Massive star formation in Wolf-Rayet galaxies: II. Optical spectroscopy results
aa r X i v : . [ a s t r o - ph . C O ] O c t Astronomy & Astrophysics manuscript no. MSFinWRG˙II˙main˙ACCEPTED˙26sep09 c (cid:13)
ESO 2018September 1, 2018
Massive star formation in Wolf-Rayet galaxies ⋆ II. Optical spectroscopy results ´Angel R. L´opez-S´anchez , and C´esar Esteban , CSIRO / Australia Telescope National Facility, PO-BOX 76, Epping, NSW 1710, Australia Instituto de Astrof´ısica de Canarias, C/ V´ıa L´actea S/N, E-38200, La Laguna, Tenerife, Spain Departamento de Astrof´ısica de la Universidad de La Laguna, E-38071, La Laguna, Tenerife, SpainReceived March 12, 2009; Accepted September 26, 2009
ABSTRACT
Aims.
We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that showthe presence of a substantial population of very young massive stars, most of them classified as Wolf-Rayet (WR)galaxies. In this paper, the second of the series, we present the results of the analysis of long-slit intermediate-resolutionspectroscopy of star-formation bursts for 16 galaxies of our sample.
Methods.
We study the spatial localization of the WR stars in each galaxy. We analyze the excitation mechanism andderive the reddening coefficient, physical conditions and chemical abundances of the ionized gas. We study the kinematicsof the ionized gas to check the rotation/turbulence pattern of each system. When possible, tentative estimates of theKeplerian mass of the galaxies have been calculated.
Results.
Aperture effects and the exact positioning of the slit onto the WR-rich bursts seem to play a fundamental rolein their detection. We checked that the ages of the last star-forming burst estimated using optical spectra agree withthose derived from H α imagery. Our analysis has revealed that a substantial fraction of the galaxies show evidences ofperturbed kinematics. With respect to the results found in individual galaxies, we remark the detection of objects withdifferent metallicity and decoupled kinematics in Haro 15 and Mkn 1199, the finding of evidences of tidal streams inIRAS 08208+2816, Tol 9 and perhaps in SBS 1319+579, and the development of a merging process in SBS 0926+606 Aand in Tol 1457-262. Conclusions.
All these results –in combination with those obtained in Paper I– reinforce the hypothesis that interactionswith or between dwarf objects is a very important mechanism in the triggering of massive star formation in starburstgalaxies, specially in dwarf ones. It must be highlighted that only deep and very detailed observations –as these presentedin this paper– can provide clear evidences that these subtle interaction processes are taking place.
Key words. galaxies: starburst — galaxies: interactions — galaxies: dwarf — galaxies: abundances — galaxies: kinematicsand dynamics— stars: Wolf-Rayet
1. Introduction
Wolf-Rayet (WR) stars are the evolved descendants of themost massive, very hot and very luminous (10 to 10 L ⊙ )O stars. In the so-called Conti (1976) and Maeder (1990,1991) scenarios, WR stars are interpreted as central He-burning objects that have lost the main part of their H-richenvelope via strong winds. Hence, their surface chemicalcomposition is dominated by He rather than H, along withelements produced by the nuclear nucleosynthesis. WN andWC stars show the products of the CNO cycle (H-burning)and the triple- α (He-burning), respectively. The most mas-sive O stars ( M ≥ M ⊙ for Z ⊙ ) became WR stars be-tween 2 and 5 Myr since their birth, spending only some few Send offprint requests to : ´Angel R. L´opez-S´anchez, e-mail:
[email protected] ⋆ Based on observations made with NOT (Nordic OpticalTelescope), INT (Isaac Newton Telescope) and WHT (WilliamHerschel Telescope) operated on the island of La Palma jointlyby Denmark, Finland, Iceland, Norway and Sweden (NOT)or the Isaac Newton Group (INT, WHT) in the SpanishObservatorio del Roque de Los Muchachos of the Instituto deAstrof´ısica de Canarias. hundreds of thousands of years ( t W R ≤ × yr) in thisphase (Meynet & Maeder 2005). A review of the physicalproperties of WR stars was recently presented by Crowther(2007).The broad emission features that characterized the spec-tra of WR stars are often observed in extragalactic H ii re-gions. Actually, the so-called Wolf-Rayet galaxies make upa very inhomogeneous class of star-forming objects: giantH ii regions in spiral arms, irregular galaxies, blue compactdwarf galaxies (BCDGs), luminous merging IRAS galaxies,active galactic nuclei (AGNs), Seyfert 2 and low-ionizationnuclear emission-line regions (LINERs). All objects have incommon ongoing or recent star formation which has pro-duced stars massive enough to evolve to the WR stage(Shaerer et al. 1999). There are two important broad fea-tures that reveal the presence of WR stars in the integratedspectra of an extragalactic H ii region:1. A blend of He ii λ iii /C iv λ iii λ blueWR Bump . This feature is mainly due to the pres-ence of WN stars. The broad, stellar, He ii λ L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results its main feature. Although rarely strong, the narrow,nebular He ii λ iii λ iv λ red or yellow WR Bump . C iv λ UV , optical and NIR colorsat a given age (i.e., Leitherer et al. 1999), nucleosyntheticyields (e.g., Woosley & Weaver 1995), the dust-to-gas ra-tio (e.g., Hirashita et al 2001), the shape of the interstellarextinction curve (e.g., Piovan et al. 2006), or even the WRproperties (Crowther 2007).The most robust method to derive the metallicity instar-forming and starburst galaxies is via the estimationof metal abundances and abundance ratios, in concretethrough the determination of the gas-phase oxygen and ni-trogen abundances and the nitrogen-to-oxygen ratio. Therelationships between current metallicity and other galaxyparameters, such as colors, luminosity, neutral gas content,star formation rate, extinction or total mass, constraintgalaxy evolution models and give clues about the currentstage of a galaxy. For example, is still debated if massivestar formation result in the instantaneous enrichment ofthe interstellar medium of a dwarf galaxy, or if the bulkof the newly synthesized heavy elements must cool beforebecoming part of the ISM that eventually will form thenext generation of stars. Accurate oxygen abundance mea-surements of several
H ii regions within a dwarf galaxy willincrease the understanding of its chemical enrichment andmixing of enriched material. The analysis of the kinemat-ics of the ionized gas will also help to understand the dy-namic stage of galaxies and reveal recent interaction fea-tures. Furthermore, detailed analysis of starburst galaxiesin the nearby Universe are fundamental to interpret the ob-servations of high-z star forming galaxies, such as LymanBreak Galaxies (Erb et al. 2003), as well as quantify theimportance of interactions in the triggering of the star-formation bursts, that seem to be very common at higher redshifts (i.e., Kauffmann & White 1993; Springer et al.2005).We have performed a detailed photometric and spec-troscopic analysis of a sample of 20 WR galaxies. Ourmain aim is the study of the formation of massive starsin starburst galaxies, their gas-phase metal abundance andits relationships with other galaxy properties, and therole that the interactions with or between dwarf galaxiesand/or low surface brightness objects have in the trigger-ing mechanism of the star-formation bursts. In Paper I(L´opez-S´anchez & Esteban 2008) we exposed the motiva-tion of this work, compiled the list of the analyzed WRgalaxies (Table 1 of Paper I) and presented the results of op-tical/
NIR broad-band and H α photometry. In this secondpaper we present the results of the analysis of intermediate-resolution long slit spectroscopy of 16 objects of our sampleof WR galaxies –the results for the other 4 objects havebeen published separately. In many cases, two or more slitpositions have been used in order to analyze the most inter-esting zones, knots or morphological structures belongingto each galaxy or even surrounding objects. In particular,these observations have the following aims:1. Study the content and spatial location of the WR starsin each galaxy. We examine the spectra for the presenceof the He ii λ tidal dwarf galaxy (TDG)or pre-existing dwarf galaxy nature of nearby diffuse ob-jects surrounding the main galaxy.4. Determine the radial velocities of different star-formation bursts, galaxies in the same system and/orobjects in possible interaction. The distance to the maingalaxy is also calculated.5. Study the velocity field via the analysis of position-velocity diagrams in order to understand the kinemat-ics of the ionized gas associated to different members inthe system in order to know their evolution (rotation,interactions features, fusion evidences, movements asso-ciated to superwinds...). The Keplerian mass has beenestimated in objects showing solid-body rotation.6. Obtain independent estimations of the age of the laststar-forming burst via the comparison with stellar pop-ulation synthesis models.7. Study the stellar population underlying the bursts us-ing the analysis of absorption lines (i.e. Ca ii H,K, Mg i λλ i λλ ii triplet).8. Finally, the spectral energy distribution ( Sed ) has beenanalyzed in some cases in order to constrain the prop-erties of the underlying stellar population.This paper mainly presents the analysis of the ionizedgas within our WR galaxy sample. In § ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 3 observations, some details of the data reduction processes,and describe some useful relations. In § T e , n e , reddening coefficient, excitationmechanism), the chemical abundances and the kinematicsof the ionized gas for each galaxy. Finally, the most impor-tant results derived from our spectroscopic study, includinga comparison with the ages derived from the H α photom-etry and an estimation of the age of the underlying stellarcomponent using the Sed , are summarized in §
4. A detailedanalysis of the O and WR populations and the comparisonwith theoretical models is presented in Paper III. The globalanalysis of our optical/NIR data will be shown in PaperIV. The final paper of the series (Paper V) will compilethe properties derived using data from other wavelengths(UV, FIR, radio and X-ray) and complete the global analy-sis combining all available multiwavelength data of our WRgalaxy sample. It is, so far, the most complete and exhaus-tive data set of this kind of galaxies, involving multiwave-length results and analyzed following the same procedures.
2. Observations and data reduction and analysis
We obtained intermediate-resolution long slit spectroscopyfor all our sample WR galaxies except for NGC 5253, forwhich high-resolution echelle spectroscopy was taken (seeL´opez-S´anchez et al. 2007 for details). We used three tele-scopes to carry out these observations: 2.5m
Isaac NewtonTelescope (INT), 2.56m
Nordical Optical Telescope (NOT),and 4.2m
William Herschel Telescope (WHT), all locatedat Roque de los Muchachos Observatory (ORM, La Palma,Spain). The details of these observations are the following:1.
Observations at the 2.5m INT . We used the IDS(
Intermediate Dispersion Spectrograph ) instrument at-tached at the Cassegrain focus in December 1999. AEEV CCD 2148 × µ m, was used, that corresponds to an spatial resolu-tion of 0.40 ′′ pix − . The slit was 2.8 ′ long and 1 ′′ wide.We used the R400V grating, that has a dispersion of104.5 ˚A mm − (1.40 ˚A pix − ) and an effective spectralresolution of 3.5 ˚A. The spectra cover the wavelengthrange from 3200 to 7700 ˚A. The absolute flux calibra-tion was achieved by observations of the standard starsFeige 56, Hiltner 600 and Feige 110 (Massey et al. 1988).2. Observations at the 4.2m WHT . We completedtwo observation runs in this telescope on December2000 and December 2002. In both cases, the double-armISIS (
Intermediate dispersion Spectrograph and ImagingSystem ) instrument located at the Cassegrain focus ofthe telescope was used. The dichroic used to separatethe blue and red beams was set at 5400 ˚A. The slit was3.7 ′ long and 1 ′′ wide. We used different configurationsin each observing run:(a) December 2000 : – Blue arm : an EEV CCD with a 4096 × µ m size was used. The spa-tial resolution was 0.20 ′′ pix − . The gratingwas R600B, giving a dispersion of 33 ˚A mm − (0.45 ˚A pix − ) and an effective spectral resolu-tion of 1.8 ˚A. The observed spectral range was3600 – 5200 ˚A. – Red arm : we used a TEX CCD with a con-figuration of 1024 × µ m pixel size, having a spatial resolution of 0.36 ′′ pix − .The grating R316R, that has a dispersion of66 ˚A mm − (0.93 ˚A pix − ) and an effective spec-tral resolution of 2.6 ˚A, was used, covering thespectral range 5400 – 6800 ˚A.(b) December 2002 : – Blue arm : We used the same CCD that previ-ously indicated but the R1200B grating, thatgives a dispersion of 17 ˚A mm − (0.23 ˚A pix − )and an effective spectral resolution of 0.86 ˚A.The spectral range was 4450 – 5480 ˚A. – Red arm : A Marconi CCD with 4700 × µ m pixel size was used. Thespatial resolution was 0.20 ′′ pix − , hence iden-tical to that provided in the blue arm. We usedthe R316R grating covering the spectral range5700 – 8600 ˚A.The absolute flux calibration was achieved by ob-servations of the Massey et al. (1988) standard stars G191B2B and Feige 34 (December 2000) and Feige15, Feige 110, Hiltner 600 and Hz44 (December 2002).3.
Observations at the 2.56m NOT . We completedthree observation run at this telescope, always usingthe ALFOSC (
Andaluc´ıa Faint Object Spectrograph andCamera ) instrument and a Loral/Lesser CCD detector(2048 × µ m andspatial resolution of 0.19 ′′ pixel − . The slit was 6.4 ′ longand 1 ′′ wide. We used several configurations:(a)
20 March 2004 . We used grism − (1.5 ˚A pix − ) and a spec-tral resolution of 7.5 ˚A, covering the spectral range3200 – 6800 ˚A.(b) . We used twodifferent grisms to obtain the blue and the red rangesof the optical spectrum. Grism − (1.4 ˚A pix − ) and a spectralresolution of 7.0 ˚A, was used to cover the spectralrange 3300 – 6100 ˚A. This grism has a low efficiencyfor λ ≤ − (1.3 ˚A pix − ), a spectral resolution of6.5 ˚A and covers the spectral range 5800 – 8300 ˚A.The spectrophotometric standard star Feige 56(Massey et al. 1988) was used for flux calibrating allthe spectra obtained with this telescope.In all observations, three or four exposures for each slitposition were taken to get a good S/N ratio and to re-move cosmic rays. Table 1 compiles all the intermediate-resolution long-slit spectroscopy observations performed forthe 16 WR galaxies included in this paper. All the data reduction were completed at the IAC. IRAF software was used to reduce the CCD frames (bias cor-rection, flat-fielding, cosmic-ray rejection, wavelength andflux calibration, sky subtraction) and extract the one-dimensional spectra. The correction for atmospheric extinc-tion was performed using an average curve for the contin-uous atmospheric extinction at Roque de los Muchachos IRAF is distributed by NOAO which is operated by AURAInc., under cooperative agreement with NSF. L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
Table 1.
Journal of the intermediate-resolution long-slit spectroscopy observations, carried out using the 2.5m INT, 2.56m NOTand 4.2m WHT telescopes and with the instrumentation explained in text. In last column, sec z represents the average airmass inwhich the observations were made. Galaxy Tel. Date Exp. Time Spatial R. Grism P.A. Spectral R. ∆ λ sec z [s] [ ′′ /pix] [ ◦ ] [˚A] [˚A]Haro 15 INT 99/12/27 3 × × × × ×
700 0.20 R1200 90 0.86 4300-5100 1.90WHT 02/12/27 3 ×
700 0.20 R136R 90 2.6 5700-7800 1.90INT 99/12/29 3 × × × × × × × × × × ×
600 0.20 R1200 27 0.86 4300-5100 1.18WHT 02/12/27 3 ×
600 0.20 R136R 27 2.6 5700-7800 1.18SBS 0948+532 WHT 00/12/31 3 × × × × × ×
600 0.20 R1200 39 0.86 4300-5100 1.48WHT 02/12/27 3 ×
600 0.20 R136R 39 2.6 5700-7800 1.48SBS 1415+437 WHT 02/12/27 3 ×
600 0.20 R1200 20 0.86 4300-5100 1.54WHT 02/12/27 3 ×
600 0.20 R136R 20 2.6 5300-6650 1.54III Zw 107 INT 99/12/28 3 × × × ×
900 0.19 g8 109 6.5 5800-8300 1.92Tol 1457-262a NOT 05/04/04 3 ×
900 0.19 g14 155 7.0 3300-6100 1.92NOT 05/04/04 3 ×
900 0.19 g8 155 6.5 5800-8300 1.75Arp 252 NOT 05/04/04 3 ×
900 0.19 g14 342 7.0 3300-6100 1.67NOT 05/04/04 3 ×
900 0.19 g8 342 6.5 5800-8300 1.56
Observatory. For each two-dimensional spectra severalapertures were defined along the spatial direction to ex-tract the final one-dimensional spectra of each galaxy oremission knot. The apertures were centered at the brightestpoint of each aperture and the width was fixed to obtain agood signal-to-noise spectrum. In case of having the opticalspectrum separated in two different wavelength intervals,identical apertures in both spectral ranges were used.
IRAF software was also used to analyze the one-dimensional spectra. Line intensities and equivalent widthswere measured by integrating all the flux in the line betweentwo given limits and over a local continuum estimated byeye. In the cases of line blending, a multiple Gaussian profilefit procedure was applied to obtain the line flux of each in-dividual line. We used the standard assumption I (H β )=100to compute the line intensity ratios. The identification andadopted laboratory wavelength of the lines, as well as theirerrors, were obtained following Garc´ıa-Rojas et al. (2004);Esteban et al. (2004). We computed the distance to the galaxies using the bright-est emission lines (H α and [O iii ] λ H =75 km s − Mpc − , q =0.5, and corrected for Galactic Standard of Rest. Thedistance we derived for each galaxy is listed in Table 1 inPaper 1. All values agree well within the errors with thedistances quoted by the NED, except for Tol 9. For thisgalaxy, we measure a radial velocity of v r = 3441 km s − ,while previous observations suggested v r = 3190 km s − (Lauberts & Valentijn 1989). We consider that our value ismore appropriate because the maximum of the H i emis-sion detected in Tol 9 (L´opez-S´anchez et al. 2010) showsthe same radial velocity than that our optical spectrumprovides.
The reddening coefficient, c (H β ), was derived from theBalmer decrement. However, in extragalactic objects thefluxes of nebular Balmer lines may be affected by absorp-tions produced by the underlying stellar population (mainly ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 5 B and A stars). We performed an iterative procedure to de-rive simultaneously the reddening coefficient and the equiv-alent widths of the absorption in the hydrogen lines, W abs ,to correct the observed line intensities for both effects. Weassumed that W abs is the same for all the Balmer lines andused the relation given by Mazzarrella & Boronson (1993)to the absorption correction, c ( Hβ ) = 1 f ( λ ) log " I ( λ ) I ( Hβ ) × (cid:16) W abs W Hβ (cid:17) F ( λ ) F ( Hβ ) × (cid:16) W abs W λ (cid:17) , (1)for each detected hydrogen Balmer line. In this equation, F ( λ ) and I ( λ ) are the observed and the theoretical fluxes(unaffected by reddening or absorption); W abs , W λ and W Hβ are the equivalent widths of the underlying stellar ab-sorption, the considered Balmer line and H β , respectively,and f ( λ ) is the reddening curve normalized to H β usingthe Cardelli, Clayton & Mathis (1989) law. We always con-sidered the theoretical ratios between these pairs of H i Balmer lines expected for case B recombination given byStorey & Hummer (1995) assuming the electron tempera-ture and density computed independently for each region.Using three different Balmer lines (i.e., H α , H β and H γ ) anunique value for c (H β ) and W abs is computed. However, inthe case of using four or more Balmer lines, several solu-tions are derived, so we considered the values that providethe best match between the corrected and the theoreticalBalmer line ratios as representative of the region. An ex-ample of this method using 5 Balmer line ratios is shownin Figure 14 when analyzing UM 420.Tables A.1, A.3, A.5, A.7, A.9, A.11, A.13, A.15 andA.17 show the dereddened line intensity ratios and theirassociated errors for all the regions and galaxies, as wellas the adopted f ( λ ) for each emission line. In these tables,we also include other important quantities as: the size ofthe extracted aperture, its relative distance to the mainregion of the galaxy, the observed H β flux (uncorrected forextinction), the adopted values of c (H β ) and W abs and theequivalent widths of H α , H β , H γ and [O iii ] λ We studied the physical conditions and chemical abun-dances of the ionized gas from the 1-D spectra ofeach galaxy or knot. We used a two-zone approxima-tion to define the temperature structure of the nebu-lae, assuming the electron temperature, T e , provided bythe [O iii ] ion, T e (O iii ), as the representative tempera-ture for high ionization potential ions and T e (N ii ) or T e (O ii ) for the low ionization potential ones. T e (O iii )is ob-tained from the [O iii ] ( λ λ λ T e (N ii )from [N ii ] ( λ λ λ T e (N ii ) from [O ii ]( λ λ λ λ T e s are calculated makinguse of the five-level program for the analysis of emission-line nebulae that is included in IRAF NEBULAR task(Shaw & Dufour 1995). Notice that we used an updatedatomic dataset for O + , S + , and S ++ for NEBULAR. Thereferences are indicated in Table 4 of Garc´ıa-Rojas et al.(2004).When one of the high/low ionization electron temper-atures can not be computed, we used the linear relation between T e (O iii ) and T e (O ii ) provided by Garnett (1992), T e ( O ii ) = T e ( N ii ) = 0 . × T e ( O iii ) + 3000 , (2)to estimate the unknown electron temperature. In the casethat no direct estimate of the electron temperature can beobtained, we considered the T e (O iii ) and T e (O ii ) pairs thatreproduce the total oxygen abundance obtained by apply-ing the Pilyugin (2001a,b) empirical method (see below),also assuming the Garnett (1992) relation, that is the sameequation that Pilyugin uses in his empirical calibrations.The electron density of the ionized gas, n e , was usu-ally computed via the [S ii ] λλ ii ] λλ n e <
100 cm − are below the low-densitylimit and hence we considered n e =100 cm − in those cases.Veilleux & Osterbrock (1987) proposed diagnostic di-agrams plotting two different excitation line ratios, suchas [O iii ]/H β versus [S ii ]/H α , for classifying the excita-tion mechanism of ionized nebulae. H ii regions (or
H ii or starburst galaxies) lie into a narrow band within thesediagrams, but when the gas is ionized by shocks, accretiondisks or cooling flows (in the case of AGNs or LINERs)its position is away from the locus of
H ii regions. Weused the analytic relations given by Dopita et al. (2000)and Kewley et al. (2001) between different line ratios tocheck the nature of the excitation mechanism of the ion-ized gas within the bursts. Figure 5 shows an example ofthese diagrams applied to the regions analyzed in the galaxyMkn 1199.
Once the electron density and temperature are estimated,the ionic abundances can be derived for each region. All theionic abundances except He + and Fe ++ were calculated us-ing the IRAF NEBULAR task (Shaw & Dufour 1995) fromthe intensity of collisionally excited lines. We assumed atwo-zone scheme for deriving the ionic abundances, adopt-ing T e (O iii ) for the high ionization potential ions O ++ ,Ne ++ , S ++ , Ar ++ , Ar and Cl ++ ; and T e (N ii ) or T e (O ii )for the low ionization potential ions O + , N + , S + and Fe ++ .The He + /H + ratio was computed from the He i linesintensities and using the predicted line emissivities calcu-lated by Smith, Shara & Moffat (1996) for the T e (O iii )and n e assumed for each region. We also correctedfor collisional contribution following the calculations byBenjamin, Skillman & Smits (2002). Self-absorption effectswere not considered.Fe ++ abundances were derived via the [Fe iii ] λ + /H + + O ++ /H + to determine the total oxygen abundance. We detect theHe ii λ ++ to the total amount of helium is neg-ligible, implying that O has also a very low abundancein the nebula, thus we did not consider its contributionto the total O/H ratio. For deriving the nitrogen abun-dance we assumed the standard ionization correction factor(icf) by Peimbert & Costero (1969): N/O = N + /O + , whichis a reasonably good approximation for the excitation de-gree of the observed galaxies. We used the icf provided by L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
Peimbert & Costero (1969) to derive the total neon abun-dance. We computed the total sulphur abundance whenboth S + /H + and S ++ /H + ratios are available using the icfgiven by the photoionization models by Stasi´nska (1978).The total argon abundance was calculated considering theicfs proposed by Izotov, Thuan, & Lipovetski (1994). Thetotal iron abundances were obtained from the Fe ++ /H + ratio and the icf given by Rodr´ıguez & Rubin (2005).As we said before, when direct estimations of the elec-tron temperature were not available, we resorted to em-pirical calibrations. Pilyugin (2001a,b) performed a de-tailed analysis of observational data combined with pho-toionization models to obtain the oxygen abundance fromthe relative intensities of strong optical lines. Pilyugin(2001a) gives the empirical calibration between the R and P (an indicator of the hardness of the ionizing ra-diation) parameters and the oxygen abundance in mod-erately high-metallicity H ii regions, 12+log(O/H) ≥ T e was possi-ble. Sometimes, we estimated the total oxygen abundancemaking use of the Denicol´o, Terlevich & Terlevich (2002)or Pettini & Pagel (2004) empirical calibrations, which in-volve the [N ii ] λ α ratio. The bidimensional spectra have been used to perform aposition-velocity diagram via the analysis of the brightestemission line profiles (H α and/or [O iii ] λ M Kep ) of the galaxies assuming thatthe kinematics are representative of circular rotation. Weusually considered the half of the maximum velocity differ-ence, ∆ v , the half of the spatial separation correspondingto these maxima, r , and applied the equation: M kep [ M ⊙ ] ∼ × r [pc] (cid:18) ∆ v [km s − ]sin i (cid:19) , (3)assuming circular orbits and Keplerian dynamics(L´opez-S´anchez et al. 2004b). We remark that the re-sult of this equation is not the total dynamical mass of thegalaxy but only the total mass contained within a circle ofradius r . The inclination angle, i , is defined as that foundbetween the plane of the sky and the plane of the galaxy(hence, i =90 ◦ in an edge-on galaxy and i =0 ◦ in a face-ongalaxy). We usually estimated this angle assuming thatthe elliptical shape of the galaxy is just a consequence ofits orientation.When 21-cm H i data were available in the literature, wecomputed both the neutral gas mass, M H I , and the dynam-ical mass, M Dyn , using the typical relations (i.e., Dahlem etal. 2005). We estimated the rotation velocity of the neutralgas considering ∆ v = W H I i and adopted the maximum ra-dius observed in our deep optical images. Therefore, as theextension of the neutral gas is usually larger than the exten-sion of the stellar component, our estimations of M Dyn maybe underestimated. The gas depletion timescale defined bySkillman et al. (2003) was computed using M H I and the total star-formation rate (SFR) derived for each galaxy inPaper I. When FIR data are available, we estimated themass of the warm dust, M dust , using the equations given byHuchtmeier et al. (1995). Although M Kep , M Dyn and M dust should be considered only tentative values, their compar-ison gives important clues about the galaxy type, its dy-namic and the fate of the neutral gas.
3. Results
The spectroscopic analysis of NGC 1741 (member AC inthe galaxy group HCG 31) was presented in detail inL´opez-S´anchez et al. (2004a). Our spectra show an evidentbroad blue WR bump and the He ii λ ii λ Mkn 1087 is a luminous blue compact galaxy and the maingalaxy in a group in interaction. Although some authorsdid observe WR features in this galaxy Kunth & Joubert(1985); Vaceli et al. (1997), others did not (Vacca & Conti1992), and therefore it was classified as suspected
WRgalaxy by Schaerer, Contini & Pindao (1999). Our analysisof Mkn 1087 was presented in L´opez-S´anchez et al. (2004b);we did not detected any WR feature (Figure 36) in any im-portant star-forming region in or surrounding Mkn 1087.
Haro 15 is a blue compact galaxy well studied in all frequen-cies, including optical spectroscopy (Hunter & Gallagher1985; Mazzarrella, Bothum & Boronson 1991; Kong et al.2002; Shi et al. 2005). Schaerer et al. (1999) listed Haro 15as a WR galaxy because of the detection of the He ii λ ii λ H ii region D at the WNW (both observed with the slit withPA 117 ◦ ) and the knot B (at the NE, observed with theslit with PA 41 ◦ ). Figure 1 shows the spectra of the threebrightest objects, whereas Table A.1 compiles the emissionline ratios and other properties of the spectra of each region.The spectrum of the center of Haro 15 shows both nebu-lar emission lines and stellar absorptions; these absorptionsare observed mainly in the H i
Balmer lines. However, thespectrum of region A is entirely dominated by nebular emis-sion lines, where we detect [O iii ] λ ii λ ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 7 [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ Haro 15 C O b s e r v ed f l u x [ - e r g s - c m - Å - ] He II [O I] [SIII][Ne III] [S II][N II]He I He I H α [O II] H ε H δ [O III] [O III]H β H γ Haro 15 A - PA 117º [S II]H α [O II] [O III]H β Haro 15 B - PA 41º
Wavelength [ Å ]
Fig. 1.
IDS INT spectra for regions C, A and B of Haro 15.Fluxes are not corrected for reddening. The most importantemission lines have been labeled. See Figure 3 in Paper I forthe identification of the regions.
150 175 200 225 250 275 3000200040006000800010000120001400016000
150 175 200 225 250 275 3000100200300400500600
Center A
PA 117º R e l a t i v e f l u x distance [ pixels ] [O III] 5007 Center A
PA 117º distance [ pixels ] [O II] 3727 [Ne III][O II] H δ [O III]H β H γ Haro 15 C O b s e r v ed f l u x [ - e r g s - c m - Å - ] [Ne III][O II] H ε H δ [O III]H β H γ Haro 15 A [O II] [O III]H β Haro 15 B
Wavelength [ Å ]
Fig. 2.
Lower panels: zoom of the spectra of regions C (cen-ter), A and B of Haro 15 (bottom). Fluxes are not correctedfor reddening. Upper panels: spatial distribution of the relativeflux of [O ii ] λ iii ] λ ◦ . A direct estimation of T e (O iii ) was computed in knot Abecause of the detection of [O iii ] λ iii ] λ T e [O iii ] ∼ ii ] λ n e .We assumed a value of 100 cm − for all regions. Comparingthe line ratios with diagnostic diagrams, we found that allknots can be classified as typical H ii regions.
Table A.2 compiles all the chemical abundances computedfor each region of Haro 15. We found a significant differencebetween the oxygen abundance at the center of Haro 15, -15-10-50510152025-60 -50 -40 -30 -20 -10 0 10 20 30 -150 -100 -50 0 50 100 150-20-15-10-5051015202530 D H α A C relative velocity [ km s -1 ] d i s t an c e [ a r cs e c ] Haro 15 - PA 117º H α C B d i s t an c e [ a r cs e c ] relative velocity [ km s -1 ] Haro 15 - PA 41º [O III] 5007 [O III] 5007
Fig. 3.
Position-velocity diagrams for the two slit positionsobserved in Haro 15. Both the H α and the [O iii ] λ ii ] λ iii ] λ Figure 3 shows the position-velocity diagrams for the twoslit positions observed in Haro 15. We analyzed both theH α and the [O iii ] λ ◦ shows an apparentrotation pattern, although some divergences are found atthe SW. Knot B is clearly decoupled from the rotation ofthe disk, suggesting that it is an external object. The inter-action between knot B and Haro 15 could be the responsibleof the distortion observed at the SW of the diagram of PA41 ◦ because the object and this zone of the galaxy disk showsimilar radial velocities. Another possibility is that knot Bis a TDG but, in this case, the material that formed Bshould come from the external parts of the disk of Haro 15because its chemical abundances are more similar to thoseof knot D than to those of the central region. L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results [O II][Ar III]WR [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ Mkn 1199 C O b s e r v ed f l u x [ - e r g s - c m - Å - ] Wavelength [ Å ] [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ Mkn 1199 NE - PA 36º
Fig. 4.
IDS INT spectra for the center of Mkn 1199 and thedwarf companion galaxy at its NE. Fluxes are not corrected forreddening. The most important emission lines have been labeled.See Figure 5 in Paper I for the identification of the regions.
The diagram with PA 117 ◦ shows a clear sinusoidal pat-tern with differences of around 40 km s − . This feature iscommon in processes involving galaxy interaction or merg-ing. Although it is not completely clear, region A seems tobe kinetically coupled with the rotation of the galaxy. Thisfact and the different chemical composition indicate thatknot A is probably an external object suffering a mergingprocess with Haro 15.Assuming that the position-velocity diagram of PA 41 ◦ from − ∼
80 km s − within a distance of ∼
13 arcsec (5.46 kpc) and assum-ing i =55 ◦ from our optical images [Gordon & Gottesman(1981) found i =57 ◦ ], we derive M kep =1.21 × M ⊙ and M Kep /L B = 0 .
35. The neutral and dynamical masses ofHaro 15 are M H I =5.54 × M ⊙ and M Dyn =3.65 × M ⊙ (Gordon & Gottesman 1981). Thus, our Keplerianmass is ∼
33% of the the total mass. The M H I /L ⊙ and the M dust /L ⊙ ratios (0.16 and 0.6 × − , respec-tively) suggest that Haro 15 is a Sb or Sc spiral(Bettoni, Galleta & Garc´ıa-Murillo 2003). The gas deple-tion timescale is 2.3 Gyr, showing that the system still pos-sesses a huge amount of material available to form newstars. Izotov & Thuan (1998) observed Mkn 1199 but they didnot detect the [O iii ] λ ii λ ◦ but region A, for which the slit with PA 53 ◦ was used.The spectra of the two brightest regions (C and NE) areshown in Figure 4, while the emission line intensities andother important properties of the spectra are compiled inTable A.3. The spectrum of the center of Mkn 1199 showsstellar absorptions in both the H i
Balmer and the He i lines, -1,5 -1,0 -0,5 0,0 0,5-1,5-1,0-0,50,00,51,01,5 -2,0 -1,5 -1,0 -0,5 0,0 0,5 1,0-1,5-1,0-0,50,00,51,01,5 H II regions Shock contribution
D00 K01 log ( [S II] 6717,6731 / H α ) AMkn 1199 C Mkn 1199 DBMkn 1199 NE
Obs. data
H II regions Shock contribution l og ( [ O III] / H β ) log ( [N II] 6584 / H α ) D00 K01
BMkn 1199 C Mkn 1199 DAMkn 1199 NE
Obs. data
Fig. 5.
Comparison of some line intensity ratios in several re-gions of Mkn 1199 with the diagnostic diagrams proposed byDopita et al. (2000) and Kewley et al. (2001). as it was previously noticed by Izotov & Thuan (1998). Theblue WR bump and a probable He ii λ The [O iii ] λ ii ] λ T e [N ii ] ∼ ii ] λλ T e (O ii ) ∼ T e ( low )=6800 K. The high ionization electron temperaturewas computed using Garnett’s relation, T e ( low )=5400 K.Except for region C, the electron densities derived from the[S ii ] λλ n e ∼
100 cm − . The values for the reddening coeficient arerelatively high for C and B, suggesting the presence of animportant amount of dust in those regions. The comparisonof some line intensity ratios with the diagnostic diagramsproposed by Dopita et al. (2000) and Kewley et al. (2001)indicates that all regions can be classified as H ii regions(see Figure 5).
Table A.4 compiles all the chemical abundances computedfor the knots analyzed in Mkn 1199. The oxygen abun-dance found at the center of the system is very high,12+log(O/H)=8.75 ± (indeed, it highest metallicity region ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 9 -20-15-10-505101520253035-60 -50 -40 -30 -20 -10 0 10 20 30 40 -60 -40 -20 0 20 40 60 80 100-20-15-10-505101520253035 NE companionB D center relative velocity (km s -1 ) d i s t an c e ( a r cs e c ) Mkn 1199 - PA 36º centerA d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) Mkn 1199 - PA 53º
Fig. 6.
Position-velocity diagrams for the two slit positions ob-served in Mkn 1199 using the H α profiles. NE is up in bothdiagrams. See Figure 5 in Paper I for the identification of theregions. found in this work). However, the oxygen abundance of thedwarf companion galaxy at the NE is almost 0.3 dex lower,12+log(O/H)=8.46. The N/O ratios derived for them arealso very different. Hence, this result reinforces the hypoth-esis that they are independent galaxies in the first stagesof a minor merging.It is worthy to notice that, although the line inten-sities observed by Izotov & Thuan (1998) in Mkn 1199are quite similar to those compiled in Table A.3,the oxygen abundance derived by these authors is12+log(O/H)=8.19 ± ii ] λ α ra-tio proposed by van Zee, Salzer & Haynes (1998) obtain12+log(O/H)=9.13. We consider that 12+log(O/H)=8.75is a more appropriate value for the oxygen abundance atthe center of Mkn 1199, that also agrees with our tentativeestimations of electron temperatures.Knots A, B and D show oxygen abundances slightlylower than C, between 8.6 and 8.7, but not as low as thatreported for the companion galaxy. Interestingly, A, B andD show a N/O ratio very similar to that observed at thecenter of the galaxy, but much higher than in the dwarfcompanion. This confirms the different chemical evolutionof the disk of Mkn 1199 and the companion galaxy. Weconsider that knots A, B and D are giant H ii regions lo-cated in the spiral arms of Mkn 1199 and that this galaxycould have a slight radial metallicity gradient along its disk.Finally, we think that the triggering mechanism of the in-tense star-formation activity found in both, Mkn 1199 andits dwarf companion, is a likely consequence of the stronginteraction they are experiencing.
Figure 6 shows the position-velocity diagrams for the twoslit positions in Mkn 1199. They were obtained extracting 3pixels bins (1.2 arcsec) across the H α profile and consider-ing the center of Mkn 1199 as reference. The diagram withPA 36 ◦ , that crosses the center of Mkn 1199 and the com-panion galaxy, may be explained by the rotation of the disk of Mkn 1199, because it shows a velocity gradient between30 km s − at SW (knot D) and −
30 km s − at NE (aroundknot B). However, a sinusoidal pattern is also seen in thebrightest region of Mkn 1199, within the central 10 ′′ . Thisfeature may be consequence of an entity kinetically decou-pled from the disk, such a bar or the bulge of the galaxy,that seems to be counter-rotating, although it may be alsoan interaction feature. On the other hand, although thereare only four points in the diagram in that zone, the dwarfgalaxy at the NE seems to be rotating. This object, that hasan elliptical shape in the optical images, may be observededge-on, and hence it would explain both its kinematics andmorphology. The diagram with PA 53 ◦ also seems to showa velocity gradient from the SW to the NE (where knot Ais located) but this gradient is broken at the central 10 ′′ .The maximum velocity difference is around 80 km s − .It is difficult to estimate the Keplerian mass ofMkn 1199 because it is seen almost face-on. However, con-sidering that the velocity difference between its center andthe external regions within a radius of 10 ′′ (2.6 kpc) isaround 30 km s − and adopting an inclination angle of 15 ◦ ,we derive M kep ∼ × M ⊙ . For the companion dwarfobject, assuming that it is edge-on ( i =90 ◦ ) and consider-ing a velocity difference of 10 km s − within a radius of2.5 ′′ (1.3 kpc), we compute M kep ∼ × M ⊙ .Using the H i data given by Davoust & Contini (2004),we compute M H I =1.22 × M ⊙ and ∆ v H I =85 km s − .Assuming a radius of 25 ′′ (6.55 kpc) and i =15 ◦ , the dy-namical mass of Mkn 1199 is M dyn =1.7 × M ⊙ . Thewarm dust mass using the FIR data is M dust =3.1 × M ⊙ . Following the classification provided by Bettoni et al.(2003), the M H I / L ⊙ =0.042 and M dust / L ⊙ =1.1 × − ra-tios are not compatible: while the the first corresponds tothe typical values for S0 galaxies, the second ratio indicatesthat Mkn 1199 should be an Sc or Sd, more similar to theactual morphological classification of Sb. Furthermore, lessthan 1% of the total mass of the system is neutral hydrogenand the gas depletion timescale is very low (0.4 Gyr). Allthese facts suggest that a substantial fraction of the neutral H i gas has been expelled to the intergalactic medium, per-haps as a consequence of the interaction with the dwarfcompanion galaxy. An
H i map obtained using a radio-interferometer would be needed to confirm such hypothesis.
Conti (1991) included Mkn 5 in his catalogue of WRgalaxies because of the detection of the nebular He ii λ ii emission (Schaerer et al. 1999).Guseva, Izotov & Thuan (2000) detected N iii λ ◦ (INT-1), 354 ◦ (WHT) and349 ◦ (INT-2). All slit positions cover region A but only twocross knot B. We analyzed the three spectra extracted forregion A independently to check the quality of the results.Figure 7 shows the spectra of the region A obtained withthe slit positions with PA 349 ◦ and PA 354 ◦ . For region Bwe only analyzed the spectrum extracted using the slit po- [O II] [Ne III] + H7 [O III] [Ar III]WR [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ Mkn 5 A - INT2 - AP 349º O b s e r v ed f l u x [ - e r g s - c m - Å - ] WR [O II][Ar III]He I[O III] [O I] [S III]
Wavelength [ Å ] [S II][N II]He I H α [O III]H β H γ Mkn 5 A - WHT - AP 354º
Fig. 7.
Spectra for the region A of Mkn 5 obtained using IDS atthe INT (PA 349 ◦ ) and ISIS at the WHT (PA 354 ◦ ). Fluxes arenot corrected for reddening. The most important emission lineshave been labeled. See Figure 7 in Paper I for the identificationof the regions. sition with PA 0 ◦ (INT-1) because of its higher signal-to-noise. Table A.5 compiles all the emission line fluxes andother characteristics of each spectrum. Our spectra confirmthe presence of a nebular He ii λ H i
Balmer lines are also detected; theseare more evident in knot B.
The three spectra obtained for region A show the [O iii ] λ T e values are in agreement within theerrors, being the average value T e [O iii ] ∼ T e [O ii ] ∼ ii ] λλ T e [O ii ] ∼ T e in knot B,but this determination is very much uncertain and per-haps overestimated due to the faintness of the spectrum.Electron densities are always below the low-density limit(100 cm − ) except for knot B (although it also has a largeerror).The values of the reddening coefficient derived for regionA are somewhat different in the different spectra. Perhaps,this apparent inconsistency is a consequence of an irregulardistribution of dust within Mkn 5, as we suggested in ouranalysis of the optical/NIR colors (see § W abs =1.5 ˚A and the c (H β ) com-puted via the H α /H β ratio. The diagnostic diagrams forknots A and B agree with the loci of typical H ii regions.
The WHT spectrum does not cover the [O ii ] λλ ii ] λλ + /H + ratio. As we see in Table A.5, the agreementbetween the emission line ratios for all the three spectraextracted for region A is very good. Table A.6 compiles -20-15-10-50510152025303540-80 -60 -40 -20 0 20 40 60 80 100 -80 -60 -40 -20 0 20 40 60 80 10020151050-5-10-15-20-25-30-35-40 -80 -60 -40 -20 0 20 40 60 80 100-20-15-10-50510152025303540 A relative velocity (km s -1 ) d i s t an c e ( a r cs e c ) Mkn 5 - INT 2 - PA 349º
B A d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) Mkn 5 - WHT - PA 174º AB d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) Mkn 5 - INT 1 - PA 0º
Fig. 8.
Position-velocity diagrams for the slit positions observedin Mkn 5 using the [O iii ] λ the chemical abundances obtained for Mkn 5; for region Aall values are quite similar and in agreement with previousresults found in the literature. Averaging all data and min-imizing errors, we derive for A the following chemical abun-dances: 12+log(O/H)=8.07 ± − ± − ± − ± − ± ± Figure 8 shows the position-velocity diagrams obtained forthe three slit positions observed in Mkn 5. For the 4.2mWHT spectrum we analyzed the H α profile considering 6pixels bins (1.2 arcsec), while we used the [O iii ] λ α profile) extracting 4 pixels bins(1.6 arcsec) from the 2.5m INT spectra. As we see, theagreement between the three diagrams is very good. Thebest diagram is that provided by the analysis of the 4.2mWHT spectrum, that shows a velocity gradient of around ∼
50 km s − . The velocity of knot B with respect to thatfound in region A is ∼
40 km s − . Although the uncertain-ties are important, we detect a slight reverse in the velocityof region A, with an amplitude of ∼
20 km s − (it is moreevident in the 4.2m WHT diagram), that seems to show asinusoidal pattern in that area.Assuming that the global velocity gradient is mainlya consequence of the rotation of the galaxy, we mayestimate the Keplerian mass of the system. We found M Kep ∼ . × M ⊙ assuming i =90 ◦ , ∆ v ∼
27 km s − and r ∼ ′′ (1.22 kpc). Using H i data (Paturel et al.2003), we derive M H I = (7.2 ± × M ⊙ and M Dyn ∼ × M ⊙ . Although both M Kep and M Dyn aresimilar, notice that they are low limits because we are as-suming that Mkn 5 is an edge-on galaxy. The mass-to-luminosity ratios are M Kep /L ⊙ =7.98, M Dyn /L ⊙ =13.7 and M H I /L ⊙ =0.27. The H i mass is quite low for a dwarf orirregular galaxy, being only 2% of the total mass. The gasdepletion timescale is ∼ ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 11 H ε [Ar III][O III]He I WR [O I] [SIII][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ IRAS 08208+2816 C - PA 345º O b s e r v ed f l u x [ - e r g s - c m - Å - ] H δ [N II][O I]He IH γ [S II]H α [O II] [O III]H β IRAS 08208+2816
Wavelength [ Å ]
Fig. 9.
IDS INT spectra for the center of IRAS 0828+2816 andknot days it has decreased) or it has been expelled to the in-tergalactic medium. Indeed, Thuan & Martin (1981) foundindications of
H i gas at slightly different radial velocities( ∼
300 km s − ) using single-dish data. An interferometric H i map will be crucial to understand the fate of the neutralgas in this blue compact dwarf galaxy.
The first spectroscopic data of IRAS 08208+2816 were ob-tained by Huang et al. (1999), who reported the detec-tion of both the nebular and broad He ii λ iv λ ◦ crosses a bright star and the center of IRAS 0828+2816, theslit position with PA 355 ◦ covers the center and knots ◦ crosses ◦ . This spectrum and that obtained for knot H i
Balmer lines, we detect a slight decreaseof the continuum level on the blue range of the spectra ofthe faintest objects . This fact may be explained by bothan important extinction in these objects and by the pos-sible presence of an evolved underlying stellar population.The broad He ii λ ii λ ∼ We detect the weak auroral [O iii ] λ T e , all results are compiled in Table A.8. [O iii ] λ H ii regionsfollowing the typical diagnostic diagrams.The reddening coefficient was computed using all avail-able
H i
Balmer lines in each spectrum. We obtained verydifferent values: while the central region and knots in thenorthern tail have a low reddening coefficient, c (H β ) ∼ c (H β ) ∼ α images (see Figure 10of Paper I). Table A.8 compiles all chemical abundances computed forthe different knots analyzed in IRAS 08208+2816. The oxy-gen abundance of the central region, derived using the di-rect method, is 12+log(O/H)=8.33 ± − . ± .
11. This value ishigher than the N/O ratio expected for a galaxy with anoxygen abundance of 12+log(O/H) ∼ ∼ − ∼ .
64 (i.e.,almost the solar value). Knots iii ] λ ∼ ∼ ∼ / O) ∼ − .
84, and consistent with the value expectedfor a galaxy with almost solar metallicity. This result in-dicates that knot 2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
84, and consistent with the value expectedfor a galaxy with almost solar metallicity. This result in-dicates that knot 2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results -20-15-10-5051015202530-400 -300 -200 -100 0 100 200 -400 -300 -200 -100 0 100 200-20-15-10-5051015202530 -400 -300 -200 -100 0 100 200-20-15-10-5051015202530 center relative velocity (km s -1 ) d i s t an c e ( a r cs e c ) IRAS 08208+2816 - PA 10º
C star center d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) IRAS 08208+2816 - PA 355º C starcenter d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) IRAS 08208+2816 - PA 345º
Fig. 10.
Position-velocity diagrams for the three slit positionsobserved of IRAS 08208+2816 using the [O iii ] λ ◦ and 355 ◦ is because of the contamination by a brightstar. ing with another galaxy which nucleus coincides with knotC of IRAS 08208+2816. Figure 10 shows the position-velocity diagrams obtained forthe three slit positions. The [O iii ] λ ′′ ) and taking the centerof IRAS 08208+2816 as reference. As it was clearly seenin the bidimensional spectra, this object possesses very in-teresting kinematics, remarking a probable tidal stream inthe northern tail. In these areas we observe a velocity gra-dient larger than 300 km s − within 12 ′′ (11 kpc) in theslit position crossing the northern region of the system(PA=355 ◦ ). This velocity difference of 300 km s − is thesame that Perryman et al. (1982) reported between theseobjects. As we commented before, the spectra crossing thesouthern tail is affected by the contamination from thenearby bright star, but the slit position with PA=10 ◦ isfree of such contamination allowing the kinematic analysisof the southern zone. Again, we found an important velocitygradient towards negative values, that cannot be explainedby a rotating disk. Furthermore, this diagram shows anevident sinusoidal pattern with an amplitude larger than50 km s − at the center of the galaxy. This is an additionalevidence that we are observing a process of merging. Ingeneral, the agreement between the three diagrams is verygood, for example, the velocity found for knot ◦ is ∼ −
250 km s − , that corre-sponds quite well with the velocity observed at the end ofthe southern tail (knot ∼ −
200 km s − ) using theslit position with PA 10 ◦ .Although the kinematics of the system is not supportedby rotation, we have computed a tentative estimation of theKeplerian mass of the system. Assuming i =90 ◦ , ∆ v ∼ − (using the diagram with PA 345 ◦ that seems tobe less affected by the tidal tails) and a radius of ∼ ′′ (18.4 kpc), we derive M Kep ∼ × M ⊙ . The mass-to-luminosity ratio is quite low, M Kep /L B ∼ H i data available for this galaxy inthe literature, but it should be really interesting to com-pare the kinematics of the neutral gas with that found
POX 4 C - PA 25º
H8 [S II]He IH ε He I He I[Ar IV][Fe III]He II [O III] [O III]WR[Ne III] He I[O II] H δ [O III] H β H γ O b s e r v ed f l u x [ - e r g s - c m - Å - ] [N II] He I[Cl III] [O I] [S III] Wavelength [ Å ] [S II][N II]He I H α [O I] POX 4 C - PA 25º
Fig. 11.
ISIS 4.2m WHT spectrum for the center of POX 4.Fluxes are not corrected for reddening. The most importantemission lines have been labeled. See Figure 11 in Paper I forthe identification of the regions. here for the ionized gas. The warm dust mass is high, M dust =8.84 × M ⊙ , giving M dust /L B ∼ × − . Our complete analysis of the physical properties, chem-ical abundances and kinematics of the ionized gas inIRAS 08339+6517 and its companion dwarf galaxy was pre-sented in L´opez-S´anchez, Esteban & Garc´ıa-Rojas (2006).We reported weak spectral features that could be attributedto the blue WR bump at the center of the galaxy. The kine-matics of the ionized gas showed an interaction pattern thatindicates that the
H i tidal tail detected by Cannon et al.(2004) in the direction of the dwarf companion galaxy hasbeen mainly formed from material stripped from the maingalaxy. A star-forming region in the outskirts of the galacticdisk may be a TDG candidate.
The first indications of WR features in POX 4were noticed by Kunth & Joubert (1985) andCampbell, Terlevich & Melnick (1986) , because bothdetected the broad He ii λ . Vacca & Conti (1992) confirmed the presenceof a high number of O and WN stars in the brightestregion of POX 4 and detected the He ii λ ii λ ii λ iv λ ◦ . The spectrum of the cen-ter of POX 4 (Figure 11) is dominated by intense emissionlines and does not show any evidence of stellar absorption These authors named POX 4 as C 1148-203. These authors named POX 4 as C 1148-2020 in their Table 1and as Tol 1148-202 in their Table 2. Following the NED, theappropriate name of this galaxy is IRAS 11485-2018 = POX 4.´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 13 in the
H i or He i lines. As it was previously noticed byM´endez & Esteban (1997), broad low-intensity assymetricwings are detected in the profiles of the brightest emissionlines (H α and [O iii ]). The clear detection of the nebularHe ii λ H i
Balmer lines, indicating the presence of evolved stel-lar populations underlying the starburst. Table A.9 com-piles all emission lines detected in the center of POX 4 andin the dwarf companion object, as well as other importantproperties of their spectra.
The physical conditions and chemical abundances of theionized gas in POX 4 are compiled in Table A.10. Theelectron temperature calculated for the center of POX 4using [O iii ] λ T e (O iii ) = 14000 ±
600 K, suggestingthat it is a low metallicity object. Although [N ii ] λ T e (O ii ) from T e (O iii ). The electron densitywas computed using the [S ii ] λ λ n e ∼
250 cm − . Notice that the estimation of n e usingthe [Ar iv ] λ λ ∼
270 cm − , although it has a higher uncertainty. For thedwarf companion object we used empirical calibrations toestimate T e , being its electron density in the low-densitylimit. The comparison of the observed [O iii ] λ β and[N ii ] λ α ratios with the diagnostic diagrams let toclassify all regions as starbursts.The reddening coefficient at the center of POX 4was determined using 7 H i
Balmer lines, obtaining c (H β )=0.08 ± The oxygen abundance derived for the center of POX 4,12+log(O/H)=8.03 ± ± ii λ +3 contribution is expectedin the nebular gas, however this contribution is found tobe marginal, ∼ ± − . ± .
06 and − .
60. Despite the uncer-tainties, the resemblance of the chemical abundances maysuggest that the dwarf companion is not an independent ob-ject, as M´endez & Esteban (1999) concluded, but a TDGcandidate.
The position-velocity diagram obtained for the slit-positiontaken for POX 4 is shown in Figure 12. Because of itshigher intensity, we used the [O iii ] λ -40 -30 -20 -10 0 10 20 30 40 50 60-30-25-20-15-10-505101520 D i s t an c e [ a r cs e c ] relative velocity [ km s -1 ] POX 4 - PA 25º
Fig. 12.
Position-velocity diagram for the slit position observedin POX 4 using the [O iii ] λ of the H α profile, extracting 4 pixel bins (0.8 arcsec) andtaking as reference the velocity found in the center of thegalaxy (knot ∼
130 km s − between POX 4 and thecompanion galaxy, but this is not confirmed in our deeperspectroscopic data.Taking into account the complex kinematic structureshown in Figure 12, it is clear that we can not obtain a con-fident estimate of the Keplerian mass of POX 4. However,assuming a ratio of M Kep /L B ∼ M Kep ∼ × M ⊙ . The H i
Parkes All-SkySurvey (HIPASS; Barnes et al. 2001) shows a tentative de-tection of
H i emission. Our group has performed 21-cmobservations of POX 4 using the interferometer
AustraliaTelescope Compact Array (ATCA). Although a detailed de-scription and analysis of such observations will be presentedelsewhere (L´opez-S´anchez et al. 2010), the very preliminaryanalysis suggests that the system possesses a lot of neutralgas. The
H i kinematic are perturbed in the position of thedwarf companion object but it shows the same radial ve-locity we found using optical spectroscopy. An independent
H i cloud, that has the same radial velocity that POX 4,is found at ∼ ′ ( ∼
60 kpc) at the south. It shows a clearalignment with both the bright center of POX 4 and thedwarf companion object, suggesting a very probable inter-action in the past. A detailed analysis of the
H i observa- [Fe III][Ne III]H ε H δ He II[O II] He IHe I [O III] [O I] [S III][N II]He I H α [O III]H β H γ UM 420 - PA 90º O b s e r v ed f l u x [ - e r g s - c m - Å - ] Mg ICa II K H
G bandO II
Wavelength [ Å ] H δ O IIIH β H γ UGC 1809 - PA 90º
Fig. 13.
ISIS 4.2m WHT spectrum for the center of UM 420( top ) and the galaxy UGC 1809 ( bottom ). Fluxes are not cor-rected for reddening. The most important emission lines havebeen labeled. tions will confirm o discard the TDG nature of the dwarfcompanion object surrounding POX 4.
Izotov & Thuan (1998) reported the detection of thebroad He ii λ iv λ iv λ ii H,K, G-band and Mg i λ z =0.0243.The radial velocity of UGC 1809 is v r =7290 km s − , inexcellent agreement with the value given by the NED( v r =7306 km s − ) but much lower than the radial velocityof UM 420 ( v r =17507 km s − ). That confirms that bothgalaxies are not physically related.The spectra of UM 420 do not show absorption features.We observe, although with large error, the nebular He ii λ ii ] λ Using the [O iii ] λ T e (O iii )=13200 ±
600 K in UM 420. Although it has a largeerror, the detection of the auroral [N ii ] λ T e (N ii ) ∼ T e (O iii ) and T e (O ii ). The electron density computed us-ing the [S ii ] λλ ii ] λλ n e ∼
140 cm − . The reddening coefficient and the underly-ing stellar absorption in the H i
Balmer lines were computedusing 5 ratios between the
H i
Balmer lines and give veryconsistent results, which mean values are c (H β )=0.09 ± C (H β ) = 0.09 W abs = 2.0 UM 420
H9 / H β H ε / H β H δ / H β H γ / H β C ( H β ) Wabs H α / H β Fig. 14.
Interactive estimation of c (H β ) and W abs using the sixbrightest H i
Balmer lines detected in the spectrum of UM 420.Note the excellent agreement in the behaviour of all lines. and W abs =2.0 ± Table A.10 lists all the chemical abundances com-puted for UM 420. The derived oxygen abundance is12+log(O/H)=7.95 ± +3 (that should ex-ist because of the detection of the He ii ) but it is smallerthan 0.01 dex. The N/O ratio, log(N/O)= − . ± . − ± − ± Figure 15 shows the position-velocity diagram obtained forthe slit position with PA 90 ◦ observed in UM 420. Boththe H α and H β profiles were analyzed, extracting 3 pixelbins (1.08 ′′ ) for H α and 4 pixel bins (0.8 ′′ ) for H β . Thediagram is identical in both cases. Although the numberof points is small, we notice a velocity gradient of around30 km s − from the eastern region to the center of thegalaxy, but this tendency is reversed in the western areasof UM 420. Indeed, in this region a negative velocity gra-dient of ∼
70 km s − is found within only 4 ′′ ( ∼ H i data of the galaxyto derive its neutral or dynamical masses. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 15 -100 -90 -80 -70 -60 -50 -40 -30 -20 -10 0 10 20 30 40 50-8-6-4-202468 H α H β center D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] UM 420 - PA 90º
Fig. 15.
Position-velocity diagram for the slit position observedin UM 420 using both the H α and H β profiles. W is up. [Ar IV][Fe III]He II [O III]WRHe I[O III] H β H γ SBS 0926+606 A - PA 14º O b s e r v ed f l u x [ - e r g s - c m - Å - ] He I[Ar III] [O II][N II] He I[O I] [S III]
Wavelength [ Å ] [S II][N II]He I H α [O I] SBS 0926+606 A - PA 14º
Fig. 16.
ISIS 4.2m WHT spectrum for SBS 0926+606 A. Fluxesare not corrected for reddening. The most important emissionlines have been labeled. See Figure 15 in Paper I for the identi-fication of this region.
SBS 0926+606 actually is a pair of compact nearby ob-jects (see Figure 15 of Paper I). SBS 0926+606 A hasbeen spectroscopically studied by Izotov and collaboratorswith the aim of determining the primordial helium abun-dance (Izotov, Thuan & Lipovetski 1997; Izotov & Thuan1998) and the chemical abundances of heavy elementsin BCDGs (Izotov & Thuan 1999). Subsequent spectro-scopic studies were performed by P´erez-Montero & D´ıaz(2003) and Kniazev et al. (2004). Izotov et al. (1997) in-dicated the presence of broad low-intensity components inboth the H α and [O iii ] λ ii λ ◦ . Itmainly crosses the subregion A2 defined using our opticaldata (see Figure 15 in Paper I). As we see in Figure 16, thespectrum has a good spectral resolution, but it only covers between 4200 and 5000 ˚A in the blue range and between5600 and 7400 ˚A in the red range. Therefore, we did notobserve the [O ii ] λ iii ] λ ii λ α profile. Table A.10 compiles the values found for the electron tem-peratures of the ionized gas within these objects. For mem-ber A, T e (O ii ) was computed using the direct method be-cause of the detection of [O iii ] λ T e (O ii ) was es-timated using Garnett’s relation. The spectrum obtainedfor member B does not allow to calculate R , so weused the N ratio and the empirical calibrations given byDenicol´o et al. (2002) and Pettini & Pagel (2004) to esti-mate T e . As empirical calibrations involving the N ratioseem to overestimate the actual abundance (we will discussthis in Paper III), the electron temperatures derived forgalaxy B may be underestimated.The electron density, computed using the[S ii ] λλ H i
Balmer lines. The com-parison of the emission lines ratios with the diagnosticdiagrams indicates that both galaxies are starbursts.
Table A.10 compiles the chemical abundances derived inthis galaxy pair. Because the [O ii ] λλ ii ] λλ + /H + ratio. The chemical abundances de-rived for SBS 0926+606 A are 12+log(O/H) = 7.94 ± − ± − ± − ± N empirical calibrations is12+log(O/H) ∼ M B (see Paper I). Figure 17 shows the position-velocity diagram derivedfrom the slit position observed in SBS 0926+606. We ex-tracted 4 pixel bins (0.8 arcsec) across the H α profile.SBS 0926+606 A shows a clear sinusoidal pattern, withand amplitude of around 50 km s − , suggesting that thedouble nucleus we found in this galaxy (see Figures 15 and16 of Paper I) may be a consequence of an advanced merg-ing process between two objects. The northern outskirts -150 -100 -50 0 50 100-30-20-1001090100110 center AB d i s t an c e [ a r cs e c ] relative velocity [ km s -1 ] SBS 0926+606 - PA 14º
Fig. 17.
Position-velocity diagram for the slit position with PA14 ◦ observed in SBS 0926+606 using the H α profile. Notice thatthe y-axis is broken. NE is up. See Figure 15 in Paper I for theidentification of the regions. of SBS 0926+606 A seems to be partially decoupled fromthe sinusoidal pattern (there is a difference of around 60km s − with respect to the central velocity). On the otherhand, SBS 0926+606 B also shows a perturbed kinematics,because both its northern and southern edges have similarradial velocities. The elongated shape of SBS 0926+606 B,the two tails towards the west detected in our deep im-ages and the disturbed kinematics suggest that the interac-tion that this galaxy is experiencing –most probably withSBS 0926+606 A– is very close to the plane perpendicularto the line of sight, and therefore SBS 0926+606 B is ob-served almost edge-on. In any case, we do not detect anymorphological feature, such as the debris of a tidal tail or adiffuse non-stellar object, between both galaxies thus theirpossible interaction is nowadays no very strong.Finally, the complexity of the position-velocity dia-gram shown in Figure 17 does not allow to determine theKeplerian mass of the galaxy. Using the H i data given byPustilnik et al. (2002), we derive M H I =(9.6 ± × M ⊙ and M H I =(8.1 ± × M ⊙ for A and B, respectively,that indicate mass-to-luminosities ratios of M H I /L B =0.75and 0.59. The gas depletion timescales are 1.7 Gyr for Aand 5.5 Gyr for B. Assuming half of the amplitude of the H i velocity ( ∼
60 km s − for both galaxies) and effectiveradii of ∼ ′′ (2.71 kpc) for A and ∼ ′′ (5.42 kpc) forB, we estimate dynamical masses of M Dyn ∼ × M ⊙ and M Dyn ∼ × M ⊙ for A and B, respectively. Themass-to-luminosity ratios, M Dyn /L B =1.8 and 3.3 for Aand B, respectively, which are values typical for BCDGs(Huchtmeier, Krishna & Petrosian 2005). However, the M H I /M Dyn ratios are high, 0.42 and 0.18 for A and B,respectively, indicating that a considerable amount of thetotal mass of the galaxies ( ∼
42% for A) is neutral hydrogen.All these values indicate that the system still possesses ahuge amount of fresh material from which new stars may beborn. Indeed, the
H i profile obtained by Thuan et al (1999)shows two peaks, that coincide with the optical velocitiesof the galaxies, embedded in an common
H i envelope. Thisfact strongly suggests that a lot of neutral gas should befound between both galaxies. An HI map obtained using a [Ne III] [S II][O II] He I [Ar IV][Fe III]He II [O III]WRHe I[O III] H β H γ SBS 0948+532 - PA 114º O b s e r v ed f l u x [ - e r g s - c m - Å - ] [Cl III] [N II] He I[O I] [S III] Wavelength [ Å ] [N II]He I H α [O I] SBS 0948+532 - PA 114º
Fig. 18.
ISIS 4.2m WHT spectrum for SBS 0948+532 using aslit with PA 114 ◦ . Fluxes are not corrected for reddening. Themost important emission lines have been labeled. radio-interferometer would be necessary to study the distri-bution and kinematics of the neutral gas, giving key cluesabout the evolution of the system. SBS 0948+532 was studied by Izotov and collaborators(Izotov et al. 1994; Thuan et al. 1995; Izotov & Thuan1998; Guseva et al. 2000; Izotov & Thuan 2004).Schaerer et al. (1999) included this BCDG in theircatalogue of WR galaxies because of the detection of boththe broad and nebular He ii λ ◦ . The emis-sion line fluxes of the detected lines and other properties ofthe spectrum are compiled in Table A.11. No stellar absorp-tions are observed in this spectrum. We detect the broadand nebular He ii λ The intensity of [O iii ] λ n e ∼
250 cm − , was derived using the [O ii ] λλ n e estimated from the [Ar iv ] λ λ ∼
260 cm − . The reddening coefficient was estimatedwith a good precision because of the detection of many H i
Balmer lines. The equivalent width of the
H i
Balmerstellar absorption lines, W abs , is very small, suggesting thatthe underlying population of evolved stars is not important.The comparison of the observed emission line fluxes withthe diagnostic diagrams confirms that the gas is ionized bythe strong U V emission of the massive stars.
Table A.12 lists all the chemical abundances computedfor SBS 0948+532. The value of the oxygen abun-dance is 12+log(O/H)=8.03 ± ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 17 -70 -60 -50 -40 -30 -20 -10 0 10 20 30 40 50 60 70-10-8-6-4-20246810 Center d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) SBS 0948 + 532 - PA 114º
Fig. 19.
Position-velocity diagram for the slit position with PA114 ◦ observed in SBS 0948+532 using the [O iii ] λ log(N/O)= − . ± .
08, the typical found for objects withthe metallicity of this galaxy. In general, all our chemicalabundances for this BCDG agree with those estimated byIzotov & Thuan (1999) within the errors.
Figure 19 shows the position-velocity diagram obtained forthe slit position observed in SBS 0948+532 (see Figure 17in Paper I). We analyzed the [O iii ] λ − , is observed. This feature seems to be disturbed inthe SE areas, that is precisely the region where the faintarc is seen in our optical images (see Figure 17 of Paper I).Assuming that the kinematics of the galaxy is due to ro-tation, i =90 ◦ and ∆ v ∼
50 km s − within a radius of ∼ ′′ (3.63 kpc), we derive a Keplerian mass of M Kep ∼ × M ⊙ and M Kep /L ⊙ ∼ H i or FIR data are available for this galaxy.
The first spectroscopic study of SBS 1054+365 was per-formed by Izotov, Thuan & Lipovetski (1997), who de-tected the broad He ii λ ◦ , crossing the main body of thegalaxy along its major axis (see Figure 19 of Paper I).Hence, we observed knot b , the central bright region Cand part of the star-forming semi-ring located at the west(knot a ). We only analyzed the physical conditions andthe chemical abundances of the ionized gas in region Cand knot b . Figure 20 shows the spectrum of the center He I He I[Ar III]WR[O III] O b s e r v ed f l u x [ - e r g s - c m - Å - ] Wavelength [ Å ] [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ SBS 1054+365 C - PA 55º
Fig. 20.
IDS 2.5m INT spectrum for SBS 1054+364 using aslit with PA 55 ◦ . Fluxes are not corrected for reddening. Themost important emission lines have been labeled. See Figure 19in Paper I for identification of the region. of SBS 1054+365. Table A.11 compiles the dereddened lineintensity ratios and other properties of the spectrum. As wesee, the spectrum is dominated by nebular emission withoutany features of stellar absorptions. We do detect the nebu-lar He ii λ ii line(Figure 36). The red WR bump is not detected (Figure 37)perhaps because of the lacking of enough S/N in our spec-trum. The electron temperatures in the central region werecomputed using the [O iii ] λ T e (O iii ) and T e (O ii ). Forknot b we used the Pilyugin (2001a,b) empirical calibra-tions. The electron density was estimated using the [S ii ] λλ iii ] λ β ,[N ii ] λ α and [S ii ] λλ α ratios with thediagnostic diagrams allows to classify all knots as typical H ii regions.
Table A.12 compiles the chemical abundances derived forSBS 1054+364. The oxygen abundance computed for thecenter of the galaxy is 12+log(O/H)=8.00 ± − . ± .
09, in excellent agree-ment with the values obtained by Izotov & Thuan (1999).Despite of its higher error, the oxygen abundance and N/Oratio estimated for knot b are very similar to those foundat the center of the galaxy. We used our bidimensional spectrum for the slit positionwith PA 55 ◦ to build the position-velocity diagram shownin Figure 21. We extracted 3 pixel bins (1.2 arcsec) acrossthe H α profile and took as reference the brightest region ofthe galaxy. The diagram does not show a clear rotation pat-tern but several changes in the velocity distribution. Thecentral region seems to show a velocity gradient of around40 km s − between − ′′ and 10 ′′ . This feature was previ-ously noticed by Zasov et al. (2000) but their lower spatialresolution did not permit to see the small amplitude veloc-ity variations (see their Figure 3b). These authors suggestedthat this velocity gradient is consequence of the rotation ofthe galaxy. Our diagram also indicates that the SW region,the partial ring where knot a is located, does not followthe kinematics of the center of the galaxy, showing a veloc- -50 -40 -30 -20 -10 0 10 20 30 40 50 60-15-10-5051015202530 b2 b1aC D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] SBS 1054 + 365 - PA 55º
Fig. 21.
Position-velocity diagram for the slit position observedin SBS 1054+365 using the H α profile. NE is up. See Figure 19in Paper I for identification of the regions. ity variation of 40 km s − in 7.2 ′′ . On the other hand, thekinematics of the region b , that has an inverted velocitygradient of around 40 km s − , also seem to be decoupledfrom the movement of the gas in the central region, thatshows a positive velocity gradient of ∼
25 km s − in its NEarea. However, the amplitude of all the velocity variationsseen in Figure 21 are rather small and cover spatial exten-sions of the order of several hundred pc. Therefore, it ispossible that they could be due to local movements of thebulk of the ionized gas due to the combined action of windsor supernova explosions.In any case, assuming that rotation is present at thecenter of the galaxy and considering ∆ v ∼
20 km s − within a radius of ∼ ′′ (624 pc) and an inclinationof i ∼ ◦ (determined using its optical sizes) we havecomputed a tentative value for the Keplerian mass of M Kep ∼ × M ⊙ , that indicates M Kep /L ⊙ ∼ H i data provided by Zasov et al. (2000), we es-timated M H I =(6.08 ± × M ⊙ and M H I /L B ∼ ′′ (1.37 kpc) and the sameinclination angle, we estimated a dynamical mass of M Dyn ∼ × M ⊙ . The neutral gas to total mass ratio, M H I /M Dyn ∼ M Dyn /L B ∼ H i map can confirm this issue. The gas depletiontimescale is higher than 2.6 Gyr, indicating that the galaxystill possesses a huge amount of fresh material available tocreate new generations of stars.
SBS 1211+540 was included in the study of chemical abun-dances in BCDGs performed by Izotov and collabora-tors (Izotov et al. 1991; Thuan et al. 1995; Izotov & Thuan1998; Guseva et al. 2000; Izotov & Thuan 2004). WR fea-tures was firstly reported by Izotov et al. (1994), who de-tected the nebular He ii λ [Ne III] [Ne III][O II] [Ar IV] [O III]He I[O III] H β H γ SBS 1211+540 C - PA 138º O b s e r v ed f l u x [ - e r g s - c m - Å - ] [S III] [S II]He I[O I] Wavelength [ Å ] [N II]H α SBS 1211+540 C - PA 138º
Fig. 22.
ISIS 4.2m WHT spectrum for SBS 1211+540 using aslit with PA 138 ◦ . Fluxes are not corrected for reddening. Themost important emission lines have been labeled. See Figure 21in Paper 1 for identification of the region. (2000) only indicates the presence of the broad emissionline.Figure 22 shows our ISIS 4.2m WHT spectrum of thecenter of SBS 1211+540; Table A.11 compiles all its derivedproperties. The spectrum is dominated by the nebular emis-sion showing no traces of stellar absorptions. We do not de-tect the blue WR bump or the nebular He ii λ We derive a very high electron temperature, T (O iii )=17100 ±
600 K, using the direct method (seeTable A.12). The low ionization temperature was es-timated considering Garnett’s relation. Both the [O ii ] λλ ii ] λλ n e = 320 ±
50 cm − .The reddening coefficient was determined using all avail-able H i
Balmer lines with errors lower than 20%. Thecomparison of the emission line ratios with the diagnosticdiagrams confirms the starbursting nature of this BCDG.
Table A.12 compiles all the chemical abundancesderived for SBS 1211+540. The oxygen abun-dance, 12+log(O/H)=7.65 ± − . ± .
10, are in excellent agreement withthe values given by Izotov & Thuan (1999). Hence, itis the lowest metallicity object analyzed in this work.The rest of chemical abundances, log(S/O) ∼ − ∼ −
Figure 23 shows the position-velocity diagram obtained us-ing the bidimensional spectrum of SBS 1211+540. The slitposition we used, with a PA of 138 ◦ , crosses the center of thegalaxy but not knot a that, as it was explained in § iii ] λ ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 19 -60 -50 -40 -30 -20 -10 0 10 20 30-8-7-6-5-4-3-2-1012345678 Center d i s t an c e ( a r cs e c ) relative velocity (km s -1 ) SBS 1211 + 540 - PA 138º
Fig. 23.
Position-velocity diagram for the slit position withPA 138 ◦ observed in SBS 1211+532 using the [O iii ] λ reference. The position-velocity diagram does not show aclear rotation pattern, only a reverse of the velocity gradi-ent at the center of the galaxy. In any case, the amplitudeof the velocity variations are very small. If we do not con-sider the four lowest points at the SE of Figure 23 wherewe detected two faint plumes (see Figure 21 of Paper I), wemay assume that the kinematics is explained by rotationwith a velocity gradient of ∼ − .Considering that the gas is rotating with the param-eters described above, we may derive the Keplerian massof the galaxy. Assuming ∆ v ∼
30 km s − within a radiusof ∼ ′′ (315 pc) and an inclination angle of i ∼ ◦ (valuederived from the optical shape of the galaxy), we find M Kep ∼ × M ⊙ and M Kep /L B ∼ H i data provided by Huchtmeier et al. (2005), we derive M H I =(2.4 ± × M ⊙ and M Dyn ∼ × M ⊙ ,and hence M H I /M Dyn ∼ M H I /L B ∼ M Dyn /L B ∼ The only bibliographic references of SBS 1319+579are from Izotov and collaborators (Izotov et al. 1997;Izotov & Thuan 1998, 1999; Guseva et al. 2000;Izotov & Thuan 2004). Schaerer et al. (1999) included
He I [O II]He I[Ar IV][Ar IV]He I [O III] [Ar III][O I] [S III] [S II][N II]He I H α [O III]H β H γ SBS 1319 + 579 A O b s e r v ed f l u x [ - e r g s - c m - Å - ] [O II][Ar III]He I[O III] [O I] [S III] Wavelength [ Å ] [S II][N II]He I H α [O III]H β H γ SBS 1319 + 579 C
Fig. 24.
ISIS 4.2m WHT spectrum for regions A ( top ) and C( bottom ) of SBS 1319+579 obtained with a slit with PA 49 ◦ .Fluxes are not corrected for reddening. The most importantemission lines have been labeled. See Figure 23 in Paper I foridentification of the regions. SBS 1319+579 in their WR galaxies catalogue becauseIzotov et al. (1997) reported the detection of the broad andnebular He ii λ ◦ to analyze the ionized gasalong the main axis of the galaxy (see Figure 23 of Paper I).We got spectroscopic data of regions A, B, C, d and e butwe only analyzed A, B and C because of their higher S/Nratio. The spectra of the two brightest regions A and C areshown in Figure 24, and Table A.13 compiles the dered-dened flux ratios for all knots. Region B is the only onethat shows some stellar absorptions in its spectra. We donot have a clear detection of the blue WR bump or the neb-ular He ii λ ii in knot A (see Paper III). Wedo not detect the red WR bump in that region (Figure 37). The [O iii ] λ T e (O iii ) via the di-rect method. The low ionization electron temperatureswere estimated using Garnett’s relation. We found a sig-nificant difference in the electron temperatures found forthe brightest regions A and C, T e (O iii ) ∼ ii ] λλ Because of the lacking of [O ii ] λλ ii ] λλ + abundance.All the results for the chemical abundances derived inSBS 1319+579 are compiled in Table A.14. The oxygenabundance found in all regions are similar within the er-rors, 12+log(O/H) ∼ .
10, although that computed in re-gion A, 12+log(O/H)=8.05 ± ± -150 -100 -50 0 50 100-20-15-10-50510152025303540455055 H α emission ed A BC
Flux [ counts ] e d
CB A D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] SBS 1319 + 579 - PA 39º
Fig. 25.
Position-velocity diagram for the slit position observedin SBS 1319+579 using the H α profile. The relative intensity ofthe H α emission along the spatial direction is also shown, iden-tifying all observed regions. NE is up. See Figure 23 in Paper Ifor identification of the regions. a high excitation degree, log(O ++ /O + )=0.77, somethingthat is not observed in the other regions. The N/O ra-tios are very similar in A and B [log(N/O)= − ± − ± Figure 25 shows the position-velocity diagram obtainedfrom our bidimensional spectrum using a slit position withPA 39 ◦ . We extracted 4 pixel bins (0.8 arcsec) across theH α profile. The relative intensity of the H α emission alongthe spatial direction in also shown in this figure. Althoughthe velocity continuously decreases from the eastern regions( v ∼ −
105 km s − ) to the western areas ( v ∼
65 km s − ) ofthe galaxy, the velocity gradient is not the same across thesystem. We observe two tendencies: from region C to re-gion B (velocity difference of ∼
40 km s − in 30 ′′ ) and fromregion B to region A (velocity difference of ∼
130 km s − in 28 ′′ ). This behavior may suggest that there is a tidalstream moving from B to A in the direction away fromthe observer, but our deep images do not show such tailor any morphological feature that support this hypothesis.Another explanation to this feature may be the assump-tion that they are two systems, as we suggested from themorphology of the H α images and the chemical abundancesmay indicate, with different kinematics and in interaction.If this idea is correct, we should expect to observe distor-tions in the kinematics of the gas with higher amplitudesthat those we see. However, because of the high inclinationangle that the galaxy seems to have, i ∼ ◦ , we cannotdiscard any of both hypothesis.Considering that the kinematic pattern is consequenceof the rotation of the galaxy and assuming i ∼ ◦ and∆ v ∼
88 km s − within a radius of ∼ ′′ (4.2 kpc), wederive a Keplerian mass of M Kep ∼ × M ⊙ . The cor-responding mass-to-luminosity ratio, M Kep /L B ∼ v ∼
45 km s − , we now find M Kep ∼ × M ⊙ and M Kep /L B ∼ [Fe III]He II [O II]He I[Ar IV]He I [O III] [Ar III][O I] [S III] [S II][N II]He I H α [O III]H β H γ SBS 1415 + 437 C - PA 20º O b s e r v ed f l u x [ - e r g s - c m - Å - ] [O II][Ar III]He I[O III] [O I] [S III] Wavelength [ Å ] [S II][N II]He I H α [O III]H β H γ SBS 1415 + 437 A - PA 20º
Fig. 26.
ISIS 4.2m WHT spectrum for regions C ( top ) and A( bottom ) of SBS 1415+437 using a slit with PA 20 ◦ . Fluxes arenot corrected for reddening. The most important emission lineshave been labeled. See Figure 25 in Paper I for identification ofthe regions. BCDGs (Huchtmeier et al. 2005). This fact seems to con-firm that the kinematics surrounding region A are disturbedand not produced by rotation. Using the
H i data providedby Huchtmeier et al. (2007), we derive M H I =1 . × M ⊙ and M Dyn ∼ . × M ⊙ , assuming a rotation velocity of109 km s − within 4.5 kpc and the same inclination angle.The neutral gas accounts for only the 12% of all the massof the system. Its M Dyn /L B value, ∼ τ ∼ H i gas has been expelled from the galaxy. An
H i map obtained using a radio-interferometer that includesboth SBS 1319+579 and the nearby spiral NGC 5113 wouldbe fundamental to understand the dynamics and evolutionof this system.
The first spectroscopic data of SBS 1415+437 werereported by Thuan et al. (1995), who determined anoxygen abundance of 12+log(O/H)=7.51, being one ofthe less-metallicity galaxies known. A later reanaly-sis of the same spectrum raised this value to 7.59(Izotov & Thuan 1998, 1999; Thuan, Izotov & Foltz 1999).Their spectrum shows the broad and nebular He ii emis-sion lines, and therefore SBS 1415+437 was includedin the latest WR galaxies catalogue (Schaerer et al.1999). Subsequent spectroscopic analysis were pub-lished by Melbourne & Salzer (2002); Melbourne et al.(2004); Guseva et al. (2003); Izotov & Thuan (2004) andLee, Salzer & Melbourne (2004).Figure 26 shows the spectra of regions A and C obtainedusing the instrument ISIS at the 4.2m WHT and a slitwith PA 20 ◦ that crosses the main body of the galaxy (seeFigure 25 of Paper I). Although we detected some emissionlines in knot B, we have not analyzed its properties be-cause of the low S/N ratio of its spectrum. All spectra aredominated by nebular emission; no stellar absorptions aredetected. Table A.13 compiles all the line intensities ratiosand other properties of the spectra. Although we do not seethe broad blue WR bump, the nebular He ii λ ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 21 -80 -70 -60 -50 -40 -30 -20 -10 0 10 20 30 40 50 60 70 80 90-20-15-10-505101520253035 A B C D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] SBS 1415 + 457 - PA 20º
Fig. 27.
Position-velocity diagram for the slit position withPA 20 ◦ observed in SBS 1415+437 using the H α profile. NEis up. See Figure 25 in Paper I for identification of the regions. The electron temperatures were computed using the directmethod and are very high, T (O iii )=16400 and 15500 Kfor C and A, respectively. T e (O ii ) was estimated usingGarnett’s relation. The electron density was derived usingthe [S ii ] λ c (H β ) ∼ c (H β )in region A, c (H β )=0.16, suggests an inhomogeneous dis-tribution of dust in the galaxy. The comparison of the ob-served [O iii ] λ β and [N ii ] λ α ratios with thepredictions given by the diagnostic diagrams confirm theirstarbursting nature. Table A.14 compiles our results of the chemical abundancesderived in SBS 1415+437. These data confirm the verylow-metallicity of the galaxy, being the oxygen abundances12+log(O/H)=7.58 ± ± − . ± . The position-velocity diagram shown in Figure 27 was ob-tained extracting 4 pixel bins (0.8 ′′ ) along the H α pro-file of our bidimensional spectrum. We took the velocityof the brightest object C as reference. We observe thatthe velocity continuously decreases from the SW regions( v ∼
30 km s − , where region A is located) to the NE areas( v ∼ −
30 km s − , where region B is found), that may beattributed to the rotation of the galaxy. Some kinematic di-vergences are detected between regions A and C. However,because of the low amplitude of such variations (less than15 km s − ), they may just be a consequence of local move-ments in the ionized gas. Our position-velocity diagram [Ar III]WR [O I][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ III Zw 107 B - PA 0º O b s e r v ed f l u x [ - e r g s - c m - Å - ] [O II]He I[Ar III]He I[O III] Wavelength [ Å ] [O I] [S III][Ne III] [S II][N II]He I H α [O II] H δ [O III]H β H γ III Zw 107 A - PA 0º
Fig. 28.
IDS INT spectra for the regions A ( bottom ) and B( top ) of IRAS III Zw 107. Fluxes are not corrected for redden-ing. The most important emission lines have been labeled. SeeFigure 27 in Paper I for identification of the regions. is similar in both shape and values to that obtained byThuan, Izotov & Foltz (1999) using a slit with a PA of 22 ◦ (see their Figure 10). They also reported the peculiar kine-matic behavior we observe between A and C.The Keplerian mass we estimate for this galaxy, as-suming a rotation velocity of ∼
30 km s − within a ra-dius of ∼ ′′ (1.13 kpc) and an inclination angle of i ∼ ◦ (determined using the shape of the galaxy wesee in our optical images), is M Kep ∼ × M ⊙ , andits mass-to-luminosity ratio M Kep /L B =2.5. The H i massestimated by Huchtmeier, Krishna & Petrosian (2005) is M H I =(9.64 ± × M ⊙ . Using their data of W H I andconsidering a radius of 40 ′′ (1.8 kpc), we estimate a dynam-ical mass of M Dyn =4.9 × M ⊙ . With these data, we de-rive M H I /L B =0.96, M Dyn /L B =4.9 and M H I /M Dyn =0.20,that are the typical values found for BCDGs (Salzer et al.2002; Huchtmeier et al. 2005). These estimations are morereliable that those given by Thuan, Izotov & Foltz (1999)because we are using recent data with a higher sensibility.Both the gas depletion timescale ( ∼ III Zw 107 was analyzed using spectroscopy by Sargent(1970); Gallego et al. (1997) and Kunth & Joubert (1985).The last authors detected a continuum excess in the spec-tral region of the blue WR bump in the southern object,and hence Schaerer et al. (1999) included this BCDG intheir catalogue of WR galaxies.A slit with a PA of 0 ◦ was used in the IDS spectrographat the 2.5m INT to observe III Zw 107 (see Figure 27 ofPaper I). Three different regions, A, B and C, were ex-tracted. The optical spectra of regions A and B are shownin Figure 28. Region A possesses important underlying stel-lar absorption features. The faint region C is not very ev-ident from our optical images, but it is clearly identifiedat the north of region B in our bidimensional spectrum.Table A.15 compiles all the emission intensity ratios andother properties of the spectra of III Zw 107. We clearly de-tect the broad He ii λ2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
30 km s − within a ra-dius of ∼ ′′ (1.13 kpc) and an inclination angle of i ∼ ◦ (determined using the shape of the galaxy wesee in our optical images), is M Kep ∼ × M ⊙ , andits mass-to-luminosity ratio M Kep /L B =2.5. The H i massestimated by Huchtmeier, Krishna & Petrosian (2005) is M H I =(9.64 ± × M ⊙ . Using their data of W H I andconsidering a radius of 40 ′′ (1.8 kpc), we estimate a dynam-ical mass of M Dyn =4.9 × M ⊙ . With these data, we de-rive M H I /L B =0.96, M Dyn /L B =4.9 and M H I /M Dyn =0.20,that are the typical values found for BCDGs (Salzer et al.2002; Huchtmeier et al. 2005). These estimations are morereliable that those given by Thuan, Izotov & Foltz (1999)because we are using recent data with a higher sensibility.Both the gas depletion timescale ( ∼ III Zw 107 was analyzed using spectroscopy by Sargent(1970); Gallego et al. (1997) and Kunth & Joubert (1985).The last authors detected a continuum excess in the spec-tral region of the blue WR bump in the southern object,and hence Schaerer et al. (1999) included this BCDG intheir catalogue of WR galaxies.A slit with a PA of 0 ◦ was used in the IDS spectrographat the 2.5m INT to observe III Zw 107 (see Figure 27 ofPaper I). Three different regions, A, B and C, were ex-tracted. The optical spectra of regions A and B are shownin Figure 28. Region A possesses important underlying stel-lar absorption features. The faint region C is not very ev-ident from our optical images, but it is clearly identifiedat the north of region B in our bidimensional spectrum.Table A.15 compiles all the emission intensity ratios andother properties of the spectra of III Zw 107. We clearly de-tect the broad He ii λ2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results low spectral resolution of these 2.5m INT spectra. [O iii ] λλ α show broad wings in their profiles,more evident in the spectrum of region A. In region A we computed both T e (O iii ) and T e (O ii ) us-ing the direct method because of the detection of [O iii ] λ ii ] λλ T e (O iii )=10900 ±
900 K and T e (O ii )=10500 ±
800 K. Theelectron temperatures for knots B and C were estimatedusing empirical calibrations. All results are compiled inTable A.16. The electron density was computed using the[S ii ] λλ H i
Balmer lines of thisregions gives a much higher value in region A, c (H β ) ∼ c (H β ) ∼ H ii regions.
Table A.16 lists all the chemical abundances computedfor the bursts analyzed in III Zw 107. The oxygen abun-dance of the region A, derived using the direct method,is 12+log(O/H)=8.23 ± ∼ The position-velocity diagram obtained using our bidimen-sional spectrum is shown in Figure 29. We extracted 3 pixelbins (1.2 arcsec) along the H α profile, taking as referencethe velocity of region A, the brightest knot. We observe anegative velocity gradient between the southern regions ofthe galaxy ( ∼
30 km s − ) and region A ( ∼ −
20 km s − ), butbetween this knot and region B a reverse of the velocity of ∼
40 km s − is found within 4 ′′ . The velocity difference be-tween region B and C is ∼ −
40 km s − . Hence, althoughthe velocity amplitudes are not large and the spatial reso-lution is not very high, the position-velocity diagram seemsto show a sinusoidal pattern This feature may suggest in-teraction or merging phenomena between the two brightestknots seen in III Zw 107. This hypothesis would explain theexistence of the tail found in the deep optical images (seeFigure 27 of Paper I). It is also possible that the velocitygradient observed at the south of the galaxy is a conse- -60 -50 -40 -30 -20 -10 0 10 20 30 40 50 60 70 80-15-10-505101520 C BA D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] III Zw 107 - PA 0º
Fig. 29.
Position-velocity diagram for the slit position withPA ◦ observed in III Zw 107 using the H α profile. N is up.See Figure 27 in Paper I for identification of the regions. quence of the movement of the ionized gas within/towardsthat tail.We have performed a tentative estimation of theKeplerian mass of III Zw 107 using the position-velocity diagram shown in Figure 29. Considering thatwe observe the galaxy edge-on ( i ∼ ◦ ) and assum-ing ∆ v ∼
30 km s − within a radius of ∼ ′′ (3.9 kpc),we estimate M Kep ∼ × M ⊙ and M Kep /L B ∼ M H I =(6.7 ± × M ⊙ , M Dyn ∼ × M ⊙ (the dynamical mass was estimatedusing a radius of 20 ′′ =7.8 kpc and the half of the H i width, ∼
100 km s − ), and the mass-to-luminosity ratios de-rived from them, M H I /L B =0.38 and M Dyn /L B ∼
1. Indeed,if all these values are right, around 37% of the mass ofthe system is neutral gas. We consider that, because of thedetection of an important population of old stars withinthe galaxy and its relatively high metallicity, the dynami-cal mass of III Zw 107 has been probably underestimated.Hence, the
H i distribution should be several times largerthan the optical extent. The comparison of the velocity am-plitudes between the optical ( ∼
30 km s − ) and the radio( ∼
100 km s − ) data strongly supports this idea. Perhaps,the neutral gas has been expelled and/or dispersed as aconsequence of the possible interaction or merging betweenthe two main objects observed in III Zw 107. An interfer-ometric H i map of this galaxy is needed to answer to allthese issues. In any case, the gas depletion timescale is high, τ ∼ The WR nature of Tol 9 has been controversial.Penston et al. (1977) indicated a probable detection of afaint emission line around λ ii emission linebut suggested the detection of the red WR bump. BothConti (1991) and Schaerer et al. (1999) included Tol 9 intheir lists of candidate WR galaxies. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 23 [O III] He IHe IH ε [Ne III] WR [O II]H δ [O II] [Ar III][O I] [S II]+ [N II] He IH α [O III] H β H γ Tol 9 - INT - PA 49º O b s e r v ed f l u x [ - e r g s - c m - Å - ] [Ar III] [O II]He I[O I][Ne III] H δ He I [O III] [O I]
Wavelength [ Å ] [S II]+ [N II]H α [O II] H β H γ Tol 9 - NOT - PA 109º
Fig. 30.
IDS INT spectrum ( bottom ) using a slit with PA 49 ◦ and ALFOSC 2.56m NOT spectrum ( top ) using a slit of PA109 ◦ of Tol 9. Fluxes are not been corrected for reddening. Themost important emission lines have been labeled. See Figures 29and 30 in Paper I for identification of the regions. We observed two slit positions of Tol 9 (see Figure 29of Paper I). For the first one, we used the IDS instrumentat the 2.5m INT, using a slit with a PA of 49 ◦ that crossesthe center of Tol 9 and the dwarf companion galaxy lo-cated at the SW. The second slit position was taken withthe ALFOSC instrument at the 2.56m NOT, choosing a slitposition almost perpendicular to that used at the 2.5m INT.This last slit position was centered at the SW of the centerof the galaxy and the PA was set to 109 ◦ , our main ob-jective was to analyze the kinematics and properties of thefilamentary structure of ionized gas found in our deep H α images (see Figure 30 of Paper I). The list of all the emissionlines observed using both slits, as well as other importantproperties of the spectra, are shown in Table A.15. As wesee, both spectra show very similar line intensities. Noticethat the radial velocity obtained for Tol 9 using our opticalspectra is ∼
200 km s − more positive than the value previ-ously reported and listed in the NED (Lauberts & Valentijn1989).The spectra obtained for Tol 9 show nebular emissionand a continuum dominated by stellar absorptions in the H i
Balmer lines. We also observe an important decrementin the continuum at the blue range of the spectra. Thisfeature can be explained by both the contribution of theolder stars and an high extinction. We detect both the blueWR bump and the nebular He ii λ The 2.5m INT spectrum of Tol 9 shows [O iii ] λ ii ] λ ii ] λλ T e (O iii ) ∼ T e (low) ∼ T e (O ii ) using the [O ii ] lines and estimate the T e (O iii ) us-ing Garnett’s relation. In all cases, the electron densities -90 -60 -30 0 30 60 90 120 150 180-30-20-1001020100120 -150 -120 -90 -60 -30 0 30 60-30-20-100102030 Tol 9 - NOT PA 109º D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] H α [O III] 5007C Tol 9 - INT - PA 49º D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] H α [O III] 5007 Fig. 31.
Position-velocity diagrams for the slit positions ob-served in Tol 9: PA 49 ◦ ( left , using 2.5m INT data) and PA109 ◦ ( right , using 2.56m NOT data). Both the H α (circles) andthe [O iii ] λ were in the low-density limit. As we said, the values forthe reddening coefficients are high, about c (H β ) ∼ H i
Balmer lines, com-puted interactively with c (H β ), are large and give valuesof W abs ∼ Table A.16 compiles all the chemical abundances com-puted in Tol 9, showing almost identical results for bothspectra. The average value of the oxygen abundance de-rived using the direct method is 12+log(O/H)=8.57 ± − ± − . ± .
12, log(Ne/O)= − . ± . − . ± . The kinematics of the ionized gas in Tol 9 were analyzedusing our bidimensional spectra. We extracted 3 pixel bins(1.2 ′′ ) and 5 pixel bins (0.95 ′′ ) along the H α and the [O iii ] λ α and the [O iii ] λ ◦ , that crosses the cen-ter of Tol 9, does not show a rotation pattern. Indeed, weobserve two velocity gradients: while the velocity changes ∼
120 km s − from the NE regions to the center, this ten-dency is completely reversed in the SW region, that showsa velocity variation of ∼ −
120 km s − . Notice that the center of Tol 9 is not located in the (0,0) position ofthis diagram because the maximum of the H α emissionis displaced ∼ ′′ towards the SW of the center. Perhapsthe velocity pattern we seen at this PA is a combinationof rotation in the center and NE areas and the veloc-ity gradient produced by the optical tail we detected inour deep optical images that connects Tol 9 with a dwarfcompanion galaxy located at the SW (see Figure 29 ofPaper I). Assuming ∆ v ∼
70 km s − within a radius of ∼ ′′ (2.5 kpc) and an inclination angle of i ∼ ◦ , weestimate a Keplerian mass of ∼ × M ⊙ and a mass-to-luminosity radio of M Kep /L B ∼ H i measurements are neededto get a reliable estimation of the dynamical mass of thisgalaxy because single-dish observations (such those pro-vided by HIPASS) would blend the
H i gas in Tol 9 and thenearby spiral ESO 436-46. This analysis will be performedusing the new
H i
ATCA data for the galaxy group whereTol 9 resides obtained by our group (L´opez-S´anchez et al.2010). Using the
FIR fluxes, the warm dust mass is M dust ∼ × M ⊙ . Following the analysis performed byBettoni et al. (2003), the derived mass-to-luminosity ratio, M dust /L B ∼ × − , agree with the typical value found inspiral galaxies ( M dust /L B ∼ × − ).On the other hand, the position-velocity diagram ob-tained using the slit with PA 109 ◦ , shows several reversesin the velocity of the ionized gas. Regions α map (see Figure 30 of Paper I), are labeled in this dia-gram. They show a similar kinematic behavior (a variationof ∼ −
90 km s − in their velocities within the same dis-tance, ∼ ′′ (2.5 kpc), with respect to the center of thesystem. This kinematic structure, which is not coincidentwith any stellar distribution, reminds an expanding bipo-lar bubble, reinforcing the hypothesis that the H α envelopesurrounding Tol 9 is consequence of some kind of galacticwind. Knot ′′ (21 kpc)from the maximum of H α emission, but it seems to showa radial velocity similar to that observed at the ending ofthe filament Tol 1457-262 was studying using spectroscopy by Winkler(1988); Terlevich et al. (1991); Kewley et al. (2001);Westera et al. (2004) and Buckalew, Kobulnicky & Dufour(2005). Schaerer et al. (1999) included Tol 1457-262 in theirWR galaxies catalogue because Contini (1996) detected thebroad He ii λ ii λ Object 1 , seeFigure 31 of Paper I) using the instrument ALFOSC at
PA 155º [O III][Ne III]H ε H γ He I [O II]He IHe II[Fe III]He I [O II] [Ar III][O I] [S III] [S II][N II]H α [O III]H β H δ Tol 1457-262 A O b s e r v ed f l u x [ - e r g s - c m - Å - ] PA 155º
He IHe II[Ne III][O II] H ε H δ H γ [O II][Ar III]He I[O III] [S II][N II]H α [O III]H β Tol 1457-262 B H ε [Ne III][O II] H δ [O II][Ar III]He I[O III] [O I] Wavelength [ Å ] [S II][N II]H α [O III]H β H γ Tol 1457-262 C - PA 155º
Fig. 32.
ALFOSC 2.56m NOT spectra of the regions analyzedin Tol 1457-262. Fluxes are not corrected for reddening. Themost important emission lines have been labeled. See Figure 31in Paper I for identification of the regions. the 2.56m NOT and a slit with a PA of 155 ◦ . Table A.17compiles all the line intensities ratios and other propertiesof all the spectra analyzed in this galaxy. Stellar absorp-tions are barely detected, indicating that the nebular emis-sion strongly dominates their spectra. Broad low-intensitywings are detected in the H α profile in region A. We ob-serve the nebular He ii λ The electron temperatures were computed using the directmethod because of the detection of [O iii ] λ ii ] λλ T e (O ii ) and T e (O iii ) provided byGarnett (1992). Despite their similar ionization degree, re-gion B has an electron temperature, T e (O iii ) ∼ T e (O iii ) ∼ n e ∼
200 cm − , but n e was inthe low-density limit in region B. The determination ofthe reddening coefficient for region A was done using 5 H i
Balmer ratios, yielding consistently a very high value, c (H β )=0.83 ± c (H β ) found in region Bgave a negative value. This result may not attributed to abad flux calibration because adjacent regions A and C donot have this problem. Hence, we assumed c (H β ) ∼ and scaled the blue and red spectra considering thetheoretical ratio between the H α and H β fluxes for the elec-tron temperature estimated for this region. The comparisonof the emission line ratios with the diagnostic diagrams in-dicates that all regions can be classified as starbursts. However, the Galactic value usingSchlegel, Finkbeiner & Davis (1998) is c (H β )=0.23 ± -200 -150 -100 -50 0 50 100 150 200 250-25-20-15-10-5051015 H α emission d ABC
Flux [ counts ] d CB A D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] Tol 1457-262
Obj 1 - PA 155º
Fig. 33.
Position-velocity diagram for the slit position observedin
Object 1 of Tol 1457-262 using the H α profile. The relativeintensity of the H α emission along the spatial direction is alsoshown, identifying all observed regions. NW is up. See Figure 31in Paper I for identification of the regions. The chemical abundances computed for the regions an-alyzed in Tol 1457-262 are listed in Table A.18. Theoxygen abundance derived for the brightest knot (regionA) is 12+log(O/H)=8.05 ± ± ± ii ] emission lines, that givelog(N/O)= − ± − ± Figure 33 shows the position-velocity diagram of the slitposition with PA 155 ◦ analyzed in Object 1 in Tol 1457-262. We extracted 5 pixel bins (0.95 ′′ ) along the H α profileand took as reference the center of region A (the maxi-mum of the H α emission). Figure 33 includes the relativeintensity of the H α emission along the spatial direction.Although a velocity gradient between the northern regions( v ∼ −
120 km s − ) and the southern areas ( v ∼
190 km s − )is seen, we notice several velocity reverses along the sys-tem, being the most important that found between regionsA and C, that has an amplitude larger than 100 km s − .These features indicate that the kinematics of the differentstar-forming regions found within Object 1 of Tol 1457-262are decoupled. Indeed, the observed sinusoidal pattern sug-gests that the system is experiencing a merging process.The fast increasing of the velocity between knot B and thesouthernmost regions may be related to the movement ofthe material within the faint tail we detected in our deepoptical images (see Figure 31 of Paper I). We consider thatknot C is a TDG candidate because of its kinematics andchemical abundances are similar to those derived in brightregion A.
He IHe IH ε [Ne III] He II [O II][Ar IV]H δ [O II] [Ar III][O I] [OI] [S II]+ [N II] He IH α [O III]H β H γ ESO 566-8 - PA 342º O b s e r v ed f l u x [ - e r g s - c m - Å - ] WR Bump? He I [O III] [O I]
Wavelength [ Å ] [S II][N II]H α [O II] H β H γ ESO 566-7 - PA 342º
Fig. 34.
ALFOSC 2.56m NOT spectra for the galaxy pair ESO566-8 ( top ) and ESO 566-7 ( bottom ) that constitute Arp 256.Fluxes are not been corrected for reddening. The most impor-tant emission lines have been labeled. See Figure 34 in Paper Ifor identification of the galaxies.
Assuming that the general kinematics pattern is conse-quence of rotation, we derived a tentative Keplerian massof M Kep ∼ × M ⊙ and a mass-to-luminosity ratio of M Kep /L B ∼ Object 1 in Tol 1457-262. We consid-ered ∆ v ∼
60 km s − within a radius of r ∼ ′′ (4.95 kpc)and an inclination angle of i ∼ ◦ (from our optical im-ages). HIPASS provides a detection of the H i gas in thisgalaxy, for which we derive M H I =4.7 × M ⊙ . However,this estimation for the neutral gas mass is for all the ob-jects that compose Tol 1457-262 and therefore an interfer-ometric H i map is needed to quantify the amount of
H i and the dynamical mass of each member. The total lumi-nosity of the system, computed from our optical data, is L B =1.82 × L ⊙ , and hence the neutral hydrogen mass-to-luminosity radio of Tol 1457-262 is M H I /L B ∼ M H I /L B found in simi-lar galaxies, indicating the large amount of neutral gas inTol 1457-262. Arp 252 is a pair of interacting galaxies de-signed ESO 566-8 (galaxy A) and ESO 566-7(galaxy B). Their spectroscopic properties wereanalyzed by Pe˜na, Ruiz & Maza (1991) andMasegosa, Moles & del Olmo (1991). These authorsdetected the blue WR bump in ESO 566-7 . The WRfeature was confirmed by Pindao (1999), and henceESO 566-7 was included in the latest catalogue of WRgalaxies (Schaerer et al. 1999). Contini (1996) reported atentative detection of the nebular He ii λ ◦ (see Figure 34 inPaper I). The spectrum of ESO 566-8 shows many emissionlines, but the spectrum of ESO 566-7 only shows a few.All line intensity ratios are listed in Table A.17. Althoughthe spectra are dominated by the emission of the ion-ized gas, they also show some stellar absorptions, that are Masegosa et al. (1991) named this object C 0942-1929A, butit is incorrect following Schaerer et al. (1999). Notice that the real slit position in Figure 34 of Paper I is342 ◦ = − ◦ and not 18 ◦ .6 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results more evident in the weakest H i
Balmer lines observed inESO 566-7. We detect the nebular He ii λ The detection of the [O ii ] λλ ii ] λ T e (O ii ) ∼ T e (N ii ) ∼ T e (low)=9100 ±
800 K. The high ionization electron tem-perature was computed using Garnett’s relation. The elec-tron temperatures in ESO 566-7 were estimated via empir-ical calibrations. The electron densities were derived usingthe [S ii ] λλ c (H β ) ∼ c (H β ) ∼ Table A.18 compiles all the chemical abundances derivedfor the galaxy members of Arp 252. The oxygen abun-dance in ESO 566-8 (galaxy A) is 12+log(O/H)=8.46 ± ± − ± Figure 35 shows the position-velocity diagram obtained forArp 252 analyzing the bidimensional spectrum using a slitwith PA 342 ◦ . The H α profile was analyzed extracting 5pixel bins (0.95 ′′ ) and taking as reference the maximumof the emission in ESO 566-8. As we see, both galaxieshave a slight velocity difference of 50 km s − between theircenters. However, we appreciate important differences inthe kinematics of the galaxies: while A seems to be rotat-ing in its central regions (there is a velocity variation of −
130 km s − at the north and 100 km s − at the south), Bshows a distorted kinematic pattern. The velocity gradientsobserved in the upper and lower regions of galaxy A maycorrespond to the tidal streams induced in the long tails weobserve in the optical images (see Figure 34 of Paper I). -150 -100 -50 0 50 100 150 200-40-35-30-25-5051015 (Galaxy B)(Galaxy A) ESO 566-8 ESO 566-7 D i s t an c e [ a r cs e c ] Relative velocity [ km s -1 ] Arp 252 - PA 342º
Fig. 35.
Position-velocity diagram for the slit position observedin Arp 252 using the H α profile. Notice that the y-axis is brokenin two parts. N is up. See Figure 34 in Paper I for identificationof the galaxies. We performed a tentative determination of theKeplerian mass of the galaxies. We assumed a velocity of∆ v ∼
100 km s − within a radius of r ∼ ′′ (3.15 kpc) inESO 566-8 and a velocity of ∆ v ∼
30 km s − within a ra-dius of r ∼ ′′ (1.89 kpc) in ESO 566-7. For both we consid-ered an inclination angle of i = 90 ◦ , hence our Keplerianmass determinations are low limits to the real ones. Weconsider, however, that this assumption is not bad be-cause in ESO 566-8 the northern tidal tail shows a highinclination angle with respect to the plane of the sky andESO 566-7 shows its long southern tail almost in the planeof the sky and its disk seems to be edge-on. We estimatedKeplerian masses of M Kep ∼ × M ⊙ for ESO 566-8and M Kep ∼ × M ⊙ for ESO 566-7, that indicate amass-to-luminosity ratios of M Kep /L B ∼ FIR fluxes, is M dust ∼ × M ⊙ .Arp 252 is not detected in HIPASS, and hence we cannot derive the neutral gas mass and the dynamical mass.However, considering the absolute magnitude of the maingalaxy, M B = − .
9, and despite the distance to the sys-tem ( D ∼
130 Mpc) we should expect some
H i emission.Hence, or Arp 252 does not have too much neutral gas or ithas been lost in the intergalactic medium because of tidaleffects. Finally, it would be very interesting to analyze thekinematics of knots c and d to check their probable TDGnature (see Figure 34 of Paper I). The echelle spectrophotometric analysis of the BCDGNGC 5253 was presented in L´opez-S´anchez et al. (2007).We measured the intensities of a large number of permit-ted and forbidden emission lines in four zones of the centralpart of the galaxy. The physical conditions of the ionizedgas were derived using a large number of different line inten-sity ratios. Chemical abundances of He, N, O, Ne, S, Cl, Ar,and Fe were determined following the standard methods. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 27
We detected, for the first time in a dwarf starburst galaxy,faint C ii and O ii recombination lines. We confirmed thepresence of a localized N enrichment in certain zones ofthe center of the galaxy and suggested a possible slightHe overabundance in the same areas. We shown that theenrichment pattern agrees with that expected for the pol-lution by the eyecta of WR stars. The amount of enrichedmaterial needed to produce the observed overabundance isconsistent with the mass lost by the number of WR starsestimated in the starbursts.The analysis of the H i data provided by the
LocalVolumen
H i
Survey project (Koribalski 2008) reveals thatthe neutral gas kinematics within NGC 5253 has a velocitygradient along the optical minor axis of the galaxy; it doesnot show any sign of regular rotation (L´opez-S´anchez et al.2008). Some authors suggested that this feature is an out-flow, but most likely its origin is the disruption/accretion ofa dwarf gas-rich companion (Kobulnicky & Skillman 2008)or the interaction with another galaxy in the M 83 sub-group. The finding of a distorted
H i morphology in theexternal parts of the galaxy supports this hypothesis. Acomprehensive analysis of the neutral gas within NGC 5253will be presented elsewhere (L´opez-S´anchez et al. 2010).
4. Summary
We have presented a detailed analysis of the ionized gaswithin 16 Wolf-Rayet galaxies using long-slit intermediate-resolution optical spectroscopy. In many cases, more thantwo star-forming regions have been studied per galaxy. Wehave analyzed the physical properties of the ionized gas,deriving their electron temperatures, electron density, thereddening coefficient and the stellar absorption underlyingthe
H i
Balmer lines. We have confirmed that the excita-tion mechanism of the ionized gas in all bursts is mainlyphotoionization and not due to shock excitation as it hap-pens in AGNs and LINERs. In the majority of the cases,we have computed the chemical abundances of O, N, S, Ne,Ar and Fe using the direct determination of the electrontemperature (see second column in Table 2). When thesedata were not available, we used the empirical calibration ofPilyugin (2001a,b) to get an estimation of the metallicityof the ionized gas. We have estimated the oxygen abun-dance of many new regions within the sample galaxies andrefined the chemical properties of some of them, remarkingregions in HCG 31, Mkn 1087, Mkn 1199, III Zw 107, Tol 9,Tol 1457-262 and NGC 5253.The derived physical and chemical properties were usu-ally in agreement with previous observations reported inthe literature. Table 2 compares our oxygen abundance de-terminations with those previously reported in the litera-ture. As we see, the majority of the results agree well withprevious estimations, but there are important differences inMkn 1199, III Zw 107, Tol 9 and Tol 1457-262.Including the data of the four systems analyzed in pre-vious papers, a very useful database of objects with oxygenabundances between 7.58 and 8.75 in units of 12+log(O/H)is provided. In Papers IV and V we will explore thisdatabase comparing their properties with other data de-rived from both our deep optical/NIR images and othermultiwavelength observations available in the literature.We have confirmed the detection of Wolf-Rayet featuresin the majority of the galaxies, as we should expect because
10 100 1000101001000
HCG 31 G
EW(H a ) spec = EW(H a ) phot Linear fit E W ( H a ) [ Å ] - I m age r y EW (H a ) [Å] - Spectroscopy Fig. 38.
Comparison between the H α equivalent widths derivedfrom our H α images and those obtained from the analysis ofthe optical spectrum for every particular star-forming regionanalyzed in this work. our sample was extracted from the latest WR galaxies cat-alogue (Schaerer et al. 1999). We have reported the detec-tion of broad WR features in 20 regions within 16 systems.Figure 36 shows a detail of the optical spectrum in the4600–4750 ˚A range (blue WR bump) of all important ob-jects; faint regions with very low S/N have been excluded.We have indicated the spatial localization of the massivestars in each system. WR features are sometimes found indifferent knots within a same galaxy (HCG 31, Haro 15,Tol 1457-262, NGC 5253). The He ii λ iii λ iv λ iv λ Fig. 36.
Detail of the spectra of the main regions within our galaxy sample showing the zones around the blue WR bump. Thered dotted line represents the position of the He ii λ iii ] λ iv ] λλ Fig. 37.
Detail of the spectra of the main regions within our galaxy sample showing the zones around the red WR bump. Thered dotted line represents the position of the C iv λ ii ] λ i λ is through the H α equivalent width since it decreases withtime. Our spectroscopic observations provide an indepen-dent estimation of W (H α ) within every knot. We havechecked the correspondence between the values we obtainedin our deep H α images (see Table 7 in Paper I) with thosederived from the spectroscopic data. Figure 38 plots suchcorrelation, showing the excellent agreement of both kindsof data in almost all objects. Indeed, the linear fit to thedata practically coincides with a x = y function. Small di-vergences are found in very few objects, but they can be explained because of considerable differences in the relativesizes for which the photometric/spectroscopic values wereextracted. The most evident case is member G in HCG 31,for which we only extracted the spectrum of a small knotat its NW, but the W (H α ) value derived from the imagesconsiders the flux of all the galaxy, that possesses a globalstar-formation activity lower than that observed in the NWknot.Figure 39 shows the comparison of our data with thetheoretical predictions provided by the last release of the SB99 Geneva, Z=0.4Zo E W ( H β ) [ Å ] - f r o m s pe c t r o sc op y Age [Myr] - from H α images POX 4 Comp
SB99, Padova Z=0.2Zo SB99, Padova Z=0.4Zo
Fig. 39. H β equivalent width vs. age of the most recent star-forming burst diagram comparing the predictions given by theevolutionary synthesis models provided by STARBURST 99(Leitherer et al. 1999). We include the Z/Z ⊙ =0.4 model orig-inally included in STARBURST 99 using Geneva tracks andtwo new models with Z/Z ⊙ = 0.2 and 0.4 than consider Padovatracks (Vazquez & Leitherer, 2005). The age was computed fromthe W (H α ) given by our images; W (H β ) was determined fromour optical spectra. STARBURST 99 (Leitherer et al. 1999) models which usesPadova tracks (V´azquez & Leitherer 2005). We assumed aninstantaneous burst with a Salpeter IMF, a total mass of10 M ⊙ , and a metallicity of Z/Z ⊙ = 0.2 and 0.4, themost common values according to the oxygen abundanceof the majority of the knots. For comparison, we have alsoincluded the predictions of the original STARBURST 99models, that consider Geneva tracks, for Z/Z ⊙ =0.4. Theages of the last star-formation event are those estimatedfrom the W (H α ) determined from our deep images, whilethe H β equivalent widths are those directly measured fromour spectra. Thus, both set of data come from indepen-dent observations. As we see, the agreement is excellent,and therefore we are quite confident in the determinationof the age of the star-forming regions. As it should be ex-pected, the predictions given by the new models are betterthan those obtained from the old models. The only datapoint that does not follow the models is the dwarf compan-ion object surrounding POX 4. However, as we explainedin Paper I, the values of the W (H α ) have been taken fromimages obtained by M´endez & Esteban (1999) and seem tobe somewhat overestimated.Although absorption features have been only detectedin some of the object, all of them possess an old stellarpopulation underlying the bursts, as we commented inthe analysis of the optical/ NIR colors in Paper I. Thisfact is evident from the values of the W abs derived fromour spectra using the H i
Balmer lines, that seems toincrease with increasing metallicity. A very powerfulmethod to constraint the ages of the stellar populationswithin a starburst galaxy is the analysis of its spectralenergy distributions (SED), although have the degeneracyproblem between the interstellar extinction and the ageof the old stellar population. In some objects, we havechecked the results given by this method considering ourestimation of the reddening contribution derived fromthe Balmer decrement, as we previously did in our anal- ysis of the stellar populations in IRAS 08339+6517(L´opez-S´anchez, Esteban & Garc´ıa-Rojas 2006).We have made use of the PEGASE.2 code(Fioc & Rocca-Volmerange 1997) to produce a gridof theoretical SEDs for an instantaneous burst of starformation and ages between 0 and 10 Gyr, assuming a Z ⊙ metallicity and a Salpeter IMF with lower and upper masslimits of 0.1 M ⊙ and 120 M ⊙ . Although the grid includethe ionized gas emission, we have neglected it becauseits contribution to the continuum is almost irrelevant. InFigure 40 we show the extinction-corrected spectra of theregion A in Haro 15 ( left ) and the center of POX 4 ( right )and two synthetic continuum spectral energy distributionsassuming young ( blue line ) and old ( red line ) populations.We considered the ages derived from W (H α ) representativefor the young population ages (4.5 Myr in both cases)and the ages of the underlying component estimated fromthe optical/ NIR colors (500 Myr in Haro 15, 250 Myr inPOX 4). As we expected, none of the individual syntheticspectra fitted our observed SED. We then constructed amodel than combines both young and old models. For bothcases, the best fits are found when 15% of the 4.5 Myrmodel and 85% of the old model are considered. As wesee, this combined model is in very good agreement withthe shape of our derreddened spectrum. We concludethat, although the star formation activity is very intensein these starbursts, an important underlying old stellarpopulation is usually found in the galaxies, indicatingprevious star-forming phenomena and ruling out thehypothesis that some of them are pristine dwarf galaxies.Our study of the kinematics of the ionized gas andthe morphology and environment included in this and pre-vious papers of our group has revealed that 14 up to20 of the analyzed galaxies show rather clear kinemati-cal and/or morphological evidences of interaction or merg-ing. The morphological evidences were presented and dis-cussed in Paper I. The kinematical evidences presentedhere are of diverse nature: presence of objects with veloci-ties decoupled from the main rotation pattern (Mkn 1087,Haro 15), sinusoidal velocity patterns that suggest a merg-ing process (HCG 31 AC, Mkn 1199, IRAS 08208+2816,SBS 0926+606 A, III Zw 107,
Object 1 in Tol 1457-262),reverses in the velocity distribution (Tol 9, Arp 252), in-dications of tidal streaming (HCG 31, IRAS 08208+2816,SBS 1319+579, Tol 9) or the presence of TDG candi-dates (HCG 31 F1 and F2, Mkn 1087, IRAS 08339+6517,POX 4, Tol 1457-262). The interaction could be betweena spiral –or, in general, a non-dwarf– galaxy (HCG 31,IRAS 08208+2816, Tol 1457-262, III Zw 107 and Arp 252),between and spiral or non-dwarf galaxy and a dwarfone (Mkn 1087, Haro 15, Mkn 1199, IRAS 08339+6517,Tol 9), between two dwarf galaxies (POX 4, SBS 0926+606,SBS 1319+579). In the case of NGC 5253, we have a dwarfstarburst galaxy that has suffered a possible interactionwith a galaxy in the M 83 subgroup or with the spiralgalaxy M 83 itself (L´opez-S´anchez et al. 2008). These re-sults reinforce the hypothesis that interaction with or be-tween dwarf objects is an important mechanism to triggerthe massive star formation in this kind of starbursts. Theseneighboring interacting dwarf or low-luminosity objects areonly detected when a systematic and detailed analysis ofthe morphology, environment, chemical composition andkinematics of the objects are carried out. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 31
Fig. 40.
Spectra of region A in Haro 15 ( left ) and the center of POX 4 ( right ) compared with synthetic continuum spectralenergy distributions obtained using the PEGASE.2 (Fioc & Rocca-Volmerange 1997) code. The gray/green continuous line is theextinction-corrected spectrum, the upper continuous line corresponds to a model with an age of 4.5 Myr (young population model),whereas the lower continuous line is a 500 Myr (for Haro 15 A) or a 250 Myr (for POX 4 C) model (old population model). Theshape of our observed derredened spectra fit in both cases with a model with a contribution of 15% for the young population and85% for the old population (continuous black line over the galaxy spectrum).
Finally, our detailed spectroscopical analysis has pro-vided with some clear evidences of chemical differenceswithin the objects or of interacting objects of differ-ent metallicities. For example, Mkn 1087, Haro 15 andMkn 1199 are in clear interaction with dwarf galaxies withlower O/H and N/O ratios. NGC 5253, IRAS 08208+2816,and Tol 1457-262 show zones with different chemical com-positions. In the case of NGC 5253 this is produced bylocalized pollution by massive stars, but in the cases ofIRAS 08208+2816 and Tol 1457-262 the different chemi-cal composition seem to be because the regions correspondto different galaxies in interaction. Apart from NGC 5253,two more galaxies: IRAS 08208+2816 and UM 420, show alocalized high N/O that could be a signature of contami-nation by WR winds.
Acknowledgements. ´A.R. L-S thanks C.E. (his PhD supervisor) forall the help and very valuable explanations, talks and discussionsalong these years. He also acknowledges Jorge Garc´ıa-Rojas, SergioSim´on-D´ıaz and Jos´e Caballero for their help and friendship duringhis PhD, extending this acknowledge to all people at Instituto deAstrof´ısica de Canarias (Spain). The authors thank B¨arbel Koribalski(CSIRO/ATNF) for her help analyzing HIPASS data. ´A.R. L-S. deeply thanks to Universidad de La Laguna (Tenerife, Spain) for force him totranslate his PhD thesis from English to Spanish; he had to translateit from Spanish to English to complete this publication. This workhas been partially funded by the Spanish Ministerio de Ciencia yTecnolog´ıa (MCyT) under project AYA2004-07466. This research hasmade use of the NASA/IPAC Extragalactic Database (NED) whichis operated by the Jet Propulsion Laboratory, California Institute ofTechnology, under contract with the National Aeronautics and SpaceAdministration.
References
Asplund, M., Grevesse, N. & Sauval, A. J. 2005, in ASP Conf.Ser. 335,
Cosmic Abundances as Records of Stellar Evolution andNucleosynthesis , ed. F.N. Bash & T.G. Barnes (San Francisco: ASP),25Barnes, D.G. et al. 2001, MNRAS, 322, 486Benjamin, R.A., Skillman, E.D. & Smits, D.P. 2002, ApJ,, 569, 288Bettoni, D., Galleta, G. & Garc´ıa-Murillo, S. 2003, A&A, 405, 5Brinchmann, J., Kunth, D., & Durret, F. 2008, A&A, 485, 657Buckalew, B.A., Kobulnicky, H.A. & Dufour, R.J. 2005, ApJS, 157, 30Campbell, A.W., Terlevich, R., & Melnick, J. 1986, MNRAS, 223, 811 Cannon, J.M., Skillman, E.D., Kunth, D., Leitherer, C., Mas-Hesse, M.,¨Ostlin, M. & Petrosian, A. 2004, ApJ, 608, 768Cardelli, J. A., Clayton, G. C. & Mathis J. S. 1989, ApJ, 345, 245Conti, P.S., 1976, MSRSL, 9, 193Conti, P.S., 1991, ApJ, 377, 115Contini, T. 1996,
Liege International Astrophysical Colloquia 33: Wolf-Rayet stars in the framework of stellar evolution , Liege: Universitede Liege, Institut d’Astrophysique, Edited by J.M. Vreux, A. Detal, D.Fraipont-Caro, E. Gosset, and G. Rauw, p.619Crowther, P.A. 2007, ARAA, 45, 177Dahlem, M., Ehle, M., Ryder, S. D., Vlaji´c, M. & Haynes, R. F. 2005,A&A, 432, 475Davoust, E. & Contini, T. 2004, A&A, 416, 515Denicol´o, G., Terlevich, R. & Terlevich, E. 2002, MNRAS, 330, 69Dopita, M.A., Kewley, L. J., Heisler, C.A. & Sutherland, R.S. 2000, ApJ,542, 224Erb, D.K., Shapley, A.E., Steidel, C.C., Pettini, M. et al. 2003, ApJ, 591,101Esteban, C., Peimbert, M., Garc´ıa-Rojas, J., Ruiz, M. T., Peimbert, A. &Rodr´ıguez, M., 2004, MNRAS, 355, 229Fernandes, I.F., de Carvalho, R., Contini, T. & Gal, R.R. 2004, MNRAS355, 728Fioc, M. & Rocca-Volmerange, B. 1997, A&A 326, 950French, H. B. 1980, ApJ, 240, 41Gallego, J.;,Zamorano, J., Rego, M. & Vitores, A.G. 1997, ApJ, 475, 502Garc´ıa-Rojas, J., Esteban, C., Peimbert, M., Rodr´ıguez, M., Ruiz, M. T.,& Peimbert, A. 2004, ApJS, 153, 501Garnett, D.R. 1992, AJ, 103, 1330Garnett, D.R., Kennicutt, R. C.Jr., Chu, Y.-H. & Skillman E. D. 1991,ApJ, 373, 458Garnett, D.R. 2003, lectures on
Cosmochemistry: The melting pot ofthe elements . XIII Canary Islands Winter School of Astrophysics,Puerto de la Cruz, Tenerife, Spain, November 19-30, 2001, edited by C.Esteban, R. J. Garc´ıa L´opez, A. Herrero, F. S´anchez. Cambridge con-temporary astrophysics. Cambridge, UK: Cambridge University Press,ISBN 0-521-82768-X, 2004, p. 171Gordon, D. & Gottesman, S.T. 1981, AJ, 86, 161Guseva, N., Izotov, Y. I. & Thuan, T.X. 2000, ApJ, 531, 776Guseva, N.G., Izotov, Y.I., Papaderos, P., Chaffee, F.H., Foltz, C.B.,Green, R.F., Thuan, T.X., Fricke, K.L. & Noeske, K.G. 2001, A&A,378, 756Guseva, N.G., Papaderos, P., Izotov, Y.I., Green, R.F., Fricke, K.J.,Thuan, T.X. & Noeske, K.G. 2003, A&A, 407, 105Hirashita, H., Inoue, A.K., Kamaya, H. & Shibai, H. 2001, A&A, 366, 83Huang, J.H., Gu, Q.S., Ji,L., Li, W.D., Wei, J.Y. & Zheng, W. 1999, ApJ513, 215Huchtmeier, W.K., Sage, L.J. & Henkel, C. 1995, A&A, 300, 675Huchtmeier, W. K., Krishna, G. & Petrosian, A. 2005, A&A, 434, 887Huchtmeier, W. K., Petrosian, A., Krishna, G. & Kunth, D. 2007, A&A,462, 919Hunter, D.A. & Gallagher, J.S., 1985, AJ, 90, 1457Izotov, Y.I., Guseva, N.G., Lipovetskii, V.A., Kniazev, A.Y. & Stepanian,J.A. 1991, A&A, 247, 303Izotov, Y.I., Thuan, T.X., & Lipovetski, 1994, ApJ, 435, 6472 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
Cosmochemistry: The melting pot ofthe elements . XIII Canary Islands Winter School of Astrophysics,Puerto de la Cruz, Tenerife, Spain, November 19-30, 2001, edited by C.Esteban, R. J. Garc´ıa L´opez, A. Herrero, F. S´anchez. Cambridge con-temporary astrophysics. Cambridge, UK: Cambridge University Press,ISBN 0-521-82768-X, 2004, p. 171Gordon, D. & Gottesman, S.T. 1981, AJ, 86, 161Guseva, N., Izotov, Y. I. & Thuan, T.X. 2000, ApJ, 531, 776Guseva, N.G., Izotov, Y.I., Papaderos, P., Chaffee, F.H., Foltz, C.B.,Green, R.F., Thuan, T.X., Fricke, K.L. & Noeske, K.G. 2001, A&A,378, 756Guseva, N.G., Papaderos, P., Izotov, Y.I., Green, R.F., Fricke, K.J.,Thuan, T.X. & Noeske, K.G. 2003, A&A, 407, 105Hirashita, H., Inoue, A.K., Kamaya, H. & Shibai, H. 2001, A&A, 366, 83Huang, J.H., Gu, Q.S., Ji,L., Li, W.D., Wei, J.Y. & Zheng, W. 1999, ApJ513, 215Huchtmeier, W.K., Sage, L.J. & Henkel, C. 1995, A&A, 300, 675Huchtmeier, W. K., Krishna, G. & Petrosian, A. 2005, A&A, 434, 887Huchtmeier, W. K., Petrosian, A., Krishna, G. & Kunth, D. 2007, A&A,462, 919Hunter, D.A. & Gallagher, J.S., 1985, AJ, 90, 1457Izotov, Y.I., Guseva, N.G., Lipovetskii, V.A., Kniazev, A.Y. & Stepanian,J.A. 1991, A&A, 247, 303Izotov, Y.I., Thuan, T.X., & Lipovetski, 1994, ApJ, 435, 6472 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
Izotov, Y.I., Thuan, T.X., & Lipovetski, 1997, ApJS, 108, 1Izotov, Y.I. & Thuan, T.X., 1998, ApJ, 500, 188Izotov, Y.I. & Thuan, T.X. 1999, ApJ, 511, 639Izotov, Y.I. & Thuan, T.X. 2004, ApJ, 616, 768Kauffmann, G. & White, S.D.M. 1993, MNRAS, 261, 921Kewley, L.J., Dopita, M.A., Sutherland, R.S., Heisler, C.A. & Trevena, J.2001, ApJS, 556, 121Kobulnicky, H.A, & Skillman, E.D. 1996, ApJ, 471, 211Kobulnicky, H.A., Skillman, E.D., Roy, J.-R., Walsh, J.R. & Rosa, M.R.,1997, ApJ, 277, 679Kobulnicky, H.A, & Skillman, E.D. 2008, AJ, 135, 527Kong, X., Cheng, F. Z., Weiss, A. & Charlot, S. 2002, A&A, 396, 503Koribalski, B.S. 2008, proceedings of
Galaxies in the Local Volume con-ference, Springer, p.41Kovo, O. & Contini, T. 1999, in
Wolf-Rayet Phenomena in MassiveStars and Starburst Galaxies. Proceedings of the 193rd symposiumof the International Astronomical Union held in Puerto Vallarta,Mexico, 3-7 November 1998 . Edited by Karel A. van der Hucht,Gloria Koenigsberger, and Philippe R. J. Eenens. San Francisco, Calif.: Astronomical Society of the Pacific, p.604Kniazev, A.Y., Pustilnik, S.A., Grebel, E.K., Lee., H. & Pramsku, A.G.2004, ApJS, 153, 429Kunth, D. & Joubert, M. 1985, A&A, 142, 411Kunth, D. & Schild, H. 1986, A&A, 169, 71Lauberts, A. & Valentijn, E.A. 1989,
The surface photometry catalogue ofthe ESO-Uppsala galaxies , Garching: European Southern Observatory(ESO)Lee, J.C., Salzer J.J. & Melbourne, J. 2004, ApJ, 616, 752LLeitherer, C., Schaerer, D., Goldader, J.D., Gonz´alez-Delgado, R.M.,Robert, C., Kune, D.F., de Mello, D.F., Devost, D. & Heckman, T.M.1999, ApJS, 123, 3 (
STARBURST 99 )L´opez-S´anchez, ´A.R. 2006, PhD Thesis, Universidad de la Laguna(Tenerife, Spain)L´opez-S´anchez, ´A.R., Esteban, C. & Rodr´ıguez, M. 2004a, ApJS, 153, 243L´opez-S´anchez, ´A.R., Esteban, C. & Rodr´ıguez, M. 2004b, A&A 428,445L´opez-S´anchez, ´A.R., Esteban, C. & Garc´ıa-Rojas, J. 2006, A&A, 449, 997L´opez-S´anchez, ´A.R., Esteban, C., Garc´ıa-Rojas, J., Peimbert, M. &Rodr´ıguez, M. 2007, ApJ, 656, 168L´opez-S´anchez, ´A.R., Koribalski, B., Esteban, C. & Garc´ıa-Rojas, J. 2008,proceedings of
Galaxies in the Local Volume , Springer, p.53L´opez-S´anchez, ´A.R. & Esteban, C. 2008, A&A, 491, 131, Paper IL´opez-S´anchez, ´A.R. et al. 2010, in prepMaeder, A. 1990, A&AS, 84, 139Maeder, A. 1991, A&A, 242, 93Masegosa, M., Moles, M. & del Olmo, A. 1991, A&A, 244, 273Masegosa, M., Moles, M. & Campos-Aguilar, A. 1994, ApJ, 420, 576Massey P., Strobel K., Barnes J.V. & Anderson E. 1988, ApJ 328, 315Mazzarella, J.M., Bothun, G.D. & Boroson, T.A. 1991, AJ, 101, 2034Mazzarrella, J.M. & Boronson, T.A. 1993, ApJS, 85, 27Melbourne, J. & Salzer, J.J. 2002, AJ, 123, 2302Melbourne, J., Phillips, A., Salzer, J.J., Gronwall, C. & Sarajedini, V.L.2004, AJ, 127, 686M´endez, D.I. & Esteban, C. 1997, ApJ, 488, 652M´endez, D.I. & Esteban, C. 1999, AJ 118, 2733Meynet G. & Maeder A. 2005. A&A, 429, 581Paturel, G., Theureau, G., Bottinelli, L., Gouguenheim, L., Coudreau-Durand, N., Hallet, N. & Petit, C. 2003, A&A, 412, 57Peimbert, M. & Costero, R. 1969, Bol. Obs. Tonantzintla y Tacubaya, 5,3Penston, M.V., Fosbury, R.A.E., Ward, M.J. & Wilson, A.S. 1977,MNRAS, 180, 19Pe˜na, M., Ruiz, M.T. & Maza, J. 1991, A&A, 251, 417P´erez-Montero, E. & D´ıaz, A. I. 2003, MNRAS, 346, 105Perryman, M.A.C., Longair, M.S., Allington-Smith, J.R. & Fielden, J.1982, MNRAS, 201, 957Pettini, M. & Pagel, B.E.J. 2004, MNRAS, 348, 59Pilyugin, L.S. 2001a, A&A, 369, 594Pilyugin, L.S. 2001b, A&A, 374, 412Pindao, M. 1999, in
IAU Symp. 193: Wolf-Rayet Phenomena in MassiveStars and Starburst Galaxies , 193, 614Pindao, M., Schaerer, D., Gonz´alez-Delgado, R.M. & Stasi´nska, G. 2002,A&A 394, 443Piovan, L., Tantalo, R. & Chiosi, C. 2006, MNRAS, 366, 923Pustilnik, S. A., Martin, J.-M., Huchtmeier, W. K., Brosch, N., Lipovetsky,V. A. & Richter, G. M. 2002, A&A, 389, 405Pustilnik, S., Kniazev, A., Pramskij, A., Izotov, Y., Foltz, C., Brosch, N.,Martin, J.-M. & Ugryumov, A. 2004, A&A 419, 469Quinet, P. 1996, A&AS, 116, 573Richer, M.G., Georgiev, L., Rosado, M., Bullejos, A., Valdez-Guti´errez,M. & Dultzin-Hacyan, D. 2003, A&A, 397, 99Rodr´ıguez, M. & Rubin, R.H. 2005, ApJ, 626, 900Salzer, J. J., Rosenberg, J. L., Weisstein, E. W., Mazzarella, J. M. &Bothun, G. D. 2002, AJ, 124, 191Sargent, W.L.W. 1970, ApJ, 160, 405Schaerer, D., Contini, T. & Pindao, M. 1999, A&AS 136, 35Schaerer, D., Guseva, N G., Izotov, Yu.I. & Thuan, T.X. 2000, A&A 362,53 Schlegel, D.J., Finkbeiner, D.P. & Davis, M. 1998, ApJ, 500, 525Shaw, R.A. & Dufour, R.J. 1995, PASP, 107, 896Shi, F., Kong, X., Li, C. & Cheng, F. Z. 2005, A&A, 437, 849Skillman, E.D., Cˆot´e, S. & Miller, B.W. 2003, AJ 125, 593Smith, L.F., Shara, M.M. & Moffat, A.F.J. 1996, MNRAS, 281, 163Springel, V., White, S. et al. 2005, Nature, 435, 629Stasi´nska, G. 1978, A&A, 66, 257Storey, P.J. & Hummer, D.G. 1995, MNRAS 272, 41Terlevich, R., Melnick, J., Masegosa, J., Moles, M. & Copetti, M.V.F.1991, A&AS, 91, 285Thuan, T.X. & Martin, G.E. 1981, ApJ, 247, 823Thuan, T.X., Izotov, Y.I., & Lipovetsky, V.A. 1995, ApJ, 445, 108Thuan, T.X., Izotov Y.I. & Foltz, C.B. 1999, ApJ, 525, 105Thuan, T.X., Lipovetski, V.A., Martin, J.-M. & Pustilnik, S.A. 1999A&ASS, 139, 1Vacca, W.D. & Conti, P.S., 1992, ApJ, 401, 543Vaceli, M.S., Viegas, S.M., Gruenwald, R. & De Souza, R.E. 1997, AJ,114, 1345van Zee, L., Salzer, J.J. & Haynes, M.P. 1998, ApJ, 497, 1V´azquez, G.A. & Leitherer, C. 2005, ApJ, 621, 695Veilleux, S. & Osterbrock, D.E. 1987, ApJS, 63 295Westera, P., Cuisinier, F., Telles, E. & Kehrig, C. 2004, A&A, 423, 133Winkler, H. 1988, MNRAS, 234, 703Woosley S. E. & Weaver, T.A. 1995, ApJS, 101, 181Woosley, S.E. & Bloom, J.S. 2006. ARAA, 44, 507Zasov, A. V., Kniazev, A. Y., Pustilnik, S. A., Pramsky, A. G., Burenkov,A. N., Ugryumov, A. V. & Martin, J.-M. 2000, A&AS, 144, 429Zhang, H. L. 1996, A&AS, 119, 523Zhang, W., Kong, X., Li, C., Zhou, H.-Y., Cheng, F.-Z. 2007, ApJ, 655,851
List of Objects ‘NGC 1741’ on page 6‘Mkn 1087’ on page 6‘Haro 15’ on page 6‘Mkn 1199’ on page 8‘Mkn 5’ on page 9‘IRAS 08208+2816’ on page 11‘IRAS 08339+6517’ on page 12‘POX 4’ on page 12‘UM 420’ on page 14‘SBS 0926+606’ on page 15‘SBS 0948+532’ on page 16‘SBS 1054+365’ on page 17‘SBS 1211+540’ on page 18‘SBS 1319+579’ on page 19‘SBS 1415+437’ on page 20‘III Zw 107’ on page 21‘Tol 9’ on page 22‘Tol 1457-262’ on page 24‘Arp 252’ on page 25‘NGC 5253’ on page 26
Appendix A: Tables ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 33
Table 2.
Comparison between the oxygen abundance of theregions analyzed here and their previous estimations compiledfrom the literature. The second column indicates if T e was com-puted using the direct method (D) or via empirical calibrations(EC) in this work. T e This work Previous W. RefHCG 31 AC D 8.22 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± a D 8.05 ± ± ± a D 8.12 ± ± a D 8.15 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± a We follow ITL97 names in SBS 1319+579. Notice, however, thanIT99 named region A to SBS 1319+579 C and region B toSBS 1319+579 A; they did not consider SBS 1319+579 B because ofits higher uncertainties. REFERENCES: G97: Gallego et al. (1997);GIT00: Guseva et al. (2000); G03: Guseva et al. (2003); IT98:Izotov & Thuan (1998); IT99: Izotov & Thuan (1999); KS96:Kobulnicky & Skillman (1996); K97: Kobulnicky et al. (1997); KJ85:Kunth & Joubert (1985); KS86: Kunth & Schild (1986); M91;Masegosa et al. (1991); M94: Masegosa et al. (1994); R03:Richer et al. (2003); S05: Shi et al. (2005); VC92: Vacca & Conti(1992).
Table A.1.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in Haro 15. Line f ( λ ) C A B D3705.04 He I 0.260 1.05: ... ... ...3728.00 [O II] 0.256 294 ±
19 113 ±
21 402 ±
106 336 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
39 100 ± ± ±
11 81 ±
30 132 ± ±
13 648 ±
36 232 ±
63 383 ± ± ± ± ± ± ± ± ± ± ±
19 202 ±
46 284 ±
75 288 ± ± ± ± ± ± ± ± × × × × b (arcsec) 0 13 23 12 F (H β ) a ± ± ± ± C (H β ) 0.11 ± ± ± ± W abs (˚A) 2.4 ± ± ± − W (H α ) (˚A) 75.2 ± ± ± ± − W (H β ) (˚A) 16.4 ± ± ± ± − W (H γ ) (˚A) 5.5 ± ± ± − W ([O III]) 5007 (˚A) 29.4 ± ± ± ± a In units of 10 − erg s − cm − and not corrected for extinction. b Relative distance with respect to the center of Haro 15. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 35
Table A.2.
Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Haro 15.
Region C A B D T e (O III) (K) 9500 ± a ±
700 11500 ± a ± a T e (O II) (K) 9600 ±
600 12000 ±
500 11000 ±
700 11260 ± n e (cm − ) 100 100 100 10012+log(O + /H + ) 8.16 ± ± ± ± ++ /H + ) 7.94 ± ± ± ± ± ± ± ± ++ /O + ) − ± ± − ± − ± + /H + ) 7.13 ± ± ± ± ± ± ± ± − ± − ± − ± − ± + /H + ) 6.05 ± ± ± ± ++ /H + ) 6.52 ± ± ± ± − ± − ± ++ /H + ) 7.29 ± ± ± ± ± ± − ± − ± − ± +3 /H + ) ... 4.92 ± ++ /H + ) ... 4.26 ± ++ /H + ) 5.2: 5.5: ... ...12+log(Fe/H) 6.2: 6.5: ... ...log(Fe/O) − − + /H + ) 10.97 ± ± ± − − ± − − a Estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± Table A.3.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in Mkn 1199. Line f ( λ ) C NE A B D3728.00 [O II] 0.256 124.6 ± ±
20 204: 145: 192 ± ± ± ± ± ± ± ± ± ± ± ± ±
11 100: 100: 100 ± ± ± ± ±
15 62: 85: 60 ± ± ± ± ± ± ± ± ± ±
17 290 ±
25 294: 296: 298 ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × × b (arcsec) 0 26 18 14 8.4 F (H β ) a ± ± ± ± ± C (H β ) 0.30 ± ± ± ± ± W abs (˚A) 1.8 ± ± ± ± − W (H α ) (˚A) 129.1 ± ±
10 21 ± ± ± − W (H β ) (˚A) 21.4 ± ± ± ± ± − W (H γ ) (˚A) 6.7 ± ± ± ± − W ([O III]) 5007 (˚A) 6.8 ± ± ± ± ± a In units of 10 − erg s − cm − and not corrected for extinction. b Relative distance with respect to the center of Mkn 1199. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 37
Table A.4.
Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Mkn 1199.
Region Center NE a A a B a D a T e (O III) (K) 5400 ±
700 8450 ±
800 6950 ±
800 6300 ±
800 6750 ± T e (O II) (K) 6800 ± b ±
600 7850 ±
600 7400 ±
600 7700 ± n e (cm − ) 300 ±
100 100 100 100 10012+log(O + /H + ) 8.59 ± ± ± ± ± ++ /H + ) 8.24 ± ± ± ± ± ± ± ± ± ± ++ /O + ) − ± − ± − ± − ± − ± + /H + ) 7.98 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 6.70 ± ± ± ± ± ++ /H + ) 7.05 ± ± ± ± − ± − ± ++ /H + ) 7.65 ± ± ± ± − ± − ± +3 /H + ) 6.07 ± ± ++ /H + ) 6.75 ± ± − ± + /H + ) 10.79 ± b +0.09 ± − − − a Electron temperatures estimated using empirical relations. b Derived from [N ii ] and [O ii ] ratios, see § c [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± Table A.5.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in Mkn 5. Region A was observedusing three slit positions with a PA of 0 ◦ (INT-1), 349 ◦ = − ◦ (INT-2) and 354 ◦ = − ◦ (WHT). Line f ( λ ) A-INT-1 A-INT-2 A-WHT B3666.10 H I 0.267 1.99 ± ± ±
12 213 ±
12 ... 252:3750.15 H I 0.253 1.79 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
22 430 ±
21 374 ±
19 214 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
16 284 ±
15 284 ±
14 284 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × b (arcsec) 0 0 0 16 F (H β ) a ± ± ± ± C (H β ) 0.36 ± ± ± ± W abs (˚A) 1.1 ± ± ± − W (H α ) (˚A) 449 ±
26 435 ±
23 678 ±
35 43 ± − W (H β ) (˚A) 75 ± ± ± − W (H γ ) (˚A) 43 ± ± ± − W ([O III]) 5007 (˚A) 320 ±
17 360 ±
18 530 ±
28 33 ± a In units of 10 − erg s − cm − and not corrected for extinction. b Relative distance with respect to the center of Mkn 5. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 39
Table A.6.
Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Mkn 5.
Region A-INT-1 A-INT-2 A-WHT B a T e (O III) (K) 12400 ±
700 12450 ±
650 12700 ±
600 13250 ± T e (O II) (K) 11700 ±
500 11700 ±
450 11900 ±
400 12300 ± n e (cm − ) ≤ ≤ ≤ ≤ + /H + ) 7.62 ± ± ± ± ++ /H + ) 7.88 ± ± ± ± ± ± ± ± ++ /O + ) 0.25 ± ± ± − ± + /H + ) 6.27 ± ± ± ± ± ± ± ± − ± − ± − ± − ± + /H + ) 5.78 ± ± ± ± ++ /H + ) 6.37 ± ± ± ± ± ± − ± − ± − ± ++ /H + ) 7.02 ± ± ± ± − ± − ± ++ /H + ) 5.69 ± ± ± +3 /H + ) ... 4.69 ± ± ± ± − ± − ± ++ /H + ) 5.73 ± ± − ± − − + /H + ) 10.96 ± ± ± b − ± − ± − ± − a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± Table A.7.
Dereddened line intensity ratios with respect to I (H β )=100 for regions analyzed in IRAS 08208+2816. The slitpositions we used for each knot are: PA 345 ◦ for C, PA 355 ◦ for ◦ for Line f ( λ ) C ± ±
20 251 ±
24 164 ±
22 324 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
11 100 ±
12 100 ± ± ± ±
10 44.7 ± ± ±
24 228 ±
14 305 ±
24 89 ±
13 151 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
15 283 ±
20 281 ±
25 288 ±
33 286 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × × b - 12.6 6.8 8.8 16 F (H β ) a ± ± ± ± ± C (H β ) 0.11 ± ± ± ± ± W abs (˚A) 3.2 ± ± ± ± ± − W (H α ) (˚A) 331 ±
18 202 ±
15 346 ±
32 89 ±
10 98 ± − W (H β ) (˚A) 80 ± ± ± ± ± − W (H γ ) (˚A) 30 ± ± ± ± ± − W ([O III]) (˚A) 370 ±
19 130 ± ±
16 18 ± ± a In units of 10 − erg s − cm − and not corrected for extinction. b Relative distance with respect to the center of IRAS 08208+2816. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 41
Table A.8.
Physical conditions and chemical abundances of the ionized gas of the regions analyzed in IRAS 08208+2816.
Region C a a a a T e (O III) (K) 10100 ±
700 9400 ±
900 9600 ± ± ± T e (O II) (K) 10100 ±
500 9600 ±
700 9700 ±
800 7700 ±
800 9650 ± n e (cm − ) < < < < < + /H + ) 7.77 ± ± ± ± ± ++ /H + ) 8.20 ± ± ± ± ± c ± d ± ± ± ± ± ++ /O + ) 0.43 ± ± ± − ± − ± + /H + ) 6.88 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.94 ± ± ± ± ± ++ /H + ) 6.52 ± ± ± ± − ± − ± ++ /H + ) 7.52 ± ± ± ± ± ± ± ± − ± − ± − ± − ± +2 /H + ) 5.78 ± ± ± ± − ± − ± ++ /H + ) 5.90 ± ± − ± + /H + ) 10.97 ± ± ± ± b − ± − − − − a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± c Oxygen abundance computed using the direct method. d Oxygen abundance computed using the empirical calibrations given by Pilyugin (2001a,b).2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
800 9650 ± n e (cm − ) < < < < < + /H + ) 7.77 ± ± ± ± ± ++ /H + ) 8.20 ± ± ± ± ± c ± d ± ± ± ± ± ++ /O + ) 0.43 ± ± ± − ± − ± + /H + ) 6.88 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.94 ± ± ± ± ± ++ /H + ) 6.52 ± ± ± ± − ± − ± ++ /H + ) 7.52 ± ± ± ± ± ± ± ± − ± − ± − ± − ± +2 /H + ) 5.78 ± ± ± ± − ± − ± ++ /H + ) 5.90 ± ± − ± + /H + ) 10.97 ± ± ± ± b − ± − − − − a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± c Oxygen abundance computed using the direct method. d Oxygen abundance computed using the empirical calibrations given by Pilyugin (2001a,b).2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results
Table A.9.
Dereddened line intensity ratios with respect to I (H β )=100 for regions analyzed in POX 4, UM 420 and SBS0926+606. Line f ( λ ) POX 4 POX 4 Comp UM 420 SBS 0926+606A SBS 0926+606B3679.36 H I 0.265 0.30 ± ± ± ± ± ± ± ± ± ± ± ± ±
35 85.9 ± ± ±
40 140.2 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
16 47.0 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
19 100.0 ± ± ± ± ± ±
10 93 ±
19 106.8 ± ± ± ± ±
32 255 ±
55 312 ±
17 ... ...5015.68 He I -0.038 2.57 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
12 282 ±
50 281 ±
14 286 ±
15 286 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ′′ ) 7.2 × × × × × ′′ ) a – 20.4 – – 74.4 F (H β ) b ± ± ± ± ± C (H β ) 0.08 ± ± ± ± ± W abs (˚A) 2.0 ± ± ± ± ± − W (H α ) (˚A) 1075 ±
48 380 ±
65 1076 ±
55 613 ±
33 92 ± − W (H β ) (˚A) 200 ± ± ±
10 125 ± ± − W (H γ ) (˚A) 71 ± ± ± ± ± − W ([O III]) 5007 (˚A) 1366 ±
60 19 ± ±
32 ... ... a Relative distance of the knot with respect to the main region in the galaxy. b In units of 10 − erg s − cm − and not corrected for extinction. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 43 Table A.10.
Physical conditions and chemical abundances of the ionized gas in POX 4, UM 420 and SBS 0926+606.
Object POX 4 POX 4 Comp a UM 420 SBS 0926+606A SBS 0926+606B a T e (O III) (K) 14000 ±
500 12400 ± ±
600 13600 ±
700 11500 ± T e (O II) (K) 12800 ±
400 11800 ±
600 12200 ±
500 12500 ±
500 11000 ± n e (cm − ) 250 ± <
100 140 ± < < + /H + ) 7.21 ± ± ± ± ± ++ /H + ) 7.96 ± ± ± ± ± ± c ± ± c ± c ± ++ /O + ) 0.74 ± − ± ± ± ± + /H + ) 5.68 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.28 ± ± ± ± ± ++ /H + ) 6.03 ± ± ± ± ± ± − ± − ± − ± ++ /H + ) 7.18 ± ± ± ± − ± − − ± +2 /H + ) ... ... ... 5.52 ± ± +3 /H + ) 5.03 ± ± ± − ± ++ /H + ) 3.83 ± ± ++ /H + ) 5.14 ± ± ± ± ± ± − ± − ± − ± + /H + ) 10.91 ± ± ± b − ± − − ± − ± − a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± c Considering the existence of O +3 because of the detection of He ii λ ∼ Table A.11.
Dereddened line intensity ratios with respect to I (H β )=100 for regions analyzed in SBS 0948+532, SBS 1054+365and SBS 1211+540. Line f ( λ ) SBS 0948+532 SBS 1054+365 SBS 1054+365 b SBS 1211+5403697.15 H I 0.262 0.48: ... ... ...3703.86 H I 0.260 1.96 ± ± ± ± ± ± ±
95 ...3728.82 [O II] 0.256 65.7 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
39 100.0 ± ± ±
13 75 ±
23 163 ± ± ±
28 623 ±
37 183 ±
53 481 ± ± ± ± ± ± ± ± ± ±
14 277 ±
17 279 ±
74 280 ± ± ± ± ± ± ± ± ± ± ± ± ± a (arcsec) 0 0 17.8 0 F (H β ) b ± ± ± ± C (H β ) 0.35 ± ± ± ± W abs (˚A) 0.3 ± ± ± ± − W (H α ) (˚A) 788 ±
43 422 ±
27 32 ± ± − W (H β ) (˚A) 213 ±
11 89 ± ± ± − W (H γ ) (˚A) 57 ± ± ± ± − W ([O III]) 5007 (˚A) 689 ±
34 567 ±
35 12 ± ± a Relative distance of the knot with respect to the main region in the galaxy. b In units of 10 − erg s − cm − and not corrected for extinction. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 45 Table A.12.
Physical conditions and chemical abundances of the ionized gas in SBS 0948+532, SBS 1054+365 and SBS 1211+540.
Object SBS 0948+532 SBS 1054+365 SBS 1054+365b a SBS 1211+540 T e (O III) (K) 13100 ±
600 13700 ±
900 11800 ± ± T e (O II) (K) 12200 ±
400 12600 ±
700 11300 ±
900 15000 ± N e (cm − ) 250 ± <
100 300 ±
200 320 ± + /H + ) 7.33 ± ± ± ± ++ /H + ) 7.94 ± ± ± ± ± ± ± ± ++ /O + ) 0.61 ± ± − ± ± + /H + ) 5.91 ± ± ± ± ± ± ± ± − ± − ± − ± − ± + /H + ) 5.43 ± ± ± ± ++ /H + ) 6.16 ± ± ± ± ± ± − ± − ± − ± ++ /H + ) 7.21 ± ± ± ± ± ± − ± − ± − ± +2 /H + ) ... 5.62 ± +3 /H + ) 4.79 ± ± ± ± − ± ++ /H + ) 3.97 ± ++ /H + ) 5.64 ± ± − ± + /H + ) 10.88 ± ± ± b − ± − ± − − ± a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± Table A.13.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in SBS 1319+579 (regions A, B andC) and SBS 1415+457 (regions C and A). SBS 1319+579 SBS 1415+437Line f ( λ ) A B C C A4340.47 H I 0.127 47.2 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
12 132 ±
10 129.9 ± ± ± ± ±
14 286 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
13 279 ±
19 281 ±
15 274 ±
13 278 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × × a - 10 29 - 6 F (H β ) b ± ± ± ± ± C (H β ) 0.03 ± ± ± ± ± W abs (˚A) 0.0 ± ± ± ± ± − W (H α ) (˚A) 1530 ±
75 162 ±
11 295 ±
23 1300 ±
65 1187 ± − W (H β ) (˚A) 285 ±
14 42 ± ± ±
11 130 ± − W (H γ ) (˚A) 84 ± ± ± ± ± − W ([O III]) 5007 (˚A) ... ... ... 542 ±
26 574 ± a Relative distance of the knot with respect to the main region in the galaxy. b In units of 10 − erg s − cm − and not corrected for extinction. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 47 Table A.14.
Physical conditions and chemical abundances of the ionized gas for the regions analyzed in SBS 1319+579 andSBS 1415+437.
Object SBS 1319+579A SBS 1319+579B SBS 1319+ 579C SBS 1415+ 437C SBS 1415+ 437A T e (O III) (K) 13400 ±
500 11900 ±
800 11500 ±
600 16400 ±
600 15500 ± T e (O II) (K) 12400 ±
400 11300 ±
600 11050 ±
400 14500 ±
400 13850 ± N e (cm − ) < < < < < + /H + ) 7.22 ± ± ± ± ± ++ /H + ) 7.98 ± ± ± ± ± ± ± ± ± b ± ++ /O + ) 0.77 ± ± ± ± ± + /H + ) 5.69 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.29 ± ± ± ± ± ++ /H + ) 6.09 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± +2 /H + ) 5.44 ± ± ± ± ± +3 /H + ) 5.19 ± ± ± ± − ± − ± ++ /H + ) ... ... 5.46: 5.23 ± ± − − ± − + /H + ) 10.94 ± ± ± ± ± a − ± − ± − ± − ± − ± a [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± b Considering the existence of O +3 because of the detection of He ii λ ∼ Table A.15.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in III Zw 107 (regions A, B and C)and Tol 9 (spectra obtained with 2.5m INT and 2.56m NOT). Line f ( λ ) III Zw 107 A III Zw 107 B III Zw 107 C Tol 9 INT Tol 9 NOT3554.42 He I 0.283 ... ... ... 3.8 ± ±
12 306 ±
23 20.32: 142 ±
10 177 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
15 100.0 ± ± ± ± ±
14 78.3 ± ± ±
18 293 ±
19 257 ±
32 236 ±
13 225 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
16 273 ±
18 280 ±
35 267 ±
18 284 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × × a - 7.2 12.4 - - F (H β ) b ± ± ± ± × C (H β ) 0.68 ± ± ± ± ± W abs (˚A) 2.0 ± ± ± ± ± − W (H α ) (˚A) 306 ±
18 76 ± ± ±
12 186 ± − W (H β ) (˚A) 44 ± ± ± ± ± − W (H γ ) (˚A) 16.4 ± ± ± ± − W ([O III]) 5007 (˚A) 172 ± ± ± ± ± a Relative distance of the knot with respect to the main region in the galaxy. b In units of 10 − erg s − cm − and not corrected for extinction. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 49 Table A.16.
Physical conditions and chemical abundances of the ionized gas for the regions analyzed in III Zw 107 and Tol 9.
Object III Zw 107 A III Zw 107 B a III Zw 107 C a Tol 9 INT Tol 9 NOT T e (O III) (K) 10900 ±
900 10400 ± ± ± ± T e (O II) (K) 10500 ±
800 10300 ±
800 10250 ±
800 8300 ±
700 8500 ± N e (cm − ) 200 ± < <
100 180 ±
60 260 ± + /H + ) 7.87 ± ± ± ± ± ++ /H + ) 7.99 ± ± ± ± ± ± ± ± ± ± ++ /O + ) 0.12 ± − ± − ± ± ± + /H + ) 6.70 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.87 ± ± ± ± ± ++ /H + ) 6.23 ± ± ± ± ± ± − ± − ± − ± ++ /H + ) 7.26 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± +2 /H + ) 5.94 ± ± ± ± ± ± ± ± − ± − ± − ± − ± ++ /H + ) 4.32: ... ... 5.31: ...12+log(Fe ++ /H + ) 5.61: ... ... 6.14: ...12+log(Fe/H) 5.92: ... ... 6.51: ...log(Fe/O) − − + /H + ) 10.94 ± ± ± ± b − ± − − − ± − ± a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± Table A.17.
Dereddened line intensity ratios with respect to I (H β )=100 for knots analyzed in Tol 1457-262 (regions A, B andC) and Arp 252 (galaxy A, ESO 566-8, and galaxy B, ESO 566-7). Line f ( λ ) Tol 1457-262A Tol 1457-262B Tol 1457-262C ESO 566-8 ESO 566-73728.00 [O II] 0.256 224 ±
17 187 ±
16 270 ±
25 256 ±
18 280 ± ± ± ± ± ± ± ±
11 4.77: ...3889.05 H I 0.226 12.0 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
10 100.0 ± ± ±
10 190 ±
12 163 ±
11 56.9 ± ± ±
27 522 ±
29 455 ±
28 204 ±
12 75 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
15 286 ±
17 279 ±
18 286 ±
17 287 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± × × × × × a - 9.7 5.5 - 49.7 F (H β ) b ± ± ± ± ± C (H β ) 0.57 ± ± ± ± ± W abs (˚A) 1.4 ± ± ± ± ± − W (H α ) (˚A) 603 ±
32 390 ±
28 342 ±
23 472 ±
29 79 ± − W (H β ) (˚A) 101 ± ± ± ± ± − W (H γ ) (˚A) 31 ± ± ± ± ± − W ([O III]) 5007 (˚A) 560 ±
27 430 ±
25 411 ±
26 197 ±
12 10 ± a Relative distance of the knot with respect to the main region in the galaxy. b In units of 10 − erg s − cm − and not corrected for extinction. ´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results 51 Table A.18.
Physical conditions and chemical abundances of the ionized gas for the regions analyzed in Tol 1457-262 and Arp252 (galaxy A, ESO 566-8, and galaxy B, ESO 566-7).
Object Tol 1457-262A Tol 1457-262B Tol 1457-262C ESO 566-8 ESO 566-7 a T e (O III) (K) 14000 ±
700 15200 ±
900 13400 ± ±
900 7900 ± T e (O II) (K) 12500 ±
600 14200 ±
700 12400 ± ±
800 8500 ± N e (cm − ) 200 ± <
100 200 ±
100 300 ±
100 100 ± + /H + ) 7.59 ± ± ± ± ± ++ /H + ) 7.87 ± ± ± ± ± ± c ± ± ± c ± ++ /O + ) 0.27 ± ± ± − ± − ± + /H + ) 6.02 ± ± ± ± ± ± ± ± ± ± − ± − ± − ± − ± − ± + /H + ) 5.65 ± ± ± ± ± ++ /H + ) 5.95 ± ± ± ± − ± − ± ++ /H + ) 6.99 ± ± ± ± ± ± ± ± − ± − ± − ± − ± +2 /H + ) 5.73 ± ± ± ± ± +3 /H + ) ... ... ... 5.48 ± ± ± ± ± ± − ± − ± − ± − ± − ± ++ /H + ) 5.43: 5.62 ± ± − − ± − + /H + ) 10.99 ± ± ± ± b − ± − ± − ± − ± − a Electron temperatures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) ⊙ , using 12+log(O/H) ⊙ = 8.66 ± c Considering the existence of O +3 because of the detection of He ii λ ∼2 L´opez-S´anchez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results