Reaching the oldest stars beyond the Local Group: ancient star formation in UGC 4483
Elena Sacchi, Alessandra Aloisi, Matteo Correnti, Francesca Annibali, Monica Tosi, Alessia Garofalo, Gisella Clementini, Michele Cignoni, Bethan James, Marcella Marconi, Tatiana Muraveva, Roeland van der Marel
DDraft version March 1, 2021
Preprint typeset using L A TEX style emulateapj v. 12/16/11
REACHING THE OLDEST STARS BEYOND THE LOCAL GROUP:ANCIENT STAR FORMATION IN UGC 4483 (cid:63)
Elena Sacchi , , , Alessandra Aloisi , Matteo Correnti , Francesca Annibali , Monica Tosi , AlessiaGarofalo , Gisella Clementini , Michele Cignoni , , , Bethan James , Marcella Marconi , TatianaMuraveva , and Roeland van der Marel , Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA Leibniz-Institut für Astrophysik Potsdam, An der Sternwarte 16, 14482 Potsdam, Germany; [email protected] INAF–Osservatorio di Astrofisica e Scienza dello Spazio di Bologna, Via Gobetti 93/3, I-40129 Bologna, Italy Dipartimento di Fisica, Università di Pisa, Largo Bruno Pontecorvo, 3, 56127 Pisa, Italy INFN, Sezione di Pisa, Largo Pontecorvo 3, 56127 Pisa, Italy INAF–Osservatorio Astronomico di Capodimonte, Salita Moiariello 16, 80131, Naples, Italy Center for Astrophysical Sciences, Department of Physics & Astronomy, Johns Hopkins University, Baltimore, MD 21218, USA
Draft version March 1, 2021
ABSTRACTWe present new WFC3/UVIS observations of UGC 4483, the closest example of a metal-poor bluecompact dwarf galaxy, with a metallicity of Z (cid:39) / Z (cid:12) and located at a distance of D (cid:39) . Mpc.The extremely high quality of our new data allows us to clearly resolve the multiple stellar evolutionaryphases populating the color-magnitude diagram (CMD), to reach more than 4 mag deeper than thetip of the red giant branch, and to detect for the first time core He-burning stars with masses (cid:46) M (cid:12) ,populating the red clump and possibly the horizontal branch (HB) of the galaxy. By applying thesynthetic CMD method to our observations, we determine an average star formation rate over thewhole Hubble time of at least (7 . ± . × − M (cid:12) / yr , corresponding to a total astrated stellarmass of (9 . ± . × M (cid:12) , 87% of which went into stars at epochs earlier than 1 Gyr ago.With our star formation history recovery method we find the best fit with a distance modulus ofDM = . ± . , slightly lower than previous estimates. Finally, we find strong evidence of anold ( (cid:38) Gyr) stellar population in UGC 4483 thanks to the detection of an HB phase and theidentification of six candidate RR Lyrae variable stars.
Keywords: galaxies: dwarf – galaxies: irregular – galaxies: evolution – galaxies: individual (UGC 4483)– galaxies: star formation – galaxies: stellar content – galaxies: starburst INTRODUCTIONStar formation (SF) studies beyond the Local Grouphave been pushed to their current limits in the past fewyears. Thanks to the spatial resolution and sensitivityof the
Hubble Space Telescope ( HST ), it is possible toresolve single stars in galaxies up to ∼ Mpc, with thelimitation of characterizing only the brightest stellar evo-lutionary features, i.e., the upper main sequence (MS),the He-burning phase of massive and intermediate-massstars (blue loops, BLs), the asymptotic giant branch(AGB), and the red giant branch (RGB). Even thoughthe RGB can host stars with any age between − and13 Gyr, it is unfortunately affected by an age-metallicitydegeneracy; this implies that only the most recent starformation history (SFH), back to − Gyr ago, can bederived with good time resolution ( ∼ ), while a pre-cise characterization of the SF behaviour is not possibleat the oldest epochs.So, why do we even go through the trouble of study-ing such distant systems, with all the uncertainties asso-ciated with them? Beyond removing the environmentaleffects that the big spirals inside the Local Group have onsmaller systems, and studying dwarf galaxies in isolation,a key reason is that a particular sub-group of the dwarf (cid:63) Based on observations obtained with the NASA/ESA
HubbleSpace Telescope at the Space Telescope Science Institute, whichis operated by the Association of Universities for Research inAstronomy under NASA Contract NAS 5-26555. class is not present inside the Local Group: blue com-pact dwarf (BCD) galaxies. These are extremely inter-esting systems, most often characterized by bluer colors,more intense star formation activities, and higher cen-tral surface brightness compared to regular star-formingdwarfs. Because of their recent or ongoing bursts of SFand typical low metallicities, for a long time they were be-lieved to be young galaxies, but all the systems studied indetails through their color-magnitude diagrams (CMDs)and SFHs revealed a population of RGB stars, thus atleast − Gyr old, and possibly as old as 13.7 Gyr (see,e.g., Schulte-Ladbeck et al. 2001; Annibali et al. 2003;Aloisi et al. 2007; McQuinn et al. 2015).All the studies conducted so far on BCDs are limitedby the depth of the photometry, reaching − magni-tudes below the RGB tip (TRGB), which allows to ro-bustly reconstruct the SFH up to a few Gyr ago only. Tomake progress, it is necessary to reach fainter featureswith discriminatory power at older ages. The bright-est signatures of several Gyr old stars are the red clump(RC) and horizontal branch (HB), both correspondingto core He-burning evolution phases of stars with masses (cid:46) M (cid:12) . In particular, despite the uncertainties on whichother parameters actually affect the color extension ofthe HB (e.g., alpha-element content, internal dynamics,etc.), the potential detection of a blue HB (i.e., starswith M (cid:46) M (cid:12) at M I (cid:39) mag) would unambiguouslyindicate the presence of a population that is both old( (cid:38) Gyr) and metal-poor ([Fe/H] (cid:46) − . dex). Old a r X i v : . [ a s t r o - ph . GA ] F e b Sacchi et al.
Figure 1.
Left panel.
HST /WFC3 observations: blue corresponds to F606W (broad V ),red to F814W ( I ), while the green channel was obtained using the mean of the two. H i contours from the VLA-ANGST survey (Ott et al.2012) corresponding to column densities of N H = . × cm − have been superimposed to the HST image.The displayed field of view is . (cid:48) × . (cid:48) . North is up, East is left. Right panel.
F656N image of UGC 4483 with superimposed the sameH i contours. metal-poor HB stars can cross the instability strip andproduce RR Lyrae (RRLs) type pulsating variables, sothe detection of such stars allows to unambiguously tracethe signature of a ∼ Gyr old population throughout agalaxy (see, e.g., Clementini et al. 2003) even when themagnitude level of the HB is close to the detection limit.Here we present a detailed analysis of UGC 4483, theclosest example of metal-poor BCD galaxy, located be-tween the bright spirals M81 and NGC 2403 at a distanceof D = 3 . ± . Mpc, corresponding to a distance mod-ulus of DM = 27 . ± . (Izotov & Thuan 2002), andwith an oxygen abundance of O/H ) = 7 . ± . (van Zee & Haynes 2006), corresponding to Z (cid:39) / Z (cid:12) (using 8.73 for the solar oxygen abundance, Caffau et al.2015). Its cometary shape (see Figure 1) resembles thatof SBS 1415+437 (Aloisi et al. 2005), which also hasa very similar metal content. From H i
21 cm obser-vations, Thuan & Seitzer (1979) derived a gas mass of M H i = 4 . × M (cid:12) , and a gas fraction which corre-sponds to 39% of all visible mass. The H i mass fromLelli et al. (2012b) is M H i = 2 . × M (cid:12) , and they alsofind a steeply-rising rotation curve that flattens in theouter parts, making UGC 4483 the lowest-mass galaxywith a differentially rotating H i disk. This steep rise ofthe rotation curve indicates a strong central concentra-tion of mass, a property which seems to be typical ofBCDs. UGC 4483 has already been resolved into stars withthe HST /WFPC2 (Dolphin et al. 2001; Izotov & Thuan2002; Odekon 2006; McQuinn et al. 2010). The I vs. V − I CMDs reveal a young stellar population of blueMS stars as well as blue and red supergiants associatedwith the bright H ii region at the northern tip of thegalaxy (see Figure 1 and Region 0a of Figure 2). An olderevolved population was found throughout the whole lowsurface brightness body of the galaxy, as indicated byvery bright AGB stars and the tip of a very blue RGB.However, those data were not deep enough to discrimi-nate between a relatively metal-poor population with anage of ∼ Gyr for the RGB/AGB stars, or a somewhathigher metallicity and an age of ∼ Gyr.We present here new WFC3/UVIS observations reach-ing more than mag deeper than the TRGB, to detectand characterize the RC and/or HB populations of thegalaxy. These new data strongly constrain both the ageand metallicity properties of the SFH of UGC 4483 backto many, possibly 10, Gyr ago. The RC absolute magni-tude (near M I (cid:39) − . mag) depends sensitively on theage of the stellar population, and it is ∼ mag brighterfor a 1 Gyr old population than for a 10 Gyr old popu-lation (see, e.g., figure 23 of Rejkuba et al. 2005). Thedependence of the RC magnitude on metallicity is notstrong, and either way, is different than the dependenceof the RGB color on age and metallicity. Therefore, a tar Formation History of UGC 4483 OBSERVATIONS AND DATA REDUCTIONObservations of UGC 4483 were performed on January2019 (Visits 1 and 2) and February 2019 (Visit 3) usingthe UVIS channel of the WFC3 as part of the
HST pro-gram GO-15194 (PI: A. Aloisi). Despite the smaller fieldof view, we preferred WFC3/UVIS over ACS/WFC be-cause of the slightly higher resolution that provides abetter photometry in the most crowded central regionsof the galaxy. The target was centered on one of thetwo CCD chips of the WFC3/UVIS camera, in order tominimize the impact of chip-dependent zeropoints andto avoid the loss of the galaxy central region due to theCCD chip gap. The observations were obtained in thetwo broad-band filters F606W and F814W. We also re-quired narrow-band imaging in F656N to study the dis-tribution of the ionized gas. We selected the broaderF606W instead of F555W as the blue filter, because itprovides the best compromise between a reasonable ex-posure time (a factor of ∼ shorter for F606W than forF555W) and the achievement of our science goals. Dur-ing Visits 1 and 2, 14 long exposures (2580 sec each) weretaken in the F606W filter (for a total exposure time of36120 sec), and 8 in the F814W filter (for a total ex-posure time of 20640 sec). Twelve additional exposureswith the same integration time (2580 sec each, for a totalexposure time of 30960 sec) were collected during Visit 3in the F814W filter. Those exposure times were chosenin order to achieve a signal-to-noise ratio of ∼ , at theHB magnitude level, which corresponds to I ∼ . magat the distance of UGC 4483. The exposures were exe-cuted with a spatial offset, using a dithering pattern ofan integer+fraction of pixels, in order to move across thegap between the two chips, to simplify the identificationand removal of bad/hot pixels and cosmic rays, and toimprove the point spread function (PSF) sampling.It is worth mentioning that the observations were per-formed in the Continuous Viewing Zone (CVZ) of HST ,which allowed us to observe for the entire 96 minute or-bit. This resulted in a very deep photometry even in just18 orbits.To reduce the images, we followed the same proce-dure outlined in Annibali et al. (2019). First, the cali-brated .flc science images were retrieved from the
HST archive. The .flc images are the products of the calwf3 data reduction pipeline and constitute the bias-corrected,dark-subtracted, flat-fielded, and charge transfer ineffi-ciency corrected images. Then, we combined togetherthe individual .flc images into a single drizzled, stacked,and distortion-corrected image ( .drc image) using the
Drizzlepac software (Gonzaga et al. 2012). To do so,all the images in the same filter were first aligned us-ing the software
TweakReg and then, using the software
AstroDrizzle , bad pixels and cosmic rays were flaggedand rejected from the input images. Finally, those inputundistorted and aligned images were combined together into a final stacked image.In the left panel of Figure 1 we show a 3-color com-posite image of UGC 4483 from our
HST /WFC3 obser-vations: blue corresponds to F606W (broad V ), red toF814W ( I ), while the green channel was obtained usingthe mean of the two. The right panel shows instead theF656N image of the galaxy (please note that given itsvery low S/N, the H α image has not been continuumsubtracted.) H i contours from the VLA-ANGST survey(Ott et al. 2012) corresponding to column densities of N H = . × cm − havebeen superimposed to both images. Nebular gas emissioncan contaminate broad-band photometry, but as shownby the right panel of Figure 1, outside of the most active0a and 0b regions (as defined in Figure 2), the contami-nation is negligible.PSF stellar photometry was performed using the latestversion of Dolphot (Dolphin 2000, and numerous subse-quent updates). After the usual pre-processing requiredby the software and performed on each single science im-age (i.e., creation of bad pixel mask and generation ofsky frame), iterative PSF photometry was performed si-multaneously on the .flc images using the F606W .drc image as the reference frame for alignment. The differentDolphot parameters that govern alignment and photom-etry were set as in Annibali et al. (2019), adopting ahybrid combination between the values recommended bythe Dolphot manual and those adopted in Williams et al.(2014).Together with the positions and photometry of the in-dividual stars, the final photometric catalog contains sev-eral diagnostic parameters which are useful to excluderemaining artifacts and spurious detections. Hence, weselected from the total catalog all the objects with theDolphot “object type” flag = 1, and then we applied aseries of consecutive selection cuts using the remainingdiagnostics (i.e., error, sharpness, roundness, and crowd-ing). The final clean catalog contains ∼
14 000 stars, andthe corresponding CMD is shown in Figure 3.We also analyzed the coordinated parallel images ob-tained with the ACS/WFC, which in principle could giveus rich information about the stellar population of thehalo of the galaxy. However, the resulting CMD does notcontain any sources we could link to UGC 4483, but onlybackground contamination; unfortunately, these fieldsare probably already too far away from the galaxy to in-clude its halo. This suggests that the halo of UGC 4483does not reach as far as the ACS field, i.e. is smaller than ∼ kpc.To properly analyze the spatial variations of the SFwithin the galaxy, we divided our final catalog in sixsub-regions, as shown in Figure 2, following the isophotalcontours of the F606W image. The innermost regions, 0aand 0b, correspond to the most active areas of the galaxy,where we see H α emission from the H ii regions. STELLAR POPULATIONSFigure 1 clearly shows the very young stellar popula-tion of this BCD galaxy. The bright H ii region at thenorthern edge of UGC 4483 is evident in both the com-posite 3-color image in the left panel and in H α emissionin the right panel, which also reveals another active re-gion more to the south. The H i distribution, shown withthe contours overlapped on the images, does not exactly Sacchi et al.
800 1050 1300 1550 1800 2050 2300 2550 2800 3050 x y Figure 2.
F606W image with over-plotted the isophotal contours used to divide the galaxy into different regions (0a, 0b, 1, 2, 3, and 4)and to study the radial behaviour of the stellar populations and SFH. Notice that the image is rotated by 90 degrees to the right withrespect to the ones in Figure 1. m F606W m F814W m F W MS blueHe-b redHe-b TP-AGBRGBRCHB
Figure 3.
F814W versus F606W − F814W color-magnitude dia-gram of UGC 4483 corresponding to the field of view covered byour UVIS imaging (after the quality cuts, see Section 2). The mainstellar evolutionary phases are indicated (see Section 3). follow the H α , which can be due to the fact that the re-cently formed massive stars ionized the gas (thus we seeH α emission) leaving less neutral gas in their surround-ings (and also external effects like a recent interactionmight have influenced the H i gas distribution).Figure 3 shows the F814W versus F606W − F814WCMD of UGC 4483. We notice the extremely high quality of these data, with the main stellar evolutionary phasesclearly recognizable and well separated in the diagram.The CMD shows a continuity of stellar populations, sug-gesting a SF that spans from ancient to recent epochs;we can identify two well separated sequences for MS starsat m F606W − m F814W (cid:46) and stars at the blue edge ofthe BLs (core He-burning intermediate- and high-massstars) at m F606W − m F814W ∼ . , the diagonal fea-ture produced by stars at the red edge of the BLs at m F606W − m F814W (cid:38) . (well separated from the RGB),and a few carbon stars and thermally pulsing asymptoticgiant branch (TP-AGB) stars, outlining the horizontalfeature around m F606W − m F814W (cid:38) and m F814W (cid:46) ;older stars are enclosed in the RGB with its clearly de-fined tip at m F814W ∼ . , in the RC, and possibly inthe HB, our oldest age signature.To guide the eye and explore the age and metallicityof these populations, we over-plot the MIST isochrones(Choi et al. 2016; Dotter 2016) for four different metal-licities ([Fe/H] = − . , − . , − . , − . ) on the CMD(Figure 4). As a reference, from the oxygen abundance ofthe H ii regions we obtain a current metallicity [Fe/H] = − . . The isochrones were shifted according to galaxy’sdistance modulus of DM = 27.45 (which is what we ob-tain from out best-fitting technique, see Section 5), andforeground extinction of E(B − V) = 0 . (Schlafly &Finkbeiner 2011). The different metallicities are labeledat the top of each plot, while ages are labeled in thelegend. We can see how the most metal-poor isochroneset is too blue to fit the youngest stellar populations,in particular to reproduce the color of the BLs, whichinstead starts to be compatible with a metallicity of[Fe/H] = − . . Also RGB stars are significantly red-der than the two most metal-poor isochrone sets shown tar Formation History of UGC 4483 m F606W m F814W m F W [Fe/H] = 2.9
10 Myr20 Myr50 Myr100 Myr200 Myr500 Myr1 Gyr2 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W [Fe/H] = 2.0
10 Myr20 Myr50 Myr100 Myr200 Myr500 Myr1 Gyr2 Gyr10 Gyr1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W [Fe/H] = 1.5
10 Myr20 Myr50 Myr100 Myr200 Myr500 Myr1 Gyr2 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W [Fe/H] = 1.0
10 Myr20 Myr50 Myr100 Myr200 Myr500 Myr1 Gyr2 Gyr10 Gyr
Figure 4.
Same CMD of Figure 3 with MIST isochrones (Choi et al. 2016; Dotter 2016) of 4 different metallicities over-plotted: [Fe/H] = − . (top left), [Fe/H] = − . (top right), [Fe/H] = − . (bottom left), and [Fe/H] = − . (bottom right), shifted to match the distanceand foreground extinction of the galaxy. We can recognize the HB feature in both the [Fe/H] = − . and the [Fe/H] = − .
10 Gyr oldisochrones. here, even though they touch its blue edge. We cannotinstead exclude such low metallicities for older, fainterstars.To understand how the different stellar populations aredistributed within the galaxy, we select age intervals inthe CMD using the same MIST isochrones as a guide, andplot the spatial map of the selected stars. The results areshown in Figure 5. We choose to select upper MS stars(in blue), with ages (cid:46) Myr, blue BL stars (in green),with intermediate ages between ∼ and ∼ Myr,RGB stars (in red), older than ∼ Gyr, RC stars (in magenta), older than ∼ Gyr, and HB stars (in orange),older than ∼ Gyr. The top left panel shows theseselections in the CMD, while in the other panels we plotthe spatial maps of the age-selected stars, as labeled,with contours showing the different regions we selectedwithin the galaxy (see Figure 2).The general trend is that younger stars have more con-centrated and clumpy distributions compared to olderstars, a behaviour expected as stars move out of theirnatal structures as they age. In particular, very youngMS stars are mainly found in the two inner regions (0a
Sacchi et al. m F606W m F814W m F W x y MS ( < 30 Myr) 1000 1500 2000 2500 3000 x y BL (50 500 Myr)1000 1500 2000 2500 3000 x y RGB (2 13.7 Gyr) 1000 1500 2000 2500 3000 x y RC (2 13.7 Gyr) 1000 1500 2000 2500 3000 x y HB (10 13.7 Gyr)
Figure 5.
Top left panel.
Selection of different stellar populations in the CMD: in blue, upper MS stars with ages (cid:46) Myr; in green, BLstars with intermediate ages between ∼ and ∼ Myr; in red, RGB stars older than ∼ Gyr; in magenta, RC stars older than ∼ Gyr; in orange, HB stars older than ∼ Gyr. The other panels show the corresponding spatial distribution of the age-selected stars, aslabeled, with contours showing the different regions of the galaxy (see Figure 2). and 0b) which host the very bright SF regions visible inFigures 1 and 2, where the H ii regions are located. Wealso notice the central holes in these same regions in theRGB, RC, and HB maps, a sign of the strong incomplete-ness of these fainter stellar populations there. In reality,these faint old stars are likely uniformly distributed overthe galaxy. ARTIFICIAL STAR TESTSTo properly characterize the photometric errors and in-completeness of the data, we perform artificial star tests(ASTs) on our images using the dedicated DOLPHOTroutine.We follow a general standard procedure where we adda fake star (for which we know exactly the input positionand magnitudes) to the real images, re-run the photome-try, and check whether the source is detected and its out-put magnitudes. We then repeat the process many times(2 millions in this case) varying the position and mag-nitudes of the input fake stars, to fully map the wholeimage and explore the whole range of magnitudes andcolors covering the observed CMD. We inject each indi-vidual star simultaneously in both F606W and F814Wimages, in order to reproduce a realistic situation and toaccount for the error and completeness correlation in thetwo bands. We consider a star “recovered” in the outputcatalog if the measured magnitude is within 0.75 mag from its input value in both filters, and satisfy all the se-lection cuts applied to the real photometry. Adding onefake star at a time guarantees not to artificially alter thecrowding of the images.Given the very different crowding conditions within thegalaxy, we build the input distribution of stars followingthe surface brightness of the F606W image, to obtaina more accurate estimate of the incompleteness in themost crowded regions that would be under-sampled witha uniform input distribution.From the output distribution of the recovered stars wederive an estimate of the photometric error (from the m output − m input versus m input distribution, Figure 6)and completeness (from the ratio between the numberof output and input stars, Figure 7) as a function ofboth space and magnitude. Generally speaking, a star isconsidered recovered if its output flux agrees within 0.75mag with its input value. We also consider “lost” thestars that do not pass the same quality tests as the realdata. The distribution of the photometric error displayedin Figure 6 also allows to take into account the effect ofblending of multiple sources on the photometry; in fact,a systematic negative skewness in the output − input flux(stars found brighter than their input) is a signature ofoverlapping of artificial stars with other ones. We takethis skewness into account to consider the blending effectwhen we create the synthetic CMDs. tar Formation History of UGC 4483
18 20 22 24 26 28 30 m F606W, inp m o u t m i n p
20 22 24 26 28 m F814W, inp m o u t m i n p Figure 6.
Photometric errors in F606W (top panel) and F814W(bottom panel) from our artificial star tests; the contours indicatethe 1 σ , 2 σ and 3 σ levels of the distributions.
18 20 22 24 26 28 30 m F606W, inp C o m p l e t e n e ss Region 0aRegion 0bRegion 1Region 2Region 3Region 4
20 22 24 26 28 30 m F814W, inp
Region 0aRegion 0bRegion 1Region 2Region 3Region 4
Figure 7.
Completeness in F606W (left panel) and F814W (rightpanel) from our artificial star tests in the various regions of thegalaxy, highlighting the very different crowding conditions frominside out. The dashed horizontal line marks the 50% completenesslevel.
In Figure 7 we plot the completeness as a function ofmagnitude for the different regions in which we dividedthe galaxy (see Figure 2). The very different crowd-ing conditions from inside out are reflected in the dif-ferent completeness behaviours: indeed, the most in-ternal regions are the most crowded, thus incompleteones, while the completeness increases as we move out-wards. The completeness is 50% at m F606W (cid:39) . and m F814W (cid:39) . in the central regions of the galaxy (0aand 0b), and at m F606W (cid:39) . and m F814W (cid:39) in theouter field (Region 4). STAR FORMATION HISTORYWith the photometric catalog and the artificial starcatalog in our hands, we can perform detailed studies ofthe stellar populations and recover the SFH of UGC 4483using the synthetic CMD method (see, e.g., Tosi et al.1991; Gallart et al. 2005; Tolstoy et al. 2009; McQuinnet al. 2010; Weisz et al. 2011; Cignoni et al. 2015; Sacchiet al. 2018): the observed CMD is compared to synthetic ones built from a set of stellar evolution models (evolu-tionary tracks or isochrones) adequately treated to matchthe distance, extinction, and photometric properties ofthe galaxy. Synthetic CMDs, each representing a sim-ple stellar population of fixed age and metallicity, arecreated and used as “basis” functions (BFs), and a linearcombination of these BFs creates a composite populationwhich can represent, with the appropriate weights, anySFH. The weight associated with each BF is proportionalto the number of stars formed at that age and metallicity,and the best-fit SFH is described by the set of weightsproducing a composite model CMD most similar to theobserved one. The best-fitting weights are determinedby using a minimization algorithm to compare data andmodels.We build our models from both the PARSEC-COLIBRI (Bressan et al. 2012; Marigo et al. 2017)and MIST (Choi et al. 2016; Dotter 2016) isochrone li-braries using the following parameters: Kroupa initialmass function (Kroupa 2001) from 0.1 to 350 M (cid:12) ; 30%binary fraction; [Fe/H] from − . to . in steps of . ;DM = 27.60 (Izotov & Thuan 2002); foreground extinc-tion E(B − V) = 0 . (Schlafly & Finkbeiner 2011);photometric errors and incompleteness from our ASTs(see Section 4, and Figures 6 and 7). We also exploredifferent distance/reddening combinations to obtain thebest match with the data.We do not impose a metallicity distribution function,but we require that the metallicity at each time bin can-not be more than 25% lower than the metallicity of theadjacent older bin. This is to include a soft constraintwithout forcing an unknown distribution a priori . Thisis a looser boundary condition with respect to a metallic-ity monotonically increasing with time, often imposed inthis kind of studies; however, in a galaxy as metal-pooras UGC 4483, nobody knows exactly how the metallicityvaries with time, for instance because of the accretionof large amounts of metal poor gas, or because of sig-nificant losses of heavy elements, through galactic windstriggered by powerful supernova explosions.We use the hybrid genetic code SFERA (Cignoni et al.2015) to build the synthetic CMDs and perform the com-parison between models and data. SFERA is built to per-forms a complete implementation of the synthetic CMDmethod, from the MonteCarlo extraction of syntheticstars from the adopted stellar evolution models, to theinclusion of observational effects in the models throughASTs, and finally, performing the minimization of theresiduals between observed and synthetic CMDs, withan accurate estimate of the uncertainties. For the mini-mization, we bin both the synthetic and observed CMDsto compare star counts in each grid cell. Given that weneed to take into account the possible low number countsin some CMD cells, we choose to follow a Poissonianstatistics, looking for the combination of synthetic CMDsthat minimizes a likelihood distance between model anddata: χ P = Nbin (cid:88) i =1 obs i ln obs i mod i − obs i + mod i (1) Notice that the PARSEC-COLIBRI models adopt Z (cid:12) =0 . , while the MIST models adopt Z (cid:12) = 0 . . Sacchi et al. m F606W m F814W m F W Region 0a125 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W Region 0b196 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W Region 1906 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W Region 22368 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W Region 34720 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W Region 45612 stars
10 Myr100 Myr500 Myr1 Gyr10 Gyr
Figure 8.
CMDs of the regions we identified in the galaxy (see Figure 2) and used to recover the SFH (0a and 0b being the innermostones, 4 the outermost one, see Figure 2). Over-plotted in colors are the MIST isochrones (Choi et al. 2016; Dotter 2016) with [Fe/H] = − . and ages as labeled, shifted to match the distance and foreground extinction of the galaxy. where mod i and obs i are the model and the data his-tograms in the i − bin. This likelihood is minimized withthe hybrid-genetic algorithm, i.e. a combination of aclassical genetic algorithm (Pikaia ) and a local search(Simulated Annealing).Statistical uncertainties are computed with a bootstraptechnique on the data, i.e. applying small shifts in colorand magnitude to the observed CMD and re-deriving theSFH for each new version. We average the different so-lutions and take the rms deviation as statistical error onthe mean value. Systematic uncertainties are accountedfor by re-deriving the SFH with different grids to bin ageand CMDs, and using different sets of isochrones. Theseare added in quadrature to the statistical error.Before discussing the details of the SFH in the variousregions, we can have a look at their CMDs which alreadyreveal profound differences within the body of the galaxy.Figure 8 shows the 6 CMDs of these regions (0a and 0b Routine developed at the High Altitude Observatory and pub-licly available: being the innermost ones, 4 the outermost one, see Figure2), with over-plotted in colors the MIST isochrones with[Fe/H] = − . and ages in the range 10 Myr −
10 Gyras labeled, shifted to match the distance and foregroundextinction of the galaxy.Regions 0a and 0b include the most actively starforming areas of the galaxy, where we see H α emissionfrom the H ii regions. Their CMDs show very young( (cid:46) Myr) MS stars, with some BL stars, and veryfew RGB stars. The observed MS in these two regions isa bit redder with respect to the plotted isochrones, andpresumably needs a higher metallicity and/or larger red-dening, which is most likely in these highly star-formingregions. As already pointed out in Section 3, the pres-ence of very bright young stars makes the incomplete-ness severe here, hindering a proper characterization ofall the stellar populations present in these regions. This,together with the low number of stars observed, preventsus from running our SFH procedure in a statistically sig-nificant way.Going outwards from Region 1 to Region 4, we see a tar Formation History of UGC 4483 − m F606W − m F814W m F W data : region 1 − m F606W − m F814W model : MIST − m F606W − m F814W residuals − [ F e / H ] Age [yr]0.00.51.01.52.02.5 S F R [ M / y r ] region1 - colibriregion1 - mist − m F606W − m F814W m F W data : region 2 − m F606W − m F814W model : MIST − m F606W − m F814W residuals − [ F e / H ] Age [yr]0.00.20.40.60.81.01.2 S F R [ M / y r ] region2 - colibriregion2 - mist − m F606W − m F814W m F W data : region 3 − m F606W − m F814W model : MIST − m F606W − m F814W residuals − [ F e / H ] Age [yr]0.00.10.20.30.40.50.6 S F R [ M / y r ] region3 - colibriregion3 - mist − m F606W − m F814W m F W data : region 4 − m F606W − m F814W model : MIST − m F606W − m F814W residuals − [ F e / H ] Age [yr]0.000.050.100.150.200.250.300.350.40 S F R [ M / y r ] region4 - colibriregion4 - mist Figure 9.
Left panels.
Hess diagrams of the CMDs for the different regions (1 to 4 from top row to bottom row) of UGC 4483: theobservational diagram is on the left and the one reconstructed on the basis of the MIST models in the middle, while on the right we showthe residuals between the two in terms of the likelihood used to compare data and models with SFERA; the black line shows the 50%completeness limit, used as a boundary for the SFH recovery.
Right panels.
Recovered SFH from the two adopted sets of models (COLIBRIin blue, MIST in orange) with 1 σ error bars which include both random and systematic uncertainties. Sacchi et al. clear evolution of the CMD which becomes progressivelyolder: the MS is less and less populated, while the RGBand RC features start to stand out. In Region 4 we reacha 50% completeness at m F814W ∼ . These differenceswill be reflected in the radial SFH of the galaxy.5.1. Region 1
Region 1 includes the area just outside the H ii regionsof UGC 4483. Its CMD contains a variety of stellar pop-ulations, suggesting a continuous SF from ancient to re-cent epochs. Figure 9, top row, shows the Hess diagram(i.e. the density of points) of the CMD, with the obser-vational diagram being on the left, the one reconstructedon the basis of the MIST models in the middle, and theresiduals between the two on the right; the black lineshows an estimate of the 50% completeness limit, usedas a boundary for the SFH recovery. The rightmost panelcontains the best-fit SFH from the two adopted sets ofmodels (blue for COLIBRI, orange for MIST).As expected from the CMD analysis, we see a preva-lence of young SF, even though the low number of starsand severe incompleteness result in huge errors on ourrates. The models perfectly reproduce the good sepa-ration between the MS and blue BL sequences, and be-tween the red BL and RGB sequences, a great value ofthese excellent data, and a significant improvement withrespect to previous observations.The two sets of models reasonably agree, in particularin the best constrained intermediate-age bins, between ∼ Myr and ∼ Gyr. In both cases, the best-fit SFHswere obtained with an additional internal reddening of
E(B − V) = 0 . , i.e. A F606W = 0 . and A F814W =0 . . 5.2. Region 2
The results for Region 2 are shown in the second rowof Figure 9. From the Hess diagram we can see a clearshortening of the bright MS and a higher density of olderBLs, while the RGB starts to become more populatedand more carbon stars and TP-AGB stars appear. Thisis well reflected in the SFH, whose peak is between ∼ and ∼ Myr ago, with rates that stay high until ∼ Gyr ago, in both solutions. We still see SF in older timesbins, but the increasing incompleteness and the worsetime resolution make our results at the older epochs lessreliable.The internal reddening we find in this region is
E(B − V) = 0 . , i.e. A F606W = 0 . and A F814W =0 . . 5.3. Region 3
The decreasing young-to-old SFR trend continues inRegion 3, as shown in Figure 9, third row. We find aroughly constant SF, with some ups and downs in par-ticular in the COLIBRI solution. The internal reddeningfor this region is
E(B − V) = 0 . , i.e. A F606W = 0 . and A F814W = 0 . , which is lower that in the Regions 1and 2, where the higher recent SF activity creates moredust, thus we can expect a more severe extinction.Interestingly, the two solutions show a significant dis-crepancy in the last time bin, with the COLIBRI modelsproviding a rate about 3 times higher than the MISTmodels. This is most likely a systematic difference in − m F606W − m F814W m F W data : region 4 − m F606W − m F814W model : MIST − m F606W − m F814W residuals − Figure 10.
Best-fit solution for Region 4, but with the SFH forcedto stop 8 Gyr ago (to be compared with Figure 9, bottom row). the stellar tracks, and in the way that the most uncer-tain stellar evolution parameters (like overshooting) aretreated. It is worth noticing that the COLIBRI tracksalso produce a dip in the SFH between ∼ and ∼ Gyr ago that is much less pronounced in the MIST mod-els, maybe balancing that old peak at ages older than10 Gyr ago. This is why it is important to use differentsets of models, to compare their systematics and providefeedback to keep improving our knowledge on stellar evo-lution.We also stress that this age range is constrained bystars at the edge of the 50% completeness, and shouldbe interpreted with caution.5.4.
Region 4
Region 4 includes the most external part of UGC 4483,and our most complete data set. The SFH here, shownin the bottom row of Figure 9, is compatible with 0 backto ∼ Myr ago, and its peak is at ages older than ∼ Gyr ago, according to both solutions. We find an inter-nal reddening of
E(B − V) = 0 . , i.e. A F606W = 0 . and A F814W = 0 . . This is lower than in Regions 1 and2 but counter-intuitively a bit higher than in Region 3;however, the sensitivity of the CMD-fitting procedure toreddening can be degenerate with distance, so these val-ues should be taken as an indication. Moreover, this red-dening difference corresponds to a magnitude differenceof about 0.05 mag, which gives a distance difference ofabout 10 pc, consistent with a distance difference alongthe line-of-sight between the body and the halo of thegalaxy. The very deep CMD reaches 50% of complete-ness at m F814W ∼ , thus fainter than where we expectto see the HB of the galaxy (see the isochrones in Figure4). Despite the systematic differences we just discussed,both sets of models show a consistent SFR at all ages,including those older than 10 Gyr ago, putting an im-portant constraint on the age of this galaxy outside theLocal Group.Since we are working at the edge of the data complete-ness, we tested the SFH recovery forcing the SF to startonly 8 Gyr ago, instead of 13.7, to compare the resultsand check if such a SFH would still be compatible withthe observed CMD. The result of this test is shown in Fig-ure 10. Even though the difference with Figure 9 mightnot seem significant, the recovered CMD, residuals, andoutput likelihood of this solution are significantly worsethan the case shown in Figure 9, bottom row. Not onlythe faintest part of the RC/HB is not well reproduced, tar Formation History of UGC 4483 . ± . , which is slightly lower than theone provided by Izotov & Thuan (2002), but consis-tent with the one provided by Dolphin et al. (2001), i.e. . ± . .A summary of the derived SFRs and stellar massesformed at various epochs in the different regions of thegalaxy is given in Table 1. In particular for Regions 3and 4, we can notice a significant difference in the epochof the main SF peak derived with the COLIBRI or theMIST models. As already discussed in Section 5.3, the50% completeness in Region 3 is at the edge of the HB orthe faintest part of the RC phases, thus systematic uncer-tainties in the stellar evolution models can greatly affectthe resulting age and rate determination, depending onhow much of the synthetic HB falls within or outsidethe CMD region covered by the photometry. For Region4, the situation is better, thanks to the lower crowding;despite the systematics still affecting the fit, Figure 9(bottom right panel) shows a very good agreement be-tween the two solutions, and the difference in age peak reported in Table 1 ( − Gyr for the MIST models,versus − . Gyr for the COLIBRI models) is simplythe result of small variations of the SFRs in the last 3age bins.5.5.
Searching for RRL variable stars in UGC 4483
As already mentioned, the detection of RRL stars ina galaxy can unambiguously confirm the presence of anHB, and unveil a stellar population at least ∼ Gyrold. RRLs are ∼ mag brighter than coeval turnoffstars, and the typical form of their light variation makesthem easily recognized even in very crowded fields andin galaxies where a 10 Gyr population may be buriedinto the younger stars. To look for these variables inUGC 4483, the observations have been reduced as a timeseries to provide a full mapping of the classical instabil-ity strip of pulsating variables with periods from abouthalf a day to a couple of days. A detailed analysis ofthe candidate RRLs and theirs light curves will be thesubject of a forthcoming paper (Garofalo et al., in prepa-ration), while here we simply mention what is relevantto the present paper.In our analysis, we found six stars with light curvescompatible with those of RRL variables. In particular,two of them are in the most external region of the galaxywe analyzed, Region 4, where crowding conditions areless severe and the completeness much more favorablethan in more internal regions. Within the uncertain-ties of the HB modeling, these candidates also have col- m F606W m F814W m F W [Fe/H] = 2.5 dataMIST, 10 GyrCOLIBRI, 10 Gyrcandidate RRLs 1.0 0.5 0.0 0.5 1.0 1.5 2.0 m F606W m F814W m F W [Fe/H] = 2.0 dataMIST, 10 GyrCOLIBRI, 10 Gyrcandidate RRLs Figure 11.
F606W versus F606W − F814W color-magnitude di-agram of UGC 4483 in grey, the six RRL candidates we found ingreen, and the MIST (in red) and COLIBRI (in orange) 10 Gyrold isochrones for a metallicity of [Fe/H] = − . (left panel) and[Fe/H] = − . (right panel). The isochrones were shifted to matchthe distance and foreground extinction of the galaxy. Age [yr]0.000.250.500.751.001.251.501.752.00 S F R [ M / y r / k p c ] mist region1region2region3region4 Figure 12.
SFR surface densities (SFR/area) as a function oftime in different regions of UGC 4483, from the MIST solutions. ors and magnitudes compatible with being indeed RRLs,as shown in Figure 11. In the left panel, we plot theF606W versus F606W − F814W color-magnitude diagramof UGC 4483 in grey, the six candidates in green, andthe MIST (in red) and COLIBRI (in orange) 10 Gyrold isochrones for a metallicity of [Fe/H] = − . . Wecan notice how different the models are, with the MISTHB being more than half a magnitude brighter than theCOLIBRI one. At this metallicity, our candidates fallbetween the two sets of HB. As shown in the right panel,the disagreement between the isochrones starts to dis-appear at [Fe/H] = − . , and most of our candidateswould fall slightly brighter and bluer than the HB, butstill compatible with it once the uncertainties on bothdata and models are considered. DISCUSSIONThe Local Group is an incredible environment, fullof interesting processes and galaxies of many differenttypes. However, since there are not known BDCs withinit (except for IC 10, sometimes considered as a such),in order to study this class of very actively star-forming,low-metallicity, systems the only possibility is to pushour limits beyond its borders. Even though the dis-tance makes observations challenging even for
HST , withthe photometry reaching only the brightest evolutionaryphases, many studies tried to characterize the formationand evolution of these intriguing objects.2
Sacchi et al.
Table 1
Summary of the derived star formation rates and stellar masses in the different regions of UGC 4483. ( ∗ ) The last column is the ratio between the two previous ones.region (cid:104)
SFR (cid:105)
SFR peak age peak M ∗ (age ≤
50 Myr) M ∗ (age > young/old SFR ( ∗ ) [M (cid:12) / yr / kpc ] [M (cid:12) / yr / kpc ] [10 M (cid:12) ] [10 M (cid:12) ] [ − ] COLIBRI models (1 . ± . × − (2 . ± . × − − Myr . ± .
79 1 . ± . (1 . ± . × − (4 . ± . × − − Myr . ± .
92 1 . ± . (6 . ± . × − (1 . ± . × − − . Gyr . ± .
45 3 . ± . (4 . ± . × − (7 . ± . × − − . Gyr . ± .
30 2 . ± . MIST models (1 . ± . × − (1 . ± . × − − Myr . ± .
94 1 . ± . (8 . ± . × − (5 . ± . × − − Myr . ± .
93 1 . ± . (5 . ± . × − (1 . ± . × − − Myr . ± .
50 2 . ± . (4 . ± . × − (8 . ± . × − − Gyr . ± .
23 2 . ± . Among these works, some embarked on the challengeof applying the synthetic CMD method to the observedCMDs, trying to reconstruct the galaxies’ SFHs withinthe reachable lookback time (see, e.g., Aloisi et al. 1999for I Zw 18, Schulte-Ladbeck et al. 2001 for I Zw 36, orAnnibali et al. 2003 for NGC 1705). Despite the obviousdifferences among the individual SFHs, all these dwarfsshow a qualitatively similar behavior, with a strong on-going SF activity and a moderate and rather continuousstar formation at older epochs. Most notably, they allshow an RGB, meaning that they contain stars as oldas the lookbacktime reached by the photometry. Find-ing old stars in these extremely metal-poor galaxies isa key information to understand how they compare toother systems and to place them in the general contextof galaxy formation and evolution. For this reason, itis crucial to try to reach even older stellar evolutionaryfeatures, to finally figure out whether BCDs are reallyold systems as currently believed.We presented here a detailed analysis of UGC 4483,the closest example of the BCD category, which we stud-ied thanks to deep new data obtained with the exactpurpose of reaching stars older than a few Gyr. Our newWFC3/UVIS observations reach more than 4 mag deeperthan the TRGB, and allow us to detect and characterizefor the first time the population of core He-burning starswith masses (cid:46) M (cid:12) .To take into account the different crowding and SF con-ditions across the galaxy, we used isophotal contours ofthe F606W image to divide it into sub-regions, as shownin Figure 2; the resulting SFHs from the MIST solutionare shown in Figure 12, where the rates have been nor-malized to each region’s area to facilitate the compari-son. As expected for a BCD, the main recent activityis in the central Region 1, which, despite the uncertain-ties, shows a continuous increase of SF in the most re-cent time bins, with a rate about 10 times higher thanthe average. Going outwards, the rate densities gener-ally decrease, together with the ratio of present-to-pastaverage SF (see the last column of Table 1 for the de-tails). Indeed, Regions 3 and 4 had already formed 50%of their total stellar mass around ∼ Gyr ago, while thisoccurred between 2 and 3 Gyr ago for Regions 1 and 2, confirming that the SF proceeded from outside in.Even if we could not derive the SFHs there, it is clearthat this trend still holds when we consider the two in-nermost regions of the galaxy, 0a and 0b, sites of verybright and active H ii regions. In particular, Region 0ahosts a super star cluster with absolute magnitudes of M V = − . ± . and M I = − . ± . , corre-sponding to an age of ∼ Myr and an initial mass of ∼ M (cid:12) , as derived by Dolphin et al. (2001). For com-parison, this is just slightly fainter than 30 Doradus, thestar forming region within the Large Magellanic Cloud.Using the equivalent width of the H β emission line, Izo-tov & Thuan (2002) derived a much younger age of 4Myr.In the outer part of the galaxy, Region 4, we find theSF peak at ages (cid:38) Gyr, and the rates are consistentwith a SFH starting 13.7 Gyr ago. This is the first timethat we can put such a strong constraint on a galaxyoutside the Local Group, and despite the uncertaintiesand systematics of the stellar evolution models, we be-lieve this to be a robust result which is supported alsoby the detection of a number of candidate RR Lyraestars. The implication is that UGC 4483, and possiblyall BCD galaxies, are indeed as old as a Hubble time,and not young systems as previously believed. Indeed,despite the very young recent activity, about 87% of thetotal stellar mass of the galaxy was formed at ages olderthan 1 Gyr.The color extension of the identified HB region includesthe expected location of the RR Lyrae gap, correspond-ing to the pulsation instability strip. Follow-up time-series observations would be very useful to definitely con-firm our detection in UGC 4483 (Garofalo et al., in prepa-ration) of candidate pulsating stars of this class (unam-biguous tracers of a (cid:38) Gyr old population) and tobe able to use the inferred pulsation characteristics (pe-riods, amplitudes, etc.) to trace the properties of theoldest stellar population in UGC 4483.It is interesting to notice that, as far as the SFH isconcerned, UGC 4483 is fairly similar to many dIrrs, inspite of being classified a BCD: it shows a rather contin-uous SF, with no evidence of short, intense bursts, andno long quiescent phases, at least within the limits of our tar Formation History of UGC 4483
HST /WFPC2, Dolphin et al. (2001) derived aSFR of . × − M (cid:12) / yr , which is almost a factor oftwo higher than ours, (7 . ± . × − M (cid:12) / yr . How-ever, their CMD is not particularly deep (the fainteststars have an I magnitude of 25, with much larger pho-tometric errors), so it is likely that they are biased bythe brightest young stars in the upper part of the CMD.On the other hand, our rate is a lower limit as it is basedon Regions 1 to 4, thus excluding the most active Re-gions 0a and 0b which have not enough stars to run afull SFH derivation. Based on the same data, McQuinnet al. (2010) derived a more detailed SFH, which shows avery recent SF peak (within the last 50 Myr), consistentwith what we find in Region 1, and a rather continu-ous activity up to 1 Gyr ago. Their peak rate, how-ever, is about . × − M (cid:12) / yr , while we reach at most . × − M (cid:12) / yr when considering the uncertainties.This can be again the result of the different sampling ofthe stellar populations, in particular in the most centralregions of the galaxy. They also find an old population,which is consistent with the RGB phase revealed by theirCMD, and confirmed by our deeper one. Also Weisz et al.(2011) used the same data and code to derive the SFH ofthis galaxy, obtaining similar results (their average SFRof . × − M (cid:12) / yr ).From the H α luminosity (Gil de Paz et al. 2003;Kennicutt et al. 2008), we can infer the current ( (cid:46) Myr) SFR following the prescription of Murphy et al.(2011), which assumes a Kroupa IMF. We obtain . × − M (cid:12) / yr , compatible within the error bars with ourresults in the most recent time bin of Region 1, but stillhigher, as expected as the H α mainly traces the emissionfrom the two most active Regions 0a and 0b (see Figure1, right panel).Finally, from the analysis of VLA observations, Lelliet al. (2012b) found that UGC 4483 has a steeply-risingrotation curve, making it the lowest-mass galaxy witha differentially rotating H i disk. These rotation-velocitygradients are directly related to the dynamical mass sur-face densities, and signal a strong central concentrationof mass, also found in other BCDs, like I Zw 18 (Lelliet al. 2012a), NGC 2537 (Matthews & Uson 2008), andNGC 1705 (Meurer et al. 1998; see also Figure 10 of Lelli et al. 2012b). They also showed that the centralmass concentration cannot be explained by the newlyformed stars or by the concentration of the H i , imply-ing that either the progenitors of BCDs are compact,gas-rich dwarfs, or there must be a mechanism (exter-nal, such as interactions and mergers, or internal, suchas torques from massive star-forming clumps, Elmegreenet al. 2012) leading to a concentration of gas, old stars,and/or dark matter, causing the SF increase. SUMMARY AND CONCLUSIONSHere we summarize the main results of this paper,where we presented new WFC3/UVIS data of the BCDgalaxy UGC 4483, and investigated its resolved stellarpopulations and radial SFH.- The CMD of UGC 4483 is populated by many gen-erations of stars, from young MS and BL stars, tointermediate-age AGB and RGB stars, and olderRC and possibly HB stars. In particular, we wereable to reach and detect for the first time thepopulation of core He-burning stars with masses (cid:46) M (cid:12) .- The stellar populations in the galaxy have the typ-ical distribution of star-forming irregular galaxies,with the youngest stars being more centrally con-centrated and close to the inner star-forming re-gions, while the distribution becomes more andmore uniform as the stars age.- From the SFH recovered in different radial regionsof the galaxy, we found a declining trend of theoverall SF activity and of the present-to-past SFRratio going from inside out. In all regions, wefound the best fit SFH using a distance modulusof DM = . ± . , slightly lower than previousestimates.- Using the synthetic CMD method, we determinedan average SFR over the whole Hubble time of (7 . ± . × − M (cid:12) / yr , corresponding to a to-tal astrated stellar mass of (9 . ± . × M (cid:12) ,87% of which went into stars at epochs earlier than1 Gyr ago. These are lower limits, as they do notinclude the two innermost regions of the galaxy,hosting very active H ii regions and a super starcluster almost as luminous as 30 Doradus.- We found strong evidence of a (cid:38) Gyr old popu-lation, which might be responsible for the presenceof a blue HB, as also suggested by our detectionof a number of candidate RRL variable stars inUGC 4483.These data are associated with the
HST
GO Program15194 (PI A. Aloisi). Support for this program was pro-vided by NASA through grants from the Space Tele-scope Science Institute. F.A., M.C., and M.T. acknowl-edge funding from the INAF Main Stream program SSH1.05.01.86.28. We thank the anonymous referee for theuseful comments and suggestions that helped to improvethe paper. REFERENCES4