Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://www.atnf.csiro.au/people/Angel.Lopez-Sanchez/papers/MSFinWRG_II_main_ACCEPTED_26sep09.pdf Äàòà èçìåíåíèÿ: Mon Sep 28 10:26:02 2009 Äàòà èíäåêñèðîâàíèÿ: Wed Nov 25 05:56:16 2009 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: m 106

Astronomy & Astrophysics manuscript no. MSFinWRGIImainACCEPTED26sep09 September 28, 2009

c ESO 2009

Massive star formation in Wolf-Rayet galaxies
I I. Optical sp ectroscopy results
´ Angel R. L´ ez-S´ hez1,2 and C´ op anc esar Esteban2
1 2 3

,3

CSIRO / Australia Telescope National Facility, PO-BOX 76, Epping, NSW 1710, Australia Instituto de Astrof´ isica de Canarias, C/ V´ L´ ia actea S/N, E-38200, La Laguna, Tenerife, Spain Departamento de Astrof´ isica de la Universidad de La Laguna, E-38071, La Laguna, Tenerife, Spain

Received March 12, 2009; Accepted September 26, 2009 ABSTRACT Aims. We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that show the 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-resolution spectroscopy 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 and derive the reddening coefficient, physical conditions and chemical abundances of the ionized gas. We study the kinematics of the ionized gas to check the rotation/turbulence pattern of each system. When possible, tentative estimates of the Keplerian 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 role in their detection. We checked that the ages of the last star-forming burst estimated using optical spectra agree with those derived from H imagery. Our analysis has revealed that a substantial fraction of the galaxies show evidences of perturbed kinematics. With respect to the results found in individual galaxies, we remark the detection of ob jects with different metallicity and decoupled kinematics in Haro 15 and Mkn 1199, the finding of evidences of tidal streams in IRAS 08208+2816, Tol 9 and perhaps in SBS 1319+579, and the development of a merging process in SBS 0926+606 A and in Tol 1457-262. Conclusions. All these results ­in combination with those obtained in Paper I­ reinforce the hypothesis that interactions with or between dwarf ob jects is a very important mechanism in the triggering of massive star formation in starburst galaxies, specially in dwarf ones. It must be highlighted that only deep and very detailed observations ­as these presented in 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: kinematics and dynamics-- stars: Wolf-Rayet

1. Introduction
Wolf-Rayet (WR) stars are the evolved descendants of the most massive, very hot and very luminous (105 to 106 L ) O stars. In the so-called Conti (1976) and Maeder (1990, 1991) scenarios, WR stars are interpreted as central Heburning ob jects that have lost the main part of their H-rich envelope via strong winds. Hence, their surface chemical composition is dominated by He rather than H, along with elements produced by the nuclear nucleosynthesis. WN and WC stars show the products of the CNO cycle (H-burning) and the triple- (He-burning), respectively. The most massive O stars (M 25 M for Z ) became WR stars between 2 and 5 Myr since their birth, spending only some few
´ Send offprint requests to : Angel R. L´ ez-S´ hez, e-mail: op anc Angel.Lopez-Sanchez@csiro.au Based on observations made with NOT (Nordic Optical Telescope), INT (Isaac Newton Telescope) and WHT (William Herschel Telescope) operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway and Sweden (NOT) or the Isaac Newton Group (INT, WHT) in the Spanish Observatorio del Roque de Los Muchachos of the Instituto de Astrof´ isica de Canarias.

hundreds of thousands of years (tW R 5 â 105 yr) in this phase (Meynet & Maeder 2005). A review of the physical properties of WR stars was recently presented by Crowther (2007). The broad emission features that characterized the spectra of WR stars are often observed in extragalactic H ii regions. Actually, the so-called Wolf-Rayet galaxies make up a very inhomogeneous class of star-forming ob jects: giant H ii regions in spiral arms, irregular galaxies, blue compact dwarf galaxies (BCDGs), luminous merging IRAS galaxies, active galactic nuclei (AGNs), Seyfert 2 and low-ionization nuclear emission-line regions (LINERs). All ob jects have in common ongoing or recent star formation which has produced stars massive enough to evolve to the WR stage (Shaerer et al. 1999). There are two important broad features that reveal the presence of WR stars in the integrated spectra of an extragalactic H ii region: 1. A blend of He ii 4686, C iii/C iv 4650 and N iii 4640 emission lines originated in the expanding atmospheres of the most massive stars, the so-called blue WR Bump. This feature is mainly due to the presence of WN stars. The broad, stellar, He ii 4686 is its main fea-


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

ture. Although rarely strong, the narrow, nebular He ii 4686 is usually associated with the presence of these massive stars, but its origin is still somewhat controversial (Garnett et al. 1991; Garnett 2004; Brinchmann, Kunth & Durret 2008). 2. The C iii 5698 and C iv 5808 broad emission lines, sometimes called the red or yellow WR Bump. C iv 5808 is the strongest emission line in WC stars but it is barely seen in WN stars. The red WR bump is rarely detected and it is always weaker than the blue WR bump (Guseva, Izotov & Thuan 2000; Fernandes et al. 2004). Making use of population synthesis models it is possible to determine the age of the bursts, the number of O and WR stars, the WN/WC ratio, the initial mass function (IMF) or the mass of the burst. Therefore, the study of WR galaxies helps to widen our knowledge about both the massive star formation and the evolution of starbursts: they allow to study the early phases of starbursts and are the best direct measure of the upper end of the IMF, a fundamental ingredient for studying unresolved stellar populations, (Schaerer et al. 2000; Guseva, Izotov & Thuan 2000; Pindao et al. 2002; Fernandes et al. 2004; Buckalew, Kobulnicky & Dufour 2005; Zhang et al. 2007), giving key constraints to stellar evolution models. On the other hand, the knowledge of the chemical composition of galaxies, in particular in dwarf galaxies, is vital for understanding their evolution, star formation history, stellar nucleosynthesis, the importance of gas inflow and outflow and the enrichment of the intergalactic medium. Indeed, metallicity is a key ingredient for modelling galaxy properties, such it determines UV, optical and NIR colors at a given age (i.e., Leitherer et al. 1999), nucleosynthetic yields (e.g., Woosley & Weaver 1995), the dust-to-gas ratio (e.g., Hirashita et al 2001), the shape of the interstellar extinction curve (e.g., Piovan et al. 2006), or even the WR properties (Crowther 2007). The most robust method to derive the metallicity in star-forming and starburst galaxies is via the estimation of metal abundances and abundance ratios, in concrete through the determination of the gas-phase oxygen and nitrogen abundances and the nitrogen-to-oxygen ratio. The relationships between current metallicity and other galaxy parameters, such as colors, luminosity, neutral gas content, star formation rate, extinction or total mass, constraint galaxy evolution models and give clues about the current stage of a galaxy. For example, is still debated if massive star formation result in the instantaneous enrichment of the interstellar medium of a dwarf galaxy, or if the bulk of the newly synthesized heavy elements must cool before becoming part of the ISM that eventually will form the next generation of stars. Accurate oxygen abundance measurements of several H ii regions within a dwarf galaxy will increase the understanding of its chemical enrichment and mixing of enriched material. The analysis of the kinematics of the ionized gas will also help to understand the dynamic stage of galaxies and reveal recent interaction features. Furthermore, detailed analysis of starburst galaxies in the nearby Universe are fundamental to interpret the observations of high-z star forming galaxies, such as Lyman Break Galaxies (Erb et al. 2003), as well as quantify the importance of interactions in the triggering of the starformation 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 spectroscopic analysis of a sample of 20 WR galaxies. Our main aim is the study of the formation of massive stars in starburst galaxies, their gas-phase metal abundance and its relationships with other galaxy properties, and the role that the interactions with or between dwarf galaxies and/or low surface brightness ob jects have in the triggering mechanism of the star-formation bursts. In Paper I (L´ ez-S´ hez op anc & Esteban 2008) we exposed the motivation of this work, compiled the list of the analyzed WR galaxies (Table 1 of Paper I) and presented the results of optical/NIR broadband and H photometry. In this second paper we present the results of the analysis of intermediate-resolution long slit spectroscopy of 16 ob jects of our sample of WR galaxies ­the results for the other 4 ob jects have been published separately. In many cases, two or more slit positions have been used in order to analyze the most interesting zones, knots or morphological structures belonging to each galaxy or even surrounding ob jects. In particular, these observations have the following aims: 1. Study the content and spatial location of the WR stars in each galaxy. We examine the spectra for the presence of the He ii 4686 emission line and/or the blue-WR bump as well as for the red-WR bump. The characteristics of the WR population can be derived by comparison with theoretical population synthesis models. 2. Know the physical properties of the ionized gas: excitation mechanism, electron density, high and low ionization electron temperatures and reddening coefficient. 3. Analyze the ionization structure and the chemistry of the gas (abundances of He, O, N, S, Ne, Ar, Fe and Cl) associated with different morphological zones in each galaxy, especially in those areas in which WR features are detected. This analysis is specially relevant in cases of interaction or merging processes because the regions may have different chemical composition and also would allow to discern between the tidal dwarf galaxy (TDG) or pre-existing dwarf galaxy nature of nearby diffuse objects surrounding the main galaxy. 4. Determine the radial velocities of different starformation bursts, galaxies in the same system and/or ob jects in possible interaction. The distance to the main galaxy is also calculated. 5. Study the velocity field via the analysis of positionvelocity diagrams in order to understand the kinematics of the ionized gas associated to different members in the system in order to know their evolution (rotation, interactions features, fusion evidences, movements associated to superwinds...). The Keplerian mass has been estimated in ob jects showing solid-body rotation. 6. Obtain independent estimations of the age of the last star-forming burst via the comparison with stellar population synthesis models. 7. Study the stellar population underlying the bursts using the analysis of absorption lines (i.e. Ca ii H,K, Mg i 5167,5184, Na i 5890,5896, Ca ii triplet). 8. Finally, the spectral energy distribution (Sed) has been analyzed in some cases in order to constrain the properties of the underlying stellar population. This paper mainly presents the analysis of the ionized gas within our WR galaxy sample. In §2 we describe our


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

3

observations, some details of the data reduction processes, and describe some useful relations. In §3 we describe the physical properties (Te , ne , reddening coefficient, excitation mechanism), the chemical abundances and the kinematics of the ionized gas for each galaxy. Finally, the most important results derived from our spectroscopic study, including a comparison with the ages derived from the H photometry and an estimation of the age of the underlying stellar component using the Sed, are summarized in §4. A detailed analysis of the O and WR populations and the comparison with theoretical models is presented in Paper III. The global analysis of our optical/NIR data will be shown in Paper IV. The final paper of the series (Paper V) will compile the properties derived using data from other wavelengths (UV, FIR, radio and X-ray) and complete the global analysis combining all available multiwavelength data of our WR galaxy sample. It is, so far, the most complete and exhaustive data set of this kind of galaxies, involving multiwavelength results and analyzed following the same procedures.

2. Observations and data reduction and analysis
2.1. Spectroscopic observations We obtained intermediate-resolution long slit spectroscopy for all our sample WR galaxies except for NGC 5253, for which high-resolution echelle spectroscopy was taken (see L´ ez-S´ hez et al. 2007 for details). We used three teleop anc scopes to carry out these observations: 2.5m Isaac Newton Telescope (INT), 2.56m Nordical Optical Telescope (NOT), and 4.2m Wil liam Herschel Telescope (WHT), all located at 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 attached at the Cassegrain focus in December 1999. A EEV CCD 2148â4200 pixel array, with a pixel size of 13.5 µm, was used, that corresponds to an spatial resolution of 0.40 pix-1 . The slit was 2.8 long and 1 wide. We used the R400V grating, that has a dispersion of 104.5 ° mm-1 (1.40 ° pix-1 ) and an effective spectral A A resolution of 3.5 ° The spectra cover the wavelength A. range from 3200 to 7700 ° The absolute flux calibraA. tion was achieved by observations of the standard stars Feige 56, Hiltner 600 and Feige 110 (Massey et al. 1988). 2. Observations at the 4.2m WHT. We completed two observation runs in this telescope on December 2000 and December 2002. In both cases, the double-arm ISIS (Intermediate dispersion Spectrograph and Imaging System ) instrument located at the Cassegrain focus of the telescope was used. The dichroic used to separate the blue and red beams was set at 5400 ° The slit was A. 3.7 long and 1 wide. We used different configurations in each observing run: (a) Decemb er 2000: ­ Blue arm : an EEV CCD with a 4096 â 2048 pixels array and 13 µm size was used. The spatial resolution was 0.20 pix-1 . The grating was R600B, giving a dispersion of 33 ° mm-1 A (0.45 ° pix-1 ) and an effective spectral resoluA tion of 1.8 ° The observed spectral range was A. 3600 ­ 5200 ° A. ­ Red arm : we used a TEX CCD with a configuration of 1024â1024 pixels of 24 µm pixel

size, having a spatial resolution of 0.36 pix-1 . The grating R316R, that has a dispersion of 66 ° mm-1 (0.93 ° pix-1 ) and an effective specA A tral resolution of 2.6 ° was used, covering the A, spectral range 5400 ­ 6800 ° A. (b) Decemb er 2002: ­ Blue arm : We used the same CCD that previously indicated but the R1200B grating, that gives a dispersion of 17 ° mm-1 (0.23 ° pix-1 ) A A and an effective spectral resolution of 0.86 ° A. The spectral range was 4450 ­ 5480 ° A. ­ Red arm : A Marconi CCD with 4700â2148 pixels array and 14.5 µm pixel size was used. The spatial resolution was 0.20 pix-1 , hence identical to that provided in the blue arm. We used the R316R grating covering the spectral range 5700 ­ 8600 ° A. The absolute flux calibration was achieved by observations of the Massey et al. (1988) standard stars G191B2B and Feige 34 (December 2000) and Feige 15, Feige 110, Hiltner 600 and Hz44 (December 2002). 3. Observations at the 2.56m NOT. We completed three observation run at this telescope, always using the ALFOSC (Andaluc´ Faint Object Spectrograph and ia Camera ) instrument and a Loral/Lesser CCD detector (2048 â 2048 pixels) with a pixel size of 13.5 µm and spatial resolution of 0.19 pixel-1 . The slit was 6.4 long and 1 wide. We used several configurations: (a) 20 March 2004. We used grism #7 that has a dispersion of 111 ° mm-1 (1.5 ° pix-1 ) and a specA A tral resolution of 7.5 ° covering the spectral range A, 3200 ­ 6800 ° A. (b) 4 April 2005 and 26­27 April 2006. We used two different grisms to obtain the blue and the red ranges of the optical spectrum. Grism #14, which has a dispersion of 104 ° mm-1 (1.4 ° pix-1 ) and a spectral A A resolution of 7.0 ° was used to cover the spectral A, range 3300 ­ 6100 ° This grism has a low efficiency A. for 4000 ° Spectra in the red range were obA. tained using the grism #8, that has a dispersion of 96 ° mm-1 (1.3 ° pix-1 ), a spectral resolution of A A 6.5 ° and covers the spectral range 5800 ­ 8300 ° A A. The spectrophotometric standard star Feige 56 (Massey et al. 1988) was used for flux calibrating all the spectra obtained with this telescope. In all observations, three or four exposures for each slit position were taken to get a good S/N ratio and to remove cosmic rays. Table 1 compiles all the intermediateresolution long-slit spectroscopy observations performed for the 16 WR galaxies included in this paper. 2.2. Reduction of the spectra All the data reduction were completed at the IAC. IRAF1 software was used to reduce the CCD frames (bias correction, flat-fielding, cosmic-ray rejection, wavelength and flux calibration, sky subtraction) and extract the onedimensional spectra. The correction for atmospheric extinction was performed using an average curve for the continuous atmospheric extinction at Roque de los Muchachos
1 IRAF is distributed by NOAO which is operated by AURA Inc., under cooperative agreement with NSF.


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table 1. Journal of the intermediate-resolution long-slit spectroscopy observations, carried out using the 2.5m INT, 2.56m NOT and 4.2m WHT telescopes and with the instrumentation explained in text. In last column, sec z represents the average airmass in which the observations were made.
Galaxy Haro 15 Mkn 1199 Mkn 5 Tel. INT INT INT INT WHT WHT INT INT WHT WHT WHT WHT INT INT INT INT WHT WHT WHT WHT INT WHT WHT WHT WHT WHT WHT INT INT NOT NOT NOT NOT NOT NOT Date 99/12/27 99/12/27 99/12/28 99/12/28 02/12/27 02/12/27 99/12/29 99/12/29 00/12/30 00/12/30 00/12/30 00/12/30 99/12/28 99/12/28 99/12/28 99/12/28 02/12/27 02/12/27 00/12/31 00/12/31 99/12/29 00/12/31 00/12/31 02/12/27 02/12/27 02/12/27 02/12/27 99/12/28 99/12/27 06/04/28 06/04/27 05/04/04 05/04/04 05/04/04 05/04/04 Exp. Time [s] 3â1200 3â1200 3â1200 3â1200 3â700 3â700 3â1200 3â1200 3â1800 3â1800 4â1800 4â1800 3â1200 3â1200 3â1200 3â1200 3â600 3â600 3â1800 3â1800 3â1200 3â1800 3â1800 3â600 3â600 3â600 3â600 3â1200 4â1200 4â1200 3â900 3â900 3â900 3â900 3â900 Spatial R. Grism [ /pix] 0.40 0.40 0.40 0.40 0.20 0.20 0.40 0.40 0.20 0.36 0.20 0.36 0.40 0.40 0.40 0.40 0.20 0.20 0.20 0.36 0.40 0.20 0.36 0.20 0.20 0.20 0.20 0.40 0.40 0.19 0.19 0.19 0.19 0.19 0.19 R400V R400V R400V R400V R1200 R136R R400V R400V R600B R136R R600B R136R R400V R400V R400V R400V R1200 R136R R600B R136R R400V R600B R136R R1200 R136R R1200 R136R R400V R400V g14 g8 g14 g8 g14 g8 P.A. [ ] 41 117 32 53 90 90 0 349 25 25 90 90 90 10 345 355 27 27 114 114 55 138 138 39 39 20 20 0 49 109 109 155 155 342 342 Spectral R. [° A] 3.5 3.5 3.5 3.5 0.86 2.6 3.5 3.5 1.8 2.6 1.8 2.6 3.5 3.5 3.5 3.5 0.86 2.6 0.86 2.6 3.5 1.8 2.6 0.86 2.6 0.86 2.6 3.5 3.5 7.0 6.5 7.0 6.5 7.0 6.5 [° A] 3500-7700 3500-7700 3500-7700 3500-7700 4300-5100 5700-7800 3500-7700 3500-7700 3650-5100 5300-6650 3650-5100 5300-6650 3500-7700 3500-7700 3500-7700 3500-7700 4300-5100 5700-7800 3650-5100 5300-6650 3500-7700 3650-5100 5300-6650 4300-5100 5700-7800 4300-5100 5300-6650 3500-7700 3500-7700 3300-6100 5800-8300 3300-6100 5800-8300 3300-6100 5800-8300 sec z 1.67 1.40 1.02 1.00 1.90 1.90 1.54 1.47 1.60 1.60 1.14 1.14 1.23 1.11 1.01 1.32 1.18 1.18 1.10 1.10 1.15 1.12 1.12 1.48 1.48 1.54 1.54 1.18 1.90 1.85 1.92 1.92 1.75 1.67 1.56

POX 4 UM 420

IRAS 08208+2816

SBS 0926+606A SBS 0948+532 SBS 1054+365 SBS 1211+540 SBS 1319+579 SBS 1415+437 III Zw 107 Tol 9

Tol 1457-262a Arp 252

Observatory. For each two-dimensional spectra several apertures were defined along the spatial direction to extract the final one-dimensional spectra of each galaxy or emission knot. The apertures were centered at the brightest point of each aperture and the width was fixed to obtain a good signal-to-noise spectrum. In case of having the optical spectrum separated in two different wavelength intervals, identical apertures in both spectral ranges were used. 2.3. Analysis of the spectra IRAF software was also used to analyze the onedimensional spectra. Line intensities and equivalent widths were measured by integrating all the flux in the line between two given limits and over a local continuum estimated by eye. In the cases of line blending, a multiple Gaussian profile fit procedure was applied to obtain the line flux of each individual line. We used the standard assumption I (H )=100 to compute the line intensity ratios. The identification and adopted laboratory wavelength of the lines, as well as their errors, were obtained following Garc´ ia-Ro jas et al. (2004); Esteban et al. (2004).

2.3.1. Distance to the galaxies We computed the distance to the galaxies using the brightest emission lines (H and [O iii] 5007) in our optical spectra. We assumed a Hubble flow with H0 =75 km s-1 Mpc-1 , q0 =0.5, and corrected for Galactic Standard of Rest. The distance we derived for each galaxy is listed in Table 1 in Paper 1. All values agree well within the errors with the distances quoted by the NED, except for Tol 9. For this galaxy, we measure a radial velocity of vr = 3441 km s-1 , while previous observations suggested vr = 3190 km s-1 (Lauberts & Valentijn 1989). We consider that our value is more appropriate because the maximum of the H i emission detected in Tol 9 (L´ ez-S´ hez et al. 2010) shows op anc the same radial velocity than that our optical spectrum provides. 2.3.2. Correction for reddening The reddening coefficient, c(H ), was derived from the Balmer decrement. However, in extragalactic ob jects the fluxes of nebular Balmer lines may be affected by absorptions produced by the underlying stellar population (mainly


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

5

B and A stars). We performed an iterative procedure to derive simultaneously the reddening coefficient and the equivalent widths of the absorption in the hydrogen lines, Wabs , to correct the observed line intensities for both effects. We assumed that Wabs is the same for all the Balmer lines and used the relation given by Mazzarrella & Boronson (1993) to the absorption correction, 1 c(H ) = log f ()
I () I (H ) F () F (H )

â 1+ â 1+

Wabs WH Wabs W

,

(1)

for each detected hydrogen Balmer line. In this equation, F () and I () are the observed and the theoretical fluxes (unaffected by reddening or absorption); Wabs , W and WH are the equivalent widths of the underlying stellar absorption, the considered Balmer line and H , respectively, and f () is the reddening curve normalized to H using the Cardelli, Clayton & Mathis (1989) law. We always considered the theoretical ratios between these pairs of H i Balmer lines expected for case B recombination given by Storey & Hummer (1995) assuming the electron temperature and density computed independently for each region. Using three different Balmer lines (i.e., H, H and H ) an unique value for c(H ) and Wabs is computed. However, in the case of using four or more Balmer lines, several solutions are derived, so we considered the values that provide the best match between the corrected and the theoretical Balmer line ratios as representative of the region. An example of this method using 5 Balmer line ratios is shown in Figure 14 when analyzing UM 420. Tables A.1, A.3, A.5, A.7, A.9, A.11, A.13, A.15 and A.17 show the dereddened line intensity ratios and their associated errors for all the regions and galaxies, as well as the adopted f () for each emission line. In these tables, we also include other important quantities as: the size of the extracted aperture, its relative distance to the main region of the galaxy, the observed H flux (uncorrected for extinction), the adopted values of c(H ) and Wabs and the equivalent widths of H, H , H and [O iii] 5007. Colons indicate errors of the order or larger than 40%. 2.3.3. Physical conditions of the ionized gas We studied the physical conditions and chemical abundances of the ionized gas from the 1-D spectra of each galaxy or knot. We used a two-zone approximation to define the temperature structure of the nebulae, assuming the electron temperature, Te , provided by the [O iii] ion, Te (O iii), as the representative temperature for high ionization potential ions and Te (N ii) or Te (O ii) for the low ionization potential ones. Te (O iii)is obtained from the [O iii] (4959+5507)/4363 ratio, Te (N ii) from [N ii] (6548+6583)/5755 and Te (N ii) from [O ii] (3727+3729)/(7319+7330). Te s are calculated making use 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 updated atomic dataset for O+ , S+ , and S++ for NEBULAR. The references are indicated in Table 4 of Garc´ ia-Ro jas et al. (2004). When one of the high/low ionization electron temperatures can not be computed, we used the linear relation between Te (O iii) and Te (O ii) provided by Garnett (1992), Te (O ii) = Te (N ii) = 0.7 â Te (O iii) + 3000, (2)

to estimate the unknown electron temperature. In the case that no direct estimate of the electron temperature can be obtained, we considered the Te (O iii) and Te (O ii) pairs that reproduce the total oxygen abundance obtained by applying the Pilyugin (2001a,b) empirical method (see below), also assuming the Garnett (1992) relation, that is the same equation that Pilyugin uses in his empirical calibrations. The electron density of the ionized gas, ne , was usually computed via the [S ii] 6716,6731 doublet, although sometimes the [O ii] 3726,3729 doublet was also used. Regions showing ne <100 cm-3 are below the low-density limit and hence we considered ne =100 cm-3 in those cases. Veilleux & Osterbrock (1987) proposed diagnostic diagrams plotting two different excitation line ratios, such as [O iii]/H versus [S ii]/H, for classifying the excitation mechanism of ionized nebulae. H ii regions (or H ii or starburst galaxies) lie into a narrow band within these diagrams, but when the gas is ionized by shocks, accretion disks or cooling flows (in the case of AGNs or LINERs) its position is away from the locus of H ii regions. We used the analytic relations given by Dopita et al. (2000) and Kewley et al. (2001) between different line ratios to check the nature of the excitation mechanism of the ionized gas within the bursts. Figure 5 shows an example of these diagrams applied to the regions analyzed in the galaxy Mkn 1199. 2.3.4. Chemical abundances Once the electron density and temperature are estimated, the ionic abundances can be derived for each region. All the ionic abundances except He+ and Fe++ were calculated using the IRAF NEBULAR task (Shaw & Dufour 1995) from the intensity of collisionally excited lines. We assumed a two-zone scheme for deriving the ionic abundances, adopting Te (O iii) for the high ionization potential ions O++ , Ne++ , S++ , Ar++ , Ar3+ and Cl++ ; and Te (N ii) or Te (O ii) for the low ionization potential ions O+ , N+ , S+ and Fe++ . The He+ /H+ ratio was computed from the He i lines intensities and using the predicted line emissivities calculated by Smith, Shara & Moffat (1996) for the Te (O iii) and ne assumed for each region. We also corrected for collisional contribution following the calculations by Benjamin, Skillman & Smits (2002). Self-absorption effects were not considered. Fe++ abundances were derived via the [Fe iii] 4658 emission line. We used a 34 level model-atom that includes the collision strengths of Zhang (1996) and the transition probabilities of Quinet (1996) We computed the total abundances of O, N, S, Ne, Ar and Fe. We always adopt O/H = O+ /H+ + O++ /H+ to determine the total oxygen abundance. We detect the He ii 4686 line in several ob jects, but the relative contribution of He++ to the total amount of helium is negligible, implying that O3+ has also a very low abundance in the nebula, thus we did not consider its contribution to the total O/H ratio. For deriving the nitrogen abundance we assumed the standard ionization correction factor (icf ) by Peimbert & Costero (1969): N/O = N+ /O+ , which is a reasonably good approximation for the excitation degree of the observed galaxies. We used the icf provided by Peimbert & Costero (1969) to derive the total neon abundance. We computed the total sulphur abundance when both S+ /H+ and S++ /H+ ratios are available using the icf given by the photoionization models by Stasinska (1978). The total ar´


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

gon abundance was calculated considering the icfs proposed by Izotov, Thuan, & Lipovetski (1994). The total iron abundances were obtained from the Fe++ /H+ ratio and the icf given by Rodr´ iguez & Rubin (2005). As we said before, when direct estimations of the electron temperature were not available, we resorted to empirical calibrations. Pilyugin (2001a,b) performed a detailed analysis of observational data combined with photoionization models to obtain the oxygen abundance from the relative intensities of strong optical lines. Pilyugin (2001a) gives the empirical calibration between the R23 and P (an indicator of the hardness of the ionizing radiation) parameters and the oxygen abundance in moderately high-metallicity H ii regions, 12+log(O/H)8.3. Pilyugin (2001b) provides the empirical calibration for the low-metallicity branch. Unless indicated, we always used Pilyugin empirical calibrations to derive the electron temperatures and the chemical abundances of the ionized gas in those regions where no direct determination of Te was possible. Sometimes, we estimated the total oxygen abundance making use of the Denicol´ Terlevich & Terlevich (2002) or Pettini & o, Pagel (2004) empirical calibrations, which involve the [N ii] 6583/H ratio. 2.3.5. Estimation of the Keplerian and Dynamical masses The bidimensional spectra have been used to perform a position-velocity diagram via the analysis of the brightest emission line profiles (H and/or [O iii] 5007) for all the ob jects. Although the main ob jective is the analysis of the kinematics of the ionized gas, in some cases we estimated the Keplerian mass (MKep ) of the galaxies assuming that the kinematics are representative of circular rotation. We usually considered the half of the maximum velocity difference, v , the half of the spatial separation corresponding to these maxima, r, and applied the equation: Mk
ep

ison gives important clues about the galaxy type, its dynamic and the fate of the neutral gas.

3. Results
3.1. NGC 1741 - HCG 31 AC The spectroscopic analysis of NGC 1741 (member AC in the galaxy group HCG 31) was presented in detail in L´ ezop S´ hez et al. (2004a). Our spectra show an evident broad anc blue WR bump and the He ii 4686 emission line in this ob ject (Figure 36). A careful re-analysis of the data reveals an evident and broad red WR bump in the same region (Figure 37), in agreement with previous observations of the same ob ject (Guseva, Izotov & Thuan 2000). The spectrum of members F1 and F2 (TDG s candidates) show the He ii 4686 emission line; member F2 also seems to show a blue WR bump. The analysis of the kinematics of the group reveals almost simultaneous interaction processes involving several ob jects. 3.2. Mkn 1087 Mkn 1087 is a luminous blue compact galaxy and the main galaxy in a group in interaction. Although some authors did observe WR features in this galaxy Kunth & Joubert (1985); Vaceli et al. (1997), others did not (Vacca & Conti 1992), and therefore it was classified as suspected WR galaxy by Schaerer, Contini & Pindao (1999). Our analysis of Mkn 1087 was presented in L´ ez-S´ hez et al. (2004b); op anc we did not detected any WR feature (Figure 36) in any important star-forming region in or surrounding Mkn 1087. 3.3. Haro 15 Haro 15 is a blue compact galaxy well studied in all frequencies, including optical spectroscopy (Hunter & Gallagher 1985; Mazzarrella, Bothum & Boronson 1991; Kong et al. 2002; Shi et al. 2005). Schaerer et al. (1999) listed Haro 15 as a WR galaxy because of the detection of the He ii 4686 emission line by Kovo & Contini (1999). Our analysis confirms the presence of WR stars (nebular and broad He ii 4686) in the bright star-forming region A (Figure 36) . Our long-slit spectroscopy covers the four main regions observed in Haro 15 (see Figure 3 of Paper I): the center (C), the bright region A at the ESE, the relatively bright H ii region D at the WNW (both observed with the slit with PA 117 ) and the knot B (at the NE, observed with the slit with PA 41 ). Figure 1 shows the spectra of the three brightest ob jects, whereas Table A.1 compiles the emission line ratios and other properties of the spectra of each region. The spectrum of the center of Haro 15 shows both nebular emission lines and stellar absorptions; these absorptions are observed mainly in the H i Balmer lines. However, the spectrum of region A is entirely dominated by nebular emission lines, where we detect [O iii] 4363 and He ii 4686. Because of the faintness of B and D, few emission lines are observed in these regions. 3.3.1. Physical conditions of the ionized gas A direct estimation of Te (O iii) was computed in knot A because of the detection of [O iii] 4363. We used empirical calibrations to estimate the electron temperatures for the

[M ] 233 â r [pc]

v [km s sin i

-1

]

2

,

(3)

assuming circular orbits and Keplerian dynamics (L´ ezop S´ hez et al. 2004b). We remark that the result of this anc equation is not the total dynamical mass of the galaxy but only the total mass contained within a circle of radius r. The inclination angle, i, is defined as that found between 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-on galaxy). We usually estimated this angle assuming that the elliptical shape of the galaxy is just a consequence of its orientation. When 21-cm H i data were available in the literature, we computed both the neutral gas mass, MH I , and the dynamical mass, MDyn , using the typical relations (i.e., Dahlem et al. 2005). We estimated the rotation velocity of the neutral WH gas considering v = 2 sinIi and adopted the maximum radius observed in our deep optical images. Therefore, as the extension of the neutral gas is usually larger than the extension of the stellar component, our estimations of MDyn may be underestimated. The gas depletion timescale defined by Skillman et al. (2003) was computed using MH I and the total star-formation rate (SFR) derived for each galaxy in Paper I. When FIR data are available, we estimated the mass of the warm dust, Mdust , using the equations given by Huchtmeier et al. (1995). Although MKep , MDyn and Mdust should be considered only tentative values, their compar-


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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10 20 30
-1

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50

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Wavelength [ å ]

relative velocity [ km s ]

relative velocity [ km s ]

Fig. 1. IDS INT spectra for regions C, A and B of Haro 15. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 3 in Paper I for the identification of the regions.
Relative flux
1 1 1 1 6 4 2 0 8 6 4 2 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 PA 117º 0 0 0 0 0 0 0 150 175
Center A

Fig. 3. Position-velocity diagrams for the two slit positions observed in Haro 15. Both the H and the [O iii] 5007 profiles have been analyzed. East is on top in both diagrams. See Figure 3 in Paper I for the identification of the regions.

[O III] 5007

600 500 PA 117º 400 300 200 100 0

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[Ne III] [O II] H H H H

H

H H

100 50 0 3500

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4500

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3500

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Wavelength [ å ]

Fig. 2. Lower panels: zoom of the spectra of regions C (center), A and B of Haro 15 (bottom). Fluxes are not corrected for reddening. Upper panels: spatial distribution of the relative flux of [O ii] 3727 (upper, right) and [O iii] 5007 (upper, left) emission lines along the slit of PA 117 .

experienced a different chemical evolution. The values for the oxygen abundance in knots B and D ­at larger distances from the center of the galaxy­ are slightly lower than that at the center of Haro 15, perhaps indicating a possible radial abundance gradient in the disk of the galaxy. However, this may be no applicable to knot B because it could be an independent ob ject due to its decoupled kinematics (see below). Regions A, B and D show a much lower N/O ratio than the center of the galaxy. Figure 2 shows the relative flux of the [O ii] 3727 and [O iii] 5007 emission lines along the slit. As we can see, knot A has a much higher ionization degree compared to the other regions. This should be a consequence of the extreme youth of the star formation in this zone an its lower metallicity. Both things suggest the different nature of this ob ject. 3.3.3. Kinematics of the ionized gas Figure 3 shows the position-velocity diagrams for the two slit positions observed in Haro 15. We analyzed both the H and the [O iii] 5007 profiles, extracting 4 pixels bins (1.6 arcsec) and considering the velocity of the center of the galaxy as reference. Both emission lines give almost identical results. The diagram of PA 41 shows an apparent rotation pattern, although some divergences are found at the SW. Knot B is clearly decoupled from the rotation of the disk, suggesting that it is an external ob ject. The interaction between knot B and Haro 15 could be the responsible of the distortion observed at the SW of the diagram of PA 41 because the ob ject and this zone of the galaxy disk show similar radial velocities. Another possibility is that knot B is a TDG but, in this case, the material that formed B should come from the external parts of the disk of Haro 15 because its chemical abundances are more similar to those of knot D than to those of the central region. The diagram with PA 117 shows a clear sinusoidal pattern with differences of around 40 km s-1 . This feature is common in processes involving galaxy interaction or merging. Although it is not completely clear, region A seems to be kinetically coupled with the rotation of the galaxy. This

rest of the regions; the results are shown in Table A.2. The spectrum of the central region of Haro 15 has a marginal detection of [O iii] 4363. Using this value we derived Te [O iii]9700 K, which is very similar to that obtained using empirical calibrations. [S ii] 6717 is blended with a skyline in all spectra, and hence we can not compute ne . We assumed a value of 100 cm-3 for all regions. Comparing the line ratios with diagnostic diagrams, we found that all knots can be classified as typical H ii regions. 3.3.2. Chemical abundances Table A.2 compiles all the chemical abundances computed for each region of Haro 15. We found a significant difference between the oxygen abundance at the center of Haro 15, 12+log(O/H)=8.37, and that found in the ESE region (knot A), 12+log(O/H)=8.10, which is larger than the uncertainties. This fact suggests that, although their pro jected distance is very small (5.5 kpc), both ob jects seem to have

erg s cm å ]

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8
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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
1,5
[O II] [Ne III] H H WR H [O III]

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Shock contribution Shock contribution

erg s cm å ]

60 50 40 30 20

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He I [O I]

H [N II] [S II] [Ar III] [O II]

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Observed flux [ 10

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H II regions

Mkn 1199 C

-0,5

4000

4500

5000

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6000

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Wavelength [ å ]

H II regions

Fig. 4. IDS INT spectra for the center of Mkn 1199 and the dwarf companion galaxy at its NE. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 5 in Paper I for the identification of the regions.

-1,0

Obs. data D00 K01 -1,5 -1,0 -0,5 0,0 0,5 1,0

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Obs. data D00 K01

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fact and the different chemical composition indicate that knot A is probably an external ob ject suffering a merging process with Haro 15. Assuming that the position-velocity diagram of PA 41 from -5 to 15 arcsec from the center is consequence of circular rotation, we can estimate the Keplerian mass of Haro 15. Considering a radial velocity of 80 km s-1 within a distance of 13 arcsec (5.46 kpc) and assuming i=55 from our optical images [Gordon & Gottesman (1981) found i=57 ], we derive Mkep =1.21â1010 M and MKep /LB = 0.35. The neutral and dynamical masses of Haro 15 are MH I =5.54â109 M and MDyn =3.65â1010 M (Gordon & Gottesman 1981). Thus, our Keplerian mass is 33% of the the total mass. The MH I /L and the Mdust /L ratios (0.16 and 0.6â10-4 , respectively) suggest that Haro 15 is a Sb or Sc spiral (Bettoni, Galleta & Garc´ ia-Murillo 2003). The gas depletion timescale is 2.3 Gyr, showing that the system still possesses a huge amount of material available to form new stars. 3.4. Mkn 1199 Izotov & Thuan (1998) observed Mkn 1199 but they did not detect the [O iii] 4363 emission line. However, they reported the blue and red WR bumps and nebular He ii 4686 emission. Hence Mkn 1199 was classified as a WR galaxy (Schaerer, Contini & Pindao 1999). Guseva et al. (2000) revisited its WR properties indicating the existence of WNL, WNE and WCE stellar populations. We obtained intermediate-resolution long slit spectroscopy of Mkn 1199 using two different slit positions (see Figure 5 of Paper I). We analyzed five different regions: the center of Mkn 1199 (C), the companion dwarf ob ject (NE) and knots A, B and D, all extracted with a slit with PA 36 but region A, for which the slit with PA 53 was used. The spectra of the two brightest regions (C and NE) are shown in Figure 4, while the emission line intensities and other important properties of the spectra are compiled in Table A.3. The spectrum of the center of Mkn 1199 shows stellar absorptions in both the H i Balmer and the He i lines, as it was previously noticed by Izotov & Thuan (1998). The blue WR bump and a probable He ii 4686 are detected in C (Figure 36). Our spectrum also shows a tentative detection of the red WR bump in this region (Figure 37). The spectrum of the dwarf companion galaxy at the NE has a lower S/N ratio than the obtained for Mkn 1199, but

Fig. 5. Comparison of some line intensity ratios in several regions of Mkn 1199 with the diagnostic diagrams proposed by Dopita et al. (2000) and Kewley et al. (2001).

all relevant emission lines are clearly identified. Stellar absorptions are also found in this region but are fainter than those observed in C. The spectra of the rest of the knots only show the brightest emission lines. 3.4.1. Physical conditions of the ionized gas The [O iii] 4363 line was not detected in any region, hence we computed the electron temperatures and chemical abundances using empirical calibrations (see Table A.4). The electron temperatures found in C are rather low. In this zone, we detected [N ii] 5755, which gives a Te [N ii]6740 K, a value very similar to that estimated using the empirical method. We also detect the [O ii] 7319,7330 doublet in the spectrum of C; we obtain Te (O ii)6910 K. Hence, we consider that the electron temperature values obtained using the empirical calibrations are entirely reliable for this ob ject, assuming Te (low)=6800 K. The high ionization electron temperature was computed using Garnett's relation, Te (low)=5400 K. Except for region C, the electron densities derived from the [S ii] 6717,6730 doublet were in the low-density limit, ne 100 cm-3 . The values for the reddening coeficient are relatively high for C and B, suggesting the presence of an important amount of dust in those regions. The comparison of some line intensity ratios with the diagnostic diagrams proposed by Dopita et al. (2000) and Kewley et al. (2001) indicates that all regions can be classified as H ii regions (see Figure 5). 3.4.2. Chemical abundances Table A.4 compiles all the chemical abundances computed for the knots analyzed in Mkn 1199. The oxygen abundance found at the center of the system is very high, 12+log(O/H)=8.75± (indeed, it highest metallicity region found in this work). However, the oxygen abundance of the dwarf companion galaxy at the NE is almost 0.3 dex lower, 12+log(O/H)=8.46. The N/O ratios derived for them are also very different. Hence, this result reinforces the hypothesis that they are independent galaxies in the first stages of a minor merging.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
35 30 25 20
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9

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Mkn 1199 - PA 53º

distance (arcsec)

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15
B

15 10 5 0 -5 -10 -15 -20 -60 -40 -20 0 20 40 60 80 100
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c e n te r

c e n te r

D

Mkn 1199 - PA 36º -20 -60 -50 -40 -30 -20 -10 0 10 20 30 40

relative velocity (km s )

-1

relative velocity (km s )

Fig. 6. Position-velocity diagrams for the two slit positions observed in Mkn 1199 using the H profiles. NE is up in both diagrams. See Figure 5 in Paper I for the identification of the regions.

It is worthy to notice that, although the line intensities observed by Izotov & Thuan (1998) in Mkn 1199 are quite similar to those compiled in Table A.3, the oxygen abundance derived by these authors is 12+log(O/H)=8.19±0.18, extremely low when compared with our values. However, Guseva, Izotov & Thuan (2000) using the same data that Izotov & Thuan (1998) and the empirical calibration based on the [N ii]6583/H ratio proposed by van Zee, Salzer & Haynes (1998) obtain 12+log(O/H)=9.13. We consider that 12+log(O/H)=8.75 is a more appropriate value for the oxygen abundance at the center of Mkn 1199, that also agrees with our tentative estimations of electron temperatures. Knots A, B and D show oxygen abundances slightly lower than C, between 8.6 and 8.7, but not as low as that reported for the companion galaxy. Interestingly, A, B and D show a N/O ratio very similar to that observed at the center of the galaxy, but much higher than in the dwarf companion. This confirms the different chemical evolution of the disk of Mkn 1199 and the companion galaxy. We consider that knots A, B and D are giant H ii regions located in the spiral arms of Mkn 1199 and that this galaxy could have a slight radial metallicity gradient along its disk. Finally, we think that the triggering mechanism of the intense star-formation activity found in both, Mkn 1199 and its dwarf companion, is a likely consequence of the strong interaction they are experiencing. 3.4.3. Kinematics of the ionized gas Figure 6 shows the position-velocity diagrams for the two slit positions in Mkn 1199. They were obtained extracting 3 pixels bins (1.2 arcsec) across the H profile and considering the center of Mkn 1199 as reference. The diagram with PA 36 , that crosses the center of Mkn 1199 and the companion galaxy, may be explained by the rotation of the disk of Mkn 1199, because it shows a velocity gradient between 30 km s-1 at SW (knot D) and -30 km s-1 at NE (around knot B). However, a sinusoidal pattern is also seen in the brightest region of Mkn 1199, within the central 10 . This feature may be consequence of an entity kinetically decoupled from the disk, such a bar or the bulge of the galaxy, that seems to be counter-rotating, although it may be also

an interaction feature. On the other hand, although there are only four points in the diagram in that zone, the dwarf galaxy at the NE seems to be rotating. This ob ject, that has an elliptical shape in the optical images, may be observed edge-on, and hence it would explain both its kinematics and morphology. The diagram with PA 53 also seems to show a velocity gradient from the SW to the NE (where knot A is located) but this gradient is broken at the central 10 . The maximum velocity difference is around 80 km s-1 . It is difficult to estimate the Keplerian mass of Mkn 1199 because it is seen almost face-on. However, considering that the velocity difference between its center and the external regions within a radius of 10 (2.6 kpc) is around 30 km s-1 and adopting an inclination angle of 15 , we derive Mkep 8.2â109 M . For the companion dwarf ob ject, assuming that it is edge-on (i=90 ) and considering a velocity difference of 10 km s-1 within a radius of 2.5 (1.3 kpc), we compute Mkep 2.9â107 M . Using the H i data given by Davoust & Contini (2004), we compute MH I =1.22â109 M and vH I =85 km s-1 . Assuming a radius of 25 (6.55 kpc) and i=15 , the dynamical mass of Mkn 1199 is Mdyn =1.7â1011 M . The warm dust mass using the FIR data is Mdust =3.1â106 M . Following the classification provided by Bettoni et al. (2003), the MH I /L =0.042 and Mdust /L =1.1â10-4 ratios are not compatible: while the the first corresponds to the typical values for S0 galaxies, the second ratio indicates that Mkn 1199 should be an Sc or Sd, more similar to the actual morphological classification of Sb. Furthermore, less than 1% of the total mass of the system is neutral hydrogen and the gas depletion timescale is very low (0.4 Gyr). All these facts suggest that a substantial fraction of the neutral H i gas has been expelled to the intergalactic medium, perhaps as a consequence of the interaction with the dwarf companion galaxy. An H i map obtained using a radiointerferometer would be needed to confirm such hypothesis. 3.5. Mkn 5 Conti (1991) included Mkn 5 in his catalogue of WR galaxies because of the detection of the nebular He ii 4686 line by French (1980). However, Izotov & Thuan (1998) only observed the blue WR bump, without any trace of the nebular He ii emission (Schaerer et al. 1999). Guseva, Izotov & Thuan (2000) detected N iii 4640, implying the presence of WNL stars within the starburst. We used three slit positions to get the spectroscopic data of Mkn 5 (see Figure 7 of Paper I), two of them using the 2.5m INT and an additional position using 4.2m WHT. All have a very similar PA, 0 (INT-1), 354 (WHT) and 349 (INT-2). All slit positions cover region A but only two cross knot B. We analyzed the three spectra extracted for region A independently to check the quality of the results. Figure 7 shows the spectra of the region A obtained with the slit positions with PA 349 and PA 354 . For region B we only analyzed the spectrum extracted using the slit position with PA 0 (INT-1) because of its higher signal-tonoise. Table A.5 compiles all the emission line fluxes and other characteristics of each spectrum. Our spectra confirm the presence of a nebular He ii 4686 line on top of a broad emission line in region A (Figure 36). Although the WHT spectrum has high S/N and spectral resolution, we do not detect the red WR bump in this region (Figure 37). All


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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Wavelength [ å ]

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Fig. 7. Spectra for the region A of Mkn 5 obtained using IDS at the INT (PA 349 ) and ISIS at the WHT (PA 354 ). Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 7 in Paper I for the identification of the regions.

Fig. 8. Position-velocity diagrams for the slit positions observed in Mkn 5 using the [O iii] 5007 profile. N is up in all diagrams. See Figure 7 in Paper I for the identification of the regions.

spectra are dominated by nebular emission but some absorptions in the H i Balmer lines are also detected; these are more evident in knot B. 3.5.1. Physical conditions of the ionized gas The three spectra obtained for region A show the [O iii] 4363 emission line and hence we computed the electron temperature using the direct method. As it is seen in Table A.6 all Te values are in agreement within the errors, being the average value Te [O iii]12500 K. The low ionization temperature was derived using Garnett's relation, Te [O ii]11700 K. The spectrum INT-2 shows a tentative detection of [O ii] 7318,7330, that gives Te [O ii]11950 K, in agreement with the electron temperature derived using Garnett's relation. We used Pilyugin (2001b) empirical calibration to derive Te in knot B, but this determination is very much uncertain and perhaps overestimated due to the faintness of the spectrum. Electron densities are always below the low-density limit (100 cm-3 ) except for knot B (although it also has a large error). The values of the reddening coefficient derived for region A are somewhat different in the different spectra. Perhaps, this apparent inconsistency is a consequence of an irregular distribution of dust within Mkn 5, as we suggested in our analysis of the optical/NIR colors (see § 3.5.1 of Paper I). For knot B, we assumed Wabs =1.5 ° and the c(H ) comA puted via the H/H ratio. The diagnostic diagrams for knots A and B agree with the loci of typical H ii regions. 3.5.2. Chemical abundances The WHT spectrum does not cover the [O ii] 3726,29 doublet, hence we used [O ii] 7318,7330 to compute the O+ /H+ ratio. As we see in Table A.5, the agreement between the emission line ratios for all the three spectra extracted for region A is very good. Table A.6 compiles the chemical abundances obtained for Mkn 5; for region A all values are quite similar and in agreement with previous results found in the literature. Averaging all data and minimizing errors, we derive for A the following chemical abundances: 12+log(O/H)=8.07±0.04, log(N/O)=-1.38±0.07, log (S/O)=-1.62±0.11, log(Ne/O)=-0.80±0.13 and log(Ar/O)=-2.31±0.12. These values are very similar to

the results provided by Izotov & Thuan (1999). On the other hand, the oxygen abundance estimated for knot B, 12+log(O/H)=7.89±0.17, lower than that derived for the main star-forming region but consistent within the errors. If real, this difference may suggests different chemical evolutions between the two regions. 3.5.3. Kinematics of the ionized gas Figure 8 shows the position-velocity diagrams obtained for the three slit positions observed in Mkn 5. For the 4.2m WHT spectrum we analyzed the H profile considering 6 pixels bins (1.2 arcsec), while we used the [O iii] 5007 profile (brighter than the H profile) extracting 4 pixels bins (1.6 arcsec) from the 2.5m INT spectra. As we see, the agreement between the three diagrams is very good. The best diagram is that provided by the analysis of the 4.2m WHT spectrum, that shows a velocity gradient of around 50 km s-1 . The velocity of knot B with respect to that found in region A is 40 km s-1 . Although the uncertainties are important, we detect a slight reverse in the velocity of region A, with an amplitude of 20 km s-1 (it is more evident in the 4.2m WHT diagram), that seems to show a sinusoidal pattern in that area. Assuming that the global velocity gradient is mainly a consequence of the rotation of the galaxy, we may estimate the Keplerian mass of the system. We found MK ep 2.1 â 109 M assuming i=90 , v 27 km s-1 and r 21 (1.22 kpc). Using H i data (Paturel et al. 2003), we derive MH I = (7.2±0.9)â107 M and MDyn 3.6â109 M . Although both MKep and MDyn are similar, notice that they are low limits because we are assuming that Mkn 5 is an edge-on galaxy. The mass-toluminosity ratios are MK ep /L =7.98, MDyn /L =13.7 and MH I /L =0.27. The H i mass is quite low for a dwarf or irregular galaxy, being only 2% of the total mass. The gas depletion timescale is 1.8 Gyr, high for a starburst galaxy. All these facts suggest that something has happened with the atomic gas of Mkn 5: or it has been consumed forming stars at a high rate until some few hundred Myr ago (nowadays it has decreased) or it has been expelled to the intergalactic medium. Indeed, Thuan & Martin (1981) found indications of H i gas at slightly different radial velocities ( 300 km s-1 ) using single-dish data. An interferometric H i map will be crucial to understand the fate of the neutral gas in this blue compact dwarf galaxy.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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Fig. 9. IDS INT spectra for the center of IRAS 0828+2816 and knot #8. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 9 in Paper I for the identification of the regions.

All electron temperatures derived using empirical methods, remarking those found in knot #8, are systematically lower than those computed in the central region following the direct method. All ob jects can be classified as H ii regions following the typical diagnostic diagrams. The reddening coefficient was computed using all available H i Balmer lines in each spectrum. We obtained very different values: while the central region and knots in the northern tail have a low reddening coefficient, c(H )0.12, knots located in the southern tail show a higher value, c(H )0.43. This fact seems to indicate an inhomogeneous distribution of the dust within the galaxy, being the southern regions dustier than the rest of the system. Huang et al. (1999) also derived low extinction values in the center of IRAS 08208+2816; they explained this considering the presence of a galactic wind in the galaxy. But we do not detect such structures in our deep H images (see Figure 10 of Paper I). 3.6.2. Chemical abundances Table A.8 compiles all chemical abundances computed for the different knots analyzed in IRAS 08208+2816. The oxygen abundance of the central region, derived using the direct method, is 12+log(O/H)=8.33±0.08, and its nitrogento-oxygen ratio is log(N/O)= -0.89 ± 0.11. This value is higher than the N/O ratio expected for a galaxy with an oxygen abundance of 12+log(O/H)8.3, which should be log(N/O) -1.2. If this effect is real, it may be due to nitrogen pollution by the winds of the WR stars, as we confirm to occur in the case of NGC 5253 (L´ ez-S´ hez op anc et al. 2007). For the rest of the ob jects, the oxygen abundances were calculated using Pilyugin (2001a,b) empirical calibrations. Although all estimations are slightly higher than the value found in the central region, we notice a significant difference in the case of knot #8, that has 12+log(O/H) 8.64 (i.e., almost the solar value). Knots #3 and #5 show a tentative detection of [O iii] 4363 in their spectra, for which we derive an oxygen abundance 0.12 ­ 0.15 dex lower than that estimated using empirical calibrations (see Table A.8). This trend is also found in the central region, for which we derive 12+log(O/H)8.41 following the empirical calibrations. Hence, the values obtained using the Pilyugin method may be somewhat overestimated for this galaxy. In any case, the chemical abundance differences seem to be real in knot #8, first because its oxygen abundance is 0.3 dex higher than that found in the central region, and second because its N/O ratio is also the highest, log(N/O) -0.84, and consistent with the value expected for a galaxy with almost solar metallicity. This result indicates that knot #8 could be an ob ject more chemically evolved that the others. Because of this and its position within the system, knot #8 may even correspond to the center of an independent galaxy that is in a process of merging with another galaxy which nucleus coincides with knot C of IRAS 08208+2816. 3.6.3. Kinematics of the ionized gas Figure 10 shows the position-velocity diagrams obtained for the three slit positions. The [O iii] 5007 profile was analyzed considering 3 pixel bins (1.2 ) and taking the center

-1

-2

3.6. IRAS 08208+2816 The first spectroscopic data of IRAS 08208+2816 were obtained by Huang et al. (1999), who reported the detection of both the nebular and broad He ii 4686 emission line. They also detected the red WR bump, C iv 5808, suggesting the presence of both WNL and WCE populations in the starburst. Schaerer et al. (1999) included IRAS 08208+2816 in the latest catalogue of WR galaxies . We used three slit positions with the IDS spectrograph at the 2.5m INT to cover all bright knots within the galaxy (see Figure 9 of Paper I). The slit position with PA 345 crosses a bright star and the center of IRAS 0828+2816, the slit position with PA 355 covers the center and knots #8, #10 and #1 (that is very affected by the light of the bright star) and the slit position with PA 10 crosses #3, #4 (very weak) and #5, although it also may be contaminated by some emission from the center and knot #8. Table A.7 compiles the properties of the five regions that we analyze spectroscopically. The center of the galaxy, C, corresponds to the brightest region extracted using the slit position with PA 345 . This spectrum and that obtained for knot #8 are shown in Figure 9. Although all spectra are dominated by nebular emission and do not show any evidence of absorptions in the H i Balmer lines, we detect a slight decrease of the continuum level on the blue range of the spectra of the faintest ob jects . This fact may be explained by both an important extinction in these ob jects and by the possible presence of an evolved underlying stellar population. The broad He ii 4686 line is weakly detected in the central region of the galaxy but the nebular He ii 4686 line is not identified in this spectrum (Figure 36). A faint red WR bump at around 5800 ° seems also to be observed in A this region (Figure 37). Spectra with higher S/N ratio and spectral resolution are needed to get a proper value of the WR bump fluxes. 3.6.1. Physical conditions of the ionized gas We detect the weak auroral [O iii] 4363 emission line at the center of the galaxy, its flux value was used to estimate the electron temperature following the direct method. For the rest of the ob jects we used empirical calibrations to compute Te , all results are compiled in Table A.8. [O iii] 4363 is barely detected in knots #3 and #5, but they were not considered in the analysis because of their large errors.


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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relative velocity (km s )

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Fig. 10. Position-velocity diagrams for the three slit positions observed of IRAS 08208+2816 using the [O iii] 5007 profile. N is up in all diagrams. We have included the position of the regions (see Figure 9 in Paper I for their identification), emphasizing those analyzed by spectroscopy. Notice that the lacking of data at the southern edge of the southern tail in the diagrams with PA 345 and 355 is because of the contamination by a bright star.

Fig. 11. ISIS 4.2m WHT spectrum for the center of POX 4. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 11 in Paper I for the identification of the regions.

of IRAS 08208+2816 as reference. As it was clearly seen in the bidimensional spectra, this ob ject possesses very interesting kinematics, remarking a probable tidal stream in the northern tail. In these areas we observe a velocity gradient larger than 300 km s-1 within 12 (11 kpc) in the slit position crossing the northern region of the system (PA=355 ). This velocity difference of 300 km s-1 is the same that Perryman et al. (1982) reported between these ob jects. As we commented before, the spectra crossing the southern tail is affected by the contamination from the nearby bright star, but the slit position with PA=10 is free of such contamination allowing the kinematic analysis of the southern zone. Again, we found an important velocity gradient towards negative values, that cannot be explained by a rotating disk. Furthermore, this diagram shows an evident sinusoidal pattern with an amplitude larger than 50 km s-1 at the center of the galaxy. This is an additional evidence that we are observing a process of merging. In general, the agreement between the three diagrams is very good, for example, the velocity found for knot #1 using the slit position with PA 355 is -250 km s-1 , that corresponds quite well with the velocity observed at the end of the southern tail (knot #3, with -200 km s-1 ) using the slit position with PA 10 . Although the kinematics of the system is not supported by rotation, we have computed a tentative estimation of the Keplerian mass of the system. Assuming i=90 , v 30 km s-1 (using the diagram with PA 345 that seems to be less affected by the tidal tails) and a radius of 20 (18.4 kpc), we derive MKep 3.9â109 M . The mass-toluminosity ratio is quite low, MKep /LB 0.08, hence we are very probably underestimating the mass of the system. There are no H i data available for this galaxy in the literature, but it should be really interesting to compare the kinematics of the neutral gas with that found here for the ionized gas. The warm dust mass is high, Mdust =8.84â106 M , giving Mdust /LB 1.73â10-4 . 3.7. IRAS 08339+6517 Our complete analysis of the physical properties, chemical abundances and kinematics of the ionized gas in IRAS 08339+6517 and its companion dwarf galaxy was pre-

sented in L´ ez-S´ hez, Esteban & Garc´ op anc ia-Ro jas (2006). We reported weak spectral features that could be attributed to the blue WR bump at the center of the galaxy. The kinematics of the ionized gas showed an interaction pattern that indicates that the H i tidal tail detected by Cannon et al. (2004) in the direction of the dwarf companion galaxy has been mainly formed from material stripped from the main galaxy. A star-forming region in the outskirts of the galactic disk may be a TDG candidate. 3.8. POX 4 The first indications of WR features in POX 4 were noticed by Kunth & Joubert (1985) and Campbell, Terlevich & Melnick (1986)2 , because both detected the broad He ii 4686 emission line. Therefore, Conti (1991) included POX 4 in his catalogue of WR galaxies. Masegosa, Moles & del Olmo (1991) also suggested the presence of WR stars in one or two regions of the galaxy3 . Vacca & Conti (1992) confirmed the presence of a high number of O and WN stars in the brightest region of POX 4 and detected the He ii 4686 emission line in other knot (Schaerer et al. 1999). We do detect the nebular He ii 4686 line with a very good S/N ratio in the brightest knot of POX 4. Both the broad He ii 4686 and C iv 5808 lines are clearly identified in this region too (Figures 36 and 37). We used ISIS at the 4.2m WHT to obtain the spectroscopic data of POX 4 and its dwarf companion galaxy (knot #18 following Figure 11 in Paper I). The position angle of the slit was set to 25 . The spectrum of the center of POX 4 (Figure 11) is dominated by intense emission lines and does not show any evidence of stellar absorption in the H i or He i lines. As it was previously noticed by M´ endez & Esteban (1997), broad low-intensity assymetric wings are detected in the profiles of the brightest emission lines (H and [O iii]). The clear detection of the nebular He ii 4686 emission on the top of a broad feature indicates the presence of WR stars at the center of the galaxy. However, we do not see the red WR bump despite the high S/N and spectral resolution of our spectrum. The spectrum of the dwarf companion ob ject has a rather low S/N ratio and only the brightest emission lines
These authors named POX 4 as C 1148-203. These authors named POX 4 as C 1148-2020 in their Table 1 and as Tol 1148-202 in their Table 2. Following the NED, the appropriate name of this galaxy is IRAS 11485-2018 = POX 4.
3 2


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13

Distance [ arcsec ]

are detected. However, it shows clear absorption features in the H i Balmer lines, indicating the presence of evolved stellar populations underlying the starburst. Table A.9 compiles all emission lines detected in the center of POX 4 and in the dwarf companion ob ject, as well as other important properties of their spectra. 3.8.1. Physical conditions of the ionized gas The physical conditions and chemical abundances of the ionized gas in POX 4 are compiled in Table A.10. The electron temperature calculated for the center of POX 4 using [O iii] 4363 is Te (O iii) = 14000±600 K, suggesting that it is a low metallicity ob ject. Although [N ii] 5755 is barely detected, we preferred to use Garnett's relation to determine Te (O ii) from Te (O iii). The electron density was computed using the [S ii] 6717/6731 ratio, yielding ne 250 cm-3 . Notice that the estimation of ne using the [Ar iv] 4711/4740 ratio gives a very similar value, 270 cm-3 , although it has a higher uncertainty. For the dwarf companion ob ject we used empirical calibrations to estimate Te , being its electron density in the low-density limit. The comparison of the observed [O iii]5007/H and [N ii]6584/H ratios with the diagnostic diagrams let to classify all regions as starbursts. The reddening coefficient at the center of POX 4 was determined using 7 H i Balmer lines, obtaining c(H )=0.08±0.01. A similar low value of the reddening has been found for the companion ob ject. 3.8.2. Chemical abundances The oxygen abundance derived for the center of POX 4, 12+log(O/H)=8.03±0.04, agrees with that found in the literature (e.g. Kobulnicky & Skillmann 1996 reported 7.97±0.02). The nebular He ii 4686 emission line is clearly detected and, therefore, some O+3 contribution is expected in the nebular gas, however this contribution is found to be marginal, 0.01­0.02 dex. The oxygen abundance derived for the companion ob ject using empirical calibrations, 12+log(O/H)=8.03±0.14, is the same than that found in POX 4. The values of the N/O ratio are also similar, log(N/O)= -1.54 ± 0.06 and -1.60. Despite the uncertainties, the resemblance of the chemical abundances may suggest that the dwarf companion is not an independent object, as M´ endez & Esteban (1999) concluded, but a TDG candidate. 3.8.3. Kinematics of the ionized gas The position-velocity diagram obtained for the slit-position taken for POX 4 is shown in Figure 12. Because of its higher intensity, we used the [O iii] 5007 profile instead of the H profile, extracting 4 pixel bins (0.8 arcsec) and taking as reference the velocity found in the center of the galaxy (knot #9, see Figure 11 in Paper I). The diagram shows a clear irregular pattern without any rotation evidence, indicating that the movement of the ionized gas of the system is rather chaotic but of small amplitude. M´ endez & Esteban (1999) proposed that the companion (#18) has actually gone through POX 4, being the origin of the peculiar ring-morphology of the galaxy and the strong starforming bursts observed throughout all the system. The

20 15 10 5 0 -5 -10 -15 -20

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# 18 - Companion
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40

50

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Fig. 12. Position-velocity diagram for the slit position observed in POX 4 using the [O iii] 5007 profile. NE is up. See Figure 11 in Paper I for the identification of the regions.

diffuse dwarf companion, that possesses a radial velocity similar to that observed at the center of POX 4, also shows peculiar kinematics, being in some sense coupled with the movement of the ionized gas in the main galaxy. Our optical data do not suggest that the dwarf companion ob ject is a TDG because we do not detect any other dwarf galaxy that interacted with POX 4 and stripped some material from it. M´ endez & Esteban (1999) indicated a velocity difference of 130 km s-1 between POX 4 and the companion galaxy, but this is not confirmed in our deeper spectroscopic data. Taking into account the complex kinematic structure shown in Figure 12, it is clear that we can not obtain a confident estimate of the Keplerian mass of POX 4. However, assuming a ratio of MK ep /LB 0.1, it would be of the order of MK ep 5 â 108 M . The H i Parkes Al l-Sky Survey (HIPASS; Barnes et al. 2001) shows a tentative detection of H i emission. Our group has performed 21-cm observations of POX 4 using the interferometer Australia Telescope Compact Array (ATCA). Although a detailed description and analysis of such observations will be presented elsewhere (L´ ez-S´ hez et al. 2010), the very preliminary op anc analysis suggests that the system possesses a lot of neutral gas. The H i kinematic are perturbed in the position of the dwarf companion ob ject but it shows the same radial velocity we found using optical spectroscopy. An independent H i cloud, that has the same radial velocity that POX 4, is found at 4.5 (60 kpc) at the south. It shows a clear alignment with both the bright center of POX 4 and the dwarf companion ob ject, suggesting a very probable interaction in the past. A detailed analysis of the H i observations will confirm o discard the TDG nature of the dwarf companion ob ject surrounding POX 4. 3.9. UM 420 Izotov & Thuan (1998) reported the detection of the broad He ii 4686 emission line in UM 420, being therefore included in the latest catalogue of WR galaxies (Schaerer et al. 1999). The reanalysis of their spectra performed by Guseva et al. (2001) also suggests the presence of the C iv


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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Fig. 13. ISIS 4.2m WHT spectrum for the center of UM 420 (top ) and the galaxy UGC 1809 (bottom ). Fluxes are not corrected for reddening. The most important emission lines have been labeled.

H9 / H 0,00 0,0 0,5 1,0 1,5 2,0 2,5 3,0

Wa b s

4658 and C iv 5808 emission lines, indicating the possible existence of WCE stars in the starburst. Figure 13 shows the spectra of UM 420 and UGC 1809 for our slit position (see Figure 13 of Paper I). Notice the huge difference between both spectra: while the spectrum of UM 420 is dominated by emission lines, the spectrum of UGC 1809 only shows stellar absorption features as Ca ii H,K, G-band and Mg i 5167. Hence, we may classify this ob ject as a S0 spiral galaxy with redshift of z =0.0243. The radial velocity of UGC 1809 is vr =7290 km s-1 , in excellent agreement with the value given by the NED (vr =7306 km s-1 ) but much lower than the radial velocity of UM 420 (vr =17507 km s-1 ). That confirms that both galaxies are not physically related. The spectra of UM 420 do not show absorption features. We observe, although with large error, the nebular He ii 4686 line on top of a very faint broad feature (Figure 36). We do not detect the red WR bump despite the good S/N ratio and the clear detection of the weak [N ii] 5755 auroral line (Figure 37). The list of the emission lines observed in UM 420 is compiled in Table A.9. 3.9.1. Physical conditions of the ionized gas Using the [O iii] 4363 line intensity we compute Te (O iii)=13200±600 K in UM 420. Although it has a large error, the detection of the auroral [N ii] 5755 emission line indicates Te (N ii)11800 K, that is similar to the low ionization temperature given by Garnett's relation between Te (O iii) and Te (O ii). The electron density computed using the [S ii] 6716,6731 doublet is in the low-density limit, but that estimated using the [O ii] 3726,3729 lines gives ne 140 cm-3 . The reddening coefficient and the underlying stellar absorption in the H i Balmer lines were computed using 5 ratios between the H i Balmer lines and give very consistent results, which mean values are c(H )=0.09±0.01 and Wabs =2.0±0.1 (see Figure 14). The comparison of the observed line flux ratios with the diagnostic diagrams clearly identify UM 420 as a starburst galaxy. 3.9.2. Chemical abundances Table A.10 lists all the chemical abundances computed for UM 420. The derived oxygen abundance is 12+log(O/H)=7.95±0.05, in excellent agreement with the value given by Izotov & Thuan (1998). Our estimation does

Fig. 14. Interactive estimation of c(H ) and Wabs using the six brightest H i Balmer lines detected in the spectrum of UM 420. Note the excellent agreement in the behaviour of all lines.
8 6 4

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H H
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Fig. 15. Position-velocity diagram for the slit position observed in UM 420 using both the H and H profiles. W is up.

not consider the small contribution of O+3 (that should exist because of the detection of the He ii) but it is smaller than 0.01 dex. The N/O ratio, log(N/O)=-1.11 ± 0.08, is also similar to the value reported by these authors. It is interesting to notice that its N/O ratio is higher than that expected for its oxygen abundance, which should be around 0.4 dex lower. This fact was previously reported by Pustilnik et al. (2004), who suggested that the overabundance of nitrogen may be produced by pollution by the large amount of WR stars present in the violent starburst triggered by galaxy merging. The value of the neon abundance in UM 420, log(Ne/O)=-0.71±0.13, is also similar to that given by Izotov & Thuan (1998). We derive log(S/O)=-1.66±0.15, that is the typical value for BCDGs with the same oxygen abundance (Izotov & Thuan 1999). 3.9.3. Kinematics of the ionized gas Figure 15 shows the position-velocity diagram obtained for the slit position with PA 90 observed in UM 420. Both


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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are not observed despite the high S/N and spectral resolution. There are not substantial stellar absorptions in the spectrum of SBS 0926+606 A, but we detect broad lowintensity assymetric wings in the H profile.
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[Ar III] He I [O II]

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Fig. 16. ISIS 4.2m WHT spectrum for SBS 0926+606 A. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 15 in Paper I for the identification of this region.

the H and H profiles were analyzed, extracting 3 pixel bins (1.08 ) for H and 4 pixel bins (0.8 ) for H . The diagram is identical in both cases. Although the number of points is small, we notice a velocity gradient of around 30 km s-1 from the eastern region to the center of the galaxy, but this tendency is reversed in the western areas of UM 420. Indeed, in this region a negative velocity gradient of 70 km s-1 is found within only 4 (4.6 kpc). The diagram does not allow to get any determination of the Keplerian mass. There are not H i data of the galaxy to derive its neutral or dynamical masses. 3.10. SBS 0926+606 SBS 0926+606 actually is a pair of compact nearby objects (see Figure 15 of Paper I). SBS 0926+606 A has been spectroscopically studied by Izotov and collaborators with the aim of determining the primordial helium abundance (Izotov, Thuan & Lipovetski 1997; Izotov & Thuan 1998) and the chemical abundances of heavy elements in BCDGs (Izotov & Thuan 1999). Subsequent spectroscopic studies were performed by P´ erez-Montero & D´ (2003) iaz and Kniazev et al. (2004). Izotov et al. (1997) indicated the presence of broad low-intensity components in both the H and [O iii] 5007 profiles. These authors also detected the blue WR bump, strongly contaminated by nebular emission, observing both the nebular and broad He ii 4686 emission lines. Hence, Schaerer et al. (1999) included SBS 0926+606 A in their catalogue of WR galaxies. Guseva, Izotov & Thuan (2000) revisited the properties of the massive stars in this galaxy. Until now, there were no spectroscopic data for SBS 0926+606 B. Figure 16 shows our ISIS 4.2m WHT spectrum obtained for SBS 0926+606 A using a long-slit with a PA of 14 . It mainly crosses the subregion A2 defined using our optical data (see Figure 15 in Paper I). As we see in Figure 16, the spectrum has a good spectral resolution, but it only covers between 4200 and 5000 ° in the blue range and between A 5600 and 7400 ° in the red range. Therefore, we did not A observe the [O ii] 3726,29 doublet and the bright [O iii] 5007 emission line. The spectrum obtained for the galaxy B is very noisy and only shows the brightest emission lines. Table A.9 compiles all the emission line fluxes and the main properties of each spectrum. The nebular He ii 4686 emission line is clearly detected in SBS 0926+606 A (Figure 36). Both the blue and red (Figure 37) WR bumps

Table A.10 compiles the values found for the electron temperatures of the ionized gas within these ob jects. For member A, Te (O ii) was computed using the direct method because of the detection of [O iii] 4363 line. Te (O ii) was estimated using Garnett's relation. The spectrum obtained for member B does not allow to calculate R23 , so we used the N2 ratio and the empirical calibrations given by Denicol´ o et al. (2002) and Pettini & Pagel (2004) to estimate Te . As empirical calibrations involving the N2 ratio seem to overestimate the actual abundance (we will discuss this in Paper III), the electron temperatures derived for galaxy B may be underestimated. The electron density, computed using the [S ii] 6717,31 doublet, was always at the low-density limit. For both galaxies, the reddening coefficient was derived using the 3 brightest H i Balmer lines. The comparison of the emission lines ratios with the diagnostic diagrams indicates that both galaxies are starbursts. 3.10.2. Chemical abundances Table A.10 compiles the chemical abundances derived in this galaxy pair. Because the [O ii] 3726,29 doublet is not covered in our spectra, we used [O ii] 7318,7330 fluxes to compute the O+ /H+ ratio. The chemical abundances derived for SBS 0926+606 A are 12+log(O/H) = 7.94±0.08, log(N/O) = -1.45±0.09, log (S/O) = -1.60±0.13 and log(Ar/O) = -2.34±0.13; very similar to those obtained by Izotov & Thuan (1999). The oxygen abundance found for member B using the N2 empirical calibrations is 12+log(O/H)8.15, somewhat higher than that found in galaxy A, but as we commented above, it may be overestimated. The N/O ratios of both galaxies are rather similar, so both ob ject should have suffered a similar chemical evolution. In fact, they are dwarf galaxies with a very similar MB (see Paper I). 3.10.3. Kinematics of the ionized gas Figure 17 shows the position-velocity diagram derived from the slit position observed in SBS 0926+606. We extracted 4 pixel bins (0.8 arcsec) across the H profile. SBS 0926+606 A shows a clear sinusoidal pattern, with and amplitude of around 50 km s-1 , suggesting that the double nucleus we found in this galaxy (see Figures 15 and 16 of Paper I) may be a consequence of an advanced merging process between two ob jects. The northern outskirts of SBS 0926+606 A seems to be partially decoupled from the sinusoidal pattern (there is a difference of around 60 km s-1 with respect to the central velocity). On the other hand, SBS 0926+606 B also shows a perturbed kinematics, because both its northern and southern edges have similar radial velocities. The elongated shape of SBS 0926+606 B, the two tails towards the west detected in our deep images and the disturbed kinematics suggest that the interaction that this galaxy is experiencing ­most probably with


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Fig. 18. ISIS 4.2m WHT spectrum for SBS 0948+532 using a slit with PA 114 . Fluxes are not corrected for reddening. The most important emission lines have been labeled.
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Fig. 17. Position-velocity diagram for the slit position with PA 14 observed in SBS 0926+606 using the H profile. Notice that the y-axis is broken. NE is up. See Figure 15 in Paper I for the identification of the regions.

SBS 0926+606 A­ is very close to the plane perpendicular to the line of sight, and therefore SBS 0926+606 B is observed almost edge-on. In any case, we do not detect any morphological feature, such as the debris of a tidal tail or a diffuse non-stellar ob ject, between both galaxies thus their possible interaction is nowadays no very strong. Finally, the complexity of the position-velocity diagram shown in Figure 17 does not allow to determine the Keplerian mass of the galaxy. Using the H i data given by Pustilnik et al. (2002), we derive MH I =(9.6±3.6)â108 M and MH I =(8.1±3.6)â108 M for A and B, respectively, that indicate mass-to-luminosities ratios of MH I /LB =0.75 and 0.59. The gas depletion timescales are 1.7 Gyr for A and 5.5 Gyr for B. Assuming half of the amplitude of the H i velocity (60 km s-1 for both galaxies) and effective radii of 10 (2.71 kpc) for A and 20 (5.42 kpc) for B, we estimate dynamical masses of MDyn 2.3â109 M and MDyn 4.5â109 M for A and B, respectively. The mass-to-luminosity ratios, MDyn /LB =1.8 and 3.3 for A and B, respectively, which are values typical for BCDGs (Huchtmeier, Krishna & Petrosian 2005). However, the MH I /MDyn ratios are high, 0.42 and 0.18 for A and B, respectively, indicating that a considerable amount of the total mass of the galaxies (42% for A) is neutral hydrogen. All these values indicate that the system still possesses a huge amount of fresh material from which new stars may be born. Indeed, the H i profile obtained by Thuan et al (1999) shows two peaks, that coincide with the optical velocities of the galaxies, embedded in an common H i envelope. This fact strongly suggests that a lot of neutral gas should be found between both galaxies. An HI map obtained using a radio-interferometer would be necessary to study the distribution and kinematics of the neutral gas, giving key clues about the evolution of the system. 3.11. SBS 0948+532 SBS 0948+532 was studied by Izotov and collaborators (Izotov et al. 1994; Thuan et al. 1995; Izotov & Thuan 1998; Guseva et al. 2000; Izotov & Thuan 2004). Schaerer

et al. (1999) included this BCDG in their catalogue of WR galaxies because of the detection of both the broad and nebular He ii 4686 emission lines in the spectra presented by Izotov et al. (1994). The reanalysis performed by Guseva et al. (2000) indicated the presence of WNL stars and a rather noisy red WR bump. Figure 18 shows our ISIS 4.2m WHT spectrum of SBS 0948+532 using a slit position with PA 114 . The emission line fluxes of the detected lines and other properties of the spectrum are compiled in Table A.11. No stellar absorptions are observed in this spectrum. We detect the broad and nebular He ii 4686 emission lines (Figure 36), but the red WR bump (Figure 37) is not seen besides the good S/N ratio of our spectrum. 3.11.1. Physical conditions of the ionized gas The intensity of [O iii] 4363 was used to compute the high ionization electron temperature; the low ionization electron temperature was estimated using Garnett's relation. These values are compiled in Table A.12. The electron density, ne 250 cm-3 , was derived using the [O ii] 3726,3729 doublet. Despite of its high error, the value of the ne estimated from the [Ar iv] 4711/4740 ratio is similar, 260 cm-3 . The reddening coefficient was estimated with a good precision because of the detection of many H i Balmer lines. The equivalent width of the H i Balmer stellar absorption lines, Wabs , is very small, suggesting that the underlying population of evolved stars is not important. The comparison of the observed emission line fluxes with the diagnostic diagrams confirms that the gas is ionized by the strong U V emission of the massive stars. 3.11.2. Chemical abundances Table A.12 lists all the chemical abundances computed for SBS 0948+532. The value of the oxygen abundance is 12+log(O/H)=8.03±0.05. The N/O ratio is log(N/O)=-1.42 ± 0.08, the typical found for ob jects with the metallicity of this galaxy. In general, all our chemical abundances for this BCDG agree with those estimated by Izotov & Thuan (1999) within the errors. 3.11.3. Kinematics of the ionized gas Figure 19 shows the position-velocity diagram obtained for the slit position observed in SBS 0948+532 (see Figure 17


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Fig. 20. IDS 2.5m INT spectrum for SBS 1054+364 using a slit with PA 55 . Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 19 in Paper I for identification of the region.
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Fig. 19. Position-velocity diagram for the slit position with PA 114 observed in SBS 0948+532 using the [O iii] 5007 profile. NW is up (see Figure 17 in Paper I).

C

in Paper I). We analyzed the [O iii] 5007 profile extracting 4 pixel bins (0.8 arcsec). Although the diagram only shows 11 points because of the compact size of this BCDG, a kind of rotation pattern, with an velocity gradient of around 100 km s-1 , is observed. This feature seems to be disturbed in the SE areas, that is precisely the region where the faint arc is seen in our optical images (see Figure 17 of Paper I). Assuming that the kinematics of the galaxy is due to rotation, i=90 and v 50 km s-1 within a radius of 4 (3.63 kpc), we derive a Keplerian mass of MK ep 2.1â109 M and MK ep /L 0.57. Unfortunately, no additional H i or FIR data are available for this galaxy. 3.12. SBS 1054+365 The first spectroscopic study of SBS 1054+365 was performed by Izotov, Thuan & Lipovetski (1997), who detected the broad He ii 4686 emission line. Guseva et al. (2000), Izotov & Thuan (2004) and Buckalew et al. (2005) confirmed the WR feature of this low-metallicity galaxy. This BCDG is included in the latest WR galaxies catalogue (Schaerer et al. 1999). We used the IDS spectrograph attached at the 2.5m INT to get the spectroscopy of SBS 1054+365. The slit position was set at 55 , crossing the main body of the galaxy along its ma jor axis (see Figure 19 of Paper I). Hence, we observed knot b, the central bright region C and part of the star-forming semi-ring located at the west (knot a ). We only analyzed the physical conditions and the chemical abundances of the ionized gas in region C and knot b. Figure 20 shows the spectrum of the center of SBS 1054+365. Table A.11 compiles the dereddened line intensity ratios and other properties of the spectrum. As we see, the spectrum is dominated by nebular emission without any features of stellar absorptions. We do detect the nebular He ii 4686 emission line onto a broad stellar He ii line (Figure 36). The red WR bump is not detected (Figure 37) perhaps because of the lacking of enough S/N in our spectrum.

a

30

40

50

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Relative velocity [ km s ]

Fig. 21. Position-velocity diagram for the slit position observed in SBS 1054+365 using the H profile. NE is up. See Figure 19 in Paper I for identification of the regions.

3.12.1. Physical conditions of the ionized gas The electron temperatures in the central region were computed using the [O iii] 4363 emission line intensity and Garnett's relation between Te (O iii) and Te (O ii). For knot b we used the Pilyugin (2001a,b) empirical calibrations. The electron density was estimated using the [S ii] 6716,6731 doublet, being in the low-density limit in the central region. The values of the reddening coefficient are rather low. The comparison of the [O iii]5007/H , [N ii]6584/H and [S ii]6716,6730/H ratios with the diagnostic diagrams allows to classify all knots as typical H ii regions. 3.12.2. Chemical abundances Table A.12 compiles the chemical abundances derived for SBS 1054+364. The oxygen abundance computed for the center of the galaxy is 12+log(O/H)=8.00±0.07, and its N/O ratio is log(N/O)=-1.41± 0.09, in excellent agreement with the values obtained by Izotov & Thuan (1999). Despite of its higher error, the oxygen abundance and N/O ratio estimated for knot b are very similar to those found at the center of the galaxy. 3.12.3. Kinematics of the ionized gas We used our bidimensional spectrum for the slit position with PA 55 to build the position-velocity diagram shown


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3.13. SBS 1211+540 SBS 1211+540 was included in the study of chemical abundances in BCDGs performed by Izotov and collaborators (Izotov et al. 1991; Thuan et al. 1995; Izotov & Thuan 1998; Guseva et al. 2000; Izotov & Thuan 2004). WR features was firstly reported by Izotov et al. (1994), who detected the nebular He ii 4686 emission line (Schaerer et al. 1999). However, the re-analysis performed by Guseva et al. (2000) only indicates the presence of the broad emission line. Figure 22 shows our ISIS 4.2m WHT spectrum of the center of SBS 1211+540; Table A.11 compiles all its derived properties. The spectrum is dominated by the nebular emission showing no traces of stellar absorptions. We do not detect the blue WR bump or the nebular He ii 4686 emission, although the spectrum has a low S/N ratio (Figure 36). The spectral range where the red WR bump is located was not observed, but this feature should not be expected because of the very low metallicity of SBS 1211+540. 3.13.1. Physical conditions of the ionized gas We derive a very high electron temperature, T (O iii)=17100±600 K, using the direct method (see Table A.12). The low ionization temperature was estimated considering Garnett's relation. Both the [O ii] 3726,3729 and [S ii] 6717,6731 doublets were used to compute the electron density, that is ne = 320 ± 50 cm-3 . The reddening coefficient was determined using all available H i Balmer lines with errors lower than 20%. The comparison of the emission line ratios with the diagnostic diagrams confirms the starbursting nature of this BCDG. 3.13.2. Chemical abundances Table A.12 compiles all the chemical abundances derived for SBS 1211+540. The oxygen abundance, 12+log(O/H)=7.65±0.04, and the N/O ratio, log(N/O)=-1.62 ± 0.10, are in excellent agreement with the values given by Izotov & Thuan (1999). Hence, it is the lowest metallicity ob ject analyzed in this work. The rest of chemical abundances, log(S/O) -1.47 and log(Ne/O) -0.75, are also similar to those determined by these authors. 3.13.3. Kinematics of the ionized gas Figure 23 shows the position-velocity diagram obtained using the bidimensional spectrum of SBS 1211+540. The slit position we used, with a PA of 138 , crosses the center of the galaxy but not knot a that, as it was explained in §3.13.2 of Paper I, also shows nebular emission. We extracted 4 pixel bins (0.8 arcsec) across the [O iii] 5007 profile (the brightest line) and considered the velocity of the center as reference. The position-velocity diagram does not show a clear rotation pattern, only a reverse of the velocity gradient at the center of the galaxy. In any case, the amplitude of the velocity variations are very small. If we do not consider the four lowest points at the SE of Figure 23 where we detected two faint plumes (see Figure 21 of Paper I), we may assume that the kinematics is explained by rotation with a velocity gradient of 20­30 km s-1 . Considering that the gas is rotating with the parameters described above, we may derive the Keplerian

SBS 1211+540 C - PA 138º
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Fig. 22. ISIS 4.2m WHT spectrum for SBS 1211+540 using a slit with PA 138 . Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 21 in Paper 1 for identification of the region.

in Figure 21. We extracted 3 pixel bins (1.2 arcsec) across the H profile and took as reference the brightest region of the galaxy. The diagram does not show a clear rotation pattern but several changes in the velocity distribution. The central region seems to show a velocity gradient of around 40 km s-1 between -5 and 10 . This feature was previously noticed by Zasov et al. (2000) but their lower spatial resolution did not permit to see the small amplitude velocity variations (see their Figure 3b). These authors suggested that this velocity gradient is consequence of the rotation of the galaxy. Our diagram also indicates that the SW region, the partial ring where knot a is located, does not follow the kinematics of the center of the galaxy, showing a velocity variation of 40 km s-1 in 7.2 . On the other hand, the kinematics of the region b, that has an inverted velocity gradient of around 40 km s-1 , also seem to be decoupled from the movement of the gas in the central region, that shows a positive velocity gradient of 25 km s-1 in its NE area. However, the amplitude of all the velocity variations seen in Figure 21 are rather small and cover spatial extensions of the order of several hundred pc. Therefore, it is possible that they could be due to local movements of the bulk of the ionized gas due to the combined action of winds or supernova explosions. In any case, assuming that rotation is present at the center of the galaxy and considering v 20 km s-1 within a radius of 16 (624 pc) and an inclination of i60 (determined using its optical sizes) we have computed a tentative value for the Keplerian mass of MK ep 7.8â107 M , that indicates MK ep /L 1.19. Using the H i data provided by Zasov et al. (2000), we estimated MH I =(6.08±0.59)â107 M and MH I /LB 0.93. Assuming a radius of 35 (1.37 kpc) and the same inclination angle, we estimated a dynamical mass of MDyn 1.5â109 M . The neutral gas to total mass ratio, MH I /MDyn 0.04, is typical for BCDGs but its total massto-light ratio, MDyn /LB 22.9, is quite high (Salzer et al. 2002; Huchtmeier et al. 2005). That may suggest that the dynamic of the system is perturbed, but only an interferometric H i map can confirm this issue. The gas depletion timescale is higher than 2.6 Gyr, indicating that the galaxy still possesses a huge amount of fresh material available to create new generations of stars.


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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 important emission lines have been labeled. See Figure 23 in Paper I for identification of the regions.

relative velocity (km s )

Fig. 23. Position-velocity diagram for the slit position with PA 138 observed in SBS 1211+532 using the [O iii] 5007 profile. NW is up. See Figure 21 in Paper I for identification of the regions.

mass of the galaxy. Assuming v 30 km s-1 within a radius of 5 (315 pc) and an inclination angle of i 50 (value derived from the optical shape of the galaxy), we find MKep 1.13â108 M and MKep /LB 3.6. Using the H i data provided by Huchtmeier et al. (2005), we derive MH I =(2.4±0.4)â107 M and MDyn 1.14â108 M , and hence MH I /MDyn 0.21, MH I /LB 0.76 and MDyn /LB 3.6. The fact that the dynamical mass (determined using radio data) and the Keplerian mass (estimated using our optical data) completely agree indicates that we have probably overestimated the rotation velocity of the ionized gas and/or the actual extension of the neutral gas is much larger that the optical extent. If the first assumption is true, it would indicate that, besides the rotation, there is an additional velocity component, that may be connected with the detection of a very faint plume at the NW in our deep optical images. In any case, the high proportion of neutral mass estimated for this BCDG (21%) and the high value for the gas depletion timescale, 2.5 Gyr, indicate that SBS 1211+540 possesses a huge amount of fresh material available for the birth of new stars. 3.14. SBS 1319+579 The only bibliographic references of SBS 1319+579 are from Izotov and collaborators (Izotov et al. 1997; Izotov & Thuan 1998, 1999; Guseva et al. 2000; Izotov & Thuan 2004). Schaerer et al. (1999) included SBS 1319+579 in their WR galaxies catalogue because Izotov et al. (1997) reported the detection of the broad and nebular He ii 4686 emission lines. Guseva et al. (2000) indicated the presence of WNL and WCE populations in the galaxy. We used the ISIS spectrograph at the 4.2m WHT with a slit position with a PA of 39 to analyze the ionized gas along the main axis of the galaxy (see Figure 23 of Paper I). We got spectroscopic data of regions A, B, C, d and e but we only analyzed A, B and C because of their higher S/N ratio. The spectra of the two brightest regions A and C are

shown in Figure 24, and Table A.13 compiles the dereddened flux ratios for all knots. Region B is the only one that shows some stellar absorptions in its spectra. We do not have a clear detection of the blue WR bump or the nebular He ii 4686 in any region (Figure 36). A careful analysis of the spectrum indicates a tentative detection of both the broad and the nebular He ii in knot A (see Paper III). We do not detect the red WR bump in that region (Figure 37). 3.14.1. Physical conditions of the ionized gas The [O iii] 4363 line is measured in all the regions and therefore we could determine Te (O iii) via the direct method. The low ionization electron temperatures were estimated using Garnett's relation. We found a significant difference in the electron temperatures found for the brightest regions A and C, Te (O iii)13400 and 11500 K, respectively. The electron density, computed using the [S ii] 6317,31 doublet, was always in the low-density limit. The reddening coefficient derived for regions A and C is low and similar to the Galactic reddening. All regions can be classified as starbursts following the results given by the analysis of the diagnostic diagrams. 3.14.2. Chemical abundances Because of the lacking of [O ii] 3726,29 flux values, we used [O ii] 7318,30 to compute the O+ abundance. All the results for the chemical abundances derived in SBS 1319+579 are compiled in Table A.14. The oxygen abundance found in all regions are similar within the errors, 12+log(O/H) 8.10, although that computed in region A, 12+log(O/H)=8.05±0.06, is slightly lower to that found in region C, 12+log(O/H)=8.15±0.07. Knot A shows a high excitation degree, log(O++ /O+ )=0.77, something that is not observed in the other regions. The N/O ratios are very similar in A and B [log(N/O)=-1.53±0.10 in A] but also slightly different than in region C, which has log(N/O)=-1.38±0.10. All the chemical abundances are consistent with those reported by Izotov & Thuan (1999). 3.14.3. Kinematics of the ionized gas Figure 25 shows the position-velocity diagram obtained from our bidimensional spectrum using a slit position with


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Fig. 25. Position-velocity diagram for the slit position observed in SBS 1319+579 using the H profile. The relative intensity of the H emission along the spatial direction is also shown, identifying all observed regions. NE is up. See Figure 23 in Paper I for identification of the regions.

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 are not corrected for reddening. The most important emission lines have been labeled. See Figure 25 in Paper I for identification of the regions.

PA 39 . We extracted 4 pixel bins (0.8 arcsec) across the H profile. The relative intensity of the H emission along the spatial direction in also shown in this figure. Although the velocity continuously decreases from the eastern regions (v -105 km s-1 ) to the western areas (v 65 km s-1 ) of the galaxy, the velocity gradient is not the same across the system. We observe two tendencies: from region C to region B (velocity difference of 40 km s-1 in 30 ) and from region B to region A (velocity difference of 130 km s-1 in 28 ). This behavior may suggest that there is a tidal stream moving from B to A in the direction away from the observer, but our deep images do not show such tail or any morphological feature that support this hypothesis. Another explanation to this feature may be the assumption that they are two systems, as we suggested from the morphology of the H images and the chemical abundances may indicate, with different kinematics and in interaction. If this idea is correct, we should expect to observe distortions in the kinematics of the gas with higher amplitudes that those we see. However, because of the high inclination angle that the galaxy seems to have, i 70 , we cannot discard any of both hypothesis. Considering that the kinematic pattern is consequence of the rotation of the galaxy and assuming i 70 and v 88 km s-1 within a radius of 30 (4.2 kpc), we derive a Keplerian mass of MKep 8.6â109 M . The corresponding mass-to-luminosity ratio, MKep /LB 2.14, is high for that expected for a dwarf galaxy with the properties observed in SBS 1319+579. However, if we consider that only the NE region (from C to B) is rotating with a v 45 km s-1 , we now find MKep 1.7â109 M and MKep /LB 0.42, similar to the values found in other BCDGs (Huchtmeier et al. 2005). This fact seems to confirm that the kinematics surrounding region A are disturbed and not produced by rotation. Using the H i data provided by Huchtmeier et al. (2007), we derive MH I =1.64 â 109 M and MDyn 1.4 â 1010 M , assuming a rotation velocity of 109 km s-1 within 4.5 kpc and the same inclination angle. The neutral gas accounts for only the 12% of all the mass of the system. Its MDyn /LB value, 3.5, is the expected in BCDGs. However, the gas depletion timescale is huge for a starburst galaxy, 10.8 Gyr. This may suggest that the star formation is not very efficient in the system, perhaps

because the H i gas has been expelled from the galaxy. An H i map obtained using a radio-interferometer that includes both SBS 1319+579 and the nearby spiral NGC 5113 would be fundamental to understand the dynamics and evolution of this system. 3.15. SBS 1415+437 The first spectroscopic data of SBS 1415+437 were reported by Thuan et al. (1995), who determined an oxygen abundance of 12+log(O/H)=7.51, being one of the less-metallicity galaxies known. A later reanalysis 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 emission lines, and therefore SBS 1415+437 was included in the latest WR galaxies catalogue (Schaerer et al. 1999). Subsequent spectroscopic analysis were published by Melbourne & Salzer (2002); Melbourne et al. (2004); Guseva et al. (2003); Izotov & Thuan (2004) and Lee, Salzer & Melbourne (2004). Figure 26 shows the spectra of regions A and C obtained using the instrument ISIS at the 4.2m WHT and a slit with PA 20 that crosses the main body of the galaxy (see Figure 25 of Paper I). Although we detected some emission lines in knot B, we have not analyzed its properties because of the low S/N ratio of its spectrum. All spectra are dominated by nebular emission; no stellar absorptions are detected. Table A.13 compiles all the line intensities ratios and other properties of the spectra. Although we do not see the broad blue WR bump, the nebular He ii 4686 emission line is well detected in the spectrum of region C (Figure 36). We do not observe the red WR bump (Figure 37) besides the good S/N ratio and spectral resolution. 3.15.1. Physical conditions of the ionized gas The electron temperatures were computed using the direct method and are very high, T (O iii)=16400 and 15500 K for C and A, respectively. Te (O ii) was estimated using Garnett's relation. The electron density was derived using the [S ii] 6716,31 doublet and was below the low-density limit. The reddening coefficient found in region C is extremely low, c(H )0.01, and identical to that determined by Guseva et al. (2003). However, the higher value of c(H ) in region A, c(H )=0.16, suggests an inhomogeneous dis-


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Fig. 28. IDS INT spectra for the regions A (bottom ) and B (top ) of IRAS III Zw 107. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 27 in Paper I for identification of the regions.

Relative velocity [ km s ]

Fig. 27. Position-velocity diagram for the slit position with PA 20 observed in SBS 1415+437 using the H profile. NE is up. See Figure 25 in Paper I for identification of the regions.

tribution of dust in the galaxy. The comparison of the observed [O iii]5007/H and [N ii]6584/H ratios with the predictions given by the diagnostic diagrams confirm their starbursting nature. 3.15.2. Chemical abundances Table A.14 compiles our results of the chemical abundances derived in SBS 1415+437. These data confirm the very low-metallicity of the galaxy, being the oxygen abundances 12+log(O/H)=7.58±0.05 (for C) and 7.61±0.06 (for A). The N/O ratio for both ob jects, log(N/O)=-1.57 ± 0.08, is the expected for such low metallicities. Our results are in very good agreement with the abundances obtained by Guseva et al. (2003). 3.15.3. Kinematics of the ionized gas The position-velocity diagram shown in Figure 27 was obtained extracting 4 pixel bins (0.8 ) along the H profile of our bidimensional spectrum. We took the velocity of the brightest ob ject C as reference. We observe that the velocity continuously decreases from the SW regions (v 30 km s-1 , where region A is located) to the NE areas (v -30 km s-1 , where region B is found), that may be attributed to the rotation of the galaxy. Some kinematic divergences are detected between regions A and C. However, because of the low amplitude of such variations (less than 15 km s-1 ), they may just be a consequence of local movements in the ionized gas. Our position-velocity diagram is similar in both shape and values to that obtained by Thuan, Izotov & Foltz (1999) using a slit with a PA of 22 (see their Figure 10). They also reported the peculiar kinematic behavior we observe between A and C. The Keplerian mass we estimate for this galaxy, assuming a rotation velocity of 30 km s-1 within a radius of 25 (1.13 kpc) and an inclination angle of i 75 (determined using the shape of the galaxy we see in our optical images), is MKep 2.5â108 M , and its mass-to-luminosity ratio MKep /LB =2.5. The H i mass

estimated by Huchtmeier, Krishna & Petrosian (2005) is MH I =(9.64±0.65)â107 M . Using their data of WH I and considering a radius of 40 (1.8 kpc), we estimate a dynamical mass of MDyn =4.9â108 M . With these data, we derive MH I /LB =0.96, MDyn /LB =4.9 and MH I /MDyn =0.20, that are the typical values found for BCDGs (Salzer et al. 2002; Huchtmeier et al. 2005). These estimations are more reliable that those given by Thuan, Izotov & Foltz (1999) because we are using recent data with a higher sensibility. Both the gas depletion timescale (3.2 Gyr) and the fact that 1/5 of the mass of the system is neutral gas indicate that SBS 1415+437 possesses a huge reservoir of fresh material available for new star-forming phenomena. 3.16. III Zw 107 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 spectral region of the blue WR bump in the southern ob ject, and hence Schaerer et al. (1999) included this BCDG in their catalogue of WR galaxies. A slit with a PA of 0 was used in the IDS spectrograph at the 2.5m INT to observe III Zw 107 (see Figure 27 of Paper I). Three different regions, A, B and C, were extracted. The optical spectra of regions A and B are shown in Figure 28. Region A possesses important underlying stellar absorption features. The faint region C is not very evident from our optical images, but it is clearly identified at the north of region B in our bidimensional spectrum. Table A.15 compiles all the emission intensity ratios and other properties of the spectra of III Zw 107. We clearly detect the broad He ii 4686 feature in the spectrum of knot A (Figure 36), the same region where Kunth & Joubert (1985) indicated as WR-rich. However, we do not see the red WR bump (Figure 37) perhaps because of the relatively low spectral resolution of these 2.5m INT spectra. [O iii] 4959,5007 and H show broad wings in their profiles, more evident in the spectrum of region A. 3.16.1. Physical conditions of the ionized gas In region A we computed both Te (O iii) and Te (O ii) using the direct method because of the detection of [O iii] 4363 and the [O ii] 7319,7330 doublet. The values are Te (O iii)=10900±900 K and Te (O ii)=10500±800 K. The


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Fig. 29. Position-velocity diagram for the slit position with PA observed in III Zw 107 using the H profile. N is up. See Figure 27 in Paper I for identification of the regions.

electron temperatures for knots B and C were estimated using empirical calibrations. All results are compiled in Table A.16. The electron density was computed using the [S ii] 6716,31 doublet and it was in the low-density limit for regions B and C. The comparison of the spectra shown in Figure 28 indicates that the spectral energy distribution of the continuum in region A is less steeper than that observed in region B. This effect may be explained by the contribution of a more evolved stellar population in A. However, the analysis of the reddening coefficient using the H i Balmer lines of this regions gives a much higher value in region A, c(H )0.68, than in region B, c(H )0.15, and hence it seems that a difference in extinction is the explanation to the different slope of the continuum. The diagnostic diagrams for all regions agree with the loci of typical H ii regions. 3.16.2. Chemical abundances Table A.16 lists all the chemical abundances computed for the bursts analyzed in III Zw 107. The oxygen abundance of the region A, derived using the direct method, is 12+log(O/H)=8.23±0.09. This value is similar to that reported by Kunth & Joubert (1985), but more than 0.3 dex higher than that reported by Gallego et al. (1997). The oxygen abundances of regions B and C were estimated using the Pilyugin (2001a,b) empirical calibrations, yielding values of 12+log(O/H)8.31, similar to the abundance of region A within the errors. The rest of the chemical abundances are also similar in all regions, just slightly lower in A. The results of the chemical abundances found between B and C are essentially identical besides all the uncertainties involved in their determination. 3.16.3. Kinematics of the ionized gas The position-velocity diagram obtained using our bidimensional spectrum is shown in Figure 29. We extracted 3 pixel bins (1.2 arcsec) along the H profile, taking as reference the velocity of region A, the brightest knot. We observe a

negative velocity gradient between the southern regions of the galaxy (30 km s-1 ) and region A ( -20 km s-1 ), but between this knot and region B a reverse of the velocity of 40 km s-1 is found within 4 . The velocity difference between region B and C is -40 km s-1 . Hence, although the velocity amplitudes are not large and the spatial resolution is not very high, the position-velocity diagram seems to show a sinusoidal pattern This feature may suggest interaction or merging phenomena between the two brightest knots seen in III Zw 107. This hypothesis would explain the existence of the tail found in the deep optical images (see Figure 27 of Paper I). It is also possible that the velocity gradient observed at the south of the galaxy is a consequence of the movement of the ionized gas within/towards that tail. We have performed a tentative estimation of the Keplerian mass of III Zw 107 using the positionvelocity diagram shown in Figure 29. Considering that we observe the galaxy edge-on (i 90 ) and assuming v 30 km s-1 within a radius of 10 (3.9 kpc), we estimate MKep 8.2â108 M and MKep /LB 0.05. These values are very low compared with the neutral gas mass and the dynamical mass derived using the radio data (Paturel et al. 2003), MH I =(6.7±1.2)â109 M , MDyn 1.8â1010 M (the dynamical mass was estimated using a radius of 20 =7.8 kpc and the half of the H i width, 100 km s-1 ), and the mass-to-luminosity ratios derived from them, MH I /LB =0.38 and MDyn /LB 1. Indeed, if all these values are right, around 37% of the mass of the system is neutral gas. We consider that, because of the detection of an important population of old stars within the galaxy and its relatively high metallicity, the dynamical mass of III Zw 107 has been probably underestimated. Hence, the H i distribution should be several times larger than the optical extent. The comparison of the velocity amplitudes between the optical (30 km s-1 ) and the radio (100 km s-1 ) data strongly supports this idea. Perhaps, the neutral gas has been expelled and/or dispersed as a consequence of the possible interaction or merging between the two main ob jects observed in III Zw 107. An interferometric H i map of this galaxy is needed to answer to all these issues. In any case, the gas depletion timescale is high, 3.5 Gyr, indicating that there is still a lot of neutral gas available to form new stars. 3.17. Tol 9 The WR nature of Tol 9 has been controversial. Penston et al. (1977) indicated a probable detection of a faint emission line around 4686. Kunth & Schild (1986) did not find the blue WR bump or the He ii emission line but suggested the detection of the red WR bump. Both Conti (1991) and Schaerer et al. (1999) included Tol 9 in their lists of candidate WR galaxies. We observed two slit positions of Tol 9 (see Figure 29 of Paper I). For the first one, we used the IDS instrument at the 2.5m INT, using a slit with a PA of 49 that crosses the center of Tol 9 and the dwarf companion galaxy located at the SW. The second slit position was taken with the ALFOSC instrument at the 2.56m NOT, choosing a slit position almost perpendicular to that used at the 2.5m INT. This last slit position was centered at the SW of the center of the galaxy and the PA was set to 109 , our main objective was to analyze the kinematics and properties of the

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Fig. 30. IDS INT spectrum (bottom ) using a slit with PA 49 and ALFOSC 2.56m NOT spectrum (top ) using a slit of PA 109 of Tol 9. Fluxes are not been corrected for reddening. The most important emission lines have been labeled. See Figures 29 and 30 in Paper I for identification of the regions.

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filamentary structure of ionized gas found in our deep H images (see Figure 30 of Paper I). The list of all the emission lines observed using both slits, as well as other important properties of the spectra, are shown in Table A.15. As we see, both spectra show very similar line intensities. Notice that the radial velocity obtained for Tol 9 using our optical spectra is 200 km s-1 more positive than the value previously reported and listed in the NED (Lauberts & Valentijn 1989). The spectra obtained for Tol 9 show nebular emission and a continuum dominated by stellar absorptions in the H i Balmer lines. We also observe an important decrement in the continuum at the blue range of the spectra. This feature can be explained by both the contribution of the older stars and an high extinction. We detect both the blue WR bump and the nebular He ii 4686 emission line in the 2.5m INT spectrum (Figure 36), indicating the presence of WNL stars at the center of Tol 9. However, we do not have a clear detection of the red WR bump (Figure 37) although there are some evidences that it is there. If WNL stars are seen in Tol 9, WCE stars should also exist because of the high metallicity of the ionized gas. Deep spectroscopy with higher S/N ratio and spectral resolution is therefore needed to confirm this feature. 3.17.1. Physical conditions of the ionized gas The 2.5m INT spectrum of Tol 9 shows [O iii] 4363, [N ii] 5755 and the [O ii] 7319,7330 doublet, and hence we derived the electron temperatures using the direct method. The results (that are compiled in Table A.16) are Te (O iii)7600 K and Te (low)8300 K. For the spectrum obtained using the 2.56m NOT data we computed Te (O ii) using the [O ii] lines and estimate the Te (O iii) using Garnett's relation. In all cases, the electron densities were in the low-density limit. As we said, the values for the reddening coefficients are high, about c(H )0.45. The absorption equivalent widths in the H i Balmer lines, computed interactively with c(H ), are large and give values of Wabs 6­8 ° The comparison of the data with the diA. agnostic diagrams classifies all regions as starbursts, but it seems that there is a small shock contribution to the ionization of the gas in the region analyzed with the 2.56m NOT spectrum.

Fig. 31. Position-velocity diagrams for the slit positions observed in Tol 9: PA 49 (left, using 2.5m INT data) and PA 109 (right, using 2.56m NOT data). Both the H (circles) and the [O iii] 5007 (triangles) profiles were analyzed. Notice that the y-axis is broken in two parts in the right diagram. NE is up in the left diagram and NW is up in the right diagram. See Figures 29 and 30 in Paper I for identification of the regions.

3.17.2. Chemical abundances Table A.16 compiles all the chemical abundances computed in Tol 9, showing almost identical results for both spectra. The average value of the oxygen abundance derived using the direct method is 12+log(O/H)=8.57±0.10. This value is around 0.8 dex higher than that provided by Kunth & Schild (1986). The average N/O ratio, log(N/O)=-0.81±0.11, is the expected for a galaxy with the oxygen abundance found in Tol 9. The rest of the chemical abundances computed for this galaxy averaging both set of data are log(S/O)=-1.62±0.12, log(Ne/O)=-0.72±0.14 and log(Ar/O)=-2.45 ± 0.15. 3.17.3. Kinematics of the ionized gas The kinematics of the ionized gas in Tol 9 were analyzed using our bidimensional spectra. We extracted 3 pixel bins (1.2 ) and 5 pixel bins (0.95 ) along the H and the [O iii] 5007 profiles for the 2.5m INT and the 2.56m NOT spectra, respectively. The reference velocity was always chosen in the brightest region. The position-velocity diagrams are shown in Figure 31. Although the diagrams have an excellent agreement between the H and the [O iii] 5007 results, their interpretation is not easy. First, the diagram with AP 49 , that crosses the center of Tol 9, does not show a rotation pattern. Indeed, we observe two velocity gradients: while the velocity changes 120 km s-1 from the NE regions to the center, this tendency is completely reversed in the SW region, that shows a velocity variation of -120 km s-1 . Notice that the center of Tol 9 is not located in the (0,0) position of this diagram because the maximum of the H emission is displaced 10 towards the SW of the center. Perhaps the velocity pattern we seen at this PA is a combination of rotation in the center and NE areas and the velocity gradient produced by the optical tail we detected in our deep optical images that connects Tol 9 with a dwarf companion galaxy located at the SW (see Figure 29 of Paper I). Assuming v 70 km s-1 within a radius of 12 (2.5 kpc) and an inclination angle of i 50 , we


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
80 60
[O II] [Ne III] H H H

estimate a Keplerian mass of 1.2â109 M and a massto-luminosity radio of MKep /LB 0.16. This last value is very low compared with that found in similar ob jects, and therefore we think that the mass of the galaxy has been underestimated. Interferometric H i measurements are needed to get a reliable estimation of the dynamical mass of this galaxy because single-dish observations (such those provided by HIPASS) would blend the H i gas in Tol 9 and the nearby spiral ESO 436-46. This analysis will be performed using the new H i ATCA data for the galaxy group where Tol 9 resides obtained by our group (L´ ez-S´ hez op anc et al. 2010). Using the FIR fluxes, the warm dust mass is Mdust 2.41â106 M . Following the analysis performed by Bettoni et al. (2003), the derived mass-to-luminosity ratio, Mdust /LB 3.1â10-4 , agree with the typical value found in spiral galaxies (Mdust /LB 2â10-4 ). On the other hand, the position-velocity diagram obtained using the slit with PA 109 , shows several reverses in the velocity of the ionized gas. Regions #3 and #5, that we identified as two filamentary structures using our deep H map (see Figure 30 of Paper I), are labeled in this diagram. They show a similar kinematic behavior (a variation of -90 km s-1 in their velocities within the same distance, 12 (2.5 kpc), with respect to the center of the system. This kinematic structure, which is not coincident with any stellar distribution, reminds an expanding bipolar bubble, reinforcing the hypothesis that the H envelope surrounding Tol 9 is consequence of some kind of galactic wind. Knot #6 is barely detected at around 100 (21 kpc) from the maximum of H emission, but it seems to show a radial velocity similar to that observed at the ending of the filament #5. This fact suggests that knot #6 has been kinematically coupled to the main filamentary structure. Maybe, knot #6 has been expelled from the expanding bubble. 3D optical spectroscopy is needed to clarify all these issues, as well as to compare the kinematics of the stellar and ionized gas components. 3.18. Tol 1457-262 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 their WR galaxies catalogue because Contini (1996) detected the broad He ii 4686 emission line in the brightest region of the western ob ject. This feature was also reported by Pindao (1999). Both authors also detected WR features in another region of the same ob ject, although while Contini (1996) observed the nebular He ii 4686 emission line, Pindao (1999) only reported the detection of the blue WR bump. Figure 32 shows the spectra extracted for regions A, B and C of the western ob ject in Tol 1457-262 (Object 1, see Figure 31 of Paper I) using the instrument ALFOSC at the 2.56m NOT and a slit with a PA of 155 . Table A.17 compiles all the line intensities ratios and other properties of all the spectra analyzed in this galaxy. Stellar absorptions are barely detected, indicating that the nebular emission strongly dominates their spectra. Broad low-intensity wings are detected in the H profile in region A. We observe the nebular He ii 4686 emission line in the spectrum of region A; this feature is also observed in the spectrum of region B (Figure 36). We do not see the blue or the red

Tol 1457-262 A PA 155º
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Fig. 32. ALFOSC 2.56m NOT spectra of the regions analyzed in Tol 1457-262. Fluxes are not corrected for reddening. The most important emission lines have been labeled. See Figure 31 in Paper I for identification of the regions.

(Figure 37) WR bumps in any region besides the good S/N of the spectra. 3.18.1. Physical conditions of the ionized gas The electron temperatures were computed using the direct method because of the detection of [O iii] 4363 and [O ii] 7318,7330 lines in all the spectra. The derived values, that are compiled in Table A.18, agree well with the empirical relation between Te (O ii) and Te (O iii) provided by Garnett (1992). Despite their similar ionization degree, region B has an electron temperature, Te (O iii)15200 K, that is higher than that found in regions A and C, Te (O iii)14000 K. The electron densities found in regions A and C are similar, ne 200 cm-3 , but ne was in the low-density limit in region B. The determination of the reddening coefficient for region A was done using 5 H i Balmer ratios, yielding consistently a very high value, c(H )=0.83±0.03. However, the c(H ) found in region B gave a negative value. This result may not attributed to a bad flux calibration because adjacent regions A and C do not have this problem. Hence, we assumed c(H )0 in region B4 and scaled the blue and red spectra considering the theoretical ratio between the H and H fluxes for the electron temperature estimated for this region. The comparison of the emission line ratios with the diagnostic diagrams indicates that all regions can be classified as starbursts. 3.18.2. Chemical abundances The chemical abundances computed for the regions analyzed in Tol 1457-262 are listed in Table A.18. The oxygen abundance derived for the brightest knot (region A) is 12+log(O/H)=8.05±0.07, similar to that found in adjacent region C, 12+log(O/H)=8.06±0.11, but higher than the oxygen abundance computed in region B, 12+log(O/H)=7.88±0.07. This would suggest that regions A and B have experienced different chemical evolution.
4 However, the Galactic value using Schlegel, Finkbeiner & Davis (1998) is c(H )=0.23±0.02.


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-1

H emission
200 250 0 2000 4000 6000

Wavelength [ å ]

Relative velocity [ km s ]

Flux [ counts ]

Fig. 33. Position-velocity diagram for the slit position observed in Object 1 of Tol 1457-262 using the H profile. The relative intensity of the H emission along the spatial direction is also shown, identifying all observed regions. NW is up. See Figure 31 in Paper I for identification of the regions.

Fig. 34. ALFOSC 2.56m NOT spectra for the galaxy pair ESO 566-8 (top ) and ESO 566-7 (bottom ) that constitute Arp 256. Fluxes are not been corrected for reddening. The most important emission lines have been labeled. See Figure 34 in Paper I for identification of the galaxies.

However, the rest of the chemical abundances are relatively similar in all regions and only slightly higher in A and C. Remarkably weak are the [N ii] emission lines, that give log(N/O)= -1.57±0.11 and -1.61±0.12 for A and B, respectively. Our results do not agree with those reported by Masegosa et al. (1994) in region A, for which they computed 12+log(O/H)=8.23. These are the first abundance determinations for the rest of the ob jects. 3.18.3. Kinematics of the ionized gas Figure 33 shows the position-velocity diagram of the slit position with PA 155 analyzed in Object 1 in Tol 1457262. We extracted 5 pixel bins (0.95 ) along the H profile and took as reference the center of region A (the maximum of the H emission). Figure 33 includes the relative intensity of the H emission along the spatial direction. Although a velocity gradient between the northern regions (v -120 km s-1 ) and the southern areas (v 190 km s-1 ) is seen, we notice several velocity reverses along the system, being the most important that found between regions A and C, that has an amplitude larger than 100 km s-1 . These features indicate that the kinematics of the different star-forming regions found within Object 1 of Tol 1457-262 are decoupled. Indeed, the observed sinusoidal pattern suggests that the system is experiencing a merging process. The fast increasing of the velocity between knot B and the southernmost regions may be related to the movement of the material within the faint tail we detected in our deep optical images (see Figure 31 of Paper I). We consider that knot C is a TDG candidate because of its kinematics and chemical abundances are similar to those derived in bright region A. Assuming that the general kinematics pattern is consequence of rotation, we derived a tentative Keplerian mass of MKep 6.2â109 M and a mass-to-luminosity ratio of MKep /LB 0.34 for Object 1 in Tol 1457-262. We considered v 60 km s-1 within a radius of r 15 (4.95 kpc) and an inclination angle of i 55 (from our optical images). HIPASS provides a detection of the H i gas in this galaxy, for which we derive MH I =4.7â109 M . However, this estimation for the neutral gas mass is for all the objects 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 luminosity of the system, computed from our optical data, is LB =1.82â1010 L , and hence the neutral hydrogen massto-luminosity radio of Tol 1457-262 is MH I /LB 0.26. This value is higher than the typical MH I /LB found in similar galaxies, indicating the large amount of neutral gas in Tol 1457-262. 3.19. Arp 252 Arp 252 is a pair of interacting galaxies designed ESO 566-8 (galaxy A) and ESO 566-7 (galaxy B). Their spectroscopic properties were analyzed by Pena, Ruiz & ~ Maza (1991) and Masegosa, Moles & del Olmo (1991). These authors detected the blue WR bump in ESO 566-75 . The WR feature was confirmed by Pindao (1999), and hence ESO 566-7 was included in the latest catalogue of WR galaxies (Schaerer et al. 1999). Contini (1996) reported a tentative detection of the nebular He ii 4686 emission line in ESO 566-8, and therefore it was listed as a suspected WR galaxy (Schaerer et al. 1999). Figure 34 shows the optical spectra of ESO 566-8 and ESO 566-7 obtained using the instrument ALFOSC at the 2.56m NOT using a slit with PA 342 (see Figure 346 in Paper I). The spectrum of ESO 566-8 shows many emission lines, but the spectrum of ESO 566-7 only shows a few. All line intensity ratios are listed in Table A.17. Although the spectra are dominated by the emission of the ionized gas, they also show some stellar absorptions, that are more evident in the weakest H i Balmer lines observed in ESO 566-7. We detect the nebular He ii 4686 emission line onto a faint broader feature in the spectrum of ESO 5668 (Figure 36). The red WR bump is also observed in this ob ject (Figure 37), confirming the presence of both WNL and WCE stars in this galaxy. Besides both Masegosa et al. (1991) and Pindao (1999) reported a detection of the blue WR bump in ESO 566-7, our spectrum does not show this feature. This may be consequence of the slit position chosen to get the spectrum of this galaxy, that does not cross along
5 Masegosa et al. (1991) named this ob ject C 0942-1929A, but it is incorrect following Schaerer et al. (1999). 6 Notice that the real slit position in Figure 34 of Paper I is 342 = -18 and not 18 .


26

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
15 10 5

the ma jor body of ESO 566-7 but almost perpendicular to it (see Figure 34 in Paper I). 3.19.1. Physical conditions of the ionized gas
Distance [ arcsec ]

Arp 252 - PA 342º

The detection of the [O ii] 7318,7329 doublet and the [N ii] 5755 emission line in the spectrum of ESO 566-8 allowed the direct determination of the low ionization electron temperature, that gave values of Te (O ii)9300 K and Te (N ii)9000 K. Averaging these numbers, we estimate Te (low)=9100±800 K. The high ionization electron temperature was computed using Garnett's relation. The electron temperatures in ESO 566-7 were estimated via empirical calibrations. The electron densities were derived using the [S ii] 6617,6730 doublet. All results are tabulated in Table A.18. The reddening coefficient are c(H )0.49 in ESO 566-8 and c(H )0.27 in ESO 566-7. The comparison of the observed emission line ratios with the diagnostic diagrams indicates that the nature of the ionization of the gas in ESO 566-7 is photoionized, but some shocks contribution seems to be present in ESO 566-8. 3.19.2. Chemical abundances Table A.18 compiles all the chemical abundances derived for the galaxy members of Arp 252. The oxygen abundance in ESO 566-8 (galaxy A) is 12+log(O/H)=8.46±0.11 and was computed using the direct method. The empirical calibration provided by Pilyugin (2001a) suggests that the oxygen abundance in ESO 566-7 (galaxy B), 12+log(O/H)=8.50±0.16, similar to that found in ESO 566-8. The N/O ratio is also similar in both galaxies, log(N/O)= -0.76±0.12 in ESO 566-8. The value of the oxygen abundance in ESO 566-7 reported by Masegosa et al. (1991), 12+log(O/H)=8.54, agrees with our estimations. 3.19.3. Kinematics of the ionized gas Figure 35 shows the position-velocity diagram obtained for Arp 252 analyzing the bidimensional spectrum using a slit with PA 342 . The H profile was analyzed extracting 5 pixel bins (0.95 ) and taking as reference the maximum of the emission in ESO 566-8. As we see, both galaxies have a slight velocity difference of 50 km s-1 between their centers. However, we appreciate important differences in the kinematics of the galaxies: while A seems to be rotating in its central regions (there is a velocity variation of -130 km s-1 at the north and 100 km s-1 at the south), B shows a distorted kinematic pattern. The velocity gradients observed in the upper and lower regions of galaxy A may correspond to the tidal streams induced in the long tails we observe in the optical images (see Figure 34 of Paper I). We performed a tentative determination of the Keplerian mass of the galaxies. We assumed a velocity of v 100 km s-1 within a radius of r 5 (3.15 kpc) in ESO 566-8 and a velocity of v 30 km s-1 within a radius of r 3 (1.89 kpc) in ESO 566-7. For both we considered an inclination angle of i = 90 , hence our Keplerian mass determinations are low limits to the real ones. We consider, however, that this assumption is not bad because in ESO 566-8 the northern tidal tail shows a high inclination angle with respect to the plane of the sky and ESO 566-7 shows its long southern tail almost in the plane

0 -5

ESO 5 6 6 - 8
(Galaxy A)

-25 -30 -35 -40 -150

ESO 5 6 6 - 7
(Galaxy B)

-100

-50

0

50
-1

100

150

200

Relative velocity [ km s ]

Fig. 35. Position-velocity diagram for the slit position observed in Arp 252 using the H profile. Notice that the y-axis is broken in two parts. N is up. See Figure 34 in Paper I for identification of the galaxies.

of the sky and its disk seems to be edge-on. We estimated Keplerian masses of MKep 7.3â109 M for ESO 566-8 and MKep 4.0â108 M for ESO 566-7, that indicate a mass-to-luminosity ratios of MKep /LB 0.21 and 0.05 for ESO 566-8 and ESO 566-7, respectively. We consider that the Keplerian mass in ESO 566-7 has been highly underestimated. The warm dust mass of Arp 252, computed using the FIR fluxes, is Mdust 6.3â106 M . Arp 252 is not detected in HIPASS, and hence we can not derive the neutral gas mass and the dynamical mass. However, considering the absolute magnitude of the main galaxy, MB = -20.9, and despite the distance to the system (D 130 Mpc) we should expect some H i emission. Hence, or Arp 252 does not have too much neutral gas or it has been lost in the intergalactic medium because of tidal effects. Finally, it would be very interesting to analyze the kinematics of knots c and d to check their probable TDG nature (see Figure 34 of Paper I). 3.20. NGC 5253 The echelle spectrophotometric analysis of the BCDG NGC 5253 was presented in L´ ez-S´ hez et al. (2007). op anc We measured the intensities of a large number of permitted and forbidden emission lines in four zones of the central part of the galaxy. The physical conditions of the ionized gas were derived using a large number of different line intensity ratios. Chemical abundances of He, N, O, Ne, S, Cl, Ar, and Fe were determined following the standard methods. We detected, for the first time in a dwarf starburst galaxy, faint C ii and O ii recombination lines. We confirmed the presence of a localized N enrichment in certain zones of the center of the galaxy and suggested a possible slight He overabundance in the same areas. We shown that the enrichment pattern agrees with that expected for the pollution by the eyecta of WR stars. The amount of enriched material needed to produce the observed overabundance is consistent with the mass lost by the number of WR stars estimated in the starbursts.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

27

The analysis of the H i data provided by the Local Volumen H i Survey pro ject (Koribalski 2008) reveals that the neutral gas kinematics within NGC 5253 has a velocity gradient along the optical minor axis of the galaxy; it does not show any sign of regular rotation (L´ ez-S´ hez et al. op anc 2008). Some authors suggested that this feature is an outflow, but most likely its origin is the disruption/accretion of a dwarf gas-rich companion (Kobulnicky & Skillman 2008) or the interaction with another galaxy in the M 83 subgroup. The finding of a distorted H i morphology in the external parts of the galaxy supports this hypothesis. A comprehensive analysis of the neutral gas within NGC 5253 will be presented elsewhere (L´ ez-S´ hez et al. 2010). op anc

Table 2. Comparison between the oxygen abundance of the regions analyzed here and their previous estimations compiled from the literature. The second column indicates if Te was computed using the direct method (D) or via empirical calibrations (EC) in this work.
12+log(O/H) This work Previous W. 8.22±0.05 8.22± 8.14± 8.13± 8.07± 8.03± 8.15± 8.3± 8.57± 8.23± 8.54± 8.37± 8.10± 8.21± 8.75± 8.46± 8.07± 7.89± 8.33± 8.64± 8.45± 8.38± 8.03± 8.03± 7.95± 7.94± 8.15± 8.03± 8.00± 8.13± 7.65± 8.05± 0.10 0.08 0.09 0.06 0.10 0.07 0.2 0.10 0.10 0.10 0.10 0.06 0.14 0.12 0.13 0.04 0.17 0.08 0.15 0.10 0.10 0.04 0.14 0.05 0.08 0.16 0.05 0.07 0.16 0.04 0.06 8.31 8.3±0.2 ... 8.34±0.20 ... 8.1±0.2 ... ... ... 8.55 ... ... 8.33 ... ... 8.19±0.18 9.13 ... 8.04±0.04 ... ... ... ... ... 7.97±0.02 ... 7.93±0.05 7.95±0.01 ... 8.00±0.01 7.97±0.02 ... 7.64±0.01 8.13±0.01 8.09±0.03 7.95±0.10 8.15±0.03 8.11±0.01 7.61±0.01 7.62±0.03 8.20 7.90 ... 7.73 8.23 ... ... ... 8.54 8.12±0.06 8.19±0.07 8.16±0.12 ...

Galaxy HCG 31 AC HCG HCG HCG HCG HCG HCG HCG Mkn Mkn Mkn Haro Haro Haro Mkn 31 A1 31 B 31 E 31 F1 31 F2 31 G 31 H 1087 1087 N 1087 #7 15 C 15 A 15 B 1199

Te D EC D D D D D EC EC EC EC EC D EC D EC D EC D EC EC EC D EC D D EC D D EC D D D D D D D EC D D D D D EC D D D D

Ref VC92 R03 ... R03 ... R03 ... ... ... VC92 ... ... S05 ... ... IT98 GIT00 ... IT99 ... ... ... ... ... KS96 ... IT98 IT98 ... IT98 IT98 ... IT98 ITL97 IT99 ITL97 ITL97 IT99 G03 G03 KJ85 G97 ... KS86 M94 ... ... ... M91 K97 K97 K97 ...

4. Summary
We have presented a detailed analysis of the ionized gas within 16 Wolf-Rayet galaxies using long-slit intermediateresolution optical spectroscopy. In many cases, more than two star-forming regions have been studied per galaxy. We have analyzed the physical properties of the ionized gas, deriving their electron temperatures, electron density, the reddening coefficient and the stellar absorption underlying the H i Balmer lines. We have confirmed that the excitation mechanism of the ionized gas in all bursts is mainly photoionization and not due to shock excitation as it happens in AGNs and LINERs. In the ma jority of the cases, we have computed the chemical abundances of O, N, S, Ne, Ar and Fe using the direct determination of the electron temperature (see second column in Table 2). When these data were not available, we used the empirical calibration of Pilyugin (2001a,b) to get an estimation of the metallicity of the ionized gas. We have estimated the oxygen abundance of many new regions within the sample galaxies and refined the chemical properties of some of them, remarking regions 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 usually in agreement with previous observations reported in the literature. Table 2 compares our oxygen abundance determinations with those previously reported in the literature. As we see, the ma jority of the results agree well with previous estimations, but there are important differences in Mkn 1199, III Zw 107, Tol 9 and Tol 1457-262. Including the data of the four systems analyzed in previous papers, a very useful database of ob jects with oxygen abundances between 7.58 and 8.75 in units of 12+log(O/H) is provided. In Papers IV and V we will explore this database comparing their properties with other data derived from both our deep optical/NIR images and other multiwavelength observations available in the literature. We have confirmed the detection of Wolf-Rayet features in the ma jority of the galaxies, as we should expect because our sample was extracted from the latest WR galaxies catalogue (Schaerer et al. 1999). We have reported the detection of broad WR features in 20 regions within 16 systems. Figure 36 shows a detail of the optical spectrum in the 4600­4750 ° range (blue WR bump) of all important obA jects; faint regions with very low S/N have been excluded. We have indicated the spatial localization of the massive stars in each system. WR features are sometimes found in different knots within a same galaxy (HCG 31, Haro 15, Tol 1457-262, NGC 5253). The He ii 4686 emission line is unambiguously detected in 14 regions (HCG 31 AC and

Mkn 1199 NE Mkn 5 A Mkn 5 B IRAS 08208+2816 C IRAS 08208+2816 #8 IRAS 08339+6517 IRAS 08339+6517 c. POX 4 POX 4 comp. UM 420 SBS 0926+606 A SBS 0926+606 B SBS 0948+532 SBS 1054+365 SBS 1054+365 b SBS 1211+540 SBS 1319+579 Aa SBS 1319+579 Ba SBS 1319+579 Ca SBS 1415+437 C SBS 1415+437 A III Zw 107 A III Zw 107 B,C Tol 9 Tol 1457-262 A Tol 1457-262 B Tol 1457-262 C ESO 566-8 ESO 566-7 NGC 5253 A NGC 5253 B NGC 5253 C NGC 5253 D
a

8.12±0.10 8.15±0.07 7.58±0.05 7.61±0.06 8.23±0.09 8.31± 8.57± 8.05± 7.88± 8.06± 8.46± 8.50± 8.18± 8.19± 8.28± 8.31± 0.12 0.10 0.07 0.07 0.11 0.11 0.16 0.04 0.04 0.04 0.07

We follow ITL97 names in SBS 1319+579. Notice, however, than IT99 named region A to SBS 1319+579 C and region B to SBS 1319+579 A; they did not consider SBS 1319+579 B b ecause of its 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).


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Fig. 36. Detail of the spectra of the main regions within our galaxy sample showing the zones around the blue WR bump. The red dotted line represents the position of the He ii 4686 emission line, the blue dotted lines indicate the position of [Fe iii] 4658 and [Ar iv] 4711,4740 emission lines. The black dotted line represents the continuum level fitted by eye.


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29

Fig. 37. Detail of the spectra of the main regions within our galaxy sample showing the zones around the red WR bump. The red dotted line represents the position of the C iv 5808 emission line, the blue dotted lines indicate the position of [N ii] 5755 and He i 5875 emission lines. The black dotted line represents the continuum level fitted by eye. The red WR bump is clearly identified in HCG 31 AC and POX 4, detected Mkn 1199 and ESO 566-8 and it also seems to be observed in IRAS 08208+2816 and Tol 9. We do not have data for SBS 1211+540 and NGC 5253 in this spectral range.

F1, Haro 15 A, Mkn 5, POX 4, UM 420, SBS 0926+606 A, SBS 0948+532, SBS 1415+437 C, Tol 1457-262 A and B, ESO 566-8 and NGC 5253 A and D), being specially strong in POX 4. Only in three ob jects previously listed as WR galaxies (Mkn 1087, SBS 1211+540 and ESO 566-7) we do not detect any feature that can be attributed to the presence of this sort of massive stars. We consider that aperture effects and the exact positioning of the slit onto the WRrich bursts seem to play a fundamental role in the detection of the WR features.

As it was expected, the red WR bump (the broad C iii 5696 and C iv 5808 emission lines) is much more difficult to observe. Figure 37 shows a detail of the optical spectrum in the 5550­6000 ° range (red WR bump) of A all important ob jects for which data is this spectral range is available. The broad C iv 5808 emission line is clearly identified in 2 galaxies (HCG 31 AC and POX 4), detected in 2 galaxies (Mkn 1199 and ESO 566-8) and it also seems to be observed in other 2 galaxies (IRAS 08208+2816 and Tol 9). However, we do not see this feature in galaxies


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L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
1000

1000
HCG 31 G

EW (H) [å] - from spectroscopy

EW (H) [å] - Imagery

100

100

POX 4 Comp

10

EW(H)spec= EW(H) Linear fit
10 10 100 1000

phot

SB99, Padova Z=0.2Zo SB99, Padova Z=0.4Zo SB99 Geneva, Z=0.4Zo
1 3 4 5 6 7 8 9 10

EW (H) [å] - Spectroscopy

Age [Myr] - from H images

Fig. 38. Comparison between the H equivalent widths derived from our H images and those obtained from the analysis of the optical spectrum for every particular star-forming region analyzed in this work.

for which previous detection have been reported [Mkn 5, UM 420, SBS 0926+606, SBS 0948+532, SBS 1054+365 and SBS 1319+579A, Guseva, Izotov & Thuan (2000)]. We will discuss this result in Paper III. A detailed quantitative analysis of the WR/(WR+O) and WC/WN ratios and the comparison with theoretical evolution models will be also performed in Paper III. As it was commented in Paper I, one of the best methods of determining the age of the last star-forming burst is through the H equivalent width since it decreases with time. Our spectroscopic observations provide an independent estimation of W (H) within every knot. We have checked the correspondence between the values we obtained in our deep H images (see Table 7 in Paper I) with those derived from the spectroscopic data. Figure 38 plots such correlation, showing the excellent agreement of both kinds of data in almost all ob jects. Indeed, the linear fit to the data practically coincides with a x = y function. Small divergences are found in very few ob jects, but they can be explained because of considerable differences in the relative sizes for which the photometric/spectroscopic values were extracted. The most evident case is member G in HCG 31, for which we only extracted the spectrum of a small knot at its NW, but the W (H) value derived from the images considers the flux of all the galaxy, that possesses a global star-formation activity lower than that observed in the NW knot. Figure 39 shows the comparison of our data with the theoretical predictions provided by the last release of the STARBURST 99 (Leitherer et al. 1999) models which uses Padova tracks (V´ azquez & Leitherer 2005). We assumed an instantaneous burst with a Salpeter IMF, a total mass of 106 M , and a metallicity of Z/Z = 0.2 and 0.4, the most common values according to the oxygen abundance of the ma jority of the knots. For comparison, we have also included the predictions of the original STARBURST 99 models, that consider Geneva tracks, for Z/Z =0.4. The ages of the last star-formation event are those estimated from the W (H) determined from our deep images, while the H equivalent widths are those directly measured from

Fig. 39. H equivalent width vs. age of the most recent starforming burst diagram comparing the predictions given by the evolutionary synthesis models provided by STARBURST 99 (Leitherer et al. 1999). We include the Z/Z =0.4 model originally included in STARBURST 99 using Geneva tracks and two new models with Z/Z = 0.2 and 0.4 than consider Padova tracks (Vazquez & Leitherer, 2005). The age was computed from the W (H) given by our images; W (H ) was determined from our optical spectra.

our spectra. Thus, both set of data come from independent observations. As we see, the agreement is excellent, and therefore we are quite confident in the determination of the age of the star-forming regions. As it should be expected, the predictions given by the new models are better than those obtained from the old models. The only data point that does not follow the models is the dwarf companion ob ject surrounding POX 4. However, as we explained in Paper I, the values of the W (H) have been taken from images obtained by M´ endez & Esteban (1999) and seem to be somewhat overestimated. Although absorption features have been only detected in some of the ob ject, all of them possess an old stellar population underlying the bursts, as we commented in the analysis of the optical/NIR colors in Paper I. This fact is evident from the values of the Wabs derived from our spectra using the H i Balmer lines, that seems to increase with increasing metallicity. A very powerful method to constraint the ages of the stellar populations within a starburst galaxy is the analysis of its spectral energy distributions (SED), although have the degeneracy problem between the interstellar extinction and the age of the old stellar population. In some ob jects, we have checked the results given by this method considering our estimation of the reddening contribution derived from the Balmer decrement, as we previously did in our analysis of the stellar populations in IRAS 08339+6517 (L´ ez-S´ hez, Esteban & Garc´ op anc iaRo jas 2006). We have made use of the PEGASE.2 code (Fioc & Rocca-Volmerange 1997) to produce a grid of theoretical SEDs for an instantaneous burst of star formation and ages between 0 and 10 Gyr, assuming a Z metallicity and a Salpeter IMF with lower and upper mass limits of 0.1 M and 120 M . Although the grid include the ionized gas emission, we have neglected it because its contribution to the continuum is almost irrelevant. In Figure 40 we show the extinction-corrected spectra of the region A in Haro 15 (left ) and the center of POX 4 (right ) and two synthetic


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

31

Fig. 40. Spectra of region A in Haro 15 (left ) and the center of POX 4 (right ) compared with synthetic continuum spectral energy distributions obtained using the PEGASE.2 (Fioc & Rocca-Volmerange 1997) code. The gray/green continuous line is the extinction-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). The shape of our observed derredened spectra fit in both cases with a model with a contribution of 15% for the young population and 85% for the old population (continuous black line over the galaxy spectrum).

continuum spectral energy distributions assuming young (blue line ) and old (red line ) populations. We considered the ages derived from W (H) representative for the young population ages (4.5 Myr in both cases) and the ages of the underlying component estimated from the optical/NIR colors (500 Myr in Haro 15, 250 Myr in POX 4). As we expected, none of the individual synthetic spectra fitted our observed SED. We then constructed a model than combines both young and old models. For both cases, the best fits are found when 15% of the 4.5 Myr model and 85% of the old model are considered. As we see, this combined model is in very good agreement with the shape of our derreddened spectrum. We conclude that, although the star formation activity is very intense in these starbursts, an important underlying old stellar population is usually found in the galaxies, indicating previous star-forming phenomena and ruling out the hypothesis that some of them are pristine dwarf galaxies. Our study of the kinematics of the ionized gas and the morphology and environment included in this and previous papers of our group has revealed that 14 up to 20 of the analyzed galaxies show rather clear kinematical and/or morphological evidences of interaction or merging. The morphological evidences were presented and discussed in Paper I. The kinematical evidences presented here are of diverse nature: presence of ob jects with velocities decoupled from the main rotation pattern (Mkn 1087, Haro 15), sinusoidal velocity patterns that suggest a merging 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), indications of tidal streaming (HCG 31, IRAS 08208+2816, SBS 1319+579, Tol 9) or the presence of TDG candidates (HCG 31 F1 and F2, Mkn 1087, IRAS 08339+6517, POX 4, Tol 1457-262). The interaction could be between a 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 dwarf one (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 dwarf starburst galaxy that has suffered a possible interaction with a galaxy in the M 83 subgroup or with the spiral galaxy M 83 itself (L´ ez-S´ hez et al. 2008). These results reinop anc force the hypothesis that interaction with or between dwarf ob jects is an important mechanism to trigger the massive star formation in this kind of starbursts. These neighboring interacting dwarf or low-luminosity ob jects are only detected when a systematic and detailed analysis of the morphology, environment, chemical composition and kinematics of the ob jects are carried out. Finally, our detailed spectroscopical analysis has provided with some clear evidences of chemical differences within the ob jects or of interacting ob jects of different metallicities. For example, Mkn 1087, Haro 15 and Mkn 1199 are in clear interaction with dwarf galaxies with lower O/H and N/O ratios. NGC 5253, IRAS 08208+2816, and Tol 1457-262 show zones with different chemical compositions. In the case of NGC 5253 this is produced by localized pollution by massive stars, but in the cases of IRAS 08208+2816 and Tol 1457-262 the different chemical composition seem to be because the regions correspond to different galaxies in interaction. Apart from NGC 5253, two more galaxies: IRAS 08208+2816 and UM 420, show a localized high N/O that could be a signature of contamination by WR winds.
´ Acknow ledgements. A.R. L-S thanks C.E. (his PhD supervisor) for all the help and very valuable explanations, talks and discussions along these years. He also acknowledges Jorge Garc´ ia-Ro jas, Sergio Sim´ on-D´ and Jos´ Caballero for their help and friendship during iaz e his PhD, extending this acknowledge to all people at Instituto de Astrof´ isica de Canarias (Spain). The authors thank B¨ el Koribalski arb ´ (CSIRO/ATNF) for her help analyzing HIPASS data. A.R. L-S. deeply thanks to Universidad de La Laguna (Tenerife, Spain) for force him to translate his PhD thesis from English to Spanish; he had to translate it from Spanish to English to complete this publication. This work has been partially funded by the Spanish Ministerio de Ciencia y Tecnolog´ (MCyT) under pro ject AYA2004-07466. This research has ia made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.


32

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc
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33

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 13 `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

App endix A: Tables


34

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.1. Dereddened line intensity ratios with respect to I (H )=100 for knots analyzed in Haro 15.
Line 3705.04 3728.00 3770.63 3797.90 3835.39 3868.75 3889.05 3967.46 3970.07 4026.21 4068.60 4101.74 4340.47 4363.21 4471.48 4658.10 4686.00 4711.37 4740.16 4754.83 4861.33 4921.93 4958.91 5006.84 5015.68 5158.81 5197.90 5517.71 5537.88 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6730.85 He I [O II] HI HI HI [Ne III] HI [Ne III] HI He I [S II] HI HI [O III] He I [Fe III] He II [Ar IV] [Ar IV] [Fe III] HI He I [O III] [O III] He I [Fe I I] [N I] [Cl II I] [Cl II I] He I [O I] [S II I] [O I] [N II] HI [N II] He I [S II] f ( ) 0.260 0.256 0.249 0.244 0.237 0.230 0.226 0.210 0.210 0.198 0.189 0.182 0.127 0.121 0.095 0.050 0.043 0.037 0.030 0.026 0.000 -0.015 -0.024 -0.036 -0.038 -0.073 -0.082 -0.154 -0.158 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.344 C 1.05: 294±19 0.68: 0.33: 1.69: 14.2±2.8 7.6±3.0 ... 19.6±1.7 ... 2.69: 28.1±2.7 46.2±3.5 1.22: 3.56: 1.75: 0.49: ... ... ... 100.0±6.7 ... 67.2±5.0 210±13 1.02: 1.95: 2.76: ... ... 12.8±1.6 7.5±1.1 1.10: 2.18: 20.1±1.8 288±19 64.8±4.6 2.81: 18.0±2.0 6â1 0 23.25 ± 1.08 0.11 ± 0.03 2.4 ± 0.4 75.2 16.4 5.5 29.4 ± ± ± ± 5.0 1.1 0.4 1.8 A ... 113±21 1.69±0.52 3.10±1.00 4.59±0.99 48.6±8.4 15.5±2.8 29.6±4.8 ... 1.05: ... 26.1±3.5 47.0±4.8 8.59±0.87 4.17±0.68 0.90: 1.65±0.42 1.10: 0.73: 0.25: 100.0±5.9 0.55: 212±11 648±36 ... ... ... 0.33: 0.26: 9.95±1.70 2.21±0.76 1.16±0.31 0.65: 2.82±0.71 202±46 7.76±1.84 2.55±0.74 5.09±1.32 8.4â1 13 23.42 ± 0.89 0.33 ± 0.03 1.3 ± 0.3 423.6 ± 22.5 75.7 ± 4.2 26.9 ± 1.5 462.7 ± 23.1 B ... 402±106 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 100±39 ... 81±30 232±63 ... ... ... ... ... 46.1: ... ... ... ... 284±75 23.4: ... 28.5: 8â1 23 0.52 ± 0.10 0.06 ± 0.03 0.5 43.9 ± 10.0 8.4 ± 3.2 ... 20.2 ± 5.1 D ... 336±46 ... ... ... 39.4: ... ... ... ... ... 26.1±7.0 46.7±9.9 3.9: ... ... ... ... ... ... 100±19 ... 132±19 383±44 ... ... ... ... ... 10.5±2.8 7.7±3.1 ... ... 6.0: 288±35 23.2±6.1 3.9: 17.6±8.3 4.4 â1 12 1.32 ± 0.12 0.37 ± 0.02 2.2 ± 0.2 48.8 20.8 7.5 77.9 ± ± ± ± 2.7 1.3 0.9 3.9

Aperture size (arcsec) Distanceb (arcsec) F (H )a C (H ) Wabs (° A) -W -W -W -W
a b

(H) (° A) (H ) (° A) (H ) (° A) ([O III]) 5007 (° A)

In units of 10-15 erg s-1 cm-2 and not corrected for extinction. Relative distance with respect to the center of Haro 15.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc Table A.2. Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Haro 15.
Region Te (O III) (K) Te (O II) (K) ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar+3 /H+ ) 12+log(Cl++ /H+ ) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]
a Estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

35

C 9500 ± 800a 9600 ± 600 100 8.16 ± 0.11 7.94 ± 0.10 8.37 ± 0.10 -0.23 7.13 7.34 -1.03 6.05 6.52 6.65 -1.71 ± ± ± ± ± ± ± ± 0.16 0.07 0.10 0.15 0.10 0.25 0.20 0.18

A 12900 ± 700 12000 ± 500 100 7.35 ± 0.08 8.01 ± 0.06 8.10 ± 0.06 0.66 6.00 6.75 -1.35 5.26 6.02 6.20 -1.89 ± ± ± ± ± ± ± ± 0.10 0.06 0.10 0.11 0.08 0.13 0.11 0.15

B 11500 ± 1000a 11000 ± 700 100 8.04 ± 0.14 7.72 ± 0.13 8.21 ± 0.14 -0.32 6.55 6.72 -1.49 ± ± ± ± 0.18 0.19 0.21 0.20

D 11800 ± 800a 11260 ± 600 100 7.93 ± 0.13 7.90 ± 0.10 8.22 ± 0.11 -0.03 6.47 6.76 -1.46 ± ± ± ± 0.15 0.14 0.15 0.16

6.08 ± 0.22 ... ... ... ... ... ... ... ... ... ... ... 10.96: -0.45

5.85 ± 0.21 ... ... ... 7.14 ± 0.20 7.46 ± 0.20 -0.76 ± 0.18 ... ... ... ... ... 10.96 ± 0.12 -0.44

7.29 ± 0.15 7.72 ± 0.15 -0.65 ± 0.18 ... ... 5.2: 6.2: -2.2: 10.97 ± 0.05 -0.29

7.33 ± 0.10 7.42 ± 0.10 -0.68 ± 0.12 4.92 ± 0.17 4.26 ± 0.28 5.5: 6.5: -1.6: 10.88 ± 0.06 -0.56±0.11

= 8.66±0.05 (Asplund, Grevesse & Sauval 2005).


36

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.3. Dereddened line intensity ratios with respect to I (H )=100 for knots analyzed in Mkn 1199.
Line 3728.00 3835.39 3868.75 3889.05 3967.46 3970.07 4101.74 4340.47 4471.48 4658.10 4686.00 4861.33 4958.91 5006.84 5055.98 5197.90 5754.64 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7318.39 7329.66 [O II] HI [Ne III] HI [Ne III] HI HI HI He I [Fe I II] He II HI [O III] [O III] Si II [N I] [N II] He I [O I] [S III] [O I] [N II] HI [N II] He I [S II] [S II] He I [Ar III] [O II] [O II] f () 0.256 0.237 0.230 0.226 0.210 0.210 0.182 0.127 0.095 0.050 0.043 0.000 -0.024 -0.036 -0.048 -0.082 -0.194 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.418 -0.420 C 124.6±7.8 1.23: 1.39±0.46 9.2±1.9 ... 18.0±1.2 27.0±1.9 45.9±3.4 2.48: 3.44±0.68 0.24: 100.0±6.0 10.7±1.4 30.6±2.4 0.81: 2.49±0.60 0.61: 9.46±0.93 3.57±0.80 0.14: 0.85: 48.2±3.1 293±17 149.8±8.9 1.90±0.49 35.8±2.3 32.1±6.4 1.30±0.42 1.91±0.56 0.87±0.32 0.39: 10â1 0 74.2 ± 3.1 0.30 ± 0.03 1.8 ± 0.4 129.1 21.4 6.7 6.8 ± ± ± ± 7.9 1.3 0.5 0.5 NE 254±20 ... 6.64: 9.40: 6.22: ... 26.0±5.9 46.7±6.2 5.46: ... ... 100±11 60.0±8.4 167±15 2.37: ... ... 12.7±3.1 6.6±2.4 1.19: ... 13.9±3.5 290±25 41.2±5.6 3.88: 35.4±5.2 25.5±4.4 ... 6.29: ... ... 6â1 26 3.17 ± 0.21 0.16 ± 0.03 0.6 ± 0.3 110 20.2 8.4 35.5 ± ± ± ± 10 2.3 1.1 3.4 A 204: ... ... ... ... ... ... 31: ... ... ... 100: 23: 62: ... ... ... ... ... ... ... 51: 294: 142: ... 66: 40: ... ... ... ... 8â1 18 0.33 ± 0.08 0.17 ± 0.04 1.7 ± 0.3 21 5.0 ± 1.5 ± 2.8 ± ±7 2.6 1.7 1.6 B 145: ... ... ... ... ... ... ... ... ... ... 100: ... 85: ... ... ... ... ... ... ... 51: 296: 141: ... 74: 44: ... ... ... ... 8â1 14 0.42 ± 0.10 0.44 ± 0.06 2 22 ± 8 5.2 ± 2.4 ... 4.2 ± 1.8 D 192±45 ... ... ... ... ... 30.5±7.6 47±18 ... ... ... 100±27 23.5: 60±20 ... ... ... 11.6: 16.8: ... ... 57±14 298±65 172±35 ... 74±19 49±14 ... ... ... ... 5.6â1 8.4 0.98 ± 0.14 0.27 ± 0.04 2.5 ± 0.3 34.7 8.3 2.1 5.1 ± ± ± ± 7.6 2.3 0.8 1.7

Aperture size (arcsec) Distanceb (arcsec) F (H )a C (H ) Wabs (° A) ° -W (H) (A) -W (H ) (° A) -W (H ) (° A) -W ([O III]) 5007 (° A)
a b

In units of 10-15 erg s-1 cm-2 and not corrected for extinction. Relative distance with respect to the center of Mkn 1199.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc Table A.4. Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Mkn 1199.
Region Te (O III) (K) Te (O II) (K) ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar+3 /H+ ) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]b Center 5400 ± 700 6800 ± 600b 300 ± 100 8.59 ± 0.11 8.24 ± 0.12 8.75 ± 0.12 -0.36 7.98 8.14 -0.62 6.70 7.05 7.22 -1.54 ± ± ± ± ± ± ± ± 0.16 0.09 0.11 0.10 0.09 0.17 0.15 0.14 NEa 8450 ± 800 8900 ± 600 100 8.25 ± 0.12 8.05 ± 0.11 8.46 ± 0.13 -0.19 7.05 7.26 -1.20 6.28 6.80 6.92 -1.54 ± ± ± ± ± ± ± ± 0.09 0.10 0.13 0.11 0.09 0.22 0.18 0.17 Aa 6950 ± 800 7850 ± 600 100 8.43 ± 0.20 7.99 ± 0.23 8.57 ± 0.21 -0.44 7.76 7.90 -0.67 ± ± ± ± 0.18 0.14 0.18 0.20 Ba 6300 ± 800 7400 ± 600 100 8.43 ± 0.21 8.32 ± 0.26 8.68 ± 0.23 -0.10 7.85 8.10 -0.58 ± ± ± ± 0.32 0.18 0.28 0.30 D
a

37

6750 ± 800 7700 ± 600 100 8.44 ± 0.16 8.05 ± 0.19 8.59 ± 0.17 -0.37 7.86 8.00 -0.59 ± ± ± ± 0.22 0.11 0.15 0.17

6.67 ± 0.16 ... ... ... ... ... ... ... ... ... ... ... -0.09

6.80 ± 0.16 ... ... ... ... ... ... ... ... ... ... ... +0.02

6.76 ± 0.13 ... ... ... ... ... ... ... ... ... ... 10.9: -0.07

7.65 ± 0.17 8.17 ± 0.19 -0.58 ± 0.17 6.07 6.75 6.89 -1.86 ± ± ± ± 0.16 0.13 0.13 0.26

7.40 ± 0.20 7.81 ± 0.20 -0.65 ± 0.18 5.95 ± 0.22 ... ... ... 10.9: -0.20

10.79 ± 0.07 +0.09±0.17 relations.

a Electron temp eratures estimated using empirical b Derived from [N ii] and [O ii] ratios, see §3.4.1 c [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

= 8.66±0.05 (Asplund et al. 2005).


38

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.5. Dereddened line intensity ratios with respect to I (H )=100 for knots analyzed in Mkn 5. Region A was observed using three slit positions with a PA of 0 (INT-1), 349 = -11 (INT-2) and 354 = -6 (WHT).
Line 3666.10 3697.15 3705.04 3711.97 3728.00 3750.15 3770.63 3797.90 3819.61 3835.39 3868.75 3889.05 3967.46 4026.21 4068.60 4101.74 4168.97 4340.47 4363.21 4387.93 4416.27 4471.48 4658.10 4686.00 4711.37 4713.14 4740.16 4861.33 4921.93 4958.91 4985.90 5006.84 5015.68 5197.90 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7281.35 7318.39 7329.66 7751.10 HI HI He I HI [O II] HI HI HI He I HI [Ne III] HI [Ne III] He I [S I I] HI He I HI [O III] He I [Fe II] He I [Fe III] He II [Ar IV] He I [Ar IV] HI He I [O III] [Fe III] [O III] He I [N I] He I [O I] [S III] [O I] [N II] HI [N II] He I [S II] [S II] He I [Ar III] He I [O II] [O II] [Ar III] f () 0.267 0.262 0.260 0.259 0.256 0.253 0.249 0.244 0.240 0.237 0.230 0.226 0.210 0.198 0.189 0.182 0.167 0.127 0.121 0.115 0.109 0.095 0.050 0.043 0.037 0.037 0.030 0.000 -0.015 -0.024 -0.031 -0.036 -0.038 -0.082 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.414 -0.418 -0.420 -0.467 A-INT-1 1.99±0.71 1.16: 1.65: ... 191±12 1.79±0.68 ... ... ... 7.4±1.7 23.1±3.3 13.6±2.9 18.7±1.9 ... 2.24±0.70 26.0±2.8 ... 47.0±3.2 5.14±0.91 ... ... 4.43±0.85 ... 0.92±0.19 ... ... ... 100.0±6.3 ... 144.1±8.5 2.33±0.65 423±22 ... ... 8.46±0.83 5.12±0.68 2.31±0.55 1.90±0.51 4.56±0.56 283±16 14.7±1.2 3.43±0.63 21.5±1.5 15.6±1.2 2.13±0.50 8.49±0.61 ... ... ... ... 14.4â1 0 17.66 ± 0.73 0.36 ± 0.02 1.1 ± 0.2 449 ± 26 75 ± 5 43 ± 3 320 ± 17 A-INT-2 ... ... ... 2.03±0.76 213±12 ... 2.47±0.79 1.54: 0.65: 5.6±1.3 31.0±2.4 17.7±2.1 21.3±1.8 1.55: 2.29±0.76 26.8±2.1 1.06: 47.0±3.1 5.29±0.93 0.56: ... 3.96±0.84 1.10: 1.01: 0.59: ... ... 100.0±5.7 1.70: 144.7±7.8 2.02: 430±21 ... 0.46: 8.41±0.88 4.57±0.66 1.92±0.43 1.52: 5.25±0.69 284±15 14.0±1.1 3.49±0.66 22.2±1.7 16.0±1.4 2.40±0.64 8.56±0.97 ... 2.87: 2.57: ... 16â1 0 18.13 ± 0.69 0.17 ± 0.02 0.8 ± 0.2 435 ± 23 80 ± 5 34 ± 2 360 ± 18 A-WHT ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 47.0±2.8 4.93±0.74 1.32±0.40 0.38: 4.16±0.60 0.72: 0.82: ... 0.28: 0.31: 100.0±5.2 0.95±0.36 133.8±7.2 2.20: 374±19 2.13±0.62 ... 11.0±1.0 4.05±0.58 1.79±0.47 1.19±0.38 5.49±0.63 284±14 14.9±1.1 3.05±0.51 24.3±1.5 18.1±1.2 2.17±0.45 7.70±0.52 0.41: 3.62±0.54 2.74±0.44 1.64±0.38 3.6â1 0 10.83 ± 0.40 0.03 ± 0.02 0 678 ± 35 135 ± 7 44 ± 3 530 ± 28 B ... ... ... ... 252: ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 100: ... 70.43: ... 214±69 ... ... ... ... ... ... ... 284±85 10.13: ... 67.4: 51.6: ... ... ... ... ... ... 6 â1 16 0.35 ± 0.08 0.30 ± 0.06 1.5 ± 0.5 43 ± 12 10: ... 33 ± 11

Aperture size (arcsec) Distanceb (arcsec) F (H )a C (H ) Wabs (° A) ° -W (H) (A) -W (H ) (° A) -W (H ) (° A) -W ([O II I]) 5007 (° A)
a b

In units of 10-15 erg s-1 cm-2 and not corrected for extinction. Relative distance with respect to the center of Mkn 5.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc Table A.6. Physical conditions and chemical abundances of the ionized gas of the regions analyzed in Mkn 5.
Region Te (O III) (K) Te (O II) (K) ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar++ /H+ ) 12+log(Ar+3 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]b
a Electron temp eratures estimated using empirical b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

39

A-INT-1 12400 ± 700 11700 ± 500 100 7.62 ± 0.09 7.88 ± 0.06 8.07 ± 0.07 0.25 6.27 6.72 -1.35 5.78 6.37 6.51 -1.56 ± ± ± ± ± ± ± ± 0.10 0.06 0.09 0.10 0.04 0.13 0.12 0.13

A-INT-2 12450 ± 650 11700 ± 450 100 7.67 ± 0.08 7.87 ± 0.06 8.08 ± 0.07 0.21 6.29 6.71 -1.38 5.79 6.28 6.44 -1.64 ± ± ± ± ± ± ± ± 0.09 0.05 0.08 0.10 0.04 0.13 0.12 0.12

A-WHT 12700 ± 600 11900 ± 400 100 7.71 ± 0.11 7.80 ± 0.06 8.06 ± 0.08 0.09 6.30 6.65 -1.41 5.82 6.23 6.40 -1.67 ± ± ± ± ± ± ± ± 0.11 0.05 0.08 0.10 0.04 0.12 0.12 0.13 ... ... ... 5.60 4.54 5.73 -2.33 ± ± ± ± 0.09 0.21 0.10 0.12

Ba 13250 ± 900 12300 ± 700 110 7.66 ± 0.19 7.49 ± 0.13 7.89 ± 0.17 -0.17 6.08 6.30 -1.58 ± ± ± ± 0.18 0.22 0.22 0.20

6.24 ± 0.11 ... ... ... ... ... ... ... ... ... ... ... ... ... ... -0.77

7.02 ± 0.12 7.22 ± 0.12 -0.85 ± 0.14 5.69 ± 0.07 ... ... ... 5.73 ± 0.13 6.11 ± 0.16 -1.96 ± 0.18 10.96 ± 0.06 -0.59±0.12

7.14 ± 0.10 7.35 ± 0.11 -0.74 ± 0.13 5.69 4.69 5.79 -2.29 ± ± ± ± 0.09 0.25 0.12 0.15

5.66: 6.00: -2.08: 10.91 ± 0.06 -0.58±0.12

5.68: 5.98: -2.08: 10.92 ± 0.05 -0.60±0.13

relations. = 8.66±0.05 (Asplund et al. 2005).


40

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.7. Dereddened line intensity ratios with respect to I (H )=100 for regions analyzed in IRAS 08208+2816. The slit positions we used for each knot are: PA 345 for C, PA 355 for #8 and #10 and PA 10 for #3 y #5.
Line 3728.00 3750.15 3770.63 3797.90 3835.39 3868.75 3889.05 3967.46 4026.21 4068.60 4101.74 4340.47 4363.21 4471.48 4658.10 4861.33 4958.91 5006.84 5197.90 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 [O II] HI HI HI HI [Ne II I] HI [NeIII]H7 He I [S II] HI HI [O III] He I [Fe III] HI [O III] [O III] [N I] He I [O I] [S III] [O I] [N II] HI [N II] He I [S II] [S II] He I [Ar III] f () 0.256 0.253 0.249 0.244 0.237 0.230 0.226 0.210 0.198 0.189 0.182 0.127 0.121 0.095 0.050 0.000 -0.024 -0.036 -0.082 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 C 146.8±9.8 1.74±0.51 1.55±0.49 2.86±0.60 5.53±0.79 32.8±2.2 15.1±1.7 22.1±1.9 0.92: 1.14±0.45 26.1±2.1 46.7±2.7 3.12±0.48 4.00±0.67 1.68±0.58 100.0±5.7 152.6±8.3 470±24 1.09±0.42 12.8±1.2 4.74±0.62 1.46±0.42 1.70±0.46 14.2±1.1 285±15 36.9±2.6 3.26±0.63 22.0±1.5 16.6±1.3 2.42±0.57 6.37±0.78 #3 279±20 ... ... 3.17: ... 15.1±3.4 13.5±3.2 15.9±2.9 ... ... 26.1±2.9 46.8±4.6 1.66: 3.7±1.3 1.72: 100.0±8.4 76.6±5.8 228±14 ... 12.4±1.9 5.6±1.1 1.13: ... 16.0±2.4 283±20 39.2±4.2 3.04±0.93 40.2±4.2 28.5±3.3 ... 11.8±3.0 #5 251±24 ... ... ... ...... 26.3±7.0 12.9±5.0 17.0±5.3 ... ... 24.1±5.1 43.8±6.8 2.26: ... ... 100±11 104±10 305±24 2.46: 14.6±2.9 9.8±2.3 ... 3.21: 16.2±3.9 281±25 37.8±5.3 ... 52.9±6.8 34.5±5.2 ... ... #8 164±22 ... ... ... ... 5.27: 6.36: 22.0±5.7 ... ... 26.1±5.7 43.5±7.4 ... ... ... 100±12 44.7±8.9 89±13 ... 12.3±3.4 10.1±3.1 ... 2.24: 30.2±5.3 288±33 75.3±9.3 2.93: 66.4±9.6 46.7±7.5 ... ... #10 324±64 ... ... ... ... ... 26.23: ... ... 26.23: 46±15 ... ... ... 100±22 54±15 151±30 ... 11.28: 9.00: ... ... 14.95: 286±51 42±11 3.65: 50±13 36±10 ... ...

Aperture size (arcsec) Distance (arcsec)b F (H )a C (H ) Wabs (° A) ° -W (H) (A) -W (H ) (° A) -W (H ) (° A) -W ([O III]) (° A)
a b

1.4â1 2.8â1 4.0â1 6.4â1 4.0â1 12.6 6.8 8.8 16 12.9 ± 0.5 5.97 ± 0.28 2.98 ± 0.19 4.05 ± 0.34 1.36 ± 0.18 0.11 ± 0.02 0.47 ± 0.04 0.41 ± 0.04 0.17 ± 0.03 0.12 ± 0.02 3.2 ± 0.1 1.6 ± 0.2 1.4 ± 0.2 4.9 ± 0.1 1.9 ± 0.1 331 ± 18 80 ± 5 30 ± 2 370 ± 19 202 ± 15 56 ± 5 20 ± 2 130 ± 8 346 ± 32 62 ± 7 27 ± 4 200 ± 16 89 ± 10 24 ± 3 10 ± 2 18 ± 3 98 ± 18 17 ± 4 9±3 19 ± 4

In units of 10-15 erg s-1 cm-2 and not corrected for extinction. Relative distance with respect to the center of IRAS 08208+2816.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc Table A.8. Physical conditions and chemical abundances of the ionized gas of the regions analyzed in IRAS 08208+2816.
Region Te (O III) (K) Te (O II) (K) ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H)c 12+log(O/H)d log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) C 10100 ± 700 10100 ± 500 <100 7.77 8.20 8.33 8.41 0.43 6.88 7.44 -0.89 5.94 6.52 6.69 -1.64 ± ± ± ± ± ± ± ± ± ± ± ± 0.09 0.07 0.08 0.10 #3a 9400 ± 900 9600 ± 700 <100 8.15 ± 0.13 8.01 ± 0.11 8.24: 8.39 ± 0.11 ± ± ± ± ± ± ± ± #5a 9600 ± 1000 9700 ± 800 <100 8.08 ± 0.15 8.10 ± 0.12 8.27: 8.39 ± 0.13 ± ± ± ± #8a 6750 ± 1000 7700 ± 800 <100 8.39 ± 0.15 8.28 ± 0.14 ... 8.64 ± 0.15 ± ± ± ± #10a 9500 ± 1000 9650 ± 800 <100 8.20 ± 0.18 7.82 ± 0.17 ... 8.35 ± 0.17 ± ± ± ± 0.22 0.14 0.17 0.18

41

0.12 0.15 0.05 6.98 0.10 7.22 0.11 -1.17 0.05 6.25 0.16 6.56 0.13 6.75 0.16 -1.64

0.17 0.02 0.08 6.96 0.11 7.27 0.13 -1.12 0.08 0.28 0.23 0.25

0.19 -0.10 0.10 7.54 0.13 7.80 0.14 -0.84

0.20 -0.38 0.10 6.97 0.16 7.12 0.15 -1.22

6.33 ± 0.09 ... ... ...

6.73 ± 0.12 ... ... ...

6.34 ± 0.13 ... ... ... ... ... ... ... ... ... ... ... ... 10.94: -0.31

12+log(Ne++ /H+ ) 7.52 ± 0.11 7.37 ± 0.19 7.53 ± 0.22 7.60 ± 0.25 12+log(Ne/H) 7.66 ± 0.11 7.75 ± 0.18 7.82 ± 0.22 7.96 ± 0.25 log(Ne/O) -0.67 ± 0.13 -0.64 ± 0.21 -0.57 ± 0.25 -0.68 ± 0.26 12+log(Ar+2 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]b 5.78 ± 0.08 6.11 ± 0.15 5.86 ± 0.11 5.92 ± 0.15 -2.51 ± 0.15 -2.47 ± 0.18 5.90 ± 0.15 6.47 ± 0.15 -1.95 ± 0.17 10.97 ± 0.04 -0.33 ± 0.13 ... ... ... 10.95 ± 0.07 -0.27 ... ... ... ... ... ... 11.03 ± 0.09 -0.27 ... ... ... ... ... ... 10.90 ± 0.12 -0.02

a Electron temp eratures estimated using empirical b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

relations. = 8.66±0.05 (Asplund et al. 2005). c Oxygen abundance computed using the direct metho d. d Oxygen abundance computed using the empirical calibrations given by Pilyugin (2001a,b).


42

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.9. Dereddened line intensity ratios with respect to I (H )=100 for regions analyzed in POX 4, UM 420 and SBS 0926+606.
Line 3679.36 3682.81 3686.83 3691.56 3697.15 3703.86 3711.97 3721.83 3726.03 3728.82 3734.17 3750.15 3770.63 3797.90 3819.61 3835.39 3868.75 3889.05 3967.46 3970.07 4009.22 4026.21 4068.60 4076.35 4101.74 4143.76 4168.97 4276.83 4287.40 4340.47 4363.21 4387.93 4413.78 4471.48 4562.60 4571.20 4658.10 4686.00 4701.53 4711.37 4713.14 4740.16 4754.83 4861.33 4881.00 4921.93 4958.91 4985.90 5006.84 5015.68 5041.03 5047.74 5197.90 5200.26 5270.40 5517.71 5537.88 5754.64 5875.64 6300.30 6312.10 6363.78 6371.36 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7281.35 7318.39 7329.66 7751.10 HI HI HI HI HI HI HI [S I II] [O II] [O II] HI HI HI HI He I HI [Ne III] HI [Ne III] HI He I He I [S I I] [S I I] HI He I He I [Fe II] [Fe II] HI [O III] He I [Fe II I] He I Mg I] Mg I] [Fe II I] He I I [Fe II I] [Ar IV] He I [Ar IV] [Fe II I] HI [Fe II I] He I [O III] [Fe II I] [O III] He I Si I I He I [N I] [N I] [Fe II I] [Cl I II] [Cl I II] [N II] He I [O I] [S I II] [O I] Si I I [N II] HI [N II] He I [S I I] [S I I] He I [Ar III] He I [O II] [O II] [Ar III] f () 0.265 0.264 0.263 0.263 0.262 0.260 0.259 0.257 0.257 0.256 0.255 0.253 0.249 0.244 0.240 0.237 0.230 0.226 0.210 0.210 0.202 0.198 0.189 0.187 0.182 0.172 0.167 0.142 0.139 0.127 0.121 0.115 0.109 0.095 0.073 0.071 0.050 0.043 0.039 0.037 0.037 0.030 0.026 0.000 -0.005 -0.015 -0.024 -0.031 -0.036 -0.038 -0.044 -0.046 -0.082 -0.083 -0.100 -0.154 -0.158 -0.194 -0.215 -0.282 -0.283 -0.291 -0.292 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.414 -0.418 -0.420 -0.467 POX 4 POX 4 Comp 0.30±0.07 0.40±0.08 0.74±0.10 0.92±0.11 1.08±0.12 2.02±0.17 1.86±0.16 3.68±0.24 42.8±2.0 57.1±2.7 2.58±0.19 2.88±0.21 3.77±0.25 5.30±0.37 0.91±0.11 7.18±0.40 51.7±2.3 17.75±0.95 16.17±0.78 15.36±0.71 0.15±0.06 1.57±0.14 0.92±0.11 0.28±0.07 26.0±1.2 0.13: ... 0.08: 0.11: 51.2±2.3 11.89±0.61 0.42±0.08 0.07: 4.12±0.26 0.19±0.06 0.13: 0.64±0.09 1.21±0.16 0.14: 1.86±0.15 ... 1.34±0.12 ... 100.0±4.3 0.15±0.06 0.49±0.10 237±10 0.57±0.09 731±32 2.57±0.21 0.14: 0.16±0.06 0.23±0.06 0.30±0.07 ... 0.31±0.07 0.23±0.08 0.08: 10.93±0.56 1.81±0.12 1.60±0.12 0.61±0.09 0.05: 1.58±0.11 285±12 4.20±0.23 2.93±0.20 7.80±0.38 6.03±0.30 ... ... ... ... ... ... 7.2â1 ­ 56.0 ± 1.8 0.08 ± 0.01 2.0 ± 0.1 1075 ± 48 ... ... ... ... ... ... ... ... 142±35 165±40 ... ... ... ... ... ... 35.6: ... ... ... ... ... ... ... ... ... ... ... ... 49±16 ... ... ... ... ... ... ... ... ... ... ... ... ... 100±19 ... ... 93±19 ... 255±55 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 282±50 11.7: ... 15.9: 10.2: ... ... ... ... ... ... UM 420 SBS 0926+606A ... ... ... ... 0.53: 0.93±0.36 1.14±0.38 1.30±0.39 85.9±4.7 140.2±7.5 2.19±0.45 2.32±0.46 2.41±0.47 3.36±0.75 0.87: 5.96±0.70 29.7±2.0 16.8±1.5 6.45±0.72 16.1±1.2 ... 1.55±0.40 2.40±0.46 0.71: 26.1±2.1 ... ... ... 0.42: 47.0±3.0 4.48±0.82 0.21: ... 3.07±0.50 ... ... 1.31±0.46 1.04±0.28 ... ... 0.45: ... ... 100.0±5.5 ... 0.22: 106.8±6.0 0.57: 312±17 ... ... ... ... ... 0.48: 0.25: 0.26: 0.59: 10.51±0.79 7.49±0.49 1.75±0.18 2.06±0.41 ... 9.14±0.71 281±14 28.4±1.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.77±0.25 ... ... 47.5±2.8 7.00±0.66 ... ... 3.69±0.41 0.44: ... 1.20±0.40 0.72±0.21 ... 0.49: 0.72±0.24 0.49: 0.29: 100.0±4.9 0.29: 0.38: 157.8±7.7 ... ... ... ... ... ... ... ... ... ... ... 10.77±0.85 3.13±0.30 1.98±0.24 0.92±0.23 ... 2.47±0.30 286±15 7.50±0.64 2.72±0.34 16.8±1.0 12.12±0.77 2.45±0.32 6.54±0.45 0.61±0.20 2.11±0.23 1.71±0.22 1.75±0.25 4.0â1 ­ 16.20 ± 0.58 0.12 ± 0.03 0.7 ± 0.1 613 ± 33 SBS 0926+606B ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 47.4±9.1 ... ... ... ... ... ... ... ... ... ... ... ... ... 100±14 ... ... 134±18 ... ... ... ... ... ... ... ... ... ... ... 12.4±4.0 5.32: ... 2.24: ... 4.51: 286±34 18.2±4.6 2.63: 42.8±7.0 29.5±5.6 ... 5.09: ... 2.22: 2.05: ... 5.6â1 74.4 1.58 ± 0.15 0.18 ± 0.04 1.0 ± 0.3 92 ± 11

Ap erture size ( ) Distance ( )a F (H )b C (H ) Wabs (° A) -W (H) (° A)

3.6â1 3.6â1 20.4 ­ 0.11 ± 0.02 6.88 ± 0.27 0.06 ± 0.03 0.09 ± 0.01 1.4 ± 0.2 2.0 ± 0.1 380 ± 65 1076 ± 55


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc Table A.10. Physical conditions and chemical abundances of the ionized gas in POX 4, UM 420 and SBS 0926+606.
Ob ject Te (O II I) (K) Te (O II) (K) ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar+2 /H+ ) 12+log(Ar+3 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Cl++ /H+ ) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]
b

43

POX 4 14000 ± 500 12800 ± 400 250 ± 80 7.21 ± 0.04 7.96 ± 0.04 8.03 ± 0.04c 0.74 5.68 6.50 -1.54 5.28 6.03 6.24 -1.80 ± ± ± ± ± ± ± ± 0.06 0.04 0.06 0.06 0.03 0.08 0.07 0.10

POX 4 Compa 12400 ± 1000 11800 ± 600 <100 7.85 ± 0.14 7.55 ± 0.11 8.03 ± 0.13 -0.30 6.25 6.43 -1.60 ± ± ± ± 0.22 0.18 0.22 0.20

UM 420 13200 ± 600 12200 ± 500 140 ± 80 7.63 ± 0.05 7.67 ± 0.05 7.95 ± 0.05c 0.00 6.52 6.84 -1.11 5.61 6.16 6.29 -1.66 ± ± ± ± ± ± ± ± 0.08 0.05 0.06 0.07 0.11 0.09 0.10 0.13

SBS 0926+606A 13600 ± 700 12500 ± 500 <100 7.38 ± 0.10 7.80 ± 0.08 7.94 ± 0.08c 0.42 5.93 6.48 -1.45 5.61 6.17 6.34 -1.60 ± ± ± ± ± ± ± ± 0.12 0.05 0.10 0.09 0.04 0.11 0.10 0.13 ... ... ... 5.52 4.61 5.60 -2.34 ± ± ± ± 0.07 0.18 0.11 0.13

SBS 0926+606Ba 11500 ± 1000 11000 ± 800 <100 7.73 ± 0.16 7.04 ± 0.15 8.15 ± 0.16 0.21 6.39 6.80 -1.35 ± ± ± ± 0.14 0.09 0.13 0.12

5.62 ± 0.20 ... ... ... 6.97: 7.40: -0.60: ... ... ... ... ... ... ... ... ... -0.63

6.12 ± 0.11 ... ... ... ... ... ... 5.54 ± 0.15 ... ... ... ... ... ... ... 11.0: -0.51

7.18 ± 0.06 7.26 ± 0.06 -0.78 ± 0.10 ... 5.03 ± 0.07 ... ... 3.83 5.14 5.86 -2.17 ± ± ± ± 0.14 0.10 0.10 0.11

6.96 ± 0.09 7.24 ± 0.09 -0.71 ± 0.13 ... ... ... ... 4.18 5.52 5.79 -2.16 ± ± ± ± 0.26 0.12 0.13 0.13

... 5.47 ± 0.14 5.95 ± 0.15 -1.99 ± 0.16 10.94 ± 0.04 -0.72 ± 0.13

10.91 ± 0.03 -0.63 ± 0.09

10.88 ± 0.04 -0.71±0.10

a Electron temp eratures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H) = 8.66±0.05 c Considering the existence of O+3 b ecause of the detection of He

(Asplund et al. 2005). ii 4686, this value should be 0.01­0.02 dex higher.


44

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.11. Dereddened line intensity ratios with respect to I (H )=100 for regions analyzed in SBS 0948+532, SBS 1054+365 and SBS 1211+540.
Line 3697.15 3703.86 3711.97 3721.83 3726.03 3728.00 3728.82 3734.17 3750.15 3770.63 3797.90 3819.61 3835.39 3868.75 3889.05 3967.46 3970.07 4026.21 4068.60 4101.74 4340.47 4363.21 4471.48 4658.10 4686.00 4711.37 4713.14 4733.93 4740.16 4754.83 4861.33 4958.91 4985.90 5006.84 5015.68 5197.90 5517.71 5537.88 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 HI HI HI [S I II] [O II] [O II] [O II] HI HI HI HI He I HI [Ne I II] HI [Ne I II] HI He I [S I I] HI HI [O III] He I [Fe III] He II [Ar IV] He I [Fe III] [Ar IV] [Fe III] HI [O III] [Fe III] [O III] He I [N I] [Cl I II] [Cl I II] He I [O I] [S I II] [O I] [N II] HI [N II] He I [S I I] [S I I] He I [Ar III] f () SBS 0948+532 0.262 0.260 0.259 0.257 0.257 0.256 0.256 0.255 0.253 0.249 0.244 0.240 0.237 0.230 0.226 0.210 0.210 0.198 0.189 0.182 0.127 0.121 0.095 0.050 0.043 0.037 0.037 0.031 0.030 0.026 0.000 -0.024 -0.031 -0.036 -0.038 -0.082 -0.154 -0.158 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 0.48: 1.96±0.38 1.25±0.32 1.57±0.35 45.4±2.7 ... 65.7±3.8 2.52±0.43 2.56±0.43 2.88±0.45 5.11±0.71 0.65±0.25 7.74±0.77 41.6±2.6 19.7±1.5 15.0±1.2 16.2±1.2 1.45±0.32 1.53±0.33 26.2±1.7 47.2±2.9 8.13±0.82 3.81±0.49 1.13: 1.3±0.3 0.85±0.25 ... 0.18: 0.61±0.24 0.34: 100.0±5.2 189.5±9.2 1.49±0.30 584±28 ... ... 0.39: 0.24: 10.76±0.81 3.00±0.31 1.74±0.24 1.07±0.23 2.45±0.30 278±14 6.21±0.60 2.81±0.36 ... ... ... ... 3.6 0 8.44 ± 0.32 0.35 ± 0.03 0.3 ± 0.1 788 ± 43 213 ± 11 57 ± 4 689 ± 34 SBS 1054+365 SBS 1054+365 b SBS 1211+540 ... ... ... ... ... 100.2±7.5 ... ... ... 1.53: 2.85: ... 6.8±1.3 49.6±4.8 19.8±2.8 27.4±2.8 ... ... 2.05: 26.4±3.0 47.2±4.1 9.7±1.5 3.31±0.81 ... 0.61±0.23 1.18: ... ... ... ... 100.0±7.4 210±13 ... 623±37 ... 1.03: ... ... 8.7±1.6 1.12: 1.36: 0.65: 2.02±0.74 277±17 5.6±1.0 2.94±0.96 10.1±1.2 7.03±0.97 2.08±0.83 8.8±1.5 6.4 0 14.57 ± 0.68 0.02 ± 0.02 0.8 ± 0.1 422 ± 27 89 ± 7 43 ± 4 567 ± 35 ... ... ... ... ... 350±95 ... ... ... ... ... ... ... ... ... ... ... ... ... ... 28.51: ... ... ... ... ... ... ... ... ... 100±39 75±23 ... 183±53 ... ... ... ... ... ... ... ... 7.66: 279±74 20.2: 7.40: 22.4: 19.6: ... ... 5.8 17.8 0.66 ± 0.13 0.6 ± 0.1 0.3 ± 0.1 32 8 2 12 ± ± ± ± 8 3 1 3 ... ... ... ... 33.0±3.8 ... 44.8±4.6 2.61: 3.7±1.3 3.2±1.2 3.99: ... 5.9±1.5 37.6±4.0 19.9±3.0 13.3±2.1 16.1±2.2 ... ... 26.2±3.0 47.3±4.1 12.2±1.8 4.6±1.3 ... ... ... 0.75: ... 0.98: ... 100.0±7.5 163±10 ... 481±29 2.24: ... ... ... ... 2.36: 2.8±1.1 0.57: 0.82: 280±17 2.24±0.89 3.6±1.1 5.65±0.88 4.91±0.82 ... ... 3.6 0 1.84 ± 0.09 0.12 ± 0.01 1.3 ± 0.1 705 ± 45 135 ± 10 74 ± 7 618 ± 38

Aperture size (arcsec) Distancea (arcsec) F (H )b C (H ) Wabs (° A) - - - -
a b

W W W W

° (H) (A) (H ) (° A) (H ) (° A) ([O III]) 5007 (° A)

Relative distance of the knot with respect to the main region in the galaxy. In units of 10-15 erg s-1 cm-2 and not corrected for extinction.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

45

Table A.12. Physical conditions and chemical abundances of the ionized gas in SBS 0948+532, SBS 1054+365 and SBS 1211+540.
Ob ject Te (O II I) (K) Te (O II) (K) Ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar+2 /H+ ) 12+log(Ar+3 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Cl++ /H+ ) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]
b

SBS 0948+532 13100 ± 600 12200 ± 400 250 ± 80 7.33 ± 0.05 7.94 ± 0.05 8.03 ± 0.05 0.61 5.91 6.61 -1.42 5.43 6.16 6.34 -1.69 ± ± ± ± ± ± ± ± 0.08 0.05 0.07 0.08 0.12 0.11 0.11 0.14

SBS 1054+365 13700 ± 900 12600 ± 700 <100 7.22 ± 0.10 7.92 ± 0.07 8.00 ± 0.07 0.70 5.81 6.59 -1.41 5.37 5.99 6.21 -1.79 ± ± ± ± ± ± ± ± 0.11 0.08 0.09 0.08 0.07 0.22 0.18 0.15

SBS 1054+365ba 11800 ± 1100 11300 ± 900 300 ± 200 7.97 ± 0.18 7.62 ± 0.12 8.13 ± 0.16 -0.35 6.49 6.65 -1.47 ± ± ± ± 0.20 0.20 0.21 0.20

SBS 1211+540 17100 ± 600 15000 ± 400 320 ± 50 6.88 ± 0.05 7.57 ± 0.04 7.65 ± 0.04 0.69 5.26 6.03 -1.62 5.04 6.02 6.18 -1.47 ± ± ± ± ± ± ± ± 0.07 0.12 0.13 0.10 0.06 0.14 0.12 0.14

5.89 ± 0.16 ... ... ... ... ... ... ... ... ... ... ... ... ... ... 11.30: -0.53

7.21 ± 0.09 7.30 ± 0.09 -0.73 ± 0.12 ... 4.79 ± 0.15 ... ... 3.97 5.64 6.25 -1.78 ± ± ± ± 0.18 0.09 0.09 0.10

7.25 ± 0.09 7.33 ± 0.12 -0.67 ± 0.11 5.62 4.90 5.71 -2.29 ± ± ± ± 0.10 0.20 0.17 0.14 ... ... ... ... 10.88 ± 0.07 -0.66 ± 0.12

6.82 ± 0.08 6.90 ± 0.08 -0.75 ± 0.10 ... 4.77 ± 0.22 ... ... ... ... ... ... 10.90 ± 0.15 -1.01 ± 0.09

10.88 ± 0.04 -0.63 ± 0.10

a Electron temp eratures estimated using empirical b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

relations. = 8.66±0.05 (Asplund et al. 2005).


46

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.13. Dereddened line intensity ratios with respect to I (H )=100 for knots analyzed in SBS 1319+579 (regions A, B and C) and SBS 1415+457 (regions C and A).
SBS 1319+579 B 47.4±4.6 4.2±1.0 ... ... 4.4±1.6 ... ... ... ... ... 100.0±8.5 ... ... 132±10 ... ... ... 11.7±3.2 3.08: 1.14: 1.28: 3.8±1.5 279±19 13.1±2.4 2.7±1.1 26.0±3.2 17.8±2.9 ... 6.6±1.6 ... 3.1±1.0 2.01: ... 2.8â1 10 1.97 ± 0.10 0.11 ± 0.03 0.4 ± 0.1 162 ± 11 42 ± 4 15 ± 1 ... SBS 1415+437 C 47.5±2.8 7.11±0.61 ... ... 3.86±0.49 1.05±0.32 2.35±0.23 ... 0.30: 0.56: 100.0±4.9 0.21: 1.41±0.33 107.5±5.5 1.35±0.33 301±14 1.52±0.34 9.61±0.92 3.44±0.43 1.34±0.31 0.96±0.31 1.13±0.35 274±13 4.30±0.45 2.31±0.41 13.15±0.83 9.60±0.63 1.49±0.34 3.77±0.38 ... 2.05±0.33 1.73±0.35 0.96±0.31 6â1 18.51 ± 0.66 0.01 ± 0.02 0.8 ± 0.1 1300 ± 65 222 ± 11 73 ± 4 542 ± 26

Line 4340.47 4363.21 4387.93 4437.55 4471.48 4658.10 4686.00 4711.37 4713.14 4740.16 4861.33 4881.00 4921.93 4958.91 4985.90 5006.84 5015.68 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7281.35 7318.39 7329.66 7751.10 HI [O III] He I He I He I [Fe III] He II [Ar IV] He I [Ar IV] HI [Fe III] He I [O III] [Fe III] [O III] He I He I [O I] [S II I] [O I] [N II] HI [N II] He I [S II] [S II] He I [Ar III] He I [O II] [O II] [Ar III]

f () 0.127 0.121 0.115 0.104 0.095 0.050 0.043 0.037 0.037 0.030 0.000 -0.005 -0.015 -0.024 -0.031 -0.036 -0.038 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.414 -0.418 -0.420 -0.467

A 47.2±2.6 9.98±0.62 0.43±0.17 0.20: 3.85±0.36 ... 0.8: 2.14±0.35 ... 1.61±0.26 100.0±4.9 ... 0.96±0.23 232±12 ... ... ... 11.90±0.86 1.74±0.26 1.56±0.29 0.60±0.18 1.44±0.27 280±13 4.08±0.42 2.97±0.34 7.79±0.56 5.76±0.44 2.69±0.36 4.76±0.47 0.62: 1.30±0.23 1.17±0.24 1.55±0.25 6â1 14.57 ± 0.53 0.03 ± 0.01 0.0 ± 0.1 1530 ± 75 285 ± 14 84 ± 5 ...

C 47.4±3.1 3.73±0.64 1.12: ... 3.88±0.68 0.72: ... ... ... ... 100.0±6.2 ... 0.75: 129.9±7.4 ... ... ... 11.2±1.2 4.77±0.72 1.80±0.52 1.48±0.52 5.50±0.72 281±15 14.8±1.3 2.87±0.66 27.9±1.9 18.9±1.3 2.09±0.54 6.06±0.72 ... 2.69±0.33 1.94±0.27 ... 5.6â1 29 8.18 ± 0.32 0.02 ± 0.02 0.2 ± 0.1 295 ± 23 94 ± 6 23 ± 1 ...

A

47.4±3.8 6.1±1.2 ... ... 2.80±0.66 0.99: ... ... ... ... 100.0±6.3 ... ... 107.3±6.5 1.19: 286±16 ... 7.7±1.1 2.41±0.56 0.98: 0.86: 1.07: 278±17 3.49±0.98 2.28±0.66 9.2±1.1 6.71±0.94 2.09±0.66 3.28±0.61 ... 1.56±0.62 1.35: 0.94: 3.4â1 6 4.07 ± 0.17 0.16 ± 0.03 1.0 ± 0.2 1187 ± 75 130 ± 8 58 ± 5 574 ± 32

Aperture size (arcsec) Distance (arcsec)a F (H )b C (H ) Wabs (° A) -W -W -W -W
a b

° (H) (A) (H ) (° A) (H ) (° A) ([O III]) 5007 (° A)

Relative distance of the knot with respect to the main region in the galaxy. In units of 10-15 erg s-1 cm-2 and not corrected for extinction.


L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

47

Table A.14. Physical conditions and chemical abundances of the ionized gas for the regions analyzed in SBS 1319+579 and SBS 1415+437.
Ob ject Te (O I II) (K) Te (O II) (K) Ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ar+2 /H+ ) 12+log(Ar+3 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]a
a [O/H]=log(O/H)-log(O/H) b Considering the existence of

SBS 1319+579A 13400 ± 500 12400 ± 400 <100 7.22 ± 0.07 7.98 ± 0.06 8.05 ± 0.06 0.77 5.69 6.52 -1.53 5.29 6.09 6.29 -1.76 5.44 5.19 5.65 -2.41 ± ± ± ± ± ± ± ± ± ± ± ± 0.12 0.05 0.09 0.10 0.03 0.12 0.11 0.10 0.07 0.09 0.08 0.11 ... ... ... 10.94 ± 0.04 -0.61 ± 0.11

SBS 1319+579B 11900 ± 800 11300 ± 600 <100 7.73 ± 0.12 7.89 ± 0.10 8.12 ± 0.10 0.16 6.24 6.63 -1.49 5.88 6.13 6.36 -1.76 ± ± ± ± ± ± ± ± 0.19 0.08 0.12 0.12 0.05 0.18 0.15 0.14

SBS 1319+ 579C 11500 ± 600 11050 ± 400 <100 7.75 ± 0.07 7.93 ± 0.08 8.15 ± 0.07 0.18 6.37 6.77 -1.38 5.93 6.39 6.55 -1.60 ± ± ± ± ± ± ± ± 0.13 0.05 0.07 0.10 0.04 0.15 0.12 0.11

SBS 1415+ 437C 16400 ± 600 14500 ± 400 <100 7.07 ± 0.07 7.42 ± 0.04 7.58 ± 0.05b 0.35 5.50 6.01 -1.57 5.38 5.75 5.96 -1.62 5.13 4.56 5.27 -2.31 ± ± ± ± ± ± ± ± ± ± ± ± 0.08 0.05 0.07 0.08 0.03 0.11 0.08 0.12 0.10 0.16 0.12 0.13

SBS 1415+ 437A 15500 ± 700 13850 ± 500 <100 7.05 ± 0.10 7.47 ± 0.05 7.61 ± 0.06 0.42 5.48 6.04 -1.57 5.26 5.68 5.89 -1.72 ± ± ± ± ± ± ± ± 0.14 0.10 0.11 0.09 0.05 0.15 0.13 0.14

5.61 ± 0.12 ... ... ... ... ... ... 10.94 ± 0.11 -0.54 ± 0.15

5.61 ± 0.09 ... ... ... 5.46: 5.80: -2.35: 10.92 ± 0.05 -0.51 ± 0.13

5.14 ± 0.14 ... ... ... 5.24: 5.72: -1.89: 10.77 ± 0.07 -1.05 ± 0.12

5.23 ± 0.12 5.67 ± 0.12 -1.91 ± 0.13 10.75 ± 0.06 -1.08 ± 0.10

, using 12+log(O/H) = 8.66±0.05 (Asplund et al. 2005). O+3 because of the detection of He ii 4686, this value should be 0.01­0.02 dex higher.


48

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.15. Dereddened line intensity ratios with respect to I (H )=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 3554.42 3728.00 3770.63 3797.90 3835.39 3868.75 3889.05 3967.46 4068.60 4101.74 4340.47 4363.21 4471.48 4658.10 4686.00 4814.55 4861.33 4921.93 4958.91 5006.84 5197.90 5200.26 5517.71 5537.88 5754.64 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7318.39 7329.66 7751.10 He I [O II] HI HI HI [Ne III] HI [NeIII]H7 [S II] HI HI [O III] He I [Fe II I] He II [Fe II] HI He I [O III] [O III] [N I] [N I] [Cl III] [Cl III] [N II] He I [O I] [S III] [O I] [N II] HI [N II] He I [S II] [S II] He I [Ar III] [O II] [O II] [Ar III] f () 0.283 0.256 0.249 0.244 0.237 0.230 0.226 0.210 0.189 0.182 0.127 0.121 0.095 0.050 0.043 0.012 0.000 -0.015 -0.024 -0.036 -0.082 -0.083 -0.154 -0.158 -0.198 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.418 -0.420 -0.467 I II Zw 107 A ... 213±12 1.00: 2.37: 3.60±0.95 23.3±2.1 12.6±1.7 17.1±1.7 2.26±0.82 26.2±1.8 46.7±2.8 3.14±0.67 4.08±0.75 1.06: ... ... 100.0±5.3 0.42: 123.7±6.5 375±18 1.08: 0.36: 0.22: 0.20: ... 12.5±1.1 4.56±0.61 1.02±0.23 1.55±0.42 9.98±0.75 282±16 28.7±1.8 3.17±0.50 19.6±1.3 15.9±1.1 2.27±0.40 11.4±0.9 2.12±0.58 2.48±0.66 ... 7.2â1 22.3 ± 0.8 0.68 ± 0.04 2.0 ± 0.3 306 ± 18 44 ± 3 16.4 ± 1.0 172 ± 8 III Zw 107 B ... 306±23 0.54: ... 1.69: 21.1±4.0 11.1±4.2 22.4±2.9 4.69: 26.1±4.2 46.9±4.8 1.97: 3.40: ... ... ... 100.0±8.2 ... 99.5±8.1 293±19 2.07: ... ... ... ... 13.1±2.3 6.1±2.0 ... 2.01: 12.1±1.4 273±18 37.0±3.2 4.01: 36.4±4.3 27.5±3.6 ... 9.4±1.7 ... ... ... 5.6â1 7.2 8.6 ± 0.4 0.15 ± 0.02 1.30 ± 0.10 76 15 5.5 ± 41 ±5 ±2 0.6 ±3 III Zw 107 C ... 20.32: ... ... 5.02: 20.64: 8.84: ... ... 27.5±6.5 47.5±7.9 ... ... ... ... ... 100±15 ... 93±14 257±32 ... ... ... ... ... 12.94: 7.57: ... 3.58: 13.6±4.5 280±35 41.4±8.3 ... 52.2±8.7 35.6±6.9 ... ... ... ... ... 5.2â1 12.4 1.56 ± 0.14 0.22 ± 0.03 0.50 ± 0.10 30 ± 4 4.7 ± 0.7 1.8 ± 0.3 11.6 ± 1.5 Tol 9 INT 3.8±1.4 142±10 0.56: 0.83: 1.94: 10.9±2.0 6.9±2.1 21.6±1.9 3.3±1.1 25.3±2.5 45.8±3.2 0.55: 3.94±0.65 1.09: 0.33: 0.85: 100.0±6.0 ... 78.3±5.2 236±13 ... 2.63±0.78 0.57: 0.64: 0.41: 12.6±1.4 7.78±0.71 0.66: 2.14±0.59 23.8±1.8 267±18 72.2±4.9 3.16±0.64 37.0±2.6 29.8±2.2 1.93±0.51 10.5±1.0 1.96: 1.08: ... 6.4â1 23.4 ± 0.9 0.50 ± 0.05 7.5 ± 0.8 178 ± 12 33 ± 2 12.3 ± 0.8 77 ± 4 Tol 9 NOT ... 177±27 ... ... ... 12.08: ... ... ... 26.3±4.0 46.9±5.8 ... 3.96: ... ... ... 100±10 ... 75.5±9.1 225±21 ... ... ... ... ... 16.3±3.8 7.32±0.96 0.66: 2.30±0.91 24.3±2.3 284±30 82.1±7.3 4.1±1.3 42.1±3.9 35.7±3.4 2.41±0.94 10.8±1.5 1.57±0.43 1.30±0.34 2.68±0.68 3.8â1 3.8â0.3 0.40 ± 0.05 6.2 ± 0.6 186 ± 15 17 ± 3 32 ± 4

Aperture size (arcsec) Distance (arcsec)a F (H )b C (H ) Wabs (° A) - - - -
a b

W W W W

° (H) (A) (H ) (° A) (H ) (° A) ([O III]) 5007 (° A)

Relative distance of the knot with respect to the main region in the galaxy. In units of 10-15 erg s-1 cm-2 and not corrected for extinction.


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49

Table A.16. Physical conditions and chemical abundances of the ionized gas for the regions analyzed in III Zw 107 and Tol 9.
Ob ject Te (O III) (K) Te (O II) (K) Ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) III Zw 107 A 10900 ± 900 10500 ± 800 200 ± 60 7.87 ± 0.11 7.99 ± 0.08 8.23 ± 0.09 0.12 6.70 7.07 -1.16 5.87 6.23 6.42 -1.82 ± ± ± ± ± ± ± ± 0.14 0.06 0.08 0.10 0.06 0.16 0.13 0.15 III Zw 107 Ba 10400 ± 1000 10300 ± 800 <100 8.05 ± 0.14 7.96 ± 0.10 8.31 ± 0.12 -0.09 6.82 7.08 -1.23 ± ± ± ± 0.18 0.08 0.10 0.15 III Zw 107 Ca 10350 ± 1000 10250 ± 800 <100 8.08 ± 0.15 7.92 ± 0.11 8.31 ± 0.13 -0.15 6.88 7.11 -1.20 ± ± ± ± 0.21 0.11 0.12 0.16 Tol 9 INT 7600 ± 1000 8300 ± 700 180 ± 60 8.15 ± 0.18 8.38 ± 0.14 8.58 ± 0.15 0.22 7.37 7.80 -0.78 6.41 6.78 6.97 -1.61 ± ± ± ± ± ± ± ± 0.18 0.08 0.14 0.15 0.08 0.24 0.20 0.17 Tol 9 NOT 7850 ± 1000 8500 ± 800 260 ± 80 8.21 ± 0.19 8.29 ± 0.14 8.55 ± 0.16 0.09 7.37 7.72 -0.84 6.46 6.70 6.92 -1.63 ± ± ± ± ± ± ± ± 0.20 0.09 0.14 0.17 0.08 0.25 0.21 0.18

6.14 ± 0.08 ... ... ... 7.30 ± 0.16 7.65 ± 0.16 -0.66 ± 0.20 5.87 ± 0.14 5.69 ± 0.14 -2.52 ± 0.18 ... ... ... ... 10.99 ± 0.08 -0.35

6.28 ± 0.09 ... ... ... 7.30 ± 0.19 7.68 ± 0.20 -0.62 ± 0.20 ... ... ... ... ... ... ... 10.99: -0.35

12+log(Ne++ /H+ ) 7.26 ± 0.13 12+log(Ne/H) 7.51 ± 0.13 log(Ne/O) -0.73 ± 0.15 12+log(Ar+2 /H+ ) 5.94 ± 0.08 12+log(Ar/H) 5.77 ± 0.09 log(Ar/O) -2.46 ± 0.13 12+log(Cl++ /H+ ) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]
b

7.64 ± 0.18 7.84 ± 0.18 -0.74 ± 0.18 6.30 ± 0.13 6.13 ± 0.13 -2.45 ± 0.20 5.31: 6.14: 6.51: -2.07: 10.93 ± 0.06 -0.08 ± 0.20

7.60 ± 0.23 7.86 ± 0.23 -0.69 ± 0.22 6.28 ± 0.15 6.12 ± 0.15 -2.44 ± 0.22 ... ... ... ... 11.04 ± 0.12 -0.11 ± 0.21

4.32: 5.61: 5.92: -2.31: 10.94 ± 0.05 -0.43 ± 0.14

a Electron temp eratures estimated using empirical b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H)

relations. = 8.66±0.05 (Asplund et al. 2005).


50

L´ ez-S´ hez & Esteban: Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results op anc

Table A.17. Dereddened line intensity ratios with respect to I (H )=100 for knots analyzed in Tol 1457-262 (regions A, B and C) and Arp 252 (galaxy A, ESO 566-8, and galaxy B, ESO 566-7).
Line 3728.00 3797.90 3835.39 3868.75 3889.05 3967.46 4068.60 4101.74 4243.97 4340.47 4363.21 4471.48 4658.10 4686.00 4711.37 4861.33 4958.91 5006.84 5041.03 5055.98 5197.90 5270.40 5754.64 5875.64 6300.30 6312.10 6363.78 6548.03 6562.82 6583.41 6678.15 6716.47 6730.85 7065.28 7135.78 7155.14 7281.35 7318.39 7329.66 7751.10 [O II] HI HI [Ne I II] HI [NeIII]H7 [S II] HI [Fe II] HI [O I II] He I [Fe III] He II [Ar IV] HI [O I II] [O I II] Si I I Si I I [N I] [Fe III] [N I I] He I [O I] [S III] [O I] [N I I] HI [N I I] He I [S II] [S II] He I [Ar III] [Fe II] He I [O II] [O I I] [Ar III] f () 0.256 0.244 0.237 0.230 0.226 0.210 0.189 0.182 0.149 0.127 0.121 0.095 0.050 0.043 0.037 0.000 -0.024 -0.036 -0.044 -0.048 -0.082 -0.100 -0.194 -0.215 -0.282 -0.283 -0.291 -0.318 -0.320 -0.323 -0.336 -0.342 -0.344 -0.387 -0.396 -0.399 -0.414 -0.418 -0.420 -0.467 Tol 1457-262A Tol 1457-262B 224±17 4.46±0.81 4.29±0.79 30.4±8.9 12.0±3.0 21.9±1.8 ... 26.0±2.3 ... 46.7±2.8 8.68±0.77 4.10±0.66 0.85: 1.9±0.4 ... 100.0±5.5 203±10 560±27 ... ... ... ... ... 12.6±2.8 2.88±0.33 1.34±0.24 0.95±0.23 3.13±0.35 285±15 9.32±0.70 3.51±0.42 17.1±1.0 13.81±0.87 3.61±0.47 12.39±0.86 0.38: 0.83±0.21 3.79±0.32 3.02±0.26 2.64±0.33 3.8â1 222 ± 8 0.57 ± 0.03 1.4 ± 0.2 603 ± 32 101 ± 6 31 ± 2 560 ± 27 187±16 ... 11.8±2.1 27.1±9.1 27.3±5.6 25.3±3.5 ... 31.0±4.6 ... 50.8±5.6 10.6±3.1 4.3±1.0 1.12±0.43 2.9±0.3 ... 100.0±9.0 190±12 522±29 ... ... ... ... ... 11.3±2.5 2.23±0.61 1.76±0.56 0.96: 2.00±0.61 286±17 5.7±1.2 3.26±0.94 17.6±2.3 13.3±1.9 3.1±1.1 11.6±1.7 ... ... 3.46±0.75 2.98±0.70 2.80: 1.5â1 9.7 47.7 ± 2.3 0.00 ± 0.05 0.0 ± 0.1 390 ± 28 82 ± 7 24 ± 3 430 ± 25 Tol 1457-262C 270±25 ... ... 23±11 23.1±5.5 23.3±4.5 ... 27.0±5.0 ... 48.7±6.7 6.8±2.3 4.9±1.5 ... ... ... 100±10 163±11 455±28 ... ... ... ... ... 13.4±2.4 4.0±1.1 ... 1.00: 3.8±1.3 279±18 10.4±2.2 3.6±1.3 26.3±3.3 21.2±2.7 3.12: 12.0±2.5 ... ... 4.7±1.3 3.4±1.3 ... ESO 566-8 256±18 ... 6.0±1.3 4.77: 10.2±2.9 14.8±2.1 1.38: 25.9±3.0 1.12: 46.7±4.2 ... 5.0±1.4 1.09: 0.74±0.25 1.25: 100.0±7.1 56.9±4.3 204±12 0.48: 0.33: 2.96±0.79 0.50: 1.23±0.32 14.8±4.4 4.45±0.54 0.23: 1.64±0.32 42.4±3.2 286±17 118.3±7.7 5.03±0.90 36.7±2.7 32.0±2.4 3.47±0.82 7.7±1.6 ... ... 2.98±0.67 2.25±0.60 1.93±0.52 ESO 566-7 280±51 ... ... ... ... ... ... ... ... 46±10 ... ... ... ... ... 100±16 25.9±5.4 75±11 ... ... ... ... ... 16.25: 4.42: ... 3.15: 40.3±7.2 287±35 121±15 ... 89±12 66±10 ... 14.9: ... ... ... ... ...

Aperture size (arcsec) Distance (arcsec)a F (H )b C (H ) Wabs (° A) -W -W -W -W
a b

2.25â1 3.2â1 3.0â1 5.5 49.7 33.6 ± 1.7 111 ± 5 10.2 ± 0.8 0.15 ± 0.02 0.49 ± 0.03 0.23 ± 0.05 0.7 ± 0.1 1.3 ± 0.1 2.7 ± 0.2 342 ± 23 92 ± 9 30 ± 4 411 ± 26 472 ± 29 95 ± 7 38 ± 3 197 ± 12 79 ± 10 13 ± 2 4±2 10 ± 3

° (H) (A) (H ) (° A) (H ) (° A) ([O I II]) 5007 (° A)

Relative distance of the knot with respect to the main region in the galaxy. In units of 10-15 erg s-1 cm-2 and not corrected for extinction.


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51

Table A.18. Physical conditions and chemical abundances of the ionized gas for the regions analyzed in Tol 1457-262 and Arp 252 (galaxy A, ESO 566-8, and galaxy B, ESO 566-7).
Ob ject Te (O III) (K) Te (O II) (K) Ne (cm-3 ) 12+log(O+ /H+ ) 12+log(O++ /H+ ) 12+log(O/H) log(O++ /O+ ) 12+log(N+ /H+ ) 12+log(N/H) log(N/O) 12+log(S+ /H+ ) 12+log(S++ /H+ ) 12+log(S/H) log(S/O) 12+log(Ne++ /H+ ) 12+log(Ne/H) log(Ne/O) 12+log(Ar+2 /H+ ) 12+log(Ar+3 /H+ ) 12+log(Ar/H) log(Ar/O) 12+log(Fe++ /H+ ) 12+log(Fe/H) log(Fe/O) 12+log(He+ /H+ ) [O/H]b Tol 1457-262A 14000 ± 700 12500 ± 600 200 ± 80 7.59 ± 0.10 7.87 ± 0.05 8.05 ± 0.07c 0.27 6.02 6.48 -1.57 5.65 5.95 6.18 -1.88 ± ± ± ± ± ± ± ± 0.11 0.06 0.09 0.11 0.05 0.14 0.10 0.13 Tol 1457-262B 15200 ± 900 14200 ± 700 <100 7.32 ± 0.10 7.74 ± 0.06 7.88 ± 0.07 0.43 5.70 6.27 -1.61 5.53 5.95 6.16 -1.72 ± ± ± ± ± ± ± ± 0.11 0.09 0.12 0.12 0.06 0.16 0.15 0.18 Tol 1457-262C 13400 ± 1100 12400 ± 1000 200 ± 100 7.69 ± 0.15 7.82 ± 0.09 8.06 ± 0.11 0.14 6.10 6.47 -1.59 ± ± ± ± 0.16 0.10 0.15 0.16 ESO 566-8 8700 ± 900 9100 ± 800 300 ± 100 8.23 ± 0.13 8.04 ± 0.10 8.46 ± 0.11c -0.19 7.49 7.71 -0.76 ± ± ± ± 0.17 0.06 0.08 0.12 ESO 566-7a 7900 ± 1000 8500 ± 900 100 ± 150 8.39 ± 0.18 7.82 ± 0.15 8.50 ± 0.16 -0.57 7.57 7.67 -0.82 ± ± ± ± 0.22 0.10 0.11 0.16

5.84 ± 0.09 ... ... ... 6.98 ± 0.20 7.22 ± 0.20 -0.84 ± 0.22 5.77 ± 0.17 ... 5.61 ± 0.17 -2.45 ± 0.20 ... ... ... 10.99 ± 0.08 -0.60 ± 0.16

6.32 ± 0.07 ... ... ... 7.48 ± 0.16 7.90 ± 0.16 -0.56 ± 0.19 6.01 5.48 6.30 -2.17 ± ± ± ± 0.14 0.20 0.17 0.19

6.74 ± 0.12 ... ... ... ... ... ... 6.41 ± 0.18 ... 6.00 ± 0.20 -2.49 ± 0.25 ... ... ... 11.06: -0.16

6.99 ± 0.15 7.17 ± 0.15 -0.88 ± 0.18 5.73 ± 0.08 ... 5.55 ± 0.08 -2.50 ± 0.13 5.43: 5.81: -2.24: 10.99 ± 0.06 -0.61 ± 0.12

6.87 ± 0.18 7.00 ± 0.18 -0.88 ± 0.20 5.66 ± 0.14 ... 5.44 ± 0.14 -2.44 ± 0.18 5.62 ± 0.19 5.98 ± 0.19 -1.90 ± 0.22 10.95 ± 0.10 -0.78 ± 0.12

5.80: 6.00: -2.46: 11.09 ± 0.07 -0.20 ± 0.16

a Electron temp eratures estimated using empirical relations. b [O/H]=log(O/H)-log(O/H) , using 12+log(O/H) = 8.66±0.05 c Considering the existence of O+3 b ecause of the detection of He

(Asplund et al. 2005). ii 4686, this value should be 0.01 dex higher.