# A stellar overdensity associated with the Small Magellanic Cloud

A. Pieres, B.X. Santiago, A. Drlica-Wagner, K. Bechtol, R.P. van Der Marel, G. Besla, N.F. Martin, V. Belokurov, C. Gallart, D. Martinez-Delgado, et al.

► To cite this version:

A. Pieres, B.X. Santiago, A. Drlica-Wagner, K. Bechtol, R.P. van Der Marel, et al. A stellar overdensity associated with the Small Magellanic Cloud. Monthly Notices of the Royal Astronomical Society, 2017, 468 (2), pp.1349-1360. 10.1093/mnras/stx507. hal-01582789

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# A stellar overdensity associated with the Small Magellanic Cloud

A. Pieres,¹,²* B. X. Santiago,¹,² A. Drlica-Wagner,³ K. Bechtol,⁴

R. P. van der Marel,⁵ G. Besla,⁶ N. F. Martin,⁷,⁸ V. Belokurov,⁹ C. Gallart,¹⁰,¹¹

D. Martinez-Delgado,¹² J. Marshall,¹³ N. E. D. Nöel,¹⁴ S. R. Majewski,¹⁵

M.-R. L. Cioni,¹⁶,¹⁷ T. S. Li,³ W. Hartley,¹⁸ E. Luque,¹,² B. C. Conn,¹⁹

A. R. Walker,²⁰ E. Balbinot,¹⁴ G. S. Stringfellow,²¹ K. A. G. Olsen,²² D. Nidever,²²

L. N. da Costa,²,²³ R. Ogando,²,²³ M. Maia,²,²³ A. Fausti Neto,² T. M. C. Abbott,²⁰

F. B. Abdalla,²⁴,²⁵ S. Allam,³ J. Annis,³ A. Benoit-Lévy,²⁵,²⁶,²⁷ A. Carnero Rosell,²,²³

M. Carrasco Kind,²⁸,²⁹ J. Carretero,³⁰,³¹ C. E. Cunha,³² C. B. D'Andrea,³³,³⁴ S. Desai,³⁵

H. T. Diehl,³ P. Doel,²⁵ B. Flaugher,³ P. Fosalba,³⁰ J. García-Bellido,³¹ D. Gruen,³²,³⁶

R. A. Gruendl,²⁸,²⁹ J. Gschwend,²,²³ G. Gutierrez,³ K. Honscheid,³⁷,³⁸ D. James,²⁰,³⁹

K. Kuehn,⁴⁰ N. Kuropatkin,³ F. Menanteau,²⁸,²⁹ R. Miquel,⁴¹,⁴² A. A. Plazas,⁴³

A. K. Romer,⁴⁴ M. Sako,⁴⁵ E. Sanchez,⁴⁶ V. Scarpine,³ M. Schubnell,⁴⁷

I. Sevilla-Noarbe,⁴⁶ R. C. Smith,²⁰ M. Soares-Santos,³ F. Sobreira,²,⁴⁸

E. Suchyta,⁴⁹ M. E. C. Swanson,²⁹ G. Tarle,⁴⁷ D. L. Tucker³ and W. Wester³

*Affiliations are listed at the end of the paper*

Accepted 2017 February 23. Received 2017 February 23; in original form 2016 December 12

## ABSTRACT

We report the discovery of a stellar overdensity 8° north of the centre of the Small Magellanic Cloud (SMC; Small Magellanic Cloud Northern Over-Density; SMCNOD), using data from the first 2 yr of the Dark Energy Survey (DES) and the first year of the MAGellanic Satellite LITEs Survey (MagLiteS). The SMCNOD is indistinguishable in age, metallicity and distance from the nearby SMC stars, being primarily composed of intermediate-age stars (6 Gyr, Z=0.001), with a small fraction of young stars (1 Gyr, Z=0.01). The SMCNOD has an elongated shape with an ellipticity of 0.6 and a size of ~6° × 2°. It has an absolute magnitude of $M_V \approx -7.7$, $r_h = 2.1$ kpc, and $\mu_V(r < r_h) = 31.2$ mag arcsec$^{-2}$. We estimate a stellar mass of ~$10^5$ M$_\odot$, following a Kroupa mass function. The SMCNOD was probably removed from the SMC disc by tidal stripping, since it is located near the head of the Magellanic Stream, and the literature indicates likely recent Large Magellanic Cloud-SMC encounters. This scenario is supported by the lack of significant H I gas. Other potential scenarios for the SMCNOD origin are a transient overdensity within the SMC tidal radius or a primordial SMC satellite in advanced stage of disruption.

**Key words:** galaxies: interactions – Magellanic Clouds.

## 1 INTRODUCTION

The Magellanic Clouds (MCs) are a rich and nearby system where we can observe dynamic evolution as well as the results of star formation throughout time. The system also includes the Magellanic Stream (MS), an H I gas stream (Mathewson, Cleary &

Murray 1974) connected to the MCs spanning at least 200° on the sky (Nidever et al. 2010), where no stellar counterpart has yet been identified (Recillas-Cruz 1982; Guhathakurta & Reitzel 1998). Other important structures belonging to this system are the Magellanic Bridge, containing neutral hydrogen, stars and star clusters linking the MCs (Irwin, Demers & Kunkel 1990; Grondin, Demers & Kunkel 1992; Bica et al. 2015) and the Leading Arm (LA) or Leading Arm Feature, a gas stream on the opposite side of the MS.

<footnote>* E-mail: adriano.pieres@gmail.com; adriano.pieres@ufrgs.br</footnote>
Given the higher velocities (than previously estimated) for the MCs in recent works (Kallivayalil et al. 2006a; Kallivayalil, van der Marel & Alcock 2006b; Vieira et al. 2010; Kallivayalil et al. 2013), it is thought that the MCs are completing their first passage around the Milky Way (MW). This conclusion is supported by proper motion measurements using *HST* (Kallivayalil et al. 2013) and *Gaia* data release 1 (van der Marel & Sahlmann 2016). Thus, the gravitational interaction between the Small and Large Magellanic Clouds (SMCs and LMCs, respectively) may be playing a larger role than the MW in triggering star formation.

In the recent decades, a wide range of dynamical simulations of the MCs has improved our understanding of their substructures, taking advantage of more reliable proper motion measurements, among other enhanced initial conditions (e.g. masses, gas fraction, ellipticity, stellar disc scalelength). Using N-body simulations, Connors, Kawata & Gibson (2006) reproduced the MS and LA as substructures formed through tidal interaction between the clouds. Their work reproduced for the first time the spatial and kinematic bifurcations in the LA and in the MS. The MCs simulations of Bekki & Chiba (2007) over the last 800 Myr are able to reproduce the off-centre bar and the H1 spirals in the LMC. They also predict that a substantial number of SMC stars could be transferred to the LMC to form diffuse halo components around that galaxy. Restricting their study to the SMC, Bekki & Chiba (2009) designed chemodynamical simulations, using an SMC ‘dwarf spheroidal model’ (an extended H1 gas disc and a spherical distribution for old stars), which they argue is a better description of the stellar and gas kinematic properties. In their fiducial model, the final distribution of old stars is more regular (spherical) than that of the younger stars (which form basically a bar-like structure). Diaz & Bekki (2012) simulate a large set of models based on proper motion data from Vieira et al. (2010) and Kallivayalil et al. (2006b,a), predicting two main encounters between the SMC and the LMC (260 Myr and 1.97 Gyr ago), suggesting a joint history for these galaxies. In their simulations, the first encounter forms two substructures: the Magellanic Bridge and a less obvious structure called Counter-Bridge. Besla et al. (2012) present two models for the Magellanic System, designed to explain the MS as the action of LMC tides on the SMC. In their models, the LMC is a one-armed spiral and features as well as a warped, off-centre stellar bar as a result of the gravitational interaction.

The possible association between the MCs and ultrafaint dwarf galaxies recently discovered in the Dark Energy Survey (DES; The Dark Energy Survey Collaboration 2005) footprint has revived the search for dwarf galaxy satellites of the LMC or SMC (Deason et al. 2015; Drlica-Wagner et al. 2015, 2016; Koposov et al. 2015; Jethwa, Erkal & Belokurov 2016; Sales et al. 2016). The recent discovery by Mackey et al. (2016) of a stellar cloud with a length of 10 kpc within the LMC tidal radius (and an additional extension farther west of the LMC) shows that the exploration of the outer area of the MCs has an important potential for new discoveries. It also reinforces the idea that newly discovered structures can be used to trace the gravitational interaction history of the MCs. Besla et al. (2016) suggest that the existence of stellar arcs and multiple spiral arms in the northern LMC periphery (without comparable counterparts in southern regions of the SMC) could be attributed to repeated close interactions between the LMC and the SMC. A large number of simulations predict clumpy substructures formed by a spheroidal distribution surrounding the SMC (see e.g. the references listed in Section 4.3), though there is no specific prediction of overdensities as large as those presented here. Nevertheless, the discovery of this structure reinforces the scenario where the LMC and the SMC have had recent and drastic encounters.

<table><caption>**Table 1.** List of bands and central positions for each MagLiteS DECam exposures used in this work.</caption><thead><tr><th>Band</th><th>α(°)</th><th>δ(°)</th></tr></thead><tbody><tr><td rowspan="6">g</td><td>14.092</td><td>-65.826</td></tr><tr><td>10.690</td><td>-65.029</td></tr><tr><td>7.280</td><td>-65.861</td></tr><tr><td>10.804</td><td>-66.803</td></tr><tr><td>9.789</td><td>-65.893</td></tr><tr><td>13.324</td><td>-66.744</td></tr><tr><td rowspan="6">r</td><td>16.504</td><td>-65.664</td></tr><tr><td>13.055</td><td>-64.994</td></tr><tr><td>14.114</td><td>-65.815</td></tr><tr><td>10.699</td><td>-65.036</td></tr><tr><td>7.306</td><td>-65.856</td></tr><tr><td>10.807</td><td>-66.783</td></tr><tr><td></td><td>9.792</td><td>-65.893</td></tr><tr><td></td><td>13.322</td><td>-66.729</td></tr><tr><td></td><td>16.524</td><td>-65.653</td></tr><tr><td></td><td>13.077</td><td>-65.004</td></tr></tbody></table>

In what follows, we report a stellar overdensity located 8° north of the SMC centre, hereafter referred to as the Small Magellanic Cloud Northern Over-Density (SMCNOD). The SMCNOD was discovered in data from DES and follow-up imaging was performed with the Dark Energy Camera (DECam; Flaugher et al. 2015) as part of the MAGellanic SatelliteLITES Survey – MagLiteS. The data sets and criteria used to select stellar sources are discussed in Section 2. In Section 3, we describe the analysis of the stellar populations and the structure of the SMCNOD. We conclude by discussing the SMCNOD stellar population and gas content, as well as its formation and fate, in Section 4.

## 2 DATA

The DES data used in this work is the year-2 quick release (DES-Y2Q1) catalogue, constructed using 26 590 DECam exposures. The DES-Y2Q1 images were taken between 2013 August and 2014 February and between 2014 August and 2015 February, in the first 2 yr of the survey. The images cover most of the DES footprint (5000 square degrees), with the exception of a few hundred square degrees in the region near the South Galactic Pole. We refer to section 2 of Drlica-Wagner et al. (2015) for a detailed description of the data and the star selection criteria.

MagLiteS is a National Optical Astronomy Observatories (NOAO) community survey (NOAO proposal 2016A-0366) that is using the DECam to complete an annulus of contiguous imaging around the periphery of the Magellanic System (Drlica-Wagner et al. 2016).

The MagLiteS data used here are composed of 16 DECam 90 s exposures in the *g* and *r* bands. The positions for each MagLiteS DECam exposure are listed in Table 1. The MagLiteS exposures were taken in 2016 June 27, in an effort to enlarge DECam coverage in the SMCNOD region. MagLiteS images were reduced using the Dark Energy Survey Data Management (DESDM) pipeline, and source detection was performed separately on each exposure.

To assemble a combined DES-Y2Q1 and MagLiteS source catalogue, we first set the zero-points by comparing DES-Y2Q1 bright stars to individual MagLiteS DECam single-exposure catalogues. Since the DES-Y2Q1 is a de-reddened catalogue, we applied an extinction correction to each DES-Y2Q1 source following Schlegel, Finkbeiner & Davis (1998). Comparison
<figure><img src="image_3.png" /><figcaption>**Figure 1.** DES-Y2Q1 stellar density map using the HEALPIX scheme with NSIDE = 64 (0.839 square degrees), corrected by a survey coverage map and saturating at 2 stars arcmin⁻<fnref n="2" />. Objects detected in the search for extended structures are explicitly labelled in the DES-Y2Q1 footprint. In this density map, we are counting all of the DES-Y2Q1 stars. Inserted box: zoomed view of the SMCNOD region.</figcaption></figure>

stars were selected in the magnitude range of 17 < *g* < 21, |*wavg_spread_model*| < 0.003 and *flags* < 4 in each band (g and *r*) in the DES-Y2Q1 catalogue. SPREAD MODEL<fnref n="1" /> is a morphological output from SExTRACTOR<fnref n="2" /> used to distinguish stars from galaxies. The prefix *wavg* means we used the weighted average of *spread_model* from individual single-epoch detections. The maximum positional deviation (object matching between DES-Y2Q1 and MagLiteS sources) was set to 1 arcsec. After adding the photometric zero-points, we joined all sources from the MagLiteS fields into a single catalogue. We then subtracted the extinction and incorporated the final MagLiteS catalogue into the DES-Y2Q1 catalogue, to create the final DES-MagLiteS stars list used in this paper. We applied the same criteria used to select DES-Y2Q1 stars to filter our final sample of stars, namely using a star/galaxy separation criterion of |*wavg_spread_model_r*| < 0.003 + *spreaderr_model_r*, *flags*-{g, *r*} < 4, and *magerr-psf*-{g, *r*} < 1.

Moreover, we applied a magnitude cut of 17 < *g* < 23 to ensure high source detection efficiency on the DES-MagLiteS catalogue. Also, we applied a colour cut to select stellar sources with −0.5 < *g* − *r* < 1.2.

A single 90 s DECam exposure (the DECam exposure time for DES in *g* and *r* bands) reaches point sources with magnitudes as faint as {*g*, *r*} ≅ {23.6,23.2} with a signal-to-noise ratio (S/N) equal to 10. Therefore, the faint magnitude cut adopted here results in uniform depth at least down to this S/N level. We emphasize that the quoted limiting magnitudes and S/N may change slightly due to seeing and weather conditions during the observing nights.

## 3 ANALYSIS

The SMCNOD was discovered during a search for extended (*r*ᵢ ≅ 30 arcmin) and low surface brightness structures in the DES-Y2Q1 catalogue. We initially built density maps for the DES-Y2Q1 stars, partitioning the sky into equal area HEALPIX.<fnref n="3" /> We set the pixel area to 0.839 square degrees (NSIDE = 64). We then counted all DES-Y2Q1 stars within each pixel, correcting the density in each pixel by the respective survey coverage for both *g* and *r* bands. The coverage maps were created by using a finer grid of pixels (pixel size ≅ 1.7 arcmin on a side, NSIDE = 2048), then checking whether or not a pixel contains any star or galaxy and finally grouping into pixels with 0.839 square degrees (NSIDE = 64), where the effective survey coverage area was computed. We then calculated the average number of stars in the eight immediately neighbouring pixels (with NSIDE = 64). The significance of any overdensity was calculated by subtracting the average counts in the neighbouring pixels and dividing the result by the square root of that average, thus yielding the number of standard deviations (following a Poisson distribution) of the star counts. For example, the least significant candidate has a star count equal to 3781, whereas the average counts of the neighbouring pixels is 1.0 star per arcmin<fnref n="2" /> (3600 stars per square degree). Its significance is then only 3σ ((3781−3600)/√3600), presenting an excess of 5 per cent above the mean star counts. We examined all candidates with significance greater than 3σ, which results in a list of 314 candidates.

The highest significance candidates were mostly known globular clusters and dwarf galaxies (see Fig. 1). However, one candidate

<footnote n="1">¹ SPREAD MODEL is a ‘normalized simplified linear discriminant between the best-fitting local point spread function model and a slightly more extended model’ as described in Desai et al. (2012).</footnote>

<footnote n="2">² http://www.astromatic.net/software/sexttractor</footnote>

<footnote n="3">³ HEALPIX is an equal-area pixelization scheme for spherical surfaces (in our case, the sky) in a certain number of pixels. This number of pixels is given by 12 times the square of the parameter NSIDE, chosen by the user. See more details in http://healpix.sourceforge.net/.</footnote>
<figure><img src="image_4.png" /><figcaption>**Figure 2.** First panel: Hess diagram for stars in the range $6^{\circ} < \alpha < 15^{\circ}$ and $-66^{\circ} < \delta < -63^{\circ}$, which covers most of the SMCNOD. A stellar population with a turnoff at $g \approx 22$ can be clearly seen. A main sequence, sub-giant branch, RGB and RC are all discernible against the foreground Galactic stars even with no subtraction. Second panel: Hess diagram for a region at the same Galactic latitude ($b = -52^{\circ}58$) as the previous one ($32^{\circ} < \alpha < 40^{\circ}$, $-64^{\circ} < \delta < -58^{\circ}25$). Third panel: subtracted Hess diagram (object minus field). Fourth panel: same as the third panel, but featuring isochronal masks of intermediate-age (young) SMC populations bounded by the solid black (dashed blue) line, encompassing most of the SMCNOD stars by displacing the isochrone. The black rectangle denotes the RC stars. A PARSEC model with $\tau \simeq 6$ Gyr and $Z = 0.001$ ($\tau \simeq 1$ Gyr, $Z = 0.01$) is shown in dotted black (solid thin blue) line. Both models (intermediate-age and young) are displaced by a distance modulus equal to 18.96, following the SMC distance modulus obtained by de Grijs & Bono (2015). Fifth panel: sample of SMC field stars with $\delta < -67^{\circ}$ and $10^{\circ} < \alpha < 15^{\circ}$ minus the second Hess diagram (MW foreground stars). The data show a spread in magnitude that is comparable between the SMC stars and SMCNOD stars. The isochronal mask for intermediate-age stars is reproduced in the last panel.</figcaption></figure>

located at $\alpha \simeq 12^{\circ}$ and $\delta \simeq -65^{\circ}$ was significantly higher (with a significance of $8\sigma$ at the highest density pixel) than the local background, and several of its neighbouring pixels emerged in the significance list, suggesting that the overdensity spans multiple pixels (insert in Fig. 1). Given its proximity to the SMC, we refer to this overdensity as the SMCNOD. This object is located in the border of the DES survey and so we performed follow-up imaging with MagLiteS to cover an extra area around SMCNOD.

In the leftmost panel of Fig. 2, we plot the *g* versus *g*-*r* colour-magnitude Hess diagram for the region surrounding the SMCNOD in the DES-MagLiteS catalogue, to analyse the photometric features of that putative stellar population. The second Hess diagram samples stars in a field with 20 square degrees, centred on $l = 304^{\circ}60$ and $b = -52^{\circ}60$ at the same Galactic latitude as the SMCNOD ($l = 284^{\circ}72, b = -52^{\circ}60$). Subtracting the first two Hess diagrams (and weighting by their respective areas) results in the third Hess diagram in Fig. 2. It is dominated by a stellar population with age $\tau \simeq 6$ Gyr and metallicity $Z = 0.001$ as attested by the overlaid PARSEC model (Bressan et al. 2012) represented by the dotted black line in the fourth panel. This model was chosen by a visual comparison to

the Hess diagram in the fourth panel. A colour-magnitude diagram (CMD) mask is drawn (solid black lines) displacing the PARSEC model for intermediate-age stars in *g* − *r* colour and *g* magnitude. The SMCNOD distance modulus is indistinguishable from that of the SMC (18.96 ± 0.02 following de Grijs & Bono 2015). A PARSEC model is also shown in the fourth panel of Fig. 2 (thin blue line) to represent the blue plume of younger stars ($\tau \simeq 1$ Gyr and $Z=0.01$). We note that there is some overlap between both populations (younger and intermediate-age) in the lower main sequence, red giant branch (RGB) and red clump (RC) CMD regions. In the last Hess diagram, stars from the DES-MagLiteS catalogue with $-68^{\circ} < \delta < -67^{\circ}$ and $10^{\circ} < \alpha < 15^{\circ}$ (the closest region to the SMC in the DES-MagLiteS catalogue) are sampled and the CMD mask for intermediate-age stars is reproduced to compare the SMC and SMCNOD stellar populations.

The PARSEC model for intermediate age is a good description of the SMCNOD population, and we selected stars that are more likely to belong to the object using the CMD filter described above. As a young population is also visible in the third and fourth subtracted Hess diagrams, we added an extra filter box to include the younger
<figure><img src="image_5.png" /><figcaption>**Figure 3.** Top left: density map (each pixel has 14 × 14 arcmin, NSIDE=256) for stars filtered by the isochrone mask shown in the fourth panel of Fig. 2, in the field surrounding the SMCNOD. Top right: the same density map as in the left, but corrected by the coverage map. Bottom: density map for young (left) and intermediate-age (right) stellar population. Both bottom panels are corrected by the coverage map. The object in the top-right corner of each panel at α, δ ≃ {1°, −61°} is the dwarf galaxy candidate Tucana IV (Drlica-Wagner et al. 2015). Top-left panel and both right-hand panels are sharing the rightmost colourbar, while the young population is shown in a different colourbar scale, to highlight its weak contribution.</figcaption></figure>

main-sequence stars. Using both CMD filters described above (for intermediate-age and young stars), we reanalyzed the stellar density distribution in the DES-MagLiteS catalogue. First, we built the stellar density map (top-left panel in Fig. 3) using NSIDE=256 (pixel size ≅ 14 arcmin on a side). Dividing this stellar density map by the coverage map results in the stellar density map shown in top-right panel in Fig. 3. We now see a stellar overdensity with a roughly elliptical shape, mainly composed of intermediate-age stars (comparing both bottom panels of Fig. 3) at a distance of 8° from the SMC centre (α = 13:000, δ = −72:817).

We follow the model from Noël & Gallart (2007) to compare the SMCNOD brightness to the expected SMC surface brightness extrapolated to that position. They fit the SMC surface brightness profile (in B and R bands) using three 34 arcmin × 33 arcmin fields located southwards of the SMC, at a distance of 4.7, 5.6 and 6.5 kpc (respectively, 4.2, 4.9 and 5.8). Extrapolating their surface brightness profile out to a radial distance of 8° from the SMC centre, we derive an expected B-band surface brightness of μ<sub>B</sub> = 32.4 ± 0.3 mag arcsec<sup>-2</sup>. To compare to the SMCNOD surface brightness, we first applied a transformation of stellar magnitudes from g and r DES bands to g and r SDSS bands, following Bechtol et al. (2015):

$$g_{\text{SDSS}} = g_{\text{DES}} + 0.104(g_{\text{DES}} - r_{\text{DES}}) - 0.01 \quad (1)$$

$$r_{\text{SDSS}} = r_{\text{DES}} + 0.102(g_{\text{DES}} - r_{\text{DES}}) - 0.02. \quad (2)$$

We then converted the SDSS magnitudes from the CMD filtered stars to the B band, using the transformation equation from Jester et al. (2005):

$$B = g_{\text{SDSS}} + 0.390(g_{\text{SDSS}} - r_{\text{SDSS}}) + 0.21. \quad (3)$$

We evaluate the integrated B flux at the SMCNOD centre, in the same HEALPIX pixels (NSIDE = 256) applied before, obtaining a surface brightness of μ<sub>B</sub> = 29.7 ± 0.17 mag arcsec<sup>-2</sup>. This is almost 3 mag brighter than expected from extrapolating the main body of the SMC based on Noël & Gallart (2007). The uncertainties were estimated using a bootstrap method, where the stars in the central pixel were randomly sorted (with replacement) to make up a new estimate of the brightness in the B band. A total of 1000 such bootstrap realizations were carried out.

Nidever et al. (2011) explored the SMC RGB distribution, using data from the MAGellanic Periphery Survey (MAPS), sampling 36 arcmin × 36 arcmin fields with the MOSAIC II Camera mounted on the CTIO 4 m Blanco telescope, and reaching as far as 12° from the SMC centre. They observed stars with Washington photometry in three bands (DDO51, M and T<sub>2</sub>), as these bands are useful to discriminate MW foreground dwarfs from SMC RGB stars. The best-fitting elliptical density profile for the SMC giants sampled presents a ‘break’ at 7.5 from the fitted centre (α = 15°, 129, δ = −72°, 720), where the density slope abruptly decreases and the distribution of giants begins to scatter around this flatter profile. In Fig. 4, we
<figure><img src="image_6.png" /><figcaption>**Figure 4.** RGB density profile reproduced from fig. 3 of Nidever et al. (2011). The circles are the densities (as found by those authors) colour coded by position angle (from N to E). Also shown are the foreground contamination level and fitted profiles. The brown square is the density of giants as sampled in the SMCNOD centre, while the brown circle at 8.4 (indicated by a black arrow) is the field 84S341, which overlaps with DES-MagLiteS and which was used to re-normalize DES-MagLiteS density scale. The SMC centre adopted here is at $\alpha = 15.129, \delta = -72.720$ and a distance of 61.94 kpc from the Sun (de Grijs & Bono 2015).</figcaption></figure>

reproduce fig. 3 from Nidever et al. (2011) using colour-coded circles according to the field position angles. The figure also reproduces the best-fitting models, both internal and external to the profile break.

To compare the density of RGB stars at the centre of the SMCNOD to the measurements of Nidever et al. (2011), we subtracted the Galactic foreground dwarf stars contaminating the RGB locus and we normalized our densities to the density profile shown in Fig. 4. This second goal is achieved with the use of the only field from Nidever et al. (2011) that overlaps the DES-MagLiteS footprint, which they name 84S341 and which is located 2.2 away from the SMCNOD centre. Table 2 lists the positions and DES-MagLiteS stellar densities at the SMCNOD centre, at the 84S341 field and at a field far from the SMCNOD. The three fields are at nearly the same Galactic latitude and we assume that the difference in MW dwarf counts is negligible. The densities listed in the second to last column correspond to the stars falling within the intermediate-age isochrone CMD mask described earlier with an additional colour–magnitude cut of $g < 21$ applied and with no RC stars included. We refer to that filter as the RGB box. The first line in the last column represents the resulting RGB density after subtracting the foreground

<figure><img src="image_7.png" /><figcaption>**Figure 5.** Top: stellar density (black dots) versus angular separation (corrected for elliptical shape and position angle) of the SMC centre. The stellar density was determined in each HEALPIX (NSIDE=256) in DES-MagLiteS, and the angular separation corresponds to the elliptical exponential model from Nidever et al. (2011). The solid red line is the fit for HEALPIX between 5.5 and 7.5 and solid blue line for pixels between 10.0 and 14.0. The grey circles are the boxes within the SMCNOD position (cells within truncation radius). Bottom: stellar density data divided by the model for all HEALPIX pixels. The SMCNOD resides at the interface between exponential models, but is discrepant from both. The radial scalelength is 1.33 for the inner fit and 10.13 for the outer fit.</figcaption></figure>

contamination. The bracketed density value for the 84S341 field is the RGB density actually measured by Nidever et al. (2011). The final SMCNOD RGB density (also shown in brackets) is then obtained by applying the same ratio as in the 84S341 field (166 giants degree$^{-2}$), placing it clearly above the density profile of any of the individual fields analysed by Nidever et al. (2011) at that angular distance (Fig. 4).

To compare the SMCNOD stellar density to the surrounding areas, we fit two models: a profile closer to the SMC than the SMCNOD (called the inner profile) and a profile more distant of the SMC than the SMCNOD (the outer profile). The distances from the SMC centre to each HEALPIX pixel were set following the Nidever et al. (2011) elliptical exponential model. The density of CMD-filtered stars was calculated in HEALPIX pixels with NSIDE = 256. We fitted the inner (outer) profile for boxes between 5.5 and 7.5 (between 10° and 14°) from the SMC centre. The fits provide an independent and striking confirmation of the SMC extended profile, along with the break at $\simeq 8°$ from the SMC centre. The top panel in Fig. 5 shows both inner (red line) and outer (blue line) fits for the density profiles near the SMCNOD. Dividing the density by the respective fits (bottom panel in Fig. 5), we find that the HEALPIX

<table><caption>**Table 2.** Name (first column) and position in equatorial (columns 2 and 3) and Galactic (columns 4 and 5) coordinates for three fields: the SMCNOD centre, the field overlapping (Nidever et al. 2011) and an MW foreground field at roughly the same Galactic latitude. The numbers in the second to last column are the stellar density after applying our CMD filters. The last column presents the density of giants after subtracting the MW foreground density (247 stars degree$^{-2}$). Numbers in brackets in the last columns are normalizing to Nidever et al. (2011), used as reference. More details are given in the text.</caption><thead><tr><th>Field name</th><th>α (°)</th><th>δ (°)</th><th>l (°)</th><th>b (°)</th><th>ρ stars degree<sup>-2</sup></th><th>ρ giants degree<sup>-2</sup></th></tr></thead><tbody><tr><td>SMCNOOD centre</td><td>12.000</td><td>-64.800</td><td>303.529</td><td>-52.317</td><td>457</td><td>210 [166]<sup>a</sup></td></tr><tr><td>84S341</td><td>6.892</td><td>-64.741</td><td>307.082</td><td>-52.194</td><td>357</td><td>110 [87.2]<sup>b</sup></td></tr><tr><td>MW foreground</td><td>19.928</td><td>-64.600</td><td>297.985</td><td>-52.256</td><td>247</td><td>—</td></tr></tbody></table>

Notes. $^a$Density of giants estimated for the SMCNOD centre.

$^b$Density of giants from Nidever et al. (2011), used as reference. Both ($^a$ and $^b$) are plotted in Fig. 4.
<table><caption>**Table 3.** SMCNOD properties: equatorial (α and δ) and Galactic (l and b) coordinates of the centre, ellipticity, half-light radius, truncation radius, absolute magnitude and surface brightness. The last two properties are in the V band.</caption><thead><tr><th>Property</th><th>Value</th><th>Unit</th></tr></thead><tbody><tr><td>α</td><td>12.00<sup>+0.08</sup><sub>-0.06</sub></td><td>deg</td></tr><tr><td>δ</td><td>-64.80<sup>+0.05</sup><sub>-0.08</sub></td><td>deg</td></tr><tr><td>l</td><td>303.53</td><td>deg</td></tr><tr><td>b</td><td>-52.32</td><td>deg</td></tr><tr><td>ε</td><td>0.60<sup>+0.19</sup><sub>-0.20</sub></td><td>-</td></tr><tr><td>r<sub>h</sub></td><td>120.4<sup>+19.2</sup><sub>-3.12</sub></td><td>arcmin</td></tr><tr><td>r<sub>tr</sub></td><td>192 ± 20.0</td><td>arcmin</td></tr><tr><td>M<sub>V</sub></td><td>-7.7 ± 0.3</td><td>mag</td></tr><tr><td>μ<sub>V</sub>(r < r<sub>h</sub>)</td><td>31.23 ± 0.21</td><td>mag arcsec<sup>-2</sup></td></tr></tbody></table>

pixels within the SMCNOD truncation radius (where the densities decrease to the level of the background density) have notably higher densities than those of the surrounding areas.

The structural parameters listed in Table 3 were fit using a marginalized likelihood approach and EMCEE, an affine-invariant ensemble sampler for Markov Chain Monte Carlo models (Foreman-Mackey et al. 2013).<fnref n="4" /> We applied the marginalized likelihood fit to the HEALPIX pixels from Fig. 6, modelling the stellar density with a Plummer profile (Plummer 1911).

The absolute magnitude $M_V$ was derived by adding the V flux within the ellipse bounded by the SMCNOD truncation radius converted from the g and r bands, also using equation from Jester et al. (2005):

$$V = g_{SDSS} - 0.590(g_{SDSS} - r_{SDSS}) - 0.01. \quad (4)$$

To evaluate the V-band flux of the background, we added the flux within an ellipse with the same area shifted to 3° north of the SMCNOD and then subtracting from the SMCNOD flux in the V band. The $M_V$ uncertainty incorporates the spatial fluctuations in the background flux and is in fact dominated by them.

The stellar mass of the SMCNOD was estimated by comparing a luminosity function (LF) from a simulated simple stellar population to the SMCNOD LF. The SMCNOD LF is subtracted from a field LF immediately above the overdensity, with equal area and located 3° north of the SMCNOD centre. The subtracted LF comprises 6068 stars within the range $21.0 \le g \le 23.0$, corresponding to the mass range of 0.90–0.99 $M_\odot$. A simulated simple stellar population ($\tau=6$ Gyr and $Z=0.001$) with an evolved Kroupa mass function (MF; Kroupa 2001) was generated using GENCMD,<fnref n="5" /> populating the 0.90–0.99 $M_\odot$ mass range with the same star counts as the SMCNOD subtracted LF. The mass range between 0.1 and 1.02 $M_\odot$ amounts to a stellar mass for the SMCNOD $\simeq 1.1 \times 10^5 M_\odot$, and its resulting $M/L$ is very close to unity (1.07 $M_\odot/L_\odot$). The young population density is about one-tenth of the intermediate-age population density and thus the computed young population mass is included in the mass error range for SMCNOD. As a comparison, the SMCNOD has a stellar mass comparable to Galactic Globular cluster NGC 6287 (Gnedin & Ostriker 1997), but it is brighter (for NGC 6287, $M_V = -7.36$; following Harris 1996, updated 2010). Another estimate, using an evolved MF similar to that found for Palomar 5

by Koch et al. (2004), where the fainter stars were removed from the main body, yields an SMCNOD stellar mass $\simeq 8.0 \times 10^4 M_\odot$. These values show how the choice of an MF changes the total estimated stellar mass. Therefore, we interpret the first estimate as an upper limit for the SMCNOD stellar mass.

The dynamical mass of the SMCNOD (m) was estimated (in the case SMCNOD is bounded to the SMC) using equations (7–84) from Binney & Tremaine (2008):

$$\frac{m}{M} = 3 \left(\frac{r_J}{D}\right)^3, \quad (5)$$

where M is the SMC dynamical mass, $r_J$ is the SMCNOD tidal radius and D is the distance between the SMC and SMCNOD centres. Assuming the SMCNOD tidal radius as 1.5 (from the bottom panel in Fig. 5) and $D = 8^\circ$, we determined that $m/M \simeq 2.0 \times 10^{-2}$. Bekki & Stanimirović (2009) estimate that the SMC dark halo has a mass of $3 \times 10^9 M_\odot$ in the inner 3 kpc for a V-band mass-to-light ratio $\simeq 2$. This mass agrees reasonably well with SMC rotation curves. Using this conservative estimate for the SMC mass, the SMCNOD dynamical mass is $6 \times 10^7 M_\odot$, a six hundred times greater than calculated for the stellar mass. The large disagreement between the stellar mass and the dynamical mass calculated using equation (5) is an argument favouring the SMCNOD is a structure detached from the SMC. The uncertainties in the SMCNOD dynamical mass are dominated by the errors in the SMC dynamical mass, which is estimated as about 13 per cent (Bekki & Stanimirović 2009).

The H I gas map could provide more insight into the nature of the SMCNOD, as well as a possible connection to the SMC. The H I gas column density map from the GASS Third Data Release<fnref n="6" /> (Kalberla & Haud 2015) is shown in Fig. 7. While the LMC and the SMC H I gas contents are obvious, there is no apparent excess of gas associated with the position of the SMCNOD. We also use the GASS data to look for peaks in the velocity distribution of the gas within the velocity range from $-495$ to $495$ km s$^{-1}$ (1 km s$^{-1}$ steps). The emission for one square degree centred on the SMCNOD exhibits two main peaks: 94 and $186$ km s$^{-1}$ (which are shown in Figs 8 and 9, respectively). These velocities agree with the velocity field related to the MS at the SMCNOD position, at an MS longitude $L_{MS} \simeq -25^\circ$. See for example fig. 8 in Nidever et al. (2010). We discuss details about H I gas distribution in Section 4.

# 4 DISCUSSION

In this section, we discuss the characteristics of the SMCNOD: we compare its stellar populations to those of the SMC (Section 4.1), its gas content (Section 4.2), possible scenarios for its origin (Section 4.3) and in the last subsection we provide a brief summary of the discovery, discussing the SMCNOD fate and some prospects for future analyses (Section 4.4).

## 4.1 The SMCNOD and SMC stellar populations

The $(g-r)$ colour distributions of SMCNOD and SMC stars for 3 mag ranges in the Hess diagram from Fig. 2 are shown in the left-hand panel of Fig. 10, whereas their number counts in bins of g magnitude filtered by the CMD mask and normalized (to the areas) are presented in the right-hand panel of the same figure. The colour distributions look very similar. For RC stars (solid lines), there may be a slight preference for the SMCNOD being a little bluer than the

<footnote n="6">⁶ https://www.astro.uni-bonn.de/hisurvey/gass/</footnote>

<footnote n="4">⁴ http://dan.iel.fm/emcee/current/</footnote>

<footnote n="5">⁵ GENCMD yields position, magnitude and errors with a simple stellar population. See details in https://github.com/balbinot/gencmd.</footnote>
<figure><img src="image_8.png" /><figcaption>**Figure 6.** Ratio from Fig. 5 (bottom panel) projected on to the sky. The truncation radius from Table 3 is shown as a dotted black line.</figcaption></figure>

<figure><img src="image_9.png" /><figcaption>**Figure 7.** H I gas column density map from Kalberla & Haud (2015), in zenith equal area projection, close to the south celestial pole. A grid of equatorial coordinates (α and δ) is indicated. The SMC is close to the centre of the figure and the LMC is located near the bright spot in the lower-left corner. No significant excess of H I gas is observed at the position of the SMCNOD (dotted ellipse).</figcaption></figure>

SMC. However, a Kolmogorov–Smirnov (KS) test indicates that the two RC populations come from the same parent distribution (p = 0.42). As for the LF comparison, the two distributions look similar as well.

## 4.2 The SMCNOD and H I gas

The results of considering the H I gas velocity channels are inconclusive. The $v_{LSR}=186$ km s⁻¹ channel map (Fig. 9) shows a few links between the SMC and the SMCNOD, while the $v_{LSR}=94$ km s⁻¹ channel map (Fig. 8) shows a bar-shaped gas cloud detached from the SMC main body. A looping feature is visible in the

<figure><img src="image_10.png" /><figcaption>**Figure 8.** H I gas emission map from Kalberla & Haud (2015), in equatorial coordinates for the velocity channel 94 km s⁻¹ < $v_{LSR}$ < 95 km s⁻¹. The SMCNOD position is highlighted by an empty ellipse and the SMC centre by a plus symbol.</figcaption></figure>

$v_{LSR}=186$ km s⁻¹ channel map that could be the result of a weak gas inflow (from the SMC, counterclockwise). But in summary, the SMCNOD does not seem to contain a large amount of gas and it is currently unclear whether the H I gas features present in either velocity channel map are connected to it. It has been suggested that the drift rate of the MS gas away from the LMC is ~49 km s⁻¹, as indicated by Nidever, Majewski & Butler Burton (2008). Using these results, an age of 1.74 Gyr is expected for the MS. In this sense, the gas features surrounding the SMC should be very recent (a few hundred million years or even younger), showing a complex dynamics. Taken at face value, the gas properties around the SMC-NOD are more consistent with gas-poor dwarf spheroidals (Grebel, Gallagher & Harbeck 2003) and ultrafaint dwarf galaxies in the Local Group (Grcevich 2013).
<figure><img src="image_11.png" /><figcaption>**Figure 9.** Same as Fig. 8 but for velocity channel equal to 186 km s⁻¹ (186 km s⁻¹ < v<sub>LSR</sub> < 187 km s⁻¹).</figcaption></figure>

**4.3 The SMCNOD origin**

Regarding the formation of the SMCNOD, the most likely scenario is that this structure was formed by material pulled from the SMC disc through tidal stripping, given the recent N-body simulations as cited in Section 1. Following the classical galactic interaction theory of Toomre & Toomre (1972), many other works predict the existence of a Magellanic Counter-Bridge as a counterpart to the Magellanic Bridge (Diaz & Bekki 2012; Dobbie et al. 2014a). In the north-west part of the SMC, a kinematical substructure discovered by Dobbie et al. (2014a) is associated with the Magellanic Counter-Bridge, as an observational counterpart. The simulations also predict a spread of stars as a result of an LMC-SMC close encounters. See for example the tidal tail in fig. 5 of Gardiner & Noguchi (1996), the set of particles located south-west of the SMC in fig. 12 from

Yoshizawa & Noguchi (2003), and the SMC stellar distribution in figs 4 and 5 of Connors et al. (2006). Also, a conspicuous clump of young stars can be seen for the fiducial model simulated by Bekki & Chiba (2009) in their figs 5, 6 and 9, located 4° from the centre of the SMC, along with a stream-like feature on the opposite side. Earlier simulations present a spread of particles around the SMC, for example fig. 10 b of Fujimoto & Sofue (1976) and fig. 6 a of Murai & Fujimoto (1980). The substantial stream-like stellar overdensity in the northern periphery of the LMC centre recently discovered by Mackey et al. (2016), with characteristics similar to the SMCNOD, also corroborates this scenario for the SMCNOD formation based on close LMC-SMC interactions. Finally, a close encounter occurred ≈200 Myr ago is also claimed as an explanation for a 55 kpc stellar structure in the eastern SMC (Nidever et al. 2013), where likely the stars were tidally stripped from the SMC.

If the SMCNOD is the result of an LMC-SMC collision, a con-
temporaneous peak in star formation is expected in both galax-
ies. Unfortunately, the results of star formation history (SFH)
analyses have large uncertainties, and there is a significant spa-
tial variation for the SFH in the LMC (Holtzman et al. 1999;
Olsen 1999; Smecker-Hane et al. 2002; Harris & Zaritsky 2009;
Rubele et al. 2012) and SMC (Noël et al. 2009; Cignoni et al. 2013;
Rubele et al. 2015). SFH variations are larger when based on dif-
ferent models (isochrones) and/or stellar tracers (RGB, carbon or
variable stars) and the MCs SFH is still far from being fully char-
acterized in spite of much work. Even so, it is interesting to note
that most of the references listed above agree with a simultaneous
peak in star formation rate between 4 and 6 Gyr (also discussed
in Dobbie et al. 2014b), the age of the main stellar population of
the SMCNOD. A complete reconstruction of the SFH of the SMC-
NOD (and also a comparison to the SFH of various SMC regions)
is beyond the scope of this work.

Another possible scenario for the SMCNOD origin is a tidal
dwarf galaxy (TDG), an object formed as described by Elmegreen,
Kaufman & Thomasson (1993), where mainly gas is stripped from

<figure><img src="image_12.png" /><figcaption>**Figure 10.** Histograms in bins of colour (left) and in bins of g magnitude (right) filtered by the CMD mask for the SMCNOD (third panel in Fig. 2) and the SMC (fifth panel in Fig. 2). In the right-hand panel, the Poissonian uncertainties are shown as shaded areas. The stellar colour distribution is very similar between the SMC and SMCNOD in the left-hand panel (comparing dotted red and blue lines and dashed red and blue lines), as well as the histogram in bins of magnitude for the stars filtered with the CMD mask in the last two panels in Fig. 2 (right-hand panel).</figcaption></figure>
past mergers and resemble as dwarf galaxies, where the stars are
forming during/after the main encounter. The lack of dark mat-
ter in TDGs makes them fragile, leading to short lifetimes (a few
Gyr) as cohesive systems. This scenario for the SMCNOD forma-
tion is disfavoured due to its poor gas content and predominantly
intermediate-age stars (at least 6 Gyr), compared to an expected
young TDG stellar population (Duc 2012). On the other hand, nu-
merical simulations by Ploeckinger et al. (2014) show TDGs could
survive at least 3 Gyr, despite the lack of dark matter content. This is
corroborated by the existence of the relatively old TDG VCC2062
observed by Duc et al. (2007), where its parent galaxies have likely
merged.

The scenario where the SMCNOD is a primordial galaxy orbit-
ing and/or merging with the SMC could be favoured if the stellar
populations of the SMCNOD have narrower age and metallicity
ranges than those of the SMC. This would make the SMCNOD
more consistent with typical dwarf spheroidal galaxies.

An intriguing possibility for the origin of the SMCNOD is the
'resonant stripping' predicted by D'Onghia et al. (2009). This pro-
cess allows for an efficient removal of stars in a dwarf-dwarf en-
counter, where the smaller dwarf loses stars by a resonance be-
tween the angular frequency of its orbit and spin, changing the ratio
of the stellar to dark matter mass. Simulations also predict that
dwarf-disc galaxies will evolve into compact spheroidal systems
with stream-like and shell-like structures, resembling the SMC-
NOD shape around the SMC. Future deep photometric surveys of
the SMC/LMC outskirts could reveal similar structures, testing the
significance of this 'resonant stripping' in the model.

The relatively old age for most of the stars in the SMC-NOD rules
out its origin as being formed by H I gas from the MS. As Nidever
et al. (2008) point out, the expected age of putative MS stars should
be ≈2.0 Gyr, less than half of the characteristic age found for the
overdensity stars. The origin of the SMCNOD younger population
may or may not be attributed to the same physical mechanism as
the intermediate-age population. Since the SMC-LMC interaction
is known to have triggered star formation at recent times (as in the
case of the Magellanic Bridge, see for example Bica et al. 2015;
Noël et al. 2015), we cannot rule out that this interaction may be
responsible for the younger SMCNOD population.

4.4 Summary, prospects and fate of the SMCNOD

Using DES and MagLiteS data, we have found and analysed a
stellar overdensity located about 8° north of the SMC centre. The
stellar density and surface brightness associated with this feature
lie significantly above the values expected from the extrapolated
stellar profile of the SMC itself. This is true even when we con-
sider only the contribution from the RGB stars. Previous surveys
around the SMCNOD, such as the MAPS (Nidever et al. 2011),
Two Micron All Sky Survey (Skrutskie et al. 1997), Digitized Sky
Survey<fnref n="7" /> and Infrared Astronomical Satellite survey,<fnref n="8" /> among others,
did not reveal any stellar overdensity similar to the one measured
here. This may be due to either their non-contiguous area or their
lower photometric depth.

The fact that the structure discussed here has a density peak lying
significantly in excess of the expected SMC exponential density
profile (or above the combined SMC and Galactic background)
and follows a roughly elliptical profile encourage us to argue that

it is a distinct SMC substructure. On the other hand, the CMD
analysis indicates that the stellar populations are very similar to
those found in the main SMC body, and the SMCNOD lies at a
similar heliocentric distance as the SMC.

The fate of the SMCNOD has interesting implications for the
Magellanic System generally. If the SMCNOD is an unstable object,
such as a stellar cloud, it should dissipate into the SMC main body
or, if unbound to the SMC, be ejected and dissipate eventually into
the Galactic field. However, for the first hypothesis to hold, the
SMC truncation radius must be ≈10 kpc (see the SMCNOD limits
in Fig. 5) to encompass the entire stellar cloud presented here. As
a reference, the truncation radius derived from chemodynamical
simulations involving SMC-like objects is in the range between 5
and 7.5 kpc (Bekki & Chiba 2007, 2009; Diaz & Bekki 2012), not
enough to include the entire SMCNOD.

Radial velocities and proper motions of likely stellar members
will constrain systemic and internal kinematics of the SMCNOD,
as well as its internal motions. Metallicities and other abundance
estimates may indicate similarities and differences between the SM-
CNOD stars and those belonging to the main SMC body. An internal
age and/or metallicity gradient (or its absence) may also constrain
its nature as either a primordial or tidal object. Finally, the SMC-
NOD discovery shows that the Magellanic System, despite being
relatively well studied, still hides surprising substructures that may
be revealed with deep photometric surveys. The discovery of the
SMCNOD at such a large distance from the SMC should provide
an additional constraint for simulations of the Magellanic System.

ACKNOWLEDGEMENTS

We are grateful to the anonymous referee for valuable comments,
which contributed to improving this paper. This paper has gone
through internal review by the DES and MagLiteS Collaboration.
EB acknowledges financial support from the European Research
Council (ERC-StG-335936). DMD acknowledges support by Son-
derforschungsbereich (SFB) 881 ‘The Milky Way System’ of the
German Research Foundation (DFB), subproject A2.

Funding for the DES Projects has been provided by the US De-
partment of Energy, the US National Science Foundation, the Min-
istry of Science and Education of Spain, the Science and Technol-
ogy Facilities Council of the United Kingdom, the Higher Education
Funding Council for England, the National Center for Supercomput-
ing Applications at the University of Illinois at Urbana-Champaign,
the Kavli Institute of Cosmological Physics at the University of
Chicago, the Center for Cosmology and Astro-Particle Physics at
the Ohio State University, the Mitchell Institute for Fundamental
Physics and Astronomy at Texas A&M University, Financiadora
de Estudos e Projetos, Fundação Carlos Chagas Filho de Amparo à
Pesquisa do Estado do Rio de Janeiro, Conselho Nacional de Desen-
volvimento Científico e Tecnológico and the Ministério da Ciência,
Tecnologia e Inovação, the Deutsche Forschungsgemeinschaft and
the Collaborating Institutions in the DES.

The Collaborating Institutions are Argonne National Labora-
tory, the University of California at Santa Cruz, the University of
Cambridge, Centro de Investigaciones Energéticas, Medioambien-
tales y Tecnológicas-Madrid, the University of Chicago, Univer-
sity College London, the DES-Brazil Consortium, the University
of Edinburgh, the Eidgenössische Technische Hochschule (ETH)
Zürich, Fermi National Accelerator Laboratory, the University of
Illinois at Urbana-Champaign, the Institut de Ciències de l'Espai
(IEEC/CSIC), the Institut de Física d'Altes Energies, Lawrence
Berkeley National Laboratory, the Ludwig-Maximilians Universität

<footnote n="7">⁷ http://archive.stsci.edu/cgi-bin/dss_form</footnote>

<footnote n="8">⁸ http://irsa.ipac.caltech.edu/Missions/iras.html</footnote>
München and the associated Excellence Cluster Universe, the Uni-
versity of Michigan, the NOAO, the University of Nottingham, The
Ohio State University, the University of Pennsylvania, the Univer-
sity of Portsmouth, SLAC National Accelerator Laboratory, Stan-
ford University, the University of Sussex, Texas A&M University
and the OzDES Membership Consortium.

The DES data management system is supported by the Na-
tional Science Foundation under Grant Number AST-1138766.
The DES participants from Spanish institutions are partially sup-
ported by MINECO under grants AYA2012-39559, ESP2013-
48274, FPA2013-47986, and Centro de Excelencia Severo Ochoa
SEV-2012-0234. Research leading to these results has received
funding from the European Research Council under the European
Union's Seventh Framework Programme (FP7/2007-2013) includ-
ing ERC grant agreements 240672, 291329 and 306478.

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<footnote n="1">¹Instituto de Física, Universidade Federal do Rio Grande do Sul, 91501-900 Porto Alegre, RS, Brazil</footnote>

<footnote n="2">²Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, 20921-400 Rio de Janeiro, RJ, Brazil</footnote>

<footnote n="3">³Fermi National Accelerator Laboratory, P.O. Box 500, Batavia, IL 60510, USA</footnote>

<footnote n="4">⁴Large Synoptic Survey Telescope, 933 North Cherry Avenue, Tucson, AZ 85721, USA</footnote>

<footnote n="5">⁵Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA</footnote>

<footnote n="6">⁶Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721, USA</footnote>

<footnote n="7">⁷Observatoire astronomique de Strasbourg, Université de Strasbourg, CNRS, UMR 7550, 11 rue de l'Université, F-67000 Strasbourg, France</footnote>
This paper has been typeset from a *TEX/LATEX* file prepared by the author.

<footnote n="8">⁸Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany</footnote>

<footnote n="9">⁹Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK</footnote>

<footnote n="10">¹⁰Instituto de Astrofísica de Canarias. Vía Láctea s/n. E-38200 - La Laguna, Tenerife, Canary Islands, Spain</footnote>

<footnote n="11">¹¹Department of Astrophysics, University of La Laguna. Vía Láctea s/n. E-38200 - La Laguna, Tenerife, Canary Islands, Spain</footnote>

<footnote n="12">¹²Astronomisches Rechen-Institut, Zentrum für Astronomie der Universität Heidelberg, Mönchhofstr. 12-14, D-69120 Heidelberg, Germany</footnote>

<footnote n="13">¹³George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and Astronomy, and Department of Physics and Astronomy, Texas A&M University, College Station, TX 77843, USA</footnote>

<footnote n="14">¹⁴Department of Physics, University of Surrey, Guildford GU2 7XH, UK</footnote>

<footnote n="15">¹⁵Department of Astronomy, University of Virginia, Charlottesville, VA 22904-4325, USA</footnote>

<footnote n="16">¹⁶University of Hertfordshire, Physics Astronomy and Mathematics, College Lane, Hatfield AL10 9AB, UK</footnote>

<footnote n="17">¹⁷Leibnitz-Institut für Astrophysik Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany</footnote>

<footnote n="18">¹⁸Department of Physics, ETH Zurich, Wolfgang-PauliStrasse 16, CH-8093 Zurich, Switzerland</footnote>

<footnote n="19">¹⁹Research School of Astronomy & Astrophysics, Mount Stromlo Observatory, Cotter Road, Weston Creek, ACT 2611, Australia</footnote>

<footnote n="20">²⁰Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatory, Casilla 603, La Serena, Chile</footnote>

<footnote n="21">²¹Center for Astrophysics and Space Astronomy, Department of Astrophysical and Planetary Sciences, University of Colorado, 389 UCB, Boulder, CO 80309, USA</footnote>

<footnote n="22">²²National Optical Astronomy Observatory, 950 N. Cherry Ave, Tucson, AZ 85719, USA</footnote>

<footnote n="23">²³Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro 20921-400, RJ, Brazil</footnote>

<footnote n="24">²⁴Department of Physics and Electronics, Rhodes University, P.O. Box 94, Grahamstown 6140, South Africa</footnote>

<footnote n="25">²⁵Department of Physics & Astronomy, University College London, Gower Street, London WC1E 6BT, UK</footnote>

<footnote n="26">²⁶CNRS, UMR 7095, Institut d'Astrophysique de Paris, F-75014 Paris, France</footnote>

<footnote n="27">²⁷Sorbonne Universités, UPMC Univ Paris 06, UMR 7095, Institut d'Astrophysique de Paris, F-75014 Paris, France</footnote>

<footnote n="28">²⁸Department of Astronomy, University of Illinois, 1002 W. Green Street, Urbana, IL 61801, USA</footnote>

<footnote n="29">²⁹National Center for Supercomputing Applications, 1205 West Clark St, Urbana, IL 61801, USA</footnote>

<footnote n="30">³⁰Institut de Ciències de l'Espai, IEEC-CSIC, Campus UAB, Carrer de Can Magrans, s/n, E-08193 Bellaterra, Barcelona, Spain</footnote>

<footnote n="31">³¹Instituto de Física Teorica UAM/CSIC, Universidad Autonoma de Madrid, E-28049 Madrid, Spain</footnote>

<footnote n="32">³²Kavli Institute for Particle Astrophysics & Cosmology, P.O. Box 2450, Stanford University, Stanford, CA 94305, USA</footnote>

<footnote n="33">³³Institute of Cosmology & Gravitation, University of Portsmouth, Portsmouth PO1 3FX, UK</footnote>

<footnote n="34">³⁴School of Physics and Astronomy, University of Southampton, Southampton SO17 1BJ, UK</footnote>

<footnote n="35">³⁵Department of Physics, IIT Hyderabad, Kandi, Telangana 502285, India</footnote>

<footnote n="36">³⁶SLAC National Accelerator Laboratory, Menlo Park, CA 94025, USA</footnote>

<footnote n="37">³⁷Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus, OH 43210, USA</footnote>

<footnote n="38">³⁸Department of Physics, The Ohio State University, Columbus, OH 43210, USA</footnote>

<footnote n="39">³⁹Astronomy Department, University of Washington, Box 351580, Seattle, WA 98195, USA</footnote>

<footnote n="40">⁴⁰Australian Astronomical Observatory, North Ryde, NSW 2113, Australia</footnote>

<footnote n="41">⁴¹Institució Catalana de Recerca i Estudis Avançats, E-08010 Barcelona, Spain</footnote>

<footnote n="42">⁴²Institut de Física d'Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, E-08193 Bellaterra (Barcelona), Spain</footnote>

<footnote n="43">⁴³Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Dr, Pasadena, CA 91109, USA</footnote>

<footnote n="44">⁴⁴Department of Physics and Astronomy, Pevensey Building, University of Sussex, Brighton BN1 9QH, UK</footnote>

<footnote n="45">⁴⁵Department of Physics and Astronomy, University of Pennsylvania, Philadelphia, PA 19104, USA</footnote>

<footnote n="46">⁴⁶Centro de Investigaciones Energéticas, Medioambientales y Tecnológicas (CIEMAT), Avda. Complutense, 40 E-28040, Madrid, Spain</footnote>

<footnote n="47">⁴⁷Department of Physics, University of Michigan, Ann Arbor, MI 48109, USA</footnote>

<footnote n="48">⁴⁸Universidade Federal do ABC, Centro de Ciências Naturais e Humanas, Av. dos Estados, 5001, Santo André 09210-580, SP, Brazil</footnote>

<footnote n="49">⁴⁹Computer Science and Mathematics Division, Oak Ridge National Laboratory, Oak Ridge, TN 37831, USA</footnote>