We also compare ISM abundances obtained from H II region observations. It is particularly interesting to compare Sextans A with Pegasus. While Sextans A shows a relatively high star formation rate recently, Pegasus has been relatively quiescent. We find that the ISM in Pegasus shows a relatively high N/O ratio, while Sextans A shows a relatively low N/O ratio. These results are consistent with the hypothesis that oxygen production is dominated by high mass stars, while nitrogen production is dominated by intermediate to low mass stars. Thus, the star formation history as recorded in the stars is consistent with the star formation history inferred from the ISM abundances. It appears that a good lower limit on the ability of oxygen to mix back into the cold/warm ISM is about 10 Myr and a good upper limit may be about 100 to 200 Myr. A good upper limit on the ability of nitrogen to mix back into the cold/warm ISM may be about 400 Myr.
Much has been learned from ground-based observations (e.g., Tosi et al. 1991; Greggio et al. 1993; Gallart et al. 1996a,b,c). These researchers have developed methods of comparing synthetic color magnitude diagrams (CMDs) to the observations, and these methods are successful in re-creating the distribution of the stars. Of course, the diagnostic power of these comparisons are limited by the quality of the photometry and ground based observations are severely limited by crowding (for distances greater than 0.5 Mpc). Important progress has been made by Tolstoy & Saha (1996) who have developed a method which allows them to quantify the differences between Monte Carlo simulations and the observed distributions of stars in CMDs (e.g., Tolstoy 1996).
A tremendous advance in this field has come about due to the fantastic imaging abilities of the Hubble Space Telescope. This has resulted in great improvements in the accuracy of the photometry of individual stars. As we will demonstrate below, this allows us to produce very detailed recent SFHs for nearby galaxies. The great potential for understanding chemical evolution becomes clear. If we have galaxies for which we have detailed SFHs, then we can observe the effects of different SFHs on chemical abundances. Eventually, we hope that this will allow us to use relative chemical abundances to constrain SFHs in distant systems where individual stellar measurements are not possible.
Four nearby dwarf irregular galaxies were observed in a cycle 5 HST program. Program collaborators include Chiosi (Padua), Dufour (Rice), Gallagher, Hoessel, Cole (Wisconsin), Mateo (Michigan), Saha (NOAO), and Tolstoy (ESA). Results from these observations are published (Dohm-Palmer et al. 1997a,b, 1998a; Gallagher et al. 1998), and forthcoming (Tolstoy et al. 1998; Dohm-Palmer et al. 1998b).
HST observations were obtained on 1 December 1995, and consisted of WFPC2 images in three filters: F439W (4000 s), F555W (1800 s) and F814W (1800 s). The stellar photometry was extracted using a modified version of the PSF fitting program DoPHOT (Schechter et al. 1993; Saha et al. 1996). Artificial star tests showed the data to be 50% complete down to V = 25.9, B = 25.5, and I = 25.2.
The CMD for Sextans A is shown in Fig. 1. The remarkable feature of this CMD is that the blue plume has been resolved into two separate populations. The well-defined MS contains stars as young as 10 Myr. Just redward of the MS is a clearly separate population that corresponds with massive stars in their core Helium burning phase (HeB stars). This is the bluest extent of the so-called ``blue-loop'' phase. The magnitude of the HeB stars is well correlated with age over the past 1 Gyr (Bertelli et al. 1994). Thus, we can use the HeB stars to trace the recent SFH within the galaxy. We also see a population of massive, evolved, red supergiants that are counterparts to the blue HeB stars.
In addition to the young populations, there is an older population that includes a well-defined RGB and a few AGB stars. Just at the limit of photometry are hints of the red clump stars and horizontal branch.
We have mapped the spatial density distributions of several populations (Fig. 2). The young association in the lower right is densely populated with MS stars. This association is near to a dense clump of H I (Skillman et al. 1988) and several bright H II regions (Hodge et al. 1994). There is also a low density bridge of young stars that extends to the left. This bridge does not align with the pattern of H I gas or the H II regions.
The blue HeB stars are not as strongly clustered in the young association, but do reside in the low density bridge to the left. These stars are, on average, older than the MS stars, and thus indicate where the star formation has occurred in the recent past.
The distribution of the RGB stars is quite different from that of the younger stars. They form a bar through the middle of the galaxy, corresponding to an H I depression. The bar may be related to the dynamical bar postulated by Skillman et al. (1988). The evidence for the presence of a bar consisted of a distorted H I velocity field. Note that the optical evidence for a bar appears only when the high surface brightness youngest star populations are subtracted away.
Skillman et al. (1989) find the H II regions of Sextans A to have a metal content ∼4% of Solar (Z = 0.0008). There are two appropriate sets of models available for comparison. The Geneva group provides evolution tracks for Z = 0.001 (Schaller et al. 1992). The Padua group provides evolution tracks for Z = 0.0004 (Fagotto et al. 1994), and isochrones for Z = 0.001 (Bertelli et al. 1994). Both sets of evolution tracks need to be converted from Teff and L into the observed Johnson B and V, and the Cousins I passbands (the Padua isochrones are given in these passbands). This was done using the model atmospheres of Kurucz (1992) at the appropriate metallicity. These models do not apply to M stars, however, we are not concerned with such low mass stars in this paper.
Figure 3a compares the Padua models for Z = 0.001 with the observed stars. Galaxies have a continuous age distribution of stars, so the isochrones are not meant to represent the entire CMD. Rather, it is instructive to use the isochrones to identify the age ranges in which the galaxy may have been forming stars. Note that the red extreme of the ``blue-loop'' phase aligns better with the red supergiants, but not perfectly for the handful of very bright stars. In addition, the shape of the RGB matches well the observed shape, and the MS aligns with the locus of observed MS points.
There is clearly a range in ages for the observed stars of Sextans A. Rough comparisons with the isochrones indicate stars as young or younger than 10 Myr, as well as a population several Gyr old that produces the RGB.
Comparison of mixed age populations with theoretical isochrones can be somewhat deceiving because the stars do not equally populate each isochrone. Rather, they spend more time in certain phases. For example, stars are more likely to be found at the extremes of the ``blue-loops'', rather than in between. The isochrone tables provided by Bertelli et al. (1994) allow us to extract the color and magnitude for these populations as a function of age. Figure 3b shows the path of the MS turnoff, the blue HeB stars, and the red HeB stars through the V and I CMD. The three populations align very well with the observed stars. The exception, as noted above, is the red HeB, which is slightly bluer than the few observed stars for MV < -4. Note that the stellar models for the metallicities determined from the H II regions give a very good representation of the positions of the young stellar populations in the CMDs. This gives us increased confidence in the use of stellar models to determine star formation histories.
There are two main populations which we can use to glean the recent SFH. Traditionally the MS is used, specifically using the strength and position of MS turnoffs. This works very well for isolated generations of stars, such as in clusters. However, the MS in galaxies with continuous star formation can be difficult to interpret because each successive generation lies on top of the previous. We can, however, examine the luminosity function of the MS to look for gaps, flat regions and other features that may give a clue to the SFH (e.g., Butcher 1977). In addition, we can use statistical methods for extracting SFH's from these chronologically over-lapping populations (e.g., Tolstoy & Saha 1996).
The blue HeB stars provide a parallel track to the MS in which to observe star formation events. From the number of stars at each magnitude, we can calculate the SFR for the age corresponding to the time it takes stars to reach this phase of evolution. There are two advantages to using the blue HeB as an indicator of SFH: (1) The blue HeB stars are about 2 magnitudes brighter than the MS turnoff stars (MSTO) of the same age (e.g., Bertelli et al. 1994). This allows us to probe the recent SFH further back in time (for the same photometric limits) than we can from MS turnoffs; (2) There is little confusion from overlapping generations. All of the blue HeB stars of a certain magnitude come from the same generation of stars (there is confusion from stars leaving the MS, but only at the few percent level). In principle, the blue HeB stars can probe the SFH back to almost 1 Gyr. At older ages, the blue HeB stars blend with the red clump and horizontal branch, becoming degenerate in time.
There are four relations needed to calculate the SFR from the blue HeB stars.
We have applied the normalization correction described above to the blue HeB luminosity function in order to derive the SFH for Sextans A for the last 600 Myr (Fig. 4). There is a peak in the SFR of 7500 Msun Myr-1 kpc-2 at ∼30 Myr. Within errors, this agrees with the peak in the SFR determined from the MS stars. Beyond 50 Myr, the SFR falls off to approximately 1000 Msun Myr-1 kpc-2 by 100 Myr ago and remains roughly constant. The limit of 600 Myr is determined by both photometric errors and incompleteness corrections.
We can compare the current SFR with the average SFR over the history of the galaxy. We estimate the total number of stars in the galaxy based on the luminosity and colors. For this, we use the population synthesis models of Bressan et al. (1994). They present the integrated colors and magnitude for single age stellar populations at a variety of metallicities. We used the Z = 0.001 (5% Solar) models. To calculate the integrated colors and luminosity, we assume the SFH derived from the blue HeB stars. Sextans A has (B-V) = 0.26 and MV = -14.13, adjusted for (m-M) = 25.8 (Hunter & Plummer 1996). For this color, the population synthesis models indicate an age of 10.8 Gyr and M/LV = 0.43 Msun/LV,sun. This value is subject to systematic effects, most notably the choice of SFH can alter this by up to a factor of 2. The magnitude of Sextans A implies a luminosity LV = 3.8·107 LV,sun. So the total mass in stars is estimated to be 1.6·107 Msun.
If the galaxy has been forming stars for 10.8 Gyr, this gives an average SFR/area of 310 Msun Myr-1 kpc-2. So between the ages of 100 and 600 Myr, Sextans A appears to be forming stars at a rate which is a factor of ∼6 larger than the lifetime average. This implies that the SFR prior to 1 Gyr ago must have been much lower than which we measure for the past 600 Myr. The past 50 Myr is clearly a period of highly enhanced star formation by a factor of 20 or more compared to the average.
We conclude that Sextans A is currently undergoing a strong episode of star formation that started ∼80 Myr ago. Prior to this episode, the galaxy formed stars at a lower and relatively constant rate which was still well above the lifetime average for the galaxy.
In order to study the spatial variation of star formation in Sextans A, we have calculated the spatial distribution of blue HeB stars in a series of different age frames separated by 10 Myr. Figure 5 shows a series of these age frames. This shows us the location of star formation within Sextans A over the past 600 Myr. The oldest star formation shown, at 570 Myr, appears as a diffuse region in the lower part of the image. As time progresses, we see a patch of star formation move in from the lower left corner. This appears to propagate toward the center of the image, peaking at about 370 Myr. This event then fades away as the young association in the lower right comes to prominence at 200 Myr. When viewed as a movie (Dohm-Palmer et al. 1998a; or http://ast1.spa.umn.edu/robbie/research/sexa/movie.html) these frames suggest that the star formation may have propagated from region to region, as low level bridges connect them.
|Sex A||1.4||2.4||-13.91||0.26||1·1039||1.4||7.49||1, 2, 3, 4, 5|
|GR 8||2.2||0.69||-12.12||0.38||2·1036||0.76||7.63||2, 6, 7, 8|
|Peg DIG||0.76||1.1||-11.43||0.54||2·1036||0.47||7.93||2, 9, 10, 11, 12|
|Leo A||0.69||1.0||-11.52||0.26||1·1036||1.3||7.30||2, 5, 13, 14, 15|
|References: (1) Sakai et al. 1996; (2) de Vaucouleurs et al. 1991; (3) Hunter & Plummer 1996; (4) Skillman et al. 1988; (5) Skillman et al. 1989; (6) Tolstoy et al. 1995; (7) Carignan et al. 1990; (8) Hodge et al. 1989; (9) Gallagher et al. 1998; (10) Hunter et al. 1993; (11) Lo et al. 1993; (12) Skillman et al. 1997; (13) Tolstoy et al. 1998; (14) Strobel et al. 1991; (15) Young & Lo 1996.|
There are several noteworthy features. First, the Halpha luminosity is much higher for Sextans A than it is for the others. The Halpha is produced by the very youngest star formation (< 3 Myr; Hunter & Gallagher 1986), which is far higher in Sextans A than in the other three galaxies.
Second, there is a trend with MH I/LB. The higher the gas mass to stellar mass fraction the higher the average SFR/area. This is expected. Assuming a similar star forming efficiency for all dIs, the more gas contained in the galaxy the higher the rate of star formation.
Third, there are no large bursts in the recent past of any of the galaxies. Even the most recent event in Sextans A jumps only a factor of three in SFR/area. There are no bursts with strength matching that of the blue, compact, starburst galaxies, which have a SFR ranging from roughly 104 to 107 Msun Myr-1 (see review by Thuan 1991). With the four galaxies combined, our observations cover 2 Gyr of the recent past. Assuming a burst length of 100 Myr (NGC 1569; Vallenari & Bomans 1996), we observe no bursts during 20 possible burst episodes. We conclude that such strong bursts in the current epoch are rare, occurring in less than 5% of dI galaxies. This statistic is based on few galaxies, and observations of more galaxies are needed.
Finally, the star formation process is very similar in all four galaxies. They all show that the star formation occurs at a low level throughout the galaxy with regions of enhancement that are of order 100 pc across and last of order 100 Myr. This pattern was hinted at by the distribution of current star formation in dI galaxies (e.g., Hunter & Gallagher 1986), although such indicators, including Halpha and bright blue stars, are only sensitive to time scales of a few Myr.
Figure 7 compares the CMDs in V and I for Sextans A and Pegasus. Note, that the two galaxies are at comparable distances (1.4 and 0.7 Mpc), so the CMDs have similar photometric limits and direct comparisons are possible. The relatively faint MS and the near absence of blue HeB stars shows that while Pegasus has had some star formation over the last 100 Myrs, the recent star formation rate is clearly depressed relative to that in Sextans A.
Figurer 8 compares the SFHs for Sextans A and Pegasus constructed from the blue HeB stars. In Fig. 8 (right), we have constructed the SFH over the last 800 Myr for Pegasus from its HeB stars. Note that the SFH is probed to earlier times for Pegasus due to its closer distance. Comparing the two, we see that the average star formation rate in Pegasus, over the last 400 Myr, has been roughly about one-third that observed in Sextans A.
Figure 9 shows the position of Pegasus in the abundance - luminosity plane in comparison to the well defined sample of dIs assembled by Richer & McCall (1995) (with alterations for new distances as described in Skillman et al. 1997). The solid line in Fig. 9 is the least squares fit to the more luminous galaxies calculated by Richer & McCall (1995) (note that a new distance for Leo A determined from our WFPC2 imaging places it closer to the general trend, Tolstoy et al. 1998). In Fig. 9 it can be seen that Pegasus fits well in the trend established by other, well studied dwarf irregular galaxies, but lies near the top of the distribution. If the metallicity - luminosity relationship is physically based on a metallicity - mass relationship (cf. Skillman et al. 1989), then the position of Pegasus may be due to the lack of recent star formation. This results in a lower blue luminosity per unit mass when compared to more actively star forming galaxies.
Significant fading and reddening during a quiescent phase is most likely for the lowest mass dwarf galaxies. The higher mass galaxies have, in general, higher metallicities, higher surface brightnesses, and lower gas mass fractions, implying that a current generation of star formation will have a relatively small effect on the present luminosity. However, a burst of star formation can significantly enhance the total luminosity of a faint dwarf galaxy. Due to the lack of a strong underlying stellar population, as the current burst ages, the fading of a very low luminosity dwarf galaxy will be greatest, thus, leading to a larger scatter in the low luminosity systems (as noted by Richer & McCall 1995).
In order to test whether fluctuations in the mass/light ratios of the dIs are the dominant source of the scatter, galaxy colors can be compared versus their positions relative to the mean relationship in the metallicity - luminosity plane. Skillman et al. (1997) made such a comparison and two things became apparent. First, most of the sample galaxies have B-V colors between roughly 0.3 and 0.45 with small differences between the predicted and observed blue luminosities (which includes Sextans A). Secondly, Pegasus lies offset from this group with a redder color (0.54) and a lower than predicted luminosity. It would appear that the relatively low recent star formation rate has resulted in relatively red colors and low surface brightness.
The SFH can also have an effect on the relative abundances. Figure 10 shows the N/O abundances versus O/H for the sample of dIs and H II galaxies compiled by Kobulnicky & Skillman (1996). Only the points for those galaxies which do not show evidence of WR features in their spectra are plotted. Pegasus stands at the upper envelope of the values of N/O. Recent Calar Alto observations of Sextans A have allowed us to determine an N/O (Skillman et al. 1998), and note that Sextans A appears near the lower envelope of values.
Kobulnicky & Skillman (1996, 1997) have argued that ISM abundances mainly reflect the past chemical enrichment history and not current ``pollution''. In this interpretation, it may be possible to see depressed values of N/O in galaxies which have had recent episodes of star formation (O production happens quickly in massive stars and is over within a few tens of millions of years while N enrichment is delayed because it forms in more slowly evolving intermediate mass stars with ages in excess of 100 Myr). On the other hand, in galaxies that have been relatively quiescent for a long period of time one would expect relatively high values of N/O (Edmunds & Pagel 1978).
The positions of Pegasus and Sextans A in Fig. 10, and their recent SFHs as recorded in their CMDs support the simple picture that the N/O ratio is an indicator of recent SFR. Maintaining some skepticism, we could point out that there is a smaller range in N/O at the very low metallicity end of the scale, and, thus, maybe the position of Sextans A is not significant. However, there are very few abundance studies on quiescent or reddish dIs. If we can study some of these systems at very low metallicity, we may turn up some higher values of N/O (but note that these are not predicted in the simple scheme of nitrogen production being chiefly primary at low metallicity). van Zee et al. (1997) have shown that a number of low surface brightness dwarfs do not show elevated N/O.
At this time, it is not possible to compile an extensive list of objects to test this hypothesis, but there are some other important test cases worth mentioning. NGC 6822 has one of the lowest values of N/O (log (N/O) = -1.66; Pagel et al. 1980). Hodge (1980) found evidence for stellar cluster formation over the last 100 Myr, with a strongly enhanced period of stellar cluster formation in the interval 75 to 100 Myr ago. Marconi et al. (1995) found evidence for episodic star formation, but the resolution of their modeling does not allow for a precise star formation history over the last few 100 Myr. Gallart et al. (1996c) found an enhancement in the star formation rate over the last 100-200 Myr. If the stellar cluster formation history reflects the overall SFH, then it appears that NGC 6822 has experienced a productive period of star formation over the last 50-200 Myr, which has not yet resulted in the elevation of the ISM N abundance. This could be consistent with N production delayed by a few 100 Myr.
Finally, we consider the case of NGC 2366. This galaxy currently shows a relatively high rate of massive star formation as evidenced by its Halpha/L(B) ratio (Hunter et al. 1993). Aparicio et al. (1995) present evidence for a dominant burst of star formation roughly 20-50 million years ago. On the other hand, based on deeper photometry, Tolstoy (1995) finds a relatively constant star formation rate over the last 300 million years. Although it is important to sort out this discrepancy, both studies support the view of vigorous star formation during the last few hundred million years in NGC 2366. This may be consistent with the relatively low N/O observed (log (N/O) = -1.61; González-Delgado et al. 1994).
In the future, with more HST stellar photometry it should be possible to determine much more accurate SFHs for the last 1 Gyr for a large sample of the nearby dIs. If the preliminary evidence presented here for delayed N production holds up, then it should be possible, in principle, to calibrate the length of that delay. This would be of great value in understanding the chemical evolution of galaxies. For example, Pettini et al. (1995) and Lipman (1995) find very low values for log (N/O) of ∼-2 for several damped Lyman-α systems of low metallicity (see also Lu et al. 1998). They favor an interpretation of delayed N production which means that they are observing these systems within a few hundred million years of dominant star formation episodes. Since the nearby dIs have similar metallicities to these damped Lyman-α systems, it is reasonable to expect that the stellar populations are similar. Thus, more studies of the nearby dIs can help us to understand galaxy formation at redshifts of 2 to 3 where damped Lyman-α systems can be observed.
From the observations of Kobulnicky & Skillman (1996, 1997) it appears that a good lower limit on the ability of oxygen to mix back into the cold/warm ISM is about 10 Myr. The low N/O in NGC 6822 combined with the elevated star formation rate over the last 100-200 Myr (Gallart et al. 1996c) implies that a good upper limit on the ability of oxygen to mix back into the cold/warm ISM may be about 100 to 200 Myr. From the high N/O observations of Pegasus (Skillman et al. 1997) and the drop in the SFR about 400 Myr ago (Gallagher et al. 1998), a good upper limit on the ability of nitrogen to mix back into the cold/warm ISM may be about 400 Myr.