The Recent Star Formation in NGC 6822: a UV Study

The Recent Star Formation in NGC 6822: an Ultraviolet Study

Boryana V.Efremova11affiliation: Department of Phys.& Astron., Johns Hopkins University, 3400 N.Charles St., Baltimore, MD 21218 (boryana, , Luciana Bianchi11affiliation: Department of Phys.& Astron., Johns Hopkins University, 3400 N.Charles St., Baltimore, MD 21218 (boryana, , David A.Thilker11affiliation: Department of Phys.& Astron., Johns Hopkins University, 3400 N.Charles St., Baltimore, MD 21218 (boryana, , James D.Neill22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Denis Burgarella33affiliation: Laboratoire d’Astrophysique de Marseille, BP 8, Traverse du Siphon, 13376 Marseille Cedex 12, France , Ted K.Wyder22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Barry F.Madore44affiliation: Observatories of the Carnegie Institution of Washington, 813 Santa Barbara Street, Pasadena, CA 91101 , Soo-Chang Rey55affiliation: ”Department of Astronomy and Space Science, Chungnam National University, Daejeon 305-764, Republic of Korea , Tom A.Barlow22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA ,Tim Conrow22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Karl Forster22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Peter G.Friedman22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , D.Christopher Martin22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Patrick Morrissey22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA , Susan G.Neff66affiliation: Laboratory for Astronomy and Solar Physics, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA , David Schiminovich77affiliation: Department of Astronomy, Columbia University, New York, NY 10027 , Mark Seibert44affiliation: Observatories of the Carnegie Institution of Washington, 813 Santa Barbara Street, Pasadena, CA 91101 , Todd Small22affiliation: California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA

We characterize the star formation in the low metallicity galaxy NGC 6822 over the past few hundred million years, using far-UV (FUV,  Å) and near-UV (NUV,  Å) imaging, and ground-based H imaging. From FUV image, we define 77 star-forming (SF) regions with area  pc, and surface brightness 26.8 mag (AB) arcsec, within 0.2 (1.7kpc) of the center of the galaxy. We estimate the extinction by interstellar dust in each SF region from resolved photometry of the hot stars it contains:  ranges from the minimum foreground value of 0.22 mag up to 0.660.21 mag. The integrated FUV and NUV photometry, compared with stellar population models, yields ages of the SF complexes up to a few hundred Myr, and masses from  to . The derived ages and masses strongly depend on the assumed type of interstellar selective extinction, which we find to vary across the galaxy. The total mass of the FUV-defined SF regions translates into an average star formation rate (SFR) of  yr over the past 100 Myr, and SFR = yr in the most recent 10 Myr. The latter is in agreement with the value that we derive from the H luminosity, SFR =0.008  yr. The SFR in the most recent epoch becomes higher if we add the SFR=0.02 yr inferred from far-IR measurements, which trace star formation still embedded in dust (age a few Myr).

galaxies: individual (NGC 6822) – galaxies: stellar content – Local Group – stars: formation – ultraviolet: stars

1 Introduction

Continuum fluxes in the ultraviolet (UV) and infrared (IR) spectral regions, and H line emission, are the main indicators of star-formation activity in distant galaxies (see e.g. Kennicutt 1998). The UV flux is a direct tracer of young massive stars, whose energy is mostly emitted in this spectral region, H emission originates from interstellar gas ionized by the most massive stars, and the far-IR emission is produced by dust particles re-emitting reprocessed UV stellar light.

Integrated measurements of these fluxes can be translated into star-formation rates of galaxies, but additional information is needed. First, observed fluxes need to be corrected for extinction by interstellar dust, both foreground (by Milky Way (MW) dust along the line of sight), and internal (within the galaxy). Reddening is particularly significant at UV wavelengths (see e.g. Bianchi 2011). Stellar population models with adequate star-formation history (SFH) are then used to transform the continuum and line-emission luminosities into SFRs.

The UV photometry of star-forming galaxies is usually corrected for interstellar extinction assuming a MW-type selective extinction with (Cardelli et al., 1989) for the foreground component, and the Calzetti (2001) extinction curve for internal extinction. The amount of extinction is sometimes estimated by comparison of UV and far-IR fluxes (Calzetti et al., 2005; Cortese et al., 2006; Boissier et al., 2007; Meurer et al., 2009). Such method assumes that the intrinsic FUV-NUV color is known, however its value is strongly varying with age for young starbursts (e.g. Bianchi 2009, 2011), and that UV and far-IR fluxes are emitted by the same population, which is often not the case.

In unresolved distant galaxies only integrated measurements are possible, and a global extinction correction and star-formation history (SFH) must be assumed for interpreting such measurements. On the other hand, in nearby galaxies individual SF regions can be measured, and their stellar content studied in detail (e.g. Bianchi & Efremova 2006, Bianchi et al. 2011, 2010, 2001, Kang et al. 2009, and references therein). Therefore, the dust properties can be explored in a variety of local environments, providing information on the interplay of dust and star formation, and a calibration of star-formation indicators in distant galaxies.

Deep imaging in FUV and NUV for hundreds of nearby galaxies were obtained with the Galaxy Evolution Explorer () (Martin et al., 2005; Morrissey et al., 2007) as part of the Nearby Galaxy Survey (NGS) (Bianchi et al., 2003a, b; Bianchi, 2009; Gil de Paz et al., 2007). The wide-field UV imaging provides a characterization of the young stellar populations across the whole extent of these galaxies, and can be used, with complementary optical data, to infer their star-formation history and SFR.

In this paper we perform a comprehensive study of the young stellar populations in the Local Group low-metallicity galaxy NGC 6822, the nearest SF galaxy currently with no massive neighbor. We identify and define SF regions from wide-field imaging in FUV, where the hottest, youngest stars are more prominent, throughout the extent of the galaxy. We use integrated photometry of these regions in FUV and NUV, and complementary H emission-line imaging, as well as information from resolved stellar photometry, to investigate the star-formation in this galaxy during the past  Myr, and the characteristics of interstellar extinction. The study of this galaxy, together with results for Local Group galaxies of other types, contributes one piece to a broader puzzle, aimed at understanding the modalities of star formation in differing environments, and the role of dust.

This benchmark galaxy was chosen to complement the study by Kang et al. (2009) of M31, and of other Local Group galaxies by Bianchi et al. (2010, 2011), because of its low metallicity and vicinity ( kpc, McAlary et al. 1983) and the abundant information available from resolved stellar population studies with multi-band imaging (Bianchi et al. 2001, Bianchi & Efremova 2006), CTIO imaging (Massey et al. 2007a), imaging and extensive spectroscopy (B. Efremova et al. 2011, in preparation). NGC 6822’s metallicity is believed to be subsolar: measurements by Muschielok et al. (1999) of three B–type supergiants, and by Venn et al. (2001) of two A–type supergiants both yield .

The paper is arranged as follows. In Section 2 we define SF regions from FUV imaging, and measure their integrated fluxes in FUV and NUV; we also use the CTIO H imaging of Massey et al. (2007b) to define and measure regions of H emission. In Section 3 the integrated measurements of the SF regions are analyzed with stellar population models to derive their ages and masses, after the interstellar extinction is estimated from the massive stars within each SF region. The results are discussed in Section 4, and summarized in Section 5.

2 Observations. Detection and Photometry of the Star-Forming Regions

2.1 UV imaging

We used images in FUV ( Å, FWHM  Å), and NUV ( Å, FWHM  Å) with resolution 4.2 (FUV) and 5.3 (NUV) (Morrissey et al., 2007), corresponding to  pc at the distance of NGC 6822. The images are sampled with 1.5 pixels.

The images of NGC 6822 were taken on Aug 20th 2005 as part of the NGS program, with exposure times of 4654 sec (FUV) and 6198 sec (NUV). The data was downloaded from the MAST archive. The 1.2 degree diameter field of view is centered at RA , Dec , near the center of the galaxy (RA , Dec , FK5 2000). NGC 6822, with an optical diameter (at mag/arcsec) of 15′.6 (Karachentsev et al., 2004), is contained in the central portion of the image, which is shown in Fig. 1.

2.2 UV Source Detection

We identified the SF regions using the FUV image, which unambiguously reveals the young, hot massive stars not heavily embedded in interstellar dust. We followed the general method of Kang et al. (2009), adapted to the case of NGC 6822. We defined contours of regions with FUV surface brightness above the background. An important issue in defining extended source contours and measuring their flux is the background estimate. Several approaches were used to find the best method for background evaluation (see also the discussion in Kang et al. 2009). For the purpose of source detection only, we constructed a background image applying a two-step circular median filter (64 pixels diameter, 1.5arcmin) to the FUV intensity map (“int” file). The first pass of the filter identifies pixels which belong to localized peaks via masking pixels brighter than the local median background estimate. The second pass of the filter operates only on the final list of non-peak pixels to obtain a background image less biased by substructure than a one-pass median filter. The diameter of the median filter was chosen to provide a background image where measurements of the background for individual sources are closest to the median flux density of the intensity map images in 6-pixels wide annuli around the sources, from here on “local background”. The background image produced using the adopted median filter gives sky estimates slightly lower (by about 0.17/0.18 mag arcsec for FUV/NUV) than the local background. For comparison, the background image provided by the pipeline gives a background estimate lower than the local background by 0.52/0.60 mag arcsec on average for FUV/NUV. A background-subtracted image was constructed, subtracting our background image from the intensity map image, and used for source detection.

Source contours were defined to enclose contiguous pixels with FUV flux more than above the background image. This threshold corresponds to an FUV surface brightness of 0.0015 counts s pixel (26.8 AB mag arcsec) on the background-subtracted image, or average 0.0025 counts s pixel (26.2 AB mag arcsec) on the intensity map image.

The effect of the threshold choice on the source-contour definition is illustrated in the top panels of Figure 1, which show SF regions # 75, 57, 27, and 20 defined for thresholds of (green), (light blue), and (dark blue). A low threshold of 1 or would cause regions like # 27 and 20 (OB8 and OB6, from here on ‘OB’ designations are from Hodge 1977) to merge, and the main body of the galaxy to appear as one large region. If a threshold higher than is used, sparse associations like region # 75 (OB15) split into several sources or into individual stars (see also Kang et al. 2009 for more discussion on the procedure and the choice of parameters).

The sources defined using the threshold follow the distribution of blue stars as shown by resolved stellar photometry from imaging (Bianchi et al 2001, Bianchi & Efremova 2006), and ground-based data (Massey et al 2007a). To exclude artifacts, in the initial list we rejected sources with area less than 16 arcsec (7 pixels).

Figure 1: Contours of SF regions (blue) over the FUV image, defined with FUV surface brightness above the background level and area  arcsec ( pc). H contours () are shown in red. Enlargements of regions # 75, 57 and 27 and 20 in the top panels illustrate the effect of the threshold choice for contour definition ( in green, in light blue, in dark blue).

We further restricted the analysis sample to sources larger than 150 arcsec ( pc), in order to exclude single stars (mostly foreground) and background objects, and to examine SF complexes massive enough that stochastic effects will not be significant in deriving ages and masses by model analysis (Section 3). Stochasticity may affect the comparison of integrated star cluster photometry with stellar population models, as was first pointed out by Girardi et al. (1995). Quantitative assessment of this effect is still a matter of debate, given that more factors are relevant in such analysis, including IMF, metallicity, extinction. For example, Fatuzzo & Adams (2008) estimate that for clusters with 1000 stars the IMF is sampled well enough so that their UV flux is close to model predictions for integrated populations (but they also point out that the exact limit may vary with IFM). Their analysis concerns statistical distributions of bound, spherical, zero-age stellar clusters. Such limit corresponds to 5 stars more massive than 10  i.e. earlier than spectral type B2V (with the parameters adopted by these authors). We will return on this point again later.

We choose to restrict the analysis sample with the area cut of 150 arcsec after examining the distribution of stars with spectral type earlier than B2V111Selected from the photometry of  Massey et al. (2007a) to have and , after the reddening correction is applied, as described in Section 3.1. inside the FUV-defined contours. In Fig. 2 we plot the FUV magnitude vs area of the SF regions: those containing 5 blue massive stars are shown with dots. A cut by area of  arcsec includes 94% of these regions in the sample, and very few regions containing less than five blue stars (20%).

Figure 2: FUV magnitude vs. area of FUV-defined SF regions. Regions containing five or more blue stars (see the text) are marked with filled circles. A cut by area at 150 arcsec, adopted for our analysis sample, is shown with a vertical line. It retains 94% of the regions containing 5 blue stars in the sample. A brightness cut at  mag would retain the the same fraction of regions with 5 blue stars, however it would highly increase the fraction of regions with less than five stars included in the sample.

We examined the alternative option of a brightness cut, which is often used in studies of more distant galaxies. Such criterion would either include fewer SF regions with five or more blue stars, or more regions with less than five blue stars, in the analysis sample. For example, a brightness cut at  mag includes in the analysis sample 94% of the regions with 5 blue stars (the same fraction as our area cut of 150 arcsec), but 44% of the selected regions would contain less than five blue stars. Therefore, a cut by area better satisfies our requirement of a minimum number of blue stars within a source contour, including in the analysis sample as many as possible of the clusters having 5 massive stars, and as few as possible clusters with 5 massive stars. A brightness or luminosity cut would also strongly be affected by extinction, or extinction correction (see Section 3.2). We point out that this criterion may not necessarily be the best choice for more distant galaxies where similar data would give a lower spatial resolution, or for galaxies where star formation is less sparse. In the specific case of NGC6822, such criterion, more precise than a luminosity cut for our purpose, could be tested and tuned given the vicinity of the galaxy and the detailed information on its stellar population. Finally, a cut by area may eliminate young compact clusters, and may be undesirable in disk galaxies for example, where young compact star clusters abound (e.g. Bianchi et al. 1999, Chandar et al. 1999 for M33; Hodge et al 2010, and Kang et al in preparation, for M31). In NGC6822 there are very few such compact clusters and their exclusion would not change our results. This work aims at the detection of unbound OB associations and SF complexes, not compact star clusters. Another advantage of the area cut is that it effectively excludes foreground stars.

The resulting analysis sample includes 77 FUV-defined sources with area 150 arcsec and brightness 26.8mag arcsec, within a 0.2 deg radius (1.72 kpc) of the center of NGC 6822. The 0.2 deg radius is 1.5 times the optical semi-major axis of the galaxy (7′.8 at mag/arcsec) given by Karachentsev et al. (2004). H emission is detected out to a radius of R kpc (Hunter & Elmegreen, 2004) (see also Section 2.3), and the HI disk (de Blok & Walter, 2006) also exceeds the optical size of the galaxy (see also Bianchi 2011). Our sample extends to a slightly larger area than that of Melena et al. (2009)222 Melena et al. (2009) sample of SF regions is within 1.65 kpc, using the coordinates of their Table 2, in spite of their claim that it extends to 6 kpc..

The areas of the selected SF regions range from 150 to 5400 arcsec ( pc). Table 1 gives identification, coordinates of the “centroids” (the median  and  values of the pixels included in the contours), and areas of the FUV-defined regions, ordered by increasing R.A. The contours are shown in blue in Figure 1 over the FUV image. In the next section we describe the photometry measurements, which are used in Section 3 to derive ages and masses.

2.3 UV Photometry of the star-forming regions

For photometric measurements we used the intensity map (“int”) images (in units of counts s pixel) generated by the pipeline dividing the count map by the relative response map333 ERO_data_description_3.htm. We measured the FUV and NUV flux of each SF region within its FUV-defined contour, and the local background. The background was measured over an area defined by smoothing the source contour and expanding it by 3 pixels (inner background contour) and 9 pixels (outer contour), i.e. creating a 6 pixels wide ‘annulus’ around the source, which follows its shape. The median of the flux pixel in the background region, excluding portions of nearby sources falling in the background annulus, was then subtracted from every pixel inside the source contour. The conversion from [counts   s] to magnitudes in the AB photometric system was performed using zero-points ZP mag (FUV) and mag (NUV) (Morrissey et al., 2007). We calculated the photometric errors as = , where is the flux from the source in the aperture and is the quadratic sum of all the noises affecting the image. is expressed by , where is the flux in counts, and is the conversion factor from ADU to e, for . We consider the Poisson noise of the photon flux, and the background fluctuations to dominate, so we used the following expression to estimate the noise: , where is the area of the source (in pixels) and is the standard deviation among the pixels in the background annulus (in counts). The resulting FUV and NUV magnitudes and their errors are listed in Table 1.

The total flux from the selected SF regions is 45% of the integrated FUV flux from NGC 6822 ( erg s cm or mag AB) and 35% of the total NUV flux ( erg s cm or mag AB), measured from the pipeline sky-subtracted image in an aperture of 0.2 deg radius. The flux not included in our SF-regions comes from smaller sources excluded by our area cut (10% of ), from older diffuse populations, and from scattered emission from SF regions. The fraction of flux included in the selected SF-sites is lower in NUV than in FUV because foreground stars and diffuse light from older populations are more conspicuous at longer wavelengths.

2.4 H Emission Sources

We also used the publicly available CTIO H image from the survey of Massey et al. (2007b) to define contours of H emitting regions. The H image has an exposure of 300 sec, a scale of 0.27 pixel, and a resolution of 0.9 (2.2 pc at the distance to NGC 6822). We used the and images from the same survey (Massey et al., 2007a) to correct the H image for continuum, by subtracting from it a linear combination of the and images, scaled to match the intensity of the continuum sources. We define contours of H emitting regions using a threshold of above the background, corresponding to a surface brightness of erg scm arcsec. The H-defined contours are drawn in red in Fig. 1. They generally follow the H contours defined by Hodge et al. (1988, 1989) in a similar way and using a threshold of erg scm arcsec.

We used the calibration factor of 1 count s=ergs scm given by Massey et al. (2007b) for emission line sources. We did not attempt to correct for [N ii] emission line contamination, which we expect to not exceed a few percent of the flux. The average / in the H ii regions measured by Pagel et al. (1980) is about 6%; adopting / 1/3, we derive / 8%. Since both [N ii] 6548, 6584 fall in the wings of the 50Å -wide filter passband centered at H, we expect the actual contribution from [N ii] emission to be of the order of %.

The H measurements are listed in Table 2. In the last column we give the cross-identifications with H sources previously defined by Hubble (1925) (H followed by roman number), Hodge (1977) (Ho followed by arabic number), Kinman et al. (1979) (K followed by a Greek letter), Killen & Dufour (1982) (KD followed by arabic number), and Hodge et al. (1988) (HK followed by arabic number). For regions with H surface brightness ergs scm arcsec, our measurements agree within with those by Hodge et al. (1989), after de-correcting the latter for extinction with A (O’Dell et al 1999). The total H flux from the H emitting regions,  erg s cm (the sum of all measurements in Table 2), is higher by about 2% than the total flux from Table 1 of Hodge et al. (1989), most of the difference coming from faint H ii regions not covered by the imaging of Hodge et al. (1988). We further checked our H calibration and continuum-source subtraction by measuring the flux from the H ii regions Hubble V and Hubble X surrounding OB8 and OB13 in 42 square apertures, similar to the ones used by O’Dell et al. (1999): our measurements agree within 10% with the values given by these authors.

3 Analysis

The UV color-magnitude diagram of the FUV-selected SF regions is shown in Figure 3: in the left panel we plotted the FUV surface brightness (mag arcsec), and in the right panel the total FUV magnitude. In the right panel we also show FUV synthetic magnitudes for single-burst (coeval) populations (Bianchi 2011) at various ages, scaled for total stellar masses of  //  (solid/dashed/dash-dotted lines). Models are reddened with mag (the assumed foreground extinction) using MW-type extinction (green lines) and with an additional using the extinction curve derived for stars in the LMC2 star-forming region by Misselt et al. (1999) (blue lines).

Figure 3: UV color-magnitude diagram of the SF regions (left: FUV surface brightness, right: FUV magnitudes). The sources with associated H emission are marked with dots, the others with diamonds (filled if the background is 30% of the total flux within the source contour). The right panel shows also synthetic SSP model magnitudes at different ages, scaled to cluster masses of // (solid/dashed/dash-dotted lines). Models are reddened with MW-type exinction for a foreground reddening of (green lines), and with an additional using LMC2-type extinction (blue lines). Three age values (10, 30, and 100 Myr) are marked with triangles and labelled.

For SF regions located in the main body of the galaxy, the local background (which includes the diffuse older populations, more conspicuous in the NUV band) is significant, and therefore there is always a concern that its subtraction may lead to greater uncertainty than the formal errors indicate. An error in FUV-NUV color would propagate to an error in the derived age, and consequently on the derived mass. To verify that no bias is introduced by high background subtractions, we plotted with empty diamonds the sources for which the background amounts to more than 30% of the flux in the source contour. These high-background sources are distributed in color-magnitude space not differently from the other sources, confirming that no biases have been introduced by the critical background estimate procedure.

In the following sections we estimate ages and masses of the SF regions by comparing their UV photometry with Simple (single-burst) Stellar Population (SSP) model colors of different metallicities (e.g. Bianchi 2007, 2009, 2011), after the  is estimated for each SF region from resolved stellar photometry. The models are progressively reddened with various types of interstellar dust extinction.

3.1 Interstellar Reddening

Flux at UV wavelengths is very sensitive to extinction by interstellar dust, and in order to derive age and mass of the SF regions from model analysis, we first estimated the amount of reddening in each.

We used information from resolved stellar photometry, and estimated  of the hot massive stars (selected with and , i.e. earlier than about  B2V) within each contour. For several OB associations,  values are available for individual stars, derived by Bianchi et al (2001) and Bianchi & Efremova (2006) from multi-band photometry (from UV to optical). For the most massive stars in six OB associations, we also have spectroscopy (B. Efremova et al. 2011, in preparation), which confirms the results from photometry. For the regions without photometry, we derived  using the photometry of Massey et al. (2007a), with the standard “Q-method” (e.g. Bianchi & Efremova 2006, Kang et al. 2009), and by comparing the observed (),() colors with progressively reddened stellar model colors.

We accounted for extinction in each SF region using the median  value of the massive stars it contains. The values range from (purely foreground extinction) to 0.66 mag, with a mean of  mag, and are given in Table 1. The mean  values are similarly distributed, ranging from to 0.51 mag, with an average of  mag. The typical scatter (also given in Table 1) around the mean  in individual SF regions is 0.13 mag. The wide range of extinction values in the SF regions (not uncommon in star-forming galaxies) underscores the importance of accurate extinction corrections, particularly relevant in the UV regime, for deriving ages and masses. Other works adopt a generic assumption, for example Melena et al. (2009) used a constant extinction of 0.27 mags for all their sample regions in NGC6822, corresponding to  mag of internal extinction in addition to the  mag foreground reddening. Our results derived for individual regions (Table 1) indicate that a higher value is more typical. The model magnitudes plotted in Fig.3 illustrate how such assumptions affect the derived ages and masses; more model plots showing the effects of extinction can be found in Bianchi (2011).

We assume a foreground extinction of  mag, consistent with the minimum  estimated in this work, and with previous estimates by Bianchi et al (2001), Bianchi & Efremova (2006), and Massey et al. (2007a). Any additional extinction is considered to originate within NGC6822.

While the derived  values are mostly based on optical photometry of the stars, and do not depend significantly on the type of dust, the selective extinction A/ at UV wavelengths is known to strongly vary with environment. Therefore, the correction of UV magnitudes, and the resulting ages and masses, strongly depend on the adopted extinction curve (e.g. Bianchi 2011). In the analysis that follows, we consider four different types of internal extinction, found in the MW and in known low-metallicity star-forming environments: 1) MW-type extinction with . In this case the (FUV-NUV) color is basically reddening-free (Bianchi 2011 and references therein); 2) the average extinction curve derived by Misselt et al. (1999) for LMC stars outside the 30 Doradus region (from here on, “LMC”), which gives an average color excess ratio for hot stars ( K) of ; 3) the extinction curve in the LMC 30 Doradus region (from here on “LMC2”) derived by Misselt et al. (1999), yielding ; and 4) the extremely UV-steep extinction curve derived by Gordon & Clayton (1998) for SMC stars (from here on, “SMC”), which yields .

3.2 Ages and Masses of Star-Forming Regions

We derive ages of the SF regions from their integrated FUV-NUV colors compared with SSP models for low metallicity populations (see below), and masses from the age and UV magnitudes, accounting for reddening. We compared results obtained by adopting the different extinction curves mentioned in the previous section with information from resolved stellar photometry and H , in order to assess what type of selective extinction is more appropriate.

We found that a uniform extinction type is not adequate for the whole sample of SF regions in NGC 6822. If we assume “average MW” extinction with for the whole sample (as adopted e.g. by Hunter et al. 2010), the measured colors imply ages too old for several regions which show H emission and which appear to be a few Myr old in H-R diagrams from photometry (Bianchi et al 2001). On the other hand, if we use a UV-steep extinction curve, “LMC2” for example, to deredden all SF regions, the (FUV-NUV) color for part of the sample is over-corrected, such that it appears unrealistic when compared with model predictions at any age. We found that different extinction curves are needed to bring the ages from integrated measurements in agreement with results from resolved studies for a subsample of well studied regions. Regions # 27, 57, 75, 19, 20 (approximately corresponding to OB8, OB13, OB15, OB7, OB6 as defined by Hodge 1977), and # 52, are included in the photometric studies by Bianchi et al. (2001) and Bianchi & Efremova (2006); spectroscopy of the most massive stars confirms the ages derived from photometry. By comparing results from integrated measurements with resolved stellar photometry of these well studied SF regions (outside the central part of the galaxy, where measurements are not complicated by significant diffuse older populations), and with information from H emission, we derived a criterion for choosing the type of extinction curve, and apply it to the rest of the sample. Specifically, the age information for the best studied SF regions, with spectroscopy available for the hottest stars, was based on the presence (or absence) of O-type stars, W-R type stars, or B supergiants, as well as on the photometric H-R diagram of their stellar population, and we ruled out extinction curves giving very discrepant results from the color. For the FUV-bright regions clearly associated with H emission, we ruled out extinction curves yielding ages significantly older than 10Myr from integrated UV photometry. Finally, UV-steep extinction curves were excluded in the cases where they would yield an intrinsic FUV-NUV color bluer than any stellar population model at any age. While derivation of the actual extinction curve would require UV spectroscopy of several stars in each region, and it is not possible with broad-band photometry, the representative known curves examined give sufficiently different results (Table 3) that some of these assumptions could definitely be excluded in many cases. Some consistent trends within the subsample of SF regions with information on their stellar content allowed us to define general criteria.

For the youngest SF regions, a UV-steep extinction curve brings ages from integrated measurements into agreement with resolved H-R diagram results. These regions are associated with strong H emission (plotted with circles in Figs. 3 and 4) and typically have high FUV surface brightness, i.e. the SF is intense and the complex is very compact. Therefore, we assumed UV-steep extinction curves to correct for internal extinction of regions with surface brightness higher than mag arcsec, and MW-type extinction for the rest of the sample. Among the high surface-brightness SF regions, we found “LMC2”-type extinction to be preferable for sources with  mag, and the less steep “LMC” extinction to be better for sources fainter (in FUV) than this limit.

We stress that these criteria were derived ad hoc, to bring the results from integrated measurements into agreement with detailed studies of a subsample. photometry of sample regions for a wider sample of Local Group galaxies will be used to expand this comparison (Bianchi et al 2010, 2011). However, the results are not surprising: for example, in a study of SF sites in M51, Calzetti et al. (2005) found starburst-like extinction to be applicable only to sites with strongest star formation.

Two metallicity values, and , were explored, encompassing the estimated metallicity of young stars in NGC 6822, (Muschielok et al. 1999, Venn et al. 2001). Models with metallicity yield slightly younger ages, and consequently lower masses of the SF regions. The effect of metallicity on the derived age and mass varies with age and extinction type, as discussed by Kang et al. (2009; see their Fig.12), and Bianchi (2009, in particular Fig.9), Bianchi (2011).

Ages and masses of the SF regions, derived from integrated FUV, NUV photometry using models with metallicity , and assuming three different extinction curves for internal extinction, are given in Table 3; the last column indicates which values are adopted in our analysis. These values are plotted in Figure 4. The gap in the age distribution at  Myr is due to a slight degeneracy of the UV color in the range  Myr, due to RSGs emission (Fall et al. 2009). The uncertainty in the derived ages caused by this effect is smaller than the uncertainty introduced by the scatter in , and it does not affect our overall results significantly. We stress that the masses of the individual SF regions should not be interpreted in terms of cluster mass function, for two reasons. First, we defined irregular, mostly unbound, complexes. Second, we used a constant threshold throughout the galaxy, to derive source contours, in the interest of adhering to an objective criterion and a consistent flux limit. This choice inevitably may cause some sparse regions to break into subcomponents (but their ages, and the masses of each, would not be affected), and more diffuse regions to merge.

Figure 4: Masses vs ages of the FUV-defined SF regions, derived from integrated UV magnitudes using SSP models with metallicity . We accounted for interstellar extinction using  values estimated for each SF region, and assuming a foreground component with and MW-type dust (), and different extinction curves for the additional internal extinction as described in the text (black if “MW, ”, light blue if “LMC”, dark blue if “LMC2”). As in the previous figure, UV sources associated with H emission are plotted with dots, and the rest with diamonds (empty symbols if the background is 30% of the source flux). The black line shows the detection limit for a source with our minimum area and only foreground reddening.

The magnitude limit of our sample, from the 3 detection threshold of 26.8 mag  arcsec and the area cut of  arcsec, translates (using the SSP models) into a mass detection limit increasing with age, shown with a black line in Fig. 4 for a foreground reddening  mag. The actual mass limit of the sample is higher because most sources have a higher reddening. As can be expected, the detection limit causes incompleteness for low masses at old ages.

For comparison, we also estimated the masses of the FUV-selected SF regions from resolved stellar photometry. The mass of each SF region was derived by extrapolating the number of stars above 10  (corresponding to about B2V, and chosen with and ) and up to the most massive star still on the MS, to the interval . We assumed an IMF with in the range , and for  after Kroupa (2001). The masses derived from the H-R diagrams agree with those from integrated UV photometry for SF regions younger than 10 Myr. For older regions, the H-R diagrams give lower masses than the integrated measurements, by up to a factor of 20, probably because the most massive stars have evolved. The most discrepant regions have large areas and are in the main body of the galaxy, where the background is higher; they may be the result of merging of nearby regions expanding with age.

The possible contribution by foreground MW red dwarfs to the integrated flux of the SF regions was also estimated, since the density of foreground stars is significant in the direction of NGC 6822. We used our stellar model grids to estimate the possible contribution to the FUV and NUV fluxes from foreground stars of intermediate colors ( and , see for example Figure 6 of Bianchi & Efremova 2006). The derived potential effect on the FUV-NUV color is very small ( mag on average) and does not affect the results.

The main concerns to be addressed when deriving ages and masses of SF complexes by comparison with SSP model colors are: 1) the degree of coevality of the stellar complex and applicability of the SSP assumption; 2) stochastic effects from the top-IMF small number statistics. The latter affects both the analysis of integrated measurements and of resolved stellar counts. However, it affects only the small mass clusters. As we explained above, we restricted our analysis sample so to include sources with 5 massive stars, in order to minimize problems of stochasticity. More importantly, we do not interpret our results in terms of mass distributions of individual clusters; instead, we add the masses of SF regions in broad age bins (next section) in order to obtain the total stellar mass formed at different epochs. In this way, we derive information on global star formation with broad time-resolution, and stochastic effects on individual cluster masses average out. As for the assumption of “SSP” (or instantaneous star-formation of each region), we tested the results by using also models with exponential SFH, decaying over short time scales: the results showed no appreciable difference. The measured FUV-NUV color of most sources is incompatible with CSP (“continuous star-formation stellar populations”) model colors of any age, ruling out the CSP assumption often used to derive global galaxy SFR, as not applicable to the individual SF regions in our sample. Strict coevality is not observed even in bona fide globular clusters, the epitome of “single age” stellar population, therefore some degree of uncertainty is carried in all works by this assumption, which is however the most compatible with the observed properties of our young populations sample.

4 Results and Discussion

4.1 Recent star formation from UV fluxes

The SF regions have ages  Myr, as derived in the previous section from UV photometry, due to the FUV selection, and their masses range from  to , when individual extinction correction is applied to the sources, as explained in Sec. 3.2.

We added the UV-derived masses of the SF regions in four age ranges, to estimate the average SFR within these time intervals. We find: SFR  yr ( Myr), SFR  yr ( Myr), SFR  yr ( Myr), and SFR  yr over the whole range  Myr. The results are shown in Figure 5; the uncertainties, shown as gray boxes, take into account the photometric errors and the  scatter within the SF regions, which is typically one order of magnitude larger than the photometric errors. The uncertainties are up to a factor of four of the derived values. Additional (smaller) uncertainties may arise from the assumption that the stellar population in each SF region has formed in a single burst, and from the adopted IMF. Coevality is supported by the HR diagrams of seven young SF regions, studied with (Bianchi et al. 2001, Bianchi & Efremova 2006), but it may be more questionable for older larger regions, which may result from merging of subcomponents dissolving with time. Stochastic effects for low mass complexes, as discussed previously, could affect individual masses, but average out when masses of several clusters are summed in broad age bins. Moreover, the total mass is dominated by the most massive SF regions.

The effect of the extinction curve used to correct the UV color can be appreciated in Fig.5 where we also show results derived assuming the total reddening to be from MW-type dust for all sources (green lines), as was assumed e.g. by Wyder et al. (2007) and Hunter et al. (2010). The resulting SFR is significantly lower in the most recent epoch, because the most massive SF regions are shifted to older ages when we use this extinction curve.

Figure 5: The average SFR of NGC6822 over recent time intervals, derived by adding the masses of the FUV-defined SF regions within each epoch (solid lines; the dashed line shows the average over the wider range). The black lines show results obtained with reddening corrections as described in the text; gray boxes show the uncertainties, taking into account extinction spread within each region and photometric errors. Green lines show results obtained assuming MW-type () extinction for all sources. The red lines show the SFR estimated from H in this paper (solid line), and by Hunter & Elmegreen (2004) (dotted line). The blue line indicates the SFR derived from 24 µm emission, and the yellow line shows the SFR derived from resolved stellar photometry of the youngest regions.

For older ages, an incompleteness at low masses sets in, driven by our flux detection limit (see Fig.4). Our FUV-flux threshold translates (using our SSP models) into mass limits of  ,  , and 2900  at  Myr, Myr and  Myr respectively, if only foreground extinction were present. The actual limit is higher since most young stellar populations have additional internal extinction. The average detection limit in the  Myr bin is about 100 . The cluster mass function derived by Lada & Lada (2003) for embedded clusters in the solar vicinity has an exponent down to cluster masses of about , then it drops. If Lada & Lada’s results for solar vicinity were applicable to this dwarf irregular galaxy, the contribution of small mass clusters () to the total mass would be about 10% in this age bin, where we detect clusters with mass up to . However, star formation in NGC 6822 is patchy and may not resemble that of a massive spiral galaxy (Bianchi et al. 2001, Bianchi et al.2010, 2011); in any case this estimate should be regarded as an upper limit since the majority (%) of the embedded clusters are expected to merge and become part of larger OB associations or field stars before they are 10 Myr old (Lada & Lada, 2003). At older ages, the mass function for embedded clusters is no longer applicable: according to Lada & Lada (2003), after 100 Myr 94% of the embedded individual clusters have merged into large OB associations (if there was no cluster disruption, and if the above mass function were applicable, our detection limit would miss 45% of the clusters’ mass at this epoch: this approximate figure can be taken as a conservative upper limit). The sum of the masses of our SF regions with ages between Myr yields SFR  yr in this period, in agreement with the study of stellar populations by Gallart et al. (1996), who found SFR  yr in this epoch (adopting  mag). However, our value should be considered as a lower limit on SFR because the FUV detection threshold corresponds to a high mass detection-limit for old populations. In addition, the uncertainty factors in the UV-based method, discussed previously, become significant at ages older than 100 Myr.

We point out that our photometry of individual regions aims at isolating young SF complexes, and the older, diffuse galaxy population is subtracted from the flux. Therefore, it would not be appropriate to compare the total flux from the SF regions photometry with integrated galaxy models assuming a global SFH (see also Section 2.3).

4.2 The Very Recent Star Formation from H and far-IR Measurements

We also estimated the SFR of NGC6822 in very recent epochs from the H emission (Section 2.4). The total H flux from HII regions measured in this paper (the sum of the flux of the sources in Table 2) is erg s cm . After correcting the flux of the individual sources for reddening, using  values from the associated or nearest FUV sources, and A, the total unreddened flux is erg s cm, corresponding to a luminosity of erg s. The uncertainty takes into account photometric errors and  scatter within individual SF regions. Most of the H emission (%) comes from H regions # 5, 15, 26 (see Fig. 1 and Table 2). These include the H II regions Hubble V and Hubble X, where an excellent agreement was found (under the assumption of optically thick gas) between the H luminosity and the number of ionizing photons estimated using resolved photometry (Bianchi et al. 2001, Bianchi & Efremova 2006) and spectroscopy (B. Efremova et al. 2011, in preparation). The H luminosity translates into SFR  yr using the calibration by Panuzzo et al. (2003), which is based on the same SSP models we used to analyze the UV fluxes, the difference between the case (models) with and without dust being less than the uncertainty. Using other calibrations we obtain similar results: for example SFR  yr if we use the calibration by Hirashita et al. (2003), with =1. Among our H-defined HII regions there is one (FUV source EB-FUV #52, H source EB-H #25) which includes only two blue stars (an O-type star and an early B-type supergiant, according to our spectroscopy). This gives an indication of the sensitivity of our H source detection. In general, H-based SFRs should always be regarded as a lower limit, due to the possibility of leakage of ionizing photons.

H emission traces the hottest, most massive stars, and this estimate is in good agreement with the UV-derived SFR in the recent  Myr, as we would expect. Our result is also in agreement with previous estimates of SFR based on H luminosity: SFR yr by Hunter & Elmegreen (2004), and SFR  yr by Cannon et al. (2006).

Star formation more recent than  Myr is embedded in dust and detectable by 24 dust emission rather than in the UV, where the stellar flux is still heavily obscured. We estimated the SFR from the 24 µm emission using the /MIPS measurements by Cannon et al. (2006), and the second term in equation D10 of Leroy et al. (2008). We derived SFR(24 µm)  yr; this value is shown with a blue line in Fig. 5, over a very short time interval, since IR dust emission typically traces the youngest populations, where the massive stars have not yet dissipated the dust of the parent cloud. Therefore, the far-IR detected star formation should complement the stellar mass of young populations detected from FUV measurements.

Figure 6: The distribution of FUV-defined SF regions (blue contours) compared with the location of known Cepheids in NGC 6822 from the catalog of Pietrzyński et al. (2004) (red dots), shown over the FUV image.

5 Summary and Conclusions

We have defined regions of recent star formation in NGC6822 from wide-field FUV imaging, and derived ages and masses from their FUV, NUV photometry compared with SSP model populations. UV light is a good tracer of stellar populations up to a few hundred million years old. Extinction by interstellar dust has been estimated in each SF region from resolved photometry of the stars it contains: with an average value of =0.37mag, it exceeds in most cases the foreground reddening of =0.22mag. The characteristics of the internal (within NGC6822) selective extinction at UV wavelengths have been explored by comparing results from integrated UV photometry with H-R diagrams from high-resolution stellar photometry available for a number of well studied SF regions. We found that a UV-steep, non MW-type extinction is preferable for high surface brightness SF regions, and adopt it for sources with surface brightness 25 mag arcsec. MW-type extinction with seems adequate for less compact regions, which generally tend to be older. Such ad hoc criterion, suggested by our study for this particular case, will be verified over a larger sample of Local Group galaxies with new multi-band imaging (e.g. Bianchi et al. 2010, 2011). This study has shown quantitatively that large variations of dust characteristics as a function of environment, and related to star-formation intensity, exist within a single galaxy. The significant effect that the extinction correction bears on SFR derived from UV fluxes (Fig. 5) highlights the limitations of integrated SFR recipes, if internal extinction is not properly accounted for.

We avoided commonly used methods for deriving  from the ratio of FUV to 24µm fluxes (e.g. Burgarella et al. 2005), because most far-IR emitting regions in this galaxy are clearly not co-located with the FUV emitting regions (e.g. Bianchi 2007). UV and far-IR bands trace different populations, and using the flux ratio would consequently overcorrect the FUV fluxes. As can be seen in the 24 µm  image of NGC 6822 published by Cannon et al. (2006), most peaks of IR emission trace H emission sites, with the exception of their region 11, which is not a source of enhanced H emission. The 24 µm emission originates from dust heated by newly formed stars, still embedded in their parental clouds. The FUV-bright regions are more uniformly spread, because populations older than  Myr, no longer associated with dust nor significantly ionizing gas, still emit detectable FUV flux; only the youngest, most compact regions are bright in both FUV and H.

We estimated the total stellar mass formed in recent time intervals, by summing the masses of individual SF regions of corresponding ages. We derive an average SFR= yr over the past 100 Myr. For older ages the FUV-detection becomes incomplete and our method less robust, due to dissolving and possible merging of aging SF complexes. The uncertainty on SFR due to photometric errors, and extinction correction (the major factor) is very large (shown in Fig. 5). The overall level of star-formation activity may not be significantly variable in the last 100Myr, if we add to the FUV-detected young populations the embedded star-formation component, traced by far-IR emission from the dust which extinguishes the FUV flux. We found, similarly to Kang et al (2009), that the H SFR estimate is a good measurement of the recent ( Myr) star-formation as assessed by the FUV imaging, when it is concentrated in bright compact sources, which are likely to be optically thick.

Finally, we examine the location of the known Cepheids in NGC 6822 using the catalog by Pietrzyński et al. (2004). Such stars are the evolved descendants of populations formed mostly in earlier epochs than those sampled by this work. A few are within our FUV-source contours but mostly they avoid the FUV-bright regions and follow instead the optical appearance of the galaxy (Figure 6). While the youngest, FUV-bright SF regions are mostly found in the northern third of the galaxy, the Cepheids populate more uniformly the middle and southern part of the galaxy. We estimated their ages from the period-age relation derived by Efremov (2003) for Cepheids of similar metallicity (in the LMC). According to this relation, 85% of the Cepheids in NGC 6822 are older than 70 Myr. However, the period-age relation is constrained by very few data points at young ages (see Fig.3 of Efremov 2003), and we consider the overall spatial distribution more informative than individual ages.


We thank Philip Massey for very helpful clarifications on the calibration of the H image, Alin Tolea for initial discussions about the source contour definition, and the anonymous referee for valuable comments.

The data presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NAG5-7584 and by other grants and contracts. (The Galaxy Evolution Explorer) is a NASA Small Explorer, launched in April 2003. The H image used in this paper was obtained by Massey et al. (2007b) as a part of the Survey of Local Group Galaxies Currently Forming Stars and downloaded from

We gratefully acknowledge NASA’s support for construction, operation, and science analysis of the mission, developed in cooperation with the Centre National d’Etudes Spatiales of France and the Korean Ministry of Science and Technology. S.-C. R. is supported by the NRF of Korea to the Center for Galaxy Evolution Research.


  • Bianchi et al. (2011) Bianchi, L., 2011, Ap&SS, in press (special issue ”UV Universe 2010”, editors B. Shustov, A.I. Gomez de Castro, and M. Sachkov)
  • Bianchi et al. (2011) Bianchi, L., et al. 2011, Ap&SS, in press (special issue ”UV Universe 2010”, editors B. Shustov, A.I. Gomez de Castro, and M. Sachkov)
  • Bianchi et al. (2010) Bianchi, L. , et al. 2010, AAS 215, 455.25
  • Bianchi (2009) Bianchi, L. 2009, Ap&SS, 320, 11
  • Bianchi (2007) Bianchi, L. 2007, in “UV Astronomy: Stars from Birth to Death”, eds. A.I. Gomez de Castro and M. Barstow, UCM Editorial Complutense, p. 65
  • Bianchi & Efremova (2006) Bianchi, L., & Efremova, B. V. 2006, AJ, 132, 378
  • Bianchi et al. (2003a) Bianchi, L., et al. 2003a, Bulletin of the American Astronomical Society, 35, 1354
  • Bianchi et al. (2003b) Bianchi, L., et al. 2003b, in “The Local Group as an Astrophysical Laboratory”, STScI publication (M.Livio and T.Brown eds.), p. 10
  • Bianchi et al. (2001) Bianchi, L., Scuderi, S., Massey, P., & Romaniello, M. 2001, AJ, 121, 2020
  • Bianchi et al. (1999) Bianchi, L., Chandar, R., Ford, H. 1999, Mem. SAIT, eds. Martino & L. Buson, Vol. 70 N.2, p. 629
  • Boissier et al. (2007) Boissier, S., et al. 2007, ApJS, 173, 524
  • Burgarella, D. et al. (2005) Burgarella, D., Buat, V.& Iglesias-Páramo, J. 2005, MNRAS, 360, 1413
  • Calzetti (2001) Calzetti, D. 2001, PASP, 113, 1449
  • Calzetti et al. (2005) Calzetti, D., et al. 2005, ApJ, 633, 871
  • Cannon et al. (2006) Cannon, J. M., et al. 2006, ApJ, 652, 1170
  • Cardelli et al. (1989) Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245
  • Chandar et al. (1999) Chandar, R., Bianchi, L., Ford, H. and Salasnich, B. 1999, PASP, 111, 794
  • Cortese et al. (2006) Cortese, L., et al. 2006, ApJ, 637, 242
  • de Blok & Walter (2006) de Blok, W. J. G., & Walter, F. 2006, AJ, 131, 343
  • Gallart et al. (1996) Gallart, C., Aparicio, A., Bertelli, G., & Chiosi, C. 1996, AJ, 112, 2596
  • Gil de Paz et al. (2007) Gil de Paz, A., et al. 2007, ApJS, 173, 185
  • Girardi et al. (1995) Girardi, L., Chiosi, C., Bertelli, G., & Bressan, A. 1995, A&A, 298, 87
  • Efremov (2003) Efremov, Y. N. 2003, Astron. Rep., 47, 1000
  • Fall et al. (2009) Fall, S. M., Chandar, R., & Whitmore, B. C. 2009, ApJ, 704, 453
  • Fatuzzo & Adams (2008) Fatuzzo, M., & Adams, F. C. 2008, ApJ, 675, 1361
  • Gordon & Clayton (1998) Gordon, K. D., & Clayton, G. C. 1998, ApJ, 500, 816
  • Hirashita et al. (2003) Hirashita, H., Buat, V., & Inoue, A. K. 2003, A&A, 410, 83
  • Hodge (1977) Hodge, P. W. 1977, ApJS, 33, 69
  • Hodge et al. (1988) Hodge, P., Lee, M. G., & Kennicutt, R. C., Jr. 1988, PASP, 100, 917
  • Hodge et al. (1989) Hodge, P., Lee, M. G., & Kennicutt, R. C., Jr. 1989, PASP, 101, 32
  • Hubble (1925) Hubble, E. P. 1925, ApJ, 62, 409
  • Hunter & Elmegreen (2004) Hunter, D. A., & Elmegreen, B. G. 2004, AJ, 128, 2170
  • Hunter et al. (2010) Hunter, D. A., Elmegreen, B. G., & Ludka, B. C. 2010, AJ, 139, 447
  • Kang et al. (2009) Kang, Y., Bianchi, L., & Rey, S.-C. 2009, ApJ, 703, 614
  • Karachentsev et al. (2004) Karachentsev, I. D., Karachentseva, V. E., Huchtmeier, W. K., & Makarov, D. I. 2004, AJ, 127, 2031
  • Kennicutt (1998) Kennicutt, R. C., Jr. 1998, ARA&A, 36, 189
  • Killen & Dufour (1982) Killen, R. M., & Dufour, R. J. 1982, PASP, 94, 444
  • Kinman et al. (1979) Kinman, T. D., Green, J. R., & Mahaffey, C. T. 1979, PASP, 91, 749
  • Kroupa (2001) Kroupa, P. 2001, MNRAS, 322, 231
  • Lada & Lada (2003) Lada, C. J., & Lada, E. A. 2003, ARA&A, 41, 57
  • Leroy et al. (2008) Leroy, A. K., Walter, F., Brinks, E., Bigiel, F., de Blok, W. J. G., Madore, B., & Thornley, M. D. 2008, AJ, 136, 2782
  • Martin et al. (2005) Martin, D. C., et al. 2005, ApJ, 619, L1
  • Massey et al. (2007a) Massey, P., Olsen, K. A. G., Hodge, P. W., Jacoby, G. H., McNeill, R. T., Smith, R. C., & Strong, S. B. 2007a, AJ, 133, 2393
  • Massey et al. (2007b) Massey, P., McNeill, R. T., Olsen, K. A. G., Hodge, P. W., Blaha, C., Jacoby, G. H., Smith, R. C., & Strong, S. B. 2007b, AJ, 134, 2474
  • McAlary et al. (1983) McAlary C. W., Madore, B. F., C. W.,McGonegal, R., McLaren, R. A., & Welch, D. L.  1983, ApJ, 273, 539
  • Melena et al. (2009) Melena, N. W., Elmegreen, B. G., Hunter, D. A., & Zernow, L. 2009, AJ, 138, 1203
  • Meurer et al. (2009) Meurer, G. R., et al. 2009, ApJ, 695, 765
  • Misselt et al. (1999) Misselt, K. A., Clayton, G. C., & Gordon, K. D. 1999, ApJ, 515, 128
  • Morrissey et al. (2007) Morrissey, P., et al. 2007, ApJS, 173, 682
  • Muschielok et al. (1999) Muschielok, B., et al. 1999, A&A, 352, L40
  • O’Dell et al. (1999) O’Dell, C. R., Hodge, P. W., & Kennicutt, R. C., Jr. 1999, PASP, 111, 1382
  • Pagel et al. (1980) Pagel, B. E. J., Edmunds, M. G., & Smith, G. 1980, MNRAS, 193, 219
  • Panuzzo et al. (2003) Panuzzo, P. Bressan, A., Silva, L. et al. 2003, A&A, 409, 99
  • Pietrzyński et al. (2004) Pietrzyński, G., Gieren, W., Udalski, A., Bresolin, F., Kudritzki, R.-P., Soszyński, I., Szymański, M., & Kubiak, M. 2004, AJ, 128, 2815
  • Venn et al. (2001) Venn, K. A., et al. 2001, ApJ, 547, 765
  • Wyder et al. (2007) Wyder, T. K., et al. 2007, ApJS, 173, 293
#aaNumbers correspond to the blue labels in Fig. 1. (J2000)bbCoordinates of the centroids of the contours (see Sec. 2.2). (J2000)bbCoordinates of the centroids of the contours (see Sec. 2.2). Area FUVccPhotometric errors are estimated as explained in the text, without adding uncertainties of calibration zero points. NUVccPhotometric errors are estimated as explained in the text, without adding uncertainties of calibration zero points. Background fractionddFraction of background flux over total flux in the aperture. eeThe median of values derived for the massive stars (earlier than B2V) in each SF region (see Sec. 3.2). CommentffNames of OB associations after Hodge (1977).
[arcsec/pc] [AB mag] [AB mag] FUV NUV [mag] and scatter
EB-FUV 1 19 44 14.46 -14 46 39.1 190 / 1110 0.08 0.13
EB-FUV 2ggThe FUV source is associated with H emission. 19 44 31.00 -14 47 23.5 190 / 1090 0.21 0.30
EB-FUV 3ggThe FUV source is associated with H emission. 19 44 31.63 -14 44 04.0 340 / 1950 0.19 0.31 part of OB2
EB-FUV 4ggThe FUV source is associated with H emission. 19 44 32.45 -14 44 45.2 410 / 2350 0.15 0.25 part of OB2
EB-FUV 5ggThe FUV source is associated with H emission. 19 44 33.02 -14 47 32.5 190 / 1090 0.20 0.31
EB-FUV 6ggThe FUV source is associated with H emission. 19 44 33.13 -14 42 01.7 3440 / 19810 0.12 0.16 OB1 and OB3
EB-FUV 7 19 44 33.95 -14 42 57.2 450 / 2600 0.28 0.38
EB-FUV 8ggThe FUV source is associated with H emission. 19 44 37.82 -14 50 57.3 440 / 2540 0.24 0.36 OB4
EB-FUV 9 19 44 40.88 -14 43 40.0 160 / 910 0.21 0.39
EB-FUV 10ggThe FUV source is associated with H emission. 19 44 40.98 -14 51 19.8 160 / 900 0.19 0.28
EB-FUV 11 19 44 41.70 -14 52 17.5 300 / 1720 0.20 0.31
EB-FUV 12 19 44 42.74 -14 50 25.0 280 / 1630 0.23 0.30
EB-FUV 13ggThe FUV source is associated with H emission. 19 44 43.67 -14 46 30.3 280 / 1630 0.15 0.24
EB-FUV 14 19 44 43.82 -14 51 36.3 360 / 2090 0.32 0.45
EB-FUV 15 19 44 46.05 -14 51 28.1 2700 / 15550 0.30 0.41
EB-FUV 16ggThe FUV source is associated with H emission. 19 44 47.24 -14 46 59.6 160 / 920 0.18 0.28
EB-FUV 17 19 44 47.91 -14 52 40.8 190 / 1100 0.28 0.38
EB-FUV 18ggThe FUV source is associated with H emission. 19 44 48.43 -14 46 16.8 150 / 870 0.19 0.32
EB-FUV 19ggThe FUV source is associated with H emission. 19 44 48.90 -14 45 30.3 1160 / 6680 0.17 0.25 OB7
EB-FUV 20ggThe FUV source is associated with H emission. 19 44 49.36 -14 44 00.3 1500 / 8620 0.17 0.25 OB6
EB-FUV 21ggThe FUV source is associated with H emission. 19 44 49.52 -14 43 15.3 200 / 1130 0.28 0.38
EB-FUV 22 19 44 50.96 -14 44 52.8 360 / 2060 0.35 0.49
EB-FUV 23 19 44 51.06 -14 50 43.1 590 / 3380 0.38 0.52
EB-FUV 24 19 44 51.53 -14 45 39.3 200 / 1160 0.37 0.51
EB-FUV 25ggThe FUV source is associated with H emission. 19 44 52.35 -14 52 38.6 3480 / 20030 0.28 0.38 part of OB5
EB-FUV 26 19 44 53.13 -14 48 37.1 330 / 1930 0.28 0.44
EB-FUV 27ggThe FUV source is associated with H emission. 19 44 53.39 -14 43 02.6 2760 / 15900 0.14 0.18 OB8
EB-FUV 28 19 44 53.44 -14 44 59.6 270 / 1560 0.33 0.45
EB-FUV 29 19 44 54.48 -14 50 16.1 5060 / 29160 0.39 0.51
EB-FUV 30 19 44 54.79 -14 46 16.8 5400 / 31100 0.37 0.48
EB-FUV 31 19 44 55.00 -14 51 49.1 520 / 3010 0.30 0.46
EB-FUV 32 19 44 55.77 -14 42 04.8 160 / 920 0.23 0.35
EB-FUV 33 19 44 56.29 -14 52 01.1 250 / 1460 0.31 0.38
EB-FUV 34 19 44 56.34 -14 44 05.6 3920 / 22560 0.31 0.39 OB9
EB-FUV 35ggThe FUV source is associated with H emission. 19 44 57.48 -14 47 28.8 730 / 4180 0.30 0.39
EB-FUV 36 19 44 57.89 -14 50 04.1 410 / 2380 0.38 0.49
EB-FUV 37 19 44 58.10 -14 51 11.6 210 / 1190 0.34 0.45
EB-FUV 38 19 44 58.26 -14 45 57.3 530 / 3040 0.46 0.63
EB-FUV 39 19 44 58.36 -14 48 41.6 4220 / 24320 0.32 0.43
EB-FUV 40 19 44 58.97 -14 52 41.6 290 / 1700 0.43 0.56
EB-FUV 41 19 44 59.29 -14 45 31.8 180 / 1040 0.44 0.57
EB-FUV 42 19 44 59.60 -14 56 22.1 240 / 1410 0.23 0.28
EB-FUV 43 19 45 00.07 -14 49 47.6 260 / 1520 0.39 0.59
EB-FUV 44 19 45 00.22 -14 50 38.6 330 / 1900 0.33 0.48
EB-FUV 45 19 45 00.58 -14 52 46.8 170 / 990 0.52 0.64
EB-FUV 46 19 45 00.73 -14 44 43.8 4170 / 24030 0.29 0.40 OB11
EB-FUV 47 19 45 00.73 -14 45 34.8 190 / 1120 0.38 0.55
EB-FUV 48 19 45 00.94 -14 58 53.6 200 / 1140 0.13 0.22
EB-FUV 49 19 45 01.05 -14 52 58.8 230 / 1320 0.45 0.56
EB-FUV 50 19 45 01.20 -14 54 33.3 570 / 3270 0.26 0.28
EB-FUV 51 19 45 01.41 -14 49 46.8 220 / 1250 0.25 0.37
EB-FUV 52ggThe FUV source is associated with H emission. 19 45 01.93 -14 46 55.8 170 / 990 0.16 0.22
EB-FUV 53 19 45 02.23 -14 52 33.3 170 / 1010 0.35 0.38
EB-FUV 54 19 45 03.02 -14 54 33.3 400 / 2280 0.39 0.49
EB-FUV 55 19 45 03.22 -14 53 07.1 250 / 1450 0.39 0.48
EB-FUV 56 19 45 04.05 -14 56 14.6 580 / 3360 0.22 0.31
EB-FUV 57ggThe FUV source is associated with H emission. 19 45 04.77 -14 43 25.8 2150 / 12400 0.08 0.13 OB13
EB-FUV 58 19 45 05.03 -14 54 45.3 770 / 4460 0.25 0.31
EB-FUV 59 19 45 05.44 -14 52 59.6 270 / 1560 0.48 0.57
EB-FUV 60 19 45 06.53 -14 55 52.8 3630 / 20880 0.29 0.38 OB12
EB-FUV 61 19 45 06.74 -14 51 55.1 150 / 870 0.36 0.46
EB-FUV 62 19 45 08.66 -14 55 41.6 190 / 1110 0.20 0.31
EB-FUV 63 19 45 08.91 -14 50 17.6 230 / 1310 0.31 0.44
EB-FUV 64 19 45 09.32 -14 55 04.1 210 / 1210 0.33 0.48
EB-FUV 65 19 45 09.48 -14 53 23.5 520 / 3010 0.39 0.51
EB-FUV 66ggThe FUV source is associated with H emission. 19 45 10.31 -14 48 55.0 2100 / 12090 0.17 0.23 OB14
EB-FUV 67 19 45 10.31 -14 54 25.8 290 / 1670 0.36 0.50
EB-FUV 68 19 45 10.94 -14 55 26.5 150 / 870 0.40 0.59
EB-FUV 69 19 45 11.14 -14 53 40.0 200 / 1130 0.33 0.46
EB-FUV 70 19 45 11.33 -14 44 27.3 270 / 1570 0.26 0.42
EB-FUV 71ggThe FUV source is associated with H emission. 19 45 11.60 -14 54 58.8 530 / 3070 0.27 0.39
EB-FUV 72 19 45 12.17 -14 52 52.0 410 / 2370 0.36 0.48
EB-FUV 73 19 45 12.80 -14 57 34.0 260 / 1510 0.16 0.24
EB-FUV 74 19 45 13.47 -14 58 49.0 240 / 1370 0.21 0.28
EB-FUV 75ggThe FUV source is associated with H emission. 19 45 14.85 -14 45 01.8 3410 / 19630 0.14 0.24 OB15
EB-FUV 76 19 45 20.28 -14 45 37.0 240 / 1370 0.21 0.29
EB-FUV 77 19 45 30.58 -14 50 05.4 310 / 1800 0.24 0.40
Table 1: FUV-selected SF regions
ID (J2000)aaCoordinates of the ”centroids” of the H contours (see Sec. 2.2). (J2000)aaCoordinates of the ”centroids” of the H contours (see Sec. 2.2). Area H Flux CommentbbCross-identification with H sources in previous works: HI-X, Hubble (1925); Ho followed by arabic number, Hodge (1977); K followed by a greek letter, Kinman et al. (1979); KD followed by arabic number, Killen & Dufour (1982); HK followed by arabic number, Hodge et al. (1988)
[arcsec/pc] [ ergs sec cm]
EB-H 1 19 44 30.89 -14 47 19.7 410 / 2380 0.24  0.03 HK1
EB-H 2 19 44 30.93 -14 48 29.8 250 / 1420 1.77  0.11 K, KD2e
EB-H 3 19 44 32.29 -14 44 14.7 4420 / 25440 7.92  0.19 HII, Ho1, Ho3, HK2, HK4D, KD2, KD3
EB-H 4 19 44 32.33 -14 47 40.9 540 / 3090 1.22  0.18 K, KD5e
EB-H 5 19 44 32.90 -14 42 02.6 5990 / 34510 42.01  0.36 HI, HIII, Ho2,KD1, Ho4,KD4
EB-H 6 19 44 38.72 -14 42 32.7 380 / 2220 0.21  0.22 HK5D
EB-H 7 19 44 38.75 -14 51 10.2 4850 / 27940 6.94  1.57 Ho5, KD8, KD9, HK6, KD7
EB-H 8 19 44 39.98 -14 52 00.3 340 / 1950 0.17  0.03 HK7, HK8
EB-H 9 19 44 43.68 -14 47 00.2 160 / 930 0.03  0.02
EB-H 10 19 44 44.36 -14 45 57.9 1160 / 6710 0.62  0.46 HK11D, HK12
EB-H 11 19 44 47.46 -14 46 57.2 810 / 4660 0.61  0.08 HK23
EB-H 12 19 44 47.92 -14 51 29.5 320 / 1830 0.22  0.05
EB-H 13 19 44 48.44 -14 46 12.4 540 / 3110 0.75  0.06 Ho7, KD11
EB-H 14 19 44 48.45 -14 45 24.9 670 / 3840 0.43  0.38 HK22, HK27, HK34
EB-H 15 19 44 49.34 -14 43 48.0 9500 / 54710 50.74  1.14 HV, Ho6, KD12, KD21, HK13, HK15, HK16, HK17, HK19, HK20, HK21, HK32, HK33, HK35, HK36, HK40, HK42, HK44, HK53, Ho9, Ho11, KD19, KD10
EB-H 16 19 44 49.73 -14 52 57.8 170 / 970 1.05  0.05 KD13, KD13e
EB-H 17 19 44 50.52 -14 52 06.9 3210 / 18520 5.41  0.34 Ho10, K, KD18 , KD11e
EB-H 18 19 44 50.57 -14 52 45.9 550 / 3150 2.81  0.27 Ho8, KD14, KD15, HK48, KD16, KD17
EB-H 19 19 44 52.54 -14 44 11.0 300 / 1750 0.16  0.04 HK55D
EB-H 20 19 44 55.23 -14 42 09.5 150 / 880 0.08  0.02
EB-H 21 19 44 57.66 -14 47 25.9 630 / 3620 0.83  0.16 KD24, HK66, HK67
EB-H 22 19 45 00.34 -14 43 40.1 300 / 1700 0.24  0.25 HK73
EB-H 23 19 45 01.23 -14 54 20.8 320 / 1840 0.08  0.04
EB-H 24 19 45 01.98 -14 54 00.2 190 / 1090 0.07  0.02
EB-H 25 19 45 02.38 -14 46 55.9 970 / 5570 1.32  0.08
EB-H 26 19 45 03.88 -14 43 35.3 3980 / 22910 37.68  3.49 HX, HK77, HK78, HK79, HK80, Ho14, KD26
EB-H 27 19 45 05.54 -14 57 20.6 1270 / 7330 1.50  0.07 KD27, KD28
EB-H 28 19 45 09.73 -14 44 39.8 640 / 3710 0.27  0.05 HK97D
EB-H 29 19 45 10.91 -14 45 35.4 1210 / 6960 0.52  0.19 HK98D
EB-H 30 19 45 10.94 -14 48 53.4 3130 / 18000 3.74  0.28 Ho15, KD30
EB-H 31 19 45 11.41 -14 43 47.7 580 / 3340 0.35  0.18
EB-H 32 19 45 11.95 -14 54 47.7 440 / 2540 0.16  0.11
EB-H 33 19 45 14.07 -14 58 50.5 160 / 910 0.18  0.02
EB-H 34 19 45 14.27 -14 43 47.6 260 / 1510 0.09  0.10
EB-H 35 19 45 15.83 -14 44 58.4 3050 / 17560 3.24  0.20 Ho16, HK105D, HK106, HK107, KD31
EB-H 36 19 45 16.03 -14 49 13.5 260 / 1470 0.13  0.03
EB-H 37 19 45 17.40 -14 44 01.5 260 / 1500 0.09  0.06
EB-H 38 19 45 17.89 -14 44 24.8 180 / 1040 0.04  0.02
EB-H 39 19 45 18.87 -14 44 57.0 190 / 1100 0.09  0.19 HK108
Table 2: H emission regions
MW with Rv=3.1bbFor all sources, a foreground extinction of was applied with MW-type dust, and the additional internal extinction (from the total  given in column 9 of Table 1) using three different dust types: ‘Rv3.1’ stands for average MW extinction with Rv=3.1 (Cardelli et al., 1989); ‘LMC’ indicates the average LMC extinction for stars outside LMC2 by (Misselt et al., 1999); ‘LMC2’ indicates that the extinction curve by (Misselt et al., 1999) was used. LMCbbFor all sources, a foreground extinction of was applied with MW-type dust, and the additional internal extinction (from the total  given in column 9 of Table 1) using three different dust types: ‘Rv3.1’ stands for average MW extinction with Rv=3.1 (Cardelli et al., 1989); ‘LMC’ indicates the average LMC extinction for stars outside LMC2 by (Misselt et al., 1999); ‘LMC2’ indicates that the extinction curve by (Misselt et al., 1999) was used. LMC2bbFor all sources, a foreground extinction of was applied with MW-type dust, and the additional internal extinction (from the total  given in column 9 of Table 1) using three different dust types: ‘Rv3.1’ stands for average MW extinction with Rv=3.1 (Cardelli et al., 1989); ‘LMC’ indicates the average LMC extinction for stars outside LMC2 by (Misselt et al., 1999); ‘LMC2’ indicates that the extinction curve by (Misselt et al., 1999) was used. ext. curve
IDaaCorresponding to the IDs in Table 1 and the blue labels in Fig. 1. Age[Myr] Mass[] Age[Myr] Mass[] Age[Myr] Mass[] adoptedccThe results for the extinction curve given in this column were adopted in the analysis. only for source # 6 we adopted an SMC-type extinction curve, which gives age = Myr and mass =
EB-FUV 1 . . . LMC
EB-FUV 2 . . . Rv3.1
EB-FUV 3 . . . Rv3.1
EB-FUV 4 . . . LMC
EB-FUV 5 . . . Rv3.1
EB-FUV 6 . . . SMC
EB-FUV 7 . . . Rv3.1
EB-FUV 8 . . . Rv3.1
EB-FUV 9 . - -. - -. Rv3.1
EB-FUV 10 . - -. - -. Rv3.1
EB-FUV 11 . . . Rv3.1
EB-FUV 12 . . . Rv3.1
EB-FUV 13 . . . LMC
EB-FUV 14 . . . Rv3.1
EB-FUV 15 . . . Rv3.1
EB-FUV 16 . . . Rv3.1
EB-FUV 17 . . . Rv3.1
EB-FUV 18 . . . Rv3.1
EB-FUV 19 . . . LMC2
EB-FUV 20 . . . LMC2
EB-FUV 21 . . . Rv3.1
EB-FUV 22 . . . Rv3.1
EB-FUV 23 . . - -. Rv3.1
EB-FUV 24 . . . Rv3.1
EB-FUV 25 . . . Rv3.1
EB-FUV 26 . . - -. Rv3.1
EB-FUV 27 . .