AGES OF STAR CLUSTERS IN THE TIDAL TAILS OF MERGING GALAXIES

Observations of NGC 520 and NGC 2623 were taken with the WFC on the ACS on HST as part of program GO-9735 (PI: Whitmore). NGC 520 was observed in 2004 October using the filters F435W (≈B in the Johnson–Cousins system; exposed for 2400 s), F555W (≈V; 1000 s), and F814W (≈I; 1200 s). The entire galaxy, including the majority of its tidal tails, was covered in a single WFC frame. Observations of NGC 2623 were taken in 2004 February using the same filters with exposure times of 3730, 1206, and 2460 s, respectively. It was also covered in a single WFC frame. Images of NGC 3256 come from two separate observing proposals. F555W filter observations of NGC 3256 were obtained in 2003 November as part of the program GO-9735 (2552 s). WFC observations using F435W and F814W filters were taken in 2005 November as part of program GO-10592 (PI: Evans) for 1320 and 760 s, respectively. Both fields cover most of NGC 3256's eastern tail, as well as its main body. Figure 1 shows the WFC fields of view used in this study.

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The raw data were processed through the standard ACS pipeline. The reduced, multidrizzled images were taken from the Hubble Legacy Archive and have a scale of 005 per pixel. Figure 2 shows F814W images of our targets, with the tidal tail portions studied in this work labeled. We will refer to specific tails within a galaxy by abbreviating the location of each tail (i.e., the northern tail of NGC 520 will be called NGC 520N). We define the tail edges as the locations where the counts drop to 3σ above the general background level.

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Figure 2 shows that NGC 520 has two tidal tails, commonly referred to as the north and south tails. Both of these tails extend beyond the field of view in Figure 2. Not seen in Figure 2 is the dwarf galaxy UGC 957, at about 6' north–west of NGC 520. NGC 2623 is in the middle panel of Figure 2. It exhibits two nearly symmetric tidal tails, although the northern tail is wider than the southern tail. An extremely bright, triangle-shaped region south of the nucleus of the galaxy, which we call the pie wedge, is also highlighted. NGC 3256 is on the right in Figure 2. While it has two tails (east and west), only part of the eastern tail was observed. The three mergers are next to each other on the Toomre Sequence, in the order NGC 520, NGC 2623, and NGC 3256 (Toomre 1977). General properties of these systems are listed in Table 1.

Star cluster candidates were found by running the IRAF⁴ task DAOFIND on the I band images. A few thousand objects were detected in each galaxy. We perform BVI aperture photometry on all sources in NGC 520 and NGC 3256 using an aperture size of 3 pixels in radius and a background area of radii 8–11 pixels using the IRAF task PHOT. Because of the greater distance to NGC 2623, we use a smaller aperture of 2 pixels (with the same background area) in order to minimize contamination from nearby sources. In the pie wedge region, we estimate the background level using annuli between 5–8 pixels, in order to minimize the impact of crowding. Magnitudes were converted to the VEGAMAG system using zeropoints calculated from the ACS zeropoint calculator.⁵ Aperture corrections were performed in NGC 3256 by selecting the brightest ∼20 isolated clusters in each filter and finding their mean 3–10 pixel magnitude difference. We then applied the 10 pixel to infinity aperture corrections taken from the encircled energy catalog for ACS/WFC in Table 3 of Sirianni et al. (2005). The final corrections were 0.45, 0.42, and 0.54 magnitudes for B, V, and I, respectively, and they were applied to each cluster. Unfortunately, the background level was too high in NGC 520 and NGC 2623 for reliable aperture correction determinations. We find that clusters in NGC 520 have a similar range of sizes as those in NGC 3256 and therefore applied the aperture corrections derived from NGC 3256 to the NGC 520 clusters. Clusters in the more distant NGC 2623 are nearly indistinguishable from point sources; we assume the 2 pixel to infinity aperture corrections from Table 3 of Sirianni et al. (2005).

We measure FWHM values for all detected sources using the ISHAPE software (Larsen 1999). ISHAPE measures FWHMs by convolving the point-spread function with a King profile, then performing a χ² calculation to test the goodness of fit to each individual cluster (King 1966). ISHAPE iterates through different values for the effective radius until a minimum χ² is found.

We select cluster candidates using a combination of automated selection criteria (listed in Table 2) and visual inspection of each candidate (see Bastian et al. 2012). The selection criteria in m_V, FWHM, and magnitude error (σ_V) remove most foreground stars and background galaxies. We made a cut at the bright end of m_V = 20, in order to eliminate saturated foreground stars (brighter than any cluster expected in these mergers). The faint end of our cluster catalog likely suffers from incompleteness. We restrict our sample to luminosities brighter than the point where the luminosity function (discussed in Section 3.2) turns down rather than continuing to increase in a power-law fashion. We use cuts in FWHM (size) to separate stars and background galaxies from clusters, to the highest degree possible. In NGC 3256, we exclude sources with FWHM values less than 0.2 pixels, because the FWHM of the vast majority of likely field stars was measured to be smaller than that value. We also place an upper limit of 2.0 pixels in FWHM to exclude very extended background galaxies. The lower limit on FWHM for clusters in NGC 520 was chosen to be 0.1 pixels because there is a group of young, compact clusters in the northern portion of NGC 520 S. At the distance of NGC 2623, clusters are essentially point sources, so FWHM measurements cannot distinguish clusters from bright, individual stars. Therefore, in addition to an upper limit on FWHM of 2.0, we apply a cutoff in magnitude uncertainty, σ_V. In the tails, σ_V < 0.1 did a good job removing obvious non-clusters, but in the pie wedge, we relax our σ_V cut to 0.2 because the area is more crowded and has a very bright background.

Finally, we visually inspect each cluster candidate. For each galaxy we identify high signal-to-noise ratio (S/N) clusters, as well as obvious foreground stars, as benchmarks for fainter objects. We choose clusters based on their fuzzy appearance, as well as wider radial profiles than those of stars. Figure 3 shows three of the brightest clusters from each galaxy's tidal tails, and Figures 4–7 show locations of clusters in each tail. All possible clusters were identified independently of the selection cuts listed above, then later compared to the sample obtained from the selection cuts. We find excellent agreement between the two methods in NGC 520 and NGC 3256, and we only include sources that both fit our selection cuts and visually appear as clusters. Due to the poor S/N in NGC 2623, several sources appeared visually as likely clusters but had FWHM and/or σ_V values outside of the selection criteria. Because of this low S/N, we rely more heavily on visual inspection than selection cuts and include some sources with FWHM or σ_V values outside of the cuts. The total number of clusters found in each tidal tail (N_clus) is compiled in Table 3, and we compile the cluster catalog in Table 4.

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Table 3.  Properties of Star Clusters

Galaxy	Tail	τtail	Nclus	Σclus	τclus	$\langle B-V\rangle$	$\langle V-I\rangle$	α
NGC 520	north	300	5	0.00	∼100	0.08 ± 0.05	0.38 ± 0.05	...
NGC 520	south	300	93	0.49	260	0.19 ± 0.09	0.60 ± 0.17	−2.32 ± 0.36
NGC 2623	north	220	75	0.54	230	0.10 ± 0.15	0.48 ± 0.28	−2.39 ± 0.34
NGC 2623	south	220	18	0.32	...	0.12 ± 0.13	0.43 ± 0.17	...
NGC 2623	pie wedge	≤220	76	8.5	∼100	0.05 ± 0.15	0.43 ± 0.19	−2.02 ± 0.21
NGC 3256	east	450	141	0.10	≥260	0.21 ± 0.12	0.58 ± 0.21	−2.61 ± 0.27

Note. τ_tail (Myr) is the predicted age of the tidal tail as estimated from simulations. N_clus is the total number of clusters found in the region. Σ_clus is the star cluster density per kpc² where M_V < −8.5. τ_clus (Myr) is the approximate age of the bulk cluster formation. Median colors ($\langle B-V\rangle$, $\langle V-I\rangle$) are for clusters younger than the age of the tidal tail. α values (from the relation $dN/dL\propto {{L}^{\alpha }}$) listed come from the ${{\alpha }_{{\rm const}-{\rm num}}}$ method.

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Distant, red spheroidal galaxies can have a similar color and appearance to compact star clusters. To assess the possible contamination from background galaxies, we examine the colors and locations of sources that made it through our automated cuts, but are located outside of the tidal tail regions, with the reasonable expectation that most of these are likely background galaxies. However, we find no such objects in any of our fields, only extended, obvious background galaxies. In addition, many background galaxies should appear elliptical rather than circular. Using ISHAPE, we measure the ellipticity of each cluster candidate and find few elongated red sources. We conclude that the fraction of background galaxies that have made it through our cluster selection criteria is negligible.

One quantitative measure of a cluster population is its surface number density, which is often used when cluster selection must be done statistically. We measure the number density of clusters, Σ_clus (clusters kpc⁻²), for each tidal tail based on the catalogs in this work. We choose a magnitude cutoff of M_V < −8.5 (see Section 3.2 for details on absolute magnitude calculation) to directly compare our results to those of M11 (although altering the cutoff to M_V = −8.0 and −9.0 yields qualitatively similar results). NGC 520 S yielded Σ_clus = 0.49, while NGC 520 N yielded Σ_clus = 0.0 because its five clusters were all fainter than M_V = −8.5. The pie wedge in NGC 2623 possesses the highest cluster density in our sample by over an order of magnitude at Σ_clus = 8.5, while the north and south tails yielded Σ_clus = 0.54 and 0.32, respectively. Lastly, NGC 3256 E possessed a relatively low density of Σ_clus = 0.10. We find that while all tidal tails in our sample contain clusters, there is a large variation in Σ_clus for each tail.

For both NGC 520 and NGC 3256, we are able to directly compare our values of Σ_clus to those values found in M11, although we note that we cover different areas of the tails for both systems. We also note that we adopt slightly different distance moduli than M11, resulting in small differences in Σ_clus. M11 found Σ_clus = −0.011 ± 0.013 in NGC 520 N and Σ_clus = 0.004 ± 0.082 for NGC 3256 E, both of which are consistent with zero. While our results for NGC 520 N agree with M11, we find a clear population of clusters in NGC 3256 E brighter than M_V < −8.5. M11 also measured Σ_clus for M_V < −7.5 and M_V < −6.5, however we do not compare Σ_clus values at these cutoffs due to incompleteness. We find that, in general, younger tails tend to exhibit higher Σ_clus, in agreement with the results of M11.

The LFs of the tail clusters are shown in Figure 8. These can be described by a power law, dN/dL ∝ L^α, where α is determined from a linear fit to log$(dN/d({{M}_{V}}))$. We fit our cluster LFs over the V band magnitude range which we believe to be reasonably complete. Magnitudes include distance moduli from Table 1 as well as aperture corrections, and they have been dereddened because of foreground extinction (see Section 3.3.1). We are unable to measure the LFs in NGC 520 N and NGC 2623 S because they contained too few clusters.

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We determine the LFs in two different ways, using both a fixed magnitude bin width (denoted α_const–mag) and a fixed number of clusters per bin (${{\alpha }_{{\rm const}-{\rm num}}}$). The results from both methods are consistent, and we report α values for the latter because it is better at handling low number statistics. In all cases, however, we find that the two values for α are the same, within the uncertainties. The typical standard deviation among α calculated using B, I, and V band magnitudes is ≈0.13, with no systematic trend. The LFs in the B and I bands have similar slopes to those in the V band.

In NGC 520 S, we find a best fit of α ∼ −2.32 ± 0.36 down to M_V < −8.0. We restrict our fit to this magnitude limit because completeness becomes an issue fainter than this. For NGC 2623, we focus on the pie wedge and northern tail. We fit the luminosity function in both regions down to M_V ≈ −9.0, and find α = −2.39 ± 0.34 for the northern tail and α = −2.02 ± 0.21 for the pie wedge. The luminosity function in NGC 3256 E is best fit with α = −2.61 ± 0.27 down to M_V = −8.0. We find no evidence for a deviation from a simple power-law at the bright end of the LF in any of our tail regions.

Several studies have shown that the luminosity function can range from −2.8 < α < −1.9 for clusters in a variety of environments (e.g., Whitmore et al. 2014). We find that our results for α are within this range.

We present the color–color diagrams for clusters in all tails in Figure 9. The color coding corresponds to different tails within each merging system. We also show with blue symbols the median V − I and B − V colors of clusters that apparently formed after the tidal tails (we assume tail ages from simulations that are discussed in Section 4.1 and compiled in Table 3). The solid line in each panel in Figure 9 is a stellar population model from G. Bruzual & S. Charlot (2006, private communication, hereafter BC06; and see also Bruzual & Charlot (2003) which predicts the evolution of star clusters from about 10⁶–10¹⁰ yr. All models assume a metallicity of Z = 0.02, and a Salpeter (1955) initial mass function. Numbers mark the logarithmic age (τ) corresponding to the population.

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3.3.1. The Impact of Internal Reddening on Cluster Colors

In order to accurately age-date star clusters, we must understand how reddening affects their measured colors. We correct each cluster for foreground reddening of the Milky Way by using A_λ values given in the NASA/IPAC Extragalactic Database, which are compiled in Table 1. The reddenings were estimated with the standard relations between them and the extinction. A reddening vector for A_V = 1 is shown in Figure 9 by the red arrow. The reddening vector is nearly parallel to the predicted evolution of cluster colors, making it challenging to disentangle the effects of age and reddening in the measured colors. For example, a young cluster with some extinction will have the same colors as an older cluster with no reddening.

Figure 9 shows that the clusters in each tail have a measured color spread. In principle, there are a few reasons this spread could exist other than photometric uncertainties, such as internal (and therefore differential) extinction or a range in cluster ages. Below, we examine each galaxy individually and determine that internal extinction does not greatly affect the majority of our cluster colors.

We first focus on NGC 2623, whose color distributions are shown in the upper- and lower-right panels of Figure 9. The spread in measured colors is the smallest for any tail studied here, and instead of spreading along the reddening vector, the cluster colors are fairly isotropically distributed. This pattern suggests that photometric uncertainties are primarily responsible for the color spread. We confirm this by creating color–color diagrams of NGC 2623 clusters in different M_V bins, and while the cluster color spread increases with fainter magnitudes, the color distribution is centered near the median colors quoted in Table 3, regardless of brightness. If all of the clusters in the tails of NGC 2623 formed at the same time, clusters obscured by dust would exhibit redder colors than unobscured clusters, causing a color spread along the reddening vector. The clusters in both tidal tails of NGC 2623 appear to have formed over a relatively short period of time. This is also true in the pie wedge, where the narrow color ranges are also indicative of coeval cluster formation.

NGC 520 and NGC 3256, on the other hand, both exhibit a spread in color along the reddening vector. Here, we use the relative locations of blue (B − V < 0.3) versus red (B − V ≥ 0.3) clusters, shown in Figures 10 and 11, to assess whether age or reddening are more likely to be responsible for the observed spread in color. In the case of NGC 520 (Figure 10), the bluer clusters tend to trace out the central, densest portion of the tail, while the redder clusters are more evenly spread throughout the tail. Reddening due to dust would be expected to give the opposite trend, with reddened clusters mostly near the center of the tail. Figure 11 shows the locations of blue versus red clusters for NGC 3256. No obvious pattern is observed, but there is no evidence that red sources are primarily found in the brightest, densest portion of the tail. These results support the interpretation that the larger color spread of the clusters in the tails of NGC 520 and NGC 3256 is primarily due to a spread of cluster ages rather than to varying amounts of cluster reddening.

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We now compare the colors of red tail clusters to clusters in the halos of NGC 520 and NGC 3256, which are shown in Figure 12. Selection and photometry of the halo clusters are performed using the same methods as in Section 2. We conclude that they are fairly similar in color, and therefore the red tail clusters are likely to be clusters that formed before the tails but were swept out during the tail formation. We will study these halo clusters and other main body clusters in more detail in a forthcoming paper. All of this evidence supports the interpretation that the red sources are older, unobscured star clusters.

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3.3.2. Ages

In addition to distinct stellar clusters, we also study the diffuse stellar light in the tails. Because the tidal tails form from material stripped from the native disks, this diffuse light should consist of a mixture of old disk stars, faint newly formed stars, and dissolved clusters. While individual stars are unresolved, the integrated colors of the diffuse light can help to constrain the age of the stellar populations (e.g., Gallagher et al. 2010). The light will be dominated by the brightest and youngest unresolved objects, so we expect that our age estimates will be lower limits.

We create color images: B − V, V − I, and B − I for all three galaxies, and we show the B − I images in Figure 13. The images were boxcar-smoothed with a 4 pixel kernel width. The main bodies of the mergers are very red because they contain large amounts of dust. The tidal tails are easily seen in Figure 13 as bluer regions where internal extinction is low (except, perhaps, in NGC 520 N, as also suggested by our analysis of clusters in Section 3.3.1). We therefore assume that colors in the tail regions are not greatly affected by internal extinction, and we believe diffuse light colors trace the ages of stellar populations relatively well.

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We measure the V − I and B − V colors in different tail regions (white boxes in Figure 13 labeled from "A" up to "E" for each tail) after subtracting the point sources. Colors are calculated from the ratio of fluxes on a pixel by pixel basis and corrected for differences in instrumental zeropoints and foreground extinction. The diffuse light regions were chosen to cover most of the optically bright portions of the tails while avoiding breaks in field coverage and saturated foreground stars. The values in Table 5 are the mean color for each region. Figure 14 plots these values on a color–color diagram overlaid with the BC06 track. Figure 13 reveals that the diffuse stellar light in all tidal tails in our sample exhibits a common pattern (with the exception of NGC 3256 E). While the tidal tails are fairly blue (B − I ≤ 1.4), we find that there is a gradient with B − I color; it is redder near the edges of the tail and bluer near its axis. As previously discussed, this is not likely a reddening effect, so we believe that we are observing a gradient in the mean ages of the stars across these tails. In Section 3.3.1 we discussed the tendency for younger clusters to be near the center of the tail axis, leaving mostly older clusters along the tail edges, especially in NGC 520 S. The diffuse light images further support a lack of young stars near the edges of the tails. We believe this to be the first time such a gradient across the tails has been observed in these galaxies.

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Table 5.  Diffuse Light Measurements

Galaxy	Region	$\langle V-I\rangle$diffuse	$\langle B-V\rangle$diffuse
NGC 520	NA	1.13	0.70
	NB	1.05	0.60
	SA	1.14	0.71
	SB	0.94	0.49
	SC	0.90	0.47
	SD	0.88	0.46
	SE	0.97	0.49
NGC 2623	NA	0.96	0.48
	NB	1.07	0.42
	NC	0.76	0.33
	ND	0.90	0.47
	SA	0.88	0.45
	SB	0.87	0.40
	SC	0.92	0.38
	SD	0.91	0.41
	SE	0.96	0.42
	P	0.68	0.31
NGC 3256	EA	0.77	0.82
	EB	0.71	0.79
	EC	0.78	0.39
	ED	0.83	0.76
	EE	0.66	0.63

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Figure 14 reveals that the northern tail of NGC 520 has the reddest diffuse light of any tail in our sample. This is likely due to dust, as the red B − I color in Figure 13 extends from the main body straight through the tail. Region A in the southern tail is also almost certainly affected by dust. We believe that the colors of the diffuse light in the other regions of the southern tail are dominated by ages of the underlying populations and not extinction because of the absence of dust lanes in this tail. These regions have an age τ of around 800 Myr.

In NGC 2623, most of the regions in the north tail have colors consistent with stellar populations of ∼700 Myr old. Region NC, however, is bluer and yields an age of ∼500 Myr. We note that the diffuse light in the pie wedge (region P) has B − V and V − I colors close to region NC, as seen in Table 5, which might provide clues to its formation, as we discuss further in Section 4.1.2.

Figure 14 shows that the diffuse light in NGC 3256 is perhaps the most peculiar of our sample. Only region EC falls near the BC06 track, with an associated age τ ≈ 650 Myr. This would indicate that the diffuse light population in region EC is similar to the majority of the diffuse populations in the tails of NGC 2623. It would be plausible that region EC falls on a local region of enhanced star formation, causing its diffuse light colors to be bluer than other regions of NGC 3256 E. However, region EC does not include any young star clusters, instead containing only one cluster whose colors suggest that it predates the merger. The rest of the regions we measured are gathered around V − I ≈ 0.7 and B − V ≈ 0.8 and are off the BC06 track. Regardless, these offset colors suggest faint population ages of about 800–1000 Myr. It is likely that most of these diffuse light regions do not have colors consistent with the BC06 track because of the low surface brightness of NGC 3256 E, resulting in large photometric errors. It is unlikely that Hα line emission could be responsible for the offset from the model, although this emission would move measurements in the observed direction relative to the models, because such recent star formation is usually quite bright and lumpy.

The age of the tidal tails, which we define as τ_tail, refers to the time since perigalacticon of the two progenitor galaxies, because they form from material ejected after the galaxies undergo their closest approach. The ages of tails can be estimated from simulations and are probably accurate to within ∼20% (∼30–100 Myr; Barnes & Hibbard 2009). In this section, we briefly summarize previous simulations of our sample of mergers, and we compare tail ages from those simulations with our measured cluster ages.

We find that, for most systems, there is a period of increased cluster formation near the time that the tidal tails formed. We quantify the time of this cluster formation (hereafter denoted τ_clus) by comparing the density peak of the color distribution with cluster evolution model predictions. These density plots are shown in the right panels of Figure 15. Contours are plotted, and the density peak is marked with a circle. The colors of the peak are then compared with a BC06 track to obtain an age (shown in left panels of Figure 15).

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4.1.1. NGC 520

4.1.2. NGC 2623

4.1.3. NGC 3256

4.1.4. The Timescale of Cluster Formation

We suggest the following sequence of merger ages: NGC 2623 is the most recent of the three mergers, followed by NGC 520, with NGC 3256 being the oldest. This sequence is supported by both our estimated cluster ages and simulated ages of tidal tails (and thus ages of the mergers themselves). We note that the age sequence suggested by these simulations is different than what the Toomre Sequence suggests, as NGC 520 and NGC 2623 are flipped (Toomre 1977).

We find that clusters in the pie wedge have ages around τ_clus ∼ 100 Myr. The north tail of NGC 2623 contains clusters with ages at τ_clus ∼ 230 Myr, while the few south tail clusters are slightly younger, although we do not attempt to quantify their ages. NGC 520, the next youngest merger, exhibits a slight age difference among clusters in different tails. NGC 520 N has with clusters around τ_clus ∼ 100 Myr, while NGC 520 S has clusters with τ_clus ∼ 260 Myr. The oldest merger, NGC 3256, contains clusters in the eastern tail with ages τ_clus ∼ 260 Myr, although we caution that the broad age distribution in NGC 3256 E suggests that cluster formation likely began earlier than 260 Myr ago. For a direct comparison, all the values of τ_tail and τ_clus are listed in Table 3.

We have measured the LFs for all tails where possible, and find tentative evidence that younger mergers exhibit shallower LFs, although our sample size is small. M11 found that a LF with −2.5 < α < −2 provided qualitatively acceptable fits to their statistically determined cluster sample in tails where the measurement was possible, although they were not able to directly fit the LF with a power law due to the poor quality of the data. Whitmore et al. (2014) measured the luminosity function in a variety of spiral galaxies in search of various correlations between α and galactic environment. A weak negative correlation was found between star formation rate (SFR) and α (i.e., galaxies with high SFR tend to have flatter LFs). For instance, the Antennae, a major merger with a relatively high SFR, had the second lowest α value in their sample, at α = −2.07 ± 0.03.

No measurements of the cluster LF were made in tidal tails, however. While α tends to decrease with the tail age, all α values in our sample are within the range observed in a wide variety of other environments (Whitmore et al. 2014 and references therein).