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Piecing together the puzzle of NGC 5253: abundances, kinematics and WR stars

We present Gemini-S/GMOS-IFU optical spectroscopy of four regions near the centre of the nearby (3.8 Mpc) dwarf starburst galaxy NGC 5253. This galaxy is famous for hosting a radio supernebula containing two deeply embedded massive super star clusters, surrounded by a region of enhanced nitrogen abundance that has been linked to the presence of WR stars. We detected 11 distinct sources of red WR bump (CIV) emission over a 20"(~350 pc) area, each consistent with the presence of ~1 WCE-type star. WC stars are not found coincident with the supernebula, although WN stars have previously been detected here. We performed a multi-component decomposition of the H\alpha\ line across all four fields and mapped the kinematics of the narrow and broad (FWHM = 100-250 km/s) components. These maps paint a picture of localised gas flows, as part of multiple overlapping bubbles and filaments driven by the star clusters throughout the starburst. We confirm the presence of a strong H\alpha\ velocity gradient over ~4.5"(~80 pc) coincident with the region of N/O enhancement, and high gas density known from previous study, and interpret this as an accelerating ionized gas outflow from the supernebula clusters. We measure the ionized gas abundances in a number of regions in the outer IFU positions and combine these with measurements from the literature to assess the radial abundance distribution. We find that the O/H and N/H profiles are consistent with being flat. Only the central 50 pc exhibits the well-known N/O enhancement, and we propose that the unusually high densities/pressures in the supernebula region have acted to impede the escape of metal-enriched hot winds from the star clusters and allow them to mix with the cooler phases, thus allowing these freshly processed chemicals to be seen in the optical.


Introduction
Local starbursts are the present-day manifestations of the intense star formation thought to have been instrumental in building galaxies in the early Universe. Starbursts occur in dense, gasrich environments where the structure and kinematics of the ionized medium reflect the intense energy input from massive stars via stellar winds and supernova explosions. These winds can disrupt the surrounding gas, drive material out from the galaxy, and trigger (or quench) further star formation.
NGC 5253 is an excellent nearby (distance of 3.8 Mpc; Sakai et al. 2004) example of a starburst. This Blue Compact Dwarf (BCD) has ongoing active star formation at a rate of ∼0.2 M yr −1 (Calzetti et al. 2007;López-Sánchez et al. 2012), and contains a large number of compact young (∼1-12 Myr) star clusters (Cresci et al. 2005;Gorjian 1996;Harris et al. 2004;Vanzi et al. 2006) located within the central ∼250 pc starburst zone. Optical studies have been able to identify only very few clusters older than 20 Myr (Harris et al. 2004;Tremonti et al. 2001), supporting the extreme youth of this starburst. Recently three potentially massive ( 10 5 M ) and old (1-2 Gyr) star clus-Based on observations with the Gemini-South telescope under programme GS-2008A-Q-25. ters have been found in the outskirts of the galaxy underlining the fact that the current starburst is just the most recent star forming episode (Harbeck et al. 2012). López-Sánchez et al. (2012) concluded that the starburst was most likely triggered by the infall of an H i cloud (Meier et al. 2002), although they could not rule out a contribution from ram pressure stripping as the galaxy moves through the dense intergalactic medium of the Centaurus A group (Karachentsev et al. 2007).
Hα images show a centrally concentrated region extending ∼30 in the north-south direction, surrounded by a chaotic ionized structure with radial filaments, bubbles and loops (Calzetti et al. 1999;Marlowe et al. 1995;Martin 1998). Dense clouds within these structures produce rapidly varying levels of extinction. The morphology of the X-ray emission that traces the hightemperature gas follows the Hα distribution closely, and is consistent with a high energy-density outflow similar to that found in the dwarf starburst galaxy NGC 1569 (Strickland & Stevens 1999;Summers et al. 2004).
NGC 5253 is classified as a low-metallicity system (∼0.2-0.3 Z ; ) with a fairly uniform oxygen distribution. Famously, however, a factor of 2-3 nitrogen enhancement in the central regions has been found in a number of targeted studies López-Sánchez et al. 2007;Walsh & Roy 1989). Evidence for the presence of nitrogen-rich WN-and carbon-rich WC-type Wolf-Rayet stars (Schaerer et al. 1997) in the same N-enriched regions led to the suggestion that these WR stars could be the cause of this enhancement Walsh & Roy 1989). Indeed, in general, WR galaxies (those with large WR populations) as observed with SDSS samples exhibit elevated N/O ratios relative to non-WR galaxies (Brinchmann et al. 2008). Monreal-Ibero et al. (2010 presented the results of the first spatially-resolved study of the ionized gas kinematics and abundances in the central regions of NGC 5253 with VLT/FLAMES integral field unit (IFU) observations. They found that, while the abundances of O, Ne and Ar remain constant within 0.1 dex over the mapped area, there are in fact two regions of elevated N/O: one strong enhancement associated with the supernebula, and one weaker one associated with two older (∼10 Myr) clusters towards the south-west. The fact that signatures of WR stars are not found at the location of this second enhancement led to the conclusion that there must be a more complex relationship between the presence of WR stars and N/O excess than previously thought. Interestingly, James et al. (2012b) find high N/O ratios in Haro 11 in a region with an integrated age just older than the WR phase (∼7 Myr), but normal N/O ratios in two regions with large WR star populations, implying that the mixing time-scales are fairly long.
In this paper we present optical IFU observations of four regions within the NGC 5253 starburst, including the supernebula region and two additional fields 150-250 pc away. We begin by presenting the distribution of carbon-rich Wolf-Rayet (WR) stars (WC-type stars; via the ∼5800 Å red WR bump) in all four regions (Section 3), since this has not been examined previously. We then present a detailed study of the kinematical properties of the ionized gas in all four regions (Section 4) and the nitrogren and oxygen distribution in the two outer regions (Section 5), and in comparison to measurements from the literature. A detailed study of the He i abundance distribution within the central regions will be presented in a forthcoming paper (Monreal-Ibero et al. in prep.). Finally we discuss the state of the ISM across the whole starburst region, including the correspondence between the kinematics, chemical composition and presence of WR stars in relation to the enigmatic supernebula.

Observations
In 2008 queue-mode observations using the Gemini-South Multi-Object Spectrograph (GMOS) Integral Field Unit (IFU; Allington-Smith et al. 2002) in one-slit mode were obtained of four regions near the centre of NGC 5253 (programme ID: GS-2008A-Q-25, PI: M. Westmoquette), in 0.5-0.7 arcsec seeing 1 (∼9-12 pc at the distance of NGC 5253) and photometric conditions. A nearby bright star was used to provide guiding and 1 Position 1 was observed in exceptional seeing of 0 . 46. tip-tilt corrections using the GMOS on-instrument wave front sensor (OIWFS).
The one-slit mode IFU gives a field-of-view (FoV) of 5 × 3.5 arcsec (∼90 × 60 pc) sampled by 500 contiguous hexagonal lenslets of 0 . 2 diameter. We took a set of dithered exposures at each position, offsetting by integer spaxel values (i.e. 0 . 2) and used the R831 grating to give a spectral coverage of 4750-6850 Å at a dispersion of 0.34 Å pix −1 . An additional block of 250 lenslets (covering 5 × 1.7 arcsec), offset by 1 arcmin from the object field, provides a dedicated sky view. Positions 1, 2 and 4 were tilted by 5 • in order to avoid light from the galaxy from contaminating the sky fibres. Table 1 lists details of the observations for each position.
The GMOS spectrograph is fed by optical fibres from the IFU which reformats the arrangement of the spectra for imaging by the detector. In order to remove the fibre-to-fibre throughput and detector pixel-to-pixel sensitivity differences, and enable the wavelength and flux calibration of the data, a number of bias frames, flat-fields, twilight flats, arc calibration frames, and observations of the photometric standard star Hiltner 600, were obtained as part of the standard baseline calibration plan.
The IFU positions were chosen to cover selected areas of the disturbed ionized interstellar medium in the inner regions of NGC 5253, covering the nebulosity surrounding the central SSCs (Alonso-Herrero et al. 2004;Turner et al. 2000) and a further two regions to the south containing a number of ionized clumps. In Fig. 1, we show the position of the IFU fields on an HST/ACS-HRC colour composite image (PID 10609, P.I. Vacca).

Reduction
The field-to-slit mapping for the GMOS IFU reformats the layout of the fibres of the sky to one long row containing blocks of 100-150 object fibres interspersed by blocks of 50 sky fibres. In this way, all 750 spectra (500 object + 250 sky) can be recorded on the detector simultaneously.
Basic data reduction was performed following the standard Gemini reduction pipeline (implemented in iraf 2 ). Briefly, a trace of the position of each spectrum on the CCD was first produced from the flat field. Then throughput correction functions and wavelength calibration solutions were created and applied to the science and standard star data, and the averaged sky spectrum (computed from the separate sky field) was subtracted, before the individual spectra were extracted using the flat-field trace. The spectra were then flux calibrated from observations of the standard star. Although all our observations were taken under photometric conditions, that the standard star was not observed on the same night as our science observations means that the accuracy of the spectrophotometric flux calibration cannot be absolute. To test the accuracy of the relative flux calibration across the field, we re-reduced one science frame for IFU position 1 without sky subtraction (but with the standard fibre-to-fibre throughput correction based on the twilight flat), and measured the flux of the prominent [O i]λ5577 sky line across the whole field. Very little variation was found with no systematic patterns; the mean flux was measured to be 1.17 × 10 −16 erg s −1 cm −2 with a standard deviation of 8.80 × 10 −18 erg s −1 cm −2 (i.e. a ∼7.5% variation). This is at a level sufficient for our science goals.  The data were converted into standard 'cube' format using gfcube with a spatial (over)sampling rate of 0 . 1, and corrected for the effects of differential atmospheric correction (using an iraf-based code written by J. R. Walsh based on an algorithm described in Walsh & Roy 1990), Finally, the individual dithered exposures for each position were combined using a custom pyraf script written by J. Turner (Gemini).
In order to determine an accurate measurement of the instrumental contribution to the line broadening, we fitted single Gaussians to isolated arc lines on a wavelength calibrated arc exposure, and found the resolution to be 1.2 Å at Hα (55 km s −1 ) and 1.3 Å at Hβ (80 km s −1 ). The fit residuals shows that the instrumental line profile is extremely close to a Gaussian. Fig. 2 shows an example reduced and flux calibrated spectrum extracted from a 13-spaxel region covering a bright ionized knot in IFU position 3 (which we define as region 1 in Section 5). A plethora of emission lines are detectable from neutral to doubly ionized species, and in this spectrum the red Wolf-Rayet bump is also clearly seen. Line identifications were made using the lists compiled by Tsamis et al. (2003) and López-Sánchez et al. (2007).

The carbon-rich Wolf-Rayet star population
Because of the link posited in previous studies between WR stars and nitrogen enhancements, it is of particular interest to understand the population and distribution of WR stars within NGC 5253. Here we discuss our results on the WC star population, together with other WR star detections from the literature, Example spectrum summed over 13 spaxels covering a bright ionized knot in IFU position 3 (defined as region 1; see Section 5). All identified emission lines are labelled, including the broad C iv λ5808 feature originating in the atmospheres of WC-type Wolf-Rayet stars.
which together provide a view of the WC and WN star population within the central ∼350 pc of the galaxy.
As exemplified in Fig. 2, we detect strong WR-star signatures in some spectra in the form of the red WR bump. This feature is a blend of the broad C iii λ5698 and C iv λ5808 emission lines originating in the atmospheres of carbon-rich WC-type WR stars. At low metallicities such as in NGC 5253, early-type WC stars (WCE-types) are expected to be the dominant sub-type (Crowther 2008).
We mapped out the strength and location of the red WR bump in our four IFU fields. First we measured the continuum level in two windows either side of the WR feature and subtracted the average. We then collapsed the datacubes over the wavelength range 5765-5875 Å to produce integrated maps of the red WR bump for all four positions. These are shown in Fig. 3. The size of each WR emission region is consistent with being a point source (i.e. the size of the seeing disk, ∼0.5-0.7 arcsecs). To estimate the number of WR stars in each source, we used the relation given by (López-Sánchez & Esteban 2010) for the C iv λ5808 luminosity of a WCE star for a given metallicity (their equation 8). At the metallicity and distance of NGC 5253, this relation predicts a WCE star luminosity of 2.3 × 10 36 erg s −1 ; given the uncertainties in both our measured fluxes and the WCE star luminosity relation, each source is consistent with ∼1-few WCE stars (see Table 4).
For a typical cluster mass of 5 × 10 3 M (Harris et al. 2004) and age of 4 Myr (Table 2), the number of WR stars predicted by Starburst99 (Leitherer et al. 1999) is ∼2. Firstly, this is in good agreement with our finding of ∼1 WR star per source. Secondly, this very low number means that the presence of WR stars in any given cluster is dominated by stochasticity. Not every cluster of the right age and mass will contain a WR star (i.e. the IMF is not fully sampled in clusters of this mass). Indeed, simple Monte Carlo sampling of a Chabrier IMF over a total mass of 5 × 10 3 M , shows that only ∼1-10 stars >30 M are expected.
We do not detect a red WR bump at the location of the supernebula (the peak of emission in the optical; cluster H04-1 3 ), but find a number of WC stars within 3 including one coincident with cluster H04-21, and a concentration at the UV emission peak (cluster H04-4). We also detect four more WR emission sources much further from the supernebula in GMOS positions 3 and 4. Schaerer et al. (1997) detected W-R features (from both WNand WC-type stars) at both near the peak of emission in the optical and the UV, equivalent to ∼10 WC stars and a few 10s of WN stars in each of the two regions. Due to a pointing error during their observations they did not in fact cover the optical and UV flux peaks (clusters H04-1 and 4), but positioned their slit ∼1 . 5 to the west (D. Schaerer, priv. comm. 2007). Sidoli (2010), however, did obtain optical long-slit spectroscopy of the optical and UV peaks with VLT/UVES, and detected 5 WN and 1 WC star in a 1 ×3 region covering the supernebula (cluster H04-1), and 10 WN and 2 WC stars at the UV peak (cluster H04-4). López-Sánchez et al. (2007) found ∼1 WN star each in their regions A, B (a few arcsecs west and north-west of the supernebula), and ∼10 WN stars in region C (equivalent to cluster H04-4). Accounting for the different distances and WR subtype luminosities adopted by López-Sánchez et al. (2007), these values are in reasonable agreement with those found by Sidoli (2010) in the corresponding areas. Monreal-Ibero et al. (2010) identified seven distinct regions containing blue-bump emission from WN-type stars in their IFU data covering the nuclear regions. The location of their detections are consistent with those found by the above authors. Fig. 4 shows a zoom-in of the central regions of NGC 5253 summarising the IFU data on the location of the WC star population from this study and the WN stars from Monreal-Ibero et al. Only certain classes of WR stars, and possibly certain O stars, are predicted to be hot enough to ionize nebular He ii (Crowther 2007;Kudritzki 2002), whose ionization potential is hν = 54.4 eV, and so far only one WC star has been associated with a He ii nebula (WC star MC45 in nebula BCLMP38b in M33; Kehrig et al. 2011). Our spectroscopy shows that both our He ii-WC regions have higher than average [O iii]/Hβ line ratios, consistent with the presence of a hard radiation field. Neither, however, show convincing evidence for the presence of shocks in the traditional BPT diagrams (Baldwin et al. 1981, not shown).
Almost all previous studies of the WR population in NGC 5253 have focussed on the central regions. However, Sidoli (2010) examined a region ∼5 to the north of the supernebula, and found evidence for 1 WN star, and López-Sánchez et al. (2007) found evidence for a further WN star in their region D (coincident with cluster H04-2, ∼10 south of the supernebula and ∼5 to the east of GMOS IFU position 3). We detect redbump WR emission consistent with the presence of single WC stars in 4 locations in GMOS positions 3 and 4. In position 3, one of these WC stars is coincident with H04-17, a known cluster of mass 2-3×10 3 M and age <4 Myr (Harris et al. 2004). The remaining 3 sources are not coincident with any previously identified clusters. Table 2 summarises the detections of WN and WC stars in the regions defined in both this study and that of Monreal-Ibero et al. (2010), where we have ordered them by decreasing EW(Hβ) (or increasing age using the predictions of Starburst99).
For all stars that pass through a WR phase, it is thought that the WN stage precedes the WC stage (Crowther 2007). Thus, one would expect WN stars to be present predominantly in the younger regions, and WC stars in the older regions. Disregarding regions 2 and 6 (where we do not detect a red WR bump, and that were not covered by the observations of Monreal-Ibero et al. 2010), there does appear to be a transition at ∼ 3.4-4.4 Myr, below which there is a higher probability of finding WN stars, and above which there is a higher probability of finding WC star. Clear exceptions to this are regions H ii 2 (very young but exhibits evidence for both WR types) and WR5 (which, according to its Hβ flux, is older but has WN star signatures). In neither of these locations did Harris et al. (2004) identify a cluster, but as we have previously mentioned, their catalogue is far from complete. Furthermore, Monreal-Ibero et al. (2010) were only able to determine the location of any WR emission to a precision of ∼1 . 5, which is not sufficient to be sure which source (if identified) is responsible. Projection effects also add to the difficulties of disentangling which source is associated with which.
In summary, we detect 11 distinct sources of C iv λ5808 emission throughout the NGC 5253 starburst, each consistent with the presence of ∼1 WCE type star. Some are associated with known young clusters catalogued by Harris et al. (2004, although this study is by no means complete), and two are coincident with the location of WN stars detected by Monreal-Ibero et al. (2010). The finding of WR stars spread out over 20 (∼350 pc) attests to the large area over which the starburst has occurred.  file, and attributed them (at least near the central SSC) to cluster winds and multiple expanding shells.

Emission line kinematics
As mentioned above, Sidoli (2010) presented VLT/UVES long-slit echelle spectroscopy of the optical and UV peaks (knots 1 and 2 from the study of Monreal-Ibero et al. 2010). They also found three Gaussian components were required to satisfactorily fit the emission line profiles: a narrow (FWHM=30-40 km s −1 ) component, a broad (FWHM=100-150 km s −1 , blueshifted by 10-30 km s −1 with respect to C1) and a third, fainter component intermediate in width and radial velocity. With a spectral resolution of R∼28 000 (∼11 km s −1 at Hα), the fact that they still needed a broad component suggests that it is not simply a convolution of many narrower, unresolved components (c.f. Melnick et al. 1999, for 30Dor).

Decomposing the line profiles
We fitted multiple Gaussian profile models to the Hα emission line using an idl χ 2 fitting package called pan (Peak ANalysis; Dimeo 2005), to recover kinematic information about the ionized gas in the four IFU positions. The high spectral resolution and signal-to-noise (S/N) of our data, together with the fact that the instrumental line profile is so close to a Gaussian (Section 2.1), have allowed us to quantify the Hα line profile shape to a high degree of accuracy. We can therefore be confident of the identification of faint components and/or those with very small ( 10 km s −1 ) velocity separations.
To account for the variable number of Gaussian components in each line profile, we fitted a single, double, and triple Gaussian component initial guess to each. Line fluxes were constrained to be positive and widths to be greater than the instrumental contribution to guard against spurious results. Multi-component fits were run several times with different initial guess configurations (widths and wavelengths) and the fit with the lowest χ 2 value was kept. However, we note that the χ 2 minimisation routine employed by pan is very robust with respect to the initial guess parameters.
To determine how many Gaussian components best fit an observed profile (one, two or three), we used a likelihood ratio test to determine whether a fit of n+1 components was more appropriate to an n-component fit. This test says that if the ratio of the χ 2 statistics for the two fits falls above a certain threshold, then the fits are considered statistically distinguishable, and the one with the lower χ 2 can be selected. Here we determine the threshold ratio by visual inspection of a range of spectra and fits. This method is a more generalised form of the formal F-test that we have used in previous work, e.g. Westmoquette et al. (2009Westmoquette et al. ( , 2011. We are aware of work that cautions against the use of the F-test to test for the presence of a line (or additional line component; e.g. Protassov et al. 2002), but given the absence of a statistically correct alternative that could be applied sensibly to the volume of data presented here, we have chosen to opt for this generalised χ 2 ratio test approach.
This test, however, only tells us which of the fits (single, double or triple component) is most appropriate for the corresponding line profile. Experience has taught us that we need to apply a number of additional, physically motivated tests to filter out well-fit but physically improbable results. We specified that the measured FWHM had to be greater than the associated error on the FWHM result (a common symptom of a bad fit), the fluxes of all components should be >0, and we rejected any fits where χ 2 single /χ 2 double = 0 (another symptom of a bad fit). The final step of the fitting procedure was to assign each fit component to a particular map in such a way as to limit the confusion that might arise during analysis of the results, such as discontinuous spatial regions arising from incorrect component assignments. In positions 1, 2 and 3, we found that the maps were most coherent when we assigned the broadest component to C2, and after that, the brightest as C1 and the faintest to C3. In position 4, we simply assigned the brightest component to C1 (regardless of width), then the faintest component to C3. This approach helps to limit the confusion that might otherwise arise during analysis of the results where discontinuous spatial regions might arise from incorrect component assignments.
Experience has also shown us that the errors that pan quotes on its fit results are an underestimate of the true uncertainties, particularly once uncertainties on our post-fitting tests and filters are taken into account. We can therefore estimate more realistic errors through a visual re-inspection of the profile+fit after knowing which one was selected by our tests, and by taking into account the S/N of the spectrum. We estimate that errors in FWHM range between 0.25-2 km s −1 for C1, 4-20 km s −1 for C2 and 15-25 km s −1 for C3, whereas those for radial velocities range between 0.1-3 km s −1 for C1, 2-5 km s −1 for C2 and 10-30 km s −1 for C3. A 5-10 percent increase in the C1 and C2 uncertainties may occur with the addition of a third component.
Figs. 5, 6 and 7 show the resulting FWHM and radial velocity maps for positions 1, 2, 3 and 4, respectively. In these figures we also include a number of example Hα line profiles and best-fitting Gaussian models from certain spaxels. These examples were selected to illustrate the discussion below, but also demonstrate the high quality of the spectra and the accuracy of the line-fitting.

Positions 1 and 2
As mentioned above, Monreal-Ibero et al. (2010) present a detailed analysis of the emission line kinematics in the central regions. However, the higher S/N and spatial resolution of our observations permits us to examine the emission line profiles and kinematical morphologies in the central regions in more detail (positions 1 and 2 were taken in seeing of ∼0.5 ). We present our results in the following.
Monreal-Ibero et al. (2010) and Sidoli (2010) find broad line components in the supernebula region with widths of FWHM∼80-150 km s −1 . In our analysis, we have been able to identify components with FWHM>250 km s −1 in a handful of spaxels. These very broad components are faint (∼5 percent of the total flux); an example is shown in Fig. 5b. In these spaxels, the need for this additional component is very clear from the residuals and χ 2 of the fit. In most of the field, however, the broad component line widths are consistent with those found in the aforementioned studies (FWHM∼100-150 km s −1 , comprising 10-50 percent of the total line flux).
Monreal-Ibero et al. (2010) identified a velocity gradient in the broad component extending over ∼4 . 7 (∼85 pc) centred on the supernebula and oriented NW-SE. This gradient is clearly seen in our C2 velocity map (Fig. 5 bottom centre panel), and there are clear hints of this gradient also in the secondary narrow component (C3). Our IFU map shows that this gradient is present across the entire C2 map, with an appearance similar to a rotation field. However, there are distinct peaks on the red and blue side, hinting that there may be other explanations.
To investigate this further, we defined a 0 . 7-wide pseudoslit crossing the supernebula, intersecting the two peaks, and oriented perpendicular to the line of nodes tracing zero velocity (i.e. along the maximum velocity gradient, as shown in Fig. 5). From this we extracted the radial velocities of the three line components to plot in Fig. 8. From this figure, it is clear that whereas C1, the dominant narrow component, remains fairly static across the whole region with velocities varying by only ∼30 km s −1 , C2 and C3 trace a linear gradient in velocity from ∼−50 to ∼+50 km s −1 within the central ∼4 (70 pc). Beyond this, C2 extends to even higher velocities, but is not followed by C3. All three components cross at zero velocity within ∼0 . 5 of the location of the supernebula. Rodríguez-Rico et al. (2007) presented observations of the radio recombination line H53α over a similar region, and found a ∼50 km s −1 gradient over ∼3 . 5 in the same NW-SE direction. That this is in very good agreement with our Hα measurements indicates that the ionized material throughout the gas column in this region all behaves similarly.
Monreal-Ibero et al. (2010) interpreted this velocity gradient (seen by them only in the broad component) as tracing an outflow from the SSCs in the supernebula, and our data here also support this conclusion. The fact that the velocity gradient exists in the broad line component is difficult to reconcile with rotation since such a turbulent rotating structure could not maintain coherence for long. We therefore conclude that C2 and C3 trace an accelerating ionized gas outflow (although we cannot rule out the possibility that the outflow is superimposed on a more general rotation-like bulk motion existing separately to the majority of the ionized gas traced by component C1). The fact that the supernebula is located very close to the zero-point of the velocity gradient implies that the clusters embedded there are the driving sources of the outflow.
Proceeding with the outflow interpretation, its inclination, i, is still unknown; assuming i=45 • the deprojected maximum outflow speed is ∼70 km s −1 . This velocity is roughly equivalent to the line width of C3 (FWHM∼30-100 km s −1 ), but is significantly less than the width of C2 (FWHM∼100-250 km s −1 ). This suggests that the turbulent energy in the outflow exceeds its bulk flow energy.

Position 3
IFU position 3 covers a region ∼6-10 to the south-west of the supernebula containing two H ii knots, the eastern of which is associated with a known cluster of mass 2-3×10 3 M and age <4 Myr (H04-17). A number of other young star clusters not catalogued by Harris et al. (2004) are also located in the field.  C2 (the broad component) is identified across the entire field, and exhibits line widths extending to FWHM>200 km s −1 . The broadest line components are generally also blueshifted (example line profile (b) in Fig. 6). Around the western Hα knot, the C2 component is fainter, narrower (∼100 km s −1 ), and redshifted compared to the surrounding material (∼30 km s −1 ); see example line profile (a). The corresponding C1 component is blueshifted, suggesting the gas is expanding or outflowing from this Hα knot. To the south of the IFU field, we see regions of clear line splitting where the two components are more equally matched in terms of flux and width (example c).
This position partially overlaps with the IFU field-of-view of Monreal-Ibero et al. (2010), and here our C1 radial velocities are in good agreement with their measurements. They find similar regions around their knot 3 in the south-west of their field where two distinct narrow components are seen. They conclude that the environment here (clusters H04-3 and H04-5; see Fig. 3) must be more evolved than around the supernebula, and that these older clusters have cleared out their gas, leaving broken shells ionized by the remaining hot stars.

Position 4
Position 4 covers an interesting arc-shaped feature seen in Hα ∼12 . 5 (220 pc) to the south of the supernebula. There are a handful of young clusters in this area that appear to be located mostly along the rim of the Hα arc, including one identified by Harris et al. (2004, H04-27). The very broad Hα component seen in much of position 3 is not seen here. Line widths mostly remain below 100 km s −1 across the field, with only a few regions where line widths extend up to ∼150 km s −1 . The Hα line shape at the peak of emission is consistent with a single Gaussian. Across the rest of the arc, the line shape is characterised by a double- Gaussian (examples (b) and (c)) with velocity differences between the components of 20-30 km s −1 . In the north of the field, C2 is fainter and is blueshifted with respect to C1 by 20-30 km s −1 . These split line components indicate low-velocity flows of gas, likely to be shells or bubbles driven by the nearby star clusters. The bright Hα knot located at the northern edge of the field does not appear distinct in the kinematics maps.

Summary of kinematics
Together with the results of Monreal-Ibero et al. (2010), our results paint a picture of the overall gas dynamics in NGC 5253 consisting of localised gas flows, as part of multiple overlapping bubbles and filaments driven by the star clusters throughout the starburst. Close to the supernebula, we have been able to identify faint (∼5 percent of the total flux) components with FWHM>250 km s −1 , although over most of positions 1 and 2, the broad component line widths are consistent with those found in the aforementioned studies (FWHM∼100-150 km s −1 , comprising 10-50 percent of the total line flux Monreal-Ibero et al. 2010;Sidoli 2010).
We confirm the presence of a strong velocity gradient in the central region seen by Monreal-Ibero et al. (2010), extending over ∼4 . 5 (∼80 pc) oriented NW-SE direction. The p-v diagram extracted along the maximum gradient shows that both the broad and secondary narrow Hα components (C2 and C3) both trace this steep gradient, which crosses zero velocity at the location of the supernebula. In line with Monreal-Ibero et al. (2010), we interpret this as an ionized gas outflow from the supernebula clusters. Although the clusters are known to be deeply embedded (A V =7-17 mag; Alonso-Herrero et al. 2004;Vanzi & Sauvage 2004), this implies that the outflow has already carved channels in the dense gas distribution through which the ionized material is escaping. Furthermore, the fact that we see a continuous velocity gradient in both the blueshifted and the redshifted side implies that the extinction in the supernebula must fall off rapidly such that both the front-and back-sides of the outflow are visible in the optical. In our IFU position 3, 6-10 to the south, the broad component contributes <15 percent, and by position 4 is not detected at all. This suggests that its presence is connected to the concentration of young clusters in the supernebula region, and most likely to the aforementioned outflow. However, since we detect broad emission >150 pc from this cluster it is unlikely to be solely due to its winds. It is possible that other young clusters are blowing winds that produce this broad line. Since the broad lines in the central regions follow the bulk motions of the more quiescent gas, whatever causes this turbulence must be a local phenomenon. Westmoquette et al. (2007) detected a very similar broad line component to the ionized gas emission lines in NGC 1569. Here we found the broadest emission correlated with the location of gas clumps in the ISM, and hence concluded that it traces turbulent mixing layers (TMLs; Binette et al. 2009;Slavin et al. 1993) on the clump surfaces set up by the impact of stellar/cluster winds. In Westmoquette et al. (2010) we tested this hypothesis by observing a gas clump near to the Galactic OB association Pismis24 (NGC 6357) at high spatial and spectral resolution. That we found broad emission originating from the clump surface in projection confirmed the idea that TMLs can explain this kind of broad component emission.

Chemical abundances
Gas phase abundances in NGC 5253 have been of particular interest since the connection between N-enriched regions (Campbell et al. 1986;Walsh & Roy 1989) and the presence of WR stars (WN-types in particular) was made (Schaerer et al. 1997). However, it is only recently that an unbiassed mapping of the abundances could be made using integral field spectroscopy (Monreal-Ibero et al. 2010. Here, they The black outlines show the 6 regions from which we extracted summed spectra for our abundance analysis. studied a ∼9×5 arcsec area encompassing the central regions. By mapping out the gas phase abundances at a resolution of ∼0 . 8, they found two locations with enhanced N/O: one significant enhancement associated with the supernebula region and another much smaller enhancement associated with two clusters ∼7 to the south-west (clusters H04-3 and H04-5).

Definition of regions
Our integral field observations provide the opportunity to investigate the gas phase abundances much further from the supernebula. However, the ionized gas surface brightness falls off rapidly with distance, and our observations are not of high enough S/N to derive spaxel-to-spaxel abundance maps as was done in Monreal-Ibero et al. (2012). Thus we have used the Hα and WR flux maps to define precise apertures from which we extract integrated spectra and perform our analysis. This already represents a significant advantage over long-slit/fixed aperture studies, since positioning uncertainties and slit/aperture losses are avoided, and the measurements can be viewed in context since the surrounding environment is known. The location of our eight regions is shown in Fig. 9.
Region 1 (∼7 . 4 or 135 pc from the supernebula) corresponds to cluster H04-17 and to region He ii 4 from (Monreal-Ibero et al. 2010). The full spectrum from this region is shown in Fig. 2. Region 2 corresponds to an Hα knot 1 . 8 to the west; region 3 contains WR star signatures and likely corresponds to a cluster visible on the HST image that was not catalogued by Harris et al. (2004, see Figs. 3 and 4); region 4 exhibits WR signatures, but no strong Hα and there is no obvious cluster at this location in the HST imaging; region 5 corresponds to an Hα knot and most likely to a cluster not catalogued by Harris et al. (2004); and region 6 corresponds to the Hα knot associated with an arc or bubble. We also defined two larger boxes covering regions of fainter, diffuse gas, d_3 and d_4, in order to sample the abundances away from the bright H ii knots.

Calculation of abundances
Abundances were derived using neat (Nebular Empirical Analysis Tool; Wesson et al. 2012) 4 , a code for calculating chemical abundances in photoionized nebulae. The code carries out a standard analysis of lists of emission lines using long-established techniques to estimate the level of interstellar extinction, calculate representative temperatures and densities, compute ionic abundances from both collisionally excited lines and recombination lines, and finally estimate total elemental abundances using an ionization correction scheme (Kingsburgh & Barlow 1994). Additionally, neat uses a Monte Carlo technique to robustly propagate uncertainties from line flux measurements and ionization correction factors (ICFs) through to the derived abundances. We note here that this more accurate treatment of the uncertainties produces reliable errors on the calculated abundances, and as such they may be larger than what can be found for equivalent measurements in the literature.
We measured the integrated line fluxes of all the major emission lines in each of the extracted integrated spectra and list them in Table 3. We made no attempt to decompose each emission line into individual components for this part of the analysis. Using these values as inputs into neat we derived the electron temperature, T e , the electron density, N e , and ionic and elemental abundances for all six regions. The code was run with ∼10 000 iterations, a number sufficient to establish a stable uncertainty distribution (Wesson et al. 2012). T e 's were derived from the temperature-sensitive [N ii]λ5755/λ6584 line ratio, whilst N e 's were derived from the [S ii]λ6716/λ6731 line ratio. For Region 4, where [N ii]λ5755 was undetected, we adopted the average T e of Regions 5 and 6, 9700 K. Since the [O ii]λ3727 line was not within our wavelength coverage we adopted measurements from the nearest corresponding regions from Walsh & Roy (1989). For our regions 1-3 we used the value I([O ii])/I(Hβ)=3.083±0.034 (corresponding to 'region 2' of Walsh & Roy 1989), whilst for our regions 4-6 we adopted I([O ii])/I(Hβ)=2.929±0.031 ('region 7' of Walsh & Roy 1989). As mentioned above, neat adopts the ionisation correction factors (ICFs) listed by Kingsburgh & Barlow (1994) to convert from ionic to elemental abundances, and propagates the statistical uncertainties associated with the ICFs accordingly. Derived O, N, S, and He nebular abundances and their uncertainties are given in Table 4.
Our derived O abundances are, within the uncertainties, consistent with a metallicity of ∼0.2 Z , as has been found in other regions of the galaxy (for a discussion see below). One point lies significantly above the others (for Region 5), but has large uncertainties that derive from the low S/N of the [N ii]λ5755 line in this region. The N abundances suggest a slightly lower metallicity. The N abundance in these outer regions is significantly lower than that found in the supernebula region, which is enhanced to ∼0.4 Z (Monreal-Ibero et al. 2012). Our He/H values are consistent with what was found by López-Sánchez et al. (2007).
It is interesting to note how similar the spectra of the diffuse regions are to the emission knots/H ii regions. This implies that a significant fraction of the ionizing photons are leaking out of the gas clouds around the clusters to ionize the intra-[star]cluster medium.

Radial abundance patterns
We have compiled all the O/H, N/H and N/O abundance measurements for different regions found in the literature, and, together with our new measurements, listed them in Table 5.4. Fig. 10 shows these values plotted versus distance from the supernebula (shown in Fig. 1 with a white cross). Note that one point at a particular radius is not necessarily associated with another point at the same/similar radius. Also included are the individual spaxel values from Monreal-Ibero et al. (2012, their fig. 8). As mentioned above, (within the uncertainties) the O/H distribution is consistent with a metallicity of ∼0.2 Z . Furthermore, we can see that it does not show any systematic trend (increasing or decreasing) within the 250 pc area sampled. The N/H profile, however, shows the well-known enhancement of ∼0.5 dex in log(N/H) within the central 50 pc. The large scatter of points within the rise result from the fact that the Nenhanced region is not circular in shape (Monreal-Ibero et al. 2012). At larger radii, both our measurements and those of Walsh & Roy (1989) 5 are systematically lower than those measured by Monreal-Ibero et al. (2012), possibly due to the systematic uncertainties introduced by the different temperature-and density-sensitive lines used. However, given these uncertainties, the plot is consistent with a flat N abundance beyond the supernebular region (i.e. R>50 pc). The N/O distribution also reflects the N/H enhancement in the supernebula region, and is also consistent with being ∼flat beyond 50 pc.
Over the years a number of low-mass dwarf galaxies and giant H ii regions have been the subject of spectroscopic chemical abundance mapping (e.g. Devost et al. 1997;Gonzalez-Delgado et al. 1994;Izotov & Thuan 1999;Kobulnicky & Skillman 1996Russell & Dopita 1990), and more recently with IFUs (e.g. James et al. 2010James et al. , 2012aJames et al. , 2009Lagos et al. 2012;Pérez-Montero et al. 2011). With very few exceptions (see below), these systems do not show significant internal chemical fluctuations or gradients, even though the expected yield of heavy elements produced by massive stars in young starbursts is substantial.  put forward two explanations: either ejecta from stellar winds and supernovae are transported throughout the galaxy and mixed instantaneously and uniformly, or freshly synthesised elements remain unmixed with the surrounding interstellar medium and reside in a hard-to-observe hot 10 6 K phase (or a cold molecular phase). In support of the latter scenario, Martin et al. (2002) find evidence for a metal enriched diffuse X-ray phase in the halo of the starbursting dwarf galaxy NGC 1569. Measurements like this, however, are extremely hard to make with existing X-ray technology, and there are no equivalent X-ray measurements available for NGC 5253.
NGC 5253 is one of the most well-known (and certainly nearest) examples amongst the rare cases of dwarf galaxies with significant chemical variations. These include Mrk996 (James et al. 2009), UM448 (James et al. 2012a) and Haro11 (James et al. 2012b), possibly IC10 (López-Sánchez et al. 2011) and IIZw70 (Kehrig et al. 2008, and those with radial oxygen abundance gradients (Lee et al. 2006;Werk et al. 2010). Of these, it is only NGC 5253 and Mrk996 that show a genuine N/H and N/O enhancement (a high N/O ratio by itsself can result from the dilution of the O abundance by an inflow of metal-poor gas). However, although they both show these enrichments, the elevated N/H ratios in Mrk996 are only found in the broad emission-line component (FWHM>300 km s −1 ), and the associated gas densities are even higher (10 6 -10 7 cm −3 ). So what is it in NGC 5253 that puts it in this extremely rare class of galaxies?

What is the origin of the N/O enhancement in the supernebula region?
Putting together the available facts we know that: -The radio supernebula (Turner & Beck 2004) 5) and pv diagram (Fig. 8) show that there is a large velocity gradient along the same NW-SE direction traced by both the broad and secondary narrow Hα components (C2 and C3).
Our new data support the conclusion of Monreal-Ibero et al. (2010) that this is an ionized gas outflow, with the dynamical centre implying that the driving source(s) are the supernebula clusters (despite being so heavily embedded).
- Beck et al. (2012) observed [S iv] 10.5 μm emission from the supernebula, and found a similar line profile to Hα, comprising a broad and narrow component. The broad component was ∼20 km s −1 blueshifted and interpreted as as indicating a gas outflow. The peak velocity of the [S iv] agrees with the Brackett lines (Turner et al. 2003) and with Hα (this work; although the Hα broad component is broader and contributes less to the total line flux than [S iv]). Furthermore, spatially resolved observations of the radio-recombination line H53α over the same region show the same velocity gradient mentioned above (Rodríguez-Rico et al. 2007). These results suggest that Hα, Brα, H53α and [S iv] are all tracing ionized material with approximately the same dynamics, despite originating at different depths within the nebula.
Since star clusters have been forming throughout the central regions of NGC 5253 for last ∼10 Myr (Harris et al. 2004), it is possible that previous generations have been responsible for the build-up in chemical abundances. However, the fact that the supernebula is located approximately in the centre of the chemically enriched region, and there is no evidence for older, sufficiently massive clusters in the vicinity strongly implies that the supernebula cluster(s) must be the source.
In line with above arguments regarding chemical fluctuations in dwarf galaxies, we propose that in NGC 5253, the chemically processed material expelled from star clusters in their early evolution remains in the hot wind phase and does not mix with cooler gas. However, the unusually high densities and pressures in the supernebula and surrounding medium have acted to stall or impede the winds from these embedded clusters, allowing time for the chemically enriched hot phase to mix with the cool/warm phases. Cluster winds stalled by high ambient pressures have been previously found in M82 (Silich et al. 2007;Smith et al. 2006).
Are the supernebula clusters old enough to be the enrichment source? The equivalent width of Paα constrains the age of the clusters to <6 Myr (Alonso- Herrero et al. 2004;Martín-Hernández et al. 2005). That the radio and millimetre continuum emission are almost entirely due to thermal emission (Beck et al. 1996;Turner et al. 2000) suggests that no supernovae have exploded in the clusters, implying an upper age limit of ∼4 Myr. Martín-Hernández et al. (2005) modelled the mid-IR spectra of the embedded SSCs and found two solutions consistent with the observations: 1) a young (<4 Myr) cluster with a "nonstandard", low, upper mass cutoff of the IMF; or 2) a cluster of ∼5-6 Myr with a "standard" high upper mass cutoff. The presence of N-enrichment strongly argues for the influence of WR stars, whose appearance in a cluster peaks between ∼3-5 Myr (Crowther 2007). As Martín-Hernández et al. (2005) comments, one expects a supernova rate of at least 100 SN/Myr after approximately 3 Myr, but no effects of supernovae are seen in the radio. We therefore suggest that the clusters must be ∼4 Myr old -old enough to have (had) a significant WR star population, but young enough to have had none or very few SNe.
That the thermal radio emission from the supernebula is optically thick classifies it as an "ultra-dense" H ii region with pressures significantly higher than typical interstellar values (Turner et al. 1998). Under normal circumstances this would imply that the clusters are extremely young ( 1 Myr) from pressure/lifetime arguments (cf. Kobulnicky & Johnson 1999;Vacca et al. 2002). However, we now know that the clusters exist in a much larger high density/pressure environment, helping to explain why they remain so embedded even at ages of a few Myr (cf. Smith et al. 2006). The elongated morphology of the enriched region is spatially coincident with the extent and orientation of the outflow, implying that this has acted to increase the efficiency of either the mixing (to bring the new chemicals into the observable warm phase), and/or the subsequent diffusion. The outflowing N-enriched material has an extent of ∼60 pc, which could have propagated there over a few Myr with diffusion speeds of 20-60 km s −1 (Monreal-Ibero et al. 2010). This is in good agreement with the the outflow speeds measured above.
Further observations are required to explain why the densities/pressures in this supernebula region are so high. One possibility is that the diffuse H i cloud infalling along the minor axis of the galaxy found by López-Sánchez et al. (2012) may be responsible for compressing the gas at this point.
It is interesting that the situation in the NGC 5253 supernebula is remarkably reminiscent of that in the only other system with a confirmed N/H enhancement: Mrk996 (D≈22 Mpc). Here the N/H enhancement is seen in the nuclear region, spatially coincident with a significant WN star population (James et al. 2009). Although the elevated N/H ratios are only found in a broad emission-line component (FWHM>300 km s −1 ), the associated gas densities are also extremely high (10 6 -10 7 cm −3 ).

Conclusions and Summary
We obtained Gemini-S/GMOS-IFU observations of four regions near the centre of the nearby dwarf starburst galaxy NGC 5253. We identify 11 distinct WCE-type stars through mapping out the strength and location of the red C iv WR bump. Some are associated with known young clusters catalogued by Harris et al. IFU Table 4. Abundance measurements for the six defined regions within IFU positions 3 and 4. Also given are the equivalent widths of Hβ (W eq (Hβ), and the integrated red WR bump luminosities that provide an estimate of the equivalent number of WCE stars.  2012) individual spaxels, red triangles: this work. Physical distances have been calculated assuming a distance of 3.8 Mpc (Sakai et al. 2004).