 Description of the NIR-MIR emitting dust

The two basic SED shapes in the NIR-MIR outlined in Sect. 4.1 can be formally fitted by a superposition of several warm dust components. For the cases with flat NIR plateau the maximum temperatures are about 100-150 K (e.g. Klaas et al. 1997, 1998a). The power-law-like SEDs can be approximated by a suite of blackbodies up to the dust grain evaporation temperatures of about 1000-1500 K (the hotter the blackbody the less dust mass is involved). Modelling of the continuum is hampered by the presence of strong spectral features like PAH emission and silicate absorption. Using higher spectral resolution, Laurent et al. (2000) investigated this spectral part quantitatively.

4.4 Luminosities

Table 2 lists the luminosities derived within various bandpasses in the rest frame of the objects by integrating the spectral energy distribution as outlined by the thick solid and dash-dotted lines shown in Fig. 1 for the indicated wavelength ranges. On the Rayleigh-Jeans tail and around the SED maximum this comprises the single blackbody curve obtained with the Eq. (1) fit, and shortward thereof the lines connecting the data values by linear interpolation.

The total IR-submm luminosity $L_{\rm 10-1000~\mu m}$ is dominated by the FIR in the wavelength range 40-150 $\mu $m, while the 150-1000 $\mu $m submm range plays a minor role ( $L_{\rm 150-1000~\mu m} < 0.1 \cdot L_{\rm 40-150~\mu m}$) as well as the 10-40 $\mu $m MIR range (except for Mrk463 and the z > 0.3 sources which are MIR dominant). The luminosities $L_{\rm 8-1000~\mu m}$ extrapolated from the four IRAS bands (formula cf. Table 1 in Sanders & Mirabel 1996) typically slightly overestimate our IR-submm luminosity values by about 15%; nevertheless this is still a good agreement.

The MIR/FIR luminosity ratio has a median value of about 0.3. Thus, the sample of bright nearby ULIRGs preferentially comprises objects with cool MIR/FIR colours (compared with quasars having $L_{\rm MIR}$/ $L_{\rm FIR} > 1$, cf. Haas et al. 2000a). Though the luminosity range of the ULIRG sample spans about one decade, there is no trend of luminosity with optical spectral type or MIR/FIR colours.

For sources without submm/mm observations available, the submm luminosity $L_{\rm submm} = L_{\rm 150-1000~ \mu m}$ is extrapolated using the average value $\beta = 1.6$ (Sect. 4.3.1), while the actual $L_{\rm submm}$ depends on the actual value of $\beta $. A check on the mm-subsample shows that, for the case of minimum $\beta \approx 1.2$ or maximum $\beta \approx 2$, $L_{\rm submm}$ obtained using the average $\beta = 1.6$ can deviate from the true value by factors of 2 and 0.5, respectively.

4.5 Dust masses

In order to derive the dust masses, we used the standard approach based on Hildebrand (1983) and further developed by various authors (e.g. Chini et al. 1986; Krügel et al. 1990):

M_{\rm dust} = \frac{ 4 \pi \cdot a^3 \cdot \rho } { 3 \cdot...{ D ^{2} \cdot S_{850 ~\mu\rm m} } { B_{850 ~\mu\rm m}(T) }
\end{displaymath} (4)

$\bullet $ average grain size a = 0.1 $\mu $m;
$\bullet $ grain density $\rho$ = 3 g cm$^{\rm -3}$;
$\bullet $ distance D;
$\bullet $ flux $S_{850~ \mu\rm m}$ (in the restframe of the object);
$\bullet $ Planck function $B_{850~ \mu\rm m}$(T) at restframe 850 $\mu $m; and
$\bullet $ dust grain emission efficiency $Q_{\rm em}$(850 $\mu $m, $\beta $) where

Q_{\rm em} (\lambda, \beta ) = Q_0 \cdot a \cdot \{ 250 / \lambda \} ^ \beta
\end{displaymath} (5)

with $Q_{\rm0} = 40$ cm$^{\rm -1}$ for $\lambda = 250~\mu$m.
The first fraction in Eq. (4) is conveniently summarized yielding

M_{\rm dust} = \frac { 1 } { \kappa_{850~ \mu\rm m} ( \beta ...
...{ D ^{2} \cdot S_{850~ \mu\rm m} } { {B_{850~ \mu\rm m}(T)} }
\end{displaymath} (6)


\kappa_{850~ \mu\rm m}(\beta) = \kappa_{850~ \mu\rm m}^{\beta = 2} \cdot \{ 250~ \mu\rm m / 850~ \mu\rm m \}^{\beta - 2}
\end{displaymath} (7)


$\bullet $      $\kappa_{850~ \mu\rm m}^{\beta = 2} = 0.865~{\rm cm}^{2}\,{\rm g}^{-1}$

which corresponds to

$\bullet $      $\kappa_{1300~ \mu\rm m}^{\beta = 2} = 0.4~{\rm cm}^{2}\,\rm g^{-1}$.
Our "choice'' of $\kappa$ is consistent with that favoured by Krügel et al. (1990, see their Eq. (10)), and Lisenfeld et al. (2000). Note that we account for $\beta $ in the wavelength dependence of $Q_{\rm em}$ and $\kappa$, respectively. In the case of $\beta < 2$ this leads to a larger value of $\kappa$, hence to dust masses which are smaller than for $\beta = 2$.


Table 2: Infrared luminosities, dust temperatures, dust masses and IR source sizes. The luminosity distance is determined as DL = c/$H_{\rm0}$*( $z + z^{\rm 2}$/2), i.e. $q_{\rm0} = 0$, with $H_{\rm0} = 75$ kms-1/Mpc. The infrared luminosities $L_{\rm MIR}$ (3-40$~\mu$m), $L_{\rm FIR}$ (40-150$~\mu$m) and $L_{\rm submm}$ (150-1000$~\mu$m) are determined from the SED curves shown in Fig. 1, by integrating over the indicated wavelength range shifted to the rest frame of the objects. The dust mass is derived for three models of the IR emission: 1)  $M_{\rm d}^{\rm\tau free}$ for the single modified blackbody with free T, $\tau $ and $\beta $ as described by Eq. (1) (minimising $\chi ^{\rm 2}$) and shown in Fig. 1 ($\beta $ values in brackets $<\,>$ means that a fixed average of $\beta = 1.6$ is used, lower limits are indicated by a > sign). 2)  $M_{\rm d(FIR)}^{\rm\beta = 2}$ for a single modified blackbody with fixed $\beta = 2$ fitted by Eq. (3) to the FIR 60-200 $\mu $m range (for details see Sect. 4.5). 3)  $M_{\rm d(total)}^{\rm\beta = 2}$ for multiple modified blackbodies with fixed $\beta = 2$. The sizes of the dust emitting regions are given for the two cases of a single blackbody as (1) brightness radius $r_{\rm b}$, and (2) as "smallest transparent radius'' $r_{\rm\tau }$, at which $\tau_{\rm 100~\mu m} = 1$ (for details see Sect. 4.6). For comparison M(H$_{\rm 2}$) is listed (using the galactic conversion factor M(H $_{\rm 2}) = 4.6 \cdot L_{\rm CO}$).
Name z D L $_{\rm MIR}$ L $_{\rm FIR}$ L $_{\rm submm}$ T $\tau $ $_{\rm 100}$ $\beta $ $\chi ^{\rm 2}$ M $_{\rm d}^{\rm\tau free}$ r$_{\rm b}$ r$_{\rm b}$ T M $_{\rm d(FIR)}^{\rm\beta = 2}$ r $_{\rm\tau}$ r $_{\rm\tau}$ M $_{\rm d(total)}^{\rm\beta = 2}$ M(H$_{\rm 2}$)
      10-40 $\mu $m 40-150 150-1000                 single BB     several BBs  
    [Mpc] [10$^{\rm 9}$ $L_{\odot}$] [10$^{\rm 9}$ $L_{\odot}$] [10$^{\rm 9}$ $L_{\odot}$] [K]       [10$^{\rm 6}$ $M_{\odot}$] [pc] [ $\hbox {$^{\prime \prime }$ }$] [K] [10$^{\rm 6}$ $M_{\odot}$] [pc] [ $\hbox {$^{\prime \prime }$ }$] [10$^{\rm 6}$ $M_{\odot}$] [10$^{\rm 9}$ $M_{\odot}$]
00199-7426 0.0963 403 307 1139 130 48 2.50 <1.6> 1.07 109 390 0.199 32 160 1097 0.6 549  
00262+4251* 0.0971 407 342 675 77 68 6.00 <1.6> 1.24 28 131 0.066 36 48 602 0.3 762 29a
00406-3127 0.3422 1602   2480 101 53 0.50 <1.6> 1.04 51 688 0.088 43 59 666 0.1 157  
03068-5346 0.0778 323 126 481 37 60 2.50 <1.6> 1.07 22 172 0.109 36 32 488 0.3 152  
03158+4227 0.1343 573 685 1724 71 77 2.00 <1.6> 1.09 26 220 0.079 43 42 563 0.2 115  
03538-6432 0.3100 1431   2877 188 68 5.00 <1.6> 1.06 72 263 0.038 39 110 910 0.1 519  
04232+1436 0.0799 332 244 540 34 51 1.00 <1.6> 1.04 23 275 0.170 37 33 495 0.3 73 41a
05189-2524 0.0425 173 460 500 33 70 2.50 1.4 1.07 15 128 0.152 38 26 443 0.5 448$^{\rm +}$ 23b
06035-7102 0.0794 330 376 785 40 49 0.50 <1.6> 1.09 29 434 0.271 38 35 513 0.3 82 38b
06206-6315 0.0924 386 271 878 54 61 2.00 <1.6> 1.09 30 225 0.120 37 50 611 0.3 130 52b
12112+0305 0.0723 299 264 1052 68 53 1.00 1.5 1.06 42 355 0.244 36 71 731 0.5 856  
Mrk 231 0.0417 170 1215 1219 58 54 1.00 1.9 1.02 22 371 0.449 50 28 456 0.5 131$^{\rm +}$ 35a
Mrk 273 0.0373 152 248 714 47 62 2.00 1.6 1.14 27 208 0.282 36 47 596 0.8 104$^{\rm +}$ 23a
Mrk 463 0.0506 207 347 136 5 52 0.50 2.0 1.06 4 188 0.187 40 5 198 0.2 12  
14348-1447 0.0811 337 294 1037 79 67 5.00 2.0 1.12 37 194 0.118 35 81 781 0.5 219$^{\rm +}$ 64c
14378-3651 0.0676 279 199 651 49 69 5.00 >1.7 1.10 26 139 0.103 36 44 579 0.4 423 15b
15245+1019 0.0756 314 138 621 39 51 1.00 <1.6> 1.15 27 300 0.197 37 38 534 0.4 97  
15250+3609 0.0553 227 304 461 17 59 0.50 1.3 1.10 10 215 0.195 44 11 289 0.3 28  
Arp 220 0.0182 73 190 820 86 61 5.00 1.7 1.09 66 214 0.600 40 32 488 1.4 1479$^{\rm +}$ 32a
15462-0450* 0.1005 422 397 763 109 59 6.00 <1.6> 1.12 52 179 0.087 34 73 741 0.4 1498  
16090-0139 0.1334 569 635 1944 85 49 0.50 >1.9 1.15 43 645 0.233 40 66 708 0.3 182 56a
NGC 6240 0.0245 99 157 347 28 57 2.50 1.5 1.07 21 154 0.320 33 36 520 1.1 580$^{\rm +}$ 37a
17208-0014 0.0424 173 235 1226 107 60 3.00 1.7 1.09 64 274 0.326 34 116 936 1.1 602$^{\rm +}$ 32a
17463+5806 0.3411 1596   2192 87 59 1.00 <1.6> 1.10 40 434 0.056 43 57 658 0.1 142  
18090+0130* 0.0660 273 322 1100 193 46 6.00 <1.6> 1.06 173 319 0.241 28 322 1558 1.2 1746  
18470+3234 0.0788 327 270 537 35 74 6.00 <1.6> 1.07 14 92 0.058 38 26 442 0.3 113  
19254-7245 0.0615 253 375 473 32 72 3.00 1.2 1.04 15 108 0.088 38 24 429 0.3 1265$^{\rm +}$ 35b
19458+0944* 0.0995 418 602 1343 307 46 6.00 <1.6> 1.21 226 373 0.184 28 388 1711 0.8 6405 55a
20046-0623 0.0845 352 344 658 47 60 3.00 >1.8 1.18 26 188 0.110 34 57 654 0.4 123  
20087-0308 0.1055 444 337 1370 103 61 4.00 >2.0 1.10 50 254 0.118 35 110 911 0.4 239$^{\rm +}$ 74a
20100-4156 0.1295 551 622 1919 135 67 3.50 >1.5 1.07 65 252 0.094 40 86 803 0.3 337  
20414-1651 0.0870 363 265 836 38 65 1.50 <1.6> 1.12 19 206 0.117 40 29 471 0.3 88  
ESO 286-19 0.0426 174 272 460 16 60 0.50 <1.6> 1.05 8 225 0.266 44 11 289 0.3 28 22b
21130-4446 0.0925 387 130 772 98 52 2.50 >1.2 1.19 71 262 0.140 32 111 916 0.5 932  
21504-0628 0.0775 322 214 425 21 75 2.50 >1.6 1.10 8 106 0.068 40 16 343 0.2 39  
22491-1808 0.0760 315 268 661 40 73 3.00 1.7 1.10 16 143 0.093 39 32 488 0.3 213$^{\rm +}$ 31b
ESO 148-2 0.0446 182 248 502 33 68 3.00 <1.6> 1.05 17 136 0.154 36 34 504 0.6 83 18b
23230-6926 0.1062 447 327 1038 88 66 4.00 1.5 1.12 42 187 0.086 36 66 705 0.3 1565  
23365+3604 0.0645 266 250 704 47 68 4.00 2.0 1.12 21 159 0.123 36 47 595 0.5 108$^{\rm +}$ 39a
23389-6139 0.0927 388 182 753 56 63 4.00 >1.9 1.12 26 173 0.092 35 59 668 0.4 135  
23515-3127 0.3347 1562   2323 105 50 0.50 <1.6> 1.04 62 759 0.100 42 58 662 0.1 184  

* Measurements partly distorted and uncertain, therefore $\tau $ $_{\rm 100~ \mu m}$ has been fixed to 6.0.
+ Good determination of total dust mass, data used for estimate of gas-to-dust ratio (see Sect. 5.1.2).
a From Solomon et al. (1997), Table 2.
b From Mirabel et al. (1990); recalculated by Gao & Solomon (1999).
c From Sanders et al. (1991); recalculated by Gao & Solomon (1999).

Table 2 lists the dust masses derived from flux values at 850 $\mu $m, where the emission is certainly transparent (and no mass is "hidden''). For the further discussion, three different kinds of dust mass are determined for each galaxy:
(1) For the assumption that the FIR emission can be modelled by one single blackbody with free $\tau $, $\beta $ and T, as described by Eq. (1) and shown in Fig. 1.
For the case of several components with fixed $\beta = 2$ (Eq. (3)) the temperature decomposition is not unique (Sect. 4.3.2), and the uncertainty in the derived total dust mass becomes large, in particular for those sources without any submm/mm data. Therefore Table 2 lists two estimates:
(2) The total dust mass $M_{\rm d(total)}^{\rm\beta = 2}$ associated with the multiple "optically thin'' blackbodies.
(3) The dust mass associated with the "bulk FIR emission component'' $M_{\rm d(FIR)}^{\rm\beta = 2}$.
While $M_{\rm d(total)}^{\rm\beta = 2}$ may be considered as a maximum dust mass, $M_{\rm d(FIR)}^{\rm\beta = 2}$ is derived under the assumption of one single optically thin $\lambda^{\rm -2}$ modified blackbody of $T \approx 35$-40 K fitted to the FIR 60-200 $\mu $m range. It contributes the bulk of the luminosity and may be considered as a firm lower limit for the dust mass, in particular if other colder dust components exist. $M_{\rm d(total)}^{\rm\beta = 2}$ contains contributions from cold dust components, it is typically much larger than $M_{\rm d(FIR)}^{\rm\beta = 2}$.

For all cases, the dust mass does not show any correlation with the total, mid- or far-infrared luminosity. But the dust mass is quite well correlated with the submm luminosity $L_{\rm 150-1000~ \mu\rm m}$.

4.6 Size of dust emitting regions

The smallest possible extent of the FIR emitting region is listed in Table 2. For the case of an opaque blackbody (Eq. (1)), the brightness radius $r_{\rm b}$ is determined via

r_{\rm b} = \left\{ \frac {D ^{2} \cdot S_{100~ \mu\rm m} } ...
...- {\rm e}^{- \tau_{100~ \mu\rm m} }) } \right\}^{\rm 1/2}\cdot
\end{displaymath} (8)

For the case of a transparent blackbody with fixed $\beta = 2$ (Eq. (3)), the low opacity condition $\tau_{ 100~\mu\rm m} < 1$ requires that the extent of the dust emission cannot be smaller than a minimum size. Such a "smallest transparent radius'' can be roughly estimated as follows: the dust mass (Eq. (6)) - as determined for the bulk FIR emitting component at $T \approx 30$-40 K - is distributed homogeneously in a "minimum face-on disk'' of radius $r_{\rm\tau }$. Thereby the following standard conversions are used: 0.1 $M_{\odot}$$\cdot$pc $^{\rm -2} \Leftrightarrow \tau_{\rm V} = 0.4$ (Whittet 1992), and $\tau_{\rm 100~\mu m} = 0.006 \cdot \tau_{\rm V}$ (Mathis et al. 1983). Thus

r_{\tau} = \left\{ \frac {M_{\rm dust} [M_{\odot}]} { \tau_{...
...~ \mu m} \cdot 41.7 \cdot \pi } \right\}^{\rm 1/2} [{\rm pc}].
\end{displaymath} (9)

Typically $r_{\tau}$, for $\tau_{ 100~ \mu\rm m} = 1$, is a factor 2-5 larger than $r_{\rm b}$ (see Table 2), and we will discuss further below whether a dust emitting region as large as that given by $r_{\tau}$ is still consistent with other data.

5 Discussion

5.1 Evidence for cold cirrus-like dust

The observations show that, though the SEDs usually peak at 60-100 $\mu $m, they can stay on a plateau before the onset of the Rayleigh-Jeans tail. The analysis of Sects. 4.3.1 and 4.3.2 showed that the observed SEDs can be modelled by either one single tepid (50-75 K) blackbody with $\beta \le 2$, partly inferring a high opacity even in the FIR, or by several cool to cold transparent dust bodies with $\beta = 2$. Which of these scenarios (or which combination) might more realistically apply for the ULIRGs?

In the following discussion, we will check for observational evidence for the existence of cold cirrus-like dust in ULIRGs being in line with the multiple blackbody model. One prerequisite to allow a fit of the SEDs with several blackbodies is that $\beta \approx 2$. Even in the case of several blackbodies the opacity cannot be low at all, since the derived dust mass and the size of the emission area exceed other constraints. In order to reduce the amount of dust mass and size, the opacity of some of the blackbodies must be increased (whereby the dust temperature also rises, resulting in a lower dust mass and a smaller region). Thus, we end up with the picture of several blackbodies with $\beta = 2$, which are partly optically thick.

5.1.1 Evidence against the single blackbody model

The following arguments can be brought forward against the concept of one single blackbody to describe the FIR-submm SEDs of most of our ULIRGs (i.e all of our mm-subsample except Arp 220):

In general, the optical morphology is disturbed and suggests a variety of patchy dust complexes with less dense regions in between. Since in such a scenario various dust temperatures are also expected, multiple blackbodies appear to provide a better physical description than one single blackbody;
For a variety of objects several authors found observational evidence that $\beta \approx 2$ in the FIR-submm range ( $\lambda > 100$ $\mu $m): $1.8 \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
...\offinterlineskip\halign{\hfil$\scriptscriptstyle ... was found for interstellar dust in our Galaxy (e.g. Mathis et al. 1983; Lagache et al. 1998), $\beta \approx 2.0$ for normal spiral galaxies (e.g. Bianchi et al. 1999) and the archetypal starburst galaxy M82 (e.g. Thuma et al. 2000), and $\beta = 2.0 \pm 0.2$ for active (Markarian) galaxies (Chini et al. 1989b). In contrary to that, Dunne et al. (2000) found an average value $\beta \approx 1.3$ from their SCUBA Local Universe Galaxy Survey. However, one weakness of this study appears to be the restriction to three wavelength points at 60, 100, and 850 $\mu $m without probing the maximum and the start of the Rayleigh Jeans branch in sufficient detail. One common object between their sample and ours is Arp220. While their single temperature fit provides $\beta = 1.2$ our single modified blackbody fit to a much better sampled SED gives $\beta = 1.7$. Also, the spread of $\beta $ between 1.2 and 2.0 in our ULIRG sample would suggest a large variety of dust properties, like grain size distribution, among the same type of object with probably similar evolution histories. It appears much more plausible to explain the SEDs by a composition of dust components all with consistent emissivity laws of $\beta \approx 2$;
We developed the following quantitative method to check the consistency of the derived $\tau_{ 100~\mu\rm m}$ values from the single blackbody fit with other extinction sensitive quantities. One is strength of the PAH 7.7 $\mu $m feature which is ubiquitous in ULIRGs (Rigopoulou et al. 1999). There must be a sufficient number of UV photons to excite the PAHs, however, an intense starburst is not a prerequisite, as can be seen from the wide distribution of PAH emission in our Galaxy (Mattila et al. 1996) and in NGC891 (Mattila et al. 1999). Mattila et al. (1999) found that the PAH distribution is similar to the one of large dust grains and neutral molecular clouds. It is reasonable to assume that the PAH carriers are mixed with other constituents of the ISM, at least on the spatial scales of kpcs we are looking at with the resolution of our observations. If the PAHs are not cospatial with the dust component emitting the bulk of the FIR emission, then there is a second noticable dust component per se.
In the case of all the FIR emission being irradiated by one single blackbody, we would expect a mixture of the PAH carriers with the FIR emitting dust and a strong decrease of the 7.7 $\mu $m PAH strength (published by Rigopoulou et al. 1999, see examples in Fig. 4 and Table 3) relative to the submm and FIR fluxes with increasing $\tau_{ 100~\mu\rm m}$. This decrease should follow the direction of the A$_{\rm V}$-vector in Fig. 5, (using the galactic extinction curve by Mathis et al. 1983). However, Fig. 5 shows that the normalised 7.7 $\mu $m PAH fluxes are largely independent of $\tau_{ 100~\mu\rm m}$, except for Arp 220 which is further discussed in Haas et al. (2001).

\par\includegraphics[width=8.8cm,clip]{ms10599f4a.eps}\\ [4mm]
\end{figure} Figure 4: Examples of ISOPHOT-S 5.8-11.6 $\mu $m spectra for the ULIRG Arp 220 and the comparison galaxy NGC23. The dotted line indicates the continuum subtracted for the estimate of the PAH 7.7 $\mu $m peak flux strength. Arp 220 has a strong silicate 9.7 $\mu $m absorption feature (e.g. Dudley 1999), which is accounted for in the continuum estimate.


Table 3: PAH 7.7 $\mu $m peak flux strength, 100 $\mu $m and 850 $\mu $m continuum fluxes used for the assessment of flux ratios in Fig. 5. Measured fluxes are in Jy. The MIR continuum around the PAH line was subtracted, also the extrapolated synchrotron continuum from the 850 $\mu $m flux. Uncertainties are less than 30%. For the ULIRG sample PAH fluxes are from Rigopoulou et al. (1999), 850 $\mu $m fluxes are from Rigopoulou et al. (1996), Lisenfeld et al. (2000) and from this work. For the sample of normal comparison galaxies PAH fluxes are from this work and 850 $\mu $m fluxes are from Dunne et al. (2000). The 100 $\mu $m fluxes are all from IRAS.
ULIRGs PAH 100 850
  7.7 $\mu $m $\mu $m $\mu $m
05189-2524 0.220 11.73 0.048
UGC5101 0.186 21.24 0.143
12112+0305 0.080 9.73 0.030
Mrk 231 0.289 30.33 0.045
Mrk 273 0.193 24.58 0.077
14348-1447 0.068 7.60 0.016
15250+3609 0.128 5.80 0.027
Arp 220 0.414 126.7 0.744
NGC6240 0.420 28.13 0.137
17208-0014 0.274 35.66 0.119
19254-7245 0.084 5.38 0.029
20100-4156 0.041 5.20 <0.019
22491-1808 0.043 5.06 0.012
23365+3604 0.086 8.00 0.014
23389-6139 0.028 4.33 <0.009
Comparison Galaxies      
Arp 148 0.250 10.99 0.092
MGC+02-04-025 0.380 9.60 0.039
Mrk 331 0.650 20.86 0.132
NGC23 0.560 14.96 0.144
NGC695 0.700 13.80 0.136
NGC1122 0.600 15.15 0.084
NGC1667 0.400 16.54 0.163
NGC5256 0.350 10.35 0.082
NGC5653 0.700 21.86 0.205
NGC5713 1.000 36.27 0.359
NGC5962 0.300 20.79 0.317
NGC6052 0.270 10.18 0.095
NGC7591 0.400 13.52 0.135
NGC7592 0.450 10.50 0.108
NGC7674 0.250 7.91 0.108
NGC7679 0.350 10.65 0.093
UGC2238 0.600 15.22 0.104
UGC2369 0.300 11.10 0.072
UGC2982 0.900 17.32 0.176
UGC8387 0.500 24.90 0.113

\par\includegraphics[width=8.8cm,clip]{ms10599f5a.eps}\\ [4mm]
\end{figure} Figure 5: Normalised 7.7 $\mu $m PAH flux versus opacity $\tau_{ 100~\mu\rm m}$ as formally derived from the single blackbody fits. Different symbols for each spectral type as in Fig. 2. The arrows show the effect of extinction, if the PAH carriers are mixed with the FIR opaque dust. The filled circle and the thick vertical bar on the left-hand side show the range for normal galaxies "NGs'' (see text).

Further clues on the relation between PAH and FIR-submm emitting dust comes from the comparison of the ULIRGs with normal galaxies: we examined the PAH/100 $\mu $m and PAH/850 $\mu $m flux ratios of the ULIRGs with that for normal galaxies, for which both 850 $\mu $m and PAH data are available. Among the SCUBA sample of 104 galaxies obtained by Dunne et al. (2000), 20 sources were also measured with ISOPHOT-S (by several observers) covering a similar area (24 $\hbox {$^{\prime \prime }$ }$) as the SCUBA observations. We derived the PAH7.7 $\mu $m fluxes from the ISO data archive products processed with OLP Version 9.0 (Laureijs et al. 2000). The values are listed in Table 3 and an example for the MIR spectra is shown in Fig. 4. Note, that now we consider only the vertical distribution in Fig. 5 and ignore the horizontal one: the range for the normal galaxies is indicated by the thick vertical bars on the left-hand side. Strikingly the normal galaxies populate a similar PAH/850 $\mu $m range ( $3.9 \pm 2.3$) as the "typical'' ULIRGs ( $3.7 \pm 1.5$, excluding Arp 220), but the PAH/100 $\mu $m range ( $0.03 \pm 0.014$) lies higher than for the ULIRGs ( $0.01 \pm 0.004$). This difference cannot be explained by PAH destruction in the ULIRGs, since it should be reflected in the PAH/850 $\mu $m ratio as well. Neither can this difference be due to extinction, because the shift of the normal galaxies with respect to the ULIRGs must then be larger at 850 $\mu $m than at 100 $\mu $m (see the length and orientation of the extinction arrows in Fig. 5). The difference of the PAH/100 $\mu $m ratio between normal galaxies and ULIRGs can be best explained by the higher dust temperatures of the ULIRGs of about T = 30-60 K compared with T = 15-30 K in normal galaxies, which increases the 100 $\mu $m flux with respect to that at 850 $\mu $m. This behaviour also suggests that the PAH carriers not only trace the dust in active regions, but are even more related to the entire dust content;

We compared the dust mass ( $M_{\rm d}^{\rm\tau free}$ in Table 2) with the molecular gas mass derived from integrated CO luminosities (M(H$_{\rm 2}$) in Table 2). The resulting gas-to-dust mass ratio lies in the range between 500 and 2500, with an average of about $1500 \,\pm \,500$. This is much higher than the canonical value of about 150 for normal galaxies (e.g. Stickel et al. 2000). This discrepancy becomes even higher, when considering the additional contribution of the atomic gas. Since it is unlikely that the CO-to H$_{\rm 2}$ conversion factor for ULIRGs is so much different from that for normal galaxies, a larger dust mass than $M_{\rm d}^{\rm\tau free}$ is required for ULIRGs in order to match the canonical value. Such a suitably high dust mass can be achieved with the combination of several dust components.
In summary, we listed several items where the single blackbody model appears to be insufficient or inconsistent. This can be solved by the application of multiple blackbodies. In order to place several blackbodies underneath the Rayleigh-Jeans tail, $\beta = 2$ seems to provide a reasonable choice.

5.1.2 Constraining the FIR opacity for the multiple blackbody model

For a low FIR opacity and several cold to cool blackbodies with $\beta = 2$, the estimated total dust mass $M_{\rm d(total)}^{\rm\beta = 2}$ is considerably higher (on average by a factor of ten) than for the single optically thick blackbody (cf. Table 2). This is because of the contribution from the additional cold component ( $M \propto T^{-6}$ for $\beta = 2$). For the mm-subsample $M_{\rm d(total)}^{\rm\beta = 2}$ lies in the range $\approx$10$^{\rm 8}$-10$^{\rm 9}$ $M_{\odot}$. As mentioned in Sect. 4.5, the uncertainty for $M_{\rm d(total)}^{\rm\beta = 2}$ is quite large, because the decomposition into several blackbodies is not unique. Nevertheless, when comparing the $M_{\rm d(total)}^{\rm\beta = 2}$ for the most reliable fits (marked by "+'' in Table 2) with the molecular gas mass derived from integrated CO luminosities (Table 2), the gas-to-dust mass ratio lies in the range between 30 and 300 with an average of $165 \pm 120$, close to the canonical value of about 150. The multiple blackbody model seems to yield a consistent total dust mass.
Stricter constraints on the dust mass can be obtained, when comparing it with the extension of CO gas inferred from interferometric data (Downes & Solomon 1998). In order to keep the opacity, i.e. column density, low, the extent of the dust region must exceed a minimum size. As a check, we distribute the dust mass associated with the bulk FIR emission, $M_{\rm d(FIR)}^{\rm\beta = 2}$ (Sect. 4.5) evenly within a disk of minimal radius $r_{\tau = 1}$, so that $\tau_{ 100~ \mu\rm m} = 1$ ( $r_{\rm\tau }$ in Table 2). In each object the radius $r(\tau_{\rm 100~\mu m} = 1)$ is smaller than the radius of the optical-NIR image sizes. With the possible exception of Arp 220, the extent of low opacity dust is consistent with the size of the CO disk, where available (from Tables 3 and 4 in Downes & Solomon 1998): Mrk 231:0 $\hbox{$.\!\!^{\prime\prime}$ }$9, Mrk 273:0 $\hbox{$.\!\!^{\prime\prime}$ }$9-3 $\hbox{$.\!\!^{\prime\prime}$ }$1, 17208-0014:1 $\hbox{$.\!\!^{\prime\prime}$ }$8, 23365+3604:1 $\hbox{$.\!\!^{\prime\prime}$ }$0. For $\tau_{ 100~ \mu\rm m} = 0.3$, corresponding to $A_{\rm V} \approx 50$, the minimal radius $r_{\rm\tau }$ is about a factor two larger than listed in Table 2, and exceeds the CO disk sizes. If the dust and the molecular gas are well mixed, then the opacity of the bulk FIR dust component cannot be low ( $\tau_{ 100~ \mu\rm m} \le 0.3$), rather it must be "moderately'' high ( $\tau^{\rm bulk}_{ 100~ \mu\rm m} \approx 1$, corresponding to $A_{\rm V} \approx 150$) in many ULIRGs. Furthermore, with regard to our previous findings on the PAH/850 $\mu $m and PAH/100 $\mu $m ratios, which required relatively low extinction in the MIR (except for Arp 220), we conclude that the bulk FIR dust component cannot contain the majority of the PAH carriers. The relative constancy of the PAH/850 $\mu $m flux ratio with regard to normal galaxies is preserved, since the cold component dominates at 850 $\mu $m, as can be seen from the examples in Fig. 3.

In summary, an emissivity exponent $\beta = 2$ works well for the ULIRGs, as in normal spirals and Markarian galaxies. In some ULIRGs the single component fits already give $\beta = 2$, and for those with $\beta_{\rm fitted} < 2$ several blackbodies with $\beta = 2$ can be used. However, in order to match the dust mass and the size of the emitting regions with the constraints from the CO observations, the opacity of the bulk FIR dust component might be moderately high ( $0.3 \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
...iptstyle ...).

5.1.3 The proposed scenario: Cold cirrus dust and cool, partly opaque starburst dust

The previous discussion suggests that the FIR-submm SEDs are composed of several (at least two) $\lambda^{\rm -2}$ modified blackbodies with cool to cold temperatures. Actually, the dust might exhibit a continuous temperature range, but, when applying a formal model, within the uncertainties, a restriction to a few components works best. These dust components represent two stages: (1) the "cirrus'' and (2) "starburst'' stages.

The nuclear regions are probably surrounded by clumpy "Super-Orion'' complexes providing the heating power for the starburst dust. The gaps in between the cloud complexes and the outer regions with lower density and weaker interstellar radiation field host the cirrus-like cold dust. In this picture the PAH emission comes preferentially from the submm emitting dust. Some fraction of the FIR peak dust emission in the nucleus or in dense clouds can be opaque in the FIR.

5.2 Starburst and AGN heated dust

In this section we assess for each optical spectral class (see Table 4) the typical shape of the IR-to-mm SEDs. While the SEDs of AGNs and SBs appear to be indistinguishable at FIR and submm wavelengths, they differ in the NIR-MIR. This suggests that the cool FIR emitting dust is not connected to the AGN, and that the AGN only powers the warm and hot dust. The difference at short IR wavelengths provides a NIR diagnostic tool based on the J-K colours, in order to reveal the presence of an AGN (with red colours, while SBs have blue colours). This tool is successful in the sense that if an AGN is seen by other diagnostics (optical or MIR spectroscopy or X-rays), then in most cases it is also seen via the NIR colours.


Table 4: Morphologies and spectroscopic classifications. The morphologies are from optical and NIR images as referenced. The extinction EB-V has been compiled from H $_{\rm\alpha }$/H $_{\rm\beta }$ ratios. The optical spectral types distinguish between Seyfert 1, Seyfert 2, LINER (galaxies with Low Ionisation Nuclear Emission Regions), and HII/starburst types, as classified in the respective references. The MIR spectral types distinguish between AGN=Seyfert1+Seyfert2 and SB=LINER+HII types (Lutz et al. 1999). The NIR classifications are derived from the J-K colours (AGN, SB, SB/n=SB for total galaxy and red 1 $\hbox {$^{\prime \prime }$ }$ nucleus, SB/f=SB for total galaxy and 1 $\hbox {$^{\prime \prime }$ }$ nucleus). The X-ray evidence for AGN is taken from Risaliti et al. (2000).

Morph EB-V Optical MIR NIR X-ray
  . [mag] Sp-type type class AGN

double m1   LINER? s1   SB  
00262+4251 merger m2,*       SB*  
00406-3127r1 merger m6   Seyf 2 s6      
03068-5346 double?DSS          
03158+4227 compactm2         No
03538-6432r1 compactDSS          
04232+1436 merger *       AGN*  
05189-2524 merger m3 2.03 Seyf 2 s3 AGNa AGN Yes
06035-7102 double m1 1.34 HII/SB s1 SB SB  
06206-6315 double m1 2.06 Seyf 2 s1 SB AGN  
12112+0305 double m3 0.92 LINER s3 SB SB/n  
Mrk 231 merger m3 0.70 Seyf 1 s3 AGN AGN Yes
Mrk 273 merger m3 1.22 Sy2 s3/LINs8 AGN SB/n No
Mrk 463 double m4 0.60 Seyf 2 s4 AGN AGN Yes
14348-1447 merger m3 1.15 LINER s3 SB SB/n  
14378-3651 merger m1 1.26 LINER s1   SB  
15245+1019 merger m2          
15250+3609 ring gal. m3 0.76 LINER s3 SB SB  
Arp 220 merger m3 1.05 LINER s3 SB SB/n No
15462-0450 merger m2 0.60 Seyf 1 s5   AGN  
16090-0139 merger m1,2 1.55 LINER s1,5   SB  
NGC 6240 double m5 1.65 LINER s5 SB SB/n Yes
17208-0014 merger m1,2 1.75 HII/SB s1 SB SB/n No
17463+5806r1 compactDSS          
18090+0130 merger *       SB*  
18470+3234 double m2,* 1.13 HII/SB s5   SB*  
19254-7245 double m1,2 2.92 Seyf 2 s1 AGN AGN Yes
19458+0944 double m2          
20046-0623 double m1,2,*   LINER *   SB  
20087-0308 merger m1,2 1.43 LINER s1   SB  
20100-4156 double m1 1.12 HII/SB s1 SB SB  
20414-1651 merger? m1,2 1.31 LINER s1      
ESO 286-19 merger m1 1.07 LINER s1 SB SB No
21130-4446 merger m1 0.50 HII/SB s1   SB  
21504-0628 merger m1,*   LINER *   SB*  
22491-1808 merger m1 0.81 HII/SB s1 SB SB/f  
ESO 148-2 merger m1 1.01 HII/SB s1 SB SB No
23230-6926 merger m1 1.49 LINER s1   SB  
23365+3604 merger m7,* 0.88 LINER s7 SB SB*  
23389-6139 double m1 3.06 Seyf 2 s1   SB  
23515-3127r1 compactDSS   Seyf 2 s6      

r1 z > 0.3; * this work, see Appendix; e0 Boksenberg et al. (1977), Krabbe et al. (1997); m1,s1 Duc et al. (1997); m2 Murphy et al. (1996); m3,s3 Sanders et al. (1988a); m4 Sanders et al. (1988b); s4 Véron-Cetty & Véron (1985); m5 Fried & Schulz (1983); s5 Kim et al. (1995) and Veilleux et al. (1995); m6 Clements et al. (1996); s6 Allen et al. (1991); m7, s7 Klaas & Elsässer (1991); s8 Colina et al. (1999); a Clavel et al. (2000).

5.2.1 Seyferts

Our sample contains eight galaxies (with z < 0.3) which have optical Seyfert spectra, i.e. are sources known to house an AGN. Two are Seyfert1s and six are Seyfert2s (see Table 4).

The SEDs of the Seyferts show a variety of shapes in the NIR and MIR range:

$\bullet $
On the one hand, a power-law-like flux increase is seen in the Seyfert1s (Mrk231 and 15462-0450) and the Seyfert2s 05189-2524, 19254-7245 and Mrk463 (and probably also 06206-6315);
$\bullet $
On the other hand, Mrk273 (a LINER nucleus with an off-nucleus Seyfert2 nebula, cf. Colina et al. 1999) has a flat NIR flux plateau with a steep rise at about 10 $\mu $m;
$\bullet $
23389-6139, showing a damped power-law flux increase, appears to be in between these extremes.
The power-law shape might be attributed to the central, relatively unobscured AGN. The red J-K colours might not be caused by extinction only, rather the hottest dust clouds at a temperature just below the evaporation temperature of the dust of about 1500 K, corresponding to a peak wavelength of about 2 $\mu $m, might be seen.

In a torus geometry the flat sources like 23389-6139 could be seen more edge-on, so that the hot AGN heated dust is more hidden. Then the flat NIR flux plateau - with blue colours - may be understood in terms of light scattered at bipolar cones. Alternatively, young hot stars in the host galaxies might contribute significantly to the NIR flux, shifting the colours towards the blue. Also, the AGN could be weak compared with the strength of the starbursts; perhaps Mrk273 belongs to this group.

A somewhat indefinite case is the double source 06206-6315, which is classified by Duc et al. (1997) as a Seyfert2, but as starburst via PAH diagnostics (Lutz et al. 1999; Rigopoulou et al. 1999), although both spectra look quite noisy. The steep optical to NIR 2.2 $\mu $m flux increase argues in favor of a relatively unobscured AGN.

In the FIR and submm range the Seyfert SEDs also exhibit some diversity, though only a relatively moderate one. Examples are Mrk463 and 19254-7245, both with double nuclei (and similar power-law NIR flux increase): while Mrk463 only shows little FIR flux compared to the MIR flux, 19254-7245 has a strong FIR bump and even a considerable amount of mm flux due to cold dust. A similar difference is found for Mrk231 and 05189-2524, both with a "single'' nucleus, i.e. probably mergers in an advanced state: both show a strong FIR bump, but Mrk231 has a steeper Rayleigh-Jeans tail and less cold dust than 05189-2524. Thus, the host galaxies of AGNs may or may not have cold dust.

5.2.2 LINERs and HII/SBs

Our sample contains 11 LINERs (Low Ionisation Nuclear Emission Regions) and 9 HII/SBs. In general, the SEDs look very similar for both optical spectral types. Therefore, we discuss them together here.

The SEDs appear quite homogeneous from the NIR to the FIR, with some diversity in the submm range: all SEDs show a relatively flat NIR flux plateau with a faint 2 $\mu $m bump, PAH emission and/or Si 9.7 $\mu $m absorption features, then a steep rise in flux at about 10 $\mu $m, i.e. a 10 $\mu $m knee. The flux peaks in the FIR at about 60-100 $\mu $m. The only variety appears in the submm range: some sources show a steep Rayleigh-Jeans tail, others exhibit a flatter tail providing evidence for additional cold dust. Typical sources with and without cold dust are 23230-6926 and 23365+3604, respectively, among the LINERs, and 17208-0014 and ESO286-19, respectively, among the HII/SBs.

The homogeneity and extreme similarity of the NIR to FIR SEDs suggests that the dust in LINERs and HII/SBs has similar properties, in particular concerning the spatial distribution and heating mechanism. In fact, this similarity places the LINERs closer to HII/SB than AGN dominated ULIRGs. A similar conclusion (that LINERs are not AGN dominated) was drawn from the PAH diagnostics (Lutz et al. 1999).

5.2.3 The SED shapes and $\mathsfsl{J}$- $\mathsfsl{K}$ colours as diagnostic tool

\par\includegraphics[width=7.7cm,clip]{ms10599f6a.eps}\\ [4mm]
\end{figure} Figure 6: Two-colour diagrams 2.2 $\mu $m/1.25 $\mu $m versus 25 $\mu $m/60 $\mu $m (top) and versus 120 $\mu $m/200 $\mu $m (bottom). Different symbols for each spectral type as in Fig. 2.

\end{figure} Figure 7: Histogram showing the correspondence of the NIR SED shape classification (via the J-K colour) with the optical emission line diagnostic and MIR (PAH class) spectroscopic classifications.

The NIR-MIR power-law flux increase for the AGN-type ULIRGs and the flat NIR flux plateau with a 10 $\mu $m knee for the SB-type ULIRGs suggests the utilization of these two different SED shapes for a classification scheme. Earlier investigations (e.g. Hill et al. 1988) showed that red NIR colours provide a high probability of finding an AGN among moderately luminous IRAS galaxies. As a quantitative characterisation measure for the SED shapes we use the J-K colours: J-$K \ge 2$ mag identifies power-law shape SEDs; this value corresponds to a 2.2 $\mu $m/1.25 $\mu $m flux ratio of about 2.34.

Although longer MIR wavelengths are less sensitive to extinction, this range suffers from possible confusion by strong PAH emission and 9.7 $\mu $m silicate absorption which complicates the analysis; detailed work has been done by several authors, e.g. Laurent et al. (2000), Imanishi & Dudley (2000) and Tran et al. (2001).

As shown in Fig. 6, the J-K colours (2.2 $\mu $m/1.25 $\mu $m flux ratios) provide a much clearer separation of AGN- and SB-ULIRGs than the MIR colours (25 $\mu $m/60 $\mu $m) used to identify warm AGN-type objects (de Grijp et al. 1985), to say nothing of the FIR colours (120 $\mu $m/200 $\mu $m).

In Fig. 7 we show in histograms the correspondence of this NIR photometry classification scheme with the spectroscopic optical emission line and the MIR PAH diagnostic classification as compiled in Table 4. With respect to the optical classification, we find that the NIR SED shape confirms for 22 out of 22 (100%) the optical classification as a HII/LINER object and for 6 out of 8 (75%) the classification as Seyfert1/2. With respect to the MIR PAH diagnostics, we find that the NIR SED shape confirms for 12 out of 14 (86%) the PAH classification as SB and for 4 out of 5 (80%) the PAH classification as AGN.

These classifications are further confirmed by comparing them with the hard X-ray classifications (Table 4). Thus, the J-K colours represent an attractive tool to explore the nature of ULIRGs and other dust rich IR galaxies. This is of particular advantage, if the galaxies are too distant or too faint for spectroscopy. Then the intrinsic J-K colours can provide constraints on the nature of cosmologically interesting ULIRGs.

All these methods have some limitations, of course, and in some cases provide indefinite classifications. For example, the J-K colour does not reveal any AGN in Mrk273 which is optically classified as Seyfert2 and has a PAH/continuum flux ratio of 1.9 favouring the starburst dominance (Rigopoulou et al. 1999). Also X-ray observations with Beppo-Sax unveiled an AGN in NGC 6240 (Vignati et al. 1999), which is optically classified as a LINER, SB dominated according to PAH diagnostics, and shows no AGN-typical power-law NIR-MIR flux increase (except for the 1 $\hbox {$^{\prime \prime }$ }$ nuclear region).

5.2.4 The nature of the 10 $\mu $m knee

Pure starburst objects have a flat NIR flux plateau followed by a steep rise longwards of 10 $\mu $m. Part of this pronounced 10 $\mu $m knee could be due to the 9.7 $\mu $m silicate absorption. In contrast, galactic compact HII regions housing O5-O9 stars show SEDs with a steep power-law-like rise. As for AGNs this rise already starts at about 1 $\mu $m (Chini et al. 1987). Also the prototype starburst galaxy M 82 has such a SED shape (e.g. Thuma et al. 2000). Thus, a power-law SED could, in principle, also occur in ULIRGs with starburst classification. But this is not observed for our sample.

HST NICMOS images (Scoville et al. 2000) reveal that the central 1 $\hbox {$^{\prime \prime }$ }$ region is redder than the outer regions. Figure 1 shows that for those 9 sources in common with our sample (05189-2524, 12112+0305, Mrk273, 14348-1447, 15250+3609, Arp220, NGC6240, 17208-0014 and 22491-1808) most of the nuclear SEDs (marked with "$\times $'' symbols) exhibit - even for LINERs and HIIs - a power-law-like flux increase in the NIR. J-K colours of these nuclear regions are similar to the colours in larger apertures found for the AGN-type ULIRGs. For 05189-2524 the NICMOS fluxes (hardly to recognize in Fig. 1) are practically identical with the ground based fluxes showing that this source is completely dominated by the central point-like source. For the others the much higher and bluer ground-based fluxes suggest that the sources are extended on the 5-10 $\hbox {$^{\prime \prime }$ }$ scale, in particular for 17208-0014 and 22491-1808 for which already the NICMOS images indicate that the SEDs are composed of many different components.

With regard to the Balmer decrement LINERs and HII-starburst galaxies have low EB-V (Table 4), as found for Seyfert1s. The highest EB-V are those for Seyfert 2s, indicating that their emission regions are the most deeply embedded ones. A deeply embedded HII region should suffer from extinction in the same way. However, deeply embedded young stars will probably be outshone by those located closer to the outer surface of the galaxy and will not dominate the NIR colours. In conclusion, the relatively flat blue NIR flux plateau comes from the practically unobscured outer regions. The strong flux rise longwards of 10 $\mu $m is caused by more deeply embedded luminous star forming regions. This naturally explains the appearance of the 10 $\mu $m knee.

5.2.5 Does the AGN power the FIR emission of ULIRGs?

It is now widely accepted that a powerful (and not obscured) AGN creates a strong NIR-MIR emission (e.g. Sanders et al. 1988a, 1988b; Pier & Krolik 1992, 1993; Rowan-Robinson 1995; Haas et al. 2000a). However, it is still a matter of debate whether the FIR/submm emission in ULIRGs (as well as quasars) is also mainly powered by the AGN (e.g. Sanders 1999) or by circumnuclear SBs (e.g. Rowan-Robinson 1995; Genzel et al. 1998).

The power-law flux increase of Mrk231 appears similar to that of the quasars PG0050+124 and PG1613+658 (Haas et al. 2000a). For these quasars it is not clear whether and in which wavelength range a starburst contributes to their power-law SEDs: energetically, the AGN radiation would be sufficient to heat the dust, and since pure starbursts have different SED shapes (with knees, as shown above), the superposition of a significant starburst component would probably destroy the smooth power-law shape. For Mrk231, however, the FIR luminosity exceeds the MIR- and the (dereddened) UV-optical-luminosity (Downes & Solomon 1998). Downes & Solomon (1998) found two rotating gas disks with radii of 460 and 1150 pc, respectively, and concluded that the outer one is starburst dominated, since it is too extended and FIR-luminous to intercept sufficient power from the central region. Although the AGN is clearly visible, it dominates neither the FIR luminosity nor the total energy output. However, the PAH/continuum diagnostics indicates a dominating AGN in Mrk231 (Genzel et al. 1998; Lutz et al. 1999). The reason for this might be that the unobscured AGN continuum diminishes the relative PAH strength, mimicking a weak starburst and overemphasising the role of the AGN. Thus, it is necessary to consider not only the ratio PAH/continuum, but also the absolute PAH and continuum levels, both of which may be affected by different amounts of extinction due to different sites of origin in the galaxies. For the diagnostics of moderately luminous Seyfert 1s and Seyfert 2s this has already been stressed by Clavel et al. (2000).

As shown in Fig. 6, the ULIRGs with and without AGN signatures cannot be distinguished via their FIR colours ( $F_{\rm 120~\mu m}$/ $F_{\rm 200~\mu m}$). This suggests that either the FIR/submm emission is largely independent of the presence of an AGN, or that every ULIRG not classified as AGN contains a hidden AGN. The relative high number of ULIRGs with signs of strong starbursts versus ULIRGs with AGN signatures in our sample (22:10 from optical emission line diagnostics, 13:5 from PAH diagnostics) would argue in favour of the first alternative. Since in Sect. 5.1.2 we concluded that the bulk FIR component might be moderately opaque ( $\tau^{\rm bulk}_{\rm 100~\mu m} \approx 1$) even IR diagnostics, in particular the PAH/850 $\mu $m flux ratio, have limited meaningfullness in this context. An exception seems to be Arp220 where the dust geometry leads to a strong deficiency of PAH strength versus 850 $\mu $m flux and dereddening would yield a quasar-like IR continuum (Haas et al. 2001). Again, this is only 1 out of 22 sources of the submm ULIRG subsample showing this behaviour. Therefore, we see some evidence from our sample that the role of the AGN for powering the FIR/submm emission in nearby ULIRGs is negligible.

5.3 Cold dust and search for evolutionary trends

ULIRGs may well have considerable amounts of cold dust at 10-30 K, a temperature typical for dust in less active spiral galaxies. If the cold dust (from the parent galaxies) is continuously heated during the merging process, then we would expect more advanced mergers to show a relatively larger amount of warm dust than mergers in a beginning phase. To check this hypothesis, we use the morphological appearance and the separation of the galaxy nuclei as a measure for the merger state. This is a simplification, since the encounter might not lead to a monotonous approach of the two galaxies. According to simulations by Dubinski et al. (1999) the merging process contains repeated approaches with semi-elastical collisions followed by a drift apart. Nevertheless, we divided the sources into (1) single and (2) double (see Table 4). We use the MIR, FIR and submm luminosities as a measure for the amount of warm, cool and cold dust, respectively. Figure 8 does not show any trend of MIR-, FIR- and submm-luminosity ratios with the simple morphological classifications. Within our sample, which covers only a small range in bolometric luminosity, the relative amounts of warm, cool and cold dust are not correlated with the merger state inferred from the morphology.

\par\includegraphics[width=8.6cm,clip]{ms10599f8a.eps}\\ [2mm]
\end{figure} Figure 8: Distribution of the MIR/(FIR+submm) luminosity ratio versus total IR luminosity (top) and FIR/submm luminosity ratio versus FIR+submm luminosity (bottom). Different symbols for each spectral type as in Fig. 2. The symbol size indicates sources with either a single nucleus (small and filled or thick) or a double nucleus (large and open or thin).

A similar independence was found between the strength of the PAH-feature/MIR-continuum and the (projected) separation of the galaxy nuclei (Lutz et al. 1998; Rigopoulou et al. 1999). Either the (projected) separation of the nuclei is not well suited to measure the progress of the merging process, or the dust heating does not steadily increase during the merging. It seems likely that while the merging nuclei approach each other relatively slowly and/or repeatedly, the dust clouds undergo several phases of compression, turning into cold proto-starburst clouds which subsequently undergo heating by starbursts (and by an AGN, if present). These distinct phases could occur simultaneously in a merger, but at different locations. Photometry of the entire galaxy then shows a mixture of dust complexes at different phases. This is actually revealed by spatially resolved submm observations of the colliding IR luminous (not yet ultra-luminous) galaxy pair NGC4038/39 (Haas et al. 2000b).

\end{figure} Figure 9: Two colour diagram (25 $\mu $m/60 $\mu $m versus 150 $\mu $m/850 $\mu $m) of ULIRGs. Different symbols for each spectral type as in Fig. 2. Note the relation between this diagram and that in Fig. 2 ($\tau $ versus $\beta $): 25 $\mu $m/60 $\mu $m is related to 1/$\tau $ and 150 $\mu $m/850 $\mu $m to $\beta $.

Irrespective of the morphology, the $F_{\rm 25}$/ $F_{\rm 60}$ versus $F_{\rm 150}$/ $F_{\rm 850}$ two-colour diagram (Fig. 9) illustrates the distribution of our ULIRG mm-subsample. Schematically, the cold cirrus-like dust is represented by the 850 $\mu $m emission, the cool SB dust by the 150 and 60 $\mu $m emission, and the warm AGN dust by the 25 $\mu $m emission. In this diagram trends are recognizable: strong AGNs are located in the upper half ( $F_{\rm 25}$/ $F_{\rm 60} \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
...\offinterlineskip\halign{\hfil$\scriptscriptstyle ...), SBs in the lower half. The galaxies with and without cold cirrus-like dust lie towards the left and right side, respectively. This diagram illustrates that despite similar luminosities the nearby ULIRG sample exhibits quite some diversity. Whether or not this is due to evolution is still a puzzle.

Finally, it is interesting to note that the four ULIRGs at medium redshift ( $z \approx 0.3$) have similar dust temperatures as the low redshift ULIRGs, but exhibit the highest FIR luminosities among the sample. While their MIR SEDs and their spectral types can be determined soon, the question of whether they also contain cold dust will have to be left to future observations with the Herschel Space Observatory or the Atacama Large Millimetre Array.

6 Summary

Infrared to millimetre spectral energy distributions (SEDs) have been obtained for 41 bright ultra-luminous infrared galaxies (ULIRGs). These are the most up-to-date and detailed photometric templates of the nearby ULIRGs, which are the fundamental ingredient for cosmological studies. Arp 220, which has been considered so far as a the archetypal ULIRG, is one of the most peculiar objects in our sample.

The SED maxima lie between 60 and 100 $\mu $m, and for those 22 sources with submm detections or upper limits the slopes of the Rayleigh-Jeans tails can be well constrained. The FIR and submm parts of the SEDs between 60 and 1300 $\mu $m can be fitted in two ways:

Firstly, with a single modified blackbody yielding large ranges for both the opacity ( $0.1 <
\tau_{\rm 100~\mu m} < 5$) and the emissivity ( $1.2 < \beta < 2.2$). The resulting temperatures range from 50 to 70 K. However, there are several reasons against the physical relevance of one single dust component:

The optical-NIR morphology is generally disturbed, suggesting the presence of diffuse outer dust layers in addition to a dense central dust body.
It is difficult to understand why $\beta $ is significantly <2 in some ULIRGs, but $\beta \approx 2$ for spiral and Markarian galaxies.
The normalised strength of the 7.7 $\mu $m PAH features shows no correlation with the extinction derived formally from $\tau_{\rm 100~ \mu m}$.
The derived dust masses lead to unusually high gas-to-dust mass ratios from 500 to 2300.
Secondly, the FIR-submm SEDs can be decomposed into two or more $\lambda^{\rm -2}$ modified blackbody spectra representing cool (30-50 K) and cold (10-30 K) dust components. While the cool component should be moderately opaque ( $0.3 \mathrel{\mathchoice {\vcenter{\offinterlineskip\halign{\hfil
...scriptstyle ...) to match size constraints inferred from interferometric CO observations, the cold component should be transparent in order to conform to the observed PAH/100 $\mu $m and PAH/850 $\mu $m ratios. Remarkably, this suggests that the bulk FIR dust component cannot contain the majority of the PAH carriers, rather they are more closely coupled with the submm emitting dust.

This second way to fit the SEDs provides evidence favouring two dust stages:

the cold dust in more quiet regions, which is mainly responsible for the submm emission, and
the cool dust powered by starbursts (SB) which dominates the FIR and may be partly optically thick even in the FIR.
The dust masses derived from such a multiple component model yield gas-to-dust mass ratios in the range between 30 and 300, hence close to the galactic value of 150. ULIRGs classified as either AGNs or SBs appear indistinguishable at FIR and submm wavelengths. We suppose that the cool FIR emitting dust is not related to the AGN, in particular since SB-type ULIRGs seem to be more than twice as frequent in our sample as AGN-type ULIRGs.

In the NIR-MIR the SEDs reveal two basic shapes:

A flat NIR-MIR flux plateau followed by a steep rise at about 10 $\mu $m. All objects optically classified as LINERs and HII/SBs exhibit this SED shape, but so do some of the Seyferts.
A power-law like flux increase starting in the NIR and peaking in the FIR or even MIR. This SED shape is only seen in Seyferts.
The NIR-MIR power-law-like emission is due to warm and hot dust powered by the AGN, and this is a third dust stage restricted to AGN-type ULIRGs. From this finding, a NIR photometry diagnostic tool is established to reveal the presence of an AGN on the basis of the J-K colours (which has red colours, while SBs have blue colours). This tool is successful in the sense that if an AGN is seen by other diagnostics (optical or MIR spectroscopy), in most cases (6 out of 8) it is also seen via the NIR colours.

The detection or non-detection of cold dust could reflect different evolutionary states of the ULIRGs associated with the merger state. However, the current means to determine the merger state, like the projected distance of the nuclei, appears to be insufficient to verify evolutionary trends.


The development and operation of ISOPHOT were supported by MPIA and funds from Deutsches Zentrum für Luft- und Raumfahrt (DLR, formerly DARA). The ISOPHOT Data Centre at MPIA is supported by Deutsches Zentrum für Luft- und Raumfahrt e.V. (DLR) with funds of Bundesministerium für Bildung und Forschung, grant No. 50QI98013. The authors are responsible for the contents of this publication. JCMT is operated by the Joint Astronomy Centre on behalf of the Particle Physics and Astronomy Research Council of the UK, The Netherlands Organisation for Scientific Research, and the National Research Council of Canada.

It is a pleasure for us to thank Dr. Uwe Graser for kindly obtaining the optical spectra during operational tests, and Drs. Robert D. Joseph and José-Míguel Rodríguez-Espinoza for stimulating discussions. We thank the referee, Dr. Suzanne Madden, for a comprehensive set of valuable comments.

For literature search and photometry we used the NASA/IPAC Extragalactic Data Base (NED) and the NASA Astrophysics Data System (ADS).



Online material

Appendix A: Near-Infrared images

\par {\hbox{
... }}
\epsfig{file=ms10599a1g.eps,width=5.65cm,clip=true} }\end{figure} Figure A.1: K-band images showing a field of 60 $\hbox {$^{\prime \prime }$ }$$\times $60 $\hbox {$^{\prime \prime }$ }$ with the sources located in the image centres, greyscales with logarithmic contours.

The J- and K-band images look very similar, therefore only the K-band images are shown (Fig. 10). Most of the emission is concentrated, but also irregular/extended tails are present. The pointlike blue source 10 $\hbox {$^{\prime \prime }$ }$ southwest of 00262+4251, and the two ones located about 6 $\hbox {$^{\prime \prime }$ }$ north and south of 18090+0130 are probably stars (and are excluded from the photometry). The photometry derived from the images is listed in Table 1.

Appendix B: Optical spectra

\par {\hbox{
...esizebox{9cm}{6cm}{\includegraphics[clip=true]{ms10599a2d.eps}} }}\end{figure} Figure B.1: Spectra of 20046-0623 (top) and 21504-0628 (bottom), blue (left) and red (right) channel of the TWIN spectrograph.

Figure 11 shows optical spectra for two sources, obtained with the TWIN spectrograph at the 3.5 mtelescope on Calar Alto, Spain. Both sources have a LINER spectrum according to the line diagnostics by Veilleux & Osterbrock (1987). Now all sources of our mm-subsample are spectroscopically classified (see Table 4).

Copyright ESO 2001