arXiv:0803.4218v1 [astro-ph] 28 Mar 2008 Stochastic ‘Beads on a String’ in the Accretion Tail of Arp 285 Beverly J. Smith Department of Physics, Astronomy, and Geology, East Tennessee State University, Johnson City TN 37614 [email protected]Curtis Struck Department of Physics and Astronomy, Iowa State University, Ames IA 50011 [email protected]Mark Hancock Department of Physics, Astronomy, and Geology, East Tennessee State University, Johnson City TN 37614 [email protected]Mark L. Giroux Department of Physics, Astronomy, and Geology, East Tennessee State University, Johnson City TN 37614 [email protected]Philip N. Appleton NASA Herschel Science Center, California Institute of Technology, Pasadena CA 91125 [email protected]Vassilis Charmandaris 1 Department of Physics, University of Crete, Heraklion Greece 71003 [email protected]William Reach Spitzer Science Center, California Institute of Technology, Pasadena CA 91125 [email protected]Sabrina Hurlock Department of Physics, Astronomy, and Geology, East Tennessee State University, Johnson City TN 37614 [email protected]and Jeong-Sun Hwang Department of Physics and Astronomy, Iowa State University, Ames IA 50011 1
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Stochastic ‘Beads on a String’ in the Accretion Tail of Arp 285
Beverly J. Smith
Department of Physics, Astronomy, and Geology, East Tennessee State University, Johnson City TN 37614
We present Spitzer infrared, GALEX UV, and SDSS and SARA optical images of the peculiarinteracting galaxy pair Arp 285 (NGC 2856/4), and compare with a new numerical model of theinteraction. We estimate the ages of clumps of star formation in these galaxies using populationsynthesis models, carefully considering the uncertainties on these ages. This system contains astriking example of ‘beads on a string’: a series of star formation complexes ∼1 kpc apart. These‘beads’ are found in a tail-like feature that is perpendicular to the disk of NGC 2856, which impliesthat it was formed from material accreted from the companion NGC 2854. The extreme bluenessof the optical/UV colors and redness of the mid-infrared colors implies very young stellar ages(∼ 4 − 20 Myrs) for these star forming regions. Spectral decomposition of these ‘beads’ showsexcess emission above the modeled stellar continuum in the 3.6 µm and 4.5 µm bands, indicatingeither contributions from interstellar matter to these fluxes or a second older stellar population.These clumps have −12.0 < MB < −10.6, thus they are less luminous than most dwarf galaxies.Our model suggests that bridge material falling into the potential of the companion overshootsthe companion. The gas then piles up at apo-galacticon before falling back onto the companion,and star formation occurs in the pile-up. There was a time delay of ∼500 Myrs between thepoint of closest approach between the two galaxies and the initiation of star formation in thisfeature. A luminous (MB ∼ −13.6) extended (FWHM ∼ 1.3 kpc) ‘bright spot’ is visible at thenorthwestern edge of the NGC 2856 disk, with an intermediate stellar population (400 − 1500Myrs). Our model suggests that this feature is part of a expanding ripple-like ‘arc’ created byan off-center ring-galaxy-like collision between the two disks.
Galaxy evolution is strongly driven by in-teractions and mergers between galaxies. In-teractions can produce tidal tails and bridges(Toomre & Toomre 1972), increase star formationrates (Kennicutt et al. 1987; Bushouse, Lamb, & Werner1988), and trigger the formation of young super-star clusters (Holtzman et al. 1992, 1996). Tidalmaterial can contribute to the intergalactic medium(Morris & van den Bergh 1994) and to intergalac-tic starlight (Feldmeier et al. 2002). Gas-richgalaxy mergers can produce ultra-luminous in-frared galaxies (Soifer et al. 1987; Smith et al.1987; Sanders et al. 1988), while concentrationsof stars and gas in tidal features may becomeindependent dwarf galaxies (Barnes & Hernquist1992; Elmegreen, Kaufman, & Thomasson 1993).
1IESL/Foundation for Research and Technology - Hel-las, GR-71110, Heraklion, Greece and Chercheur Associe,Observatoire de Paris, F-75014, Paris, France
The key to understanding these processes iscareful comparison of multi-wavelength observa-tions of nearby galaxies with dynamical models.Since interactions and mergers are even more com-mon at high redshift than in the local Universe(e.g., Abraham & van den Bergh 2001), detailedstudies of nearby interacting systems are impor-tant for interpreting high redshift surveys. Suchstudies can provide information on the timescaleof the interaction, the history of gas compres-sion in different regions, star formation trig-gering, dissipation in the gas, multiple burstsof star formation, and mass transfer betweengalaxies (Struck & Smith 2003; Struck et al. 2005;Smith et al. 2005b; Hancock et al. 2007). Com-puter simulations can provide predictions of thedistribution of star formation, which can be com-pared to observational results to estimate the ef-fects of compression strength, duration, and otherfactors (e.g., Struck & Smith 2003).
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To study star formation enhancement in pre-merger interacting systems, we obtained mid-infrared observations with the Spitzer telescope(Werner et al. 2004) for three dozen interact-ing galaxies selected from the Arp (1966) Atlasof Peculiar Galaxies (the ‘Spirals, Bridges, andTails’ (SB&T) sample; Smith et al. 2007). Wehave completed detailed multi-wavelength stud-ies of three of the galaxy pairs in the SB&Tsample, and have constructed matching hydro-dynamical models of their encounters: Arp 284(Smith, Struck, & Pogge 1997; Struck & Smith2003; Smith et al. 2005a), Arp 107 (Smith et al.2005b), and Arp 82 (Hancock et al. 2007). A sim-ilar study of the interacting pair IC 2163/NGC2207 was presented by Struck et al. (2005) andElmegreen et al. (2006), while Arp 24 was studiedby Chen & Wu (2007).
In the current paper, we describe a multi-wavelength study of another of the SB&T sys-tems, the interacting galaxy pair Arp 285 (NGC2856/4), and compare with a new numerical modelof the interaction. The more northern galaxy inthis widely separated pair, NGC 2856, has a pe-culiar tail-like feature extending out perpendic-ular to the disk (Figure 1). Toomre & Toomre(1972) suggested that this feature is material fromthe southern galaxy NGC 2854, which has ac-creted onto NGC 2856 via the bridge. The pres-ence of a massive HI counterpart to this tailand the HI velocity field support this hypothesis(Chengalur, Salpeter, & Terzian 1994, 1995). TheSpitzer 3.6 µm − 8.0 µm broadband infrared colorof the NGC 2856 tail is the reddest of all the tidalfeatures in the SB&T sample (Smith et al. 2007),implying a very young stellar population.
In the current study, we investigate star for-mation in Arp 285 by combining our Spitzermid-infrared images with ultraviolet images fromthe Galaxy Evolution Explorer (GALEX) mis-sion (Martin et al. 2005) and optical imagesfrom the Sloan Digitized Sky Survey (SDSS)(Abazajian et al. 2003) and the Southeastern As-sociation for Research in Astronomy (SARA) tele-scope2. We also compare with the 2MASS At-las near-infrared images of Arp 285 (Cutri et al.2006). Arp 285 is relatively nearby, at a distanceof 39 Mpc (H0 = 75 km s−1Mpc−1).
2http://astro.fit.edu/sara/sara.html
2. Observations
The Spitzer infrared observations and data re-ductions are described in detail in Smith et al.(2007). The data used includes broadband 3.6 µm,4.5 µm, 5.8 µm and 8.0 µm images from the In-frared Array Camera (IRAC; Fazio et al. 2004),with spatial resolutions of 1.′′5 − 2.′′0, a pixel sizeof 1.′′2, and a field of view of 7.′8 × 12.′6. A24 µm image of Arp 285 was also obtained withthe Multiband Imaging Photometry for Spitzer(MIPS; Rieke et al. 2004), however, it has pro-nounced artifacts from the point spread function(see image in Smith et al. 2007). Because of theseartifacts, this image is only useful for determin-ing total galaxian fluxes, not fluxes for individualclumps or tidal features. Thus it is not used inthis analysis.
Arp 285 was observed as part of the Sloan Dig-italized Sky Survey (SDSS) in the ugriz optical fil-ters (effective wavelengths 3560A, 4680A, 6180A,7500A, and 8870A, respectively). These imageshave a pixel size of 0.′′40 and a field of view of 13.′5× 9.′8. The two galaxies in the pair are in twodifferent SDSS fields of view. The FWHM pointspread function is ∼1.′′2, based on stars in the field.
Arp 285 was also observed with the SARA 0.9moptical telescope on 2006 Jan 29, in partly cloudyweather. An 1152×770 Apogee Alta CCD with apixel size of 0.′′64 pixel−1 was used, giving a fieldof view of 12.′3 × 8.′2. A total of three 600 sec-ond exposures were made in a broadband R filter,along with seven 600 second images in a redshiftedHα filter centered at 664 nm with a FWHM of 7nm. For Arp 285, this filter contains both Hαand the [N II] λλ6548,6583 line. The SARA datawere reduced in the standard way using the Im-age Reduction and Analysis Facility (IRAF3) soft-ware. Continuum subtraction was accomplishedusing the scaled R band image.
Arp 285 was observed in a near-ultraviolet(NUV) broadband filter (1750 − 2800 A) byGALEX as part of the GALEX Medium Imag-ing Survey (MIS) (Martin et al. 2005). The MISimage had a total integration time of 813 sec. Arp285 was also observed in the far-ultraviolet (FUV)
3IRAF is distributed by the National Optical AstronomyObservatories, which are operated by the Association ofUniversities for Research in Astronomy, Inc., under coop-erative agreement with the National Science Foundation.
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(1350 − 1705 A) as part of the GALEX All-SkySurvey (Martin et al. 2005), with a shorter expo-sure time of 112 seconds. The GALEX spatialresolution is ∼5′′, with a pixel size of 1.′′5. Thefield of view is circular, with a 1.2◦ diameter.
The total magnitudes for NGC 2854 and NGC2856 in the various filters are given in Table 1.
3. The Morphology of Arp 285
3.1. NGC 2856
In Figure 2, we present a montage of the UV,optical, and infrared images of NGC 2856, thenorthern galaxy in the Arp 285 pair. In the opti-cal images, a dusty spiral pattern and a centralbar-like feature are seen in the disk. A seriesof four clumps are visible along the northern tailin all of the optical images, except for the u im-age (only two clumps detected) and the z image(only one clump detected). These clumps are la-beled on the g image in the last panel in Figure 2.From south to north, the separations between theclumps are 7.′′6 (1.4 kpc), 6.′′9 (1.3 kpc), and 5.′′1(1.0 kpc). Clumps 1 and 2 have bright unresolvedor marginally-resolved cores (FWHM < 1.′′3 = 250pc) in the g image; clumps 3 and 4 are fainter ing, with multiple peaks.
The Spitzer 5.8 µm and 8.0 µm images are af-fected by ‘banding’, where bright point sources(such as galactic nuclei) cause horizontal ‘bands’in the images (Spitzer Infrared Array CameraData Manual, Version 2.0, 2005). As discussedin Smith et al. (2007), we corrected for these arti-facts by interpolating from nearby clean regions.Unfortunately, the correction was not perfect forNGC 2856, leaving a residual diagonal ‘stripe’across the rotated image near the tail (see Figure2). In spite of this, however, clump 3 is detectedin all four Spitzer bands, and clump 2 is detectedat 3.6 µm and 4.5 µm. Clump 1 is not detected inany of the Spitzer filters, while clump 4 has onlya marginal detection at 8.0 µm. Note that clump2 is brightest in the SDSS data, but clump 3 isbrightest at 8 µm.
The northern tail is also visible in both theFUV and NUV images. In the longer exposureNUV image, clumps 1 − 3 are bright, while clump4 is marginally detected. In the short exposureFUV image, with the low resolution spatial res-olution of GALEX the individual clumps are not
well-resolved. The northern tail is not detected inthe 2MASS near-infrared images.
In the SARA Hα map (Figure 3), only clump 2is detected in the northern tail. This implies thatclump 3, the brightest clump at 8 µm, is moreextincted. This is consistent with the ∼12′′ resolu-tion C Array HI map of Chengalur, Salpeter, & Terzian(1994). In this map, two HI peaks are clearly vis-ible in the northern tail. The brightest HI peakis approximately coincident with clump 1, whilethe second is near clump 3. Thus clump 2 maybe less extincted than these other clumps. This isconsistent with our analysis of the optical colors(see Section 4.2).
In the u through 4.5 µm images of NGC 2856, a‘bright spot’ is visible in the northwestern edge ofthe disk of NGC 2856 (see Figure 2). This ‘brightspot’ is also visible in the Arp image (Figure 1),and is marginally detected in the 2MASS H andKs images. However, it is not seen as a discretesource at 5.8 µm, 8.0 µm, or in the FUV, NUV,Hα, or 2MASS J images (Figures 2 and 3). Unlikethe clumps in the northern tail, this ‘bright spot’is smoothly extended in the SDSS images, with aFWHM ∼ 7′′ (1.3 kpc) in the g filter, without acompact core or cores.
Within the inner disk of NGC 2856, bright 8µm and Hα sources are visible at the ends of thebar and the nucleus (Figures 2 and 3). The bar isasymmetric, with the clump near the southern endof the bar being brighter than the northern sourcein both 8 µm and Hα. In the higher resolutionSDSS images, the sources at the ends of the barare resolved into 2 − 4 peaks separated by 2′′ −3′′ (0.4 − 0.6 kpc).
In Figure 4, a band-merged approximately true-color optical SDSS image of NGC 2856 is pre-sented. This shows that the clumps in the north-ern tail are bluer than the main disk of the galaxy.The dust features and the spiral pattern are alsovisible in this picture. The northeastern spiral armis bluer than the southwestern portion of the disk.This is also apparent in Figure 2. In the FUV,NUV, and u images, the northeastern portion ofthe disk is brighter than the southwestern section,but in the longer wavelength images, the disk ismore symmetric. This suggests that the differ-ence at shorter wavelengths is due to extinction.This implies that the northeastern side of NGC2856 is closest to us. This is consistent with the
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sense of rotation indicated by the HI velocity field(Chengalur, Salpeter, & Terzian 1994), assumingthe northwestern spiral arm is trailing.
A connecting bridge between the two galaxies isvisible in the smoothed g and r SDSS images (seeFigure 5), but is not seen in u, i, or z. This bridgeis aligned with the ‘bright spot’ in the disk. InFigure 5, the northern tail is visible out to ≈72′′
(14 kpc) from the disk. A bend and a sudden drop-off in brightness is evident in this tail just north ofthe four bright clumps of star formation. Anotherpossible faint clump is visible in the smoothed gimage ∼7.9′′ (1.5 kpc) northwest of clump 4, northof the bend. This bend and the bridge are alsovisible in the Arp image (Figure 1).
3.2. NGC 2854
In Figure 6, a montage of the UV, optical, andinfrared images of the southern galaxy NGC 2854is shown. On-going star formation is detectedalong the spiral arms and at the ends of the bar.The base of the northern tail/bridge appears dou-ble in the u, g, r, 5.8 µm, and 8.0 µm images.A series of clumps are visible in the spiral armsin both the optical and the infrared images, and8 µm-bright sources are seen at the ends of thebar. For some clumps, there are 1′′ − 2′′ offsetsbetween the optical and 8 µm peaks; for others,including the nuclear source, there is no clear op-tical peak associated with the 8 µm source. In thelast panel of Figure 6, we identify eight clumpsselected based on the 8 µm image.
The SARA Hα and R maps of NGC 2854 arepresented in Figure 7, with the Hα superimposedon the g and 8 µm images. All of the 8 µm clumpsexcept clump 6 were detected in Hα. In addition,possible Hα emission is seen associated with thewestern portion of the double bridge.
An approximately true-color optical SDSS im-age of NGC 2854 is displayed in Figure 8. Thedouble bridge structure is visible in this image.The knots along the northern arm are visible. Thesouthern end of the bar is bluer than the northernend, and the southern arm/tail is bluer than thenorthern arm. The southwestern end of the bar isparticularly bright in the UV, but less so in themid-infrared (Figure 6). NGC 2854 appears moresymmetric in the Spitzer images than at shorterwavelengths. This implies that the color varia-
tions seen in the optical/UV are due to extinction.The color variations suggest that the southern sideof NGC 2854 is the near side, consistent with theHI velocity field of Chengalur, Salpeter, & Terzian(1994) and trailing spiral arms.
The smoothed g image of NGC 2854 is pre-sented in Figure 9. A faint optical tail is detected,extending 1.′8 (20 kpc) to the south, coincidentwith the long HI tail seen by Chengalur, Salpeter, & Terzian(1994). This tail is also visible in the smoothedFUV and NUV images.
Approximately 3′ due north of NGC 2854, at9h 24m 2.9s, 49◦ 14′ 41′′ (J2000), a small angu-lar size galaxy is visible on the Arp image (Figure1). This galaxy is detected in all of the GALEXand SDSS bands, as well as the Spitzer 3.6 µmand 4.5 µm filters. At the present time, no red-shift is available for this source, so it is unknownwhether it is associated with Arp 285. Its prox-imity to a very bright star (see Figure 1) preventsa reliable Hα detection. The magnitudes of thisgalaxy in the various filters are given in Table 1.The UV/optical colors are quite blue, consistentwith a young stellar population.
4. Clump Analysis
4.1. Photometry
The positions, SDSS, and GALEX magnitudesand Spitzer flux densities of the clumps in theNGC 2856 tail and the two disks of Arp 285 aregiven in Table 2 in R.A. order. The photometrywas done using the IRAF daophot routine. Forthe tail and the ‘bright spot’, the positions weredetermined by eye based on the g image, whilethe disk positions are the 8 µm peaks. These po-sitions are marked in the last panels of Figures2 and 6. The apertures we utilized are given inTable 3, along with the aperture corrections usedto correct to total magnitudes. Since the clumpsin the northern tail are separated by only ∼5′′ −7′′, we used relatively small apertures for the tailclump photometry. For the disk clumps, we usedlarger apertures because of less crowding, and be-cause there are sometimes multiple optical peaksassociated with a single Spitzer source.
For background subtraction, the local galax-ian background was determined using the modein an annulus surrounding the source (see Table3). To estimate the uncertainty in the colors of
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the clumps due to background subtraction, in ad-dition to the statistical uncertainties determinedfrom the rms in the background annuli, in calcu-lating the colors we added in quadrature a seconduncertainty term, determined from comparing theclump colors obtained with the above method withthose obtained with slightly larger annulus.
We also extracted approximate J, H, and Ks
photometry for the clumps from the 2MASS Atlasimages. These near-infrared fluxes were not usedin the population synthesis modeling (Section 4.2),but were only used for comparison with the Spitzerdata (Section 4.3), thus we did not include thesecond sky annulus.
We also obtained the total FUV and NUV fluxfor a 25.′′8 × 9.′′7 region containing the four knots.These magnitudes are also given in Table 2, alongwith the SDSS and Spitzer magnitudes for thesame region. The uncertainties given in Table 2for this region were calculated as in Smith et al.(2007), including both statistical uncertainties andthe uncertainty in the sky level. The total g fluxin this rectangular region is ∼3× the sum of the gfluxes for the four clumps (see Table 2), thus thereis significant diffuse emission in this tail.
4.2. Ages
In Figure 10, we plot the SDSS g − r colorsfor both the tail and disk clumps vs. their u − gcolors. In Figure 11, NUV − g is plotted againstg − r. Similar plots of r − i vs. g − r, i − z vs.r − i, and FUV − NUV vs. g − r are presentedin the Appendix. These Figures also include thecolors for the 25.′′8 × 9.′′7 rectangular region thatincludes the four clumps in the northern tail.
To estimate the ages of these clumps, we cal-culated theoretical colors for star forming regionsusing version 5.1 of the Starburst99 populationsynthesis code (Leitherer et al. 1999). This ver-sion includes the Padova asymptotic giant branchstellar models (Vazquez & Leitherer 2005). Thesemodels assume an instantaneous burst with aKroupa (2002) initial mass function (IMF) and aninitial mass range of 0.1 − 100 M⊙. We calculatedcolors for a range of ages τ from 1 Myr to 10 Gyrs.We used a time step size of ∆τ = 1 Myrs for 0 <τ < 1 Gyr, and ∆τ = 100 Myrs for 1.1 < τ < 10Gyr. The Calzetti, Kinney, & Storchi-Bergmann(1994) starburst dust reddening law was assumed.
We also generated models using a Salpeter IMF,and found the colors differed only slightly fromthose with the Kroupa IMF. This is consistentwith earlier studies (MacArthur et al. 2004). Tothe broadband fluxes, we added in the contribu-tions from Hα emission, which can be substantialin the r filter. For a 1 Myr star forming region,Hα decreases the r magnitude by 1.1 magnitude;at 5 Myrs, Hα contributes ∼0.25 magnitudes (seeFigure 10).
To systematically estimate ages and extinctionsfor the clumps in Table 2, we used a χ2 minimiza-tion calculation (e.g., Pasquali, de Grijs, & Gallagher2003) to determine the fit of the observed colorsto that of the models:
χ2 =
N∑
i=1
(
obsi − modeliσi
)2
In this equation, N is the number of colors used inthe analysis, obsi is the observed color, modeli isthe corresponding model color, and σi is the un-certainty in the obsi color. A good fit is indicatedby χ2 < N. In these calculations, we did not in-clude filters with non-detections. In a few cases,it was not possible to find a good fit when the lowS/N FUV or z fluxes were included. In these cases,they were not used in determining the ages.
To estimate the uncertainties in the best-fit pa-rameters, we used the ∆χ2 method (Press et al.1992) to determine 68.3% confidence levels for theparameters. The best-fit parameters and their un-certainties are given in Table 4, along with thecolors used in the fits. We assumed solar metal-licity in these models. In the Appendix, we giveresults assuming 1/5 solar metallicity. The de-rived age ranges for the two metallicities are sim-ilar, thus the assumed metallicity has little effecton the derived ages. We note that, because theclump masses are relatively low (see Section 4.4),stochastic sampling of the IMF can affect the agedeterminations (e.g., Cervino et al. 2002). We donot include this effect in our calculations.
On the color-color plots of the clump colors, wesuperimpose the solar metallicity models. Usingonly optical colors, it is often difficult to distin-guish between reddening due to age and redden-ing due to extinction (see Figure 10). Fortunately,however, some clumps were detected in the UV,which in some cases can constrain the ages fur-
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ther (see Table 4). The clumps in the tail havevery blue optical/UV colors, thus they have bothlow extinctions and young ages, with E(B − V) ∼0.1 and ages ∼ 4 − 20 Myrs (Table 4). Clump 1is slightly redder than the other clumps, implyinga slightly older age and/or a higher extinction.
As noted previously, in the HI map of Chengalur, Salpeter, & Terzian(1994), two peaks are visible in the tail, nearclumps 1 and 3. The brightest peak has an HIcolumn density ≈ 2 × 1021 atoms cm−2, while thesecond has N(HI) ≈ 1021 atoms cm−2. Assumingthe standard Galactic N(H)-to-extinction ratio ofN(H)/E(B−V) = 5.8 × 1021 atoms cm−2 mag−1
(Bohlin, Savage, & Drake 1978) and neglectingpossible molecular gas, these imply E(B − V) ≈
0.3 to clump 1 and E(B − V) ≈ 0.15 to clump 3.These are consistent with the population synthesisresults (Table 4).
In the optical colors (Figure 10), the nuclei ofthe two galaxies (cyan open diamond 3 and openblack circle 7) are quite red, meaning high extinc-tions and/or age. Since they were undetected withGALEX, we are not able to tightly constrain theirages (see Table 4). In the case of the ‘bright spot’in the NGC 2856 disk, the optical colors can onlyconstrain the age to between 5 − 1500 Myrs. Thelack of a detection in the NUV, however, furtherconstrains the age to be >
∼400 Myrs (see Figure11).
We also attempted to model the ages of thediffuse emission in the two galaxian disks and thenorthern tail by subtracting the clump light fromthe total emission. We modeled the star formationhistory of the diffuse emission in two ways: as aninstantaneous burst and as continuous star forma-tion. With the exception of ruling out extremelyyoung ages (1 − 5 Myrs), we cannot strongly con-strain the age of the diffuse emission. For the in-stantaneous burst models, we get upper limits tothe ages of the diffuse emission in the NGC 2854and NGC 2856 disks of 1.5 Gyrs and 400 Myrs,respectively. This implies that there are some re-cently formed stars in the disks outside of the re-gions we defined as clumps. This does not, how-ever, rule out an additional underlying older stellarpopulation in the disks.
4.3. Spitzer Colors
Figure 12 shows the Spitzer [3.6 µm] − [4.5 µm]vs. [4.5 µm] − [5.8 µm] colors for the clumpsin the NGC 2856 tail (magenta diamonds), theNGC 2856 disk (cyan open diamonds), and theNGC 2854 disk (black open circles). A similarplot for [4.5] − [5.8] vs. [5.8] − [8.0] is providedin the Appendix. These colors are compared tothe colors of the Arp 107 and Arp 82 clumps(green diamonds, from Smith et al. 2005b andHancock et al. 2007), Galactic interstellar dust(blue X’s; from Flagey et al. 2006), M0III stars(open blue square, from Cohen 2005, private com-munication), and field stars (magenta open trian-gle, from Whitney et al. 2004), as well as the totalcolors for the NGC 2856 tail region (red filled tri-angle).
As show in Figure 12, [3.6] − [4.5] ≈ 0.0 forboth stars and Galactic dust. Global values forboth interacting and spiral galaxy disks are alsoclose to this value (Smith et al. 2007), as are theclumps in the disks of Arp 285, 107, and 82 (Figure12). The [5.8] − [8.0] colors of most of the clumpsin the three Arp systems are similar to those ofinterstellar matter and redder than stars, as ex-pected since these bands are likely dominated byinterstellar dust emission. The [4.5] − [5.8] col-ors of the Arp clumps are mainly between those ofstars and interstellar matter (Figure 12), suggest-ing contributions from both. The very red [4.5]− [5.8] color of clump 3 in the Arp 285 tail com-pared to the clumps in Arp 107 and Arp 82 andthe other clumps in Arp 285 (Figure 12) impliesmore contributions from interstellar matter. Thisis consistent with the very young age determinedfrom the optical colors (Table 4).
To disentangle the contributions from starlightand dust to the Spitzer bands, the results of ourstellar population synthesis (Section 5) are helpful.In Figure 13, we plot the full optical-mid-infraredspectral energy distribution (SED) for tail clump3. We superimpose on this plot our best-fit Star-burst99 model (4 Myrs, E(B−V) = 0.1), alongwith models that span our 1σ uncertainty (68%confidence) in the age (3 − 6 Myrs). In addi-tion, we include a theoretical dust spectrum fromDraine & Li (2007). This dust spectrum showsthe broad polycyclic aromatic hydrocarbon (PAH)mid-infrared emission features, as well as a ‘hot
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dust’ continuum. The plotted dust spectrum wascalculated with the solar neighborhood interstellarradiation field, scaled up by a factor of U = 100.It uses a PAH-to-total-dust mass ratio of qPAH =4.6%. The dust SED in our wavelength range doesnot vary much with U; lowering qPAH weakens thePAH features (Draine & Li 2007). The dust modelhas been scaled to fit the observed 8 µm flux. Thesolid line in this plot is the combined stellar-dustspectrum. The Hα contribution to the SDSS rband is clearly visible in this plot. A contributionto the broadband 3.6 µm Spitzer flux from the 3.3µm PAH feature is also apparent.
Figure 13 indicates that, in addition to the PAHcontribution to the 3.6 µm band and the modeledstellar component, the 3.6 µm and 4.5 µm fluxes oftail clump 3 also include another component. Ei-ther there is a significant ‘hot dust’ contributionto these bands, as indicated by the Draine & Li(2007) model, or there is a second underlying olderstellar component that is not revealed by the op-tical/UV population synthesis.
For comparison to tail clump 3, in Figure 14 weplot the SED of clump 2 in the NGC 2856 disk.Although this also has a very young stellar popu-lation (see Table 4), it has a much more reddenedSED than the tail clump because of much higherextinction. Also, although PAH emission is clearlypresent at 8 µm, in the 3.6 and 4.5 µm bands mostof the emission is stellar, in contrast to clump 3 inthe tail.
The ‘bright spot’ in the NGC 2856 disk (cyanclump 1) is undetected at 5.8 µm and 8.0 µm, witha very blue [4.5] − [5.8] upper limit compared tothe other clumps (Figure 12). This suggests thatthis region has an older stellar population than theother disk clumps. As noted earlier, with the avail-able optical data we could not strongly constrainthe age of this clump (Table 4), however, the lackof an NUV detection points to an older age (seeSection 4.2). The Spitzer results are consistentwith this conclusion. This shows that Spitzer mid-infrared data may be useful for breaking the age-extinction degeneracy in optical colors. In Figure15, we plot the SED for the ‘bright spot’, with thebest fit from the optical data shown. The NUVlimit plotted shows the additional constraint onthe age. The SED plot shows that the 3.6 µm and4.5 µm emission is dominated by starlight, withvery little if any dust contributing.
The two Arp 285 nuclei have [4.5] − [5.8] and[5.8] − [8.0] colors similar to the other clumps,implying nuclear starbursts. This is in contrastto the Arp 107 nuclei (Figure 12), which haveolder stellar populations (Smith et al. 2005b).The two nuclei in Arp 82, like those in Arp285, have Spitzer colors of star forming regions(Hancock et al. 2007).
4.4. Absolute Magnitudes and Masses
In Table 5, we compare the absolute opticalmagnitudes of the NGC 2856 tail clumps withdwarf galaxies, candidate tidal dwarf galaxies(TDGs), ‘super star clusters’ (SSCs), and thetail clumps in Arp 82. The Arp 285 tail clumpsare lower luminosity than most nearby irregu-lar galaxies and tidal dwarf galaxies, and arenear the lower end of the range for SSCs. Thefaintest Arp 285 tail clump, clump 4, is somewhatless luminous than R136, the bright star clus-ter in 30 Doradus in the Large Magellanic Cloud(O’Connell, Gallagher, & Hunter 1994). In con-trast, the ‘bright spot’ in the NGC 2856 disk isnear the median for dwarf irregular galaxies.
For the Arp 285 clumps, in Table 6 we givethe range of stellar masses inferred from the Star-burst99 models. In this table, we also providestellar masses of various other objects for com-parison. The tail clumps are similar in mass toGalactic globular clusters, but have lower stellarmasses than those inferred for tidal dwarf galax-ies and dwarf irregular galaxies. The mass of theNGC 2856 disk ‘bright spot’ is near the medianfor dwarf irregular galaxies.
The 3.6 µm Spitzer band is sometimes used asa tracer of stellar mass (e.g., Li et al. 2007). How-ever, our SED plots (Figures 13 − 15) show thatthe stellar mass-to-3.6 µm luminosity varies signif-icantly from clump to clump, depending upon thestar formation rate and gas-to-star ratio. This isillustrated in Figure 16, where we plot the stellarmass of the clump determined from the popula-tion synthesis model against the 3.6 µm luminos-ity. We have also included values for the clumps inArp 82 (Hancock et al. 2007). On this curve, wehave superimposed lines of constant stellar mass-to-light ratios of M/L3.6 = 1 M⊙/L⊙ (solid line)and M/L3.6 = 10 M⊙/L⊙ (dotted line), whereL⊙ is the bolometric luminosity of the Sun. Thisplot shows that the NGC 2856 tail clumps and
8
four clumps in NGC 2854 (clumps 1, 3, 5, and 6,in the outer parts of the spiral arms) have lowerM/L3.6 ratios than the other clumps, which areclose to the M/L3.6 = 10 M⊙/L⊙ line. This indi-cates that contributions from hot dust and/or the3.3 µm PAH feature to the 3.6 µm flux are signifi-cant in the tail and outer spiral arm regions. Thuscaution should be used in utilizing the Spitzer 3.6µm band to estimate stellar masses in star form-ing regions. For example, for clump 3 in the tail,the stellar mass is ∼1/30th that expected basedon the M/L3.6 = 10 M⊙/L⊙ relationship.
5. A Numerical Model of the Encounter
To interpret these observational results in termsof the dynamical and star forming history of Arp285, we have constructed a numerical simulationof the Arp 285 interaction using the smoothedparticle hydrodynamics (SPH) code of Struck(1997). This code was previously used to modelArp 284 (Struck & Smith 2003), IC 2163/NGC2207 (Struck et al. 2005), Arp 107 (Smith et al.2005b), and Arp 82 (Hancock et al. 2007).
5.1. Constraints on the Model
Arp 285 is less symmetric than the ring galax-ies or planar fly-by encounters like M51, NGC2207/IC2163, and Arp 82. The collisional mor-phology of Arp 285 appears somewhat similar tothat of Arp 284, an asymmetric ring/tail galaxy(NGC 7714) with an edge-on companion (NGC7715). The substantial bridge and tail of NGC2854, like those of NGC 7715, lead us to be-lieve that it suffered a strong prograde encounter.There are also similarities between NGC 7714 andNGC 2856. The optical images show that the‘bright spot’ in the northwestern section of theNGC 2856 disk (disk clump 1) is part of an arc-like structure (see Figures 1 and 2). This arc isreminiscent of the partial ring in NGC 7714 (seethe Arp (1966) photograph of Arp 284), whichhas been successfully modeled by an off-centercollision (Struck & Smith 2003). It is also rem-iniscent of the ‘ripples’ in Arp 227, which werealso modeled by a ring galaxy-like collision byWallin & Struck-Marcell (1988).
There are some differences between NGC 2856and NGC 7714, however. In contrast to NGC7714, NGC 2856 lacks strong tidal tails, ex-
cept for the northern tail perpendicular to thedisk and a short HI extension to the northwest(Chengalur, Salpeter, & Terzian 1994). This sug-gests that NGC 2856 did not experience the en-counter as very prograde. It also does not havethe fan-like form common to strong retrograde en-counters. This suggests that the orbital path ofthe two galaxies is at a large angle to the plane ofthe NGC 2856 disk.
These considerations give us some idea of thetype of collision that produced the current mor-phologies. In our simulation of this encounter, wehave limited ourselves to the goal of reproducingthe large-scale morphological structures, but havenot attempted to simulate internal disk structuresnor match the system kinematics in any detail.
One key feature we would like the models tohelp us understand are the beads in the tail northof NGC 2856. We have considered several con-ceptual ideas for the origin of this material. TheHI morphology suggests that this material is anextension of the bridge from NGC 2854, thoughthe optical observations look as though the bridgecurves away from that direction before connectingto the bead region. It may be that the bridge isin fact a tidal tail, which is merely projected ontoNGC 2856, not connected. However, the HI kine-matics indicate that this is unlikely. Moreover,the bead material seems strongly affected by thegravitational potential of NGC 2856.
Thus, it seems likely that the bead material isaccreting onto the halo of NGC 2856 from thebridge. There two possibilities for how this oc-curs: i) as infall through the disk of NGC 2856and out the other side, or ii) by swinging aroundthat disk to the other side. It is difficult to distin-guish between these two scenarios observationally.In option i) we can imagine that clouds pushingthrough the NGC 2856 disk are shocked and com-pressed. This may trigger star cluster formation,accounting for the beads. We would naively expectthis process to be sequential, so that the beads fur-thest from the disk are oldest. In contrast, in op-tion ii), a group of inflowing clouds pile up in thehalo of NGC 2856 and collide with material thatarrived earlier. This could trigger star formationsimultaneously at several locations. Thus optioni) would predict an age gradient, while for optionii) we would expect roughly coeval clumps. Withthe available data, we cannot distinguish between
9
these two possibilities, since the expected age gra-dient for option i) is too small to measure. As-suming a nominal velocity for the tidal materialaway from the disk of ∼300 km s−1 and motionin the plane of the sky, for scenario i) we wouldexpect an age difference of ∼12 Myrs between thefirst and fourth clumps in the tail, and ∼4 Myrsbetween clumps 2 and 3. This is smaller than theuncertainties on the ages of these clumps (Table4).
Another way to distinguish between these twoscenarios is with numerical models of the interac-tion. For option i), we were not able to constructa viable simulation with a small number of trialruns. The fundamental difficulty is that in orderto produce the spirals and other tidal structuresin NGC 2854 the collision must have a substantialprograde fly-by component with respect to NGC2854. In that case, however, material accretedonto NGC 2856 from NGC 2854 generally has toomuch relative angular momentum to fall directlyonto the NGC 2856 disk. Because of this, we sus-pect that such models occupy a small volume ofthe collision parameter space. We have thereforechosen to focus on models for option ii). These arediscussed in the next two sections.
5.2. Model Details
In the SPH code, hydrodynamical forces arecalculated on a grid with fixed spacing. Gravi-tational forces are computed between particles inadjacent cells, to capture local gravitational insta-bilities. The model galaxies have disks containingboth gas particles and collisionless star particles,as well as rigid dark halo potentials (see Struck(1997) for details). Gas particles with densitiesexceeding a constant density threshold are iden-tified as star-forming particles. These generallyexceed the local Jeans critical mass. A numberof simulations were run; we will only present theresults of the best model.
The evolution of our numerical model for theArp 285 system is presented in Figure 17, with ad-ditional timesteps provided in the Appendix. Weadopt the convention that the model primary cor-responds to the southern galaxy NGC 2854 andthe companion to NGC 2856. The particles in Fig-ure 17 are color-coded according to their galaxyof origin, with red particles originating from theprimary disk and green from the companion. A
total of 13,590 star and 42,900 gas particles wereused in the primary disk and 5640 star and 5640gas particles in the secondary disk. In this model,the length unit = 1.0 kpc, and the time unit is200 Myr. Figure 17 shows four timesteps in thesimulation. The first plot (top left) shows the ap-pearance in the plane of the sky near the timeof closest approach, where the separation betweenthe two galaxies is ∼12 kpc. The second plot (topright) shows the system 370 Myrs after closest ap-proach, while the third (bottom left) shows its ap-pearance 510 Myrs after closest approach. Thesetwo plots match approximately the observed ap-pearance at the present time. The last plot showsthe appearance 740 Myrs after closest approach,the predicted appearance in the future.
The radii of the primary star and gas disks are6.0 and 10.8 kpc, respectively. The companionstar and gas disk radii are both 3.6 kpc. The pri-mary disk was set up in the x-y plane. The com-panion disk is first set up in the x-y plane, thenrotated 40◦ around a y-axis through its center, andthen 90◦ around the z-axis passing through its cen-ter. The relative orbit of the companion is in thex-y plane, so from the point of view of the com-panion disk, the primary approaches at a fairlysteep angle. In the companion disk of the modelin Figure 17, the south side is the near side.
The orbit is counter-clockwise, as is the rotationof the primary, so it sees the encounter as veryprograde. The companion disk rotation, in the x-yplane before the tilts are applied, is clockwise. Theinitial (x,y,z) position of the companion relative tothe primary center is (−8.9, −20.0, 0.0) kpc. Itsinitial relative velocity is (250, 75, 0) km s−1.
The form of the halo potential of the two galax-ies is such that the acceleration of a test particlein this halo is
a =GMh
ǫ2r/ǫ
(1 + r2/ǫ2)nh, (1)
where Mh is a halo mass scale, ǫ is a core radius(set to 2.0 and 4.0 kpc for the primary and com-panion, respectively), and the index nh specifiesthe compactness of the halo. For the primary weuse Mh = 1.3×1010M⊙ and nh = 1.2, which givesa slightly declining rotation curve at large radii.For the companion we take Mh = 2.8 × 1010M⊙
and nh = 1.35, which gives a more rapidly declin-
10
ing rotation curve. The model includes the effectsof dynamical friction with a Chandrasekhar-likefrictional term (see Struck & Smith 2003). Theeffects of this term are small except near closestapproach.
With these potentials, the halo masses for theprimary and companion out to a radius of 12 kpc(about the separation at closest approach) are 3.7× 1010 M⊙ and 3.4 × 1010 M⊙, respectively, witha ratio of about 0.92. This is in accord with thenear equality of the r and i band luminosities ofthe two galaxies (Table 1).
5.3. Model Results
The general morphology of the system is quitewell reproduced by the model, including the mod-erate countertail on NGC 2854 and the bridge (seeFigure 17). A very close encounter is required toproduce a bridge as massive as observed. On theother hand, the moderate-sized tail of the primarygalaxy is the result of a prograde perturbation thatwas not prolonged. These facts, and the relativelylarge separation between the galaxies, argue thatthe relative orbit of the companion is quite ellip-tical, as in the model.
The model primary disk is more circular in ap-pearance than that of NGC 2854. There are sev-eral possible reasons for the difference. The first issimply that the model disk should have a greatertilt relative to the plane of observation (here thex-y plane). The primary disk in the model is in thex-y plane of the sky. However, as noted in Section3.2, based on extinction arguments and the HI ve-locity field, the real disk is somewhat inclined tothe line of sight, with the south side closest. Itis also possible that tidal stretching is responsiblefor the shape of the primary disk. However, inthat case we might expect a longer and more mas-sive tidal tail. This is a rather soft argument atpresent, but it does appear that the bar and thespiral arms of the primary disk are disproportion-ately strong relative to the tail. This suggests thatthe bar and spiral arms were present in the NGC2854 disk before the encounter. This possibilitywas not included in the modeling.
In addition to the bridge, the model compan-ion galaxy has two tidal tails, one made of mate-rial originating from the companion galaxy itself,and one from material accreted from the primary
galaxy along the bridge (see Figure 17). The tidalplume drawn off the companion disk is visible asthe green feature extending northward in the lastthree panels of Figure 17. We equate this structurewith the HI emission extending to the northwestin the Chengalur, Salpeter, & Terzian (1994) HImaps, though it is not at the same position angleas in the observations, being oriented about 45◦
too much to the north compared to the data. Thered feature extending to the northeast in Figure17 we associate with the northern HI tail contain-ing the ‘beads’ of star formation. As with theother feature, the position angle of the model tailis somewhat off from the observed orientation.
In the model, the disk of the companion wastilted relative to the direction of the encounter,so the perturbation had both an orthogonal, ring-galaxy-like component, and a retrograde compo-nent. Waves with circular arc-like components de-velop in the disk of the companion. This behaviormight account for the northwest arc-like structurein NGC 2856 containing the ‘bright spot’.
As with NGC 2854, the observed structure ofNGC 2856 also shows a bar and internal arms.However, in this galaxy the structure of the baris rather irregular. The simulation shows that alarge mass of gas loses angular momentum as aresult of the encounter, and forms a compressedinner disk or bar. Thus, the bar in NGC 2856may be the result of the collision, and may nothave existed before the encounter.
The evolution of the bridge in this model is es-pecially interesting. Because of its elliptical tra-jectory, the companion speeds past its point ofclosest approach as the bridge begins to form. Asthe bridge initially stretches outward from the pri-mary center, it lags behind the companion. Later,the companion nears its apogalacticon relative tothe primary, and slows, so the bridge catches up toit. The bridge material has significant angular mo-mentum relative to the companion center, so theoutermost points swing around to the far side ofthe companion. Shortly thereafter the bridge ma-terial begins to pile-up at an outer radius north-east of the companion. As time goes on, morebridge material streams into this pile-up region,and compression drives star formation. In Fig-ure 17, star-forming gas particles are marked withcyan astericks. In the second panel, one aster-isk is visible in the pile-up region. Comparison
11
of different timesteps and different models showsthat the star formation there is quite stochastic.Sometimes there are a number of star-forming par-ticles there, and occasionally they line up like theobserved ‘beads’. Since the model does not ac-curately represent the effects of self-gravity acrossthis pile-up region, the real environment may trig-ger more such star formation than in the model.
Material from both the accretion tail and thecompanion’s plume eventually accrete onto thecompanion. Gas in the companion is compressedby the tidal perturbation, and experiences pro-longed accretion. In the model, the density thresh-old for star formation is easily exceeded, and cen-tral star formation continues for some time. Thisis consistent with the uncertainties on the age ofthe stellar population in disk clump 3 (the nucleus)of NGC 2856 (Table 4).
Our simulation somewhat resembles models ofpolar ring formation via accretion from a com-panion (Reshetnikov et al. 2006), however, thetwo model galaxies merge before this polar ringproceeds very far in its development (see latertimesteps in Appendix). The star clusters formedin the pile-up region will eventually be carried withthe companion halo into the merger with the pri-mary. They are likely to end up orbiting in theinner halo of the merger remnant and possiblyadding to the globular cluster population there.This is in contrast to dwarf galaxies formed at theend of tidal tails, which may spend long periodsin the outer halo.
6. Discussion
The clumps in the northern tail of NGC2856 are a striking example of the ‘beads on astring’ phenomenon, in which star forming regionsare regularly spaced ∼1 kpc apart along spiralarms and in tidal features (Elmegreen & Efremov1996). Such a ‘beads on a string’ morphol-ogy may be indicative of the gravitational col-lapse of interstellar gas clouds under self-gravity(Elmegreen & Efremov 1996). Similar ‘beadstrings’ are also seen in the interacting systemsIC 2163/NGC 2207 (Elmegreen et al. 2006) andArp 82 (Hancock et al. 2007). In IC 2163/NGC2207, the ‘beads’ resolve into associations of starclusters in higher resolution Hubble Space Tele-scope (HST) images (Elmegreen et al. 2006). In
general, HST images of nearby galaxies show thatyoung star clusters themselves tend to be clus-tered into complexes with typical sizes of ∼1kpc (Zhang, Fall, & Whitmore 2001; Larsen 2004;Bastian et al. 2005). As noted earlier, clumps 3and 4 in the NGC 2856 tail have multiple peaksvisible in the SDSS images. It is possible that theother two tail clumps, clumps 1 and 2, which areunresolved in the SDSS images, will also resolveinto multiple star clusters at higher resolution.
The optical-UV colors of the clumps in theNGC 2856 tail are very blue, and imply ages ofonly ∼4 − 20 Myrs. This is much younger thanthe time since the point of closest approach be-tween the two galaxies, showing that there is atime delay between the initiation of star formationand the time of closest approach between the twogalaxies. According to our numerical model of thissystem, there should be an underlying older stellarcomponent in this tail, made of stars stripped fromthe NGC 2854 disk. Diffuse optical light is clearlypresent between the clumps, and the stellar tailextends 41′′ (7.8 kpc) to the north beyond clump4. However, we are not able to tightly constrainthe age of this diffuse stellar population. Thus it isunclear from the available data whether a stellarcomponent to the tail existed before the currentstar forming episode.
The tail clumps are lower mass than concen-trations in other tails previously classified as tidaldwarf galaxies (see Table 6). They are more sim-ilar in mass to globular clusters than dwarf irreg-ular galaxies. Because of their low mass and thelack of 24 µm detections and calibrated Hα mea-surements, it is not possible to get accurate starformation rates for these clumps. Very roughly,using the 8 µm luminosity for clump 3 in thenorthern tail, and assuming the 8 µm − 24 µm re-lationship found for M51 clumps of Calzetti et al.(2005) and their correlation between star forma-tion rate and 24 µm luminosity, we find a starformation rate for this clump of ∼10−3 M⊙ yr−1.This value is very uncertain due to the low massand the bootstrapping from the 8 µm flux.
In our model, gas from the bridge falling intothe potential of the companion overshoots thecompanion, piling up in an accretion tail on thefar side of the companion. Star formation occursin this region. Our model suggests that the ‘beadson the string’ may be the result of stochastic pro-
12
cesses, albeit in a density enhanced pileup zone. Itis possible that local self gravity is pulling clumpstogether. The spacing between the star formingregions in the model is comparable to the scale oflocal self-gravity in the code. At most timesteps,the star formation is found in a couple of isolatedclumps, without any ‘beads’ appearance. Thus itappears we are seeing this feature at a favorabletime.
The Arp 285 tail is not unique. Accretionfrom a companion along a bridge may have pro-duced the star forming ‘countertails’ in Arp 105(Duc & Mirabel 1994; Duc et al. 1997) and Arp104 (Roche 2007). In addition, the inner tail onthe western side of NGC 7714, which also hasstrong star formation (Smith, Struck, & Pogge1997), may have formed from accretion from thecompanion (Struck & Smith 2003).
Our model suggests that the so-called ‘brightspot’ in the northwestern portion of the NGC 2856disk, and its associated arc, were likely caused by aring-like perturbation of the disk by an encounterwhich was mainly perpendicular to the plane ofthe NGC 2856 disk. The age of the stellar pop-ulation in this region is estimated to be between400 − 1500 Myrs, while the interaction model in-dicates that the point of closest approach betweenthe galaxies occurred between about 300 − 500Myrs ago. This is consistent with the idea thatthe brightness of this ‘spot’ may be due to paststar formation triggered by the encounter.
7. Summary
We have investigated star formation in the in-teracting galaxy pair Arp 285 using Spitzer in-frared, GALEX ultraviolet, and ground-basedoptical data, and have constructed a numericalmodel of the interaction. The northern galaxyin this pair contains an unusual tail-like featureextending perpendicular to the disk. Our modelsuggests that this structure was created by gasfrom the companion falling into the gravitationalpotential of the disk and overshooting the disk.
A series of regularly-spaced knots of recent starformation are seen in this tail. Stellar populationsynthesis suggests that these knots have ages of∼4 − 20 Myrs and masses in the range of globularclusters. The Spitzer 3.6 and 4.5 µm fluxes fromthese tail clumps are higher than expected from
the population synthesis, indicating that either asecond older stellar population is present, or thereare significant contributions to these bands fromhot dust.
The ‘bright spot’ in the NGC 2856 disk has anintermediate-age stellar population (400 − 1500Myrs). This feature and its associated arc mayhave been caused by a ring-like disturbance froman encounter almost perpendicular to the plane ofthe disk. Its brightness might be due to past starformation triggered by the interaction.
We thank the Spitzer, GALEX, and SDSSteams for making this research possible. Thisresearch was supported by NASA Spitzer grant1263924, NSF grant AST-0097616, NASA LTSAgrant NAG5-13079 and NASA GALEX grantGALEXGI04-0000-0026. V. C. acknowledges par-tial support from the EU ToK grant 39965. Wethank Jayaram Chengalur for providing us withan electronic copy of the HI data. We also ac-knowledge Amanda Moffett and Chris Carver forhelp with system administration. This researchhas made use of the NASA/IPAC Extragalac-tic Database (NED) which is operated by theJet Propulsion Laboratory, California Instituteof Technology, under contract with the NationalAeronautics and Space Administration.
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Fig. 1.— The Arp (1966) image of the wide galaxy pair Arp 285 (NGC 2856/4). The northern galaxy NGC 2856 hasan unusual ‘tail’ like feature extending out perpendicular to the disk from the middle of the disk. Toomre & Toomre(1972) suggested that this is material from the bridge/companion, which is accreting onto NGC 2856. Note that thistail appears clumpy in this image. Note the ‘bright spot’ in the northwestern edge of the NGC 2856 disk.
Fig. 2.— A montage of the GALEX, SDSS, and Spitzer images of NGC 2856, the northern galaxy in Arp 285.North is up and east to the left. The field of view is 1.
′1 × 0.′9. Notice the series of clumps in the northern tail. The
tail clumps are enclosed by 1.′′61 black circles in the last panel on the g image, and labeled as in Table 2. The NGC
2856 disk clumps listed in Table 2 are marked in the last panel by 4′′ radius white circles.
16
Fig. 3.— Upper left: The SARA Hα map of NGC2856. Upper right: The SARA R band map of NGC2856. Lower left: The SARA Hα map of NGC 2856(contours) superimposed on the Spitzer 8 µm map(greyscale). Lower right: The SARA Hα map ofNGC 2856 (contours) plotted on the SDSS g map(greyscale). North is up and east to the left. Thismap has been smoothed by a Gaussian with FWHM= 4.
′′5. The field of view is 1.′0 × 0.
′9.
Fig. 4.— An approximately true color multi-filter op-tical SDSS image of NGC 2856. Note that the clumpsin the tail are blue, and the northeastern spiral arm isbluer than the southwestern disk. The field of view is1.′1 × 0.
′9.
Fig. 5.— The smoothed SDSS g image ofNGC 2856 (color), with 21 cm HI contours (fromChengalur, Salpeter, & Terzian 1994) superimposed.North is up and east to the left. Notice the bridgeconnecting this galaxy to its southern companion NGC2854. Also note the bend in the northern tail north ofthe clumps marked in Figure 2. The field of view is3.′0 × 3.
′3. The HI beamsize is 29′′× 29′′, and the HI
contours are (2.6, 4.6, 8.1, 14, 25, 43, and 76) × 1020
cm−2.
17
Fig. 6.— A montage of images of the southern galaxy in Arp 285, NGC 2854. North is up and east to the left. Thefield of view is 1.
′2 × 1.′0. Notice the series of clumps in the northern spiral arm. In the UV and optical, the southern
end of the bar is brighter than the northern end. At longer wavelengths, the disk is more symmetrical. The positionsof the 8 µm-selected clumps in Table 2 are circled on the 8 µm image in the last panel. The circles have 2.
′′8 radii.
18
Fig. 7.— Upper left: The SARA Hα map of NGC2854. Upper right: The SARA R band map of NGC2854. Lower left: The SARA Hα map of NGC 2854(contours) superimposed on the Spitzer 8 µm map(greyscale). Lower right: The SARA Hα map ofNGC 2854 (contours) plotted on the SDSS g map(greyscale). North is up and east to the left. Thismap has been smoothed by a Gaussian with FWHM= 4.
′′5. The field of view is 1.′1 × 1.
′1.
Fig. 8.— An approximately true color multi-filteroptical SDSS image of NGC 2854. North is up andeast to the left. Note that the clumps in the northernarm are blue, and the southeastern end of the bar isbluer than the northern end. The southern arm/tailis also bluer than that in the north. The field of viewis 1.
′2 × 1.′0.
Fig. 9.— The smoothed g image ofNGC 2854, with 21 cm HI contours fromChengalur, Salpeter, & Terzian (1994) superim-posed. North is up and east to the left. The field ofview is 3.
′4 × 4.′2. Note the long tidal tail extending
1.′8 to the south. The HI beamsize is 29′′
× 29′′, andthe HI contours are (2.6, 4.6, 8.1, 14, 25, 43, and 76)× 1020 cm−2.
19
Fig. 10.— The g − r vs. u − g colors of the clumpsin the NGC 2856 tail (magenta open diamonds), theNGC 2856 disk (cyan open diamonds), and the NGC2854 disk (black open circles). The clumps are identi-fied by their numbers in Table 2, labeled in the samecolor. These data are compared with solar metallicityKroupa IMF instantaneous burst population synthesismodels with extinction of E(B − V) = 0 (black filledtriangles), 0.5 (blue open squares), and 1.0 (red opensquares). To show the effect of Hα on the g − r color,for the zero extinction model two curves are shown:with (solid line) and without (dotted line) Hα. Themodel ages start with an age of 1 Myr for the pointon the left end of the curve, increasing by 1 Myr stepsto 20 Myr, then by 5 Myr steps to 50 Myrs, 10 Myrsteps to 100 Myrs, 100 Myr steps to 1 Gyr, and 500Myr steps to 10 Gyr. The red filled triangle shows thecolors of the 25.
′′6 × 9.′′7 region enclosing all four knots
in the northern tail.
Fig. 11.— The NUV − g color of the 25.′′6 ×
9.′′7 region in the northern tail, plotted against g −
r (red filled triangle). The clumps in the tail (ma-genta open diamonds), the NGC 2856 disk (cyan opendiamonds), and the NGC 2854 disk (black open cir-cles) are also plotted. Solar metallicity Kroupa IMFinstantaneous burst population synthesis model colorsare also shown, with extinction of E(B − V) = 0 (blackdiamonds), 0.5 (blue open squares), and 1.0 (red opencircles). The model ages start with an age of 1 Myrfor the point at lower left end of the curve, increasingby 1 Myr steps to 20 Myr, then by 5 Myr steps to 50Myrs, then 10 Myr steps to 100 Myrs, 100 Myr stepsto 1 Gyr, and 500 Myr steps to 10 Gyr. All modelsinclude Hα.
20
Fig. 12.— The Spitzer [4.5] − [5.8] vs. [3.6] − [4.5]color-color plot, showing the location of the clumpsin the NGC 2856 tail (magenta open diamonds), theNGC 2856 disk (cyan open diamonds), and the NGC2854 disk (black open circles). The clumps are labeled.The colors of M0III stars (open dark blue square),from M. Cohen (2005, private communication), andthe mean colors of the field stars of Whitney et al.(2004) (magenta open triangle) are also shown. Thecolors of normal stars all lie within 0.5 magnitudesof 0, 0 in this plot (M. Cohen 2005, private com-munication). We have also plotted the locations ofthe clumps in Arp 107 and Arp 82 as green aster-isks (Smith et al. 2005a; Hancock et al. 2007), exclud-ing likely foreground stars, background quasars, upperlimits, and point with uncertainties > 0.5 magnitudes.The observed Spitzer colors (Flagey et al. 2006) fordiffuse dust towards several positions in the Milky Wayare also plotted (blue X’s). The red diamond showsthe colors of the 25.
′′6 × 9.′′7 region enclosing all four
knots in the northern tail. The errorbars include bothstatistical uncertainties and an uncertainty in the col-ors due to varying the sky annuli (see text).
-1 -0.5 0 0.5 1
-15
-14
-13
-12
Fig. 13.— The UV-mid-infrared spectral energydistribution of clump 3 in the northern tail (filledsquares), including the upper limit in the SDSS z band(arrow). The 2MASS upper limits are also shown. Thelong dashed black curve is the best fit solar metallicitypopulation synthesis model (4 Myrs, E(B−V)=0.1).The dot-dashed blue and short dashed red curves showthe youngest and oldest solar metallicity models re-spectively, with their associated best-fit extinctions.All of the models have been normalized to the g bandflux. The red dotted curve is the Draine & Li (2007)Milky Way dust model with U = 100 and qPAH =4.6%, scaled to the 8 µm flux. The solid black curveis the sum of these three components. Note the con-tribution from the 3.3 µm PAH feature to the 3.6 µmSpitzer band, and the Hα contribution to the r band.Also note that the 3.6 µm and 4.5 µm fluxes are muchhigher than expected from the stellar population syn-thesis model, suggesting contributions from either hotdust, as in the dust model shown, or a second colderstellar population undetected in the population syn-thesis.
21
-1 -0.5 0 0.5 1-13
-12
-11
-10
Fig. 14.— The UV-mid-infrared spectral energy dis-tribution of clump 2 in the NGC 2856 disk. Symbolsand curves are as in Figure 13. Note that, althoughthe age is similar to that of tail clump 3 (Table 4), theSED is very different because of the higher extinction.Starlight contributes a higher fraction of the 3.6 and4.5 µm flux in this clump than in tail clump 3.
-1 -0.5 0 0.5 1-14
-13
-12
-11
Fig. 15.— The UV-mid-infrared spectral energy dis-tribution of the ‘bright spot’ at the edge of the NGC2856 disk. Symbols and curves are as in Figure 13.Note the upper limits at 5.8 µm and 8.0 µm, as wellas the GALEX upper limits. Starlight can account forthe 3.6 µm and 4.5 µm emission.
22
Fig. 16.— The 3.6 µm luminosity of the Arp 285clumps, compared to their masses implied by the pop-ulation synthesis models. The magenta open squaresare the tail clumps, the blue filled triangles are clumpsin the NGC 2856 disk, and the green crosses are theNGC 2854 disk clumps. The black open circles arethe Arp 82 clumps, from Hancock et al. (2007). Aconstant M/L3.6 ratio of 1 M⊙/L⊙ is represented bythe solid black line, while the dotted line is M/L3.6
= 10 M⊙/L⊙. The 3.6 µm luminosity was calculatedassuming the FWHM of the bandpass ∆ν of 16.23 ×
1012 Hz.
23
Fig. 17.— Snapshots of the model gas disks. The stellar appearance is similar. Red particles originated in theprimary galaxy, green in the companion. The top left panel shows a time near closest approach (T = 0 Myrs). Thecompanion has swung in from the lower left, and swings around to an apogalacticon point at later times. The upperright panel and the lower left are at times near the present (T = 370 Myrs and 510 Myrs, respectively). The lowerright panel is at a later time (T = 740 Myrs), when the companion begins to fall back to merge with the primary.In the first three panels every third gas particle is plotted with a dot. In the second and third panels blue astericksmark star-forming particles, except those within 5 kpc of the primary center, which were omitted for clarity. Thestar forming region in the northern tail was produced from gas accreted from the companion, while the star formingregions in the central region of the northern galaxy were produced from gas that originated in the northern galaxy.The star forming regions in the bridge, southern galaxy, and southern tail were formed from gas that originated in thesouthern galaxy. In the final panel only every fifth particle was plotted, to show the persistent spiral in the primarydisk. The motion of the companion around the point of greatest separation is very slow, so little positional change isevident in the last three panels. Later timesteps are shown in the Appendix.
24
Table 1
Total Galaxian Magnitudesa
ID FUV NUV u g r i z F3.6µm F4.5µm F5.8µm F8.0µm F24µm
aThe statistical uncertainties in the optical and UV magnitudes are typically ∼0.01 magnitudes. The Spitzer uncertainties are as given inSmith et al. (2007).
bAt 9h 24m 2.9s, +49◦ 14′ 41′′ (J2000).
25
Table 2
Magnitudes and Flux Densities for Clumps in Arp 285a
ID R.A. Dec. FUV NUV u g r i z F3.6µm F4.5µm F5.8µm F8.0µm
aFor FUV and NUV, respectively. NUV from bright stars in field. FUV from bright stars inthe Arp 65 field.
bFrom bright stars in the field.
cFor 3.6, 4.5, 5.8, & 8.0 µm, respectively. From the IRAC Data Manual, Version 3.
27
Table 4
Model Ages and Extinctions for Clumps with Solar Metallicity Modelsa
Clump Age E(B−V) Colors Used(Myr) (mag)
NGC 2856 Tail Clumps1 18 ± 149
15 0.3 ± 0.3 NUV − g, g − r, r − i2 7 ± 42
2 0.1 ± 0.20.1 NUV − g, u − g, g − r, r − i, i − z
3 4 ± 21 0.1 ± 0.1 NUV − g, u − g, g − r, r − i
4 15 ± 812 0.0 ± 0.2
0.0 NUV − g, g − r, r − iNGC 2856 Disk Clumps
1b 142 ± 1358137 0.5 ± 0.2
0.5 u − g, g − r, r − i2 6 ± 2
1 0.68 ± 0.040.1 NUV − g, u − g, g − r, r − i, i − z
3c 6900 ± 32006843 0.2 ± 0.7
0.1 u − g, g − r, r − i, i − z4 88 ± 281
38 0.6 ± 0.10.3 NUV − g, u − g, g − r, r − i, i − z
NGC 2854 Disk Clumps1 8 ± 4
5 0.4 ± 0.40.2 NUV − g, u − g, g − r, r − i, i − z
2 7 ± 23932 0.6 ± 0.2
0.6 u − g, g − r, r − i, i − z3 4 ± 5
3 0.7 ± 0.40.5 u − g, g − r, r − i, i − z
4 7 ± 12 0.6 ± 0.2
0.1 u − g, g − r, r − i5 6 ± 1394
2 0.5 ± 0.20.3 u − g, g − r, r − i, i − z
6 36 ± 316432 0.4 ± 0.5
0.3 u − g, g − r, r − i, i − z7c 7 ± 7893
2 1.3 ± 0.11.0 u − g, g − r, r − i
8 1200 ± 5500259 0.2 ± 0.1
0.2 u − g, g − r, r − i, i − z
aAll clump ages obtained with instantaneous burst models. The oldestmodels run were 10 Gyrs old.
b‘Bright spot’ in northwestern edge of disk. The upper limit on the NUVflux further constrains this age to >
∼400 Myrs (see Figure 14 and Section 5).
cNucleus
28
Table 5
Optical Absolute Magnitude Ranges for Various Objects
Class Object MB MV MR Notes and References
Arp 285 Tail Clumps −12.0 to −10.6 −12.3 to −10.8 −12.3 to −10.9 This worka
NGC 2856 ‘Bright Spot’ −13.6 −14.3 −15.0 This worka
Other Arp 285 Disk Clumps −15.9 to −12.1 −17.1 to −13.2 −17.8 to −13.3 This worka
Nearby Dwarf Irregular Galaxies −18 to −8; median=−13.2 Karachentsev et al. (2004)Local Group Irregular Leo A −11.3 Karachentsev et al. (2004)Local Group Irregular GR 8 −12.0 Karachentsev et al. (2004)
M81 Dwarf A −12.4 Patterson & Thuan (1996)Tidal Dwarf Galaxies −17 to −12.5 b
Arp 82 Tail Clumps −15.9 to −13.9 Hancock et al. (2007)Super Star Clusters −16 to −12 c
30 Dor R136 Star Cluster in LMC −11.3 O’Connell, Gallagher, & Hunter (1994)
aCalculated using the SDSS color transformations for stars given in Jester et al. (2005). Does not include correction for internalextinction.
bWeilbacher, Duc, & Fritze-v. Alvensleben (2003) and Higdon, Higdon, & Marshall (2006). Fainter end of NGC 5291 TDG range ex-trapolated from high end assuming constant B − [3.6] color.
cHoltzman et al. (1992, 1996); O’Connell, Gallagher, & Hunter (1994); Whitmore et al. (1993); Whitmore & Schweizer (1995);Schweizer et al. (1996); Watson et al. (1996).
