Draft version April 17, 2015Preprint typeset using LATEX style emulateapj v. 01/23/15

NGC5195 PDR MODELLING OF COLD GAS AND DUST USING HERSCHEL PACS SPECTROSCOPY

Christian Baker and Supervisor: Christine WilsonMcMaster University

Draft version April 17, 2015

ABSTRACT

Data from the Herschel Space Observatory Photodetector Array Camera and Spectrometer (PACS)spectroscopy as part of the Very Nearby Galaxies Survey (VNGS) has revealed ISM gas characteristicsfor dwarf galaxy NGC 5195, the companion to M51, using spectral lines [C II]158 [N II]122, [O I]63, [OI]145, and [O III]88. With these observed flux values we compare to a predicted flux from a photondominated region (PDR) model to determine the characteristics of the cold gas and then compareto other recently studied galaxies. We find an averaged [C II]/FTIR value of 4.4×10−4 which fallsan order of magnitude under comparable galaxies. We find a logG0 of 2.5-3.3 and a log(n/cm−3) of2.0-3.0.

1. INTRODUCTION

We study the S0 galaxy NGC 5195, the companionto M51, for its cooling properties. NGC 5195 is lo-cated at a distance of 7.7±1.0 Mpc away (Tonry et al.2001). By observing the cooling properties we aim to usea PDR model to determine gas characteristics of NGC5195 which will allow us to compare with other galaxies.Of particular interest is M51 as NGC 5195 is a compan-ion galaxy and has interacted with M51. This follows thegeneral goals of the (VGNS) which are to probe the prop-erties of gas and dust in the interstellar medium (ISM)of 13 galaxies. NGC 5195 is not one of these galaxies butdata was taken for it at the same time as M51. Thanksspectroscopy from the PACS instrument offers us finestructure lines of [C II]158 [N II]122, [O I]63, [O I]145, and[O III]88 at resolutions better than 12” (see Table 1).

In order to investigate the cooling properties of the in-terstellar medium (ISM), fine structure lines such as [CII]158 [N II]122, [O I]63, [O I]145, and [O III]88 can be used.These lines show gas cooling by de-excitation via photonemission. In particular the [C II] line traces both neutraland ionized gas, the [O I] lines trace neutral gas, andthe [N II] and [O III] lines trace ionized gas. Due to thevarious emission lines tracing different parts of the gas,we can use the [C II] emission, a tracer of photon domi-nated regions (PDRs), even though it also traces ionizedgas (with an ionization potential of only 11.26 eV). Incomparison the [N II] and [O III] lines have ionizationpotentials of 14.5 eV and 35 eV and thus require a hardradiation field.

We probe these lines using the Herschel PACS(Poglitsch et al. 2010) in order to investigate the ISM inNGC 5195 using a PDR model, as well as comparing itto recent studies on M51,and Centaurus A.(Parkin et al.2013; Parkin et al. 2014). We use the spectra obtainedto investigate the gas component of NGC 5195, and thePDR model of Kaufman et al. (1999, 2006) to estimatethe properties of the ISM.

A PDR model predicts the physical characteristicsthrough the hydrogen nucleus density, n, and thestrength of the far-ultraviolet (FUV) radiation field inunits of the Habing Field, G0 = 1.6×10−3 erg cm−2 s−1

bakerca2@mcmaster.ca

(Habing 1968). Kaufman’s model (Kaufman 1999, 2006)gives a method to use the line ratios and total in-frared luminosity to determine n and G0. Our datauses the readily available analysis tool the PDR Toolbox(PDRT) (Pound & Wolfire 2008); which may be foundat http://dustem.astro.umd.edu/pdrt. PDR models as-sume a plane-parallel, semi-infinite slab geometry andinclude a complex chemical network, thermal balance,and radiative transfer.

Our comparisons to Cen A and M51 will include find-ing an average [C II]/FTIR value to compare to 4× 10−3

of M51 and 8.4 × 10−3 of Cen A (Parkin et al. 2013;Parkin et al. 2014). We also will compare heating effi-ciency which is determined from ([C II]+[O I]63)/FTIR.This same value when used in conjunction with [C II]/[OI]63 will give values to compare to n and G0.

Organization of this paper is as follows: Section 2 de-scribes the data and the processing of said data, Section3 describes the gas characteristics, and Section 4 com-pares these observations to theoretical models as well asother recently studied galaxies.

2. OBSERVATIONS

2.1. PACS spectroscopy

We obtained the PACS spectroscopy of NGC 5195 fromthe Guaranteed Time Project of the VGNS (PI; C. D.Wilson). Maps were obtained for [C II]158 [N II]122, [OI]63, [O I]145, and [O III]88. Properties of the observationscan be seen in Table 1. Images were handed to us alreadyreduced and ready for analysis (Parkin T. J.). The [C II]and [N II] lines were mapped in a 3 × 3 set of overlap-ping footprints that cover a square area of 47” on a side.The [O I]63 line was mapped in a 5 × 5 grid. The [OI]145 and [O III]88 lines were mapped in a 2 × 2 grid. T.J. Parkin (private communication) describes the data re-duction as follows: ”The raw data was processed by T. J.Parkin up to Level 2 using the Herschel Interactive Pro-cessing Environment (HIPE; Ott 2010) version 9.0.2649with calibration files FM,41. The standard pipeline forunchopped gratings scans was used. The Level 2 PACSspectral cubes were then passed to the line fitting andmap making program PACSman 3.52 (Lebouteiller et al.2012), where the spectral lines in each raster were fit witha Gaussian line profile and a second order baseline. The

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[C II]

1E-08 2E-08 3E-08 4E-08 5E-08 6E-08 7E-08 8E-08

[N II]

1E-09 2E-09 3E-09 4E-09 5E-09 6E-09 7E-09 8E-09

[O I]63

1E-08 2E-08 3E-08 4E-08 5E-08 6E-08 7E-08 1E-09 2E-09 3E-09 4E-09 5E-09 6E-09 7E-09

[O III]

[O I]145

1E-09 2E-09 3E-09 4E-09 5E-09 6E-09

Fig. 1.— Maps of NGC 5195 displaying Herschel PACS spectroscopic observations of fine structure lines at native resolution and pixelscale. North is up and east is to the right. Units are W m−2.

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resulting line fits were then integrated to obtain fluxes,then maps were produced by projecting fluxes from eachraster on an oversampled common grid with final pixelscale of 3.133”.” The reduced images can be seen in Fig-ure 1.

2.2. Additional Data

A total infrared flux map was also obtained from TaraParkin which was made using reprocessed photometry(Bendo et al. 2012) at 24 µm from the Multiband Imag-ing Photometer for Spitzer (Rieke et al. 2004) instrumenton the Spitzer Space Telescope (Werner et al. 2004) incombination with published PACS photometry data at70 and 160 µm (Mentuch Cooper et al. 2012). Parkinwas then able to estimate the the total far-infrared inten-sity using the empirical equation (Galametz et al. 2013):

ITIR = (2.133 ± 0.095)v24I24 + (0.681 ± 0.028)v70I70

+(1.125 ± 0.010)v160I160(1)

The map created from this process can be seen in Figure2.

2.3. Data treatment for analysis

Before any data analysis on the maps we convolved themaps to share a full width half maximum (FWHM) ofthe [C II] map ( 11.5”) using a Gaussian function usingthe STARLINK software’s gausmooth function. Onceconvolved the images were also aligned and then testedwith large apertures covering the entire map to see if fluxwas lost in the convolving process. We lose under 1% ofthe total flux and consider flux to be conserved under thetransformations. Calibration uncertainty for PACS datais at ±30% which is primarily small offsets in pointingand drifting of the detector response (PACS OM). Un-certainties listed in tables are from the uncertainty mapsin the raw data however to give a more precise look atthe data as all of the maps are equally affected by thecalibration uncertainty. [O III] and [O I]145 were belowa mean signal-to-noise ratio (SNR) of 9 across the a 30”aperture and have been discounted from most analysis.[O I]145 has a peak SNR of 27 but falls off rapidly as youcan see in Table 2, so we use the peak flux in Section 3.2but not the mean value.

3. PHYSICAL CHARACTERISTICS OF THE GAS

3.1. Line Emission Morphology

As can be seen in Figure 1 the distributions for mostlines is similar, with peaks in the centre. As previouslydiscussed we reject the [O III]. The low signal-to-noiseis evident in the uneven distribution of the map wherewe would expect a peaked distribution like the other de-tected lines. Using starlink’s beamfit command we findit fails to determine a beam width for [O I]145 and [OIII]88. It does provide fits for the full width at half max-imum (FWHM) of 26”, 21” and 15” for [C II], [N II]122and [O I]63 which agrees with what we can see visually.

As seen in Table 3 our peak flux is 2-3 times greaterthan an average over a 30” aperture. Overall morphologylooks elliptical as there is no visible spiral arm structure.

In Table 4 we observe our well detected lines in a ratiowith FTIR. Here we have used the mean and peak values

again. [C II] contributes the most of the lines studied,with [O I]63 coming in close behind. Emission from [C II]and [O I]63 account for up to 0.004% of total emission.This is a factor of 10 lower than what was observed inNGC 5195’s companion, M51 (Parkin et al. 2013).

3.2. Heating Efficiency

We can use ([C II]+[O I]63)/FTIR as a proxy to theheating efficiency (Tielens & Hollenbach 1985). Thisvalue is a measure of the amount of interstellar FUVradiation that gets converted via the photoelectric effectinto gas heating divided by the fraction of its energythatis deposited into the dust grains. We will be using it withthe ratio of [C II]/[O I]63 to use the heating efficiency topredict characteristics of the gas. A map of ([C II]+[OI]63)/FTIR can be seen in Figure 3. With the ratios of([C II]+[O I]63)/FTIR and [C II]/[O I]63 we can startmodelling our PDR’s as seen in Section 4.1. There is adeficit in the center of Figure 3 which suggests reducedheating efficiency in this region. This is consistent withwhat was found in M51 by Parkin et al (2013).

The [O I] lines can give an estimate of the temperatureas displayed in Figure 4 of Liseau et al. (2006). Giventhat our [O I]145 is poorly detected we attempt to get arough estimate of the temperature if the gas is opticallythick or optically thin by using the peak values where thesignal-to-noise is somewhat decent. We find [O I]63/[OI]145 to be 10±1 and as the lines intersect close to therewe can estimate the temperature as being T ≥200 K andn ≥ 103 cm−3.

With the ratios of ([C II]+[O I]63)/FTIR and [C II]/[OI]63 we can start modelling our PDR’s as seen in Section4.1.

3.3. Ionized Gas Fraction

The fraction of emission originating in HII regions is ofparticular interest as opposed to emission from ionizedgas. We estimated the fraction of [C II] emission origi-nates from ionized gas versus neutral gas. The [N II]205line is commonly used to determine the [N II]122/[N II]205which can be used to determine the ionized gas density.We lack the [N II]205 line which we need. Often we es-timate the [N II] ratio when this happens, such as usingthe Galactic value when appropriate as Malhotra et al.(2001) did. Parkin found however that this disagreedwith other methods (Parkin et al. 2013; Kramer et al.2005).

We instead estimate the [N II]122/[N II]205 at variousionized gas densities constrained using the [S III]18.71/[SIII]33.48 which was provided by Tara Parkin (private com-munication) as n ≤ 102 cm−3. Parkin obtained this valueusing Spitzer low resolution IRS spectrum with the linefitting program PAHfit (Smith et al. 2007) to obtain linefluxes for the various lines giving a [S III]18.71/[S III]33.48ratio of 0.35±0.02 which compared to a theoretical curve(Snijders et al. 2007) indicates that n ≤ 102 cm−3. Bytaking 4 different estimates all within the range speci-fied by the silicon ratio we can account for how differentapproaches may lead to different estimates.

3.4. Adjustments and Corrections

Kaufman’s PDR model (Kaufman et al. 1999) requirestwo adjustments to our observed line fluxes. We must re-move the fraction of [C II] flux that comes from ionized

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TABLE 1Properties of our Herschel Observations

Line Wavelength (µm) FWHMa (”) Integration Time (s) OBSID

[O I] 63.184 9.3 5578 1342223763145.525 11 1968 1342223766

[O III] 88.356 9.3 1216 1342223765[N II] 121.898 10 3974 1342223767[C II] 157.741 11.5 2288 1342223764

aValues are from the PACS Observer’s Manual and the SPIRE Observer’s ManualAll data taken on 2011-07-07.

5E-05 0.0001 0.00015 0.0002 0.00025 0.0003

Fig. 2.— The calculated total infrared intensity using Equation 1 at a resolution of 12”. Units are 10−4 W m−2. This map was alignedwith the previous maps but not convolved to the same resolution as 11.5” is close enough to 12” that it will not cause problems for theanalysis.

TABLE 2Peak and mean signal-to-noise ratio of all observed lines

Line Mean Peak

[C II] (158 µm) 89.2 178.5[N II] (122 µm) 11.8 26.6[O I] (63 µm) 25.1 72.6[O I] (145 µm) 8.1 26.7[O III] (188 µm) 5.0 8.7

gas. We also must take into account that Kaufman’smodel is a plane-parallel slab with incident radiationfrom only one side, the side we observe emission from

TABLE 3Observed Flux in a 30” aperture and peak flux

Line Mean (10−9 W m−2) Peak (10−9 W m−2

[C II] (158 µm) 36.6 ±0.2 83.2 ±0.5[N II] (122 µm) 3.1±0.1 7.2 ±0.3[O I] (63 µm) 21.5±0.4 61 ±1

in the far-infrared cooling lines. As we cannot guaranteethe cloud’s orientation faces us we use Kaufman’s advicethat the velocity dispersion for many clouds combinedwith an assumption that the [O I]63 will become opti-cally thick much faster than either [C II] or the total

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0.001 0.002 0.003 0.004 0.005 0.006

Fig. 3.— The heating efficiency map of ([C II]+[O I]63)/FTIR.

TABLE 4Line to total infrared flux ratio in NGC5195

Line Mean 10−4 Line/FTIR Peak 10−4 Line/FTIR

[C II] (158 µm) 4.4±0.2 2.44 ±0.01[N II] (122 µm) 0.37±0.01 0.211±0.009[O I] (63 µm) 2.6±0.5 1.79±0.03

infrared flux means that we only see half the total [OI]63 flux. Thus we double this value for use in our PDRmodel.

Table 5 indicates 4 guessed ratios starting with the the-oretical lower limit for the [N II]122/[N II]205 ratio(Wrightet al. 1991; Bennett et al. 1994) and going to the theoret-ical ratio at n = 100 cm−3. By subtracting the predictedionized portion from the total flux we can divide the fluxby the total to obtain the fraction of neutral gas and thusapply this fraction to our [C II] flux. As a check that ourguessed ratios are reasonably we can quickly compare toobserved ratios for the Milky Way, which fall between 1.0and 1.6 (Wright et al. 1991).

The correction uses predicted [C II]158/[N II]205 val-ues as a function of electron density in her study ofM51 using Solar Gas abundances of C/H=1.4 × 10−4

and N/H=7.9 × 10−5 (Parkin et al. 2013). It variesfrom 3.2 ±0.3 at an assumed [N II]122/[N II]205=0.7 to3.1±0.3N II]122/[N II]205=1 and 2 and finally 3.0±0.3 atN II]122/[N II]205=3.

Kaufman’s paper (Kaufman et al. 1999) recommendsthat the total infrared flux be reduced by a factor of twoto account for the optically thin infrared continuum fluxcoming from both the front and back sides of the gascloud. We have applied this correction to our infraredflux to ensure our PDR model is accurate.

4. RESULTS AND COMPARISONS

4.1. PDR modelling

We use PDRT for our modelling (Pound & Wolfire2008; Kaufman et al. 2006). The model assumesthe PDR is a plane-parallel semi-infinite slab and isparametrized by two free variables, the hydrogen gasdensity, n, and the strength of the impinging FUVradiation field normalized by the Habing field 1.6 ×10−3 erg cm−2 s−1 (Habing 1968). This model includesthermal balance, chemical network, and radiative trans-fer, and produces grids of predicted structure in terms oftwo line ratios, [C II]/[O I]63 and ([C II]+[O I]63)/FTIR

along axes of G0 and n. Figure 4 shows the line ra-tio maps for [C II]/[O I]63 on the left and([C II]+[OI]63)/FTIR on the right. First we look at Figure 5 leftis the uncorrected mean values. We are now using the30% calibration uncertainty as it will be important indetermining the potential values. Figure 5 right showsthe corrected peak values being used. Figure 6 repeatsthis but for mean corrected values.

We can ignore the low–G0 high–n solutions by consid-ering the number of clouds emitting within our beam.If you compare the model predicted [C II] emission forlow–G0 and high–n solutions to our observed [C II] emis-sion we find that we would require multiple PDR regionsof order of magnitude 103, which is a very high numberof clouds along our line of sight (Kramer et al. 2005).Therefore, we ignore solutions in the bottom right of ourplots. We also switch to the mean value for the finalplots as our uncorrected mean at least had a crossoverregion of allowed G0 and n values.

Prior to applying the corrections, our plots were sug-gesting somewhat low–G0 and low–n solutions. This wasfixed once our corrections were made. Surprisingly, thevarying of neutral [C II] compared to ionized [C II] asbased on the estimated ionization density did not make

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TABLE 5Estimated N II ratios predict neutral C II emission fraction

Assumed [N II]122/[N II]205 Predicted [C II]158/[N II]205 Assumed ionized Neutral C IIGas Density (fraction of total)

0.7 8.1 ±0.3 1 0.63 ±0.031 11.6 ±0.5 9.2 0.75 ±0.042 23 ±1 45 0.88 ±0.053 34.8 ±1.5 100 0.92 ±0.05

Fig. 4.— [C II]/[O I]63 is on the left and([C II]+[O I]63)/FTIR is on the right. These are colour maps of constant value for the lineratios. As can be seen in Figures 5 and 6 the lines trace along paths of the same colour.

a large difference after it left the theoretical lower limit.With this and the removal of the lower right solution, wecan determine a logG0 of 2.5-3.3 and a log(n/cm−3) of2.0-3.0.

If we use Figure 1 from Kaufman et al. (1999) we candetermine the surface temperature of the gas is between200K and 300K. Prior to corrections we would have esti-mated the temperature as anywhere from 200K to 1000K.This agrees with our oxygen ratio temperature that wedetermined earlier, despite using a line that was poorlydetected.

4.2. M51 and Cen A

Unlike M51 or Cen A, we did not model arms or diskregions separately as the object is not extended. Table 6shows the logG0 and log(n/cm−3) values for the differentregions of M51, Cen A, and the single region of NGC5195.

NGC 5195 has similar values to previously studiedgalaxies and resembles the average Cen A values as wellas M51’s center values. Both Cen A and NGC 5195 have

TABLE 6Properties of the gas derived from the PDR model

Object log(n/c−3) logG0 T(K)

NGC 5195 2.0-3.0 2.5-3.3 200-300M51 nucleus 3.5-4.25 3.25-4.0 240-475M51 center 2.5-4.0 2.5-3.5 170-680M51 arms 2.0-3.75 1.75-3.0 100-760M51 interarm 2.25-3.75 1.5-3.0 80-550Cen A 2.75-3.75 1.75-2.75 110-260

References: T. J. Parkin (2013; 2014)

smaller ranges on temperature as compared to M51. Per-haps spiral arm structure contributes to the increase insurface temperature by allowing new stars to form moreeasily which in turn would heat the gas.

5. CONCLUSIONS

Using Herschel PACS observations of the importantfine-structure lines [C II]158 [N II]122, [O I]63, [O I]145,and [O III]88, we measure several diagnostic ratios tocompare NGC 5195 with M51 and Cen A. With the the-oretical lower limit for [N II]122/[N II]205 of 0.7 producing

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Fig. 5.— Uncorrected mean on the left, uncorrected peak on the right

a PDR model with no lower limits, we can assume thisis not likely the physical characteristics of the galaxy.Therefore, we know that between 8% and 25% of theobserved [C II] emission originates in ionized gas.

We determined a logG0 of 2.5-3.3 and a log(n/cm−3)of 2.0-3.0. When compared to PDR models this gives atemperature range of between 200K and 300K. This issimilar to Cen A and certain regions of M51, howeverM51 exhibits much higher max temperatures, perhaps

due to its spiral arms. This temperature range and hy-drogen density agree with the potentially unreliable re-sults from using the [O I]63/[O I]145 ratio.

I would like to extend my thanks to my supervisorChristine Wilson for her extensive help on this project, aswell as Maximilien Schirm for his help in understandingPDR models.

APPENDIX

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