Stellar Population and Structural Properties of Dwarf Galaxies and Young Stellar Systems in the M81 Group^*

Figure 4 shows the i₀ versus (g − i)₀ CMDs of stellar objects lying within the half-light radii (R_h) of the best-fit Sérsic profile of galaxies, as derived in Section 3.2. Theoretical PARSEC v1.2 isochrones of age 10 Gyr with [M/H] = −2.2, −1.75, −1.3, −0.75 for the SDSS filter system (Bressan et al. 2012) are overlaid, adjusted to the distance of each galaxy using the distance moduli in the literature (Karachentsev et al. 2002, 2006, 2013; Sharina et al. 2008; Chiboucas et al. 2013; Smercina et al. 2017), which is very close to the HSC filter system (Kawanomoto et al. 2018; Komiyama et al. 2018). We have chosen to adopt literature distances (see Table 3), instead of deriving our own, because these are generally based on high-precision photometry obtained with the HST. The error bars in each panel show the photometric error at (g − i)₀ = 1 as derived from the artificial star tests. The dotted polygons delineate the selection boxes of RGB stars for the estimation of structural parameters in Section 3.2 and metallicities in Section 3.3.

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Thanks to the photometric depth of our survey, we are able to resolve individual stars to at least 1.5 mag below the RGB tip (TRGB) of each dwarf galaxy. In Figure 4, the dominant populations are seen to be metal-poor RGBs. The four brighter dwarfs, IKN, KDG061, KDG064, and BK5N, also show populations of asymptotic giant branch (AGB) stars brighter than the TRGB, indicating these systems contain intermediate-age (0.5–5 Gyr) stars. The small handful of objects blueward of the RGB selection boxes are likely to be high-z galaxies, as discussed in the previous section.

We compute the structural parameters of the eight dwarf galaxies using the density-weighted first and second moments of the RGB star spatial distribution. The selection criteria for the RGB stars are shown as the dotted polygon in each panel of Figure 4. We define the RGB box to include stars in the color interval sandwiched by the [M/H] = −2.2 and [M/H] = −0.7 isochrones, convolved with the photometric errors. This interval is deemed to be sufficiently wide to capture the expected populations, while avoiding areas that are densely populated by contaminants.

We find values for the centroids, average ellipticities, and position angles that are consistent with those from previous studies except for the decl. of IKN, which we find to be 13 north of the value listed in the NASA/IPAC Extragalactic Database (NED).¹¹ Because IKN overlaps with two bright foreground stars located at about 05 and 27 north of the IKN center, the scattered light from these stars may have prevented accurate estimation of its center (see two holes in the lower-left panel of Figure 5). To reduce these effects, we populated the areas affected by these bright stars with stars drawn from a similarly sized axially symmetric position before estimating the structural parameters of IKN. The resulting coordinate is consistent with that in Tudorica et al. (2015), which derived the center from an HST/Advanced Camera for Surveys (ACS) stellar catalog covering the southern half of IKN.

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Table 2 lists the structural parameters of the eight dwarf galaxies, which are derived from fitting the standard King, exponential, and Sérsic profiles via least-squares minimization to the completeness-corrected RGB star counts. The completeness of a given RGB star is inferred from comparing its magnitude and color with a modified sigmoid function fit to the point-source detection ratio as functions of the magnitude and color, as given by the artificial star tests. The radial star count profiles are then constructed by calculating the average number density of RGB stars in a series of elliptical annuli, defined using the structural parameters derived previously. The field contamination level at the location of each galaxy is calculated first as an average of number density in a certain radius range of outer region by eye. Then we fit the profiles and calculate the contamination level again as a trimmed mean of number density between radii of R_t < r < 2 ∗ R_t for relatively large IKN, KDG061, KDG064, and R_t < r < 4 ∗ R_t for small galaxies. Finally the profiles are fitted using these values. Figures 5–7 show the radial profiles and spatial distributions of the RGB stars in the eight galaxies. Remarkably, this is the first estimation of structural parameters for IKN, even although it is one of the most well-studied satellites in the M81 Group. Again, this is probably due to the presence of the bright foreground stars, which hinder any simple analysis. The exponential scalelengths R_exp of KDG061 and KDG064, and the Sérsic radii R_h of d0955+70, d1014+68, and d1015+69, are all larger than the previous estimations (R_exp = 048, 028, and R_h = 031, 033, 014, respectively) based on integrated light surface photometry (Sharina et al. 2008; Chiboucas et al. 2013). The Sérsic radius R_h of BK5N is consistent with that (R_h = 04) of Caldwell et al. (1998), while the Sérsic index is much smaller than their value (n = 1.81), also derived from integrated light photometry. While these differences are likely due to the greater stellar extents of these systems as mapped in our survey, they could also reflect issues with our completeness-correction in the central regions of the brightest systems. In the case of d1005+68, the R_h we find is much larger than that (R_h = 016) estimated by Smercina et al. (2017), also using resolved stars. This is likely due to their assumption of a circular distribution; if we also assume the ellipticity of d1005+68 is  = 0, we obtain a similar value of R_h = 019.

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Table 2.  The Structural Parameters of the Early-type Dwarf Galaxies

Galaxy	α	δ	P.A.a	b	Rcc	Rtd	Rexpe	nf	Rhg
	(J2000)	(J2000)	(deg)		(arcmin)	(arcmin)	(arcmin)		(arcmin)
IKN	10h08m050	+68°25'161	174.6	0.10	0.4 ± 0.1	8.5 ± 0.4	0.85 ± 0.01	1.24 ± 0.03	1.32 ± 0.01
KDG061	09h57m032	+68°35'344	53.1	0.28	1.1 ± 0.2	5.3 ± 0.4	0.74 ± 0.04	0.40 ± 0.01	1.36 ± 0.01
KDG064	10h07m020	+67°49'434	22.3	0.35	1.0 ± 0.2	4.1 ± 0.3	0.59 ± 0.01	0.70 ± 0.07	1.11 ± 0.03
BK5N	10h04m416	+68°15'248	125.4	0.26	0.15 ± 0.04	1.62 ± 0.1	0.20 ± 0.01	0.77 ± 0.08	0.38 ± 0.02
d0955+70	09h55m143	+70°24'242	178.9	0.07	0.58 ± 0.07	1.5 ± 0.1	0.31 ± 0.04	0.51 ± 0.07	0.44 ± 0.02
d1005+68	10h05m323	+68°14'173	38.4	0.27	0.3 ± 0.2	1.3 ± 0.9	0.22 ± 0.04	0.5 ± 0.2	0.30 ± 0.05
d1014+68	10h14m547	+68°45'359	115.2	0.20	0.60 ± 0.05	3.1 ± 0.5	0.48 ± 0.02	0.9 ± 0.1	0.74 ± 0.04
d1015+69	10h15m068	+69°02'116	103.5	0.24	0.09 ± 0.03	2.5 ± 2.0	0.11 ± 0.01	2.5 ± 1.5	0.22 ± 0.04

Notes.

^aPosition angle from north to east. ^bEllipticity  = 1 − b/a where b/a is the axis ratio of a galaxy. ^cCore radius of the King profile. ^dTidal radius of the King profile. ^eScalelength of the exponential profile. ^fSérsic index. ^gHalf-light radius of the Sérsic profile.

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The radial profiles shown in Figures 5–7 are well described by the Sérsic and King profiles. It is curious that the number density of RGB stars in KDG061 gradually increases beyond the tidal radius. Because KDG061 is the nearest dSph to M81 having a projected distance D_M81,pro = 31 kpc, this is likely due to contamination by M81 (and perhaps also NGC 3077) halo stars. BK5N also has a slight excess of stars at around 5 kpc, which is likely due to stars belonging to another dSph, d1005+68, located at 5 kpc from the BK5N center. The faintest two galaxies, d1005+68 and d1015+69, show small overdensities at a radial distance of roughly 2 kpc; however there is no visible structure in the RGB maps, and these enhancements may just reflect noise. Overall, no obvious tidal features, such as S-shaped structures, are found in these galaxies, although some deviations from the King profile can be seen in Figures 5–7. We return to this topic in Section 6.1.

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Using stars that lie within the best-fit Sérsic radius R_h, we estimate the total absolute magnitudes M_g and M_i of the dwarf galaxies. The total flux of RGB stars within R_h is doubled and then corrected to account for the photometric completeness, contaminants, and stars fainter than the RGB selection box. The V-band absolute magnitudes M_V,0, (V − I)₀ colors and errors are calculated from transforming M_g and M_i (Jordi et al. 2006) and are listed in Table 3. The errors include the photometric error, the distance error, and the uncertainties of the adopted corrections. While the M_V,0 of KDG061, KDG064, BK5N, and d1005+68 are all in good agreement with the previous estimations (Caldwell et al. 1998; Karachentsev et al. 2000; Smercina et al. 2017), we find M_V,0 = −14.29 for IKN, which is almost 3 mag brighter than the previous estimation of Georgiev et al. (2009) (M_V,0 = −11.51) derived from the total B-band magnitude and assuming an average (B − V) = 0.45 and (V − I) = 0.7. To estimate the luminosity of IKN, we again populated the areas affected by the bright stars, with stars drawn from similarly sized axially symmetric positions. Because our estimation is less affected by the PSF halos of these bright foreground stars than integrated light studies, we conclude our measurement is more reliable than the previous estimation. This elevates IKN to being the brightest dwarf satellite in the M81 Group and lying in the luminosity-size realm of systems typically classified as dwarf ellipticals.

Table 3.  The Properties of the Early-type Dwarf Galaxies

Galaxy	Type	Distance	Rc	Rt	ca	Rh	MV,0	(V − I)0	$\langle {[{\rm{M}}/{\rm{H}}]}_{{\rm{w}}}\rangle$ b	σ[M/H]c
		(Mpc)	(kpc)	(kpc)		(kpc)	(mag)	(mag)	(dex)	(dex)
IKN	dE	3.75 ± 0.45(1)	0.5 ± 0.2	9.3 ± 0.4	1.27	1.44 ± 0.02	−14.29 ± 0.49	1.37 ± 0.22	−1.17 ± 0.23	0.39
KDG061	dSph	3.60 ± 0.43(1)	1.2 ± 0.2	5.6 ± 0.4	0.67	1.42 ± 0.01	−13.38 ± 0.52	1.36 ± 0.22	−1.29 ± 0.21	0.38
KDG064	dSph	3.70 ± 0.44(1)	1.1 ± 0.2	4.4 ± 0.3	0.60	1.19 ± 0.04	−13.31 ± 0.55	1.34 ± 0.23	−1.32 ± 0.30	0.37
BK5N	dSph	3.78 ± 0.45(1)	0.16 ± 0.04	1.8 ± 0.03	1.05	0.41 ± 0.02	−11.23 ± 0.58	1.26 ± 0.23	−1.60 ± 0.34	0.32
d0955+70	dSph	$3.45{\,}_{-\,0.46}^{+0.60}$(2)	0.6 ± 0.1	1.5 ± 0.1	0.40	0.44 ± 0.02	−10.36 ± 0.47	1.32 ± 0.35	−1.55 ± 0.23	0.27
d1005+68	dSph	$3.98{\,}_{-\,0.43}^{+\,0.38}$(3)	0.4 ± 0.2	1.5 ± 1.1	0.57	0.34 ± 0.05	−8.14 ± 0.55	1.25 ± 0.26	−1.79 ± 0.37	0.28
d1014+68	dSph	3.84 ± 0.33(2)	0.67 ± 0.06	3.4 ± 0.6	0.70	0.82 ± 0.05	−11.28 ± 0.48	1.28 ± 0.24	−1.50 ± 0.27	0.33
d1015+69	dSph	$3.87{\,}_{-\,0.21}^{+\,0.26}$(2)	0.10 ± 0.03	2.8 ± 2.2	1.45	0.25 ± 0.05	−9.2 ± 0.41	1.26 ± 0.20	−1.61 ± 0.28	0.31

Notes.

^aConcentration parameter c = log(R_t/R_c). ^bThe weighted mean metallicity. ^cThe metallicity dispersion.

References. (1) Karachentsev et al. (2013), (2) Chiboucas et al. (2013), (3) Smercina et al. (2017).

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Figure 8 shows the metallicity distribution functions (MDFs) of the dwarf galaxies as derived from the completeness-corrected color distributions of bright RGB stars (i₀ < 25). These stars are selected according to the criteria in Figure 4 and lying within 2 × R_h of the galaxy center. The photometric metallicity of each RGB star is derived from the linear interpolation of isochrones of [M/H] = −2.2 to [M/H] = −0.7 with a fixed age of 10 Gyr. Note that we choose a 10 Gyr population so as to cover stellar populations in dwarf galaxies having wide luminosity range from M_V,0 = −14 to −8. The uncertainty introduced by assuming a constant old age on the MDFs is quite small, as shown in Lianou et al. (2010). The difference in the resulting metallicity among the cases of fixed ages of 12.5, 10.5, and 8.5 Gyr in their study is less than 0.2 dex, and the shape of MDFs does not change significantly.

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We use RGB stars brighter than i₀ = 25 to reduce the effects of spatial variance on the photometric error and completeness, and to keep wide spacing between isochrones in the g − i color. The error bar in each panel indicates the mean error estimated by performing a Monte Carlo simulation with N = 2000 for each RGB star within its photometric uncertainty. Each star is randomly resampled from a Gaussian distribution with width equal to the photometric error, and the metallicity is redetermined. The standard deviation of the resultant distribution is adopted as the metallicity uncertainty for each RGB star. Due to the narrow range of RGB color between metal-poor isochrones, the estimated uncertainties increase with the decreasing metallicity. The mean metallicity corrected for this effect is shown as the error-weighted mean value in each panel of Figure 8 and listed in Table 3. The dashed line represents single Gaussian fits to the MDF. Foreground and background contamination is estimated using point sources drawn from the area outside of R_t of each galaxy, and excluding regions populated by other galaxies or overdensities. As can be seen in the gray shaded histograms of Figure 8, the effect of contamination on the resultant MDFs is generally negligible.

The estimated metallicity of IKN is lower than the previous estimation ($\langle [\mathrm{Fe}/{\rm{H}}]{\rangle }_{{\rm{w}}}=-1.08$) using the HST/ACS data (Lianou et al. 2010). Although their HST photometry reached a deeper magnitude than ours, they were limited to a small FOV and hence likely missed more metal-poor stars in the outer regions of the galaxy. On the other hand, our HSC photometry will definitely miss some (presumably more metal-rich) stars within the central half arcminute area due to the crowding. The weighted mean metallicities of KDG061 and KDG064 are higher than those reported in Lianou et al. (2010) ($\langle [\mathrm{Fe}/{\rm{H}}]{\rangle }_{{\rm{w}}}=-1.49$ and −1.57, respectively), but the differences are within the uncertainties. The non-weighted mean metallicities of BK5N and d1005+68 are consistent with those values $[\mathrm{Fe}/{\rm{H}}]$ = −1.7 ± 0.4 of Karachentsev et al. (2000) and $[\mathrm{Fe}/{\rm{H}}]$ = −1.9 ± 0.24 of Smercina et al. (2017), respectively.

The MDFs in Figure 8 are generally well described by a single Gaussian. There is a slight excess of metal-poor stars in d0955+70, d1014+68, and d1015+69, but this might simply reflect the fact that our metallicity estimation is limited to [M/H] > −2.5 due to the available isochrones. The overall symmetric appearance of the MDFs implies that there was no sudden truncation of star formation in the past, which would be expected to cause sharp drops on the metal-rich side (Yoshii & Arimoto 1987). Although the resulting MDFs do not contain particularly apparent metal-poor tails, the overall shapes resemble the combined MDF of similarly luminous Galactic dSphs shown in Kirby et al. (2013), which are well reproduced by an infall model.

Figure 16 shows the CMDs of young stellar systems listed previously. In the CMDs for AL and Holm IX, shown in the upper panels of Figures 16, the most prominent features are vertical distributions at (g − i)₀ ∼ −0.5 composed of MS and cHeB stars, and another vertical distribution at (g − i)₀ ∼ 1.8 composed of RSG stars. The color and magnitude distributions of RGBs in AL and Holm IX are similar to those in the reference CMDs. The red RGBs (92 stars) in AL CMD are comparable to those in its reference CMD (107 stars), while those in Holm IX are less than those in its reference due to crowding. This implies that the old populations present in these fields are simply due to the M81 halo (see also de Mello et al. 2008; Sabbi et al. 2008; Barker et al. 2009); this is discussed further in Section 6.3. Similar CMD features can be seen in the SE-blob, NW-arm, BK3N, SE-streams, and the Garland. On the SE-stream, the CMDs of the two regions are very similar to each other. In the stellar arcs near M82, the brightest young stars are fainter than those seen in other systems, suggesting older ages.

Overall, these young stellar systems show clumpy spatial distributions and appear to be embedded in much smoother distributions of old RGB stars (see Figure 17). AL appears to have a tidal structure extending to the northwest side, while Holm IX shows the relatively round shape. The center of the RSG distribution in Holm IX appears offset from that of the MS/cHeB stars, while other systems do not show such a difference. Because the CMD of Holm IX in Figure 16 shows more bright MS/cHeB stars in comparison with other systems, the spatial difference of RSG and MS/cHeB in Holm IX implies the different recent star formation history (SFH) in this area from others. The SE-blob is elongated toward the southeast, while the NW-arm, which lies at an elliptical radius of 22 kpc from the M81 center, appears to truncate abruptly. BK3N, the smallest of the young stellar systems studied here, is at the same elliptical radius as the NW-arm and may be a further fragment of that feature. The SE-stream, lying between M81 and NGC 3077, shows regions of local enhancement. The Garland has a unique shape, being highly elongated toward the south and the east. Interestingly, there is no obvious counterpart to this in RGB distribution. The stellar arcs of M82 lie perpendicular to the disk. Although the width of the M82S-arc is much broader than that of the M82N-arc, their positions appear to be axially symmetric.

To further investigate the stellar populations in these systems, we also plot the g₀ versus (g − i)₀ CMDs in Figure 18. Compared with the i₀ versus (g − i)₀ CMDs, the young MS/cHeB stars in Figure 18 have the luminosity peak at around (g − i)₀ = −0.4; then, their evolutionary track moves redder and fainter. The g-band luminosity of the brightest star is highly dependent on age but is insensitive to metallicity. It varies by about 0.25 mag between the [M/H] = −1.2 and −0.2 isochrones at constant age, while it varies about 4.5 mag between the ages of 0 and 200 Myr isochrones at constant metallicity (Bressan et al. 2012). Therefore, we investigate the g-band LF of upper MS and cHeB stars in order to set a constraint on the duration of the star formation period.

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Figure 19 shows the LF of young stars that lie in the selection box shown in Figure 18. The vertical dashed lines show the g-band magnitudes of the brightest MS/cHeB stars along isochrones of age 20–150 Myr with [M/H] = −0.75, assuming populations at the distance of M81. While the Garland and the M82S/N-arcs are likely to be more closely associated with NGC 3077 and M82, the distance moduli to these systems differ by only −0.13 mag, and 0.05 mag from that of M81 (Dalcanton et al. 2009). Given that this is smaller than our bin size of 0.2 mag, this is unlikely to cause any significant effect. Photometric incompleteness becomes important at roughly g₀ > 25, which is manifest in the LF downturns seen in Holm IX, the Garland, and BK3N. We do not correct for the incompleteness in Figure 19 because we are solely interested in the brighter part of the LF.

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AL, the SE-blob, NW-arm, and SE-stream1 all have similar LFs that decrease gradually with increasing luminosity, then flatten out at g₀ ∼ 23, with few or no stars brighter than g₀ = 22. This suggests that the youngest populations in these systems are ∼30 Myr old and that they have all experienced similar recent SFHs. In the young metal-poor population, the upper MS stars (M_g < −2.3, corresponding to g₀ < 25 at the M81 distance) and the brightest MS/cHeB stars have similar mass: M = 8.1M_⊙ to 8.8M_⊙ in the 30 Myr old population and M = 5.5M_⊙ in the 75 Myr old population, which makes a flat LF for a single stellar population at this magnitude range. Therefore, the moderate slopes of these LFs indicate continuous star formation in these young objects in the period of 50–150 Myr ago. BK3N also has similar LF at g₀ < 25. Considering the locations of these systems as well as the lack of any counterparts in the old population (see Section 6.3), they are likely to have formed near simultaneously from gaseous material stripped from M81.

The LF of Holm IX decreases gradually from g₀ = 25 to g₀ ∼ 22, where it remains flat up to g₀ ∼ 21. The presence of brighter stars in this region suggests that it has sustained star formation for a longer period than the other systems. The Garland has a similar LF to Holm IX, but it has a small bump at g₀ ∼ 22.7. Stars in the M82S-arc and M82N-arc features populate the LF to g₀ ∼ 22.5, implying that they stopped forming stars earlier than the other systems.

Compared with other nearby galaxy groups, the M81 Group is unique because of the strong gravitational interactions that are ongoing among M81, M82, and NGC 3077. While tidal stripping has displaced gas from the main bodies of these systems and triggered new star formation within them, we have found little evidence for tidal stripping in the old dwarf galaxy population (e.g., signatures of extra-tidal stars in their radial profiles or distorted isophotes).

Peñarrubia et al. (2009) investigated the effect of tidal stripping on the structural properties of a dSph galaxy. They showed that the radial surface brightness profile deviates substantially from a King profile as a result of the tidal interaction. Once the dSph equilibrates again, the outer density profile settles into a power law varying as R⁻⁴. An excess of stars with respect to a King profile has been found in the outer parts of some Milky Way dSphs, such as Fornax and Ursa Minor, and cited as evidence for tidally stripped populations (e.g., Palma et al. 2003; Coleman et al. 2005). Marginal overdensities can be seen at around R_t in the radial profiles of d1005+68 and d1015+69 shown in Figure 6, but it is arguable whether they are real features or simply due to the poor statistics. Overall, the old dwarf satellite population seems to have been unaffected by the strong tidal interactions that are shaping the properties of the most massive galaxies.

The photometric properties of the dwarf galaxies derived in this study are similar to those of early-type dwarf galaxies in the Local Group. Figure 20 shows the comparison of the total magnitudes, mean metallicities, and sizes of dwarf galaxies in the nearby universe. The M81 Group galaxies studied in Section 3 are shown as black diamonds. Other M81 Group galaxies, DDO44, DDO71, DDO78, F6D1, and F12D1, from Lianou et al. (2010), are plotted as gray diamonds. The filled circles, boxes, and triangles indicate morphological types of dwarf elliptical (dE) or spheroidal (dSph), irregular (Irr) or dwarf irregular (dIrr), and transitional type, respectively, taken from McConnachie (2012).

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We fit the luminosity–metallicity relation of the dE and dSph galaxies in the Local Group with a linear least-square fit, shown as the dashed line. Magenta, cyan, green, and yellow colors of the symbols represent Galactic satellites, M31 satellites, dwarf galaxies in the Local Group, and those in the nearby field (<3 Mpc), respectively. The radii of early-type galaxies in the Virgo cluster taken from Misgeld & Hilker (2011) and the metallicities of nuclei belonging to these galaxies, taken from Spengler et al. (2017), are also presented as purple circles. Note that the different methods adopted for these metallicity and luminosity measurements will contribute to the scatter in Figure 20.

In the luminosity–metallicity plot shown in the upper panel of Figure 20, M81 Group galaxies can be seen to obey the same relation as that of dE and dSph galaxies in the Local Group, shown as the dashed line ([M/H] = −0.12 × M_V,0 − 2.83). They appear to have somewhat higher metallicities than nearby isolated dIrr galaxies (green symbols) of the same total magnitude. The nuclei of luminous early-type galaxies in the Virgo cluster have higher mean metallicities than the luminosity–metallicity relation of fainter early-type galaxies, conceivably due to a selection effect on the calculation of the mean metallicities through the spatial extent of galaxies. Spengler et al. (2017) estimated the mean metallicities from Lick spectral indices of nuclei, where more metal-rich stars are likely to be concentrated. Lianou et al. (2010) and Tudorica et al. (2015) pointed out that IKN has a higher mean metallicity for its luminosity compared with other dE and dSph galaxies in the M81 Group, and hence that IKN may have formed as a tidal dwarf galaxy. However, the significantly brighter M_V,0 calculated for IKN in this study places it on the same metallicity–luminosity relation as for other early-type galaxies of the M81 Group. This updated luminosity for IKN also means that it has a less extreme globular cluster specific frequency (Larsen et al. 2014).

The lower panel of Figure 20 shows the size–luminosity relationship of the dwarf galaxy sample. The relationships defined by the M31 and MW dSph galaxies are statistically consistent with each other (Brasseur et al. 2011; Tollerud et al. 2012). Our measurements for the M81 Group dwarf galaxies indicate that they also obey the same trend.

A long-standing question is whether the young stellar systems in the M81 Group are purely young stellar associations that have formed in gaseous tidal debris, or whether they are gravitationally bound early-type dwarf galaxies that simply have elevated star formation rates at the present epoch (e.g., Hopp & Schulte-Ladbeck 1987; Hill et al. 1993). Deep HST images and chemical abundance measurements support the former idea for Holm IX, AL, BK3N, and the Garland (Makarova et al. 2002; de Mello et al. 2008; Sabbi et al. 2008; Weisz et al. 2008; Croxall et al. 2009). Specifically, Weisz et al. (2008) did not identify RGB overdensities at the positions of Holm IX and the Garland in deep HST data, even although such overdensities were associated with the other dwarf galaxies in their sample. Barker et al. (2009) investigated this same question using deep wide-field photometry taken with Subaru's Suprime-Cam, and also concluded there was no distinct population of RGB stars associated with AL. In this section, we revisit this question for all of the prominent young stellar systems in the core of the M81 Group using our deep HSC photometry. Although our data suffer from significant incompleteness due to crowding in the innermost regions of some systems, they still enable a useful exploration of the spatial distributions of young and old stars in the vicinity of the objects discussed in Section 5.

Figures 21 and 22 show the spatial distributions of young MS/cHeB stars, blue and red RGB stars, and stars located between (g − i)₀ = 0.75 and the boundary of the blue RGB criterion in Figure 16 (named "blue objects" in the figures). The MS/cHeB stars and blue and red RGB stars are selected using the dashed and dotted polygons shown in Figure 16. The apparent absence of RGB stars at the center of Holm IX (Figure 21) as well as at the centers of M82 and NGC 3077 (see Figure 22) is due to the severe incompleteness resulting from crowding.

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There are no obvious counterparts of the young systems in the red RGB distributions in Figures 21 and 22. In the blue RGB maps, there is marginal evidence for diffuse enhancements at the positions of AL and the Garland. However, it is unlikely that these are very metal-poor ([M/H] ∼ −2) and old (> 10 Gyr) stars associated with young objects; they are more likely to be "young" blue loop stars of ages 100–400 Myr old that fall within the blue RGB selection box of Figure 16. Indeed, similar structures can be seen in the blue object maps of these systems, which will also capture blue loop stars as well as unresolved background galaxies. Overdensities of blue objects, likely due to blue loop stars, can also be seen in Holm IX, the NW-arm, and the M82S-arc.

In summary, our spatial distributions of stars in Figures 21 and 22 reveal no evidence for old populations associated with the young systems, down to the limit of our photometry. This adds further support to the idea that these are genuinely new stellar systems that have formed in gaseous tidal debris, as previously suggested for Holm IX, the Garland, and AL, the most prominent examples of such features known (Weisz et al. 2008; Barker et al. 2009). If these structures have internal kinematics that indicate they are self-gravitating, then they can be considered tidal dwarf galaxies that will likely survive for a long time. (e.g., Duc et al. 2000; Hibbard et al. 2001).