ionized gas in the xuv disc of the ngc 1512/1510 system · ionized gas in the xuv disc of ... a...

MNRAS 450, 3381–3409 (2015) doi:10.1093/mnras/stv703

Ionized gas in the XUV disc of the NGC 1512/1510 system

A. R. Lopez-Sanchez,1,2‹ T. Westmeier,3 C. Esteban4,5 and B. S. Koribalski61Australian Astronomical Observatory, PO Box 915, North Ryde, NSW 1670, Australia2Department of Physics and Astronomy, Macquarie University, NSW 2109, Australia3International Centre for Radio Astronomy Research, The University of Western Australia, 35 Stirling Hwy, Crawley, WA 6009, Australia4Instituto de Astrofısica de Canarias, E-38200 La Laguna, Tenerife, Spain5Departamento de Astrofısica, Universidad de La Laguna, E-38205 La Laguna, Tenerife, Spain6CSIRO Astronomy and Space Science, Australia Telescope National Facility, PO BOX 76, Epping, NSW 1710, Australia

Accepted 2015 March 27. Received 2015 March 26; in original form 2014 June 4

ABSTRACTWe present deep, intermediate-resolution, optical spectroscopy of 136 genuine UV-brightregions located in both the inner and outer regions of NGC 1512. This galaxy is in close inter-action with the blue compact dwarf galaxy NGC 1510 and possesses two prominent H I armswhere extended ultraviolet complexes are found. Our data were taken using 2dF/AAOmega atthe 3.9 m Anglo-Australian Telescope and are combined with the H I data from Local VolumeH I Survey and Galaxy Evolution Explorer UV data. We detect ionized gas in 82 per cent ofthe complexes, many of them located between 1 and 6.6 R25. We found significant differencesbetween regions along the Arm 1 – 8.25 � 12+log(O/H) � 8.45 –, and knots located in the ex-ternal debris of Arm 2, –8.40 � 12+log(O/H) � 8.60–. Considering a radial and an azimuthalgradient following the H I arms, we confirm that Arm 2 has experienced an enhancement instar formation because of the interaction with NGC 1510 and flattened the radial metallicityat large radii. Arm 1 appears to retain the original and poorly disturbed radial distribution.We trace the kinematics of the system up to 78 kpc using the Hα emission, which matcheswell that provided by the H I. We estimate that the gas existing at large galactocentric radiihad a metallicity of 12+log(O/H) ∼ 8.1 before the interaction started around 400 Myr ago.The metals within the H I gas are very likely not coming from the inner regions of NGC 1512but probably from material accreted during minor mergers or outflow-enriched intergalacticmedium gas during the life of the galaxy.

Key words: galaxies: abundances – galaxies: dwarf – galaxies: evolution – galaxies: individ-ual: NGC 1510 – galaxies: individual: NGC 1512 – galaxies: kinematics and dynamics.

1 IN T RO D U C T I O N

One of the most surprising discoveries obtained by the GalaxyEvolution Explorer (GALEX) satellite (Martin et al. 2005) was thefinding of UV-bright complexes in the outskirts of nearby spiralgalaxies (Gil de Paz et al. 2005, 2007a; Thilker et al. 2005). Diffusestellar tails and shells were already found surrounding nearby galax-ies using deep optical images (e.g. Malin & Carter 1983) but thesefaint features were thought to be mainly composed by old stars. H II

regions located up to 30 kpc from the main galaxy have been alsofound recently (e.g. Ryan-Weber et al. 2004; Meurer et al. 2006).Indeed, spiral galaxies typically possess a large H I disc that reacheswell beyond their optical size (e.g. Freeman et al. 1977; Walteret al. 2008; Westmeier, Braun & Koribalski 2011; For, Koribalski

� E-mail: [email protected]

& Jarrett 2012), being this material the fuel for present and futurestar-forming events. The extended ultraviolet (XUV) emission dis-covered by GALEX is located well beyond the Hα or B25 radius ofgalaxies, and it seems to exist in ∼20–30 per cent of the local discgalaxy population (Thilker et al. 2007; Zaritsky & Christlein 2007;Lemonias et al. 2011). XUV-discs have even been found aroundE/S0 galaxies (Salim & Rich 2010; Thilker et al. 2010; Moffettet al. 2012). The origin of these UV-bright complexes seems to beyoung stellar clusters associated with a recent or still on-going starformation activity (Gil de Paz et al. 2007b; Bresolin et al. 2009a;Bresolin, Kennicutt & Ryan-Weber 2012). XUV-discs should beembedded in larger H I envelopes – a 2X-H I disc as defined byKoribalski & Lopez-Sanchez (2009) – which are providing the fuelfor their star formation activity (e.g. Koribalski & Lopez-Sanchez2009; Bigiel et al. 2010b; Werk et al. 2010a).

The study of star formation processes in galaxy outskirts giveskey clues about the formation and evolution of galaxies, not only

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3382 A. R. Lopez-Sanchez et al.

because these regions may probe physical conditions similar tothose present of the early Universe, but also because they test theinside–out scenario of galaxy disc formation. In the first case, theanalysis of XUV-complexes in the outer regions of galaxies com-plements the study of dwarf galaxies as building blocks of largerobjects but also the nature of tidal dwarf galaxies (TDGs), as wellas the viability of star formation in regions of low gas density. Inthe second case, these UV-bright regions provide a powerful tool toinvestigate the mass and chemical enrichment of galaxies from theircentre to their external regions via colours, gas-to-stars ratios, andmetallicity gradients. Following the inside–out scenario, the galac-tic stellar disc grows gradually with time as it accretes gas fromtheir gaseous, extended discs (e.g. White & Frenk 1991; Bouwens,Cayon & Silk 1997) producing a relatively quick self-chemical en-richment and an almost universal negative metallicity gradient oncethis is normalized to the galaxy optical size (Zaritsky, Kennicutt &Huchra 1994; Boissier & Prantzos 1999, 2000). Indeed, detailedobservations using integral-field spectroscopy (IFS) data from theCALIFA survey (Sanchez et al. 2012; Husemann et al. 2013) haveproven that a characteristic oxygen abundance gradient does existin galaxy discs (Rosales-Ortega et al. 2012; Sanchez et al. 2014).

However, it has been observationally found that the outskirtsof galaxies do not follow the same pattern as their main disc. Inparticular, a flattening of the metallicity gradient in the externalregions of spiral galaxies has been found (Bresolin et al. 2009a;Kewley et al. 2010; Rupke, Kewley & Barnes 2010; Werk et al.2010a,b; Bresolin, Kennicutt & Ryan-Weber 2012), and it is alsodetected in our Milky Way (e.g. Vilchez & Esteban 1996; Estebanet al. 2013). It has been suggested that the flat metallicity gradientin the outer discs is a consequence of galaxy interactions, howeverthere are still some non-interacting galaxies where this behaviour isalso observed (Werk et al. 2011; Sanchez et al. 2014). In fact, for theMilky Way (Esteban et al. 2013) suggest that a levelling out of thestar formation efficiency beyond the isophotal radius can explainthe flattening of the abundance gradients in the external Galacticdisc. In any case, the finding of a flat metallicity gradient in theexternal regions of galaxies informs about the metal redistributionin these systems, an aspect which may also have consequences inthe interpretation of the scatter in the observed mass–metallicityrelations (Lequeux et al. 1979; Tremonti et al. 2004; Kewley &Ellison 2008; Lara-Lopez et al. 2010; Lara-Lopez, Lopez-Sanchez& Hopkins 2013c; Lara-Lopez et al. 2013b).

Hence, the analysis of the ionized gas within XUV complexesfound at large galaxy radii allows us to investigate the nature, phys-ical conditions, chemical abundances, and kinematics of the inter-stellar medium (ISM) in these regions. When combining with UVand 21-cm H I observations, these data provide important constraintsto the star formation activity and star formation history of galaxies,the relationships between gas and stars, as well as the physics behindthem. We are therefore conducting such study of XUV discs foundin nearby, H I-rich galaxies included in the Local Volume H I Survey(LVHIS) project (Koribalski 2008, Koribalski et al., in preparation)using the multi-object fibre-positioner instrument Two-degree Field(2dF) and the AAOmega spectrograph installed at the 3.9 m Anglo-Australian Telescope (AAT, Siding Spring Observatory, NSW, Aus-tralia). The 2dF instrument is perfect for our analysis, as it providesa very flexible position of fibres within a wide field of view, allow-ing the simultaneous observation of regions within both the XUVand spiral discs. In particular, with these observations, we want toexplore (i) how many XUV complexes show ionized gas emission(i.e. on-going star formation activity); (ii) what is the correlationbetween the ages of the dominant young stellar population (from

UV colours) and the most recent star formation event (from Hα

emission); (iii) what is the metallicity distribution throughout thesystem, studying not only radial but azimuthal chemical gradients,and how this affects the metal redistribution in galaxies; (iv) what isthe nature of the XUV-complexes and search for tidal dwarf galaxy(TDG) or external dwarf galaxies candidates; and (v) what are therelationships between the neutral gas, the stellar mass, the metallic-ity, and the star formation in these regions and the comparison withwhat is observed in dwarf galaxies.

Here, we present the results of our analysis of the XUV emissionin the impressive galaxy pair NGC 1512/1510. At the adopted dis-tance of 9.5 Mpc, 1 arcmin corresponds to 2.49 kpc, and 2 arcsec(the diameter of a 2dF fibre) corresponds to 83 pc. This systemhosts hundreds of independent UV-bright regions associated withdense H I clouds in the outskirts, the blue compact dwarf galaxy(BCDG) NGC 1510 (which is located at just 13.8 kpc from thecentre of NGC 1512), a central star-forming ring and two TDGcandidates at very large distance (projected radius of 83 kpc) iden-tified by their H I emission. A review of the properties of theNGC 1512/1510 system, as well as a very detailed UV and H I

analysis, can be found in Koribalski & Lopez-Sanchez (2009, here-after KLS09). Recently, Bresolin et al. (2012) obtained deep FORS2– 8 m Antu unit of the European Southern Observatory Very LargeTelescope on Cerro Paranal, Chile – optical spectroscopy of 62 UV-bright complexes within NGC 1512, including star-forming regionsin the optical disc and in XUV knots of NGC 1512, although themajority (all but 16) were located within the effective radius, Re

(4.26 arcmin = 11.76 kpc), of the galaxy. In six of the observedknots they were able to determine the oxygen and nitrogen abun-dances following the direct Te method, yielding to average valuesof 12+log(O/H) ∼ 8.17 ± 0.09 and log(N/O) ∼ −1.32 ± 0.12.

2 O B S E RVAT I O N S A N D DATA R E D U C T I O N

The observations were carried out at the 3.9 m AAT at Siding SpringObservatory (NSW, Australia) between 2008 Nov 29 and Dec 2. Weused the 2dF instrument (Lewis et al. 2002) in combination with theAAOmega spectrograph (Saunders et al. 2004; Smith et al. 2004;Sharp et al. 2006) to get our spectroscopic data. The 2dF instrument,which is installed at the prime focus of the AAT, consists of a robotgantry which positions up to 400 optical fibres on a plate with afield of 2◦ diameter projected on the sky. The accuracy of fibrepositioning is 0.3 arcsec. 2dF also possesses a wide-field corrector,an atmospheric dispersion compensator (ADC), and a tumblingmechanism with two field plates which allows the next field to beconfigured while the current field is being observed. Eight fibresactually are guide fibre-bundles which are used to ensure accuratetelescope positioning during the observation. The other 392 targetfibres from 2dF, each one with a 2 arcsec diameter projected on thesky (Lewis et al. 2002), are fed to the AAOmega spectrograph.

In comparison to long-slit spectroscopy, multi-object fibre spec-troscopy allows us to be more flexible in choosing the areas ofinterest and to save observing time as all pointings can be done inone exposure. For the case of the galaxy pair NGC 1512/1510, weneeded to allocate fibres within a field of view of ∼1◦.

The AAOmega spectrograph, which is stationed in the thermallystable environment of one of the telescope’s Coude rooms, possessesa dual beam system. Each arm of the AAOmega system is equippedwith a 2k × 4k E2V CCD detector and an AAO2 CCD controller. Inorder to achieve the best compromise between sensitivity, spectralresolution and wavelength coverage, we used the 580V grating inthe blue arm and the 2000R grating in the red arm. The 5700 Å

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Ionized gas in the XUV disc of NGC 1512/1510 3383

dichroic was used to split the light. The 580V grating, which hasa dispersion of 1 Å pixel−1, was centred at 4550 Å, and hence itgave a coverage of 3500–5600 Å. The 2000R grating, which has adispersion of 0.23 Å pixel−1, was centred at 6530 Å, and hence itgave a coverage of 6300–6750 Å. With this setup, we can measureall the critical diagnostic optical lines, [O II] λ3727, Hγ , [O III]λ4363, He II λ4686, Hβ, [O III] λλ4959,5007, [N II] λ6548, Hα,[N II] λ6583, He I λ6678 and [S II] λλ6717,6731, but we will alsoget the highest dispersion available in the Hα line (0.23 Å pixel−1),giving a spectral resolution of ∼35 km s−1.

We used the GALEX images to get accurate RA and Dec. posi-tions of all the UV-bright star-forming regions observed by KLS09.We then ran the 2dF CONFIGURE software, which implements a Simu-lated Annealing algorithm (Miszalski et al. 2006) to allocate opticalfibres to the UV-rich star-forming regions. We also allocated 25fibres for sky positions. We observed two different plate configura-tions, including some objects to be observed in both plates to checkthe quality of our analysis. The 2dF CONFIGURE software allocated90 UV-selected knots in Plate 1 and 74 in Plate 2, being a total of164 regions. 28 objects (including NGC 1510 and the two TDGcandidates) were included in both plates, hence we observed 136genuine targets. For each plate, we obtained 3 × 2400 s exposures,as well as the calibration frames (arcs and flat-fields).

The spectrophotometric stars GD 108 (Oke 1990) and HR 3454(Hamuy et al. 1994) were used to get the absolute flux calibrationof the 1D spectra. At least two exposures of 300 s (for GD 108)or 10 s (for HR 3454) were taken immediately after the scienceframes were obtained. Calibration frames were also taken for thespectrophotometric stars.

The raw data were processed using software developed at theAustralian Astronomical Observatory (AAO) called 2DFDR (Croom,Saunders & Heald 2004; Sharp & Birchall 2010). The 2DFDR pro-cessing applies the standard sequence of tasks for 1D spectral extrac-tion from 2D images. This includes bias subtraction, flat-fielding,tramline fitting, wavelength calibration, and sky subtraction. A mas-ter bias was created using the available bias frames, but for eachplate configuration observed, the raw AAOmega frames were runthrough 2DFDR to provide the processed spectra. The sky subtrac-tion was performed using the 25 fibres allocated in each plate tosky positions. The continuum sky subtraction accuracy is typically2–3 per cent of the sky level.

We later used standard IRAF1 tasks to perform the absolute fluxcalibration of the 2D spectra using the data obtained for the spec-trophotometric standard star. The correction for atmospheric ex-tinction was achieved using an average curve for the continuousatmospheric extinction at Siding Spring Observatory. We estimatethat the flux-calibrated spectra are typically accurate to ∼5 per cent,although this number may reach 10 per cent at the blue end of theblue camera (wavelengths lower than 3800 Å) because of the de-crease of the sensitivity of the instrument in this spectral range.

3 A NA LY SIS AND RESULTS

This section details the analysis of the ionized gas within theUV-bright knots observed using our intermediate-resolution opti-cal spectroscopy and provides the results (Hα fluxes and equivalentwidths, nature of the ionization, reddening, oxygen and nitrogen

1 IRAF (Image Reduction and Analysis Facility) is distributed by NOAOwhich is operated by AURA Inc., under cooperative agreement with NSF.

abundances) which will be later used for the discussion of the phys-ical and chemical properties of the NGC 1512/1510 system.

3.1 Identification of the UV-bright regions

Fig. 1 shows the deep NUV image of the NGC 1512/1510 galaxypair with the identification of all the UV-rich regions observed usingthe 2dF/AAOmega instrument at the AAT. The UV-knots within thetwo TDG candidates (see fig. 8 in KLS09) were also observed butthey do not appear in this figure. We have coloured each regionaccording to the results of its analysis, as explained below andit is compiled in Table 1. Each 1D spectrum of the 164 observedtargets was carefully inspected by eye to identify the bright emissionlines such as [O II] λ3727, Hβ, [O III] λ5007 and Hα. We identified153 regions (93.3 per cent). The 11 regions for which we could notdetermine their nature are coloured in red in Fig. 1. From these 153regions, 17 were identified as background galaxies, as the brightemission lines were identified at radial velocities larger than theexpected radial velocity of the regions found in the system (between∼750 and ∼1100 km s−1). Background galaxies are coloured incyan in Fig. 1. We also identified a foreground star thanks to thepresence of strong absorption lines (Ca II H, K, G band, Hγ , Hβ,Hα) found at a radial velocity of −96.8 km s−1. This object has beencoloured in blue in Fig. 1. 135 regions (82.3 per cent of the observedregions; 88.2 per cent of the identified regions) have been classifiedas emission-line objects truly belonging to the NGC 1512/1510system. Hence, our 2dF/AAOmega observations confirm that thehuge majority of the observed UV-rich star clusters has ionized gasemission.

Table 1 compiles the statistics of the results of our 2dF/AAOmegaobservations. Colours grey, pink, orange, yellow and green indicatethe number of emission lines detected in the spectrum of eachregion. Almost half of the regions identified within the galaxy pair(47.4 per cent) show the five emission lines needed([O II] λ3727,Hβ, [O III] λ5007, Hα, and [N II] λ6583) for computing the oxygenabundance of the ionized gas using several methods. However, weare able to estimate the metallicity of the ionized gas in 120 regions(88.9 per cent of the detected regions within the NGC 1512/1510system) as these are the objects for which we identify at least Hα

and [N II] λ6583.Because of the results we found in the further analysis of the

data, instead of identifying five areas within the NGC 1512/1510system as KLS09 did, we have re-classified all UV-rich knots infour zones: the star-forming ring in NGC 1512 (Ring, Zone 1), thelong XUV arm at the east (Arm 1, Zone 2), the arm between thelong XUV arm and the ring of NGC 1512 (internal arm, Zone 3)and the knots found somewhat dispersed at the W and NW (externaldebris of Arm 2, Zone 4). We then numbered the observed knotswithin these four regions that we use to clearly identify each onethroughout our study. Table 2 compiles the coordinates of each oneof the 164 targeted regions.

3.2 Emission-line analysis

The IRAF software was used to analyse the 1D spectra. Linefluxes and equivalent widths were measured by integrating allthe flux in the line between two given limits and over a localcontinuum estimated by eye. We want to emphasize that, oppo-site to an automatic process, the visual inspection of the spec-tra is needed to get a proper estimation of the adjacent contin-uum and hence a reliable line flux estimation when emission lines

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Figure 1. Deep NUV image of the NGC 1512/1510 galaxy pair. The UV-bright regions observed using 2dF/AAOmega at the AAT that have been identifiedand analysed are indicated using a code of digits and colours. The first digit indicates the zone where the region is located. We distinguish four zones: Ring(1), Arm 1 (2), internal arm (3) and external debris of Arm 2 (4). The second digit is simply a sequential number. The colour indicates the emission linesdetected in their spectrum (see Table 1). Two regions were observed in NGC 1510: one corresponds to its centre, while the other is at the west of the galaxy.The UV-knots within the two TDG candidates (see fig. 8 in KLS09) were also observed but they do not appear here. North is up and east to the left.

are faint. The errors associated with the line flux measurementswere estimated considering the noise in the adjacent continuum,the width of each emission line, and the photon noise at the lineprofile.

Table 2 compiles the dereddened line intensity ratios and theirassociated errors for all the regions. This table includes the lineintensity ratio of [O II] λ3727, Hγ , [O III] λ5007, and Hα with

respect to Hβ and their associated errors for all the regions. The [N II]λ6583/Hα ratio is also included as it is one of the magnitudes thatwe will use to estimate the oxygen abundance Table 2 also includesthe observed Hα flux (uncorrected for extinction), the radial velocityat which the Hα line is observed, and the equivalent widths of theHα, Hβ, and Hγ lines. The error in the Hα flux was computedconsidering the quadratic sum of the error in its flux measurement

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Table 1. Statistical details of the results of our 2dF/AAOmega observations. Lastcolumn indicates the colour each region is identified in Fig. 1.

Fibre Per cent Per cent Colour innumber total class Fig. 1

Observed 164 100.0 – –Non-identified 11 6.7 – redIdentified 153 93.3 – –

Background galaxies 17 10.4 11.1 cyanForeground stars 1 0.6 0.7 blueRegions in NGC 1512/1510 135 82.3 88.2 –

Only Hα detected 15 9.1 11.1 greyHα and [N II] 30 18.3 22.2 yellowHα, Hβ, and [N II] 6 3.7 4.5 pinkHα, Hβ, [N II], and [O III] 20 12.2 18.8 orangeHα, Hβ, [N II], [O III], and [O II] 64 39.0 47.4 greenHγ detected 51 31.1 37.8 –

and the error in flux calibration, which we assumed to be 5 per cent.Colons indicate errors of the order or larger than 40 per cent.

Fig. 2 shows some examples of the optical spectra obtained inour analysis of UV-bright regions within the NGC 1512/NGC 1510system. The two top panels plot the spectrum obtained forNGC 1510 in each plate configuration. The fibre in Plate 1 waslocated in the central, brightest part of this BCDG, whereas the fi-bre in Plate 2 was located slightly towards the west. Fig. 2 also plotsthe spectrum of Zone 4 Region 17, for which we derive a relativelylow oxygen abundance following the direct Te method (see below),and the spectrum of a typical UV-bright region within the XUV discof the NGC 1512/1510 system.

3.3 Nature of the ionized gas

We checked the nature of the ionization of the UV-bright regionsusing the so-called diagnostic diagrams, as first proposed by Bald-win, Phillips & Terlevich (1981) and Veilleux & Osterbrock (1987),which are now extendedly used to distinguish between pure star-forming regions and AGN/low-ionization nuclear emission-line re-gion activity. Fig. 3 plots the typical [O III] λ5007/Hβ versus [N II]λ6583/Hα diagram for the 84 regions for which the four emissionlines are detected. We distinguish between objects observed usingPlate 1 (blue stars) and those observed using Plate 2 (purple dia-monds) in order to quantify if any offset is observed between thetwo data sets. We used the analytic relations given by Dopita et al.(2000) and Kewley et al. (2001), as well as the empirical relationprovided by Kauffmann et al. (2003), to check the nature of theexcitation mechanism of the ionized gas within these regions. Aswe expected, all regions lie below the Kewley et al. (2001) theo-retical line for starburst galaxies, which considers continuous starformation. This clearly indicates that photoionization is the mainexcitation mechanism of the gas and that there is no evidence fora significant contribution from shock excitation. Interestingly, theloci of the UV-rich regions within this diagram agree well with thepredictions given by the models provided by Dopita et al. (2000),which considered extragalactic H II regions with instantaneous starformation.

Considering that a single O7V star has a Hα luminosity ofLO7V = 1.36 × 1037 erg s−1 (Schaerer & Vacca 1998), using theextinction-corrected Hα flux we estimate that the number of mas-sive ionizing stars in these regions typically is between 1 and 5. This

value agrees with the results found in XUV complexes surroundingother nearby spiral galaxies (Gil de Paz et al. 2007b).

3.4 First estimation of oxygen abundances

The preferred technique for determining the chemical abundancesof the ionized gas is via the Te method, which uses electron temper-ature sensitive lines such as the [O III] λ4363 line (e.g. Peimbert &Costero 1969; Osterbrock & Ferland 2006). However, the aurorallines are much fainter than the bright nebular lines and therefore arerarely detected in galaxies, particularly at high metallicities. Hence,strong emission line (SEL) methods must be used to derive oxy-gen abundances in external galaxies. These empirical techniquesare widely used today. The most-common SEL calibrations involveusing the N2 and O3N2 parameters, defined as

N2 = log

(I ([N II]λ6583)

I (Hα)

), (1)

O3N2 = log

(I ([N II]λ6583)

I (Hα)

/I ([O III]λ5007)

I (Hβ)

). (2)

The huge advantage of using such ratios is that they consider emis-sion lines which are very close in wavelength, and hence both thereddening correction and the absolute flux calibration can be ne-glected. However, these SEL techniques do not take into account anyionization parameter, and hence accurate oxygen abundances (i.e.with uncertainties lower than ∼0.10 dex) cannot be derived. Fur-thermore, precaution should always be taken when using any SELmethod. Reviews of the most-common empirical calibrations andtheir limitations can be found in Kewley & Ellison (2008), Lopez-Sanchez & Esteban (2010b), and Lopez-Sanchez et al. (2012). Werefer the reader to these studies for more details.

Here, we first used the empirical calibrations provided by Pettini& Pagel (2004), which considers the N2 and O3N2 ratios, to geta first estimation of the oxygen abundance of the UV-rich regionsdetected within the NGC 1512/1510 system. We used the linearrelation between the oxygen abundance and the N2 ratio,

12 + log(O/H) = 8.90 + 0.57N2, (3)

for the 120 regions for which we detect both emission lines, and thelinear relation between O/H and the O3N2 ratio,

12 + log(O/H) = 8.73 − 0.32O3N2, (4)

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loci

tyes

timat

edus

ing

the

HIB

alm

erab

sorp

tion

lines

.

for the 84 regions for which the four lines are detected. Table 3compiles the oxygen abundance estimated from these SEL methods,including the value of the N2 and O3N2 parameters for each region.

3.5 Correction for reddening

When at least Hα and Hβ are available, we used the method de-scribed in Lopez-Sanchez & Esteban (2009) to correct all line fluxesfor both reddening and underlying stellar absorption. This is an iter-ative procedure to derive simultaneously the reddening coefficient,c(Hβ), and the equivalent widths of the absorption in the hydrogenlines, Wabs, to correct the observed line intensities for both effects.The method also assumes that Wabs is the same for all the Balmerlines and uses the relation given by Mazzarella & Boroson (1993)to perform the absorption correction,

c(Hβ) = 1

f (λ)log

⎡⎣

I (λ)I (Hβ) ×

(1 + Wabs

WHβ

)F (λ)

F (Hβ) ×(

1 + WabsWλ

)⎤⎦ , (5)

for each detected hydrogen Balmer line, where F(λ) and I(λ) arethe observed and the theoretical fluxes (unaffected by reddeningor absorption), Wabs, Wλ, and WHβ are the equivalent widths of theunderlying stellar absorption, the considered Balmer line and Hβ,respectively, and f(λ) is the reddening curve normalized to Hβ usingthe Cardelli, Clayton & Mathis (1989) extinction law.

However, we do not consider the ‘standard’ value of the theoret-ical Hα/Hβ ratio for Case B recombination, 2.86, as this numberactually depends on both the electron temperature, Te, and the elec-tron density, ne, of the gas. The same happens to other H I Balmerratios. Appendix discusses this issue and provides relationships be-tween the electron temperature and the theoretical Hα/Hβ, Hγ /Hβ,and Hδ/Hβ ratios, using the Storey & Hummer (1995) calculationsand assuming ne = 100 cm−3. More importantly, Appendix A alsogives empirical relationships between the oxygen abundance andsuch theoretical Balmer ratios. We here used those equations andthe oxygen abundance determined following the Pettini & Pagel(2004) calibrations to get the theoretical Hα/Hβ ratio (and theHγ /Hβ ratio if available) for each individual region. The resultsare listed in Table 2, as these will be the H I Balmer ratios obtainedafter applying the correction for reddening and underlying stellarabsorption. Equivalent widths are also compiled in Table 2.

The bottom panel of Fig. 4 compares the values derived for c(Hβ)following the Hα/Hβ and Hγ /Hβ ratios. In order to quantify theWabs which best matches both values, we looked for the linear fit be-tween both reddening coefficient that was providing r ∼ 0 (i.e. thereis not correlation between both variables). This analysis yielded anaverage value of Wabs = 0.85 Å, corresponding to a dispersiond = 0.071 dex. We note that the value of c(Hβ) derived from theHγ /Hβ ratio is systematically 0.05 dex higher than the reddeningcoefficient derived using the Hα/Hβ ratio. This is very probably aconsequence of the much higher uncertainty in the measurement ofthe Hγ emission line.

We then assumed an average value of Wabs = 0.85 Å and appliedequation (5) to compute the reddening coefficient c(Hβ) using theHα/Hβ ratio. If Hγ is observed, we also applied equation (5) toderive c(Hβ) using the Hγ /Hβ ratio. However, because of the lowsignal-to-noise (S/N) ratio of the Hγ line in the majority of thecases, the reddening coefficient obtained using the Hα/Hβ ratio(and not the average value between both reddening coefficients) isconsidered in all cases for our analysis. Table 3 lists the reddeningcoefficient derived using the Hα/Hβ ratio for 90 regions both lines

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Figure 2. Example of the optical spectra obtained in our analysis of UV-bright regions within the NGC 1512/NGC 1510 system. We plot four regions:NGC 1510 (which has the highest S/N ratio), NGC 1510 W, and regions 4_17 and 2_2 (which has a medium-low S/N ratio and exemplifies the typical spectrumobtained for the observed UV-bright regions). The most important emission lines are labelled. The spectra have been corrected for radial velocity and skyemission, but not for extinction.

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Figure 3. Comparison of the observed [O III] λ5007/Hβ and[N II] λ6583/Hα flux ratios obtained for the regions where all four emis-sion lines are available, with the diagnostic diagrams proposed by Dopitaet al. (2000) – D00, black continuous line – Kewley et al. (2001) – K01, reddotted line – and the empirical relation provided by Kauffmann et al. (2003)– Ka03, dashed green line. Blue stars indicate regions observed using Plate 1,whilst purple diamonds correspond to regions observed using Plate 2.

are detected, as well as c(Hβ) using the Hγ /Hβ ratio for the 51regions the three H I Balmer lines are measured.

The top panel of Fig. 4 compares the oxygen abundance andthe reddening coefficient derived for the regions these values areavailable. Although with a considerable scatter, it is evident thatthe reddening coefficient (i.e. the amount of dust and obscurationin each region) increases with increasing metallicity. A linear fit tothe data,

c(Hβ) = (−3.34 ± 0.41) + (0.449 ± 0.053)x, (6)

where x = 12+log(O/H), is shown with a dotted black line in thefigure. The correlation coefficient of the fit is r = 0.6450. This be-haviour was already observed by Lopez-Sanchez & Esteban (2010b)when analysing the physical and chemical properties of a sample ofstrong starburst galaxies. A similar behaviour, but within the Balmerdecrement and the luminosity (which scales with the oxygen metal-licity following the well-known luminosity–metallicity relation),was reported by Gunawardhana et al. (2013) when analysing theproperties of a large sample of star-forming galaxies using data ofthe Galaxy And Mass Assembly (GAMA) survey (Driver et al. 2011;Hopkins et al. 2013). Interestingly, the position of the BCDG NGC1510 lies far from the observed relation.

3.6 Recomputing oxygen abundances

Once the reddening is known and in those regions where we measurethe [O II] λ3727 doublet, we can re-compute the oxygen abundancesusing additional SEL methods. We use the R23 parameter, Ta

ble

3.R

edde

ning

coef

ficie

nt,i

mpo

rtan

tem

issi

on-l

ine

ratio

san

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rica

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tions

and

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ved

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unda

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and

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our

sam

ple

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ons.

The

full

tabl

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avai

labl

eon

line.

Zon

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nPl

ate

c(H

β)

R2

R3

R23

Py

N2

O3N

2N

2O

212

+log

(O/H

)a

log(

N/O

)H

α/H

βH

γ/H

βPP

04a

PP04

bP0

1K

K04

Ado

pted

0N

1510

10.

39±

0.01

0.55

6.10

±0.

042.

99±

0.02

9.08

±0.

060.

671

±0.

007

0.31

0±

0.00

6−1

.074

±0.

001

1.72

9±

0.00

4−1

.101

±0.

004

8.29

8.18

8.22

8.44

8.23

±0.

06−1

.34

±0.

03”

”2

0.27

±0.

010.

393.

27±

0.07

4.76

±0.

078.

03±

0.16

0.40

8±

0.01

3−0

.163

±0.

017

−0.8

76±

0.00

81.

262

±0.

017

−1.1

02±

0.01

58.

408.

338.

108.

518.

36±

0.07

−1.3

9±

0.03

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54:

––

8.24

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8.67

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67±

0.13

–1

22

0.25

±0.

010.

270.

48±

0.05

1.25

±0.

061.

74±

0.12

0.27

9±

0.03

7−0

.413

±0.

062

−0.4

23±

0.01

0−0

.021

±0.

050

−0.0

60±

0.02

98.

668.

748.

719.

108.

72±

0.07

−0.7

2±

0.03

13

10.

44±

0.08

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––

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–−0

.370

±0.

038

−0.1

22:

–8.

698.

77–

–8.

77±

0.15

–1

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0.39

±0.

050.

400.

68±

0.18

1.23

±0.

171.

91±

0.41

0.36

±0.

14−0

.26

±0.

16−0

.380

±0.

026

0.08

±0.

12−0

.008

±0.

083

8.68

8.70

8.72

9.09

8.71

±0.

08−0

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±0.

061

52

0.24

±0.

010.

290.

30±

0.02

1.36

±0.

041.

66±

0.07

0.18

1±

0.01

7−0

.656

±0.

045

−0.4

32±

0.00

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.220

±0.

037

−0.1

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0.01

78.

658.

808.

689.

128.

74±

0.07

−0.7

8±

0.03

16

1–

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7±

0.16

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8.69

––

–8.

69±

0.15

–

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R23

,y.

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Figure 4. (Top panel) Comparison between the reddening coefficient,c(Hβ) derived from the Hα/Hβ ratio (y-axis) and the oxygen abundanceestimated for each knot (x-axis). (Bottom panel) Comparison between thec(Hβ) derived from the Hα/Hβ ratio (x-axis) and the difference betweenthe c(Hβ) derived from the Hγ /Hβ and the Hα/Hβ ratio. A Wabs= 0.85 Åwas used in both panels. The dotted lines show a fit to the data, excludingNGC 1510. Blue stars indicate regions observed using Plate 1, whilst purplediamonds correspond to regions observed using Plate 2.

R3 = I ([O III])λ4959 + I ([O III])λ5007

Hβ, (7)

R2 = I ([O II])λλ3726, 3729

Hβ, (8)

R23 = R3 + R2, (9)

and the P (Pilyugin 2000) or y (McGaugh 1991) parameters,

P = R3

R23, (10)

y = logR3

R2= log

1

P −1 − 1, (11)

to take into account the hardness of the ionizing radiation. Follow-ing Lopez-Sanchez et al. (2012), we use the empirical calibrationsbetween R23 and P and the oxygen abundance for high- and low-metallicity H II regions provided by Pilyugin (2001a,b), respectively,as well as the Kobulnicky & Kewley (2004) method, which is basedon photoionization models and uses the R23 and y parameters.

We also have to consider that calibrations involving the R23 pa-rameter are bivaluated; these are given for 12+log(O/H) � 8.0 (lowmetallicity) and 12+log(O/H) � 8.3 (high metallicity). In the case ofthe UV-bright regions observed in the NGC 1512/1510 system, wedo not have any with 12+log(O/H) � 8.0, but there are a few objectsin the intermediate metallicity range, 8.0 � 12+log(O/H) � 8.3.In these cases, we just considered a weighted oxygen abundancefollowing

xmed = xlow

(cu − xPP04

cd

)+ xhigh

(xPP04 − cl

cd

), (12)

where x = 12+log(O/H), cl = 8.0, cu = 8.3, and cd = cu−cl = 0.3,and the subindices high, low, med, and PP04 refer to the oxy-gen abundance given by the high-metallicity branch, the oxygen

Figure 5. Comparison of oxygen abundances derived following the Pettini& Pagel (2004) and Pilyugin (2001a,b) methods. Top panel compares themetallicity provided by the N2 parameter (x-axis) with the metallicity differ-ence obtained between the O3N2 and the N2 calibrations (y-axis) followingPettini & Pagel (2004). The red dashed vertical line at 12+log(O/H) = 8.65indicates the upper limit the N2 calibration is valid. Bottom panel comparesthe average metallicity value provided by the Pettini & Pagel (2004) meth-ods (x-axis) with the metallicity difference between the Pilyugin (2001a,b)method (which uses the R23 and P parameters) and the average value pro-vided by the Pettini & Pagel (2004) method (y-axis). The red dashed verticalline at 12+log(O/H) = 8.30 indicates the lower limit to the high-metallicitybranch. Blue stars indicate regions observed using Plate 1, whilst purplediamonds correspond to regions observed using Plate 2.

abundance given by the low-metallicity branch, the adopted oxy-gen abundance, and the oxygen abundance computed followingthe Pettini & Pagel (2004) method we discussed in Section 3.4,respectively.

Table 3 compiles the R23, P and y parameters used for the SELcalibrations which considers [O II] λ3727 for each region, as well asthe oxygen abundances derived from those data. Fig. 5 compares theoxygen abundances derived following the Pettini & Pagel (2004)and Pilyugin (2001a,b) methods. The red dashed vertical line at12+log(O/H) = 8.65 indicates the upper limit the N2 calibration isvalid, as this parameter saturates at high metallicity (for an extendeddiscussion and important consequences when analysing metallici-ties of large galaxy samples, see Lara-Lopez et al. 2013c). We finda relatively good agreement between the three methods, with anaverage scatter of only 0.049 dex (for the N2–O3N2 comparison)and 0.060 dex (for the N2, O3N2–R23, P comparison), as it seen inFig. 5.

Hence, we adopt that the valid oxygen abundance for each regionis the average value of the N2, O3N2, and R23, P techniques. Wenote that we have not considered the value given by the N2 cali-bration when the oxygen abundance estimated by the O3N2 methodwas higher than 12+log(O/H) = 8.65. The second last column inTable 3 lists the final adopted oxygen abundance determined foreach region. The error in the oxygen abundance provided by eachcalibration was estimated from the uncertainties associated with theflux ratios and assuming that each SEL technique has an uncertaintyof ±0.10 dex. We then estimated the final uncertainty consideringthe number of SEL calibrations averaged and the dispersion of theirvalues. In the majority of the cases, the dispersion of the oxygenabundances provided by the empirical calibrations was inferior to

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±0.10 dex. In many knots, it was even within ±0.05 dex, i.e. thedispersion is of the order or inferior to the typical uncertainties ofthe individual calibrations. The final uncertainty associated withthe oxygen abundance of each knot is also listed in the second lastcolumn of Table 3. For the best cases, the final uncertainty is just±0.05 dex. Colons indicate errors larger than 0.3 dex.

3.6.1 Comparison between results of regions observed twice

Figs 3–9 distinguish between regions observed using plate 1 (bluestars) and plate 2 (purple diamonds). This has been done to inves-tigate if any correlation depending on the plate exists. As we see,a plate dependence is not observed, and indeed the scatter in thefigures is essentially the same considering each plate independently.Fig. 6 explicitly checks the validity of this argument comparing theresults obtained for the regions for which we have observationsin both plates. The top panel in Fig. 6 plots the oxygen abun-dance difference between plates for each region following the N2

method (blue stars), O3N2 method (red diamonds), and the Pilyugin(2001a,b) method (purple triangles). In all cases but NGC 1510 (forwhich we are not observing the same region within the galaxy) thedispersion is always inferior to 0.06 dex, that is, lower than thetypical ±0.1 dex uncertainty associated with empirical calibrations.

Top middle panel of Fig. 6 compares the oxygen abundancedifferences with the total absolute flux of the Hα line. Again, thescatter is 0.06 dex and no dependence in the flux calibration appears.Indeed, when we compare the absolute Hα flux computed for eachregion using different plates (bottom middle panel of Fig. 6), wefind that the flux difference is better than 15 per cent in almost all thecases, being the average scatter 7.7 per cent. Note these numbers donot include the uncertainty in the flux calibration, which we alreadyestimated to be ∼5 per cent. Finally, bottom panel of Fig. 6 showsthat the agreement between the reddening coefficients determinedfor each region that has been observed twice is excellent, with anaverage scatter of 0.026 dex without considering uncertainties.

In summary, Fig. 6 also shows that the reduction process, absoluteflux calibration, and emission-line fluxes estimation has been doneconsistently in both plates. Interestingly, the fact that the oxygenabundances derived for NGC 1510 using SEL techniques differ thatmuch exemplifies how critical the position on the fibre within thegalaxy can be in some cases, and hence detailed studies of suchobjects really need IFS observations, such those provided by theCALIFA (Sanchez et al. 2012; Husemann et al. 2013) or SAMI(Croom et al. 2012; Bryant et al. 2015) surveys. We will furtherinvestigate the reason of this mismatch below when computing theoxygen abundances following the Te method.

3.6.2 Using the Kobulnicky & Kewley (2004) calibration

To compare the results provided by the Kobulnicky & Kewley(2004) method with the other empirical techniques, we have to keepin mind that SEL methods based on photoionization models (suchas those presented in Kobulnicky & Kewley 2004) systematicallyoverestimate by 0.2–0.4 dex or more the oxygen abundances derivedusing empirical calibrations which are based on direct estimationsof the electron temperature (such as those provided by Pilyugin2001a,b; Pettini & Pagel 2004).2 As a first approach, we adopt

2 However, we note that detailed tailor-made photoionization models ofindividual galaxies (e.g. Perez-Montero et al. 2010) and some recent state-

Figure 6. Comparison of properties of the same region derived using dif-ferent plate configurations. Top panel compares the oxygen abundance dif-ferences derived following the N2 (blue stars), O3N2 (red diamonds), andR23, P (Pilyugin 2001a,b, purple triangles) calibrations. Dotted horizontallines indicate a difference of ±0.06 dex. Top middle panel plots the oxygenabundance differences with the absolute Hα flux in Plate 1. Bottom middlepanel compares the absolute Hα flux computed for each region using dif-ferent plates. Bottom panel plots the reddening coefficient differences withthe absolute Hα flux in Plate 1. In all panels, the position of NGC 1510is indicated with black squares, however in this case we are not observingexactly the same region of the galaxy, as it is discussed in the text.

the conversion between abundance scales for the high-metallicitybranch provided by Lara-Lopez et al. (2013b),

KK04uTe,LL = KK04u − 0.1026

1.0211, (13)

where KK04u refers to the oxygen abundance – in units of12+log(O/H)– given by the high-branch calibration provided byKobulnicky & Kewley (2004) and KK04uTe,LL is this oxygen abun-dance shifted to the absolute abundance scale given by the Te

method. The top panel of Fig. 7 compares the KK04uTe,LL resultswith those provided by the Pilyugin (2001a,b) calibration. As we

of-the-art photoionization models (e.g. Dors et al. 2011; Perez-Montero2014) matches well the oxygen abundances derived using the Te method.

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Figure 7. (Top panel) Comparison of the oxygen abundances determinedfollowing the Kobulnicky & Kewley (2004) method (scaled to the Te methodusing Lara-Lopez et al. 2013b) with the results provided by the Pilyugin(2001a,b) calibration using the R23 and P parameters. (Bottom panel) Com-parison of the adopted oxygen abundances with the results provided bythe high-metallicity Kobulnicky & Kewley (2004) calibration without ap-plying any shift. A dot–dashed green line indicates a linear fit to the data(equation 14). In both panels, dotted horizontal lines indicate a difference of±0.1 dex, while the red dashed vertical line at 12+log(O/H) = 8.3 indicatesthe lower limit to the high-metallicity branch. Blue stars indicate regions ob-served using Plate 1, whilst purple diamonds correspond to regions observedusing Plate 2.

see, even after performing the shift of the Kobulnicky & Kewley(2004) values to the Te scale using equation (13), a significant offsetof 0.1–0.2 dex between both data sets still remains. Equation (13)is then not valid for our objects, which are physically very differ-ent to those Lara-Lopez et al. (2013b) used to get their calibration.Therefore, we again suggest caution when blindly applying an SELtechnique to objects with different physical properties to those usedto derive the empirical calibration (see Stasinska 2010, for an ex-tended discussion).

We then use our data to try to get an absolute scale conversionwhich is valid for this kind of objects. The bottom panel of Fig. 7compares the adopted oxygen abundance (x-axis) with the resultprovided by the Kobulnicky & Kewley (2004) method without ap-plying any shift (y-axis). A linear fit to the data considering onlythose regions with 12+log(O/H) ≥ 8.3 (we only consider data inthe high-metallicity branch) provides

KK04uTe = KK04u − 1.727

0.8475, (14)

being the correlation coefficient r = 0.9773 and the dispersiond = 0.021 dex. For 12+log(O/H) = 8.8, the offset is provided byequation (14) is 0.45 dex. We now use equation (14) to convert theresults of the Kobulnicky & Kewley (2004) calibration into the Te

scale. As expected, now the agreement is excellent, with a scatterof 0.025 dex.

We remark that we do not use this calibration for getting the finaloxygen abundance for each region, but just to justify that differentSEL techniques are providing quantitatively the same results in ouranalysis. That is, neglecting the problem of the absolute oxygenabundance scale (which following equation (14) provides an offsetof 0.46 dex for KK04uTe = 8.30 but 0.39 dex for KK04uTe = 8.75),

Figure 8. Comparison of the N2O2 value with the adopted oxygen metallic-ity for our data. The vertical red dotted line at 12+log(O/H) = 8.3 indicatesthe low limit to the high-metallicity branch. The green dash–dotted line isthe quadratic relationship derived by Kewley & Dopita (2002), manuallyscaled −0.36 dex to match with the data in our high-metallicity end. Thered dashed line is a quadratic fit to our high-metallicity data. Blue starsindicate regions observed using Plate 1, whilst purple diamonds correspondto regions observed using Plate 2.

the results provided by the Kobulnicky & Kewley (2004) methodperfectly matches the oxygen abundances adopted using the Pettini& Pagel (2004) and Pilyugin (2001a,b) techniques.

3.6.3 Using the N2O2 parameter as a high-metallicity oxygenabundance estimator

To further check the goodness of our data, we use the N2O2 param-eter, which is defined as,

N2O2 = log

(I ([N II])λ6584

I ([O II])λλ3726, 3729

). (15)

Kewley & Dopita (2002) recommended to use the N2O2 ratio as anestimator of the oxygen abundance in the high-metallicity regime, asits behaviour is almost linear with the oxygen abundance. Fig. 8 plotsthe N2O2 value derived for each region (y-axis) with the adoptedoxygen metallicity (x-axis). A clear, almost linear relation appears.The green dash–dotted line is the quadratic relationship derivedby Kewley & Dopita (2002), manually scaled by −0.36 dex tomatch with the data in our high-metallicity end. Even though, thisrelationship seems not to be suitable for our objects in the range 8.3≤ 12+log(O/H) ≤ 8.5. We perform a quadratic fit to the data in thehigh-metallicity branch, yielding

x = 8.7457 + 0.3570N2O2 − 0.04229(N2O2)2, (16)

being x = 12+log(O/H). The correlation coefficient of this fit isr = 0.977 59 and the dispersion d = 0.026. We again note thatthis relationship has not been used to derive the adopted oxygenabundance of the UV-bright regions within the NGC 1512/1510system but just to quantify the validity of our analysis.

3.7 The N/O ratio

We computed the nitrogen-to-oxygen (N/O) ratio for those regionswhere both the [O II] λ3727 and [N II] λ6583 lines are observed.

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Figure 9. (Top panel) Comparison between the oxygen abundance (x-axis)and the N/O ratio (y-axis) for the knots analysed here. (Bottom panel) TheN/O ratio (y-axis) is plotted against the O2N2 ratio (x-axis). The dashed redline is a fit to the data. Blue stars indicate regions observed using Plate 1,whilst purple diamonds correspond to regions observed using Plate 2.

We assumed the relationship provided by Garnett (1992) betweenTe[O III] and Te[O II] and used the emission-line ratios to computethe electron temperatures which best reproduce the adopted oxygenabundances following the direct method using the IRAF nebularpackage (Shaw & Dufour 1995). We assumed ne = 100 cm−3 andused Te[O II] and the [N II] λ6583/Hβ ratio to obtain the N+/H+

ionic abundance and derive the N/O = (N+/H+)/(O+/H+) ratio(Peimbert & Costero 1969). This ratio does not strongly depend onthe electron temperature, as both O+ and N+ ions are assumed to bein the low-ionization region. Hence, the derived uncertainties of theN/O ratio are low (±0.03 dex in many cases) and only consider theerrors in the line-flux ratios. The N/O ratio adopted for each regionis listed in last column of Table 3.

The bottom panel of Fig. 9 compares the derived N/O ratio withthe value of the N2O2 parameter. As we should expect followingFig. 8, we observe a very tight relation between both values. Alinear fit to the data provides

y = −0.6071 + 0.6721x, (17)

with r = 0.9930, where x = N2O2 and y = N/O.

3.8 Chemical abundances using the direct method

There are four regions for which we detect the auroral [O III] λ4363emission line: the two positions of the BCDG NGC 1510, region17 of zone 4, and region 27 of zone 2. For these four cases, wecomputed the oxygen abundances of the ionized gas following theTe method, taking into account all lines detected in each spectrum.Table 4 lists the dereddened line intensity ratios with respect toI(Hβ) = 100, the equivalent widths of the brightest H I Balmerlines, and the c(Hβ) and Wabs iteratively derived using all availableH I Balmer lines and the method explained in Section 3.5.

The electron density of the ionized gas, ne, was computed viathe [S II] λλ6716,6731 doublet by making use of the five-level pro-gram for the analysis of emission-line nebulae included in the IRAF

nebular package (Shaw & Dufour 1995). The four regions werefound in the low-density limit, ne < 100 cm−3, and hence we adoptne = 100 cm−3.

We then infer the Te([O III]) from the [O III](λ4959+λ5507)/λ4363 ratio by making use of the IRAF neb-ular package. As we assumed a two-zone approximation to definethe temperature structure of the nebula, we used Te([O III]) asrepresentative of high-ionization potential ions. The electrontemperature assumed for the low-ionization potential ions wasderived from the linear relation between Te([O III]) and Te([O II])provided by Garnett (1992).

We followed the very same prescriptions and ionization correc-tion factors, icf, indicated by Lopez-Sanchez & Esteban (2009) tocompute the ionic abundances, the total O, N, and Ne abundances,and the N/O and Ne/O ratios, for the ionized gas within these fourregions. We cannot derive the total S and Ar abundances becausewe do not observe any emission line associated with S++ or Ar++.The IRAF nebular package (Shaw & Dufour 1995) was used to com-pute the ionic abundances from the intensity of collisionally ex-cited lines, assuming Te[O III] for computing the ionic abundancesof O++, Ne++, and Ar+3 and the Te[O II] value derived from theGarnett (1992) relation for O+, N+, and S+. We note that we usedthe atomic data set for O+ and S+ ions recommended by Garcıa-Rojas et al. (2005). We considered the standard icf of Peimbert &Costero (1969) to derive the total N and Ne abundances.

The results for the chemical abundances derived using the Te

method are compiled in Table 5. This table also lists the electrontemperatures derived for each region following the direct method.

Table 5 also includes the oxygen abundance and the N/O ra-tio adopted for each region following the SEL methods. Interest-ingly, in all cases except NGC 1510 W the values provided by boththe Te and SEL methods are almost identical. Following the Te

method, both regions of NGC 1510W have the same metallicity,12+log(O/H) ∼ 8.25, however SEL methods are overestimatingthis value by ∼0.1 dex. This is very likely an effect of the ionizationdegree of the gas, i.e. the fibre position used for getting the data ofNGC 1510W is located almost entirely in the low-ionization regionwithin the BCDG. Hence, some caution is needed when computingoxygen abundances in these kinds of objects using SEL methodseven if IFS data are available (see extended discussion in Lopez-Sanchez et al. 2011).

KLS09 used the emission-line intensity data of NGC 1510 pro-vided by Storchi-Bergmann, Kinney & Challis (1995), which in-cluded a detection of the auroral [O III] λ4363 line, to derive12+log(O/H) = 7.95 and log(N/O) =−1.2 following the Te method.KLS09 claimed that the high N/O ratio in NGC 1510 may be re-lated to the presence of Wolf–Rayet (WR) stars in this BCDG.Indeed, our spectrum shows a detection of the nebular He II λ4686line (see Table 4) which seems to be associated with WR stars (e.g.Lopez-Sanchez & Esteban 2010a, and references therein). However,our deeper data show that the oxygen abundance, 12+log(O/H) =8.24 ± 0.06, and the N/O ratio, log(N/O) = −1.34 ± 0.05, com-puted for NGC 1510 are perfectly compatible to those usually foundin BCDGs.

4 PRO P E RT I E S D E R I V E D F RO M U V A N DR A D I O DATA

KLS09 provided a unique data base of UV and 21 cm properties forthe regions of the NGC 1512/1510 system. We here compile andupdate those values to later perform our multiwavelength analysis.We note that we used the same aperture to derive all fluxes (UV, Hα,and H I) in each region considering maps having the same resolutionand point spread function (PSF) matched.

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Table 4. Dereddened line intensity ratios with respect to I(Hβ) = 100 for the regions for which we have a direct estimation of theelectron temperature. We also compile the equivalent widths of the brightest H I Balmer lines, the reddening coefficient, c(Hβ), andthe equivalent widths of the absorption in the hydrogen lines, Wabs, used to correct the spectra for reddening, and the equivalentwidths of the emission H I Balmer lines.

Line f(λ) NGC 1510 NGC 1510 W Zone 4 Region 17 Zone 2 Region 27

[O II] 3727 0.322 291 ± 19 452 ± 33 210 ± 17 251 ± 12H11 3771 0.313 – – 6.5 ± 1.6 –H10 3798 0.307 – – 9.3 ± 2.2 –H9 3835 0.299 – – 10.5 ± 2.3 –[Ne III] 3869 0.291 44.2 ± 4.3 38.2 ± 7.6 62.5 ± 4.6 62 ± 18H8 3889 0.286 21.7 ± 2.9 – 26.6 ± 2.6 –[Ne III] 3969+H7 0.267 20.8 ± 2.4 – 34.6 ± 3.5 39 ± 12He I 4026 0.251 – – 2.85 ± 0.54 –[S II] 4069 0.239 3.2 ± 1.0 – – –Hδ 4101 0.230 26.1 ± 2.5 26.5 ± 4.2 26.1 ± 2.4 26 ± 8Hγ 4340 0.157 45.4 ± 2.8 45.9 ± 4.0 48.1 ± 3.0 48 ± 12[O III] 4363 0.150 4.64 ± 0.80 2.43 ± 0.54 5.97 ± 0.78 4.26 ± 0.64He I 4471 0.116 4.70 ± 0.80 – 3.68 ± 0.76 –[Fe III] 4658 0.059 2.50 ± 0.38 – – –He II 4686 0.049 1.70 ± 0.54 – – –[Ar IV] 4711 0.043 0.49: – – –[Ar IV] 4740 0.034 0.56: – – –Hβ 4861 0.000 100.0 ± 3.3 100.0 ± 4.3 100.0 ± 3.9 100 ± 6[O III] 4959 −0.025 142.8 ± 6.7 84.1 ± 6.0 165 ± 9 156 ± 10[O III] 5007 −0.037 452.6 ± 21 243 ± 14 494 ± 25 462 ± 20[N II] 6548 −0.295 7.98 ± 0.51 12.7 ± 1.2 4.25 ± 0.47 8.4 ± 0.9Hα 6563 −0.297 286 ± 9 288 ± 5 284 ± 14 286 ± 12[N II] 6583 −0.300 23.7 ± 1.3 37.6 ± 2.6 13.5 ± 0.9 23.3 ± 2.0He I 6678 −0.312 2.34 ± 22 1.76 ± 0.39 2.18 ± 0.28 –[S II] 6716 −0.318 12.71 ± 0.74 20.3 ± 2.2 8.89 ± 0.62 21.5 ± 1.7[S II] 6731 −0.319 9.45 ± 0.57 21.1 ± 1.6 6.45 ± 0.49 15.2 ± 1.4

−W(Hα) [Å] 188 ± 9 78 ± 3 774 ± 38 756 ± 38−W(Hβ) [Å] 32.2 ± 0.2 17.5 ± 0.2 216 ± 2 114 ± 3−W(Hγ ) [Å] 10.1 ± 0.1 6.6 ± 0.2 66 ± 2 38 ± 3−W(Hδ) [Å] 4.6 ± 0.1 3.0 ± 0.3 28.7 ± 1.5 10.3 ± 2.1

c(Hβ) 0.36 ± 0.04 0.20 ± 0.3 0.01 ± 0.02 0.03 ± 0.02Wabs [Å] 0.3 ± 0.1 0.3 ± 0.1 0.5 ± 0.1 0.5 ± 0.2

4.1 UV luminosities and colours

As explained in section 4.2 of KLS09, the GALEX FUV and NUVimages were used to derive mFUV and mNUV magnitudes and theFUV − NUV (FN) colour. They also estimated the FUV flux density,fFUV (in units of erg s−1 cm−2 Å−1), that was used to derive the starformation rate (SFR) of each region via the calibration providedby Salim et al. (2007). We include these results for each analysedregion in Table 6. Column 7 compiles the area (in kpc2) of the region,Column 8 shows the FN colour, Column 9 indicates the fFUV, andColumn 10 lists the SFR derived from the FUV flux density.

We note our FUV-based SFR are slightly different from thoseestimated by KLS09, as we now use the reddening coefficientestimated for each region (or that given by equation 6 if Hβ

was not observed) to correct the FUV flux density for extinction,assuming E(B − V) = 0.692c(Hβ)mag, while KLS09 used theGalactic value provided by Schlegel, Finkbeiner & Davis (1998),E(B − V) = 0.011 mag.

We also note that the effect of the so-called IRX–β relation (i.e.the relation between the UV slope or the FN colour with the infraredexcess in either entire galaxies or individual regions within galaxies;see Boquien et al. 2009, and references within) seems to be smallfor the analysed regions in the NGC 1512/1510 system. We estimatethat this represents a variation of only ∼0.05 mag in the FN colour

for ages older than ∼10 Myr. Therefore, we assume that the effectof dust in the FN colour is well reproduced by a screen of dustshowing a Milky Way-like extinction law.

4.2 H I and gas masses

The H I masses of each UV-bright region were derived by KLS09 us-ing the low-resolution Australia Telescope Compact Array (ATCA)data considering the same region determined by their UV emission.Considering that the vast majority of the hydrogen in low-mass ob-jects is atomic (e.g. Leroy et al. 2005), the total gas mass can bedetermined assuming Mgas = 1.32MH I, where the factor 1.32 cor-rects the H I mass for the presence of helium. Column 11 of Table 6compiles the total gas mass derived for each region. As alreadynoted by KLS09, the gas masses derived in the central regions ofNGC 1512 are very probably underestimated, as we lack of CO ob-servations to account for the molecular gas (H2) component, whichis expected to be important in these knots (see fig. 20 and section4.4 in KLS09).

4.3 Stellar and baryonic masses

We have used the UV data to get a tentative estimation of the stel-lar mass, Mstars, within each UV-bright complex. For this, we have

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Table 5. Physical conditions and chemical abundances of the ionized gas of the regions for whichwe have a direct estimation of the electron temperature The oxygen abundance and the N/O ratioadopted for each region following the SEL methods are also included for comparison.

NGC 1510 NGC 1510 W Zone 4 Region 17 Zone 2 Region 27

Te[O III] [K] 11700 ± 600 11500 ± 900 12300 ± 600 11300 ± 700Te[O II] [K] 11200 ± 500 11050 ± 600 11600 ± 400 10900 ± 500ne [cm−3] <100 <100 <100 <100

log(O++/O+) 0.09 ± 0.11 −0.35 ± 0.14 0.28 ± 0.09 0.18 ± 0.1112+log(O+/H+) 7.88 ± 0.10 8.09 ± 0.11 7.66 ± 0.07 7.86 ± 0.0912+log(O++/H+) 7.97 ± 0.06 7.74 ± 0.09 7.94 ± 0.05 8.04 ± 0.0712+log(O/H), Te 8.23 ± 0.08 8.25 ± 0.11 8.12 ± 0.06 8.26 ± 0.0812+log(O/H), emp 8.23 ± 0.06 8.36 ± 0.07 8.12 ± 0.05 8.24 ± 0.07

12+log(N+/H+) 6.54 ± 0.05 6.76 ± 0.06 6.24 ± 0.04 6.59 ± 0.0612+log(N/H) 6.89 ± 0.08 6.92 ± 0.08 6.70 ± 0.07 6.99 ± 0.09icf(N) 2.22 ± 0.34 1.45 ± 0.17 2.90 ± 0.41 2.50 ± 0.44

log(N/O), Te −1.34 ± 0.06 −1.33 ± 0.09 −1.42 ± 0.07 −1.27 ± 0.07log(N/O), emp −1.34 ± 0.03 1.39 ± 0.03 −1.44 ± 0.03 −1.29 ± 0.03

12+log(S+/H+) 5.59 ± 0.05 5.97 ± 0.06 5.39 ± 0.04 5.84 ± 0.06

12+log(Ne++/H+) 7.43 ± 0.11 7.39 ± 0.18 7.49 ± 0.08 7.61 ± 0.08icf(Ne) 1.82 ± 0.61 3.2 ± 1.6 1.53 ± 0.40 1.67 ± 0.6412+log(Ne/H) 7.69 ± 0.16 7.9 ± 0.2 7.68 ± 0.13 7.83 ± 0.16log(Ne/O) −0.54 ± 0.11 −0.35 ± 0.19 −0.45 ± 0.09 −0.43 ± 0.09

12+log(Ar+3/H+) 4.8: – – –

neglected the contribution of the current star formation event as ob-served from the ionized gas. We then assumed that the FN coloursand the FUV luminosities are coming from a dominant stellar pop-ulation created in a past star formation event. The age of this starformation event can be determined using the FN colour, as KLS09already did (see section 4.2 in their paper). We then used theoreticalevolutionary synthesis models provided by the STARBURST99 code(Leitherer et al. 1999; Vazquez & Leitherer 2005) to compute thestellar mass that matches the observed monochromatic FUV lumi-nosity and FN colour of each region. Although the FUV luminositydepends poorly on the metallicity (e.g. see discussion in section 3of Bianchi et al. 2005), we have also considered it using the oxygenabundance determined for each region. The adopted stellar massis listed in Column 12 of Table 6. As these are just tentative esti-mations, we do not provide errors for the stellar mass of the UVregions or for any other property derived using this value.

The total baryonic masses are estimated adding the gas and thestellar masses and are listed in Column 12 of Table 6. Similarly, thegas-to-star mass ratio, Mgas/Mstars, has been also derived for eachregion, and it is compiled in Column 13 of Table 6.

5 D ISCUSSION

5.1 An oxygen abundance map of the NGC 1512/150 system

Our study allows us to map the oxygen abundance of the NGC1512/1510 system. For this, we use the data of the 92 unique UV-bright regions (see Table 6) for which we compute the oxygenabundance to create the metallicity map we show in Fig. 10. Wehave assigned different symbols to each of the four zones identifiedin the system (see Fig. 1). We also include the positions of the BCDGNGC 1510 (star) and the TDG 1 (pentagon), but this actually liesfar from the system, at 78 kpc (31.4 arcmin) from the centre ofNGC 1512. We analyse this map by zone.

Zone 1 = Ring. This area corresponds to the central regionsof NGC 1512. As expected, the star-forming knots within thisarea show the highest metallicities of the system, ranging between8.6 ≤ 12+log(O/H) ≤ 8.8. The average value is 12+log(O/H) =8.71. The northern regions of this ring show slightly higher oxygenabundances than the southern regions.

Zone 2 = Arm 1. This is the long spiral/tidally induced arm atthe east of the system. Although it seems it starts at the SW ofthe system (bridge between NGC 1512 and NGC 1510) with rela-tively high metallicities – 8.5 ≤ 12+log(O/H) ≤ 8.6 – it later hasan almost flat gradient, with relatively low oxygen abundances,8.4 ≤ 12+log(O/H) ≤ 8.3. However, some bright UV regionsare showing ever lower values, reaching 12+log(O/H) ∼ 8.2 insome knots. The average value is 12+log(O/H) = 8.39, although ifthe innermost five regions are not considered the average value ofthe metallicity of Arm 1 is 12+log(O/H) = 8.35.

Zone 3 = internal arm. This area corresponds to the inner partof an original, long arm (Arm 2) which, starting at the NE of thecentral ring, has been now partially destroyed because of the inter-action between NGC 1512 and NGC 1510 (KLS09). Because of itscloseness to the centre of the system, we should expect intermediateto high metallicities in this area, and indeed that is what we find.Metallicities range between 8.65 and 8.50, the average value be-ing 12+log(O/H) = 8.55. Interestingly, the two regions found at thewestern end of this area (3_24 and 3_25, see Fig. 1) have lower abun-dances, 12+log(O/H) ∼ 8.40, than expected, as being located veryclose to the inner areas of the system. Indeed, region 4_01, whichis found at just 2 kpc N from region 3_25, has 12+log(O/H) =8.71. This fact is evidencing that the interaction between the twomain galaxies has largely disturbed the distribution and propertiesof the star-forming regions of the system.

Zone 4 = external debris. this corresponds to the external part ofthe old Arm 2, and hence in the past it should have been connectedto Zone 3. Now the distribution of star-forming regions is not asuniform as it was probably before, which we suspect was similar

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Tabl

e6.

Prop

ertie

sof

the

UV

-ric

hre

gion

sw

ithin

the

NG

C15

12/1

510

syst

em.T

hefu

llta

ble

isav

aila

ble

onlin

e.

Zon

eR

egR

25c(

Hβ

)12

+log

(O/H

)lo

g(N

/O)

Are

aF

Na

f FU

Vlo

g(SF

R)

log(

Mga

s)lo

g(M

star

s)lo

g(M

bar)

Mga

sM

star

slo

g(SF

R/A

)lo

g(M

gas/

A)

log(

Mst

ars/

A)

log(

Mba

r/A

)

(1)

(dex

)(d

ex)

(dex

)(k

pc2)

(mag

)(2

)(

M�

yr−1

)(

M�)

(M

�)(

M�)

(M�

yr−1

kpc−2

)(M

�kp

c−2)

(M�

kpc−2

)(M

�kp

c−2)

0N

1510

1.49

0.39

±0.

018.

23±

0.05

−1.3

4±

0.03

3.74

80.

23±

0.03

3580

±11

0−0

.92

±0.

037.

44±

0.03

7.41

7.73

1.08

−1.4

9±

0.03

6.87

±0.

036.

837.

150

TD

G1

6.27

0.06

±0.

018.

24:

–2.

520

0.43

±0.

1416

.1±

2.5

−3.8

2±

0.12

7.43

±0.

037.

597.

820.

684

−4.2

2±

0.12

7.03

±0.

037.

197.

42

11

0.51

–8.

67±

0.13

–0.

258

0.20

±0.

1229

.9±

3.4

−3.5

8±

0.09

6.06

±0.

115.

136.

118.

67−2

.99

±0.

096.

65±

0.11

5.71

6.70

12

0.41

0.25

±0.

018.

72±

0.07

−0.7

2±

0.03

0.45

00.

26±

0.07

94.6

±6.

1−2

.53

±0.

056.

36±

0.09

5.97

6.50

2.46

−2.1

8±

0.05

6.70

±0.

096.

316.

851

30.

360.

44±

0.08

8.77

±0.

15–

0.43

80.

21±

0.10

47.0

±4.

2−2

.42

±0.

076.

22±

0.10

5.44

6.29

6.07

−2.0

7±

0.07

6.58

±0.

105.

806.

651

40.

310.

39±

0.05

8.71

±0.

08−0

.64

±0.

060.

481

0.21

±0.

0883

.1±

5.7

−2.2

8±

0.06

6.25

±0.

105.

716.

363.

41−1

.96

±0.

066.

57±

0.10

6.03

6.68

15

0.28

0.24

±0.

018.

74±

0.07

−0.7

8±

0.03

0.69

00.

21±

0.05

187.

3±

9.3

−2.2

6±

0.04

6.37

±0.

096.

086.

551.

95−2

.10

±0.

046.

53±

0.09

6.24

6.71

16

0.28

–8.

69±

0.15

–0.

597

0.33

±0.

0612

6.3

±7.

2−2

.96

±0.

056.

39±

0.09

6.35

6.67

1.09

−2.7

3±

0.05

6.61

±0.

096.

576.

891

70.

320.

35±

0.02

8.73

±0.

08−0

.78

±0.

030.

797

0.31

±0.

0332

0±

13−1

.78

±0.

046.

59±

0.06

6.68

6.94

0.81

−1.6

8±

0.04

6.69

±0.

066.

787.

04

Not

es.(

1)R

25=

11.7

6kp

c(2

)In

units

of10

−17

erg

s−1

cm−2

Å−1

.a

FN

=FU

V−

NU

V.

to what we now see in Arm 1. Furthermore, we note it is not easyto associate some knots to be in Arm 1 or Arm 2. Hence, regions1_01 to 1_05 may actually be in the external debris zone, whilesome of the knots 4_19, 4_20, 4_26 and 4_27 may be originallyassociated with Arm 1. Knot 4_1 may actually belong to the Ring.KLS09 already discussed that the H I morphology and kinematicsof this zone is also much more disturbed than that observed in Arm1. We now see the same behaviour in the metallicity distribution ofits star-forming regions. Although the average value in this zone is12+log(O/H) = 8.44, we find regions from 12+log(O/H) = 8.71to 12+log(O/H) = 8.12. Actually, three of the four most metal-poor star-forming regions within the NGC 1512/1510 system, 4_17,4_32, and 4_38, are located here. Particularly interesting is the knot4_17, for which we derive 12+log(O/H) = 8.12 ± 0.05 consistentlyusing both the direct and the SEL methods.

It is very interesting to note the chemical differences betweenArm 1 and the external debris of Arm 2. We should expect that,originally, these two long arms shared similar properties, as we mayassume that both structures were created from the same material atthe outskirts of NGC 1512. We will further discuss this issue below.

5.2 Radial metallicity gradient

The left-hand panel of Fig. 11 compares the oxygen abundanceof each of the analysed UV-rich regions with their distance to thecentre of NGC 1512 in units of R25 (note that we are using projectedgalactocentric distances in our analysis). Because of the complexityof the system, it is not easy to see just a radial decrease of metallicityfrom the centre to the outskirts. Considering all knots in the ringand in the internal arm and all regions of the external debris with12+log(O/H) >8.20, a linear fit yields a metallicity gradient of

12 + log(O/H) = 8.703 − 0.158R25, (18)

with r = 0.7862. This fit is plotted in the left-hand panel of Fig. 11with a red dashed line. Note that the zero-point of this fit matchesthe average oxygen abundance of the central ring. However, a linearfit to all UV-bright regions along Arm 1, neglecting knots 1–5 whichhave metallicities at least 0.2 dex higher than the rest, and includingthe TDG 1, provides a metallicity gradient with a very flat slope,

12 + log(O/H) = 8.400 − 0.023R25, (19)

with r = 0.3294, which is plotted with a blue dot–dashed line inthe figure. As was already said, note that although an almost flatmetallicity gradient along Arm 1 exists, the dispersion is ∼0.2 dex.

Using the average metallicity gradient provided by equation (18),we prepared a residual metallicity map (left-hand panel of Fig. 12)which takes into account how much this value differs with respectto the adopted oxygen abundance for each region. Besides NGC1510 and the three low-metallicity knots we already identified inthe external debris, this map identifies two more zones where knotshave oxygen abundances which are systematically lower than theirsurroundings by about 0.2–0.3 dex. These are the most southernknots of the internal arm and some of the regions within Arm 1located at the N of the system.

5.3 A metallicity gradient following the spiral H I arms

Instead of using just a radial gradient, we considered a radial plusazimuthal gradient by defining two spirals which follow the H I armsand UV-rich complexes observed in the NGC 1512/1510 system (seeFig. 10). These spirals were constructed by eye using the equation

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Figure 10. Map of the oxygen abundance in the NGC 1512/1510 system, as it is provided via the analysis of the emission lines detected in the UV-richregions. The four zones identified in the system (see Fig. 1) are shown using different symbols: Zone 1 = Ring (Diamonds), Zone 2 = Arm 1 (circles), Zone3 = internal arm (triangles), Zone 4 = external debris (squares). The position of the BCDG NGC 1510 is shown with a star, while the TDG 1 (which actuallylies far from the system, at 78 kpc from the centre of NGC 1512) is plotted using a pentagon. The scale included in the top-left corner indicates the length of5 arcmin, which corresponds to 12.44 kpc at the distance of NGC 1512. The value for R25 in NGC 1512 is 11.76 kpc (4.26 arcmin). The two H I tidal/spiralarms are sketched following a blue (Arm 1) or green (Arm 2) dotted line.

Figure 11. (Left-hand panel) Oxygen abundance versus R25 for the UV-rich regions analysed in the NGC 1512/1510 system. The red dashed line in theleft-hand panel indicates the linear fit given by equation (18), which considers al knots in ring and in the internal arm and all regions of the external debris with12+log(O/H) > 8.20. The blue dot–dashed line in the left-hand panel is a linear fit to all the regions in Arm 1, neglecting knots 1–5 which have metallicities atleast 0.2 dex higher than the rest, and the TDG 1 (which is not shown here). The data points used for getting this fit are marked with a blue circle. (Right-handpanel) N/O ratio versus R25 for the UV-rich regions analysed in the NGC 1512/1510 system. The red dashed line is a linear fit to the same regions used to getthe fit provided in equation (18). In both diagrams, colour indicates the Mgas/Mstars ratio of each region. Symbols are the same than those in Fig. 10.

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Figure 12. (Left-hand panel) Oxygen abundance residual map constructed using the average metallicity gradient given by equation (18). (Right-hand panel)N/O map of the system. In both panels, symbols are the same than those in Fig. 10.

Figure 13. Oxygen abundances of the UV-rich regions following the spiral Arm 1 (bottom panel) and the spiral Arm 2 (top panel). Blue and red dotted linesmark the lower/higher oxygen abundances found in the outskirts of Arm 1. The dashed black line shows a schematic model combining a linear gradient (forR < 10 kpc) with slope α = −0.045 dex kpc−1and zero-point z0 = 8.80 dex and a flat gradient (for R > 10 kpc) with 12+log(O/H) = 8.35 constructed toguide the eye. Colour indicates the Mgas/Mstars ratio of each region. Symbols are the same than those in Fig. 10.

θ = θ0 + 1/bln (R/a), R being the distance to the centre of thespiral, using the observed knots as a guide. For each region, wesearched for the minimum distance between its actual position andthe spiral given by Arm 1 (all regions in Zone 2, which we alreadyidentified with this arm) or Arm 2 (for regions belonging to Zones3 and 4). We are not including the regions in the central ring (Zone1), NGC 1510, and TDG 1 in this analysis.

The left-hand panel of Fig. 13 plots the oxygen abundance ofthe star-forming regions along these two spiral arms. Following thisfigure, it is quite evident that the two spiral arms have experienceddifferent star-formation histories. Assuming that the regions withR �10 kpc in Arm 1 are showing the original flat oxygen abun-dance distribution (i.e. before the two main galaxies interacted)with 12+log(O/H) ∼ 8.35, and assuming that the centre of NGC1512 has 12+log(O/H) ∼ 8.80, the slope for a linear gradient forR < 10 kpc is α = −0.045 dex kpc−1. Both the slope and the con-stant oxygen value in the outer regions are similar to those reportedby Bresolin et al. (2012) for NGC 1512.

This schematic model for the oxygen abundance distribution isplotted with black dashed lines in the figure. As we discuss below,we note that the original metallicities in the outskirts of NGC 1512

could be even lower than the 12+log(O/H) ∼ 8.20 we still see insome knots, as is observed in the outskirts of spiral galaxies suchas NGC 300 (Bresolin et al. 2009b), M 83 (Bresolin et al. 2009a),NGC 4625 (Goddard et al. 2011), and NGC 3621 (Bresolin et al.2012).

In any case, almost all regions associated with the original Arm 2now have systematically higher abundances, 12+log(O/H) ∼ 8.4–8.6, than those given by our schematic model (left-hand panel ofFig. 13). The same is happening to the five internal regions withinArm 1. This fact informs that the star formation activity and thechemical enrichment have been much more important in regionsassociated with or close to Arm 2, at the NW of NGC 1512. KLS09already showed that the interaction with NGC 1510 has modifiedthe H I morphology and kinematics of the neutral gas in this zone,and that this is the reason why a much more chaotic distributionof star-forming regions is observed. We now confirm that both theenhancement of the metallicities and the large metallicity dispersionfound within the remaining Arm 2 and nearby regions are also aconsequence of the interaction between NGC 1512 and NGC 1510.

Interestingly, we also note a slight increase of the oxygen abun-dance in the six most external regions of Arm 1. Although this

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Figure 14. Mgas/Mstars map for the UV-rich regions identified in Table 6.Symbols are the same than those in Fig. 10.

increase is only ∼0.1 dex and the metallicities of these regions havethe highest errors in the system, we note that these regions also havea relatively low Mgas/Mstars ratio when compared with nearby knots(see Fig. 14). Has the gas been consumed recently here, used byseveral star formation events, and has the material therefore beenenriched more than in other regions of the Arm 1? Interestingly,the residual H I velocity field presented by KLS09 (see their Fig. 7,right-hand panel) shows deviations up to ∼30 km s−1 around thesecomplexes, which typically have the older ages of the system ac-cording to their FN colours. This suggests that the interaction hasindeed induced the star formation activity within these regions, andconsequently enhanced their oxygen abundances and depleted theirgas reservoir.

Neglecting all regions in Arm 1, it is difficult to define the ra-dial metallicity distribution that was originally present in the sys-tem. NGC 1512 is a perfect example that illustrates the fact thatgalaxy interactions may flatten the metallicity gradients in galaxies(Kewley et al. 2010; Rupke et al. 2010; Werk et al. 2011). NGC 1512nicely shows, at the same time, both the gradient flattening becauseof its interaction with NGC 1510 in one arm (Arm 2) and the as-sumed, original, poorly disturbed radial distribution (gradient+flatin the outskirts) on the other arm (Arm 1).

5.4 The N/O ratio

The N/O ratio provides important clues to the study of the starformation history of galaxies (e.g. Molla et al. 2006). While thenitrogen stellar yield mainly proceeds from low- and intermediate-mass stars, the oxygen yield comes from massive stars. The right-hand panel of Fig. 12 shows the N/O ratio map of the system. Theexternal areas show log (N/O) values between −1.30 and −1.50,while the internal areas have −0.9 ≤ log (N/O) ≤ −0.6. Thesevalues more or less correlate with the adopted oxygen abundancesof their knots (e.g. Izotov et al. 2006; Lopez-Sanchez & Esteban2010b; Moustakas et al. 2010). The right-hand panel of Fig. 11shows the N/O gradient in NGC 1512. Again, it is not clear that aunique N/O gradient exists in the system. A fit to the data of thesame regions used to get equation (18) yields

log(N/O) = −0.62 − 0.32R25, (20)

with r = 0.6401. The two areas we identified using the oxygenabundance residual map also appear evident in the N/O residualmap created using this fit. Regions located at the N of Arm 1 have12+log(O/H) ∼ 8.3 and log (N/O) ∼ −1.4, while regions located inthe external areas of the internal arm have 12+log(O/H) ∼ 8.4 andlog (N/O) ∼ −1.2. In both cases, the N/O ratio agrees with what isexpected from their oxygen abundances. We note that, although thePettini & Pagel (2004) calibrations use the [N II] lines to estimatethe oxygen abundance, these relationships are purely based on ob-servational values that consider only Te measurements for whichO+/H+, O++/H+, and N+/H+ have been computed independently,and hence there is not any ad hoc N/O versus O/H dependenceembedded in the calibrations.

Nitrogen enrichment may be expected in objects which are expe-riencing a strong star formation event as consequence of the accre-tion of low-metallicity H I gas (Koppen & Hensler 2005; Amorın,Perez-Montero & Vılchez 2010; Perez-Montero et al. 2011; Molla& Terlevich 2012; Amorın et al. 2012). In this scenario, the oxygenabundance released by the recently formed massive stars is lowerthan that in the ISM before the starburst, while the nitrogen abun-dance remains the same until it is released by the intermediate andlow-mass stars hundreds of millions of years after the massive starsdie. The net effect is an increase of the N/O ratio, that we do notobserve in these areas.

Hence, these regions with both lower O/H and N/O are actuallyless chemically evolved than their surroundings. In the case ofthe outer regions in the internal arm, they were probably originallylocated at larger distances than they are today, and therefore they stillkeep the relatively low metallicity of these outer regions. Somethingsimilar may have happened to the low-metallicity regions at the N ofArm 1; they are now closer to NGC 1512 than they were in the past.Following the residual H I velocity field of NGC 1512 (KLS09), thisarea shows little deviation from the rotating H I curve, and hencewe may expect less disturbance from the gravitational interactionbetween the two main galaxies. If this suggestion was true, it wouldexplain that other regions along the Arm 1, which show higherdeviations from the rotating H I curve, have experienced more starformation events than those areas, and therefore enhanced theirmetallicities from 12+log(O/H) ∼ 8.20 to 12+log(O/H) ∼ 8.40and from log(N/O) ∼ −1.5 to log(N/O) ∼ −1.3, as it is observednow.

5.5 Dust in the NGC 1512/1510 system

The intrinsic extinction map of the system, once the contribution ofthe extinction of the Milky Way has been subtracted, was computedusing the values of the reddening coefficient derived for each regionand applying

E(B − V )instrinsic = 0.692c(Hβ) − E(B − V )MW, (21)

where E(B − V)MW = 0.011 mag (Schlegel et al. 1998). The E(B− V)instrinsic map is shown in Fig. 15. As this figure shows, theintrinsic extinction seems not to be very important in the system.The highest values, 0.25 < E(B − V)instrinsic < 0.30, are found inNGC 1510 and in the internal regions of NGC 1512. In the externalstar-forming regions, the colour excess is typically 0.05 < E(B −V)instrinsic < 0.10. These data suggest that the dust contribution inthe system is rather small, particularly in the external areas, whichhave extinctions similar to those found in dwarf, low-metallicitygalaxies (Lopez-Sanchez 2010). Bresolin et al. (2012) found thatsome of the brightest regions along Arm 1 show c(Hβ) values upto 0.5 dex, the mean value being c(Hβ) = 0.34 dex. These values

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Figure 15. Intrinsic E(B − V) map in the NGC 1512/1510 system. Symbolsare the same than those in Fig. 10.

are translated into intrinsic colour excesses of 0.29 and 0.19 mag,respectively, which are 0.1–0.2 mag higher than the values found inour analysis. Although this difference may be a consequence of justnot analysing exactly the same regions, we note we have had specialcare deriving the reddening coefficient of our observed knots.

5.6 Kinematics using the Hα emission

To determine the radial velocity of each Hα detection towards UV-bright stellar clusters in the galaxy pair NGC 1512/1510, we fit aGaussian to each Hα emission profile.

The resulting Hα velocity field is shown in left-hand panel ofFig. 16.

As we see, the kinematics resemble that of a rotating disc, as itwas found in the H I analysis (KLS09). These authors also noticedthat the outer disc showed a more disturbed rotating pattern thanthe inner disc. The radial velocity of the centre of NGC 1512 usingthe Hα line is vr = 879 ± 35 km s−1, which matches that measured

using the H I velocity map, vH I = 865.7 km s−1. Note that despite theuncertainty in the absolute optical radial velocity, a relative velocity(i.e. the observed radial velocity of a region with respect to theobserved radial velocity of the centre of NGC 1512) allows us toachieve resolutions of ∼0.1 Å ∼ 5 km s−1.

The right-hand panel of Fig. 16 shows the mean H I velocity fieldas determined by KLS09 using the same velocity range (both opticalheliocentric). The comparison of both kinematics maps allows us toextract some interesting results. First, we note that the kinematics ofthe ionized gas within the ring show slightly blueshifted velocitieswith respect to the H I velocity at the east, while its western knotstend to show slightly redshifted velocities with. This may be afeature of the internal kinematics of the ring.

Secondly, the mean dispersion in velocity for the regions locatedat the NW (corresponding to the internal knots of Arm 1 and theexternal debris of Arm 2), 9.8 km s−1, is slightly larger than thatfor regions located along Arm 1, which is 7.0 km s−1. This resultagrees with the hypothesis that the NW zone of NGC 1512 hasbeen disrupted because of the interaction with NGC 1510. It is alsointeresting to note some kinematic disturbances at the outer regionsof Arm 1. We also note that the H I velocity measured in the TDG1 matches well with its Hα velocity.

Finally, we measure a velocity difference of vr,Hα −vr,H I = 136 km s−1 for knot 3_20, which is located at the ex-ternal regions of the inner arm. We did not detect the [N II] λ6583emission line in this object, but a very broad Hα line is clearlyobserved. The velocity dispersion of the Hα line in knot 3_20 isσ = 3.9 Å, which means that it has a velocity dispersion of 178km s−1. Typical dispersions in the Hα line profiles of the regionsobserved in NGC 1512 are σ = 1.2–1.5 Å (55–70 km s−1). A carefulinspection of its optical spectrum reveals that both Hβ and [O III]λ5007 are barely detected at the same radial velocity than Hα. Theseemission lines also show very high-velocity dispersions, 7.4 Å =456 km s−1and 7.8 Å = 467 km s−1, respectively. Examining theregion where [N II] λ6583 should be, we constrain that the highestoxygen abundance of this knot is 12+log(O/H) � 8.1 using theN2 parameter and the Pettini & Pagel (2004) calibration. That is adifference of at least 0.4 dex with respect to the typical metallicitiesfound in this area. We suggest knot 3_20 to be an independent,

Figure 16. Velocity fields of the galaxy pair NGC 1512/1510 as determined from the Hα emission line (left; this paper) and the H I emission line (right;KLS09) for the targeted UV-bright stellar clusters. Symbols are the same than those in Fig. 10. Note the very different kinematics in knot 3_20 (marked withan open black circle at coordinates RA = 04h04m03.s61 and Dec. = −43◦24′46.′′26) when comparing the H I and Hα maps.

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dwarf galaxy (or, given to its high-velocity dispersion, its remnant),which was accreted into the NGC 1512 system in the past. The H I

emission comes from the NGC 1512 disc while the Hα emission isfrom the dwarf.

5.7 The star formation in NGC 1512/NGC 1510 revisited

The combination of the metallicities and other properties derivedusing our optical spectra with the H I and UV data allow us toexplore with more details the star formation properties and theinterplay between stars and gas in the outer regions of the system.As mention in Section 4, we remind the reader that the Hα, UV, andH I fluxes were obtained using the same aperture for each regionafter matching resolution and PSF. We also note that the definitionof the region was obtained using the UV images, and hence this mayintroduce a bias towards regions of high UV/Hα rations. Howeveras the majority of the UV-bright regions show Hα emission, thiseffect is probably small.

5.7.1 The Hα-based SFR

We have used the extinction-corrected Hα luminosity of each regionto estimate the SFR. However, we have to consider that we areprobably not observing all the emission of the ionized gas, as thefibre size used to get the spectroscopic data is 2 arcsec. Meureret al. (2006) presented flux-calibrated Hα images of a large sampleof H I-rich galaxies, including the galaxy pair NGC 1512/1510. Asdiscussed in that paper, the Hα flux has been corrected for extinction,[N II] contribution, and underlying Hα absorption. Hence, we havecompared the Hα fluxes obtained using the image provided byMeurer et al. (2006) with those derived using our spectra for the 64regions for which we do have measurements from both data sets.Although the dispersion is rather large, the Hα flux estimated usingour spectroscopy is, on average, only ∼3.4, 2.2, and 1.4 times lowerthan that using the Hα images for a Hα flux of 10−15, 10−14, and10−13 erg s−1 cm−2, respectively. As the angular areas of the knotsare typically 40–60 times more extended than the aperture size usedfor getting the spectroscopic data, this informs that the ionized gasis very localized within each UV-rich cluster.

We then scaled (for objects included in both data sets) or cor-rected (using a fit to the data for those objects only observed usingspectroscopy) our Hα flux values for the missing gas emission toderive the total Hα luminosity in each knot. We finally appliedKennicutt (1998) relationship to derive their SFR. Note that theKennicutt (1998) calibration considers an instantaneous burst, ifusing the Calzetti et al. (2007) calibration – which assumes a STAR-BURST99 model (Leitherer et al. 1999; Vazquez & Leitherer 2005) ofconstant SFR for 100 Myr at solar metallicity – the SFR values willonly be two-thirds of those adopted here following the Kennicutt(1998) calibration.

Fig. 17 compares the Hα-based SFR derived using our spectrawith the FUV-based SFR derived using the GALEX image. Althoughthe dispersion is large, we see that, on average, the Hα-based SFRare ∼4 times lower than the FUV-based SFR. This result agrees withthe finding reported by some authors (Boselli et al. 2009; Lee et al.2009; Hunter, Elmegreen & Ludka 2010) in dwarf galaxies withSFR � 0.03 M� yr−1 (log SFR � −1.5): the Hα-flux underpredictthe SFR relative to the FUV luminosity.

However, it is important to remember that while Hα emissiontraces the most massive, ionizing stars (time-scales of ∼10 Myr), theUV emission probes star formation over time-scales of ∼100 Myr,the lifetime of the massive OB stars. Do the differences between Hα-

Figure 17. Comparison between the Hα-based SFR derived using ourspectra (y-axis) with the FUV-based SFR derived using the GALEX image(x-axis). Symbols are the same than those in Fig. 10.

SFR and FUV-SFR therefore somehow indicate the changes in therecent star formation history of the UV-rich clusters? Interestingly,knots 4_17 and 4_32 show an order of magnitude higher Hα-SFRwhen compared to the FUV-SFR. This seems to indicate that, atleast in these cases, the knots are experiencing right now a verystrong star formation event.

5.7.2 Age(s) of the most-recent star-forming event(s)

An estimation of the age of the most-recent star-forming event canbe obtained from the Hα equivalent width, as it decreases with time(e.g. Leitherer & Heckman 1995). The left-hand panel of Fig. 18shows the Hα equivalent width map of the system. In contrast tothe FN map (see fig. 13 in KLS09), we do not observe any kind ofage gradient or pattern in W(Hα). In fact, the Hα equivalent widthappears randomly distributed throughout the system, with valuesranging between log[−W(Hα)] ∼ 2.9 and 0.7.

Following a STARBURST99 (Leitherer et al. 1999; Vazquez &Leitherer 2005) model experiencing a starburst at Z = 0.008, thecorresponding ages for log[−W(Hα)] ∼ 3.5, 2.5, 1.3, and 0.0 are1, 5, 10, and 25 Myr, respectively. Hence, assuming an instanta-neous burst and following their W(Hα), all UV-rich regions ob-served here have experienced their most recent star formation in thelast ∼13 Myr. In some cases, such as knot 4_17, this has happenedjust ∼3 Myr ago. All these ages are much younger than those de-rived using the FN colours. This behaviour was already observedby Boquien et al. (2009) when analysing independent star-formingregions within nearby spiral galaxies. KLS09 used the FN coloursof the UV-rich regions to conclude that their star formation typicallystarted between 100 and 400 Myr ago, being more recent at the east(brightest part of Arm 1) and west (around NGC 1510) of NGC1512 (see their Table 6).

How can we explain the age difference? The right-hand panelof Fig. 18 compares the FN colour with the Hα equivalent widthfor the UV-rich stellar clusters of NGC 1512. As already discussed,both quantities provide an estimation of the age of the star for-mation activity, although their time-scales are different. We shouldexpect that very young bursts – objects with high log[−W(Hα)] –show bluer FN colours than older bursts. This panel includes thepredictions given by two STARBURST99 models (Leitherer et al. 1999;Vazquez & Leitherer 2005), both with Z = 0.008 but one assuming

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Figure 18. (Left) Map of the Hα equivalent widths derived from our optical spectroscopy data. (Right) Comparison between the Hα equivalent width (y-axis)and the FN colour (x-axis) for the regions analysed in this work. The predictions given by two STARBURST99 models (Leitherer et al. 1999; Vazquez & Leitherer2005), both with Z = 0.008 but one assuming an instantaneous burst (purple continuous line) and the other considering a constant star formation (red dottedline) are plotted. Two linear combinations of these models, one considering a 50 per cent continuous plus 50 per cent instantaneous (blue dashed line) and a90 per cent continuous plus 10 per cent instantaneous (green dot–dashed line) are also plotted. The symbol colour indicates the gas-star ratio in each region.Symbols are the same than those in Fig. 10.

an instantaneous burst (purple continuous line) and the other con-sidering a constant star formation (red dotted line). These modelsalone can only explain the position of a few clusters in this figure.As discussed in Section 4.1, we note that we assumed a Milky Way-like extinction law to correct the FN colour for internal extinction.A different dust attenuation curve would lead to a small reddeningin the FN colour with AFUV or AHα . However, as this effect is neg-ligible, its impact should be negligible too and will not modify thisanalysis. For explaining the position of the rest of the clusters wecan assume two scenarios.

(i) The UV-bright stellar clusters have experienced two instan-taneous star formation events, one that typically happened 100–400 Myr ago (following the FN colours) and a very recent starburstthat started between 3 and 13 Myr ago (according to the Hα equiva-lent width). Following this scenario, we can use W(Hα) to constrainthe age of the most recent star formation event, obtain the FN colourcorresponding to a stellar population with that very young age, andsubtract it to the observed FN colour to estimate the age of the olderstar formation event. This scenario will provide even older ages(20–100 Myr depending on the knot) for the very first ignition ofthe star formation activity than that determined by KLS09.

(ii) The star formation activity of the UV-rich stellar clusters canbe explained as a combination of on-going star formation plus astarburst event. Indeed, the right-hand panel of Fig. 18 also plotstwo extra lines which have been constructed assuming a linearcombination of continuous star formation/starburst, with weightsof 0.5/0.5 (blue dashed line) and 0.9/0.1 (green dot–dashed line),respectively. The position of the majority of the objects in thisdiagram is then easily explained according to this scenario.

The symbol colour in the right-hand panel of Fig. 18 indicatesthe gas-to-star ratio, Mgas/Mstars, in each region. Following this plot,it also appears that regions which possess a more intense recentstar formation contribution tend to have a higher gas-to-star ratiothan knots where the continuous star formation dominates (or witholder ages for the first instantaneous starburst event following theFN colour). Either way, this fact suggests that the gas has been

Figure 19. Comparison between the surface density of the cold gas consid-ering only the atomic component (y-axis) and the oxygen abundance (x-axis)for the UV-bright knots in NGC 1512/1510. Colours indicate the FUV −NUV colour. Symbols are the same than those in Fig. 10.

consumed more in those complexes with lower log[−W(Hα)]. Thisagrees with the finding by KLS09 that objects with redder FNcolours systematically have a lower amount of gas.

5.7.3 Relationships between stars, gas, and metallicity

KLS09 showed that the UV-rich clusters in the NGC 1512/1510 sys-tem do follow the Schmidt–Kennicutt scaling laws of star formation(Schmidt 1959, 1963; Kennicutt 1998; Kennicutt et al. 2007). Forall regions but those located in Zone 1, the SFR per unit area corre-lates well with the surface density of the cold gas considering onlythe atomic component (see fig. 20 in KLS09). Fig. 19 comparesthe surface density of atomic gas, Mgas/area, with the oxygen abun-dance for regions analysed here. Besides the known fact of missingthe molecular component in knots in Zone 1 (diamonds), and the

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Figure 20. Comparison between the gas-to-star ratio, Mgas/Mstars (y-axis)and the oxygen abundance (x-axis). Colours indicate the FUV − NUVcolour. Symbols are the same than those in Fig. 10.

relatively large dispersion of Mgas/area along Arm 1, we do not seeany dependence of the surface density of atomic gas on the metal-licity. Typical values of Mgas/area are between 3 and 9 M� pc−2,numbers that agree with the SFR/area found in the external regionsof spiral galaxies (Kennicutt et al. 2007; Bigiel et al. 2008, 2010a,b).

We also check if the gas-to-star ratio, Mgas/Mstars, has any de-pendence with the oxygen abundance. Fig. 20 compares both prop-erties. Remembering that the position of all knots within Zone 1(diamonds) are lower limits because of the lack of molecular gasmeasurements, again we do not see any correlation, as the decreaseof the Mgas/Mstars ratio with increasing oxygen abundance observedin galaxies (Lopez-Sanchez 2010; Lara-Lopez et al. 2013a). We doobserve, however, that regions with a higher amount of gas system-atically have bluer UV colours, as KLS09 already pointed out.

More interestingly, we calculate the effective oxygen yield foreach UV-rich complex, yO = ZO/ln (1/μ), where μ is the gasfraction with respect to the total baryonic (stars and gas) mass,μ = Mgas/Mbar, and ZO is the oxygen mass fraction.3 Follow-ing a closed-box model (Schmidt 1963; Searle & Sargent 1972;Edmunds 1990), a galaxy, which initially consists of gas with nostars and no metals, experiences instantaneous recycling through-out its life, and the products of stellar nucleosynthesis are neitherdiluted by infalling pristine gas nor lost via outflow of enriched gas.Therefore, the metallicity at a given time is only determined by thefraction of baryons that remain in gaseous form. Fig. 21 comparesthe observed oxygen abundances with those predicted by simpleclosed-box models. The theoretical yield of oxygen expected forstars with rotation following the Meynet & Maeder (2002) modelsis yO = 0.0074 (green continuous line). In galaxies, it is usuallyfound that the oxygen abundances are lower than those predictedby the closed-box model with yO = 0.0074, their effective yieldsbeing yO = 0.003–0.005 (Lee et al. 2003, 2007; Tremonti et al.2004; van Zee & Haynes 2006; Lopez-Sanchez 2010). This meansthat galaxies tend to lie above the green continuous line plotted inFig. 21, as indeed is happening in NGC 1510 (star), which has aneffective yield of yO = 0.003. But the opposite behaviour is ob-served in the UV-bright stellar clusters of NGC 1512; almost all of

3 The oxygen mass fraction is derived as ZO = f ×(O/H), f = 11.81 beingthe conversion factor from number to mass fraction.

Figure 21. Comparison of the observed oxygen abundance with those pre-dicted by simple closed-box chemical evolution models. The model follow-ing the theoretical yield of oxygen expected for stars with rotation followingMeynet & Maeder (2002) models, yO = 0.0074, is plotted with a greencontinuous line. Closed-box models with yO = 0.3, 0.03, and 0.003 are alsoshown with a yellow dot–dashed line, a red dotted line, and a purple dashedline, respectively. Symbols are the same than those in Fig. 10, the star rep-resents the position of the BCDG NGC 1510. The symbol colour indicatesthe FUV-based SFR.

them have effective yields higher than the theoretical yO = 0.0074value. This is even true for the regions located in the internal ring,for which we should expect even higher effective yields than thosecomputed here, as we lack data on the molecular gas. Typically, theeffective yields found in XUV complexes are between 0.03 and 0.3,i.e. between 1 and 2 orders of magnitude higher than the theoreticalvalue.

Therefore, these UV-bright knots have an oxygen abundancemuch higher than expected by the closed-box model, i.e. the H I

gas that has been used for the star formation already had a largefraction of metals. The mean effective yield found within the XUVregions belonging to Arm 1 is yO = 0.133, with a dispersion of0.119, while the mean effective yield obtained for the XUV clustersin the external debris of Arm 2 is yO = 0.049, with a dispersion of0.025. Again, and besides the dispersion, these numbers agree withthe hypothesis that the material within the knots belonging to thedestroyed Arm 2 have experienced a larger chemical enrichmentthan those located along Arm 1.

5.7.4 Where do the metals of the outer regions come from?

Bresolin et al. (2009a), Werk et al. (2010b, 2011), and Bresolinet al. (2012) already reported that the outer discs of spiral galaxiesare overabundant for their large gas fraction. Similarly high valuesare reported by Torres-Flores et al. (2014) when analysing TDGslocated in the long tidal tail of NGC 92. Indeed, Lopez-Sanchez(2010) found that a TDG candidate within the compact group HCG31 (member F), showed a high effective yield, yO = 0.03, thisobject being very different from the other galaxies analysed in hisstudy. The author explains this behaviour as a consequence of theaccretion of a large fraction of pre-enriched H I gas striped from themain galaxies after a fly-by encounter (Lopez-Sanchez, Esteban &Rodrıguez 2004a). Somehow, member F of HCG 31 resembles theproperties of the XUV complexes found around NGC 1512, as wewill discuss below.

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Adapting equation (1) in Bresolin et al. (2012),

(O

H

)= yO × t

f× SFR

MH I

, (22)

and knowing from the FN colours when the star formation startedin each knot, t, we can estimate how much the UV-bright re-gions have been enriched in oxygen, ( O

H ), since then. For this,we also consider the effective yield of oxygen, yO, computed inthe previous subsection. Averaging all regions of Arm 1, we find[12+log(O/H)] = 0.21, with a dispersion of 0.10 dex, while thesame computation averaging all regions of the destroyed Arm 2(Zones 3 and 4) yields [12+log(O/H)] = 0.32, the dispersionbeing 0.25 dex. This analysis suggests that about 400 Myr ago,before the interaction with NGC 1510 started, the metallicity inthe outskirts of NGC 1512 was 12+log(O/H) ∼ 8.1 in either arm.Therefore, the neutral gas located at the outer regions of NGC1512 had experienced metal enrichment even before the processeswe observe today started. Assuming that the central regions alsohave been enriched by 0.2–0.3 dex in the last 400 Myr (this mayactually be an upper limit to the real enrichment experienced inthose central clusters), they should have had oxygen abundances of12+log(O/H) ∼ 8.5–8.6, which means a difference of 0.4–0.5 dexwith respect to the outer regions.

Given the distance of many of the XUV complexes, particularlyTDG 1, to the centre of NGC 1512, it seems quite unlikely that themetal enrichment of the outer regions of this galaxy came from itsinner regions. In order to explain the flattening of the metallicitygradients, Werk et al. (2011) discussed mechanisms to transportmetals from the galaxy centres to their outskirts. Metal mixing incold neutral gas probably happens, but it seems to be a very slowprocess; assuming radial H I outflows of ∼10 km s−1 (as thoseobserved in the extended H I disc of NGC 2915 by Elson et al.2010), some few Gyr are needed to transport metals up to 25 kpcfrom the galaxy centre. Werk et al. (2011) considered that metaltransport may be occurring predominantly in a hot gas component,although both the drivers and the amount of mixing of metals in thehot-phase gas is unclear.

In any case, the amount of metals needed to enrich an almostprimordial gas to 12+log(O/H) ∼ 8.1 is large. Considering that∼75 per cent of the neutral gas is found in the outer part of thesystem (KLS09), this gas should have had ∼6.5 × 106 M� inform of oxygen to account for that oxygen abundance. Adding thismass to that existing in the centre (which we computed assuming12+log(O/H) ∼ 8.5, although this may have been larger), we de-rive an original oxygen abundance of 12+log(O/H) ∼ 8.85 (i.e. anincreasing of ∼0.35 dex) in the centre of NGC 1512. Taking intoaccount the mass–metallicity relation (Tremonti et al. 2004; Lara-Lopez et al. 2013b), a galaxy with 12+log(O/H) ∼ 9.24 shouldhave log(Mstars) ∼ 11. This value is at least one order of magnitudehigher than the stellar mass of NGC 1512 estimated using the op-tical colours and only 1/3 of the total dynamical mass (see tables1 and 3 in KLS09). Therefore, the metals (or at least an importantfraction) found in the outskirts do not seem to come from the innerregions of NGC 1512.

Ruling out the hypothesis that the gas itself has experiencedstar formation activity for longer periods of time, we suggest thatthe metals within the diffuse H I gas found in the outer regions of

4 Note that both Tremonti et al. (2004) and Lara-Lopez et al. (2013b) use theabsolute abundance scale given by photoionization models, which provideoxygen abundances that are at least 0.3–0.4 dex higher than those given bythe Temethod used here.

NGC 1512 came from low-luminosity dwarf galaxies which havebeen slowly accreted and merged into the system. This may havehappened long time ago (several Gyr); the gas was enriched inthe past as a consequence of these minor mergers and later stayedessentially untouched for a long period of time. The remnants ofstellar component of such dwarf galaxies are now so diffuse theyare almost impossible to distinguish, unless really deep optical data–μlim(V) ∼ 28.8 mag arcsec−2 or higher – are obtained. Relicsof old interactions have been found as diffuse stellar tidal streamsaround normal disc galaxies (Martınez-Delgado et al. 2009, 2010) oreven around dwarf galaxies (Martınez-Delgado et al. 2012). Thesestructures are actually expected as a consequence of the hierarchicalframework for galaxy formation following a � cold dark matter(CDM) cosmogony, as minor mergers (which should have beenvery common in early times) do not destroy pre-existing stellardiscs (Bullock & Johnston 2005; Johnston et al. 2008; Cooper et al.2010).

Alternatively, the gas in the outskirts of NGC 1512 was accretedfrom the intergalactic medium (IGM), that has been previouslyenriched with metals by outflows that happened on other galaxies.This mechanism was also proposed by Bresolin et al. (2012) and isreproduced by cosmic chemical evolution models (e.g. Kobayashi,Springel & White 2007). Both the enrichment by satellites or theaccretion of outflow-enriched IGM gas require that a significantfraction of the metals are cooled down and mixed with cold H I

before being available for further star formation, as the one observedtoday in the XUV disc of NGC 1512.

5.8 About the TDG nature of the XUV complexes

Interactions between galaxies may induce the development of tidaltails of material expelled from the parent galaxies into the IGM(e.g. Hibbard & van Gorkom 1996; Hibbard et al. 2001b). Objectswith masses typical of dwarf galaxies are usually formed from thedebris of these tidal tails (Duc & Mirabel 1994, 1998; Hibbard et al.2001a; Knierman et al. 2003; Lopez-Sanchez, Esteban & Rodrıguez2004a,b; Hibbard et al. 2005; Neff et al. 2005). When one of theseobjects appears to be self-gravitating, it is defined to be a TDG(Braine et al. 2000; Weilbacher & Duc 2001; Bournaud et al. 2004;Bournaud & Duc 2006). Contrary to dwarf galaxies, these entitiesdo not need a dark matter (DM) halo to rotate, as indeed TDGsare expected to have little to no DM. However, they are likelyto contain a significant amount of old stars and they should havechemical enrichment histories related to those of the parent galaxy(Duc et al. 2000; Weilbacher, Duc & Fritze-v. Alvensleben 2003), astidal tails do not show evidence of abundance gradients. Hence, thedefinitive confirmation that a knot in the outer regions of a galaxyis a TDG comes from both its kinematics (a self-gravitating object)and chemistry (metallicities similar to those found in the parentgalaxy).

Can we define as TDG candidates the XUV clusters found in theouter regions of the NGC 1512/1510 system? Our data do not allowus to investigate the intrinsic kinematics of the UV-bright regions inthe NGC 1512/1510 system, but our metallicities can provide somelight to this issue. Recently, Sweet et al. (2014) identified strongTDG candidates with oxygen abundances5 of 8.4 < 12+log(O/H)< 8.7. The sample analysed by Weilbacher et al. (2003) also pro-vided oxygen abundances typically higher than 8.3 dex in TDG

5 Note these values have been scaled here to the absolute Te-scale oxygenabundance, as Sweet et al. (2014) used the recent photoionization modelsby Dopita et al. (2013) to compute metallicities.

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candidates. Similarly, two objects found within the large M 81–M 82–NGC 3077 H I tidal stream also are TDG candidates, having8.4 < 12+log(O/H) < 8.7 (Croxall et al. 2009). The majority ofthe regions we have associated with Arm 2 in NGC 1512 are withinthe typical metallicity range for TDGs. However, as we discussedbefore, our data suggest that the original metallicity of these re-gions was much lower than it is now, 12+log(O/H) ∼ 8.1, as wesee along Arm 1. Hence, this material has clearly experienced adifferent chemical evolution than that observed in the central re-gions of NGC 1512, with 12+log(O/H) = 8.7–8.8. Furthermore,according their star formation histories and their comparison withthe close-box model, all these regions seem to have been recentlyborn from the large reservoir of neutral gas previously existing inthe external areas of NGC 1512, which was already enriched inmetals that were not coming from its inner regions. Therefore, thematerial from which the external UV-rich complexes were formedhas been not expelled from the centre of the galaxy, and hence theentities created within this gas do not fulfil the characteristics to bedefined as TDG candidates.

Note, however, that the H I structures we now see in the systemand from which the XUV complexes are formed have been tidallyinduced by the interaction between NGC 1512 and NGC 1510(KLS09). Therefore, besides their chemical history, these stellarclusters seem to have a tidal origin. We propose to define theseUV-bright, young, relatively low-metallicity, gas-rich, entities to be‘tidally induced star-forming clusters’ (TSFCs) in the galaxy out-skirts. The main quantitative observational properties of the TSFCsare that they are located in the outskirts of the galaxy (R/R25 > 1),still possess SFR as traced by their FUV emission (SFRFUV > 0),they still possess a lot of gas (Mgas/Mstars > 1), and that their metal-licities are at typically 0.2–0.3 dex lower than those found in thegalaxy centre (as difference of TDG candidates, that should havethe same metallicity that the inner regions of the galaxy). Hence,many of the UV-rich complexes found along the XUV discs of spi-ral galaxies are very likely to be made up by these entities. KnotsE and F within the disturbed HCG 31 group (Lopez-Sanchez et al.2004a) may be another example of TSFCs, as they were formedfrom material in the long H I tail that the compact group possesses(Williams, McMahon & van Gorkom 1991; Verdes-Montenegroet al. 2005), although in these knots the metallicities are slightlylower, 12+log(O/H) ∼ 8.05, than that observed in the UV-brightregions of the NGC 1512 system.

What is the situation of the so-called TDG 1 (KLS09) located at agalactocentric distance of 78 kpc? Following the H I map presentedby KLS09, TDG 1 belongs to Arm 2, and hence it should be associ-ated with our Zone 4. The adopted metallicity of this star-formingknot is 12+log(O/H) = 8.24, which quantitatively is the same oxy-gen abundance as that observed throughout the long Arm 1. It isvery interesting to note that, besides their low masses and distancefrom the centre of NGC 1512, the TDG candidate and the majorityof the regions studied in Arm 1 have these intermediate metallic-ities. This fact indicates that even in these far regions the gas hasbeen processed several times in order to enrich the very metal-poorand unprocessed material that presumably existed here in the past.Interestingly, Fig. 21 shows that the data of TDG 1 agrees with aclosed-box model with an effective yield of yO = 0.004, the ex-pected for a dwarf galaxy. However, note that the H I envelope thatsurrounds this object is really large (see fig. 8 in KLS09), but onlythe neutral gas within the UV-bright knot has been considered here(once the images were matched in resolution and PSF, as describedin Section 4), which is at least an order of magnitude lower than thereal H I reservoir. Hence, the effective yield of this object should be

higher than the observed one. Taking into account all available infor-mation about TDG 1 (recent star formation activity, relatively lowmetallicity, gas-to-star ratio, and H I kinematics) and our previousdiscussion about the nature of the XUV complexes we consider thatthis object is actually not a TDG but a star-forming clump within adenser H I cloud in the outskirts of the NGC 1512/1510 system, thatis, an extreme case of TSFC. Indeed, the H I cloud associated withthis object has very likely a tidal origin, as seen in the impressiveH I map of the system (fig. 11 in KLS09).

Finally, what is the nature of regions 4_17, 4_32, and 4_38 at theNW of Arm 2? These three clusters, particularly 4_17, show the low-est metallicities of the system, 12+log(O/H) ∼ 8.10, despite theirsurroundings having oxygen abundances of 12+log(O/H) ∼ 8.45.We propose two scenarios: (i) because of unknown reasons, theseclusters have experienced a quenching in their star formation, andtherefore the recycling of the gas has not been as high as in othernearby regions, or (ii) these are actually independent dwarf galax-ies which now are pulled into the internal areas of the system as aconsequence of the strong gravitational forces induced by the inter-action. Some spectroscopic properties (see Fig. 2 and Table 4), likethe strength of the emission lines and their low reddening coefficientand H I Balmer absorptions, as well as the fact that the Hα-basedSFR is almost an order of magnitude higher than the FUV-basedSFR, agree with the hypothesis that knot 4_17 actually is a dwarfgalaxy which has not been created from the debris of Arm 2. How-ever, its Hα kinematics (see Fig. 16) is undistinguishable from thatobserved in nearby knots. The Hα kinematics of regions 4_32 and4_38 also agree with that found in their surroundings. Therefore,we cannot confirm the hypothesis that these XUV complexes areindependent dwarf galaxies. In any case, we remember that knot3_20, which has both different kinematics and metallicities than itsnearby XUV regions, may actually be an independent dwarf galaxy.New, deeper optical data of all these clusters are needed to unveiltheir real nature.

6 C O N C L U S I O N S

We have used the 2dF instrument in combination with the AAOmegaspectrograph at the AAT to get deep, intermediate-resolution op-tical spectroscopy of 136 genuine UV-bright regions located inboth the spiral and the XUV discs of the NGC 1512/1510 sys-tem. In conjunction with the H I –provided by the LVHIS project(Koribalski 2008; Koribalski et al., in preparation) – and UV datawe analyse the relationships between gas and stars, the chemicalenrichment of the ISM, and the properties of the stellar popula-tions and star formation processes within these objects. Our mainconclusions are as follows.

(i) We confirm the detection of ionized gas in the huge majorityof the UV-rich regions in the system, as well as identified 17 back-ground galaxies and 1 foreground star. We confirm that in all casesexcept the centre of NGC 1512 photoionization by massive stars isthe main excitation mechanism of the gas. In the XUV complexes,only few (1–5) O7V stars are responsible for this.

(ii) Using a comprehensive analysis of the emission lines of [O II]λ3727, Hγ , Hβ, [O III] λ5007, Hα, and [N II] λ6583, in conjunctionwith a careful study of the most-common SEL techniques, we havecomputed the oxygen abundance and the N/O ratio for the majorityof the detected UV-bright regions. We provide chemical abundancemaps reaching a projected distance of 78 kpc (6.6R25) from thecentre of NGC 1512. We found significant differences betweenregions along the Arm 1 at the east, which have oxygen abundances

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8.25 � 12+log(O/H) � 8.45, and knots located in the externaldebris of Arm 2, which typically have 8.40 � 12+log(O/H) � 8.60.

(iii) Thanks to the detection of the auroral [O III] λ4363 line, wehave computed the chemical abundances of the ionized gas using thedirect Te method in four regions. The results agree within the errorsto those obtained using SEL methods, however precaution must betaken with the ionization degree of the gas and/or the area integratedto get the spectrum when using SEL techniques. In the case of NGC1510, we derive 12+log(O/H) = 8.24 ± 0.06 and a N/O ratio oflog(N/O) = −1.34 ± 0.05, results that agree with the typical valuesfound in other BCDGs. We detect the nebular He II λ4686 emissionline in the centre of NGC 1510. This line is attributed to the presenceof massive stars in the WR stage. Contrary to what was previouslysuggested, no local N enrichment is observed in this galaxy.

(iv) We have analysed the metallicity distribution of the sys-tem. Considering just a radial gradient involving all complexes itis difficult to see a radial decrease of metallicity from the centreof NGC 1512 to the outer regions. However, assuming both a ra-dial and an azimuthal gradient following the spiral arms we areable to clearly distinguish that Arm 1 has an almost flat oxygenabundance – 12+log(O/H) ∼ 8.35 – and flat N/O ratio – log(N/O)∼−1.4 – while regions located in the disrupted Arm 2 (inner armand external debris), show a large dispersion in oxygen abundances– 8.1 � 12+log(O/H) � 8.6 – and N/O ratios. Arm 2 has experi-enced an enhancement in star formation because of the interactionwith NGC 1510, flattening the radial metallicity at large radii, whileArm 1 still shows the original and poorly disturbed radial distribu-tion (gradient+flat in the outskirts).

(v) We have quantified the extinction within the UV-bright re-gions using the Balmer decrement and the theoretical Hα/Hβ andHγ /Hβ ratios expected for the oxygen abundances estimated ineach knot. Underlying stellar absorptions with values Wabs∼0.85 Åare needed to explain the Balmer decrements. The reddening coeffi-cient correlates with the oxygen abundance, indicating that regionswith higher metallicities have higher amount of dust. We providean intrinsic extinction map of the system.

(vi) Using the Hα emission-line profile, we are able to trace thekinematics of the system out to 78 kpc, which generally matcheswell with that provided by the H I kinematics (velocity differencesof less than 15 km s−1). We observe a hint of additional rotationpattern in the Ring (Zone 1). Some small velocity discrepancies(differences between 20 and 40 km s−1) are found in some clustersof the external debris and in particular areas of Arm 1. We locate aregion, 3_20, for which its Hα radial velocity differs by 136 km s−1

with respect its H I radial velocity. The upper limit to the metallicityof this knot is 12+log(O/H) = 8.1, which is 0.4 dex lower than thatfound in nearby complexes. We suggest that this region actually isan independent dwarf galaxy or its remnant.

(vii) Comparing the Hα-SFR and the FUV-SFR, we concludethat the ionized gas seems to be very localized at the centre of eachUV-rich cluster. FUV-SFR are systematically ∼4 times higher thanHα-SFR. We cannot explain the observed W(Hα) and FN coloursas just a single, instantaneous star formation event. A combinationof (i) two instantaneous events separated by 100–400 Myr or (ii)continuous star formation plus a recent (less than 13 Myr) starburst,are needed. The gas-to-star ratio decreases with both the FN colourand the W(Hα).

(viii) We do not observe any correlation between the surfacegas density or the gas-to-mass ratio with the oxygen abundance.However, a comparison with the closed-box model indicates thatthe XUV complexes have very high effective oxygen yields, asopposited to that found in dwarf galaxies, indicating that they

possess much more oxygen than expected from their gas massfraction.

(ix) Although the XUV complexes are very likely induced bythe interaction between NGC 1512 and NGC 1510, they cannotbe strictly defined as TDGs, as they have not been formed frommaterial stripped from the main galaxies, but from the diffuse gasalready existing at large galactocentric radii. We propose to definethese UV-bright, young, relatively low-metallicity, gas-rich entitiesto be TSFCs in the galaxy outskirts.

(x) Our data suggest that the gas in the outer regions of NGC 1512already had a metallicity of 12+log(O/H) ∼ 8.1 about 400 Myr ago,before the interaction with NGC 1510 started. The metals withinthe diffuse H I gas are very likely not coming from the inner regionsof NGC 1512, but probably from material accreted during the lifeof the galaxy either by minor mergers or by outflow-enriched IGMgas. We found a probable remnant of an independent dwarf galaxy(knot 3_20) and three knots (4_17, 4_32, and 4_38) which may alsobe independent systems, all of them having 12+log(O/H) � 8.1.This hypothesis agrees with the finding around other nearby largegalaxies such as M 83 or NGC 300, as well as constraints chemicaland dynamical models of galaxy evolution. Indeed, the extended,metal enriched diffuse gas should be common in spiral galaxiesaccording to the �CDM scenario.

In summary, the knowledge of the metallicity distribution in theoutskirts of galaxies give key clues about their evolution and starformation histories. This seems to be true even in an S0 galaxylike NGC 404 (Bresolin 2013). The analysis of the ionized gas inthe outer regions of galaxies, which provides the information aboutthe chemical evolution, very nicely complements the studies of theneutral gas using the 21 cm H I emission, which constrains thedynamical evolution of the galaxies. When used together, these twoanalyses provide a very powerful tool to disentangle the nature andevolution of the galaxies we now observe in the Local Universe.

AC K N OW L E D G E M E N T S

We thank Paul Dobbie for his help as AAO Support Astronomerwhile conducting our observations at the AAT. We also thank theanonymous referee for her/his very valuable comments and sugges-tions which have increased the quality of this paper. This researchmade use of images provided by the Survey for Ionization in Neu-tral Gas Galaxies (Meurer et al. 2006) which is partially supportedby the National Aeronautics and Space Administration (NASA).GALEX is a NASA Small Explorer, and we gratefully acknowledgeNASA’s support for construction, operation, and science analysisfor the GALEX mission, developed in cooperation with the CentreNational d’Etudes Spatiales of France and the Korean Ministry ofScience and Technology. CEL thanks the funding provided by Span-ish Ministerio de Economı a y Competitividad (MINECO) underthe project AYA2011-22614. This research has made extensive useof the NASA/IPAC Extragalactic Database (NED) which is oper-ated by the Jet Propulsion Laboratory, Caltech, under contract withthe National Aeronautics and Space Administration.

This research has made extensive use of the SAO/NASA Astro-physics Data System Bibliographic Services (ADS).

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APPENDI X A : R ELATI ONSHI PS BETWEENT H E TH E O R E T I C A L H I BA L M E R R AT I O SF O R C A S E B R E C O M B I NAT I O NA N D OX Y G E N A BU N DA N C E

Although it is very common to assume that the theoretical Hα/Hβ

ratio for case B recombination is 2.86, this value actually dependson both the electron temperature and electron density of the gas(Osterbrock & Ferland 2006; Storey & Hummer 1995). In H II re-gions within star-forming galaxies, the electron densities are typi-cally found in the low-density limit, ne = 100 cm−3, but electrontemperature depends on the effective temperature of the ionizationsource as well as on the metallicity of the gas, as the H II regioncools down with increasing amount of metals. We note that thedependence between the theoretical Hα/Hβ ratio and the oxygenabundance is not very strong, but in order to get the best results andminimize the uncertainties when determining the reddening coef-ficient of the ionized gas within star-forming galaxies it should beconsidered, specially at high metallicities where Te may be ratherlow. Strictly talking, the standard value of 2.86 is only valid forTe = 10 000 K and ne = 100 cm−3, which corresponds to an oxy-gen abundance of around 8.27. However, a cooler H II region (i.e. ahigh-metallicity H II region) will have a slightly higher theoreticalHa/Hβ ratio, reaching Hα/Hβ = 3.04 for Te ∼ 5000 K. Similarly,hotter H II regions (i.e. low-metallicity H II regions) will have lowertheoretical Ha/Hβ values (e.g. Hα/Hβ = 2.75 for T ∼ 20000 K).

Here, we have used the data base of H II regions in star-forminggalaxies presented by Lopez-Sanchez & Esteban (2010b, Tables 2and 3), for which a direct estimation of the electron tempera-ture is available, to get an empirical relationship between the

Table A1. Theoretical H I Balmer ratios for Case B recombination at dif-ferent electron temperatures and ne = 100 cm−3, computed from Storey &Hummer (1995), and approximate oxygen abundance for which the ionizedgas has a particular electron temperature.

Te Hα/Hβ Hγ /Hβ Hδ/Hβ Appr. 12+log(O/H)

5000 3.04 0.458 0.251 8.886000 2.99 0.461 0.254 8.757000 2.95 0.463 0.255 8.638000 2.91 0.465 0.257 8.509000 2.89 0.467 0.258 8.4010 000 2.86 0.468 0.259 8.2711 000 2.84 0.470 0.260 8.1512 000 2.83 0.471 0.261 8.0513 000 2.81 0.472 0.261 7.8814 000 2.80 0.472 0.262 7.7515 000 2.79 0.473 0.262 7.6216 000 2.78 0.473 0.263 7.5019 000 2.76 0.475 0.264 7.2020 000 2.75 0.475 0.264 7.10

Note. The bold font values indicate the standard values assumed in theliterature.

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Figure A1. Dependence of the theoretical H I Balmer ratios, Hα/Hβ (top), Hγ /Hβ (middle), and Hδ/Hβ (bottom) on the electron temperature, Te (left) andthe oxygen abundance (right), assuming Case B recombination and ne = 100 cm−3. A fit using a quadratic (for Te) or third degree (for oxygen abundance)polynomial is shown with a continuous line. The standard theoretical values for Te = 10 000 K and ne = 100 cm−3, Hα/Hβ = 2.86, Hγ /Hβ = 0.468, andHδ/Hβ = 0.259, are shown with dotted red lines.

theoretical Hα/Hβ ratio for case B recombination and oxygenabundance. Just for completeness, for the extreme low metallicityregime, 12+log(O/H) ∼ 7.2–7.1, we have used the values given byIzotov, Thuan & Guseva (2005) and Izotov & Thuan (2007). Weconsidered the theoretical ratios of H I Balmer lines expected forcase B recombination given by Storey & Hummer (1995) assumingne = 100 cm−3. Table A1 compiles these values.

Top-left panel of Fig. A1 plots the dependence of the Hα/Hβ ratiowith the electron temperature. We fitted a quadratic polynomial

between the logarithm of the Te (independent variable) and thetheoretical Hα/Hβ ratio (dependent variable), with the result

Hα

Hβ= 10.35 − 3.254 log Te + 0.3457(log Te)2, (A1)

being the correlation coefficient r = 0.9921, and the dispersiond = 0.0024. Top-right panel of Fig. A1 shows the dependence ofHα/Hβ ratio with the oxygen abundance. Similarly, we fitted adegree 3 polynomial between O/H (independent variable) and the

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theoretical Hα/Hβ ratio (dependent variable) to find

Hα

Hβ= −24.767 + 10.882x − 1.442x2 + 0.0640x3, (A2)

being x = 12+log(O/H). The correlation coefficient is r = 0.9479,and the dispersion is d = 0.001 16.

We have repeated this analysis for the case of the Hγ /Hβ andHδ/Hβ ratios. The dependence of these ratios with both the electrontemperature and oxygen abundance is also shown in Fig. A1. Theresults are

Hγ

Hβ= 0.0254 + 0.1922 log Te − 0.0204(log Te)2, (A3)

with r = 0.9907 and d = 0.000 29;

Hδ

Hβ= −0.071 32 + 0.1436 log Te − 0.0153(log Te)2, (A4)

with r = 0.9903 and d = 0.000 23;

Hγ

Hβ= 1.7871 − 0.5248x + 0.0704x2 − 0.003 17x3, (A5)

with r = 0.9477 and d = 0.000 29;

Hδ

Hβ= 1.5626 − 0.5110x + 0.0674x2 − 0.002 98x3, (A6)

with r = 0.9982 and d = 0.000 23.

S U P P O RT I N G IN F O R M AT I O N

Additional Supporting Information may be found in the online ver-sion of this article:

Table 2. Dereddened line intensity ratios with respect to I(Hβ) = 1and other spectral properties of the regions.Table 3. Reddening coefficient, important emission-line ratios andparameters used by empirical calibrations and derived oxygen abun-dances and N/O ratio for our sample regions.Table 6. Properties of the UV-rich regions within the NGC1512/1510 system.(http://mnras.oxfordjournals.org/lookup/suppl/doi:10.1093/mnras/stv703/-/DC1).

Please note: Oxford University Press are not responsible for thecontent or functionality of any supporting materials supplied bythe authors. Any queries (other than missing material) should bedirected to the corresponding author for the paper.

This paper has been typeset from a TEX/LATEX file prepared by the author.

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