On the properties of dust and gas in the environs of V838 Monocerotis
K. M. Exter, N.L.J. Cox, B. M. Swinyard, M. Matsuura, A. Mayer, E. De Beck, L. Decin
AAstronomy & Astrophysics manuscript no. v838mon_v5 c (cid:13)
ESO 2018October 3, 2018
On the properties of dust and gas in the environs of V838Monocerotis (cid:63)
K. M. Exter , , , N.L.J. Cox , , B. M. Swinyard , , M. Matsuura , , A. Mayer , E. De Beck , , and L. Decin , Institute of Astronomy, KU Leuven, Celestijnenlaan 200D, BUS 2401, 3001 Leuven, email [email protected] Herschel Science Centre, European Space Astronomy Centre, ESA, P.O.Box 78, Villanueva de la Cañada, Spain ISDEFE, Beatriz de Bobadilla 3, 28040 Madrid, Spain present address: Université de Toulouse, UPS-OMP, IRAP, 31028, Toulouse, France; CNRS, IRAP, 9 Av. colonel Roche, BP44346, F-31028 Toulouse, France Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK RAL Space, Rutherford Appleton Laboratory, Chilton, Didcot, Oxfordshire, OX11 0QX, UK School of Physics and Astronomy, Cardiff University, The Parade, Cardiff CF24 3AA, UK Department of Astrophysics, University of Vienna, Türkenschanzstrasse 17, 1180 Vienna, Austria Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, SE-439 92 Onsala,Sweden Max-Planck Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Astronomical Institute Anton Pannekoek, Universiteit van Amsterdam, Science Park 904, NL-1090GE, Amsterdam, The Nether-landsReceived; accepted
ABSTRACT
Aims.
We aim to probe the close and distant circumstellar environments of the stellar outburst object V838 Mon.
Methods.
Herschel far-infrared imaging and spectroscopy were taken at several epochs to probe the central point source and the ex-tended environment of V838 Mon. PACS and SPIRE maps were used to obtain photometry of the dust immediately around V838 Mon,and in the surrounding infrared-bright region. These maps were fitted in 1d and 2d to measure the temperature, mass, and β of the twodust sources. PACS and SPIRE spectra were used to detect emission lines from the extended atmosphere of the star, which were thenmodelled to study the physical conditions in the emitting material. HIFI spectra were taken to measure the kinematics of the extendedatmosphere but unfortunately yielded no detections. Results.
Fitting of the far-infrared imaging of V838 Mon reveals 0.5–0.6 M (cid:12) of ≈
19 K dust in the environs ( ≈ . Herschel spectral range, it is only by incorporating data from other observatories and previous epochsthat we can usefully fit the SED; with this we explicitly assume no evolution of the point source between the epochs. We find thatwarm dust with a temperature ∼
300 K distributed over a radius of 150–200 AU.We fit the far-infrared lines of CO arising from the point source, from an extended environment around V838 Mon. Assuming a modelof a spherical shell for this gas, we find that the CO appears to arise from two temperature zones: a cold zone (T kin ≈
18 K) that couldbe associated with the ISM or possibly with a cold layer in the outermost part of the shell, and a warm (T kin ≈
400 K) zone that isassociated with the extended environment of V838 Mon within a region of radius of ≈
210 AU. The SiO lines arise from a warm/hotzone. We did not fit the lines of H O as they are far more dependent on the model assumed.
Key words. infrared: stars – novae, cataclysmic variables – stars: individual (V838 Monocerotis) – ISM: clouds – ISM: dust
1. Introduction
The object V838 Monocerotis (hereafter V838 Mon) is one ofthe most enigmatic observed in stellar astrophysics in recentdecades. It came to attention when it underwent a powerful erup-tive outburst in 2002 (Brown et al. 2002), increasing in luminos-ity by a factor of 100 over a period of three months. Immediatelyfollowing this event a spectacular light echo was formed fromthe outburst light reflecting off the surrounding dust (Bond et al.2003). Post outburst, V838 Mon varied through spectral types F,G, and K, to a very cool M-giant (Evans et al. 2003). Loebman (cid:63)
Herschel is an ESA space observatory with science instruments pro-vided by European-led Principal Investigator consortia and with impor-tant participation from NASA. et al. (2015) found the star to show features of mixed spectraltype, from M3 dwarf to mid-L supergiant: the atmosphere is ob-viously not a conformist.V838 Mon is situated in a molecular cloud and shares itsneighbourhood with a sparse cluster of young B stars (e.g. Af¸sar& Bond 2007). It has a B-type companion, located 28–250 AUfrom V838 Mon (Munari et al., 2007; Tylenda et al., 2009) whichwas engulfed by the material expanding out from V838 Mon af-ter its outburst (Bond, 2006; Munari et al., 2007; Kolka et al.,2009). Photometry from before the outburst will be a mix of thecompanion star and the precursor; and using pre-outburst pho-tometry to fix the status of the precursor requires a reliable mea-sure of the distance. Analysis of high-resolution
HST polarime-try images of the light-echo led Bond et al. (2003) to place a
Article number, page 1 of 24 a r X i v : . [ a s t r o - ph . S R ] A ug &A proofs: manuscript no. v838mon_v5 lower limit of 6 kpc to the distance of V838 Mon. Working onthe same HST material, Tylenda (2004) revised the distance to8 ± . ± . (cid:12) (Tylenda et al. 2005); and thatit was a main-sequence star which suffered a merger with plan-ets or a lower-mass star (e.g. Retter et. al 2007, Tylenda & Soker2003). Tylenda & Soker (2006) discuss the theories to date toexplain V838 Mon, and similar outbursts in two other stars, M31RV (Rich et al. 1989) and V4332 Sgr (Martini et al. 1999), andthey favour the stellar merger mode; in this paper we assume thismodel is correct, although it makes little difference to many ofour calcuations.There have been numerous studies of V838 Mon (thepoint source and the circumstellar region) in atomic lines andmolecules from the optical to the mid-IR and a few in the far-IR. The spectrum has been evolving dynamically and in terms ofcomposition from the first to the most recent observations. Theresponse of the star to the outburst seems to have been to developan expanding photosphere and a multi-component circumstellarshell (Lynch et al. 2004, Geballe et al. 2007, Tylenda et al. 2009),this material flowing out and some falling back. The velocitiesare generally from some tens to 200 km s − , and temperaturesrange from about 2000 K in the photosphere, to a few 100 K inthe extended circumstellar region some 200–300 AU out.The dust in the molecular cloud surrounding V838 Mon canbe seen as the light echo sweeps through, and HST images showa very chaotic region of loops and whirls and cavities (Bond et al.2003, Crause et al. 2005). The appearance of the surroundingdust has led to the speculation that the outburst was not an iso-lated event. van Loon et al. (2004) found that IRAS and MSXobservations show evidence for multiple mass-loss events priorto the 2002 outburst: a 7 (cid:48) × (cid:48) dust shell is visible in the IRASdata, and the MSX data also show an extended object (1.5 (cid:48) di-ameter) at the position of V838 Mon. However, estimates of thecombined mass of this dust and gas are in the range 10s–100solar masses (e.g. Banerjee et al. 2006; Kami´nski et al. 2011;Tylenda & Kami´nski 2012). This would rule out all of this mate-rial having an origin in V838 Mon itself: instead, we are lookingat the interstellar cloud from which V838 Mon, and other mem-bers of its cluster, were formed. CO studies (Kami´nski et al.2011, Kaminski et al. 2007) also show this extended material,confirming it is massive and moreover cool ( ∼
20 K).
Spitzer observations from 2004/5 at 24 µ m, 70 µ m and160 µ m show extended emission distributed over an arcminutefrom the star, and moreover that this is coincident with the HST light echo (Banerjee et al. 2006). This infra-red “light echo” isbelieved to arise from reprocessed thermal emission from dustgrains ( ∼ (cid:12) ) heated by outward propagating UV–visible pho-tons. Unresolved emission at 24 µ m and 70 µ m (PSFs of 6 (cid:48)(cid:48) and24 (cid:48)(cid:48) ) is also seen in the Spitzer data at the position of the central star; moreover the flux has increased between 2004/5 and 2007(Wisniewski et al. 2008), which is suggested to indicate the pres-ence of hot dust freshly condensed in the outburst ejecta. Geminiimages taken in 2007 in the mid-IR, at 11.2 µ m and 18.1 µ m, donot show evidence of extended emission over radial distances ofup to 15 (cid:48)(cid:48) from the central source (Wisniewski et al. 2008).It is clear that V838 Mon is an intriguing star, which is vary-ing with time and which is interesting to study across most ofthe spectrum. We decided to take advantage of the opportunityto observe V838 Mon in the further reaches of the IR with thegood spectral and spatial resolution that the instruments onboard Herschel offered. We report on our imaging and spectral studyof V838 Mon and its near-by environment.We note that in this paper we will use the word "extended" torefer to the emission we detect arising from the molecular cloudin which V838 Mon is found (i.e. the region which is coinci-dent with the HST light echo, with a distribution of an arcminuteor so from the star). "Point source" refers to anything detectedby
Herschel as a point source: the FWHM of the PACS 70 µ mbeam is ∼ (cid:48)(cid:48) (55000 AU at a 6.2 kpc distance), and that of theSPIRE 500 µ m beam is ∼ (cid:48)(cid:48) . These large sizes mean that thepoint source includes the star, its circumstellar environment, andprobably also ISM. Table 1.
The
Herschel observations.
PACSDate, epoch a Obs. id. b λ reference ( µ m) c m.d.20xx05.07.11, 1 1342220129,30 phot: 70, 1601342220131,2 phot: 100, 1601342220136 spec: 79.36, 162.81165.95, 177.771342220137 spec: 179.53, 180.78169.68, 71.071342220138 spec: 118.58,108.0705.01.12, 2 1342245206,7 phot: 70,1601342245208,9 phot: 100,16010.16.12, 3 1342253505,6 phot: 70,1601342253507,8 phot: 100,160SPIRE09.18.10, 1 1342204851 phot: 250, 350, 50004.24.11, 1 1342219551 spec: 194–67104.30.12, 2 1342245152 phot: 250, 350, 50010.14.12, 3 1342253396 phot: 250, 350, 500HIFI Setting; band ν (GHz) d Notes. ( a ) Epoch: a counter based on the PACS observation dates; forSPIRE and HIFI the closest epoch to these is indicated ( b ) Obs. Idrefers to the unique reference number given to all
Herschel observa-tions ( c ) phot=photometer, spec=spectrometer ( d ) WBS ranges in LSBand USBArticle number, page 2 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
Fig. 1. Top to bottom
PACS (top three) and SPIRE (bottom three) maps at the indicated wavelengths and epochs. The flux scaling, in increasingwavelength order, is: PACS -0.05 to 0.08/0.2/0.4 Jy beam − ; SPIRE -0.03 to 0.8/0.6/0.4 Jy beam − . The brightest point on each map is V838 Mon;the tri-lobal shape of the PACS beam can be seen on the 100 µ m maps. N–E is indicated in the central maps. Herschel observations and data processing
Photometric imaging with the
Herschel (Pilbratt et al. 2010) in-struments PACS (Poglitsch et al. 2010) and SPIRE (Griffin et al.
Article number, page 3 of 24 &A proofs: manuscript no. v838mon_v5
For PACS we use the scan map AOT (astronomer observationtemplate) with two orthogonal (45 ◦ and 135 ◦ ) scanning angles,four scan legs with a length of 12 (cid:48) .4 each and a cross-scan stepof 155 (cid:48)(cid:48) , and medium scan speed. We obtained four maps at eachscan angle: 70 and 160 µ m simultaneously, and 100 and 160 µ msimultaneously. The final maps are then produced by concate-nating the frames from both the scan and cross-scan observa-tions, resulting in maps with a uniform coverage of a square areaof just over 6 (cid:48) on a side. The separate PACS observations weretaken back-to-back and within a few months of the SPIRE ob-servations.The data were processed up to Level 1 with the HIPE (Ott2010) track 12 pipeline for scan maps and with calibrationtree 56. The data were then transferred to be processed withScanamorphos (Roussel 2012, version 23). The blue and greenmaps were processed with the default options in the “minimap”mode; for the red frames we also used the option “flat” inScanamorphos when combing the data from all four datasetsto produce a clean background. Glitch detection was done inScanamorphos. The pixel sizes were chosen for the best resultsin the subsequent data analysis: 1.0 (cid:48)(cid:48) at 70 µ m, 1.5 (cid:48)(cid:48) at 100 µ mand 3.0 (cid:48)(cid:48) at 160 µ m. Finally, we subtracted the mean of the back-ground flux to set all the maps to the same zero level.There were slight shifts of the world coordinate system(WCS) of the maps between the epochs – the position ofV838 Mon (and a faint background point source) differed bya few pixels. This could be a consequence of slight telescopeastrometry errors. Since we later make a direct comparison ofthe morphology of the extended emission between epochs, wemoved the WCS of epochs 2 and 3 to that of epoch 1, indepen-dently for each band. This was achieved by fitting the position ofV838 Mon (using the HIPE task sourceFitting ) and resetting theWCS of the maps to the appropriate sky and pixel coordinates.The calibration uncertainties for PACS photometry are 5%from the models and 2% reproducibility (information providedon the PACS documentation page on the Herschel portal), anda recent report of a comparison of different map-makers showsthat those commonly used on PACS data perform to within 10%of each other for extended emission (Paladini et al. 2013). Theerrors in the measured photometry for the point and extendedsources are in fact dominated by the sky noise and the difficultiesin disentangling the two sources.The PACS maps are shown in Fig. 1. For ease of com-parison between epochs and wavelengths we converted theunits to Jy/beam. (For the conversion we used the beam areastaken from a PACS Photometer PSF report of 4 April 2012(“bolopsf_20.pdf”) from the PACS documentation page on theHSC portal. We used beam area values of 40, 54, and 145sq. arcsec for the blue, green, and red maps.)
For SPIRE we used the standard small map AOT with a repe-tition factor three. The observations were done in three filters (PSW=250 µ m, PMW=350 µ m, PLW=500 µ m) simultaneously,and produced maps with a uniform coverage over a 5 (cid:48) diametercircle.The SPIRE data were reduced in HIPE (track 12, calibrationtree spire_cal_12_2) using the pipeline recipe for small maps(with the pipeline script parameter applyExtendedEmissionGain set to True). The Planck offsets were applied via the task zero-PointCorrection . We opted for map pixel sizes slightly smallerthan the default values as this worked better for our subsequentanalysis: 4.5 (cid:48)(cid:48) , 6.25 (cid:48)(cid:48) and 9.0 (cid:48)(cid:48) for PSW, PMW, and PLW respec-tively. We produced the maps calibrated for extended emission(units of MJy sr − , and used in the analysis of the extended emis-sion) and those calibrated for point sources (units of Jy beam − ,and used to measure the photometry of the point source). Wealso created maps with Scanamorphos, using the standard re-duction path and created maps with pixel sizes of 3 (cid:48)(cid:48) , 5 (cid:48)(cid:48) and 7 (cid:48)(cid:48) for PSW, PMW, and PLW respectively and calibrated for pointsources. We subtracted the mean of the background flux to set allthe maps to the same zero level (subsequent comparison of themaps was then easier). Background source subtraction was notnecessary. The calibration uncertainties for SPIRE photometryare 4% systematic and 1.5% random for a point source, and anadditional 4% for the extended emission (Bendo et al. 2013): weadopted a value of 6% overall. As with the PACS photometry,the errors are in fact dominated by the sky noise and the diffi-culties in disentangling the two sources. The SPIRE maps (thoseproduced in HIPE and calibrated for point sources) are shown inFig. 1. We obtained ten PACS spectral line scans, split between threeconcatenated AOTs. We used the chop-nod (medium throw)Pointed mode, performing line-scan spectroscopy with line rep-etition factors between one and three and a chop-nod repetitionfactor one at the (default ) high-density spectral sampling set-ting. The reference wavelengths of the scans are given in Table 1.The wavelength range of the scans is about 2 µ m in the red and1 µ m in the blue.Our PACS observations consist of 20 discrete segments: theten requested line scans and ten accompanying parallel scans.We used the background normalisation pipeline script for linescans to reduce the data, in HIPE track 12 with calibration treeversion 60. This pipeline uses the telescope background to cali-brate the data, and is recommended for observations of long du-ration and of low continuum flux levels, this latter being the casefor V838 Mon.Diffuse emission from the extended source aroundV838 Mon pervades the entire spectroscopy field-of-view.The PACS integral field unit (IFU) extends by ∼ (cid:48)(cid:48) × (cid:48)(cid:48) (covered by a 5 × (cid:48) –2 (cid:48) . This means that flux from the extended emission maybe present in the spaxels containing the point source: seeApp. A. To measure the level of contamination of this diffuseemission on the stellar spectrum of V838 Mon, we comparedthe flux levels of the spectra of the outermost spaxels from ouron-source pointings (i.e. the spaxels sitting on the extendedemission only, with no point-source contribution) to the spectrafrom the off-source/blank sky pointings (which contain only The instrumental resolution is about: 0.02 µ m, 100 km s − (upto 60 µ m); 0.04 µ m, 120 km s − (up to 100 µ m); and 0.11–0.12 µ m,190 km s − (up to 190 µ m)Article number, page 4 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis Fig. 2.
Examples of PACS spectra of V838 Mon, with identifications indicated. The feature at 162.36 µ m (1846.4675 GHz) has not been unam-biguously identified. telescope background emission) to see how much flux from theextended emission could have been picked up in the on-sourceobservations: there is an excess of 0.1–0.4 Jy in the outer spaxelsof the on-source pointings, which we believe comes from theextended source. Under the assumption that this flux is the sameover the whole field-of-view, we removed this contribution fromthe spectra in the final 20 cubes produced by the pipeline. Foreach cube, we averaged the spectra from the faintest ten outerspaxels; fit a low-order polynomial to the average spectrum;and subtracted the fit from the entire cube. The subsequently-extracted point source spectrum of V838 Mon should thenbe clear of contamination from the extended emission. It isimportant to do this before applying the next step – the pointsource calibration – because that step requires the point sourceto contain only emission from the itself.We discovered that V838 Mon was off-centred during ourobservations: it is located almost exactly between the centralspaxel and one of its neighbours. This has consequences for theextraction of the spectrum of the star. The standard task providedin the PACS pipeline scripts to extract a correctly-calibratedspectrum of a point source ( extractCentralSpectrum ) requires itto be located close to the centre of the central spaxel. However,we were able to modify the task to work on V838 Mon, andwhile the resulting spectrum will not be as well corrected as aproperly-centred source, it is superior to applying no correction.For details see App. B.Examples of the PACS spectra of V838 Mon are shown inFig. 2 with some spectral line identifications. The absolute fluxcalibration uncertainty (RMS, combined absolute, reproduce-ability, and pointing scatter) is about 15% peak-to-peak . Fora faint source such as V838 Mon, there is an additional uncer-tainty of the order ± ; we note that this is not also an See the PACS Spectrometer Performance Document on the PACSdocumentation pages on the HSC portal. This information will be provided at a later date also on the PACSpublic web-pages and on the HSC Legacy Library portal. additional uncertainty on the integrated line fluxes. The higheruncertainty in this continuum level due to our modification ofthe point source correction task and the background subtractionwill take this to about ± We obtained a single SPIRE-FTS pointing with sparse imagesampling in the high spectral resolution mode and with 23 scanrepetitions of the FTS mechanism to build up the signal-to-noisein the interferogram. The SPIRE FTS simultaneously covers the(SSW) short wavelength band (190–313 µ m; 31–52 cm − ; 957–1577 GHz) and (SLW) long wavelength band (303–650 µ m; 15–33 cm − ; 461–989 GHz), which both have an unapodised spec-tral resolution of 0.048 cm − . The resulting data set consists of aspectrum from the central bolometer (i.e. V838 Mon) and spec-tra from the surrounding concentric circles of bolometers. Thecentre-to-centre distance of the bolometers is 33 (cid:48)(cid:48) in the SSWband and 51 (cid:48)(cid:48) in the SLW band.We reduced the SPIRE spectra using the standard point-source pipeline described by Fulton et al. (2010) using the HIPEcalibration scheme 12_1 described by Swinyard et al. (2014).The size of the SPIRE beam (17 (cid:48)(cid:48) –40 (cid:48)(cid:48) depending on frequency–see Makiwa et al. 2013) is such that some background emis-sion will be captured as well as the source itself (see App. A).A careful inspection of the spectra in the surrounding bolome-ters showed no evidence of strong emission lines and we onlyremoved a polynomial fit to the continuum (to subtract all con-tinuum sources present), resulting in the rectified spectrum illus-trated in Fig. 3.This continuum subtraction adds an extra uncertainty to thesubsequent spectral fitting, since our source is faint and thecontinuum noisy. We estimate an uncertainty in the continuumplacement of ± . Article number, page 5 of 24 &A proofs: manuscript no. v838mon_v5 flux (1 . × − W m − ), to the errors taken into account in ourmodelling of the integrated line fluxes (Sec. 6.4).As with PACS, V838 Mon is offset from the centre of theFoV. Comparing the location of the star on the PACS and SPIREimages and the location of the central SPIRE bolometers, wedetermined an offset of close to 6 (cid:48)(cid:48) . This has consequence for thepoint-source calibration, which does not account for the flux lossdue to off-centring. We used a SPIRE script provided in HIPE(Track 14) to calculate the underestimate of the fluxes due to thisoff-centring: the loss factor is about 1.04 below 1000 GHz andfrom 1.25 to 1.40 above 1000GHz. We note that this task and thecorrection factors were available only after we had reduced andmeasured our spectra, but all fluxes and calculations presentedhere do include this offset correction, unless otherwise stated.The SPIRE FTS calibration is based on the spectrum ofUranus and detailed analysis of the flux uncertainties (Swinyardet al. 2014) which shows that the absolute flux on a source ofthis brightness is calibrated to within 6%. We acquired five double sideband pointing observations withHIFI, further referred to as settings A-E. Details are given in Ta-ble 1. Settings A (in band 1b), B (in band 2a), and C (in band 5a)were observed in the standard dual beam switch (DBS) mode,while the FastDBS mode was used for the high-frequency set-tings D (in band 7a) and E (in band 7b) in order to optimisebaseline stability. Data were obtained at nominal resolution forboth the high-resolution spectrometer (HRS) and the widebandspectrometer (WBS), in both the horizontal and vertical polari-sations. The bandwidth of the HRS is narrow, at 230 MHz, that is115 km s − and 40 km s − in settings A and E, respectively. Theantenna’s half power beamwidth is 34 (cid:48)(cid:48) .8, 31 (cid:48)(cid:48) .2, 18 (cid:48)(cid:48) .4, 12 (cid:48)(cid:48) .0,and 11 (cid:48)(cid:48) .9, for settings A through E (Roelfsema et al. 2012, seetheir eq. 3). The HIFI data obtained with the WBS were reducedup to level two using the pipeline available within HIPE. Base-line subtraction and standing-wave removal in the WBS spec-tra were severely hampered by the noise levels present in thedata. Additionally, the lack of coinciding features between thedata obtained in the vertical and horizontal polarisations led usto conclude that the HIFI spectra yielded no reliable emissionline measurements. This is likely a consequence of the lines be-ing too faint and/or too wide. The ortho-H O 7 , − , transitionat 166.8 µ m (1797.159 GHz), detected in the PACS observations,is covered by setting E, but no emission feature is seen in eitherpolarisation. Using the point source sensitivities for HIFI, listedby Roelfsema et al. (2012), we find that the flux measured byPACS can indeed not be retrieved at the sensitivities reached inour HIFI observations. Moreover, Tylenda et al. (2009) measuredan outflow with a terminal velocity of 215 km s − from numerouslines from V838 Mon, giving a total linewidth of ∼
3. V838 Mon: point source photometry
As is clear from Fig. 1, the emission on the maps consists ofa bright point source (V838 Mon itself) and fairly bright ex-tended emission surrounding the star; for SPIRE the two emis-sion sources are particularly intimate. To measure the flux den-sity from the star, aperture photometry using various tasks inHIPE was possible for the PACS data, but for SPIRE the relative strength and unevenness of the extended emission made suchmeasurements highly uncertain. We therefore took an alternativeapproach for SPIRE, which we also tested for PACS: to scale thePSF (point spread function; beam) and subtract it from the mapsuntil only residual emission remained at the position of the star.The photometry of the scaled beam is then the photometry ofV838 Mon. As a refinement of this method, we also performedPSF-subtraction on deconvolved SPIRE maps. Although the re-sults were not as useful as hoped, we believe a summary of thiswork will be useful for others (App. C).Jumping ahead to the final results for the point source(Sec 6.1): within the 2 σ errors there is no change in the point-source photometry between the three epochs, which cover 16months. The properties of the dust giving rise to this emissionhave not changed noticeably during this period. For the aperture photometry we adopted aperture sizes narrowenough to avoid including too much extended emission, andadopted the appropriate aperture and colour corrections: seeApp. D for the details. The errors in the photometry arisingfrom the scatter in the sky background on the maps are 0.02 Jy(70 µ m), 0.04 Jy (100 µ m), and 0.12 Jy (160 µ m), and these val-ues were folded into the errors arising from the photometric mea-surements. Another source of uncertainty comes from the con-taminating flux from the extended emission at the position of thestar, which while not strong – the ratio of the peak flux fromthe star to that in the immediately surrounding extended emis-sion is of the order 40-30-5 for the blue-green-red maps – it isnon-negligable. To account for this, we also measured the fluxvalues from the extended emission close to the position of thestar and consider this to be the sky flux to be subtracted from theaperture photometry. The photometry given in Table 2 is thenmean ± range of these two sets of measurements.We also performed PSF-subtraction on the PACS maps ofV838 Mon, using maps of the PSF taken from the PACS cali-bration web-page: see below and App. E for more detail. Themain reason for doing this was to obtain maps of the extendedemission only, but the method can also be used to obtain pho-tometry for V838 Mon. Details of the photometric corrections(aperture and colour corrections) made on the scaled PSF aregiven in App. D, and the method itself is explained in App. E.The photometry resulting from this are also given in Table 2. Wefind V838 Mon to have a FWHM consistent with that of the PSFsource we used, but with a 2d profile that is somewhat different:residuals in the subtraction of the point source from V838 Monindicate a difference in the core and the wings of the profile (seeApp. E). This is the main contributor to the photometry errorsarising from this process.The PSF-subtracted maps were used in the study of the ex-tended emission. After subtracting out the point source, the re-gion left behind (which contains the above-mentioned residuals)was replaced by a circle of constant flux (value taken from thelocal background). To measure the photometry of V838 Mon from the SPIRE mapswe used a PSF-subtraction method. Aperture photometry wasnot feasible because of the relative brightness of the extendedsource: indeed at 500 µ m the extended source almost completelyoverwhelms the point source (see e.g. Fig. E.1). We used the Article number, page 6 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
Fig. 3.
Rectified SPIRE spectrum (unapodised) from 600 GHz to 1550 GHz (black) overlaid with the fitted spectrum of the three major speciesCO, SiO and H O (red). We note that this spectrum does not include the scaling for the flux losses due to off-centring (see text). maps produced in HIPE and calibrated for point sources, thatis with units of Jy/beam, to obtain the point-source photometry,but performed the same procedure on the maps calibrated for ex-tended sources to create point-source-free maps of the extendedemission.The PSF-subtraction method is essentially a scale-and-subtract, and the details of this work, which was carried out inHIPE, can be found in App. E. The photometry is then mea-sured from the beam maps which are scaled sufficiently to re-move the point source emission from the V838 Mon maps whenimage subtraction is done. We used a SPIRE “useful” photom-etry script in HIPE for the photometry measurements, using thesource fitting photometry task ( sourceExtractorSussextractor ),which works directly on the point-source calibrated maps. SeeApp. D for details of the photometric corrections. The errors inthe photometry with this method are dominated by the range ofvalues that result in acceptable PSF-subtraction residuals, andfor the 500 µ m maps are especially large as the beam is largecompared to the size of the extended emission; here we add thatthe result from our PSF subtraction from the deconvolved maps(App. C) is that the fluxes at this wavelength may be slightlyoverestimated. The values and errors given in Table 2 are takenfrom the acceptable range of PSF scaling factor values.After subtracting out the point source, the region left behindwas replaced by a circle of constant flux (value taken from thelocal background). These cleaner PSF maps were then used tostudy the extended emission.
4. V838 Mon: extended source photometry
For PACS, the sum of the flux in all the pixels within a certaincontour on the maps, including the point source, was measured.Subtracting the flux of the point source then gave the extendedsource flux density values. The contours were chosen for eachmap to encompass the flux close to the level of the sky, with asecondary consideration given to having the area similar for eachepoch: the contours values taken lie between 2 σ and 4 σ abovethe mean background value on the residual maps, where σ is thestandard deviation of the sky flux as measured from about 10apertures.The SPIRE maps used were those calibrated for extendedemission (units MJy/sr). From these measurements we sub-tracted the point-source photometric measurements to obtain thephotometry of the extended emission only. We compared thesevalues to the values obtained from measuring the photometry ofthe extended source from the PSF-subtracted maps (Sec. 3.2: theextended emission with the point source subtracted): these agreeto better than ± Article number, page 7 of 24 &A proofs: manuscript no. v838mon_v5
Table 2.
PACS and SPIRE point and extended source spectro-/photometric measurements of V838 Mon (except for those not applicable “na”).All values have been colour, aperture, and beam corrected (see the text). One-sigma measurement errors including the uncertainty due to thebackground subtraction are given. The calibration uncertainties (as well as other sources of error) are discussed in the text, and are 5% for PACSand 6% for SPIRE photometry. Spectroscopy values are given for comparison to the more accurate photometry.
Point source Extended source (point-source-free)Band, Aperture PSF Spect. a Flux density Area Size b epoch photometry subtraction cont.Jy Jy Jy Jy (cid:3) (cid:48)(cid:48) (cid:48)(cid:48)
70, 1 7 . ± .
12 6.8–7.4 6 . . ± .
12 8860 90 × . ± .
12 7.0–7.4 - 2 . ± .
12 853070, 3 7 . ± .
11 6.8–7.2 - 2 . ± .
11 7270100, 1 3 . ± .
07 3.2–3.6 3 . . ± .
07 13390 130 × . ± .
04 3.5–3.7 - 11 . ± .
04 11400100, 3 3 . ± .
06 3.5–3.7 - 10 . ± .
06 12760160, 1 1 . ± .
15 1.1–1.5 2 . . ± .
15 14250 140 × . ± .
15 1.2–1.5 - 14 . ± .
15 14340160, 3 1 . ± .
15 1.1–1.4 - 14 . ± .
15 16820250, 1 na 0 . ± . ∼ . . ± .
10 20100 160 × . ± .
06 - 6 . ± .
07 17400250, 3 na 0 . ± .
08 - 5 . ± .
08 17600350, 1 na 0 . ± . ∼ . . ± .
07 20100 140 × . ± .
05 - 2 . ± .
06 14600350, 3 na 0 . ± .
06 - 2 . ± .
07 18500500, 1 na 0 . ± . ∼ . . ± .
05 17700 160 × . ± .
06 - 0 . ± .
06 15400500, 3 na 0 . ± .
07 - 1 . ± .
07 21000
Notes. ( a ) Measured from the spectral segment closest to the photometry wavelength. ( b ) Mean diameter at position angles ± ◦ . The photometric results given in Table 2 show that there isno difference, within generally less than 2 σ , in the point sourcefluxes from epoch to epoch. For the extended source, from look-ing at the maps of Fig. 4 and at the better detail in the point-source subtracted deconvolved maps (App. C), it does seem thatthe flux in the northern part of the source decreases with timecompared to the flux in the southern part of the source, at allwavelengths (excepting perhaps 500 µ m). However, these andany other differences between epochs are mostly within 3 σ ,and the derived dust properties do not change significantly withepoch (c.f. Sec. 6.1).
5. V838 Mon: point source spectroscopy
The V838 Mon PACS spectra were measured in HIPE, fittinga low-order polynomial and a Gaussian to all features with thewidth and peak intensity to appear to be emission lines. We thencompared the line list to that from the red supergiant VY CMa:this star has a well-studied PACS–SPIRE spectrum, and it hasa cool and extended atmosphere (Royer et al. 2010; Matsuuraet al. 2014), as V838 Mon is suggested to have. The comparisonwas used as a guide for a second pass over the spectra to look foradditional features and eliminate those most likely to be noise.Most of the lines we identified this way are of water. The PACSspectra have units of Jy and µ m; the line fluxes we report arethe integrated intensity of the fitted Gaussian in units of W m − (computed as peak × σ × (2 π ) . × × − /λ , where σ is thewidth of the Gaussian and the last phrase is to convert the units).Measurement errors were propagated according to the standardrules. Table 3 gives the positions, fluxes and uncertainties of allthe lines identified in the PACS (and SPIRE) spectra. The lines for which we could not assign a clear identification are given inTable F.1.The PACS lines are redshifted by ∼ . µ m (blue tored wavelengths) compared to our identifications, correspondingto 60–80 km s − . The PACS pipeline produces wavelengths cor-rected to LSR: converting to heliocentric produced redshifts of40–60 km s − which is close to previous redshift reported val-ues 58 km s − (Tylenda et al. 2009) or 71 km s − (Tylenda et al.2011). Some of our measured shift, however, will be caused byV838 Mon being located slightly off from the centre of the cen-tral spaxel (PACS Observer’s Manual ). V838 Mon lies betweenthe central spaxel (12) and its neighbour (17) by less than 1/2spaxel, and this will cause the spectrum obtained from the cen-tral spaxel to have a redshift of ∼ . µ m (blue to red wave-lengths) – although this effect is diluted somewhat since our finalspectrum is a combination of spaxel 12 (slight redshift) and 17(slight blueshift).In Table 2 we have indicated the approximate values of thespectral continuum flux at the wavelengths of the photometricbands. The spectroscopy and photometry match within their re-spective uncertainties. Fourier transform spectrometers produce spectra in which lineshave sinc profile functions (Naylor et al. 2014) that are linearlysampled in frequency. Therefore, all our FTS spectral measure-ments have been made in frequency space and we have retainedthe natural sinc profile for all spectral measurements. The inte-grated line strengths were obtained using a built-in HIPE linefitting script that fits a sinc profile to a given line position usingleast square fitting of the instrumental line profile (Naylor et al. see http://herschel.esac.esa.int/twiki/pub/Public/PacsCalibrationWebArticle number, page 8 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis Fig. 4.
SPIRE and PACS maps with the point source removed and residual replaced by a blank patch (on the 70 µ m some residuals remain). Scalingranges are chosen based on the peak flux density. PACS: -0.0005–0.001 Jy/pixel (70 µ m), -0.0004–0.006 Jy/pixel (100 µ m), -0.0001–0.02 Jy/pixel(160 µ m), with pixel sizes of 1 (cid:48)(cid:48) , 1.5 (cid:48)(cid:48) , and 3 (cid:48)(cid:48) , respectively. SPIRE: -0.005–0.03 Jy/pixel (250 µ m), -0.005–0.025 Jy/pixel (350 µ m), -0.005–0.02Jy/pixel (500 µ m), with pixels sizes of 4.5 (cid:48)(cid:48) , 6.25 (cid:48)(cid:48) , and 9 (cid:48)(cid:48) , respectively. N–E is indicated in the central maps. Article number, page 9 of 24 &A proofs: manuscript no. v838mon_v5 Table 3.
PACS and SPIRE line fluxes and errors: 1- σ signal-to-noise ratio for the SPIRE lines and the measurement error for the PACS lines. Thecalibration uncertainties are: PACS: 8(blue)–12(red)% SPIRE: 6% with an additional 1 . × − W m − arising from the continuum rectification.The PACS lines start at 1660 GHz. All lines in our line list are given here, including those for which the error is larger than the value. Species Observed ν (GHz) Rest ν (GHz) δν (GHz) Flux a δ F a H O 4 , –3 , b c O 1 , –1 , O 5 , –4 , O 2 , –2 , O 4 , –3 , O 5 , –4 , d O 5 , –4 , e O 2 , –1 , d O 2 , –1 , e O 3 , –3 , O 1 , –0 , O 3 , –2 , g O 6 , –5 , O 3 , –3 , O 4 , –4 , O 2 , –2 , O 7 , –6 , O 8 , –7 , O 6 , –5 , O 5 , –5 , O 7 , –6 , O 6 , –5 , Notes. ( a ) − W m − b ) J (cid:48) K (cid:48) a , K (cid:48) c – J K a , K c ( c ) J (cid:48) – J ( d ) Seen in SLW Detector ( e ) Seen in SSW Detector ( g ) Blended with CO 10–9 O, we fitonly these lines in the SPIRE spectra. All of these species haveclear lines present, in both the PACS and SPIRE spectral ranges.Our line-list fitting to the SPIRE spectra is shown in Fig. 3. We note that there are features clearly visible in the spectrum whichwe have not fit. Some of these could be emission lines that werenot on our line list, while others will be noise. An advantage offitting a sinc profile is that it is not just the main peak that needsto fit the data, but the minor peaks in the wings also. The fittingerrors are based on the fit to the entire profile, and hence givea more reliable estimate than if fitting the SPIRE data apodisedto a Gaussian profile. However, a disadvantage of spectra with asinc profile is that the multiple-peak nature of this profile leads
Article number, page 10 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
Table 3. – continued.
Species Observed ν (GHz) Rest ν (GHz) δν (GHz) Flux a δ F a H O 2 , –2 , O 2 , –1 , O 6 , –6 , O 7 , –6 , O 7 , –7 , O 3 , –4 , O 6 , –6 , O 7 , –7 , O 7 , –7 , O 6 , –5 , O 9 , –9 , O 7 , –7 , O 5 , –4 , O 8 , –8 , O 5 , –5 , Notes. ( a ) − W m − to some noise spikes being as bright as real lines. It is for thisreason that we fit the SPIRE spectra with a line list: we know theexact position of the expected lines and fit only those. All linesfrom the line list are included in Table 3, and where the fittingerrors exceed the integrated flux, the measured flux is clearly tobe considered unreliable.
6. Analysis
The total integrated
Herschel flux density for the central (pointsource) emission and the large-scale extended infrared emission,for each epoch, are plotted in Fig. 5. In addition, the full SPIREspectrum of the central point source (smoothed) is shown in lightgrey. The near- to far-infrared spectral energy distribution of thecentral (point source) is shown in Fig. 6 (
Herschel -only and withliterature data added). The fitting results are given in Tables 4and 5.
Cool extended source
A modified black body was fitted tothe total integrated extended far-infrared
Herschel emission (in-cluding data from all epochs) given in Table 2. The emitted fluxdensity is modelled with the functional form: S ν ∝ κ ν B ν , with κ ν = . ν/ β = . µ m /λ ) β cm /g (c.f. Hilde-brand 1983; Beckwith et al. 1990), assuming a gas-to-dust ratioof 100. The power-law index β is an indicator of dust grain com-position. β ∼ β ∼ β ranges from 1.7–2.6 (Paradis et al. 2010; Gordon et al.2014). For this cool material the Herschel wavelength coverageprovides a good constraint to the dust temperature, T dust , dustmass, M dust , and the emissivity index, β . The best-fit result isshown in Fig. 5 and values for the fitted parameters for the ex-tended emission are labelled in the figure.The dust temperature for the extended source as derivedfrom Herschel data is T dust = . ± . + − derived by Banerjee et al. (2006) from 2005 Fig. 5.
Spectral energy distribution (SED) of the unresolved centralpoint source and the extended far infrared emission observed with
Her-schel
PACS and SPIRE. The blackbody curve shown for the central(point source) emission is derived from the near- to far-infrared SEDshown in the right panel. The modified black body shown for the ex-tended emission is the best-fit to the
Herschel measurements (averagedover all epochs).
Spitzer–MIPS photometry. The fitted dust emissivity index, β , is2.43 ± ± (cid:12) (at 6.2 kpc).Banerjee et al. (2006) derived a value of 0.9 M (cid:12) (converted toa distance of 6.2 kpc). Their higher mass value is because of theirhigher measured flux densities for the extended emission takenfrom Spitzer data at 70 and 160 µ m from 2004/5. The difference( ∼
11 Jy at 70 µ m compared to our ∼ Article number, page 11 of 24 &A proofs: manuscript no. v838mon_v5
Table 4. Upper : ODR fitting (scipy.odr) results of
Herschel
SED of theextended infrared emission (Sec. 6.1).
Lower : average results from 2dfitting to the PACS and SPIRE maps of the extended emission (Sec. 6.2).
Epoch M dust T dust β M (cid:12) K1 0.54 ± ± ± ± ± ± ± ± ± ± ± ± Warm point source
Using only the
Herschel far-infrared ob-servations of the central (unresolved) point source emission, wecannot extract values for the degenerate parameters of dust massand dust temperature. Therefore we have extended the SED ofthe central point source to the near-infrared with
Spitzer andGemini mid-infrared results for the unresolved central sourcein 2007 (Wisniewski et al. 2008). We note that we excludethe
Spitzer
Spitzer µ m point-sourceflux value from 2007 is 7.34 Jy (Wisniewski et al. 2008) whichis essentially the same as our measurements, giving some jus-tification to combining the data of 2007 with our later epochs(2011/13). Also available are AKARI/IRC flux densities at 8 and19 µ m (Ishihara et al. 2010), and while these wavelengths are ex-pected to arise primarily from the central source it is not clearfrom the data how much extended source contamination couldbe present.We fit these data with a pure black body and with a “dust”emission model, that is using one or two modified black bod-ies similar to the procedure adopted for the extended emission.Since it was unclear which of the literature data we can safelyuse – given it is unknown whether there were any intrinsic vari-ations between then and now – we chose to fit all possible com-binations, and we report on the best of these here: the Gem-ini+ Herschel + Spitzer data; the Gemini+
Spitzer data; and
Her-schel +Gemini data. The results of these fittings are given in Ta-ble 5 and Fig. 6.The ranges of values presented in Table 6 show that the pureblack body fits give a temperature in the range 300–350 K anda radius of 150–200 AU. The uncertainties on fit parameters arethose provided by the least-square-fitting routine. The modifiedblack body fit yielded less satisfactory results, so those of thesimple blackbody are more secure, but with this modelling itis possible to estimate a mass of the material. The temperaturerange is a slightly lower, 230–330 K, and the mass estimates liearound 9 × − M (cid:12) . The values for β are all not far from zero:that is close to a black-body. We note that the reliability of ourmass estimate depends on the appropriateness of the dust model.The uncertainties we quote in Table 5 are the statistical onesfrom the least-square fitting. The systematic uncertainties for thisare much more difficult to ascertain and can be easily factor twoto few, especially since we have no constraints from, for exam- Table 5.
Results of fitting models to the point source using the datafrom all Herschel epochs and data from the literature, and when usingselected sets of data only
Blackbody Modified blackbodyGemini+
Herschel + Spitzer M dust (M (cid:12) ) na (8.3 ± . × − T dust (K) 353 ±
31 237 ± β na 0.43 ± ± Spitzer M dust (M (cid:12) ) na (1.1 ± . × − T dust (K) 317 ± ± β na -0.02 ± ± Herschel +GeminiM dust (M (cid:12) ) na (7.9 ± . × − T dust (K) 353 ±
36 234 ± β na 0.49 ± ±
11 naple, dust features. A more elaborate model is beyond the scope ofthis work since we lack simultaneous multi-wavelength data athigh spatial resolution and have in fact very limited informationon the actual geometry and nature of this emission.Loebman et al. (2015) studied V838 Mon with mid-IR datataken in 2008 and derive a blackbody temperature of the dustof 285 K (since they have the shorter wavelengths their result ismore secure than ours).
Dust temperature maps of the extended emission were createdfrom pixel-by-pixel fitting of the measured flux densities with amodified black body ( S ν ∝ B ν ν β , as above). First the maps at70, 100, and 160 µ m were resampled (Swarp; Bertin et al. 2002)to a common grid with identical pixel sizes, and subsequentlyconvolved to match the point spread function (PSF) of either the250 µ m or 350 µ m SPIRE photometry map. For this we adoptedthe convolution kernels and algorithm from Aniano et al. (2011).The sky background was computed in an “empty” region of themap and subtracted. A fit was then only attempted if the observedflux densities are higher than 10 σ sky of the sky background at atleast three wavelengths.We let β run first as a free parameter and used the mapsconvolved to the 350 µ m SPIRE PSF. The resulting tempera-ture, dust mass, and emissivity index ( β ) maps (epoch 1 only),as well as the corresponding percentage error maps, are shownin Fig. 7. The noise-dominated low-emission edge regions suf-fer from degeneracy between dust temperature and dust emissiv-ity, and so these regions have a higher uncertainty than the er-rors indicate. In the next iteration we omitted the 350 µ m SPIREdata to improve the spatial fidelity of the dust images. To de-rive dust properties we use all maps convolved to 250 µ m. Againwe limit the fitting to pixels with significant flux above the skynoise level at at least three wavelengths. In addition we fix β to the value found in fitting also the lower resolution 350 µ mdata. This approach improves the spatial resolution of the finaldust maps, while retaining as much as possible information onthe dust opacity (through β ). Fixing β to a single global valuerestricts the fits and limits the physical interpretation possible.This procedure was repeated for each of the three epochs. Article number, page 12 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis F l u x D en s i t y [ Jy ] Spitzer(2007)ModifiedBlackbodyBlackbodyGemini(2007)PACS/SPIRE(2011)PACS/SPIRE(2012)PACS/SPIRE(2013) Wavelength [ µ m] − F l u x D en s i t y [ Jy ] Spitzer(2007)ModifiedBlackbodyBlackbodyGemini(2007)PACS/SPIRE(2011)PACS/SPIRE(2012)PACS/SPIRE(2013)
V838Mon - central source infrared SED F l u x D en s i t y [ Jy ] Spitzer(2007)ModifiedBlackbodyBlackbodyGemini(2007) Wavelength [ µ m] − F l u x D en s i t y [ Jy ] Spitzer(2007)ModifiedBlackbodyBlackbodyGemini(2007)
V838Mon - central source infrared SED [CASE 2] F l u x D en s i t y [ Jy ] ModifiedBlackbodyBlackbodyGemini(2007)PACS/SPIRE(2011)PACS/SPIRE(2012)PACS/SPIRE(2013) Wavelength [ µ m] − F l u x D en s i t y [ Jy ] ModifiedBlackbodyBlackbodyGemini(2007)PACS/SPIRE(2011)PACS/SPIRE(2012)PACS/SPIRE(2013)
V838Mon - central source infrared SED [CASE 4]
Fig. 6.
Near- to far-infrared spectral energy distribution of the unresolved central point source, including flux densities measured with differentcombinations of
Spitzer , Gemini, and
Herschel data.
The final dust temperature, dust mass and emissivity mapsof V838 Mon for each epoch are shown in Fig. 8. Mild varia-tions are seen between the three epochs, most noticeably in thetemperature of epoch 2 (being slightly higher within the lobes)but the pixel value distribution of T dust , M dust , and β agree witheach other between the three epochs. The mean results from 2dfitting are given together with the ODR fitting results in Table 4.The comparison in results for these two approaches is acceptableand we do not consider any of the differences to be significant. Spitzer maps
The point source and extended source of V838 Mon are very inti-mate at the FIR wavelengths. Even for the relatively small beamof
Herschel it was necessary to use a point-source extractionmethod to separate the point-source from the extended source atthe longer wavelengths; at the shorter wavelengths it was neces-sary to have the knowledge obtained from point-source extrac-tion to find an appropriately small aperture with which to doaperture photometry.
For the point source:
With the 2005
Spitzer data, Banerjee et al.(2006) used a point-source extraction method to obtain 70 µ m flux of ∼ µ m ; with the 2007 Spitzer data, Wisniewski et al. (2008) used aperture photometryto obtain a 70 µ m flux of ∼ µ m and160 µ m fluxes of ∼ ∼ For the extended source:
Banerjee et al. report fluxes for the ex-tended source at 70 µ m and 160 µ m of ∼
11 Jy and ∼
17 Jy, Wis-niewski et al. make no specific statement (which implies a non-detection), and we measure 70 µ m and 160 µ m fluxes of ∼ ∼
11 Jy. The comparison of these epochs indicates that thepoint source flux increased with time and the extended sourceflux decreased.In our opinion a comparison of the
Spitzer and
Herschel µ m results should only be taken as indicative: the pointsource is so much fainter than the extended source and the Spitzer beam so large, that the reliability of the results is nothigh. But a comparison of the results at 70 µ m should be morereliable.We decided to make a direct comparison of the total source(the point plus extended source) from the two epochs of Spitzer data with our
Herschel data. The 2005
Spitzer maps that wereused by Banerjee et al. are public. We downloaded these MIPSobservations (AOR 10523648 and 10523904), and from the
Article number, page 13 of 24 &A proofs: manuscript no. v838mon_v5
Fig. 7. Top panels : Dust temperature, dust mass, and dust emissivity maps of V838 Mon corresponding to the maps convolved to 350 µ m andwith β a free parameter, for epoch 1. Colour-bars in each panel indicate the parameter range. Maps are 3.5 (cid:48) × (cid:48) . Bottom panels : Correspondingpercentage error maps. Orientation is the same as in Fig. 1.
Fig. 8.
From left to right: Dust temperature (first column), dust mass (second column), dust emissivity index (third column), and PACS 160 µ mflux density (fourth column) maps of V838 Mon created from the maps convolved to 250 µ m and β derived from the 350 µ m convolved data set.From top to bottom: Epoch 1, 2, and 3. See text for details on the procedure. Orientation is the same as in Fig. 1. "maic" FITS files we created single images at each wavelengthusing the Montage code. We then convolved the Herschel im-ages to the beam size of the
Spitzer data (18 (cid:48)(cid:48) at 70 µ m and 40 (cid:48)(cid:48) at 160 µ m) in HIPE, and converted to the MJy/sr units of the Spitzer data. A single value was subtracted from all maps to do http://montage.ipac.caltech.edu a basic background subtraction. We then measured the fluxesfrom the entire source in an aperture of 80 (cid:48)(cid:48) from each map(this being the aperture used by Banerjee et al.). We then did thesame for the 2007 Spitzer data used by Wisniewski et al. (AORs21476096, 21475840, and 21475072 from 2007, for 70 µ m and160 µ m both). Applying the appropriate colour corrections to themeasurements is made difficult by the different temperatures of Article number, page 14 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis the point and extended sources (we measure here a total flux),but in any case are only of the order a few percent. At 70 µ m wetook a straight average of the respective corrections at 300 K and20 K and at 160 µ m we used the cool value.We obtain: Spitzer (2005): 15.1 Jy at 70 µ m, 16.1 at 160 µ m; Spitzer (2007): 12.4 Jy at 70 µ m, 12.1 at 160 µ m; PACS(2011/12): 11.4 Jy at 70 µ m, 15.4 Jy at 160 µ m. (Banerjee et al.report fluxes for the 2005 data of 14.7 and 17.5 Jy; the slight dif-ference with our measurements can easily be account for by ourvery rough background subtraction.) We also compared all im-ages directly to each other. The Spitzer images are slightly largerin the blue and slightly smaller in the red, but we conclude thatthere is no significant difference in the spatial extent of the emis-sion between 2005 and 2011/12.Previous investigations by the Herschel Science Centreshowed that PACS and MIPS total flux and surface brightnessagree within 5–20% and no systematics were found . Takinginto account also a 10% calibration uncertainty quoted for bothinstruments, and our rough background subtraction, the differ-ences in the total source flux that we measure (a drop with timeat 70 µ m and drop then rise at 160 µ m) are barely significant.In conclusion: looking only at the 70 µ m results (where thepoint source and extended source are the least blended) the fluxof the point source rose between 2005 and 2007 and then stayedthe same, while the flux of the total source stayed the same(or dropped slightly) and its extent on the sky remained thesame. Therefore, the brightness of the extended source must havedropped. Without longer wavelength multi-epoch data we cannotunfortunately say more. To interpret the line emission detected from CO, we plot the en-ergy diagram (Fig. 9). The data are taken from Table 3 and to theflux errors there we have added the calibration uncertainties andSPIRE continuum rectification uncertainty (Secs 2.3 and 2.4).The CO energy diagram shows two discrete components: a coldcomponent with a peak of the line intensity at the upper stateenergy E up <
55 K (lower than J=4–3), and a warm componentwith a rising trend with increasing E up . Our working hypothesisis that the cold component is associated with ISM gas – althoughit is also possible that it arises from a cold shell (à la Lynch et al.2004) – while the warm component is associated with the cir-cumstellar envelope/environment (i.e. arising out of the outflowtriggered by the stellar impact). We stress that the spectra mea-sured by PACS and SPIRE come from what they see as the pointsource, and at the spatial resolution at our wavelengths ( ∼ (cid:48)(cid:48) inat the shortest wavelength and ∼ (cid:48)(cid:48) at the longest), this area isquite large (2.7 to 11 pc at the distance of 6.2 kpc).We first compare the distribution of the CO fluxes on theenergy-level diagram for V838 Mon with those of what canbe considered similar stars (i.e. having an extended atmo-sphere), and which are well-studied at these wavelengths: thered-supergiant, VY CMa and the AGB star, W Hya. Here weare treating VY CMa and W Hya as templates of circumstellarenvelopes created by a constant mass loss, and we wish to seehow/where their mass loss differs to that of V838 Mon. It is truethat the mass loss from these stars has been ongoing for over athousand years, a timescale much longer than that for V838 Mon,but it is general trends we are looking for. It is also true that the The calibration reports (PICC-NHSC-TR-034 and PICC-NHSC-TN-029) will be provided on the Herschel Legacy Library web-pages, whichare currently being set up mass-loss rate of VY CMa is highly variable (Decin et al., 2006;Matsuura et al., 2014). However, using the values in Decin et al. we find that only beyond ∼ (cid:63) which takes a timescale ofover 1000 years to reach, are the mass-loss rate variations likelyto have an effect on the measurements. Hence we should avoidonly this part of the comparison.The CO line intensities of these stars were taken from Khouriet al. (2014) and Matsuura et al. (2014), and these intensitieshave been scaled to the distance of V838 Mon (6.2 kpc); tak-ing 1.14 kpc for VY CMa (Choi et al. 2008) and 78 pc for WHya (Knapp et al. 2003, Justtanont et al. 2005). The currentmass-loss rates of VY CMa and W Hya are estimated to be2 × − M (cid:12) yr − , and 2 × − M (cid:12) yr − . The overall CO lineintensities of the energy diagram reflect the mass-loss rates ofthese two objects. The mass-loss rate of VY CMa is higher thanW Hya, so VY CMa has brighter CO line intensities (after ad-justing for its greater distance). The energy diagram shows thatV838 Mon has fairly strong CO lines compared to these othertwo stars (scaled to the same distance).We can also see that V838 Mon has a different CO curve tothose of W Hya and VY CMa: it has a two-component curve,with a clear peak at E up ∼
50 K, and a gradual rise with in-creasing E up , while for the comparison stars there is an initialrise followed by a flattening. Khouri et al. (2014) explained thathigher J (higher E up ) transitions trace the inner and warm partof the circumstellar envelope, while lower J transitions trace theouter and cooler part. The clear peak for V838 Mon thereforelies part of the curve that traces the outer and cooler part, whichwould presumably correspond either to the first epoch of massloss, that is, that occurring during the outburst, or to pre-existingmaterial. The difference in the slopes at the higher E up trace dif-ferences in the inner part of the circumstellar environment. Theslow rise in the CO curve for V838 Mon is because that the warmgas component is a single temperature zone, as our later fittingresults will show: see also explanation Fig. 10. For such a gas,the general driver of the curve on the CO diagram is the temper-ature, and it is the peak of the overall curve that indicates the gastemperature. The two evolved stars we compare to have a morecomplex story: for example, the temperature of gas from AGBstars is know to vary thanks to their more complex and longermass-loss history. Their curve on the CO diagram will thereforenot be the same as that for a single-temperature zone.The curve of V838 Mon shows that the molecular envelopeconsists of a cooler outer part and a warmer inner part, and thecomparison to the evolved stars suggests that its envelope nothave been made with a constant mass loss from the star. A good(and perhaps obvious) suggestion is that the material is associ-ated with material erupted in 2002, which has remained close tothe star and warm. The figure (specifically the comparison of thedistance-scaled fluxes) also shows that the scale of gas eruptionwas large: one event of an eruption could have ejected mate-rial equivalent to thousands years of constant mass loss from ourcomparison stars.We next derived physical conditions of the two CO-emittingcomponents, by modelling the energy diagrams using the non- Values we adopt for this back-of-the-envelope calculation are: dis-tance to the star= 1500 pc, expansion velocity= 35 km s − (slower atthe inner region), stellar radius=1 . × cm. The initial velocity ac-celeration region is at 10 R (cid:63) , with velocity ∼ − , and the masstakes ∼
100 years to cross. Then, from 10–100 R (cid:63) we assume a con-stant mass-loss rate, with velocity gradually accelerating to ∼
20 km s − ,and covering this distance takes 230 years. Finally, from 100–1000 R (cid:63) there are several changes in the mass-loss rate, and this has a crossingtimescale of ∼ &A proofs: manuscript no. v838mon_v5 Fig. 9. Top:
The energy diagram for the V838 Mon CO lines, com-pared with the red superigant, VY CMa and the AGB star, W Hya.The X-axis indicates the upper state energy of the CO transition ( E up )in Kelvin, and the Y-axis shows the line intensities of the CO lines inW m − . V838 Mon shows two discrete components (warm and cold) ofthe CO energy distribution. Symbols are data and dotted lines are the fits(ours to V838 Mon, and using the values from the respective papers forthe other two stars). Bottom:
The energy diagram for CO lines detectedby SPIRE and PACS (black symbols), where the X-axis shows the upperstate energy of the CO transition ( E up ) in Kelvins, and the Y-axis showsthe line intensities of the CO lines in W m − . Two discrete componentsof CO energy distributions were detected, which were modelled with18 K (blue) and 400 K (red). The sum of these warm and cold compo-nents is plotted in orange. LTE radiative transfer code RADEX (van der Tak et al. 2007).RADEX calculates the level populations of molecules, we usedCO–H and SiO–H cross-sections and Einstein A-coefficientsfrom the LAMDA molecular and atomic data base of Schöieret al. (2005), which were adopted from calculations by Yanget al. (2010) and Dayou & Balança (2006). We adopted a dis-tance to the source as 6.2 kpc.The parameters derived and the fitting errors are summarisedin Table 6, and in Fig. 9 we show the model fitted results. Wenow move our attention to a word about the fitting. The er-rors (1- σ measurements+calibration error) are greatest for thecold component, and so simple χ fitting was weighted to thewarm component and ignored the cold component. To overcome Fig. 10.
Demonstration of how the CO energy diagram changes as afunction of the kinetic temperature. This applies to the optically thingas and with a constant beam size across all the transitions, which is notthe case for V838 Mon. However, this figure is just to demonstrationhow the curve changes with temperature.
Table 6.
CO model parameters with errors derived from the fitting. warm Cold T kin (K) 400 ±
50 18 ± ∆ v(FWHM) (km s − ) 200 5 (fixed) N CO (cm − ) (1 . + . − . ) × a (1 . + . − . ) × M CO ( M (cid:12) ) 7 . × − Notes. ( a ) This range calculated for a temperature value fixed to that given inthis table this limitation, we ( χ ) fit the cold and warm components inde-pendently, and summed them. Strictly-speaking, a full radiativetransfer modelling of the cold component behind the warm com-ponent is the correct path to follow, but given the gaps in ourknowledge of V838 Mon, we consider this an unnecessary step,and we ignore the small fraction of the ISM behind the warmcomponent. The straight sum of the cold and warm componentsis therefore shown in Fig. 9. Inspection of the fits on the plots re-sulted in some adjustment so we did not underestimate the cold–warm cross-over point at E up ∼
150 K. The resulting χ -fittingerrors in the temperature are given in Table 6. The errors for N CO were calculated for a fixed temperature.We derive temperatures of 400 K and 18 K for the two com-ponents. In this and the following discussion, we assume thatthese two components originate from different gas clouds: thecold component is ISM which is ubiquitous over the observedarea (the point source diameter) and our results show it is opti-cally thin, and the warm component comes from the stellar out-flow and is smaller than the Herschel point source diameter.
Cold component
The cold component could either be coldmaterial associated with the outer parts of the molecular enve-lope, or it could be pre-existing ISM. Our data cannot distinguishbetween these two cases. The χ fitting gave a temperature of18 K, and a best value for the H density of ∼ × cm − for thecold component, although this value is very poorly constrainedand could be a factor of 100 lower. In the following we assumea constant density (which is a good assumption for ISM gas). Article number, page 16 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
The line width of the cold component has been fixed to 5 km s − ,as a typical CO line width of molecular gas towards V838 Monwas found to be 2–7 km s − from low J CO lines (Kaminski et al.2007). For the cold component, we assume that the SPIRE beamsare filled by this cold CO gas, and SPIRE beam sizes, which havea frequency dependence, are taken from Makiwa et al. (2013).The estimated column density of the cold CO component is then1 × cm − . Warm component
When fitting the warm component we findthat CO line intensities are almost independent of the (assumed)H density when it is > × cm − : anything between 10 and 10 cm − is possible. The parameters resulting from the fit-ting are given in Table 6. We note here that in the RADEX cal-culations, the escape probability used is that for a slab/cylindermorphology. Lynch et al. (2004) suggests a model of a hollowsphere, and this model is one we wish to follow as much as pos-sible in our calculations. The typical optical depth of the COlines for V838 Mon as estimated by the RADEX modelling ofthe Herschel spectra is 3. With this optical depth, the escapeprobability for a slab is about 30% different to that for a sphere:this difference is very small, especially compared to our otheruncertainties, and so we attempted no changes to the RADEXmodelling process.With the fitting results, can we estimate the mass of the gas?This is a multi-step process, since we need to estimate the emit-ting area of the gas, which is not an output of the model fitting.First, can we estimate the FWHM of the CO emission lines?One of the model input parameters is the line width (FWHM; ∆ v), for which we did not have a direct measurement (but seebelow). The ∆ v of the warm component is linked with two keymodel parameters: the CO emitting area, and the saturation limitof CO lines. In the optically thin case, CO line intensities gener-ally increase with higher column density. Once the line becomesoptically thick, the line intensities become saturated, and willnot increase as the column density increases. Among commonly-observable molecules, CO lines tend to saturate with relativelysmall column densities. The saturation limit is also coupled withthe line width at a given column density. RADEX defines theoptical depth at the line centre ( τ ) as τ ∝ N / ∆ v, where N isthe column density (van Langevelde & van der Tak 2012). Nar-rowing ∆ v results in lowering the line intensities at the saturationlimit.A key point of ∆ v is its link to the CO emitting area. The sizeof the gas shell created by the eruptive event in 2002 is limited bythe expansion velocity and the time since the eruption. Let us as-sume that the warm gas has expanded at a constant velocity (v exp )since the eruption, and, as we will discuss later, v exp = ∆ v /
2. Wewill set the CO emitting area to be Ω co = πθ / . × sr,where θ = ∆ v / ∗ t / d in arcsec, t is the time since the erup-tion, d is distance to V838 Mon, and 4 . × is the conver-sion factor between sterad and arcsec . As the RADEX outputin erg cm − s − is given as the emitting area for the whole sky(4 π ), to compare with the measured CO line intensities, whichcome from a smaller area than the whole sky, the model fluxeswere converted as f CO (W m − )= f CO (erg cm − s − ) Ω co / (4 π ).Hence, the modelled values of f CO (W m − ) are ∝ ∆ v . Thisanalysis method follows the one used for the explosive eventof supernova 1987A (Matsuura et al. in preparation; Kamenet-zky et al. 2013). This equation also gives the maximum possi-ble line intensities for a given ∆ v, because of the CO saturationlimit. Using this, we narrowed down the range of acceptable ∆ vvalues as derived from our line fluxes. By changing ∆ v over a 50 km s − grid, we established a limit to the ∆ v which translatesto FWHM >
200 km s − : narrower than 200 km s − under-predictsthe CO intensities.A value of 200 km s − is the lower limit for the intrin-sic FWHM obtained from modelling the fluxes, but what canwe measure from the spectra? The lines in the SPIRE spectra,which has an instrumental resolution of 280–970 km s − , are un-resolved. One CO line is detected in PACS at 162 µ m wherethe instrumental resolution is ∼
240 km s − . The line is rathernoisy, but can be fitted with FWHM values of 240–325 km s − – at this upper limit the intrinsic FWHM (decoupling the mea-sured and instrumental widths) is ∼
220 km s − . So the upperlimit line width measured from one CO line is close to the lowerlimit taken from the modelling of the fluxes. To be more certainwhether the emission line widths support the modelling, it wouldbe necessary to obtain higher resolution spectra. For the calcu-lations we do here, we will adopt a value of 200 km s − for theintrinsic FWHM of the gas.Next, what expansion velocity does this FWHM translateinto? The FWHM of a line arising from a hollow sphere (thinshell) is 2v exp (Robinson et al., 1982), and from a filled sphereis √ exp (although the average velocity could be lower; Mc-Cray, 1993). Taking a value of 200 km s − for the FWHM givesan expansion velocity of the gas of 100 km s − for the hollowsphere case, and over 10 years this gives a radius of 210 AU; forthe filled sphere case the radius is 297 AU. Expansion velocitiesof various values are reported in the literature as measured fromdifferent atomic and molecular lines, with values varying from100–400 km s − (Loebman et al., 2015); Tylenda et al. (2009)measure a terminal velocity of 215 km s − in 2005. However,there are no contemporaneous measurements of the CO veloc-ity for us to compare to.Finally, can we estimate the warm CO gas mass? We willestimate the CO mass as r N CO π , which equation assumes aslab geometry. Lynch et al. (2004) suggests a model of a hol-low sphere, however as long as the (averaged) density alongthe line of sight remain the same, it does not matter if the ge-ometry is a slab or a hollow sphere. With a column density of1 × and a radius of 210 AU, the CO gas mass is estimated tobe 7 . × − M (cid:12) (see Table 6). We note however that the sourcesof uncertainty in this value are many – the geometry, the radius,the uncertainty in the column density – and this value should betaken as an estimate only. SiO and H O A similar analysis is applied to SiO, withRADEX modelling (Fig. 11) to all the measured transitions(measurement+calibration error included). The collisional cross-sections and Einstein coefficients were taken from the LAMDAmolecular and atomic data base (Schöier et al. 2005), whichwere adopted from calculations by Dayou & Balança (2006).The SiO line intensities have larger uncertainties than CO ones,and the transitions covered are E up >
380 K, much higher thanCO. Fig. 11 shows that the kinetic temperatures of SiO could liein the range 400–1200 K, and that the spectral shape cannot befit well with a simple component, but needs to involve multipletemperature components. The range of the temperature suggeststhat SiO has a a similar or higher temperature than the CO warmcomponent – we cannot constrain the options further, and canonly state that the SiO-emitting region is associated with CO-emitting regions or lies closer to the heating source, which ispresumably the central star.
Article number, page 17 of 24 &A proofs: manuscript no. v838mon_v5
We did not carry out a similar analysis for the H O linesas they even more sensitive to the model used and assumptionsadopted.
Fig. 11.
The energy diagram of SiO (black symbols are data and mea-surement+calibration error) with RADEX models with 400 and 1200 Kcurves plotted.
7. Discussion and conclusion
We remind the reader that the point source includes everythinginside the PSF of PACS and SPIRE: the PACS 70 µ m beam is ∼ (cid:48)(cid:48) , or 55000 AU at the distance of V838 Mon, and that of theSPIRE 500 µ m beam is ∼ (cid:48)(cid:48) .The SPIRE and PACS spectra show CO, SiO and H Omolecules. The analysis of CO suggests presence of two com-ponent temperature, namely warm ( ∼
400 K) and cold ( ∼ µ m CO line width andCO saturation limit, and assuming the model of an thin spheri-cal shell. Our estimated warm dust radius lies in the range 150 to 200 AU, a value obtained by fitting our data with previously-taken literature data. These values are consistent with the find-ings quoted previously.Lynch et al. (2004) reported the detection of CO rotational-vibration bands in the spectra up to 1.3 years after the out-burst. The estimated rotational-vibrational temperature of COwas 650 K. Our SPIRE and PACS spectra, which were obtainedabout nine years after the eruption, showed the warm CO gas hasa kinetic temperature of 400 K. These two measurements showthat the temperature has dropped over the last eight years. If oneassumes that the eruptive shell is very thin, with a uniform den-sity within, and that this shell has been expanding constantlysince 2002, with about ten years between 2002 and our Her-schel observations, the corresponding increase in volume andadiabatic cooling would normally lead to the temperature drop-ping by a factor of 100. That we do not see this suggests that anearby heating source is present.V838 Mon has a B3 V companion, located 28–250 AU dis-tant (Munari et al., 2007; Tylenda et al., 2009). Could the com-panion be heating the gas? The obvious answer is "yes", giventhe luminosity of this type of star, however modelling its specificeffect on the gas outflow from V838 Mon is beyond the scopeof this paper. The gas temperature we measure is what we find,whatever the source of the heating. For the dust it is possible tomake a quick and simple calculation of the possible heating bythe companion. Using equations from Tielens (2005: 5.43 and5.44): G = . × ( L ∗ / L (cid:12) )(0 . / d ) (1)where we can adopt L ∗ =1000 L (cid:12) and d =
200 AU: the radiationfield, G ∼ . × . For amorphous carbon dust ( β =
1) andtaking a representative grain radius, a, of 0.1 µ m): T dust = . µ m / a) . (G / ) . → T dust ∼ , (2)or, for silicate dust ( β = T dust = / a ) . ( G / ) / → T dust ∼ K . (3)Naturally the devil lies in the details. The grain radius, for ex-ample, depends on the environment in which they form – butincreasing the radius to 1 µ m only decreases the temperature by1.6 and 1.15 for the two types of dust. These temperatures areclose to the value we measure for the warm dust ( ∼
300 K), andso we conclude that at least part of the source of the dust heatingcould be the companion star.One puzzle is why dust continued to cool from ∼
750 K to280 K in 10 years as found by Loebman et al. (2015) and thiswork ( ∼
300 K), while the gas temperature is slightly higher(400 K). One possibility is that gas and dust do not have identi-cal temperatures, although this is more commonly found in lowdensity regions with less frequent collisions, and is very unlikelyto happen in reasonably dense ( > × cm − ) gas, such asV838 Mon appears to have.In our RADEX analysis we have adopted a spherical geom-etry. In fact this is probably not really the case. As well as thefinding of Chesneau et al. (2014), evidence for a non-sphericalmorphology comes from Wisniewksi et al. (2003). They foundthat spectroscopic and spectropolarimetric variations took placeduring the outburst, this being indicative of asymmetric geome-try; this result is backed up the study of Desidera et al. (2004),who also found polarimetric changes during 2002. Further po-larisation variations were observed a few months after the out-burst by Wisniewski, Bjorkman & Magalhães (2003), again sug-gesting the presence of an asymmetrical geometry of scattering Article number, page 18 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis material close to the star, and moreover that the distribution ofthis material had experience significant changes. It is thereforemost likely that the circumstellar environment of V838 Mon isnot spherical. However, to conduct a more detailed modellingof the gas and dust in the region, its temperature and chemicalstructure, would require more information about this morphol-ogy.We estimated a dust mass of 9 × − M (cid:12) for the point source.This is based on a very simple modified blackbody fitting, andwe note that the uncertainty in this value is high since it is ob-tained by fitting to previous epochs of data.We estimated a CO mass of 7 . × − M (cid:12) , based on an as-sumed geometry and shell radius, and with factors of few un-certainty from the underlying CO column density. To turn thisinto a total mass we need a CO/H ratio: we adopt here a ratio of4 × − , which is the O-rich ratio for late-type stars (Willacy &Millar 1997, and taken because of the physical rather than chem-ical similarity to V838 Mon’s extended environment). This thengives a total gas mass of ∼ . M (cid:12) . Taken together, these twovalues give a gas-to-dust mass ratio of 200. However, this valueshould only be taken as indicative, as the sources of uncertaintyare large. If the outflowing gas from which the CO lines arise ismolecular, the conversion from CO to total mass has only a fac-tor of a few uncertainty. However, if the gas is neutral or ionised,then the total gas mass could be much higherIf the outflowing gas from which the CO lines arise is molec-ular, our adopted value has only a factor of a few uncertainty.However, if the gas is neutral or ionised, then this value is com-pletely wrong. Hence, a measurement of the gas-to-gust massratio requires more and better data. We also point out that it isnot clear that the gas and dust we have detected arise from thesame component: bear in mind the large size of the point sourcethat Herschel sees (the beam is 55,000 AU at the shortest wave-lengths). One should also consider that the dust created in theoutflow from V838 Mon – created most likely as a consequenceof an infalling low-mass star – may not conform to the usualISM. Little is known about the type and properties of dust cre-ated in these eruptive systems.
The extended source detected in our PACS and SPIRE photom-etry is shown in the maps in Fig. 4. This region extends over ≈ . Herschel epochs(1.5 years). Our multi-epoch photometry observations show thatthere has been no clear evolution in the morphology or the fluxesof either the point or the extended source over the 1.5 years. Ourfluxes also do not differ strongly from what was found prior toour
Herschel observations. Loebman et al. (2015) conclude thatno new dust has formed around the point source since a dustevent taking place between 2004/5 and 2006/7 (Wisniewski et al.2008), as they find no increase in the IR fluxes since then. Our70 µ m fluxes are the same as the Spitzer
Spitzer maps of 2005 and 2007 (as used by Banerjee et al. 2006and Wisniewski et al. 2008), and also compared the size of theemission from the extended source. There is slight evidence fora decrease in the 70 µ m flux from the total source between 2005and then 2007/2011, but no difference in the extent of the emis-sion when comparing the three epochs of maps at the same flux levels. However, Wisniewski et al. (2008) find that the pointsource flux increased between 2005 and 2007 (from about 3.8 Jyat 70 µ m to 7.3 Jy) and our PACS 2011/12 value is the same asthe Spitzer
Herschel maps we measured a dust mass of0.57 M (cid:12) , which while less than the 0.9 M (cid:12) measured by Baner-jee et al. (2006) (converted to a distance of 6.1 kpc), is still highenough to result in a large gas+dust mass (if taking the standardISM value of 100), and hence we agree with the conclusion thatthis material is ISM rather than arising from the star. We mea-sured a temperature of ∼
19 K dust for this source, which is onlyslightly warmer than the T kin ∼
12 K measured by Kaminski,Tylenda & Deguchi (2011) from CO lines from the ISM aroundV838 Mon.As mentioned in the introduction, van Loon et al. (2004)found evidence for multiple mass-loss events prior to the 2002outburst: a 7 (cid:48) × (cid:48) dust shell is visible in the IRAS data, and theMSX data also show an extended object (1.5 (cid:48) diameter) at the po-sition of V838 Mon. The latter exactly matches with the size ofthe infrared emitting region detected with Herschel . The SPIREimages of V838 Mon reveal patches of faint emission at ∼ (cid:48) .5south of V838 Mon, but the maps are not large enough to verifyif these structures are part of a larger structure corresponding tothe IRAS dust shell.With our PACS and SPIRE spectra we also detected a cold( ∼
18 K) component of CO gas, which has a temperature closeto that measured for the dust. This low temperature suggests thatthe cold component is associated with ambient ISM gas. TheISM is also responsible for millimetre CO lines which were de-tected by Kaminski et al. (2007), and are most likely associatedwith the light echo detected by Bond et al. (2003). Kami´nskiet al. (2011) found the kinetic temperature of CO to be 12–15 ± × cm − from Kami´nski et al. (2011) while we ob-tained 1 × cm − . Some of the discrepancy could be associ-ated with the uncertainty of H density, as a higher density ofthe collisional partners would require a lower column density toreproduce the same line intensity within that range (this does notapply to warm component). Another possibility is that the coldCO is not from the ISM, but rather from a cold shell that is as-sociated with the warm shell (e.g. the Lynch et al. 2004 model),and hence is part of the Herschel point source.Our
Herschel study of V838 Mon has left us with many ques-tions. What is the relationship between the gas and the dust – arethey co-spatial, and how were they each created by the outburstor within the outflow? How many temperature zones are presentin the outflow, and does the cold CO gas arise from the ISM oris it directly related to the warm gas? How much dust has beencreated in the outburst outflow? When will the star return to itspre-outburst state? All very interesting questions for the future.
Acknowledgements.
We thank John Wisniewski for a thoughtful refereeing pro-cess that resulted in an improved paper. NLJC and KME acknowledge supportfrom the Belgian Federal Science Policy Office via the PRODEX Programme ofESA. PACS has been developed by a consortium of institutes led by MPE (Ger-many) and including UVIE (Austria); KU Leuven, CSL, IMEC (Belgium); CEA,LAM (France); MPIA (Germany); INAF-IFSI/OAA/OAP/OAT, LENS, SISSA(Italy); IAC (Spain). This development has been supported by the funding agen-cies BMVIT (Austria), ESA-PRODEX (Belgium), CEA/CNES (France), DLR(Germany), ASI/INAF (Italy), and CICYT/MCYT (Spain). SPIRE has been de-veloped by a consortium of institutes led by Cardiff University (UK) and in-
Article number, page 19 of 24 &A proofs: manuscript no. v838mon_v5 cluding Univ. Lethbridge (Canada); NAOC (China); CEA, LAM (France); IFSI,Univ. Padua (Italy); IAC (Spain); Stockholm Observatory (Sweden); ImperialCollege London, RAL, UCL-MSSL, UKATC, Univ. Sussex (UK); and Caltech,JPL, NHSC, Univ. Colorado (USA). This development has been supported bynational funding agencies: CSA (Canada); NAOC (China); CEA, CNES, CNRS(France); ASI (Italy); MCINN (Spain); SNSB (Sweden); STFC and UKSA(UK); and NASA (USA). This research has made use of the SIMBAD database,operated at CDS, Strasbourg, France. Herschel is an ESA space observatory withscience instruments provided by European-led Principal Investigator consortiaand with important participation from NASA.
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Article number, page 20 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
Fig. A.1. Top : crosses mark the centre of each of the PACS spaxels onthe 160 µ m map. Bottom the non-vignetted circles of the SSW bolome-ters (black) and the SLW bolometers (white) on the 250 µ m map. Theposition of V838 Mon is marked with a square. Appendix A: Footprint of the PACS and SPIREspectrometers
As mentioned in Secs 2.3 and 2.4, the footprint of the spectrom-eters extends over the extended emission around V838 Mon, andsome of this extended emission “contaminates” the point source.In Fig. A.1 we show the footprint of the PACS and SPIRE spec-trometers plotted on two of our epoch 1 maps.
Appendix B: Producing a calibrated point-sourcePACS spectrum
The HIPE task extractCentralSpectrum works by adding to-gether the spectra of the spaxels that contain the flux of the pointsource, and then applying a correction for the shape of the beam(which is much wider than a single spaxel): effectively it per-forms a wavelength-dependent aperture correction to the fluxlevels in the spectrum. This correction, contained in the PACScalibration tree, is to be used on sources located close to the cen-tre of the central spaxel of the 5x5 IFU, in particular becauseonly for this central spaxel has the PACS beam been well cal-ibrated (private communication). The task can extract only thecentral spectrum or the sum of the central 3x3 spaxel box, andthe point-source flux density loss correction applied takes thisinto account. For our data of V838 Mon, the star was locatedalmost exactly between the central spaxel and a neighbour, andhence the task provided could not be used. However, the same procedure could be used, and the resulting spectrum would thenbe more correct than a simple combination of the brightest threespaxels which contain the flux of V838 Mon (private communi-cation with the PACS instrument team).Therefore we modified the procedure of extractCentralSpec-trum to take into account our non-central position for V838 Mon.For each of our 20 spectral segments we did the following: 1. Extracted the summed spectrum of the 3x3 spaxel box sur-rounding our brightest spaxels (those with most of the fluxfrom V838 Mon).2. Multiplied that by the values in the point-source flux-loss correction tables from the calibration tree ( cal-Tree.spectrometer.pointSourceLoss using the central-3x3-to-total values).3. Then extracted the spectrum only of the brightest two spax-els within the previous 3x3 spaxel box: this spectrum has ameasurably superior signal-to-noise ratio than that of step 1.4. Fit the continuum of the spectra from steps 2 and 3, and di-vided the first by the second.5. Then multiplied that ratio into the spectrum from step 3.The result is a spectrum with the best possible SNR but correctedfor the point-source flux density losses.
Appendix C: Using deconvolution to assist withthe point source photometry
The main difficulty with measuring the point source and ex-tended source fluxes separately from the PACS and SPIRE mapsis that the point source is located within the extended emission.For the PACS maps it is possible to still see the two separately,but for the SPIRE maps – with a fainter point source and a largerbeam – the two are particularly intimate.To try to improve the results of the PSF subtraction, wetested our PSF extraction method (App. E) on maps which hadfirst been deconvolved: the separation between the point and ex-tended source will then be greater and PSF subtraction should beeasier. As a bonus, we would also have deconvolved images ofthe extended IR bright region to study.The deconvolution was based on a maximum entropy method(MEM) and follows in general the scheme described by Holliset al. (1992) for HST images and Ottensamer et al. (2011) forHerschel/PACS. In the deconvolution task, the multiplier is con-volved with the PSF model. We tested with the PACS and SPIREbeams (as used also in our PSF-subtraction work) but chose in-stead to use observations of the point source AFGL 3068, whichshows no circumstellar emission and gave very clean results.In the beginning the multiplier is the PSF itself, which is thenscaled by the flux ratio of the image and the PSF, and finallyre-convolved. At each iteration step the fluxes are compared andthe residuals become part of the new multipliers. The iteration isstopped when artefacts become visible, which are mainly man-ifest as negative fluxes around the central source. The deconvo-lution was applied to the PACS and SPIRE maps of V838 Moncreated by Scanamorphos, and we also deconvolved the maps ofAFGL 3068 to then use them for the subsequent PSF-subtractionphotometry, using the same technique as outlined in App. E.The photometry resulting from the PSF subtraction methodon the deconvolved maps is reported in Table C.1. Unfortunately,this process did not offer as much of an improvement as we hadhoped. For the PACS 70 µ m and 100 µ m maps deconvolution didnot make a much difference, and the problem of residuals result-ing from a mis-match between the adopted PSF shape and thereal PSF for our V838 Mon observations remained (see App.. E).For the PACS 160 µ m maps the deconvolution resulted in a slightimprovement: the scatter in the resulting photometry is lower,and the results agree better with the aperture photometry (whichare the preferred results for PACS) reported in Table 2. For theSPIRE maps, doing PSF subtraction on the deconvolved mapswas a challenge, partly because the deconvolution appeared toalter the FWHM of the PSF differently to that of the V838 Mon Article number, page 21 of 24 &A proofs: manuscript no. v838mon_v5 maps – an extra smoothing step was necessary for the subsequentPSF subtraction. The photometry obtained from the PSF subtrac-tion on these maps is in the same range as reported in Table 2,however, the 500 µ m values are lower than those found previ-ously, and this probably does indicate that our previous resultsare slightly too high.In summary, using deconvolution to aid in the PSF subtrac-tion did not help as much as hoped in our work because it wasclear that a very good knowledge of the beam for your particularobservation (i.e. subjected to the same observing plan, the samedata reduction, and the same type of map-making) is crucial toachieving both a good deconvolution and a good PSF subtrac-tion. Given the aims of this paper, we did not feel it was justifiedto spend more time on the deconvolution. However, in our opin-ion this method would work if more time could be invested init. Table C.1.
PACS and SPIRE point source fluxes as obtained from PSFsubtraction on deconvolved maps. Beam and colour corrections havebeen applied.
Band ( µ m) (Epoch) Flux density (Jy)70 (1) 6.7–7.3 (2) 7.0–7.6 (3) 7.0–7.6100 (1) 3.8–4.1 (2) 3.8–4.1 (3) 3.6–3.9160 (1) 1.2–1.3 (2) 1.3–1.4 (3) 1.2–1.3250 (1–3) 0.53–0.62350 (1–3) 0.21–0.27500 (1–3) 0.08-0.15The images resulting from the deconvolution are presented inFig. C.1. For these we have subtracted the point source. It is clearthat the overall features at each wavelength do not change withepoch, with the possible exception for the PACS maps, where itappears that the northern lobe becomes fainter with time. Appendix D: Photometry corrections
For aperture photometry of the point source for the PACS andSPIRE data we used the aperture sizes and corrections given inTable D. The PACS aperture corrections were taken from pho-tometer.apertureCorrection in calibration tree 65 (the fm7 val-ues), except for the aperture photometry for the scaled PSF maps,for which we used apertures of 57.6 (cid:48)(cid:48) , and the appropriate aper-ture corrections provided in the ASCII files that accompanied thetarball of the beam files (see below for details of the beams). Thecolour corrections for PACS were taken from report on the PACScalibration wiki page (dated April 12, 2011: cc_report_v1.pdf),and the colour and beam-area corrections for SPIRE were takenfrom the SPIRE calibration tree and a SPIRE photometry scriptprovided in HIPE.
Appendix E: Details on the PSF-subtractionmethod
For performing PSF subtraction using the scaled PACS/SPIREbeams, we took the beams from: – PACS: a tarball of FITS files (“PACSPSF_PICC-ME-TN-033_v2.0”) from the PACS documentation web-site from the
Herschel portal. We used the examples created from Vesta,with scan speed 20, pixfrac 0.1, from OD 345, and witharray-to-map angle +42. – SPIRE: the normalised beams created in 2012 andprovided on the SPIRE documentation web-page from the
Herschel portal (specifically the page: her-schel.esac.esa.int/twiki/bin/view/Public/ SpirePhotometer-BeamProfile.The steps were:1. Resample the PSF maps to the pixel sizes of the V838 Monmaps.2. Rotate the PSF maps to match the position angle of theV838 Mon maps: the shape of the PSF depends on the scan-ning speed, angle, and the position angle (this is more truefor PACS than for SPIRE). We note that rotation was doneon the WCS, rather than on the map as an image.3. Measure the sky position of the PSF and of V838 Mon towithin a pixel, using sourceExtractor . Change the WCS val-ues in the PSF maps so the PSF star and the point source ofV838 Mon are in the same position.4. Scale and subtract the fluxes of the PSF maps from theV838 Mon maps (in WCS space, using the HIPE task im-ageSubtract ).5. Determine the best scale-and-subtract value via visual in-spection of the residual maps and plots of cuts taken at var-ious position angles through the residual maps. The accept-able scaling factors are those that result in residuals that arethe flattest.6. Finally, replace the fluxes in a small circle around the sub-tracted star with the value of the flux surrounding that re-gion, to create cleaner PSF-subtracted maps. These wereused to create residual-free images for this paper, and touse as a check on the extended-source photometric measure-ments (Sec. 4).Fig. E.1 shows the SPIRE maps after having scaled-and-subtracted the PSF (Sec. 2.2). Also shown are cuts taken throughthe original and the residual maps at various position angles, anda cut through the scaled beam. It is from these plots and figuresthat the decision about the best scaling factors was made.The PSF-subtraction process did not work for PACS as wellas it did for SPIRE. We always had negative residuals, aris-ing from a difference in the shape of the beam from the PSFmaps and the effective shape of beam for our V838 Mon maps.It is noted in the report accompanying the PACS beams thatthe beam-shape depends also on the data reduction and map-making so that a re-reduction of the PSF data, following thesame reduction and map-making methods used for the astronom-ical observations, is recommended. However, as we use the PSF-subtraction method for the PACS data only as a check, we didnot do this. An additional consideration is that it was necessaryto rotate the beams (the supports of the
Herschel secondary leadto a tri-lobal shape for the beams, which orientation depended onthe position angle). Rotation inevitably involves an interpolationof the images, and it is possible that this could also account forthe mismatch between our point source and the beam profiles.The tri-lobal shape (and hence the resampling of that pattern)is more prominent in the PACS than the SPIRE beams, and thiscould also explain why the mismatch is greater for the PACSbeams.
Appendix F: PACS spectral lines without a clearidentification
Lines from the PACS spectra for which we could not find a clearidentification, but for which the detection is formally significant,are given in Table F.1. We checked these all against a basic modelspectrum based on the parameters of our point source, to look
Article number, page 22 of 24. M. Exter et al.: On the properties of dust and gas in the environs of V838 Monocerotis
Fig. C.1.
The deconvolved maps with the point source subtracted and with the subtracted region replaced with a constant value. Epochs, wave-lengths, and orientation are indicated on the maps. Scaling is min–max for each image.
Table D.1.
Details of the apertures and the multiplicative corrections used in the photometric measurements of the point source band ap. radius ap. correction colour correction b beam-area correction a point extendedPACS 70 6.06 (cid:48)(cid:48) (cid:48)(cid:48) (cid:48)(cid:48) (cid:48)(cid:48) (cid:48)(cid:48) (cid:48)(cid:48) Notes. ( a ) to correct to a point source with a spectral index of α =
2, rather than -1, as is applied in the pipeline ( b ) assuming a spectral index of α = for more CO and H O lines, but no matches could be found.Other potential lines include SiO, CO+, and SO , but a propermodelling of these lines is necessary before this can be takenfurther. Table F.1.
PACS spectral lines without a clear identification.
Wavelength Flux Wavelength Flux( µ m) (10 − W/m ) ( µ m) (10 − W/m )70.74 5 . ± .
59 159.12 1 . ± . . ± .
52 159.32 1 . ± . . ± .
51 159.47 3 . ± . a . ± .
71 162.37 2 . ± . . ± .
61 165.11 0 . ± . a . ± .
00 179.96 1 . ± . . ± .
40 181.12 1 . ± . . ± . Notes. ( a ) wide line Article number, page 23 of 24 &A proofs: manuscript no. v838mon_v5 Fig. E.1.
Examples of subtracting the scaled PSF from the epoch 1 SPIRE maps.