The IRAS 08589-4714 star-forming region
H. P. Saldaño, J. Vásquez, M. Gómez, C. E. Cappa, N.U. Duronea, M. Rubio
aa r X i v : . [ a s t r o - ph . S R ] S e p Manuscript for
Revista Mexicana de Astronom´ıa y Astrof´ısica (2007)
THE IRAS 08589 − H. P. Salda˜no , J. V´asquez , M. G´omez , C. E. Cappa , N.Duronea and M. Rubio Draft version: June 21, 2018
RESUMENSe presenta un an´alisis de la regi´on de formaci´on de estrellas masivas IRAS08589 − Herschel empleandocriterios basados en ´ındices de color y distribuciones espectrales de energ´ıa(SEDs). Las SEDs de algunas de las fuentes y el perfil radial de intensidad dela fuente m´as brillante en la banda de 70 µ m de Herschel (IRS 1) son modela-dos mediante el c´odigo unidimensional de transporte radiativo DUSTY. Paraestos objetos, se estiman las masas de las envolventes, los tama˜nos y densi-dades de las mismas y las luminosidades. Estos par´ametros indican que setrata de objetos muy j´ovenes, masivos y luminosos en las primeras etapas delproceso de formaci´on. Se emplean los diagramas color-color en las bandas deWISE y de 2MASS para indentificar potenciales objetos j´ovenes en la regi´on.Aquellos identificados en las bandas de WISE estar´ıan contaminados por laemisi´on de PAHs. Se emplea la distribuci´on de la emisi´on en el infrarrojo en70 y 160 µ m, para estimar el gradiente de temperatura del polvo. Esto indicaque la regi´on de formaci´on estelar masiva cercana RCW 38, localizada ∼
10 pcde la posici´on de la fuente IRAS, podr´ıa contribuir a la fotodisociaci´on del gasmolecular y al calentamiento del polvo interestelar en las inmediaciones de lafuente IRAS. ABSTRACTWe present an analysis of the IRAS 08589 − Herschel images using color index criteria and spectral energy distributions(SEDs). The SEDs of some of the infrared sources and the 70 µ m radialintensity profile of the brightest source (IRS 1) are modeled using the one-dimensional radiative transfer DUSTY code. For these objects, we estimatethe envelope masses, sizes, densities, and luminosities which suggest that theyare very young, massive and luminous objects at early stages of the formationprocess. Color-color diagrams in the bands of WISE and 2MASS are used toidentify potential young objects in the region. Those identified in the bands ofWISE would be contaminated by the emission of PAHs. We use the emission Observatorio Astron´omico, Universidad Nacional de C´ordoba, C´ordoba, Argentina. CONICET, Consejo Nacional de Investigaciones Cient´ıficas y T´ecnicas, Argentina. Instituto Argentino de Radioastronom´ıa, CONICET, CCT La Plata, Villa Elisa, Ar-gentina. Facultad de Ciencias Astron´omicas y Geof´ısicas, Universidad Nacional de la Plata, LaPlata, Argentina. Departamento de Astronom´ıa, Universidad de Chile, Santiago de Chile, Chile. µ m, to estimate the dust temperaturegradient. This suggests that the nearby massive star-forming region RCW 38,located ∼
10 pc of the IRAS source position may be contributing to thephotodissociation of the molecular gas and to the heating of the interstellardust in the environs of the IRAS source.
Key Words:
Stars: circumstellar matter — Stars: formation — Stars: massive— ISM: dust, extinction — ISM: individual objects (IRAS 08589 − n ∼ − cm − ), cold( T <
30 K), very opaque ( τ ∼ M ∼ M ⊙ )regions, known as pre-stellar cores (Sch¨oier et al. 2002; de Wit et al. 2009;Crimier et al. 2010; Battersby et al. 2014). Chambers et al. (2009) analyzeda significant number of objects having these characteristics and classified themas active cores if they were detected in the mid-infrared ( ∼ µ m) and asquiescent cores if no emission at these wavelengths was measured. Theseauthors concluded that active cores host and form massive stars whereas in-active cores are excellent candidates for starless massive cores, prior to theonset of the star formation process.However, unlike low-mass stars ( M < ⊙ ), high-mass stars lack a de-tailed formation scenario. For isolated low mass stars, the shape of the SEDsallows to classify them in four evolutionary classes (Shu et al. 1987; Lada1987; Andre et al. 1993), widely used in the literature. Several factors canbe invoked to account for our relatively less detailed knowledge of the forma-tion process/es of high-mass stars (such as: the distances, the relatively smallnumber of massive stars, the amount of energy and winds emitted by theseobjects, their short evolutionary time, the fact that they are deeply embeddedin the cloud material, etc.). One way to help to improve our understandingof the formation scenario of massive stars is to increase the number of youngstars with well-determined parameters and to study their environs.Beltr´an et al. (2006) cataloged a large number of massive clumps belong-ing to the southern hemisphere observed in the infrared (IR) continuum at 1.2mm. From their list, we selected the source IRAS 08589 − − Vela MolecularRidge (see zoomed region in the upper panel of Figure 1), which harborshundreds of low-mass Class I objects and a large number of young massivestars (Lorenzetti et al. 1993). Beltr´an et al. (2006) estimated a luminosityof 1.8 × L ⊙ (compatible with a previous estimate by Wouterloot & Brand Optical depth at 100 µ m. RAS 08589 − M ⊙ for this IRAS source. These authors classify thesource as an ultracompact HII region (UCHII) since it satisfies the criteria byWood & Churchwell (1989) , although no compact radio-continuum sourcehas been detected (S´anchez-Monge et al. 2013). In spite of its high mass, noCH OH maser emission was found towards this source (Schutte et al. 1993).However, S´anchez-Monge et al. (2013) reported water maser emission in theregion of the IRAS source.Bronfman et al. (1996) observed the region in the CS(2-1) molecular line.They found that the line has a central velocity V LSR = +4 . − anda velocity width at half-maximum ∆ V = 2 . − . In a recent survey,Urquhart et al. (2014) detected emission from ammonium molecular tracer,NH , in the (1,1) and (2,2) transitions towards the IRAS source. The centralvelocity coincides with that of the CS line.Bearing in mind velocities in the range 4 – 5 km s − for IRAS 08589 − − for the interstellar molecular gas.Ghosh et al. (2000) and Molinari et al. (2008) observed this IRAS sourcein the mid and far infrared and in the sub-millimeter range, and using theIRAS fluxes and those in the J, H, and K bands, measured by Lorenzetti et al.(1993), they constructed and modeled the SED, and found that the IRASsource has L ∼ L ⊙ , masses between 20 – 55 M ⊙ , and a B5 spectral type.The detection of the IRAS 08589 − § Herschel and WISE data and use the WISE color-color diagram toidentify YSOs in the region. In § µ m, pointing to the west, in the direction towards the RCW 38 highstar-forming region. In § §
5. 2. INFRARED SOURCES2.1 . Infrared data
The source IRAS 08589 − Herschel space tele-scope (Pilbratt et al. 2010) in the bands of the Photodetector Array Camera These criteria are based on IRAS fluxes satellite of a sample of ∼ S µm and S µm ≥
10 Jy,log( S µm / S µm ) ≥ S µm / S µm ) ≥ S indicates the flux andthe subscript the corresponding wavelength. SALDA ˜NO ET AL.and Spectrometer (PACS, 70 and 160 µ m; Poglitsch et al. 2010) and of theSpectral and Photometric Imaging Receiver (SPIRE, 250, 350, and 500 µ m;Griffin et al. 2010) instruments, with nominal FWHM of about 5 ′′ , 12 ′′ , 18 ′′ ,25 ′′ , and 36 ′′ in the five bands, respectively. The data were gathered as part ofthe Herschel Infrared Galactic Plane Survey (Hi-GAL, Molinari et al. 2010),which mapped the Galactic plane in mosaics of ∼ . ◦ × . ◦ , taken in thefive bands simultaneously. IRAS 08589 − ℓ, b ] = [268.3 ◦ , − ◦ ]. These mosaics were obtained in theparallel mode, at a scan velocity of 60 ′′ /s on 2012 November 13. The rawimages were reduced to level 1, using the HIPE software with standard pho-tometric scripts. The final maps in level 2 were processed with the software Scanamorphos , version 24 (Roussel 2013). Finally, we applied appropriateflux correction factors and color corrections to the PACS and SPIRE maps,respectively, according to the corresponding manuals .To extract the fluxes in each of the Herschel bands we applied the aperturephotometry method, using the HIPE software. We used aperture sizes of 20 ′′ and 25 ′′ in the 70 and 160 µ m bands, respectively, and of ∼ ′′ in all threeSPIRE bands. The background was estimated in a ring with internal andexternal radii of ∼ ′′ and 90 ′′ , respectively, centered on the peak emissionposition. The photometric uncertainty assigned to each flux is given by thestandard deviation of the measured values after aperture correction. Thesevalues are taken at different positions around the central source, without back-ground subtraction, using the same aperture as in the flux measurement forthe source (Balog et al. 2013). For IRS 1, 4, 5 and 6 fluxes relative errors arebetween 10 – 25 %. The remaining sources (IRS 2, 3, and 7) are contaminatedby the emission from the environment in which they are immersed and theirfluxes for λ > µ m have relative errors of ∼
50 %, preventing us frommodeling these sources.This region was also observed in the bands of the
Wide-field Infrared Sur-vey Explorer (WISE) satellite (Wright et al. 2010), at W1(3.4 µ m), W2(4.6 µ m),W3(12 µ m), and W4(22 µ m), with FWHM of about 6.1 ′′ , 6.4 ′′ , 6.5 ′′ and12.0 ′′ , respectively. The WISE images and fluxes were obtained from theIRSA ( NASA/IPAC Infrared Science Archive ) database.2.2 . Infrared emission and identification of YSOs The bottom panel of Figure 1 is a composite image of the WISE images at4.6 µ m (blue) and 12 µ m (green) and the Herschel image at 70 µ m (red) ofthe region IRAS 08589 − µ m emissionis visible to the west. This emission is likely produced by warm dust, while tothe east, the emission is dominated by colder dust, detected at λ > µ m.A transition zone from a hotter to a cooler region is also apparent in this See Data Analysis Manual Guide
Herschel and SPIRE Data Reduction Guide. http: //wise.ssl .berkeley.edu / http://irsa.ipac.caltech.edu/frontpage/ RAS 08589 − µ m) WISE band (in green) in the bottom panel, is likely producedby the emission of warm dust and polycyclic aromatic hydrocarbons (PAHs;Tielens 2008). For more details on this curved and elongated structure see §
4. In the bands of the
Herschel telescope we detected 7 IR sources (IRS,labeled from 1 to 7 in the bottom panel of Figure 1), four of which ( IRS 1,2, 3, and, 7) are also detected in the bands of the WISE telescope. From thefluxes in these bands, we determine the color indexes of the sources. Thesefour WISE detections are potential Class I or II objects, according to thecriteria by Koenig et al. (2012) . The left panel of Figure 2 shows the posi-tion of these sources on the WISE color-color diagram. The sources withoutWISE counterparts ( IRS 4, 5, and 6) are probably much younger objects(Chambers et al. 2009). IRS 1 and 4 are projected onto two dust clumps(blue contours in Figure 1) identified by Beltr´an et al. (2006) at 1.2 mm. Ta-ble 1 lists the sources detected in the WISE and Herschel images, with theircoordinates and fluxes, and their correlation with the dust clumps detectedat 1.2 mm.Green open squares on the left panel of Figure 2 correspond to faint WISEsources that lie in the region of the WISE color-color diagram contaminatedby PAH emission, according to Koenig et al. (2012) . PAH emission hasprominent lines in the WISE band W1 and W3, producing color excesses(Wright et al. 2010). These sources are listed in Table 2. Interestingly theseobjects are spatially located around the arc-shape structures near IRS 1, 5and 6. Figure 3 shows the bottom panel of Figure 1 (4.6 µ m in blue, 12 µ m ingreen, and 70 µ m in red), indicating with red triangles the positions of thesefaint WISE sources. None of these sources has been detected in the Herschel bands, with exception with those coincident with IRS 5 and 6. All sourcesin Table 2 are potential or candidate YSOs that require deeper observationsto confirm their nature and evolutionary status.The right panel of Figure 2 shows the 2MASS near-IR color-color J − H vsH − K diagram of the IRAS 08589 − These criteria state that Class I objects have WISE color indexes satisfying the fol-lowing inequalities: W1 − W2 > − W3 > − W2 − σ > − W3 − σ > .
0, where σ and σ are the combinederrors of W1 − W2 and W2 − W3, respectively. These authors concluded that sources contaminated by PAH emission have WISE colors,such that W1 − W2 < − W3 > SALDA ˜NO ET AL.
Fig. 1.
Upper panel:
Composite image showing the emissions at 250 (red), 160(green) and 70 µ m (blue) from Herschel of the IRAS 08589 − Bottom panel:
Compositeimage of the WISE images at 4.6 µ m (blue) and 12 µ m (green), in combination withthe Herschel image at 70 µ m (red) of the IRAS 08589 − Herschel sources labeled from IRS 1 to7 (see Table 1) are also indicated. The blue contours show the 1.2 mm dust emissionfrom Beltr´an et al. (2006). (Sicilia-Aguilar et al. 2006). This is confirming that these sources are young.RAS 08589 − -1 0 1 2 3 4 5 6 7 W - W W1 - W2
Class IClass II IRS 1 ★ IRS 3 ★ IRS 2 ★ IRS 7 ★ J - H H - K
IRS 2 ★ IRS 3 ★ IRS 7 ★ Fig. 2.
Left Panel:
WISE W1 − W2 vs W2 − W3 color-color diagram. The dot-ted lines limit the region where Class I and Class II objects lie according to theKoenig et al. (2012)’s criteria. IRS 1, 2, 3 and 7 are labeled and marked withstarred symbols. Green open squares are faint WISE sources likely contaminatedby PAH emission (Koenig et al. 2012). Average photometric errors are indicatedat the upper-right corner.
Right Panel: − H vs H − K color-colordiagram. The solid red and green lines mark the loci of the main sequence and gi-ant stars (Bessell & Brett 1988). The dashed lines delineate the reddening band forall main-sequence and giant stars (Rieke & Lebofsky 1985). Average photometricerrors are indicated at the upper-right corner.Fig. 3. Composite image of the WISE images at 4.6 µ m (blue) and 12 µ m (green),in combination with the Herschel image at 70 µ m (red) of the IRAS 08589 − S A L D A ˜ N O E T A L . TABLE 1SOURCES DETECTED TOWARDS THE IRAS 08589 − HERSCHEL
IMAGES AND WISE COUNTERPARTS
Herschel
Source α (2000.0) δ (2000.0) F F F F F a mmIRS (hh:mm:ss) ( o : ′ : ′′ ) (Jy) (Jy) (Jy) (Jy) (Jy) (Jy)1 09:00:40.6 − ± ±
11 201 ±
18 88 ±
20 25 ± . ± .
172 09:00:38.4 − . ± ± ±
12 27 ± ±
33 09:00:43.1 − . ± ± ±
18 40 ±
16 19 ±
94 09:00:48.0 − . ± ± ± ± ± . ± .
035 09:00:52.0 − . ± . ± ± ± . ± .
46 09:00:14.5 − ± ± ± ± . ± .
47 09:00:36.8 − . ± ± ±
12 16 ± ± α (2000.0) δ (2000.0) W1(3.4 µ m) W2(4.6 µ m) W3(12 µ m) W4(22 µ m) ID WISEIRS (hh:mm:ss) ( o : ′ : ′′ ) mag mag mag mag1 09:00:40.9 − . ± .
03 7 . ± .
02 4 . ± .
02 0 . ± .
02 J090040.97 − − . ± .
03 9 . ± .
03 5 . ± .
03 2 . ± .
04 J090038.59 − − . ± .
02 8 . ± .
02 6 . ± .
03 3 . ± .
03 J090043.08 − − . ± .
03 8 . ± .
02 5 . ± .
03 2 . ± .
06 J090037.06 − a Beltr´an et al. (2006). R A S − S T A R - F O R M I N G R E G I O N TABLE 2WISE SOURCES DETECTED TOWARDS THE IRAS 08589 − α (2000.0) δ (2000.0) W1(3.4 µ m) W2(4.6 µ m) W3(12 µ m) W4(22 µ m) ID WISE(hh:mm:ss) ( o : ′ : ′′ ) mag mag mag magWISE 1 9:00:12.9 -47:27:22.5 11.34 ± ± ± ± − ± ± ± ± − ± ± ± ± − ± ± ± ± − ± ± ± ± − ± ± ± ± − ± ± ± ± − a ± ± ± ± − a ± ± ± ± − ± ± ± ± − ± ± ± ± − a ± ± ± ± − ± ± ± ± − ± ± ± ± − ± ± ± ± − a Source with 2MASS counterpart. S A L D A ˜ N O E T A L . TABLE 32MASS SOURCES DETECTED TOWARDS IRAS 08589 − α (2000.0) δ (2000.0) J(1.2 µ m) H(1.6 µ m) Ks(2.2 µ m) ID 2MASS(hh:mm:ss) ( o : ′ : ′′ ) mag mag mag2MASS 1 a ± ± ± − a ± ± ± − a ± ± ± − ± ± ± − ± ± ± − ± ± ± − ± ± ± − a Source in the direction towards IRS 1.
RAS 08589 − . The spectral energy distributions In this Section, we analyze the evolutionary status of some of the sourcesidentified in the
Herschel images (see § Herschel bands at 70, 160, 250,350, and 500 µ m. Since protostars lie embedded in dense and opaque mate-rial we can derive the main characteristics of the envelopes surrounding thecentral objects. Fluxes for λ < µ m were excluded from the modeling of thewhole set of sources since the code used (DUSTY, see § µ m radial intensityprofile (see Sect. 3.2) simultaneously, while for IRS 4, 5, and 6 we model thecorresponding SEDs only since they are too faint to obtain reliable profiles at70 µ m. In the SED corresponding to IRS 1, we include the 1.2 mm flux fromBeltr´an et al. (2006) and the flux at 22 µ m from WISE while in the case ofthe SED for IRS 4 we only include the flux at 1.2 mm. This source does nothave a mid-IR counterpart. The SED of IRS 1 is shown in the left panel ofFigure 4, and those of IRS 4, 5, and 6 in Figure 5. -2 -1 +0 +1 +2 +3
1 10 100 1000 F l u x ( Jy ) λ ( µ m) IRS 1 -3 -2 -1 +0 N o r m a li z ed I n t en s i t y arcsec IRS 1
Fig. 4. SED (left panel) and intensity profile at 70 µ m (right panel) of IRS 1. Thefilled diamonds correspond to the fluxes from the Herschel and WISE telescopes.The open diamonds are the IRAS fluxes, shown as reference but not fitted. Errorsin the fluxes are indicated, except when they are smaller than the size of the symbol.The solid red line indicates the best model derived by the DUSTY code. The greenlines show the 20 DUSTY models comprised within the error bars of the fluxes. Thedotted line in the right panel is the instrument profile for PACS at 70 µ m (Lutz2012), used in the convolution with the theoretical intensity profile. . The 70 µ m intensity profile for IRS 1 The average azimuthal intensity profile is commonly employed to comparemodeled and observed images. This is used to complement the analysis of2 SALDA ˜NO ET AL. -1 +0 +1 +2
10 100 1000 F l u x ( Jy ) λ ( µ m) IRS 4 +0 +1 +2
10 100 1000 F l u x ( Jy ) λ ( µ m) IRS 5 +0 +1 +2
10 100 1000 F l u x ( Jy ) λ ( µ m) IRS 6
Fig. 5. SEDs of IRS 4 (left top panel), 5 (right top panel), and 6 (left bottompanel). The filled diamonds belong to the five
Herschel bands. The empty diamondin the SED for IRS 4 is an upper limit to the 70 µ m flux. The error bars areincluded. The relative errors in the fluxes are between 10 and 25%. The solid redline indicates the best model derived by the DUSTY code. The green lines show200 DUSTY models comprised within the error bars of the fluxes. SEDs, as it allows to reduce the number of free parameters, and thus mini-mizes the degeneration of the solutions (Crimier et al. 2010). We present themodeling of the intensity profile at 70 µ m corresponding to IRS 1.Since the PSF at 70 µ m is rather elongated (due to the scan speed, whichfor 60 ′′ /s gives a PSF of 5.83 ′′ × ′′ , Poglitsch et al. 2010), to determinethe radial intensity profile, several one-dimension radial profiles are obtainedin different directions from the center of the source. In the case of IRS 1(see Figure 4), we used only the hemisphere in the opposite direction to theapex of the curved structure seen in Figure 1 to avoid contamination from thesurrounding diffuse and extended emissions. The uncertainty in the profile isgiven by the noise in the image and the non-circularity of the beam of thesource. The latter effect is taken into account by the standard deviation ofthe average azimuthal flux. The profile is normalized by its maximum value.RAS 08589 − . The fitting method The SEDs of the four sources, as well as the 70 µ m intensity profile ofIRS 1, were modeled with the DUSTY code of Ivezic & Elitzur (1997). Thecode solves the radiative transfer in one-dimension for a dusty shell that sur-rounds a central source, which emits as a black body, whose radiation is ab-sorbed, scattered and re-emitted by the dust in the envelope. This envelopehas a density profile that follows a power law ( n ∝ r − p ), and it is parame-terized by the relative size given by Y = r ext /r in , where r ext and r in are theexternal and internal radii of the shell. The temperature of the central source( T star ) and the temperature at the internal radius of the envelope ( T in ) arefixed, with values of 15 000 and 300 K, respectively. The results are ratherinsensitive to the values of T star .We tested the DUSTY SEDs for T star between 5 000 and 50 000 K, andno distinguishable differences in the results were found. On the other hand,the inner radius is relatively unconstrained by current data resolution (5 ′′ corresponds to ∼
10 000 UA at 2 Kpc), as is T in . As a matter of fact, T in determines the radius at which the code starts the calculations. Several SEDsmodeling, as well as mid-infrared spectra, of massive stars rule out inner radiustemperatures consistent with dust sublimation temperatures ( ∼ / or shocks are usually invoked to producea large cavity depleted of dust and to limit the amount of short wavelengthsemission (Churchwell et al. 1990; Faison et al. 1998; Hatchell et al. 2000). ForT in we fixed a value in a similar manner as previous works in the literature,giving an inner radius larger than it would be expected from dust sublima-tion (Hatchell et al. 2000; Jørgensen et al. 2002; Crimier et al. 2009, 2010;Vehoff et al. 2010; Hirsch et al. 2012). Finally, for very young sources it isusual to assume that the envelope is composed of particles surrounded by athin layer of ice. For this reason, we adopt the opacity corresponding to adensity of 10 cm − from Ossenkopf & Henning (1994).The DUSTY model needs to be scaled to the distance and the bolometricluminosity ( L bol ) of the central source to be able to compare the models andobserved data. The bolometric luminosity is, however, an output parameter,since it is calculated by integrating the modeled SED. Consequently, the lu-minosity is re-evaluated as input parameter iteratively until minimizing thedifference between the models and observations. We adopted a distance of 2.0kpc (see § − < ′′ ). This componentis the stellar radiation attenuated by a dusty envelope, whose width is pro-portional to the stellar radius (Ivezi´c & Elitzur 1996). Given that the pixelsize at 70 µ m is 3 . ′′
2, we re-binned the model profile using a slightly largerbin of 4 ′′ . Finally, this profile is convolved with the instrumental profile at ∼ moshe/dusty/ p Y
100 – 910 10 τ T star
15 000 K − T in
300 K − µ m, taken as the average intensity profile of PACS calibrated sources, suchas: α Tau, Red Rectangle, IK Tau and the Vesta asteroid (Lutz 2012). Thefirst two sources are used to model the core and the last two the extendedwings of the profile (Aniano et al. 2011).To systematically compare the models with the observations we built agrid of 127 100 DUSTY models. The grid was constructed for 31 values of thepower index ( p ) in the density profile, ranging from 0 . . Y ) in the range 100-910 with steps of10, and 50 values of the optical depth ( τ ), from 0.1 to 5.0, in steps of 0.1.The space of the parameters explored is similar to that used by Crimier et al.(2010), restricted to a more limited range of each parameter but includingtypical values for young stellar objects. Table 4 summarizes the ranges ofthe free parameters. The best fit is obtained by means of the goodness-of-fitcriterion. We employed the weighted minimum square method to fit the SED.The weights are inversely proportional to the square of the errors assigned toeach observed data point.In the case of IRS 1, to find a model that fits both the SED and theintensity profile simultaneously, we follow the procedure used by Crimier et al.(2010). We first fixed the value of the opacity ( τ ) and by fitting the intensityprofile we obtained an initial value for the power law density index ( p ) and forthe envelope relative size ( Y ). Then we recalculated the value of τ by fittingthe SED, using the values of p and Y derived in the previous step. We adoptedthe new value of τ in the next iteration and the process was repeated. Theconvergence was achieved in a few steps. This procedure considers the factthat the intensity profile is strongly dependent on the size of the envelopeand the density distribution, whereas the SED depends mainly on the columndensity, that, in turn, depends on the opacity (Crimier et al. 2010).3.4 . Derived parameters The solid red line in Figs. 4 and 5 shows the best DUSTY models for IRS 1,4, 5, and 6. With green continuous line we show all DUSTY models comprisedwithin the error bars of the fluxes (20 in the case of IRS 1 and 200 for IRS 4, 5RAS 08589 − L τ p Y ( L ⊙ )1 1.9 × × × × r in r ext T dust M env n ( r in )(AU) (pc) (K) ( M ⊙ ) (cm − )1 112 0.14 18 86 3 × × × × TABLE 6AVERAGE MODELED PARAMETERSAverage input parametersIRS τ p Y . ± . * ± ±
534 3 . ± . ± ±
565 3 . ± . ± ± . ± . ± ± r in r ext T dust M env n ( r in )(AU) (pc) (K) ( M ⊙ ) (cm − )1 129 ± . ± .
03 16 ± ±
33 (4.00 ± × ± . ± .
01 14 ± ±
12 (1.6 ± × ±
16 0 . ± .
08 11 ± ±
14 (8.4 ± × ±
29 0 . ± .
05 15 ± ± ± × * All models that fall within the fluxes error bars have the same τ value. L , the opacity τ , the power law density index p , and the relativesize of the envelope Y . We include in this table the physical parameters ofthe envelope: the inner ( r in ) and external ( r ext ) radii, the temperature of thedust at r ext ( T dust ), and the mass ( M env ) and the density in the inner part ofthe envelope ( n ( r in )).Since the Dusty code does not provide error determinations for the modeledparameters, Table 6 lists average values and standard deviations correspond-ing to all DUSTY models, calculated for the same luminosity, that fall withinthe error bars of the fluxes (green lines in Figs. 4 and 5). Best model param-eters listed in Table 5 are not always within the standard deviations given inTable 6. However, they are not significantly different from the average values.In the case of IRS 1, p and Y are derived modeling simultaneously the SEDand the 70 µ m intensity profile. For this reason, the differences between thevalues listed in Table 5 and 6 are larger.For the SED of IRS 1, the fluxes at 3.4, 4.6, and 12 µ m are not well fittedby the model, as we have mentioned before. These fluxes are likely contam-inated by shocks and PAH emission features detected in these WISE bands(Wright et al. 2010). Moreover, Salda˜no et al. (2016) detected an outflow as-sociated with IRS 1, which likely produces a cavity in the envelope. Dust inthe walls of the cavity is heated by UV and optical radiation coming from thestellar embryo/s, re-irradiating in the IR (see, for example, Bruderer et al.2009; Zhang et al. 2013). Thus mid-IR fluxes are considered as upper limits.In addition, the observed intensity profile is well fitted by the modeled profileout to 20 ′′ , but not beyond.The density profile for the envelope of IRS 1 derived with the DUSTY codeis constant, and consistent with the supposition introduced by Ghosh et al.(2000) in their modeling of the SED of this source. In addition, the otherparameters determined by these authors ( M = 55 M ⊙ , R max = 0 . L = 2.4 × L ⊙ ) are similar or the same order as the parameters obtainedthrough DUSTY. However, we derived an opacity for the envelope four timeslarger than the value obtained by Ghosh et al. (2000). The difference is prob-ably due to the fact that these authors used silicate and graphite dust grainswithout a layer of ice on the surface, producing less attenuation to radiation.The inclusion of a layer of ices gives a more realistic model (Ossenkopf & Henning1994). 3.5 . Virial Mass Assuming that: 1) the parent cloud is similar to that one used in theDUSTY code; 2) the source is in hydrostatic equilibrium; and 3) the volumedensity of the gas decays as a power law ( ρ ∝ r − p ), we estimate the envelopemass using the following expression for the virial theorem: M vir = 3 (cid:18) − p − p (cid:19) r ext G k σ , (1)RAS 08589 − M ( M ⊙ ) M ( M ⊙ ) M ( M ⊙ ) M ( M ⊙ )Virial Mass 67 8 34 14Envelope Mass 86 42 38 11where G is the gravitational constant and σ is the one-dimensional velocitydispersion averaged over the entire system. The k factor is approximatelyequal to 1 for Y ≫ p < p and r ext , and ∆ v = p ln (2) σ , where ∆ v is the full-width at half-maximum (FWHM) of the C O line detected towards IRS 1(Salda˜no et al. 2016), we estimate σ = 0.65 ± − and a virial massfor IRS 1 listed in Table 7. For the other three starless cores (IRS 4, 5 and6), we adopt an average velocity dispersion ∼ ± − , averagingthe values determined by S´anchez-Monge et al. (2013) and Battersby et al.(2014) for multiple starless cores observed in NH molecular transitions. Table7 lists the calculated values. Sources 1, 5 and 6 have virial masses roughlyin agreement with their DUSTY mass envelope, confirming that they aregravitationally bound. On the other hand, IRS 4 has a virial mass smallerthan the envelope mass, indicating that this source may be collapsing.4. PROBABLE SCENARIO OF THE ORIGIN OF THE ARC-LIKESTRUCTURESIn Figure 6 we show large scale WISE and Herschel images of the region ofthe IRAS 08589 − µ m. We superimpose contour lines of the emissions at12 µ m from WISE (green lines) and at 1.2 mm (blue lines) from Beltr´an et al.(2006). The seven sources identified and labeled in Figure 1 are also markedonto the 4.6 µ m image.At 12 µ m we can see three arc-like structures pointing to the west. Themore extended structure borders the western rim of IRS 1, and the two smallerones coincide with IRS 5 and 6. Salda˜no et al. (2016) observed CO(3 − CO(3 − O(3 − + (3 − −
2) molecular lines in aregion of 150 ′′ × ′′ , centered on the IRAS source with the APEX telescope.These data covered IRS 1 and 3 positions. The molecular emissions of COand CO to the west of IRS 1 are very weak. In addition, the integratedemission of CO shows a very steep intensity gradient to the west, suggestingthat the material linked to IRS 1 is being compressed. Finally, in Figure 6we can see that the emission at 12 µ m borders the western rim of the colddust emission, as it is better seen in the 160 and 500 µ m images. The WISEfilter at 12 µ m (with a bandwidth ∼ µ m) includes strong spectral lines of8 SALDA ˜NO ET AL. Fig. 6. WISE 4.6 µ m and Herschel
70, 160 and 500 µ m images of theIRAS 08589 − µ m emission. The sourceslabeled from IRS 1 to 7 correspond to those in Table 1. neutral and ionized PAH (Tielens 2008), typical of photodissociation regions(PDRs). The comparison of the emission at 12 µ m likely from ionized PAHsand at 1.2 mm from cold dust suggests the existence of hot sources located tothe west, which are creating a PDRs, bordering the molecular gas.We have used the 70 and 160 µ m Herschel images to obtain a dust temper-ature map at the highest angular resolution possible in a relatively small area.In addition, these bands are good tracers of warm dust, a likely scenario forIRAS 08589 − Herschel
70 and 160 µ m maps, i.e., T c = f ( T ) − . Assuming a dust emissivity following a power law κ ν ∝ ν β ,being β the spectral index of the thermal dust emission, in the optically thinthermal dust emission regime f ( T ) has the form: f ( T ) = S S = B (70 , T ) B (160 , T ) (cid:18) (cid:19) β (2)where B (70 , T ) and B (160 , T ) are the blackbody Planck function for a temper-ature T at the wavelengths 70 µ m and 160 µ m, respectively. The pixel-to-pixeltemperature was calculated assuming a typical value β = 2.In Figure 7 we show the color-temperature map obtained using the methodexplained above. For temperatures in the range ∼ ∼ − ∼
40% of the area shown in Figure 7 hastemperatures <
20 K where, as mentioned, uncertainties may be large. How-ever, any overestimation makes the gradient less steeper. In other words, realtemperatures below 20 K would be even lower than shown in this figure andthe gradient steeper. Thus, the gradient from west to east, exists in spiteof the uncertainties in T. However ∼
60% of the area shown in Figure 7 isdominated by a relatively warmer component that prevails in the emission at70 µ m towards the west, which likely indicates the influence of RCW 38 (seebelow). Fig. 7. Dust temperature map (in color scale) derived from
Herschel emission at 70and 160 µ m. The color-temperature scale (in K) is on the right. The white crossindicates the IRAS source position. The arrows mark the positions of IRS 5 and 6,as references. About 16 . ′ − − ± ∼
10 pc away from the center of the IRAS08589 − Fig. 8. 70 µ m Herschel images of the WCR 38 region showing the locations ofdifferent HII regions. References are included at the bottom of the figures. The crossshows the position of IRAS 08589 − µ m contours aresuperimposed. Kaneda et al. (2013) carried out a large-scale mapping of RCW 38 in the[CII] and PAH emissions. They found that [CII] (158 µ m) emission extendsabout 10 ′ (about 5 pc) away from the center of RCW 38, in particular inthe north and east directions. Sources IRS 1, 2, 7 and 6 lie about 2.5 ′ orabout 1.5 pc away from the most external [CII] contour. It is noteworthy that[CII] is considered one of the main cooling agents in low-density photodisso-ciation regions with PAHs emissions (see, for example, Hollenbach & Tielens1999; Goicoechea et al. 2003) and, therefore, a good tracer of them. In addi-tion, [CII] emission strongly correlates with PAHs emissions. Consequently,it is likely that massive stars in RCW 38 are photodissociating the molecu-lar gas producing the PDR and heating the interstellar dust to the west ofIRAS 08589 − µ m and Herschel
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