The Herschel M33 extended survey (HerM33es): PACS spectroscopy of the star forming region BCLMP 302
B. Mookerjea, C. Kramer, C. Buchbender, M. Boquien, S. Verley, M. Relano, G. Quintana-Lacaci, S. Aalto, J. Braine, D. Calzetti, F. Combes, S. Garcia-Burillo, P. Gratier, C. Henkel, F. Israel, S. Lord, T. Nikola, M. Roellig, G. Stacey, F. S. Tabatabaei, F. van der Tak, P. van der Werf
aa r X i v : . [ a s t r o - ph . GA ] J un Astronomy&Astrophysicsmanuscript no. m33pacs˙mookerjea c (cid:13)
ESO 2018October 15, 2018
The Herschel M33 extended survey (HerM33es): PACSspectroscopy of the star forming region BCLMP 302 ⋆ B. Mookerjea1, C. Kramer2, C. Buchbender2, M. Boquien3, S. Verley4, M. Rela˜no4, G. Quintana-Lacaci2 , S. Aalto5,J. Braine6, D. Calzetti3, F. Combes7, S. Garcia-Burillo8, P. Gratier6, C. Henkel9, F. Israel10, S. Lord,11, T. Nikola12,M. R ¨ollig13, G. Stacey12, F. S. Tabatabaei9, F. van der Tak15, P. van der Werf10 , (A ffi liations can be found after the references) Received . . . ; accepted . . .
ABSTRACT
Context.
The emission line of [C ii ] at 158 µ m is one of the strongest cooling lines of the interstellar medium (ISM) in galaxies. Aims.
Disentangling the relative contributions of the di ff erent ISM phases to [C ii ] emission, is a major topic of the HerM33es program, a Herschelkey project to study the ISM in the nearby spiral galaxy M33. Methods.
Using PACS, we have mapped the emission of [C ii ] 158 µ m, [O i ] 63 µ m, and other FIR lines in a 2 ′ × ′ region of the northern spiralarm of M33, centered on the H ii region BCLMP 302. At the peak of H α emission, we have observed in addition a velocity resolved [C ii ] spectrumusing HIFI. We use scatterplots to compare these data with PACS 160 µ m continuum maps, and with maps of CO and H i data, at a commonresolution of 12 ′′ or 50 pc. Maps of H α and 24 µ m emission observed with Spitzer are used to estimate the SFR. We have created maps of the [C ii ]and [O i ] 63 µ m emission and detected [N ii ] 122 µ m and [N iii ] 57 µ m at individual positions. Results.
The [C ii ] line observed with HIFI is significantly broader than that of CO, and slightly blue-shifted. In addition, there is little spatialcorrelation between [C ii ] observed with PACS and CO over the mapped region. There is even less spatial correlation between [C ii ] and the atomicgas traced by H i . Detailed comparison of the observed intensities towards the H ii region with models of photo ionization and photon dominatedregions, confirms that a significant fraction, 20–30%, of the observed [C ii ] emission stems from the ionized gas and not from the molecular cloud.The gas heating e ffi ciency, using the ratio between [C ii ] and the TIR as a proxy, varies between 0.07 and 1.5%, with the largest variations foundoutside the H ii region. Key words.
ISM: clouds - ISM: HII regions - ISM: photon-dominated regions (PDR) - Galaxies: individual: M33 - Galaxies: ISM - Galaxies: starformation
1. Introduction
The thermal balance and dynamics of the interstellar medium ingalaxies, is best studied through spectroscopic observations ofits major cooling lines: [C i ], [C ii ], and [O i ] trace the transitionregions between the atomic and molecular gas, while CO tracesthe dense molecular gas that provides the reservoir for stars toform. Most of the important cooling lines lie in the far-infraredand submillimeter regime. Therefore, it is di ffi cult or impossi-ble to study them with ground-based telescopes, while previousspace-based telescopes provided low sensitivity and coarse an-gular resolution. The infrared line emission is mostly opticallythin and can be traced throughout the densest regions in galax-ies, allowing an unhindered view of the ISM. Herschel provides,for the first time, an opportunity to image the major tracers of theISM at a sensitivity, spectral and spatial resolution that allows tostudy the interplay between star formation and the active ISMthroughout our Milky Way and in nearby galaxies.The two fine structure lines of [C ii ] at 158 µ m and [O i ]at 63 µ m, are the strongest cooling lines of the ISM, carryingup to a few percent of the total energy emitted from galaxiesin the far-infrared wavelengths. The [C ii ] line lies 92 K abovethe ground state with a critical density for collisions with H of3 × cm − (Kaufman et al. 1999). While both lines are thought ⋆ Herschel is an ESA space observatory with science instruments pro-vided by European-led Principal Investigator consortia and with impor-tant participation from NASA. to trace photon-dominated regions (PDRs) at the FUV irradiatedsurfaces of molecular clouds, it has been realized early-on that anon-negligible fraction of the [C ii ] emission may stem from theionized medium (Heiles 1994). Owing to its higher upper en-ergy level (228 K) and higher critical density of ≃ × cm − ,the [O i ] 63 µ m line is a more dominant coolant in warmer anddenser neutral regions (e.g. R¨ollig et al. 2006).The [C ii ] line together with the [O i ] line are diag-nostics to infer the physical conditions in the gas, itstemperatures, densities, and radiation fields, by compar-ing the intensities and their ratios with predictions ofPDR models (e.g., Tielens & Hollenbach 1985; Wolfire et al.1990; Kaufman et al. 1999; R¨ollig et al. 2007; Ferland et al.1998). Previous observational studies of [C ii ] emissionfrom external galaxies include the statistical studies byCrawford et al. (1985); Stacey et al. (1991); Malhotra et al.(2001). More recently the [C ii ] emission from a few individ-ual galaxies e.g., LMC (Israel et al. 1996), M51 (Nikola et al.2001; Kramer et al. 2005), NGC 6946 (Madden et al. 1993;Contursi et al. 2002), M83 (Kramer et al. 2005), NGC 1313(Contursi et al. 2002), M 31 (Rodriguez-Fernandez et al. 2006),NGC1097 (Contursi et al. 2002; Beir˜ao et al. 2010) have beenobserved. These papers explore the origin of the [C ii ] emission.Of these studies, the study of LMC by Israel et al. (1996) wasat a resolution of 16 pc and Rodriguez-Fernandez et al. (2006)resolved the spiral arms of M 31 at spatial scales of 300 pc. Fig. 1. µ m Herschel-SPIRE dust continuum image ofthe northern spiral arm of M33 along with the locationsof the prominent H ii regions NGC 604, NGC 595, IC 142,BCLMP 302, and BCLMP 691. The rectangle delineates thearea observed with PACS and HIFI, which is centered onBCLMP 302. The white circles indicate the positions and 70 ′′ beam of ISO / LWS observations done within the BCLMP 302 re-gion. The white ellipse delineates 2 kpc galacto centric distance.The two parallel lines running along the major axis of M33, markthe strip for which we plan to observe the [C ii ] and other FIRlines using PACS and HIFI.In addition to [C ii ] and [O i ] 63 µ m, there are several addi-tional FIR lines, the mid- J CO transitions, and lines of singleand double ionized N and O which provide information aboutthe gas density (multiple transitions of the same ions), hardnessof the stellar radiation field (ratio of intensities of two ionizationstates of the same species), and ionizing flux. In particular the[N ii ] lines at 122 and 205 µ m allow us to estimate the densityof the ionized gas, which is a key parameter to model the [C ii ]emission stemming from the ionized gas.Although much detailed information can be obtained bystudying nearby (Milky Way) sites of star formation, a morecomprehensive view is possible with a nearby moderately in-clined galaxy such as M33. M 33 is a nearby, moderately metalpoor late-type spiral galaxy with no bulge or ring, classified asSA(s)cd. It is the 3rd largest member of the Local Group. Itsmass, size, and average metallicity are similar to those of theLarge Magellanic Cloud (LMC). M33 hosts some of the bright-est H ii complexes in the Local Group. NGC 604 is the secondbrightest H ii region after 30 Doradus in the LMC. Its inclina-tion ( i = ◦ ) (Regan & Vogel 1994) yields a small line-of-sightdepth which allows to study individual cloud complexes not suf-fering from distance ambiguities and confusion like Galacticobservations do. Its close distance of 840 kpc (Freedman et al.1991) provides a spatial resolution of 50 pc at 12 ′′ , allowingus to resolve giant molecular associations in its disk with cur-rent single dish millimeter and far-infrared telescopes like theIRAM 30m telescope and Herschel. Its recent star formationactivity, together with the absence of signs of recent merg-ers, makes M 33 an ideal source to study the interplay of gas,dust, and star formation in its disk. This is the aim of the opentime key program “Herschel M33 extended survey” HerM33es (Kramer et al. 2010). To this end we are surveying the majorcooling lines, notably [C ii ], [O i ], [N ii ], as well as the dust spec-tral energy distribution (SED) using all three instruments on Fig. 2.
Maps of 158 µ m [C ii ] (in color and black contours)and 63 µ m [O i ] (white contours) emission observed with PACStoward BCLMP 302. The [C ii ] intensities shown in the colorwedge are in units of erg s − cm sr − . The [O i ] 63 µ m map hasbeen smoothed to 12 ′′ for easy comparison with [C ii ]. The H α peak observed with HIFI is marked with the asterisk. The whitedots show the footprint of the PACS observations. The contourlevels are at 10 to 100% (in steps of 10%) of peak [C ii ] inten-sity of 1.18 × − erg cm − s − sr − . The contour levels for [O i ]63 µ m emission are 30 to 100% (in steps of 10%) of the peak of3.0 × − erg cm − s − sr − . Both images are at a common reso-lution of 12 ′′ .board of Herschel, HIFI, PACS, and SPIRE. We also use ancil-lary observations of H α , H i , CO and dust continuum at 24 µ m.The HerM33es
PACS and HIFI spectral line observations willfocus on a number of individual regions along the major axis ofM 33 which will initially be presented individually.Here, we present first spectroscopic results obtained forBCLMP302, one of the brightest H ii regions of M33. We discussa 2 ′ × ′ ( ∼ × ii regions,BCLMP 302 (Boulesteix et al. 1974; Israel & van der Kruit1974). Using ISO / SWS, Willner & Nelson-Patel (2002) studiedNeon abundances of H ii regions in M33, including BCLMP 302.Rubin et al. (2008) used Spitzer -IRS to map the emissionlines of [S iv ] 10.51, H(7–6) 12.37, [Ne ii ] 12.8, [Ne iii ] 15.56and [S iii ] 18.71 µ m in 25 H ii regions in M33, includingBCLMP 302. ISO / LWS [C ii ] unresolved spectra at ∼ ′′ res-olution (Gry et al. 2003) are available for this region from thearchive. Here, we present PACS maps of [C ii ] and [O i ] 63 µ mat a resolution of 12 ′′ , together with a HIFI [C ii ] spectrum at2 km s − velocity resolution. We compare these data with (i) theH α emission tracing the ionized gas, (ii) dust continuum imagesat mid- and far-infrared wavelengths observed with Spitzer andHerschel, tracing the dust heated by newly formed stars and thedi ff use interstellar radiation field, and (iii) CO and H i emissiontracing the neutral molecular and atomic gas.The rest of the paper is organized as follows: Sec. 2 presentsdetails of our observations and ancillary data, Sec. 3 states thebasic results of the spectroscopic observations and a qualitativeand a quantitative comparison of the [C ii ] and [O i ] emissionwith all other available tracers and their correlation. Sec. 4 stud-ies the role of [C ii ] as an indicator of the star formation rate Fig. 3.
PACS spectra of [C ii ] (158 µ m), [O i ] (63 µ m), [N ii ](122 µ m) and [N iii ] 57 µ m at the H α peak position ofBCLMP 302. The LSR velocity is given. All lines are unre-solved, i.e. line profiles only reflect the instrumental profiles. Fig. 4.
All four spectra of [C ii ], H i and the (2–1) transition ofCO and CO at the H α peak position of the H ii region BCLMP302. All four spectra are at ∼ ′′ resolution allowing for adetailed comparison. The vertical lines denote the velocities -249.7, -252.8, -254.1 km s − corresponding to H i , CO, and [C ii ]emission respectively.(SFR) and Sec. 5 analyzes the energy balance in the mapped re-gion. Sec. 6 presents a detailed analysis of the emission from theH α peak position in BCLMP 302 in terms of models of PDRs.In Sec. 7 we summarize and discuss the major findings of thepaper.
2. Observations
A region extending over 2 ′ × ′ around the H ii region BCLMP302 in the northern arm of M 33 was observed with the 5 × ff source position outside of the galaxyat RA / Dec (J2000) = ◦ / ◦ . The field of view of theIFU is 47 ′′ × ′′ with 9 . ′′ ii ] 157.7 µ m, [O i ]63.18 µ m, [O i ] 145.52 µ m, [N ii ] 121.9 µ m, [N iii ] 57.3 µ m, and[N ii ] 205.18 µ m lines, with the shortest possible observing time(1 line repetition, 1 cycle), and with a reference position at RA = h m . s
9, Dec = ◦ ′ . ′′ i and 100 µ m ISSA maps.For all the lines a 3 × ′′ gridwith the IFU centered at R.A. = h m . s =+ ◦ ′ . ′′ = ◦ . The resulting footprintis shown in Fig. 2. The FWHM beam size of the PACS spec-trometer is 9.2 ′′ near 63 µ m and 11.2 ′′ near 158 µ m (E.Sturm,priv. comm.). The lines are unresolved, as the spectral resolu-tion of PACS is larger than 90 km s − for all lines. The obser-vations were performed on January 7, 2010 and the total ob-serving time was 1.1 hours for all the 6 lines. The PACS spectrawere reduced using HIPE version 3.0 CIB 1452 (Ott et al. 2010).The WS data reduction pipeline was custom-made by the NASAHerschel Science Center (NHSC) helpdesk. The data were ex-ported to FITS cubes which were later analyzed using internallydeveloped IDL routines to extract the line intensity maps.Using PACS, the [C ii ]158 µ m, [O i ] 63 µ m, [N ii ] 122 µ m,and the [N iii ] 57 µ m lines have been detected . The lines of [O i ]145 µ m and the [N ii ] 205 µ m were not detected. The peak and1 σ noise limits of the intensities of the [C ii ], [O i ], [N ii ] 122 µ mand [N iii ] 57 µ m spectra, at the H α peak position, are presentedin Table 1. Fig. 3 show the observed PACS spectra at the positionof the H α peak position. The peak integrated intensities were de-rived by first fitting and subtracting a polynomial baseline of 2ndorder and then fitting a Gaussian. Due to unequal coverage ofdi ff erent positions, as seen from the grid of observed positionsshown in Fig. 2, the rms achieved is not uniform. It varies byabout a factor of 3 over the entire map. Table 1.
Peak, signal to noise ratio (SNR), and sigma values ofthe integrated line intensities at the H α peak position for the linesdetected with PACS all at original resolution. Line Peak (SNR) 1 σ erg s − cm − sr − erg s − cm − sr − [C ii ] 158 µ m 1.18 × − (67) 1.72 × − [O i ] 63 µ m 7.20 × − (6) 1.20 × − [N ii ] 122 µ m 9.95 × − (7) 1.42 × − [N iii ] 57 µ m 1.30 × − (7) 1.80 × − α peak Using HIFI we have observed a single spectrum at the peak po-sition (R.A. = h m . s
79 Dec = ◦ ′ . ′′ α emission from BCLMP 302 (Fig. 4). The HIFI spectrum was IRAS Sky Survey Atlas 3ookerjea et al.: PACS spectroscopy of BCLMP 302 taken on 01 August 2010 during one hour of observing time us-ing the load chop mode with the same reference position as wasused for the PACS observation. The frequency of the [C ii ] lineis 1900536.9 MHz, known to within an uncertainty of 1.3 MHz(0.2 km s − ) (Cooksy et al. 1986). The blue shifted line requiredto tune the local oscillator to 1899.268 GHz, about the highestfrequency accessible to HIFI. The [C ii ] spectra were recordedusing the wide band acousto optical spectrometer, covering abandwidth of 2.4 GHz for each polarization with a spectral res-olution of 1 MHz. We calculated the noise-weighted averagedspectrum, combining both polarizations. A fringe fitting toolavailable within HIPE was used to subtract standing waves, sub-sequently the data were exported to CLASS for further analysis.Next, a linear baseline was subtracted and the spectrum was re-binned to a velocity resolution of 0.63 km s − (Fig. 4). We scaledthe resulting data to the main beam scale using a beam e ffi ciencyof 69%, using the Ruze formula with the beam e ffi ciency fora perfect primary mirror η mb , = .
76 and a surface accuracyof σ = . µ m (Olberg 2010). The measured peak temperaturefrom a Gaussian fit is 1.14 K and the rms (1 σ limit) is 110 mKat 0.63 kms − resolution, which is consistent within 10% of therms predicted by HSPOT. The half power beam width (HPBW)is 12 . ′′ . For comparison with the [C ii ] spectrum we have observed spec-tra of the (2–1) and (1–0) transitions of CO and CO, at theposition of the H α peak, using the IRAM 30 m telescope on 21August 2010. These observations used the backend VESPA. Thespectra were smoothed to a velocity resolution of 1 km s − forall the CO lines. The forward and beam e ffi ciencies are 95% and80% respectively for the (1–0) transition. The same quantitiesare 90% and 58% for the (2–1) transition. The half-power beamwidths for the (1–0) and (2–1) transitions are 22 ′′ and 12 ′′ , re-spectively. We compare the Herschel data of the BCLMP 302 region, withthe H i VLA and CO(2–1) HERA /
30m map at 12 ′′ resolutionpresented by Gratier et al. (2010). We refer to the latter paperfor a presentation of the noise properties.We also use the Spitzer
MIPS 24 µ m map pre-sented by Tabatabaei et al. (2007), the KPNO H α map(Hoopes & Walterbos 2000), and the HerM33es
PACS andSPIRE maps at 100, 160, 250, 350 and 500 µ m (Kramer et al.2010; Verley et al. 2010; Boquien et al. 2010a). The angularresolution of the 100 and 160 µ m PACS maps are ∼ ′′ and 12 ′′ .The rms noise levels of the PACS maps are 2.6 mJy pix − at100 µ m and 6.9 mJy pix − at 160 µ m where the pixel sizes are3 . ′′ . ′′ ii ] observations at two positions within our mapped re-gion were extracted from the ISO / LWS archival data. The twopositions are at RA = h m s Dec = ◦ ′ ′′ (J2000) andRA = h m s DEC = ◦ ′ ′′ (J2000).
3. Results
Within the 2 ′ × ′ region mapped with PACS (Fig. 2), we have de-tected extended [C ii ] emission from (a) the northern spiral armtraced e.g. by the 100 µ m emission, with the strongest emission arising towards the H ii region BCLMP 302 and (b) from the dif-fuse regions to the south-east and north-west. Comparison of the[C ii ] PACS intensities with the ISO / LWS [C ii ] data at the twopositions shows an agreement of better than 11% at both posi-tions. For the comparison of the intensities at the LWS positions,we first convolved the PACS [C ii ] map to the angular resolutionof the LWS data. The full ISO / LWS [C ii ] data set along the ma-jor axis of M33 will be published in a separate paper by Abreuet al. (in prep.).Some of the PACS [O i ] 63 µ m spectra displayed baselineproblems, attributed to the now decommissioned wavelengthswitching mode. Spectra taken along the [C ii ] emitting part ofthe spiral arm extending north-east to south-west showed prob-lems and have been blanked. Both, the [N ii ] 122 µ m and the[N iii ] 57 µ m lines were detected at only a few positions withinthe mapped region. The maps of [O i ] 63 µ m and [C ii ] 158 µ memission towards the H ii region look very similar and both peakat RA = h m . s
3, Dec = ◦ ′ . ′′
30. In addition, the [O i ]63 µ m map shows a secondary peak towards the south-west,to the south of the [C ii ] ridge, at RA = h m . s
364 Dec = ◦ ′ . ′′
55 (J2000), which is not found in the [C ii ] map. Thesecond [O i ] peak lies between the two ridges detected in CO(2–1) and coincides with an H i peak. This suggests that the [O i ]emission at this peak position arises in very dense atomic gas.Overlays of the [C ii ] map with maps of H α , H i , CO(2–1)and dust continuum in the MIR and FIR (MIPS 24 µ m, PACS100 & 160 µ m) are shown in Fig. 5. The dust continuum mapscorrelate well with the [C ii ] map. They peak towards the H ii region and show the spiral arm extending from the H ii regionin south-western direction. In contrast, CO emission shows aclumpy structure wrapping around the H ii region towards theeast. CO emission shows the spiral arm seen in the continuum,but its peaks are shifted towards the south. The H i emissionshows a completely di ff erent morphology, peaking towards thenorth and south of the H ii region and showing a clumpy fila-ment running towards the west. Further below, we will discussthe correlations in more detail.Fig. 6 shows overlays of the [O i ] 63 µ m map at a resolutionof 12 ′′ with H i and CO(2–1). Towards the south and south-westof the H ii region, the [O i ] 63 µ m emission matches the H i emis-sion well. This is surprising given the high excitation require-ments for the [O i ] line. The secondary [O i ] 63 µ m peak towardsthe south-west, is also traced by H i . It lies in between two ridgesof CO emission, the arm running north-east to south-west, and asecond ridge of emission running in north-south direction. Thenorthern part of this second ridge shows an interesting layeringof emission: both [O i ] and H i are slightly shifted towards theeast relative to this CO ridge. However, there is much less corre-spondence between [O i ] and H i emission towards the east andnorth of the H ii region. α region Fig. 4 shows the velocity resolved 158 µ m [C ii ] spectrum ob-served with HIFI at the H α peak position. In addition, we showthe spectra of H i and the J = CO.HIFI and PACS integrated intensities agree very well. The [C ii ]integrated intensity of 15.6 K km s − , observed with HIFI corre-sponds to an intensity of 1.10 × − erg s − cm − sr − and thismatches extremely well with the [C ii ] intensity observed withPACS, at the nearest PACS position, which is o ff set by only 3 ′′ .All spectra are at a common resolution of ∼ ′′ , allow-ing for a detailed comparison. All spectra show a Gaussian lineshape. However, the line widths are strikingly di ff erent between the atomic material traced by H i , the [C ii ] line, and the molec-ular gas. The H i spectrum shows a FWHM of 16.5 km s − , thatof [C ii ] is 13.3 km s − , while CO 2–1 shows a width of only8 km s − (Table 2). We take this as an indication that the H i diskalong the line of sight is much thicker than the molecular diskwhile the material traced by [C ii ] appears to lie in between. Inaddition, we find that the lines are not centered at the same ve-locity. The HI line is shifted by + . − relative to [C ii ],while the CO lines are shifted by + . − relative to [C ii ],confirming that all three tracers trace di ff erent components of theISM. The shifts are significant, as the error of the Doppler cor-rections are much smaller, for HIFI (D.Teyssier, priv. comm.) asfor the other data.Based on lower spectral resolution (6 km s − ) optical spec-troscopy of H β and [O iii ] emission lines, the velocity of theionized gas is deduced to be shifted by − . − with re-spect to the [C ii ] line, with a velocity dispersion of 11 km s − (Willner & Nelson-Patel 2002; Zaritsky et al. 1989). The opticalspectroscopic data is however at much higher angular resolution(2–4 ′′ ). A map of the emission lines originating solely from theionized medium, which would allow to smooth these data to 12 ′′ resolution for direct comparison with the other data, is not yetavailable. In Section 6.2, we will use PDR models to show thatabout 20–30% of the observed [C ii ] stems from the ionized gasof the BCLMP 302 H ii region.In summary, we find a “layering” of the line-of-sight veloci-ties tracing the di ff erent ISM components. The velocity increasesfrom H α to [C ii ] to CO to H i . Table 2.
Parameters derived from Gaussian fits to spectra ob-served with HIFI, VLA, and IRAM 30m, at the H α peak position(Figs. 2 & 4). θ b indicates the half power beamwidth. Line θ b R T dv v cen ∆ v ′′ K km s − km s − km s − [C ii ] 11 15.6 ± ± ± i
11 1446.7 ± ± ± ± ± ± ± ± ± CO(1–0) 22 0.28 ± ± ± CO(2–1) 11 0.69 ± ± ± ii ] emissionwithothertracers Table 3.
Correlation coe ffi cients for the scatter plots in Fig. 7. Tracers Correlation Coe ffi cientEntire Region Region RegionMap A B C [C ii ]–H α ii ]–CO(2–1) 0.41 0.40 0.47 . . .[C ii ]–H i < . < ii ]–[O i ] . . . 0.77 . . . . . .[C ii ]–F ii ]–F For a more quantitative estimate of the correspondence be-tween the di ff erent tracers in which the spiral arm has been mapped we have created scatterplots of intensities of tracers likeH α , CO, H i , [O i ] 63 µ m, MIPS 24 µ m and PACS 100 µ m asa function of the [C ii ] intensities (Fig. 7). We have used inten-sities from all the maps which are smoothed to a resolution of12 ′′ and gridded on a 12 ′′ grid. In order to identify any apparenttrends in the emission we have defined three sub-regions withinthe mapped region: Region A corresponds to the H ii regionBCLMP 302 itself, Region B corresponds to the south-westernmore quiescent part of the spiral arm, traced by e.g. the 100 µ mcontinuum emission, and centered on a peak of CO emission,and Region C lies outside of the prominent CO arm. These threesub-regions are marked by black rectangles in Figure 5. A sec-ond H ii region along the arm, lies just outside and to the north ofthe box defining the quiescent arm region. For the remainder ofthe paper we always analyze and compared the results for thesethree sub-regions separately.In the log-log scatter plots of Figure 7, we find that the H α ,[O i ], and continuum emission at 24 µ m and 100 µ m show pro-nounced linear correlations with the [C ii ] emission within theH ii region. In the region C , the intensities of the H α emissionand the continuum emission at 24 and 100 µ m remain almostconstant. The CO(2–1) intensity in the H ii region ( A ) is onlypoorly correlated with the [C ii ] emission. In regions B and C the CO(2–1) intensity shows no correlation with the [C ii ] inten-sity and has a large scatter. H i does not show any correlationwith [C ii ].A more quantitative analysis of the correlation between thedi ff erent tracers and [C ii ], is obtained by calculating the correla-tion coe ffi cient ( r ) (Table 3). For the entire region H α and 24 µ mand 100 µ m intensities show around 60% correlation with the[C ii ] intensities. For the region A H α , [O i ], 24 µ m, and 100 µ mare well correlated ( r > .
75) with [C ii ]. The H α emission isstrongly correlated with the [C ii ] intensity also on the south-western arm position. The [O i ] / [C ii ] intensity ratio measuredprimarily in the H ii region varies between 0.1–0.4 and this vari-ation is significantly larger than the estimated uncertainties.
4. [C ii ] as a tracer of star formation In Fig. 8, we plot the star formation rate (SFR), estimatedfrom the 24 µ m MIPS data and the KPNO H α data, as a func-tion of the [C ii ] intensity. Positions within the selected re-gions A , B , and C are marked using di ff erent symbols andcolours similar to Fig. 7. The SFR has been calculated fromSFR = [ L ( H α ) + . L (24)] × . × − in M ⊙ yr − ,where L(H α ) is the H α luminosity in Watt and L(24) is definedas ν L ν at 24 µ m in Watt (Calzetti et al. 2007). We have assumeda Kroupa (2001) initial mass function with a constant SFR over100 Myr. Here, we study the correlation on scales of 12 ′′ corre-sponding to 50 pc. On these small scales, we may start to see abreak-up of any tight correlations between the various tracers ofthe SFR.Viewing at all pixels we find a steepening of the slope inthe log–log plots from regions where both the SFR and [C ii ]are weak, to regions where both are strong. Towards the H ii re-gion ( A ), we find an almost linear relation between log(SFR)and log([C ii ]), with a good correlation, r = .
90 and the fittedslope is 1.48 ± ii ] to the power 1.48.Region B shows a correlation coe ffi cient of 0.64 and region C shows no correlation. Fig. 8.
Correlation of [C ii ] intensities with SFR (derived from24 µ m + H α ). Red triangles correspond to positions in region A region, the green squares represent positions in region B andthe cyan filled circles correspond to positions in region C . Eachmarker corresponds to one position on a 12 ′′ grid. Details aboutthe regions and points are identical to Fig. 7. The black straightline marks the fits to the region A . The errorbar corresponds toan uncertainty of 20%. Fig. 9.
Correlation plot of [C ii ] and Total Infrared (TIR) inten-sities at a 12 ′′ resolution and on a 12 ′′ grid. Markers and thestraight line are the same as in Fig. 7.
5. Energy Balance in the spiral arm
Boquien et al. (2010b) have obtained a linear fit to the totalinfrared (TIR) luminosity as a function of the luminosity inthe PACS 160 µ m band for the entire M 33 galaxy. The TIRis the total infrared intensity, integrated between 1 µ m and1 mm. It is about a factor 2 (Rubin et al. 2009) larger than theFIR continuum which is integrated between 42.7 and 122 µ m.Boquien et al. (2010b) derived the TIR from fits of Draine et al.(2007) models to the MIPS, PACS, and SPIRE FIR data. Andthey find a tight linear relation between the two quantities:log L TIR = a × log L + b with a = ± b = ± . Fig. 10.
Plot of [C ii ] / TIR and ([C ii ] + [O i ]) / TIR (blue asterisks)in the H ii region as a function of the total infrared (TIR) inten-sities at a resolution of 12 ′′ on a 6 ′′ grid. The error bars denote30% errors. Markers same as in Fig. 7.tensity at each position of the mapped region, at a resolution of12 ′′ . To check that this is a valid approach, we independentlyderived the TIR at individual positions by fitting a greybody tothe MIPS, PACS, and SPIRE data, smoothed to a common reso-lution of 40 ′′ . The resulting TIR agrees within 10% with the TIRderived from only the 160 µ m band.Figure 9 shows a scatterplot between TIR and [C ii ].Positions corresponding to the selected sub-regions are indicatedusing di ff erent markers. We find that [C ii ] and TIR are tightlycorrelated with a correlation coe ffi cient of 0.87 only in the H ii region ( A ) with the fitted slope being 0.64 ± ffi cient in region B is 0.52 and and C shows no correlation.Incident FUV photons with energies high enough to ejectelectrons from dust grains (h ν > ffi ciency of 0.1–1%(Hollenbach & Tielens 1997). E ffi ciency is defined as the energyinput to the gas divided by the total energy of the FUV pho-tons absorbed by dust grains. Based on ISO / LWS observationsof a sample of galaxies, Malhotra et al. (2001) found that (a)more than 60% of the galaxies show L [CII] / L FIR > .
2% and (b) L [CII] / L FIR decreases with warmer FIR colors and increasing starformation activity, indicated by higher L FIR / L B ratios, where L B is the luminosity in the B band. We have calculated the ratio ofthe intensities of [C ii ] / TIR, as a proxy for the heating e ffi ciency,and plotted it against with [C ii ] intensities (Fig. 10). The heatinge ffi ciency vary by more than one order of magnitude within the2 ′ × ′ region, between 0.07 and 1.5%. Considering an uncer-tainty of 20% in both the measured 160 µ m intensities and the[C ii ] intensities, for the relation between the 160 µ m luminosityand the TIR luminosity mentioned earlier, we estimate the un-certainty in the [C ii ] / TIR ratio to be ∼ ii ] / TIR intensity ratios is significantly larger thanthe uncertainty we estimate from the [C ii ] and 160 µ m intensi-ties. For the H ii region ( A ), we find heating e ffi ciencies between0.2 and 1.0%. Regions B and C show e ffi ciencies between 0.15–0.4% and 0.4–1.2%. Within the H ii region ( A ), the total heat-ing e ffi ciency, including the [O i ] 63 µ m line, ([C ii ] + [O i ]) / TIR,lies between 0.3–1.2% (Fig. 10). Outside the H ii region, reliable[O i ] data is largely missing. Table 5.
Properties of the H ii region BCLMP 302 Excitation parameter u
180 pc cm − IK74Radius r
39 pc IK74Mass 10 M ⊙ IK74rms electron density 6.2 cm − IK74Electron density n e
100 cm − R08Ionization parameter U − . ff ective temperature T e ff α luminosity 2.2 10 erg s − KPNO map
Notes.
As explained in the text, the value given for the electron density( n e )is assumed to hold for BCLMP 302. References.
IK74 Israel & van der Kruit (1974), R08 Rubin et al.(2008)
6. Modeling the PDR emission at the H α peak ofBCLMP 302 One of the key parameters of PDRs is the FUV radiationfield which heats the PDR. From energy considerations, thetotal infrared cooling emission is a measure of the irradiat-ing FUV photons of the embedded OB stars. We estimatethe FUV flux G (6 eV < h ν < = π I FIR = π . I TIR with G in units of the Habing field1 . − erg s − cm − (Habing 1968) and the intensity of thefar-infrared continuum between 42 . µ m and 122 . µ m, I FIR ,in units of erg s − cm − sr − . Here, we assume that heatingby photons with h ν < ∼ I TIR is a factor of ∼ I FIR (Dale et al.2001).At the H α peak position, the TIR intensity of1.18 × − erg s − cm − sr − translates into a FUV field of G =
46 in Habing units. Outside of the H α peak, the FUVfield, estimated from the TIR, drops by more than one order ofmagnitude (cf. Fig. 9).The FUV radiation leaking out of the clouds is measuredusing the GALEX UV data (Martin et al. 2005; Gil de Paz et al.2007) to be G =
24 for a 12 ′′ aperture.We thus find, that 66% of the FUV photons emitted by theOB stars of the H ii region are absorbed and re-radiated by thedust, and 34% are leaking out of the cloud, at the H α position.This is consistent with the FUV extinction derived from the H α and 24 µ m fluxes (Rela˜no & Kennicutt 2009) and the FUV / H α reddening curve (Calzetti 2001).It is interesting to note that H ii regions observed in theLMC and other spiral galaxies by Oey & Kennicutt (1997) andRela˜no et al. (2002) show that typically around 50% of the ion-izing stellar radiation escape the H ii regions, roughly similar tothe fraction of 36% we find in BCLMP 302 / M33. ii ] emissionfromthe ionizedgas In Table 5 we present properties of the H ii region BCLMP 302compiled from the literature and also calculated thereof. Theseproperties were used to estimate the fraction of [C ii ] emissioncontributed by this H ii region using the model calculations ofAbel et al. (2005) and Abel (2006). These authors have esti-mated the [C ii ] emission for a wide range of physical conditions in H ii regions, varying their electron density n e , ionization pa-rameter U , and e ff ective temperature T e ff .The [C ii ] PACS map is centered on the bright H ii regionBCLMP302 (Boulesteix et al. 1974), which corresponds to theH ii region no. 53 in the M33 catalog of Israel & van der Kruit(1974, IK74). Using the measured radio flux and the formula u (pc cm − ) = . S / f . u . ) / ( D / kpc) / (Israel et al. 1973),IK74 estimate the excitation parameter u =
180 pc cm − for adistance of 720 kpc. IK74 also derive an rms electron densityof 6.2 cm − , and estimate the radius r of the H ii region to be38.5 pc. We did not correct these results for the now better knowndistance, as it does not a ff ect our conclusions. For an electrontemperature T e of 10,000 K, Panagia (1973) expresses the exci-tation parameter as u (pc cm − ) = . × − h N ( L )( β − β ) − i / ,where ( β - β ) is the recombination rate to the excited levels ofhydrogen in units of cm s − . Thus we get for the total flux ofionizing photons N ( L ) = . u s − . Using the above val-ues of the parameters, the dimensionless ionization parameter U was derived via U = N ( L ) / ( c π r n H + ) (Evans & Dopita 1985;Morisset 2004).The rms electron density derived by IK74, provides only alower limit to the true electron density and observations of e.g.the [S ii ] doublet at 6754 Å are missing which would allow adirect estimate. Rubin et al. (2008) observed 25 H ii regions inM33 using Spitzer -IRS, and concluded that their electron den-sities are ∼
100 cm − . Since there is no apparent reason forthe n e in BCLMP 302 to be significantly di ff erent, we assume n e =
100 cm − . Hence, with n H + = n e , we estimate log U = − . iii ] 57 µ m / [N ii ] 122 µ m, at the H α position (Table 1), indi-cates an e ff ective temperature of the ionizing stars of about38,000 K (Rubin et al. 1994).From the model calculations of Abel (2006, Fig. 5), we es-timate the fraction of [C ii ] emission from the BCLMP 302 H ii region. The [C ii ] fraction increases with dropping electron den-sity, but is only weakly dependent on the ionization parameterand the stellar continuum model. Depending on the model, theresulting fraction varies slightly, ∼ − cm − .Observations of the [N ii ] 205 µ m line, in addition to the122 µ m line would allow to estimate more accurately the elec-tron density of the H ii region. The model calculations and ob-servations compiled by Abel (2006) show that the [C ii ] intensi-ties stemming from H ii regions and the intensities of the [N ii ]205 µ m line are tightly correlated. Abel (2006) find Log h I C + H + i = . h I NIIH + i + .
689 (erg cm − s − ). Assuming that 30%of [C ii ] emission stems from the BCLMP 302 H ii region atthe H α peak position, we estimate a [N ii ] 205 µ m intensity of ∼ . × − erg s − cm − sr − (cf. Table 1), corresponding to3 × − W m − . The expected [N ii ] 205 µ m intensity is a fac-tor of 3 below the estimated rms of our PACS observations andhence is consistent with the non-detection. Using the detected[N ii ] 122 µ m line, and the predicted [N ii ] 205 µ m intensity fromabove, the [N ii ] 122 µ m / [N ii ] 205 µ m ratio is 2.8 and this corre-sponds to an n e of 10 cm − (Abel et al. 2005, Fig. 22), which isconsistent with the n e assumed by us for the above calculations. Here we investigate whether the FIR and millimeter lines, to-gether with the TIR, observed towards the H α peak, can be con- Fig. 11.
Comparison of line intensities and intensity ratios atthe position of the H α peak with plane-parallel constant densityPDR models (Kaufman et al. 1999). Grey-scales show the esti-mated reduced χ . The horizontal line shows the FUV estimatedfrom the total FIR intensity. The contours correspond to the in-tensities / ratios of di ff erent spectral lines (as shown in the labels).The [C ii ] intensity corresponds only to the estimated contribu-tion (70% of the total) from the PDRs. Fig. 12.
Star formation rate (SFR) versus the [C ii ] / CO(1–0) in-tensity ratio. Markers are the same as in Figure 7.sistently explained in terms of emission from the PDRs at thesurfaces of the molecular clouds. We use the line and total in-frared continuum intensities at a common resolution of 12 ′′ .At the position of the H α peak we observe the follow-ing intensity ratios on the erg scale (cf. Table 6) [O i ] / [C ii ] = ii ] / CO(1–0) = × , [C ii ] / CO(2–1) = ii ] + [O i ]) / TIR = × − and CO(2–1) / CO(1–0) = =
32 in units of the Habing field and a density of 320 cm − .The fitted FUV field agrees rather well with the FUV field ofG =
46 (in Habing units) estimated from the TIR continuum.The corresponding reduced χ was estimated assuming an error Table 6.
Input to the PDR model for the HIFI position
Tracer Intensity[C ii ] ∗ − [O i ] 3.0 10 − CO (1–0) ∗∗ − CO (2–1) 6.7 10 − TIR 1.18 10 − Notes.
All intensities are in units of erg sec − cm − sr − . ∗ The observed[C ii ] intensities was multiplied by 0.7, assuming that 30% of the emis-sion stems from the H ii region. ∗∗ The CO 1–0 data at 22 ′′ resolutionwere multiplied by the beamfilling factor 1.5, estimated by smoothingthe 2–1 data from 11 ′′ to 22 ′′ resolution. of 30% on the observed intensity ratios. The reduced χ is de-fined as χ = / ( n − n P i = ( I obsi − I modeli ) /σ i ), where ( n −
2) isthe number of degrees of freedom with n being the number ofobserved quantities used for the fitting and I denotes either inte-grated intensities or ratios of integrated intensities. We find theminimum value of χ to be 4.4, indicating that the fit is not sat-isfactory. Indeed, the fitted density seems rather low, given thehigh critical densities of the [C ii ] and in particular also of the[O i ] line. On the other hand, the observed [O i ] / [C ii ] ratio, isconsistent with the low density solution. The ratios with the COlines, deviate from this solution. For instance, the CO 2–1 / cm − . Using line ratios to compare with the PDRmodel, allows to ignore beam filling e ff ects, to first order. Notethat the absolute [C ii ] intensity, reduced by the fraction of 30%stemming from the H ii region, agrees well with the best fittingsolution, indicating a [C ii ] beam filling factor of about 1.The rather poor fit of the intensity ratios towards the H α peakshows the short comings of a plane-parallel single density PDRmodel. First tests using KOSMA- τ PDR models R¨ollig et al.(2006) of spherical clumps with density gradients reconcile bet-ter to the observations. This indicates strong density gradientsalong the line of sight. In a second paper, we shall include newHIFI [C ii ] data along two cuts through the BCLMP 302 regionand explore detailed PDR models, which include the e ff ects ofgeometry and sub-solar metallicity.
7. Discussion
Mapping observations of the northern inner arm of M 33 at anunprecedented spatial resolution of 50 pc have revealed detailsof the distribution of the various components of the interstel-lar medium and their contribution to the [C ii ] emission. [C ii ] at158 µ m is one of the major cooling lines of the interstellar gas.Thus irrespective of whether hydrogen is atomic or molecular,the [C ii ] line emission is expected to be strong wherever thereis warm and photodissociated gas. We have identified emissiontowards the H ii region as well as from the spiral arm seen in thecontinuum, and from a region outside of the well-defined spi-ral arm. [C ii ] emission is strongly correlated with the H α anddust continuum emission, while there is little correlation withCO, and even less with H i . This suggests that the cold neutralmedium (CNM; Wolfire et al. 1995) does not contribute signifi-cantly to the [C ii ] emission in the BCLMP 302 region. Recently,Langer et al. (2010) found a similar poor correlation between[C ii ] and H i emission in a sample of 29 di ff use clouds usingHIFI. The lack of correlation between [C ii ] and CO found in BCLMP 302, may indicate that significant parts of the molecu-lar gas are not traced by CO because it is photo-dissociated in thelow-metallicity environment of M33. This interpretation is con-sistent with both theoretical models developed by Bolatto et al.(1999) and recent observational studies of di ff use clouds in theMilky Way by Langer et al. (2010), and of dwarf galaxies byMadden et al. (2011).Comparison of the first velocity-resolved [C ii ] spectrum ofM33 (at the H α peak of BCLMP 302) with CO line profiles showthat the [C ii ] profile is much broader, by a factor of ∼ .
6, andslightly shifted in velocity, by ∼ . − . Compared to the H i line, at the same angular resolution, the [C ii ] line is less broadby a factor ∼ . ∼ . − . These findingsindicate that the [C ii ] line is not completely mixed with the COemitting gas, but rather traces a di ff erent more turbulent outerlayer of gas with slightly di ff erent systemic velocities, which isassociated with the ionized gas.Interestingly, recent [C ii ] HIFI observations of Galactic starforming regions (Ossenkopf et al. 2010; Joblin et al. 2010) alsoshow broadened and slightly shifted [C ii ] line profiles relative toCO.The two major cooling lines of PDRs are the [O i ] line at63 µ m and the [C ii ] 158 µ m line. The intensity ratios [OI] / [CII]and ([OI] + [CII]) vs. the TIR continuum, have been used exten-sively to estimate the density and FUV field of the emitting re-gions. Using ISO / LWS, Higdon et al. (2003) observed the far-infrared spectra of the nucleus and six giant H ii regions in M 33,not including BCLMP 302, but including NGC 604, IC 142, andNGC 595 shown in Figure 1. The 70 ′′ ISO / LWS beam corre-sponds to 285 pc and therefore samples a mixture of the di ff er-ent ISM phases. They find [O i ] / [C ii ] line ratios in the range0.7 to 1.3, similar to the range of values found in the centerand spiral arm positions of M83 and M51 (Kramer et al. 2005).Towards the H ii region BCLMP 302 in the northern inner armof M 33, we measure much lower [O i ] / [C ii ] ratios between 0.1–0.4. These ratios lie towards the lower end of the values foundby Malhotra et al. (2001) in their ISO / LWS study of the unre-solved emission of 60 galaxies who find values between ∼ . i ] 63 µ m line becomes stronger than the [C ii ] emission inregions of high densities of more than about 10 cm − . At the H α peak position in BCLMP 302, we observed a ratio of 0.4, aftercorrecting [C ii ] for the contribution from the ionized gas. Thisratio indicates lower densities and a FUV field of less than about100 G (cf. Fig. 11). Still lower ratios, indicate lower impingingFUV fields.The ratio of [O i ] + [C ii ] emission over the FIR continuum,is a good measure of the total cooling of the gas relative to thecooling of the dust, reflecting the ratio of FUV energy heatingthe gas to the FUV energy heating the grains, and hence thegrain heating e ffi ciency, i.e. the e ffi ciency of the photo-electric(PE) e ff ect (Rubin et al. 2009). E ffi ciencies of up to about 5%are still consistent with FUV heating, i.e. with emission fromPDRs (Bakes & Tielens 1994; Kaufman et al. 1999). The PEheating e ffi ciency is a function of FUV field, electron density,and temperature. A high FUV field leads to a large fraction ofionized dust particles, lowering the e ffi ciency. On the other hand,low metallicities naturally lead to increased e ffi ciencies as e.g.the PDRs become larger when the dust attenuation is lowered(Rubin et al. 2009). Higdon et al. (2003) compared the [O i ] + [C ii ] emission withthe FIR(LWS) continuum integrated between 43 and 197 µ m,and obtained ratios of 0.2% to 0.7%, corresponding to 0.4 to1.4% for a rough estimate of the FIR / TIR conversion factor of 2(Dale et al. 2001; Rubin et al. 2009). Towards the 2 ′ × ′ regionpresented here, the [C ii ] / TIR ratio varies by more than a factor10 between 0.07 and 1.5%. Extra galactic observations at reso-lutions of 1 kpc or more, find e ffi ciencies of only up to ∼ . / LWS at ∼
300 pc resolution, do also find high e ffi ciencies of up to 2%.Rubin et al. (2009) analyzed BICE [C ii ] maps of the LMC at225 pc resolution, and find e ffi ciencies of upto 4% in this lowmetallicity environment. The mean values found in the LMC arehowever much lower: 0.8% in the di ff use regions, dropping to ∼ .
4% in the SF regions. A similar variation was found in theGalactic plane observations by Nakagawa et al. (1998). Heatinge ffi ciencies observed in the Milky Way span about two ordersof magnitude. Habart et al. (2001) observed e ffi ciencies as highas 3% in the low-UV irradiated Galactic PDR L1721, whereasVastel et al. (2001) also using ISO / LWS fluxes of [C ii ] and [O i ],found a very low heating e ffi ciency of 0.01% in W49N. Theglobal value for the Milky Way from COBE observations is ∼ .
15% (Wright et al. 1991). The galactic averages observedby e.g. Malhotra et al. (2001), are dominated by bright emissionfrom the nuclei where the SFR and the FUV fields are large,hence lowering the photo-electric heating e ffi ciencies.Unresolved observations of external galaxies find a tight cor-relation between CO and [C ii ] emission. This was initially seenby Stacey et al. (1991) with the KAO, and subsequently sup-ported by ISO / LWS observations. However at spatial scales of50 pc we do not detect such a tight correlation in BCLMP 302.This lack of correlation between CO and [C ii ] emission is al-ready seen in the maps of spiral arms in M 31 at ∼
300 pc res-olution (Rodriguez-Fernandez et al. 2006). As summarized bythem, the [C ii ] / CO(1–0) ratio varies from ∼ / M 33, we find that the [C ii ] / CO(1–0) ratiovaries between 200 to 6000, with the H ii region having valuesbetween 800-5000 (Fig. 12). The region C shows large valuesof the [C ii ] / CO(1–0) intensity ratios as CO is hardly detected.Thus, while the [C ii ] / CO(1–0) intensity ratios are higher for re-gions with high SFRs, there is no marked correlation betweenthe two quantities at scales below about 300 pc. The tight cor-relation between [C ii ] and CO emission breaks down at scalesresolving the spiral arms of galaxies. References
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Physikalisches Institut, Universit¨at zu K¨oln, Z¨ulpicherStrasse 77, D-50937 K¨oln, Germany Institute of Astronomy, University of Cambridge, Madingley Road,Cambridge CB3 0HA, England SRON Netherlands Institute for Space Research, Landleven 12,9747 AD Groningen, The Netherlands 19 SUPA, Institute for Astronomy, University of Edinburgh, RoyalObservatory, Blackford Hill, Edinburgh EH9 3HJ, UK 710ookerjea et al.: PACS spectroscopy of BCLMP 302
Fig. 5.
Overlay of the H α , H i and CO(2–1) emission and the dust continuum at 24 (MIPS), 100 & 160 µ m (PACS) emission withcontours of [C ii ] at 158 µ m. The white and black contours are for intensities between 10–20% (in steps of 10%) and 30–90% (insteps of 15%) of the peak [C ii ] intensity of 1.18 × − erg s − cm − sr − . The beam sizes for H α , H i , CO(2–1), 24 µ m, PACS100 µ m, PACS 160 µ m and CO(2–1) observations are 1 ′′ , 12 ′′ , 12 ′′ , 6 ′′ , 6 . ′′ . ′′ i map is integratedbetween -280 km / s to -130 km s − . and the CO(2–1) map is integrated between -270 to -220 km s − . Marked in the 24 µ m image(with black rectangles) are also the regions selected for further analysis. Fig. 6.
Overlay of color plots of [O i ] 63 µ m, H i and CO(2–1) emission with contours of [O i ] at 63 µ m. All plots are at aresolution of 12 ′′ . The white and black contours are for intensities between 30–40% (in steps of 10%) and 50–100% (in steps of10%) of peak [O i ] intensity of 3.0 × − erg s − cm − sr − . Details of the H i and CO(2–1) maps are identical to those in Fig. 5.
Fig. 7.
Correlation of [C ii ] with H α , CO(2–1), H i , [O i ] 63 µ m, MIPS 24 µ m and PACS 100 µ m. The black crosses mark allpositions on a 6 ′′ grid. The red triangles correspond to positions in region A region, the green squares represent positions in region B and the cyan filled circles correspond to positions in region C . All these regions are marked in the 24 µ m map in Fig. 5. Errorbarscorresponding to a 20% error on the plotted quantities along both axes are shown at one representative point in each panel. For[O i ] 63 µ m, only positions close to the H ii region are used. The blue dashed line in the [C ii ]–[O i ] scatter plot corresponds to equalintensities of the two tracers. Table 4.
Average values and their variation of the di ff erent tracers arising from the H ii region, the spiral arm and the region C . Tracer Region A Region B Region C [C ii ] (erg s − cm − sr − ) (4 . ± . × − (1 . ± . × − (1 . ± . × − H α (erg s − ) (7.1 ± × (5.3 ± × (7.95 ± × CO (2–1) (K km s − ) 1.88 ± ± ± i ((K km s − ) 1597 ±
358 1267 ±
325 620 ± i ] (erg s − cm − sr − ) (1.28 ± × − . . . . . .F (Jy beam − ) 101.9 ± ± ± (Jy beam − ) 0.66 ± ± ±±