aa r X i v : . [ a s t r o - ph ] O c t Astronomy&Astrophysicsmanuscript no. 0482 October 28, 2018(DOI: will be inserted by hand later)
The CO line SED and atomic carbon in IRAS F10214+4724 ⋆ Y. Ao , ⋆⋆ , A. Weiß , D. Downes , F. Walter , C. Henkel and K. M. Menten Purple Mountain Observatory, Chinese Academy of Sciences, Nanjing 210008, China MPIfR, Auf dem H¨ugel 69, 53121 Bonn, Germany IRAM, Domaine Universitaire, 38406 St-Martin-d’H`eres, France MPIA, K¨onigstuhl 17, 69117 Heidelberg, Germany
Abstract.
Using the IRAM 30m telescope and the Plateau de Bure interferometer we have detected the C I ( P → P ) andthe CO 3 −
2, 4 −
3, 6 −
5, 7 − + z = . I ( P → P ) line is detected for the first time towards this source and IRASF10214 + I ( P → P ) observation and we detect a velocity gradientalong the east-west direction. The CI line ratio allows us to derive a carbon excitation temperature of 42 + − K. The carbonexcitation in conjunction with the CO ladder and the dust continuum constrain the gas density to n (H ) = . − . cm − and thekinetic temperature to T kin = µ m dust continuum morphology is more compact than the line emitting region, which supports previous findings that thefar infrared luminosity arises from regions closer to the active galactic nucleus at the center of this system. Key words. galaxies: formation – galaxies:high-redshift – galaxies:individual: IRAS F10214 +
1. Introduction
Recent submm- and mm-wavelength dust continuum surveysfrom SCUBA at the JCMT (Ivison et al. 2000; Coppin et al.2006) and MAMBO at the IRAM 30m telescope (Bertoldiet al. 2007) have revealed a population of so-called submmgalaxies (SMGs) at high redshift with far infrared (FIR)luminosities comparable to or higher than those of localUltraluminous Infrared Galaxies (ULIRGs). Spectroscopicfollow-up studies of SMGs in the optical regime suggest thatthe volume density of these sources increases by three ordersof magnitude out to z ∼ I and HCN have ⋆ Based on observations carried out with the IRAM Plateau de BureInterferometer. IRAM is supported by INSU / CNRS (France), MPG(Germany) and IGN (Spain). ⋆⋆ email: [email protected] been presented (Weiß et al. 2005b, 2007; Alloin et al. 2007),which focus on the excitation properties of these massive gasreservoirs at high redshift.IRAS F10214 + −
2) line bymany authors (see Radford et al. 1996 for a summary), and theCO(4 −
3) (Brown & Vanden Bout 1992), CO(6 −
5) (Solomonet al. 1992) as well as the C I ( P → P ) transitions (Weiß et al.2005a). CO(1 −
0) was tentatively detected (Tsuboi & Nakai1992, 1994), but not confirmed by Barvainis (1995). An upperlimit on the C I ( P → P ) line was reported by Papadopoulos(2005). The gas distribution was also investigated in detailby the high angular resolution observations in CO(3 −
2) ofDownes et al. (1995) and Scoville et al. (1995).In this paper we present new measurements of the CO(3 − − −
5) and CO(7 −
6) lines and the first successfuldetection of the higher level fine structure line of atomic car-bon, together with the continuum emission at 3 mm and 1.2 mmtowards F10214.
Y. Ao et al.: The CO line SED and atomic carbon in IRAS F10214 +
2. Observations
We observed the CO J = − I ( P → P ) fine structure line with the PdBI in the compact D config-uration during 8 nights in 2004. The dual frequency setupwas used, and assuming a redshift, z, of 2.2854 the receiverswere tuned to the redshifted CO(3 −
2) frequency of 105.252GHz in the 3 mm band and the frequency of the redshiftedC I ( P → P ) transition at 246.345 GHz in the 1.2 mm band(both upper side band tunings). During this run the weatherwas poor and the performance of the old receivers in the1.2 mm band was unsatisfactory, thus only the 3 mm data wasuseful, yielding an equivalent 6-antenna on-source integra-tion time of 17 hours for CO(3 −
2) with an average systemtemperature of 180 K. Baselines ranged from 15.5 to 112.6meters resulting in a synthesized beam of 6.5 ′′ × ′′ (P.A. ∼ o E of N) for natural weighting and 4.9 ′′ × ′′ (P.A. ∼ o E of N) for uniform weighting. The spectral correla-tor covers ∼ − with a velocity resolution of 2 km s − .The C I ( P → P ) line was reobserved with the new PdBIreceivers on December 25th 2007 under good weather con-ditions in C configuration. Both polarizations were tuned tothe redshifted C I ( P → P ) line, giving a velocity coverage of ∼ − (1GHz band width) and a velocity resolution of3 km s − . Baselines ranged from 16.3 to 175.3 meters yieldinga synthesized beam of 1.12 ′′ × ′′ (P.A. ∼ o E of N) fornatural weighting. The on-source integration time for this runis ∼ +
392 and 0955 + − and 2.7 mJy per beamat 30 km s − resolution for the 3 mm and 1.2 mm line cubes,respectively. The noise levels of the final continuum maps,which were computed by averaging the line free channels, are0.17 mJy per beam (for natural weighting) at 3 mm and 0.7mJy per beam at 1.2 mm, respectively. We estimate the fluxdensity scale to be accurate to about 10% at 3 mm and 20% at1.2 mm. The IRAM 30m observations toward F10214 were carried outduring the period between 2004 July and 2005 February. Weused the A / B and C / D receiver combinations tuned to CO(3 − −
3) (140.330 GHz, 2 mmband), CO(6 −
5) (210.468 GHz, 1.3 mm band) and CO(7 − / antenna gains for increasing frequency are 23.4 ′′ / − , 17.5 ′′ / − , 11.7 ′′ / − and 10 ′′ / − ,respectively. Typical system temperatures were about 120 K, 255 K, 240 K and 330 K for the four CO lines, respectively.The observations were done in wobbler switching mode, witha switching frequency of 0.5 Hz and a wobbler throw of 50 ′′ in azimuth. Data were recorded using the 512 × × − ,respectively. The useful on-source integration times are 3.8,5.9, 3.2 and 9.2 hours for the four CO lines and the resultingrms noise values are 3.5, 4.0, 3.3 and 4.6 mJy, respectively.
3. Results
The C I ( P → P ) line is detected towards F10214 for the firsttime. The integrated line flux is I CI( P → P ) = ± − which is consistent with the upper limit of 7 Jy km s − reported by Papadopoulos (2005). Together with the detectionof the C I ( P → P ) line (Weiß et al. 2005a), F10214 is onlythe third extragalactic source, next to the Cloverleaf (Barvainiset al. 1997; Weiß et al. 2003) and M82 (Stutzki et al. 1997),where both carbon lines have been detected. The resultingcarbon line ratio, L ′ CI( P → P ) / L ′ CI( P → P ) , is 0.84 ± I ( P → P ) line width is 182 ±
30 km s − which agrees wellwith the linewidth derived from the CO(3 −
2) and CO(6 − § I ( P → P ) spectral line data are presented as channelmaps in Fig. 1. The emission peaks move from west to east forincreasing velocities which indicates that the gas rotates alongthe east-west direction. This rotation is also seen in the velocityfield of the C I ( P → P ) line, which is presented as a color mapin Fig. 2 (right). In the same figure we show the contours ofthe integrated line intensity. From the line intensity map it isapparent that the C I ( P → P ) distribution is spatially resolved.The deconvolved full width to half maximum (FWHM) sourcesize is < . ′′ and 1.4 ′′ ± ′′ along the minor and majoraxis respectively, and the position angle is 84 o (E of N). Thismorphology is similar to the arc-like lensed structure visibleat 2.2 µ m (Matthews et al. 1994, Graham & Liu 1995). Thecarbon distribution, however, does not show the additionalcompact 2.2 µ m component about 3 ′′ north of the arc whichargues for di ff erential lensing between the wavelengths of thenear infrared (NIR) emission and those of the molecular gas.The C I ( P → P ) distribution is similar to the morphologyderived from CO (Downes et al. 1995).The 1.2 mm (restframe 360 µ m) continuum map is shown inFig 2 (left). Interestingly, its distribution di ff ers from the mor-phology seen in the carbon line. It it much more compact andthe deconvolved FWHM source size is < ′′ . The continuumpeaks at α (J2000) = h m s .56, δ (J2000) = + o ′ ′′ .8( ± . ′′ ) which agrees well with the peak of the C I ( P → P ). Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + line. The 1.2 mm continuum flux is 9.9 ± ± − , which suggests the source is slightly resolved spa-tially. Our PdBI CO(3 −
2) line observations are represented aschannel maps in Fig. 3 and the moment maps in the rightpanel of Fig. 4. The 3 mm continuum (rest frame 930 µ m) isdetected for the first time and we find an integrated flux densityof 0.58 ± σ level. Toachieve higher sensitivity, natural weighting was used to createthe continuum map (Fig. 4 left). Although the signal is onlydetected at 3 σ significance, the good positional agreementwith the 1.2 mm peak supports the reliability of our detection.The 3mm continuum as well as the integrated line distributionremain unresolved in our 4.9 ′′ × ′′ and 6.5 ′′ × ′′ beams,consistent with the size estimates from the higher resolution1.2 mm data. The emission centroids of the integrated CO(3 − −
2) emission has a velocitygradient along east-west direction, consistent with the rotationseen in C I ( P → P ) . The CO(3 −
2) velocity field is shown inFig. 4 (right), where the ordered velocity gradient is obviousalong the east-west direction. The spatial separation betweenthe centroids of the blue and red emission in the channel mapsis about 2 ′′ , somewhat higher than, but consistent with, theextent of the C I ( P → P ) emission. This further suggests thatthe molecular gas detected in CO is cospatial with the gasdetected in the carbon line.Our 30m CO spectra are presented in Fig. 5. To reduce theuncertainty in the CO line profiles we use all CO spectrato obtain an average line profile (first panel of Fig. 5). AGaussian fit to this spectrum yields a line width of 246 ± − , and redshift of 2.28562 ± −
2) line, which is mainly because the CO(4 − −
6) appear to have broader line profiles. which ismainly because the CO(4-3) and CO(7-6) lines appear to havebroader profiles. More specifically both CO(3-2) spectra aresomewhat asymmetric with the blue line wing being moreprominent than the red wing (see Fig. 5). This asymmetry isalso marginally visible in both CI profiles but not in the higher-J CO transitions. Although the e ff ect is not very pronouncedit could indicate the presence of cold foreground material atlow density with su ffi cient optical depth to a ff ect the low-JCO and CI lines but not the high J CO transitions which havemuch higher critical densities. One could also argue that onlythe low-J CO and CI lines arise cospatially while the high-JCO transitions trace a di ff erent volume. The later explanation, however, ignores that a volume that emits in the high-J COtransitions will always be bright in the low-J lines too.The CO(3 −
2) flux of the spectrum from the 30m telescopeis 3.80 ± − , which is in agreement with previousresults (see Radford et al. 1996 for a summary) and also withour high S / N PdBI observation which gives 3.40 ± − .Previous observations of the CO(4 −
3) and CO(6 −
5) lines havebeen reported by Brown & Vanden Bout (1992) and Solomonet al. (1992). Our CO(4 −
3) integrated intensity is 5.32 ± − , which is only about one third of the value given byBrown & Vanden Bout (1992). Considering the high quality ofour new spectrum and the flux densities observed for the otherCO transitions we consider our new CO(4 −
3) measurementto be more reliable. Our CO(6 −
5) line integrated intensity is7.09 ± − , which is consistent with the previousvalue by Solomon et al. (1992).Our detection of the CO(7 −
6) transition is the first publisheddetection of this transition towards this source. Its line inte-grated intensity is 5.43 ± − . Therefore, our obser-vations constrain the peak of the CO line spectral energy dis-tribution (SED; i.e. the integrated line flux densities versus therotational quantum number) in F10214 to be at the CO(6 −
4. Discussion
To investigate the CO excitation we here use a one-componentlarge velocity gradient (LVG) analysis as described by Weiß etal. (2007) adopting a fixed CO abundance per velocity gradientof [CO] / (dv / dr) = × − pc ( km s − ) − . As discussed before,the line profiles may suggest a more complex gas excitation(for an example of a two component model, see the caseof APM08279 + ff erences in the line profiles are not verypronounced and more likely explained by foreground absorp-tion rather than a 2 component excitation we use here forsimplicity a one-component model which yields the averagegas excitation in F10214. From the LVG model we can limitthe allowed range for the H density to 10 . − . cm − . Thekinetic temperature of the gas, however, is poorly constrained.This is shown in Fig. 7 where we show the χ distribution ofthe T kin –n(H ) parameter space of the CO LVG models.For a gas kinetic temperature of T kin = ) = . − . cm − . For these gas parameters the low-J CO lines are opti-cally thick and thermalized which implies L ′ CO(1 − ≈ L ′ CO(3 − .The LVG predicted integrated intensity of the CO(1 −
0) andCO(2 −
1) line emission are 0.37 ± − and 1.55 ± − , respectively. Within the uncertainties, the CO(1 − Y. Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + flux from the LVG models is in agreement with the upper limitderived by Barvainis (1995). C I Recent studies in the Milky Way and nearby galaxies haveshown that CO and C I have very similar distributions sup-porting the interpretation that their emission arises from thesame gas volume (Fixsen et al. 1999; Ojha et al. 2001;Ikeda et al. 2002). The results from the carbon line andthe CO(3 −
2) line by the PdBI also support that the molec-ular gas detected in CO is cospatial with the gas detectedin the carbon line, as mentioned in §
3. Thus, we can useour observed C I line ratio to obtain an independent estimateof the kinetic temperature via the C I excitation temperatureto solve the ambiguity between the kinetic temperature anddensity arising from the CO LVG models. As discussed inSchneider et al. (2003) the carbon excitation temperature inthe local thermodynamic equilibrium (LTE) can be derivedfrom the C I line ratio via the formula T ex = . ln ( . ) , where R = R T mb (CI( P → P )) dv / R T mb (CI( P → P )) dv , assuming thatboth carbon lines share the same excitation temperature andare optically thin. With our observed line ratio of 0.84 ± T ex = + − K. C I abundance estimates Using our estimate of the CO(1 −
0) line luminosity andthe standard ULIRG factor X CO = ⊙ (K km s − pc ) − (Downes & Solomon 1998), we find a molecular gas mass of(8.9 ± × m − M ⊙ where m is the magnification factor bythe gravitational lens. Using Eq. 2 in Weiß et al. (2005a) andour carbon excitation temperature of 42 K, the neutral carbonmass is M C I = (3.7 ± × m − M ⊙ . This leads to a carbonabundance, [C I ] / [H ] = M [C I ] / M [H ], of (6.9 ± × − , inagreement with our previous results (Weiß et al. 2005a). C I LVG models
To obtain an independent test on the optical depth of the C I lines in F10214 as well as on the LTE assumption used to de-rive the excitation temperature, we have also calculated LVGmodels for C I (Stutzki et al. 1997) using [C I ] / (dv / dr) = . × − pc ( km s − ) − . This abundance per velocity gradient cor-responds to the C I abundance estimated above and a veloc-ity gradient of (dv / dr) = − pc − . For T kin =
60 K and n (H ) = . cm − , which provides a good fit to the observedCO SED, the resulting C I ( P → P ) / C I ( P → P ) line ratio is0.8, which is consistent with our measurement of 0.84 ± I ( P → P ) / CO(3 −
2) line ratioof 0.38, which is somewhat higher than the observed ratio of0.26 ± I ( P → P ) / CO(3 −
2) line ratio in agreement with the observations without signifi-cant changes of the excitation temperature or the optical depthof the lines. This implies that CO and CI LVG models withkinetic temperatures of ∼ −
80 K provide reasonable esti-mates of the physical gas properties. This further suggests thatthe carbon lines in F10214 are indeed not far from LTE. Giventhe underlying assumptions, the single component LVG modelfor CO and CI gives a reasonable prediction for all CO and C I line intensities. This finding supports the view that the C I andCO emission arises from the same volume on galactic scales. C I The measured size of the C I emitting region in conjunctionwith the intrinsic line brightness temperatures for CO and C I allow us to estimate the lensing magnification factor and theintrinsic size of the gas distribution (see Downes et al. 1995).From our LVG models for C I and CO we find an intrinsicbrightness temperature of 10–14 K and 38–47 K for C I ( P → P ) and CO(3 −
2) respectively. Assuming our measured arc-length of 1 . ′′ also holds for the CO lines and adopting themethod in Downes et al. (1995) to derive the lens magnifica-tion factor, this yields magnifications of 11–16 for the carbonline and 9–12 for CO(3 − ff erent magnificationsbetween both lines simply reflect the uncertainties of the un-derlying brightness temperature and do not imply di ff erentialmagnification. For a magnification factor of m =
12, the meanbetween both estimates, we derive an intrinsic source radius of ∼
490 pc.
We have included our new PdBI 3 mm and 1.2 mm measure-ments in the dust continuum analysis using a 2-component dustmodel as described in Weiß et al. (2007). Because the opticallythin approximation for the dust emission does not necessarilyhold at rest wavelength shorter than ∼ µ m for ULIRGs(Downes et al. 1993) we have used the approach described inWeiß et al 2007: the flux density of the dust emission is relatedto the dust temperature, T dust , and the apparent solid angle, Ω app by the relation S ν = [B ν (T dust ) − B ν (T bg )] (1 − e − τν )(1 + z) Ω app ,where S ν is the observed (amplified) flux density and B ν is thePlanck function. The lens magnification factor is hidden inthe apparent solid angle, Ω app . The dust optical depth, τ ν , is aproperty of the source itself and is independent of the e ff ects ofgravitational lensing: τ ν = κ d ( ν r ) M dust , app / D A Ω app where D A is the angular distance of the source, M dust , app is the apparentdust mass and κ d is the dust absorption coe ffi cient. For the dustabsorption coe ffi cient, we adopt κ d ( ν ) = . ν r /
250 GHz) β cm g − , where the dust emissivity index, β = µ m is shown inFig. 8. The contribution of the non-thermal radio emission isnegligible even at 3 mm. For the “cold” component, we find adust temperature, T cold , of 80 ±
10 K and a dust mass, M(colddust), of 1.1 × m − M ⊙ ( ± . Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + poorly characterized as only two data points are available forthe relevant frequency range. Our dust temperature for thecold component is in good agreement with the value givenin Downes et al. (1992), but higher than the value of 55 Kreported by Benford et al. (1999) although they adopteda slightly lower value of the dust emissivity index of 1.5.The physical size of the radius implied by our dust modelis >
350 pc for a magnification factor of 12. Much smallerreal radii are inconsistent with the observed slope of theRayleigh-Jeans tail of the dust SED due to the increasingoptical depth of the dust emission. For the “cold” component,the dust optical depth at 100 µ m of the rest frame is about 1.9.The dust temperature for the “cold” component provides anadditional constraint on the kinetic temperature, which helpsto narrow the ambiguity of the CO LVG models.The small observed source size ( < . ′′ , i.e. < <
240 pc for a magnificationfactor of 12) of the dust continuum emission at 1.2 mm impliesthat the FIR emission in F10214 arises from a di ff erent regionthan the C I and CO lines. This conclusion has already beenreached by Downes et al. (1995) based on surface brightnessarguments. The degree of compactness, however, is surprisingbecause it is di ffi cult to explain this structure in the contextof existing lens models. Eisenhardt et al. (1996) suggested amagnification for the FIR emission region of 30 which trans-lates into an even smaller FIR radius ( <
100 pc). Such smallregions are inconsistent with our dust SED models unless wewould adopt a dust emissivity spectral index β >
Only a limited number of sources have been observed in alarge enough number of CO lines to allow for a determinationof the peak of their CO line SEDs (Fig. 9). From this figureit is apparent that F10214 has CO excitation characteristicssimilar to the centers of the nearby galaxies NGC 253 andM82, and high-z galaxies like J1148 + −
5) or CO(7 − −
4) line (Walteret al. 2002; Weiß et al. 2005c). No such low excitationcomponent is apparent in the CO SEDs of the high-z sourcesstudied so far (Riechers et al. 2006). This implies that all themolecular gas and, in consequence, the starbursts arise fromcompact regions that dominate the (sub)mm emission. This isin agreement with the small equivalent radii determined fromthe dust and CO modelling, and also with the gas morphologyof nearby ULIRGs where the CO emission often arises fromcompact, bright nuclear regions (Downes & Solomon 1998).So far, high-J CO observations in local ULIRGs have onlybeen reported for Mrk231 where the highest observed line isthe CO(6 −
5) transition (Papadopoulos et al. 2007). Also forthis source the CO SED rises up to the CO(6 −
5) transition,which implies similar or even higher gas excitation comparedto high-z FIR luminous galaxies.While the CO excitation appears to be very similar for mostactive galaxies, there are two outstanding examples visible inFig. 9: APM08279 + +
5. Conclusion
Using the IRAM 30m telescope and the PdBI we have detectedthe C I ( P → P ) , CO J = −
2, 4 −
3, 6 − − + I ( P → P ) line is detected for the firsttime towards this source and F10214 is now one out of onlythree extragalactic objects at any redshift where both carbonfine structure lines have been detected. The CI line ratio allowsus to derive a carbon excitation temperature of 42 + − K.We have used the C I excitation temperature to narrow thetemperature-density ambiguity in the LVG modeling forCO. The CO and C I lines together with the dust continuumconstrain the gas density to n (H ) = . − . cm − and thegas kinetic temperature to T kin = I line intensities. Thisfinding supports the view that the C I and CO emission arises Y. Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + from the same volume on galactic scales.The source is well resolved by our new C I ( P → P ) lineobservation along the east-west direction, showing extendedemission and a velocity gradient along the major axis, whichis similar to the structure seen in CO. The major axis diameterin conjunction with the intrinsic brightness temperatures ofCO and C I imply a gravitational magnification of ∼
12, inagreement with previous results.The continuum emission at 1.2 mm is more compact than themorphology of the C I ( P → P ) line which shows that the FIRemitting region is smaller than the molecular gas distribution.We find a temperature of the cold dust of T cold = ±
10 K. Thehigh dust temperature together with the compact morphologysuggests that parts of the FIR emission could be due to heatingof the central AGN.
Acknowledgements.
We thank P. Cox for supporting this projectthrough DDT observing time allocation for the new PdBI receiverswhich greatly improved our C I ( P → P ) line spectrum. Y.A. ac-knowledges the financial support from Chinese Academy of Sciencefor supporting his stay as a visiting scholar at MPIfR, when most ofthis work was done. Y.A. also acknowledges the support from NSFCgrant 10733030. Finally, we appreciate the comments of the anony-mous referee which improved our manuscript. References
Alloin, D., Kneib, J., Guilloteau, S., & Bremer, M., 2007, A&A, 470,53Barvainis, R. 1995, AJ, 110, 1573Barvainis, R., Maloney, P., Antonucci, R., & Alloin, D. 1997, ApJ,484, 695Beelen, A., Cox, P., Benford, D., et al. 2006, ApJ, 642, 694Benford, D. J., Cox, P., Omont, A., Phillips, T. G., & McMahon, R. G.1999, ApJ, 518, L65Bertoldi, F., Cox, P., Neri, R., et al. 2003, A&A, 409, L47Bertoldi, F., Carilli, C., Aravena, M., et al. 2007, ApJS, 172, 132Brown, R. L., Vanden Bout, P. A. 1991, AJ, 102, 1956Brown, R. L., & Vanden Bout, P. A. 1992, ApJ, 397, L19Chapman, S. C., Blain, A. W., Smail, I., & Ivison, R. J. 2005, ApJ,622, 772Coppin, K., Chapin, E. L., Mortier, A. M. J., et al. 2006, MNRAS,372, 1621Downes, D., & Solomon, P. M. 1998, ApJ, 507, 615Downes, D., Radford, J. E., Greve, A., et al. 1992, ApJ, 398, L25Downes, D., Solomon, P. M., & Radford, S. J. E. 1995, ApJ, 453, L65Eisenhardt, P. R., Armus, L., Hogg, D. W., Soifer, B. T., Neugebauer,G. & Werner, M. W., 1996, ApJ, 461, 72Fixsen, D. J., Bennett, C. L., & Mather, J. C. 1999, ApJ, 526, 207Graham, James R., & Liu, Michael C., 1995, ApJ, 449, L29Greve, T. R., Bertoldi, F., Smail, I., et al. 2005, MNRAS, 359, 1165G¨usten, R., Phillip, S. D., Weiß A., & Klein, B. 2006, A&A, 454, 115Ikeda, M., Oka, T., Tatematsu, K., Sekimoto, Y., & Yamamoto, S.2002, ApJS, 139, 467Ivison, R. J., Smail, I., Barger, A. J., et al. 2000, MNRAS, 315, 209Kneib, J.-P., van der Werf, P. P., Kraiberg Knudsen, K., et al. 2004,MNRAS, 349, 1211Kr¨ugel, E., & Siebenmorgen, R. 1994, A&A 288, 929Lawrence, A., Rowan-Robinson, M., Oliver, S., et al. 1993, MNRAS,260, 28 Matthews, K., Soifer, B. T., Nelson, J., et al. 1994, ApJ, 420, L13Ojha, R., Stark, A. A., Hsieh, H. H., et al. 2001, ApJ, 548, 253Papadopoulos, P. P. 2005, ApJ, 623, 763Papadopoulos, P. P., Isaak, K. G., & van der Werf, P. P. 2007, arXiv,0706.0811Radford, S. J. E., Downes, D., Solomon, P. M., & Barrett, J. 1996, AJ,111, 1021Riechers, D. A., Walter, F., Carilli, C. L., et al. 2006, ApJ, 650, 604Rowan-Robinson, M., Broadhurst, T., Oliver, S. J., et al. 1991, Nature,351, 719Rowan-Robinson, M., Efstathiou, A., Lawrence, A., et al. 1993,MNRAS, 261, 513Schneider, N., Simon, R., Kramer, C., et al. 2003, A&A, 406, 915Scoville, N. Z., Yun, M. S., Brown, R. L., & Vanden Bout, P. A., 1995,ApJ, 449, L109Solomon, P. M., Downes, D., & Radford, S. J. E. 1992, ApJ, 398, L29Solomon, P. M., & Vanden Bout, P. 2005, ARA&A, 43, 677Spergel, D. N., Bean, R., Dor´e, O., et al. 2007, ApJS, 170, 377Stutzki, J., Graf, U. U., Haas, S., et al. 1997, ApJ, 477, L33Tsuboi, M., & Nakai, N. 1992, PASJ, 44, L41Tsuboi, M., & Nakai, N. 1994, PASJ, 46, L179Walter, F., Bertoldi, F., Carilli, C. L., et al. 2003, Nature, 424, 406Walter, F., Weiß A., & Scoville, N. 2002, ApJ, 580, L21Weiß A., Henkel, C., Downes, D., & Walter, F. 2003, A&A, 409, L41Weiß A., Downes, D., Henkel, C., & Walter, F. 2005a, A&A, 429, L25Weiß A., Downes, D. Walter, F., & Henkel, C. 2005b, A&A, 440, L45Weiß A., Walter, F. & Scoville, N. Z. 2005c, A&A, 438, 533Weiß A., Downes, D., Neri, R., et al. 2007, A&A, 467, 955. Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + Table 1.
Observed CO and C I line parameters towards F10214 Line Telescope ν obs S ν a ∆ V FWHMb I a V b L ′ /
10 c [GHz] [mJy] [ km s − ] [Jy km s − ] [ km s − ] [K km s − pc ]avg. profile IRAM 30m ... ... 246 ±
10 ... 20 ± −
2) IRAM 30m 105.25230 15.5 ± ±
23 3.80 ± ±
11 10.83 ± −
2) PdBI 105.25160 13.5 ± ±
12 3.40 ± ± ± −
3) IRAM 30m 140.33020 21.7 ± ±
28 5.32 ± ±
11 8.53 ± −
5) IRAM 30m 210.46850 28.8 ± ±
16 7.09 ± ± ± −
6) IRAM 30m 245.52621 22.1 ± ±
30 5.43 ± ±
14 2.84 ± I ( P → P ) IRAM 30m 149.80235 8.3 ± ±
26 2.03 ± ±
14 2.86 ± d C I ( P → P ) PdBI 246.34500 18.7 ± ±
30 4.59 ± ±
10 2.39 ± a The values are derived from Gaussian fits with the fixed line width and center velocity of the averaged CO line profile. b The values are obtained from Gaussian fits to each spectrum. The velocity o ff sets are centered at a redshift of 2.2854 (Downes et al.1995). c We use a Λ cosmology with H =
73 km s − Mpc − , Ω Λ = Ω m = d The data are from Weiß et al. (2005a), but fitted with the fixed line width and central velocity.
Fig. 1.
Channel maps of the C I ( P → P ) line of F10214 after subtracting the continuum. The contours are -3, -2, 2, 3, 4, 5, 6, 7 × − (1 σ ), with a synthesized beam of 1.12 ′′ × ′′ , which is shown in the lower left corner of the first panel. Themaps are centered on the position shown as a cross ( α (J2000) = h m s .56, δ (J2000) = + o ′ ′′ .8), here and in subsequentfigures. Y. Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + Fig. 2.
Left panel: Contour map of the 1.2mm continuum towards F10214, which was created from the line free channels.Contours are -3, 3, 5, 7, 9, 11 × − (1 σ ). Right panel: Contour map of the C I ( P → P ) line integrated intensitytowards F10214 after subtracting the continuum overlaid by the velocity field in color map. Contours are -3, 3, 6, 9, 12, 15, 20 × − beam − (1 σ ). A synthesized beam of 1.12 ′′ × ′′ is shown in the lower left corner of the figures. Fig. 3.
CO(3 −
2) channel maps of F10214 after subtracting the continuum. The contours are -3, -2, 2, 3, 4, 5, 6, 7, 8, 9 × − (1 σ ), with a synthesized beam of 4.9 ′′ × ′′ , which is shown in the lower left corner of the first panel. . Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + Fig. 4.
Left panel: Contour map of the 3mm continuum towards F10214, which was created from the line free channels. Contoursare -2, 2, 3 × − (1 σ ). Right panel: Contour map of the CO(3 −
2) line integrated intensity towards F10214 aftersubtracting the continuum overlaid by the velocity field in color map. Contours are -3, 3, 9, 15, 30, 50 × − beam − (1 σ ). The synthesized beams of 6.5 ′′ × ′′ and 4.9 ′′ × ′′ are shown in the lower left corner of the figures. Fig. 5.
Spectra of the averaged profile and the CO 3 −
2, 4 −
3, 6 −
5, 7 −
6, C I ( P → P ) and C I ( P → P ) lines towards F10214. Theresults from the PdBI are shown in the two panels at the lower right. All lines are fitted and Gaussian fits are shown as continuouslines for which the FWHM line width and the central velocity are fixed to the value determined from the averaged CO profile.The dashed horizontal lines in the PdBI spectra show the dust continuum levels. The zero and Gaussian fitted velocities from theaveraged CO profile are shown as dashed vertical lines. The C I ( P → P ) spectrum is taken from Weiß et al. (2005a). + Fig. 6.
Observed CO fluxes vs. rotational quantum number (CO line SED, filled squares) for F10214 fitted by the a singlecomponent LVG model with n (H ) = . cm − and T kin =
60 K, shown as a thick solid line. Other measurements, taken fromthe literatures, are marked by di ff erent symbols. For better visibility, previous data points for the 3 − ff sets. Fig. 7.
Reduced χ distribution for a single component LVG model fit to the observed line luminosity ratios (grey scale andwhite contours, contours: χ = / (dv / dr) = × − pc ( km s − ) − . . Ao et al.: The CO line SED and atomic carbon in IRAS F10214 + Fig. 8.
Two component dust model for F10214. Displayed flux densities were taken from Lawrence et al. 1993 (20 cm, 6 cm and3.6 cm, open squares), this work (3 mm and 1.2 mm, open circles), Downes et al. 1992 (1200, 100 µ m, filled triangles), Rowan-Robinson et al. 1993 (1100, 800, 450 µ m, filled circles), Benford et al. 1999 (350 µ m, open triangle) and Rowan-Robinson et al.1991 (60 µ m, filled square). The dashed lines show the thermal dust continuum emission for 80 and 160 K dust components. Thesolid line is the total emission from both components. See § Fig. 9.
Comparison of the CO line SEDs of selected local and high-z galaxies. The SEDs are shown for IRAS F10214 + + = = + = + = −−