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Deep Radio, Optical, and Infrared Observations of SGR 1900+14

We present HST/STIS, Keck J- and K_s-band, and VLA 332-MHz and 1400-MHz images of the region around the Soft Gamma-ray Repeater SGR 1900+14. No non-stellar sources were detected at the position of SGR 1900+14, giving 3-sigma limits of S_{332 MHz}<6 mJy, S_{1.4 GHz}<1.5 mJy, m_{50CCD}>29.0 mag, J>22.8 mag, K_s>20.8 mag (point sources), and S_{332 MHz}<6.1 mJy arcmin^{-2}, S_{1.4 GHz}<6.2 mJy arcmin^{-2} (extended emission). Given the very high extinction, the STIS and J-band non-detections are not constraining, but the K_s-band limit rules out basic accretion disk models for SGR 1900+14, and is almost comparable in depth to the K_s-band detection for the anomalous X-ray pulsar 4U 0142+61. Finally, we report the detection in this field of three new candidate supernova remnants, SNRs G043.5+00.6, G042.0-00.1, and G041.5+00.4.


Introduction
Soft γ-ray Repeaters (SGRs; see Hurley 2000 for a recent observational review) are sources both of bursts of hard X-rays/γ-rays and softer quiescent X-ray emission. They are generally thought to be young (< 10 4 yr) neutron stars with extremely strong magnetic fields, or magnetars (e.g. Duncan & Thompson 1992;Thompson & Duncan 1993). This belief was motivated by the energetics of their bursts (Paczyński 1992;Thompson & Duncan 1995), their associations with young supernova remnants (SNRs; Evans et al. 1980;Kulkarni & Frail 1993;Vasisht et al. 1994), and the detection of X-ray pulsations with long (5-10 s) periods and large (∼ 10 −11 s s −1 ) spin-down rates (Kouveliotou et al. 1998;Hurley et al. 1999d). However, more recent observational evidence (Lorimer & Xilouris 2000;Hurley et al. 1999b;Kaplan et al. 2002) has shown that in most cases the SGR/SNR associations are false and can be attributed to random chance (Gaensler et al. 2001). Instead of being constrained by the SNR associations, the small ages of the SGRs are now limited by the proximity of the SGRs to sites of very massive star formation (Fuchs et al. 1999;Vrba et al. 2000).
An alternate theory, first proposed (van Paradijs, Taam, & van den Heuvel 1995;Chatterjee, Hernquist, & Narayan 2000) for the emission from the related Anomalous X-ray Pulsars (AXPs; see Mereghetti 2000), and now for the SGRs as well (Marsden et al. 2000;Marsden et al. 2001; but see Duncan 2001), is accretion from a "fallback" disk of material produced during the supernova explosion onto a central neutron star. In this model, the neutron stars are spinning near the equilibrium period which is proportional to the dipole field strength. This model then explains the spin-down and period clustering of the AXPs: for a typical field strength of 10 12 G, the resultant spin period is ∼ 8 s.
To date, no confirmed persistent counterparts have been found for SGRs at optical/IR or radio wavelengths (e.g. Kaplan et al. 2001), although Eikenberry et al. (2001) may have found an infrared counterpart to SGR 1806−20. It generally has been the case that when counterparts are identified at other wavelengths, especially optical or IR, there is considerable progress in our understanding of these objects. For example, the identification of an optical/IR counterpart of the AXP 4U 0142+61 (Hulleman, van Kerkwijk, & Kulkarni 2000) has posed the strongest challenge to accretion models. In view of this, it is important to press upon sensitive multi-wavelength observations of SGRs. Here we present optical, IR, and radio observations of the very accurately localized SGR 1900+14 (Frail, Kulkarni, & Bloom 1999;Hurley et al. 1999d).
Searches for SGR 1900+14 at wavelengths other than X-ray and γ-ray have not detected any quiescent counterpart Lorimer & Xilouris 2000;Eikenberry & Dror 2000;Vrba et al. 2000). However, Lorimer & Xilouris (2000) did detect a young (10 4−5 yr old) radio pulsar (PSR J1907+0918) located in the vicinity 1 of both SGR 1900+14 and G042.8+00.6, suggesting that G042.8+00.6 could be associated with PSR J1907+0918 just as easily as with SGR 1900+14. Vrba et al. (2000) have suggested that SGR 1900+14 is somehow related to a compact cluster found in its vicinity (about 12 ′′ away). A similar cluster has been identified near SGR 1806−20 (Fuchs et al. 1999. We note, though, that no such cluster is found in the vicinity of SGR 0526-66 in the HST images presented by Kaplan et al. (2001). Furthermore, SGR 1806−20 and SGR 1900+14 are found at low Galactic latitudes in the inner Galaxy and the chance coincidence probability with star clusters, the rare Luminous Blue Variables (LBVs) and SNRs is not negligible (e.g. Gaensler et al. 2001). Indeed, there is now significant doubt about the association of SGR 1806−20 with the SNR G10.0−0.3 (Hurley et al. 1999b;Kaplan et al. 2002). Nonetheless, the presence of star clusters in the vicinity of SGR 1900+14 appears intriguing, but while our radio maps are deeper than previous images (e.g. Vasisht et al. 1994) we find no radio source or nebulae associated with this star cluster (see . For the X-ray properties of SGR 1900+14 in what follows, we use the power-law + blackbody fit to the second epoch (the one closer to normal quiescence) of the recent Chandra observations by Fox et al. (2001): N H = 2.3 ± 1 × 10 22 cm −2 , power-law index Γ = 1.9 ± 0.3, blackbody temperature kT ≈ 0.49 ± 0.03 keV, and an un-absorbed 0.5-10 keV X-ray flux f U X = 2.0 × 10 −11 erg s −1 cm −2 (also see Kouveliotou et al. 2001). This is consistent with the ASCA measurements of Hurley et al. (1999d). We adopt a distance d = 5 kpc (Hurley et al. 1999d), which is based on hydrogen absorption. The hydrogen column density implies a visual extinction A V = 12.8±0.8 mag, assuming N H = 1.79±0.03×10 21 A V cm −2 (Predehl & Schmitt 1995). These values differ significantly from those assumed for the star cluster (d = 12-15 kpc, A V ≈ 19.2 mag; Vrba et al. 1996), and this difference may make the case for associating SGR 1900+14 and the cluster less compelling.

Observations
We have carried out a campaign at optical/IR and radio wavelengths to detect SGR 1900+14. The radio data can also determine the existence of a low surface-brightness nebula surrounding SGR 1900+14, and localize SGR 1900+14 relative to nearby SNRs. See Table 1 for a log of all the observations.

Optical
The optical data consist of 5 exposures taken by the Space Telescope Imaging Spectrograph (STIS) aboard the Hubble Space Telescope (HST), totaling 5195 s. The data were taken in imaging mode using the clear (50CCD) filter, giving a bandpass centered at roughly 585.2 nm and a full-width at half-maximum (FWHM) of ≈ 400 nm. The first exposure had a duration of 30 s, and was used for astrometry (see below). The remaining exposures were ∼ 1300-s long and were dithered in the usual manner to improve resolution and mitigate the effects of CCD irregularities.
The reduction followed the standard procedure for STIS images. We followed the standard drizzling reduction procedure (Fruchter & Mutchler 1998) expanding the individual pixels by a factor of two. This drizzling changed the pixel scale from 0. ′′ 05 pixel −1 to 0. ′′ 025 pixel −1 . Typical stellar sources had a FWHM of 0. ′′ 095; see Figure 1.
For the astrometry, we selected from the USNO-A2.0 catalogue (Monet 1998) all 76 stars that overlapped with a 30-s R-band image taken on 28 August 1998 with the Low-Resolution Imaging Spectrograph (LRIS; Oke et al. 1995) on the Keck telescope. For the 54 stars that were not overexposed, we measured centroids and corrected for instrumental distortion using a bi-cubic function determined by J. Cohen (1997, private communication). We used these to solve for the zero-points in each coordinate, the plate scale, and the position angle on the sky. After rejecting 8 outliers (residual larger than 0. ′′ 8), the root-mean-square residuals were 0. ′′ 29 in each coordinate. This is somewhat larger than we found for other projects, but not atypical for relatively crowded Galactic fields on the Northern hemisphere. Given this, the astrometry is tied to the USNO-A2.0 system at the 0. ′′ 05 level. The typical accuracy with which the USNO-A2.0 catalogue is on the International Celestial Reference Frame (ICRF) is about 0. ′′ 2 (Deutsch 1999).
The astrometry was transferred to the STIS image using 43 stars close to the radio position, solving again for plate scale, rotation and zero-points. The rms residuals were 0. ′′ 02 in each coordinate, giving an overall astrometric uncertainty of 0. ′′ 21 in each coordinate. The 3-σ limiting magnitude, based on calibration using the photometric keywords in the image header, is 29.0, or f ν,min = 10 nJy 2 .

Infrared
The infrared data were taken with the Near Infrared Camera (NIRC; Matthews & Soifer 1994) mounted on the 10-m Keck I telescope. The observations had photometric conditions and very good seeing (0. ′′ 35 at K s -band). The data were reduced with custom IRAF software. The images were dark subtracted, flat-fielded, and corrected for bad pixels and cosmic rays. We then made object masks, which were used to construct improved flat fields for a second round of reduction. The data were finally registered, shifted, and co-added; see Figure 2.
The astrometry was performed with 50 stars from the deep STIS image. This solution had errors of ±0. ′′ 04 (in each coordinate) relative to the STIS image, or again ±0. ′′ 21 relative to the ICRF. For photometric calibration we used the standard star SJ 9101 (Persson et al. 1998). The 3-σ limiting magnitudes are K s ≈ 20.8 and J ≈ 22.8, corresponding to f ν,Ks ≈ 3.1 µJy and f ν,J ≈ 1.2 µJy, based on the calibration of Bessell, Castelli, & Plez (1998).

Radio
The 1.4-GHz data were taken at the Very Large Array (VLA) in its C and D configurations using the standard continuum mode, with 2 × 50 MHz bandwidth. The separate observations were calibrated independently, following the standard procedure in AIPS. The phase calibrator was JVAS J1925+2106, and the flux calibrator was 3C 48. The data were then self-calibrated and combined for imaging. The nominal beam size was ≈ 20 ′′ , and the rms noise is ≈ 0.5 mJy beam −1 ; see Figure 3.
To aid in the detection of diffuse structure we also imaged using a uv-taper of 2.5 kλ, giving an effective beam size of ≈ 60 ′′ . The rms noise is ≈ 2 mJy beam −1 , largely due to confusion from unresolved sources in the Galactic plane; see Figure 4.
The 332 MHz data were taken at the VLA in spectral-line mode with 32 channels of 97 kHz each to aid in the excision of radio-frequency interference (RFI). We reduced the observations in AIPS, following the procedure for low-frequency data reduction 3 (see also LaRosa et al. 2000). First we removed any egregious RFI. Then we phase (using 3C 380) and gain (using 3C 286 and 3C 48) calibrated the fringe visibilities. We again searched for and removed RFI, then averaged the channels together. We mapped and self-calibrated the data iteratively, using 3-dimensional Fourier transforms in IMAGR, the basic AIPS deconvolution and imaging task. While imaging we used 25 planar facets arranged in a 5 × 5 grid, each of which had 256 × 256 13 ′′ cells. These facets were used to mitigate the effects of noncoplanar arrays that are present at low frequencies and with wide fields (Cornwell & Perley 1992;Perley 1999). Finally, we re-assembled the facets and flattened them into a single planar map using the task FLATN. The rms noise is ≈ 2 mJy beam −1 , limited primarily by confusion due to unresolved emission in the Galactic plane and to sidelobes from the bright W49 complex. The final beam size is ≈ 50 ′′ , comparable to that of the tapered 1.4-GHz images; see Figures 5 and 6.

Optical/IR
For the infrared and optical observations the astrometric uncertainty is 0. ′′ 21. Combined with the 0. ′′ 15 uncertainty from Frail et al. (1999), this gives us a 99% confidence radius of 0. ′′ 79. We examined all sources in or near a circle of this radius around the position of Frail et al. (1999). In the deepest observation, the K s -band image, we detect 7 sources that we label A-G on Figures 1 and 2 (only sources A-D are visible in Figure 1). Of these, C and D are significantly blended on all the images, as are sources C and G. Sources E, F, and G are too faint to measure on the STIS image, and sources E and F are too faint to measure on the J-band image. Because of the degree of blending, we used a point-spread function (psf) fitting and subtraction technique to perform photometry on the infrared data, using the DAOPHOT package. First we measured the psf's of a number (≈ 30) of stars in the field without close neighbors, and constructed an average psf for each image. We then fit this psf to each source in Figures 1 and 2, subtracted the fitted stellar profile from the images, and used the subtracted images to fit the sources that were hidden by the original source. After iterating this procedure, we were able to obtain consistent magnitudes for most of the sources. For the sources that we were able to fit, all had sizes consistent with unresolved sources. For the STIS data, we adjusted the measured magnitudes to a nominal infinite aperture with a 0.13 ± 0.03 mag correction, based on the observed STIS psf 4 . We present the positions and magnitudes of sources A-G in Table 2.
In Figure 7, we show a color-color diagram for the sources in Table 2 and 37 field stars for which we were able to perform reasonably accurate photometry in all bands. The sources from Table 2 all lie in the region populated by the other sources. In Figure 7 we also show model main-sequence colors (based on Cox 2000, p. 388) for A V = 5-10 mag. These fit the majority of the points, with the exception of source D. However, since source D was blended with source C (making the IR photometry less reliable) and it is still in the region of the colorcolor diagram populated by other stars (some of whom were slightly blended themselves), we consider it unlikely that D has significantly non-stellar colors, and therefore is probably not a counterpart to SGR 1900+14 (both the disk models for AXPs [Perna, Hernquist, & Narayan 2000] and the optical/IR detection of an AXP [Hulleman, van Kerkwijk, & Kulkarni 2000] have non-stellar colors, and we expect something similar for SGRs). In Figure 8 we show color-magnitude diagrams for the sources in Table 2 and 37 field stars. Again, the sources populate the same region (although the field stars were chosen to be somewhat brighter than the faintest of the candidate sources), and are well bounded by model main sequences with d = 8 kpc and A V = 5-10 mag. Source D, in particular, has slightly strange colors (also as seen in Figure 7) but its magnitudes are well within the range of the main sequence. Main sequences with d = 12-15 kpc and A V ≈ 19.2 mag, such as those plotted in Vrba et al. (2000), are too faint and too red to match these sources. Therefore the sources in Table 2 are likely to be in the foreground of the cluster found by Vrba et al. (2000). Table 2 is too faint for accurate photometry given its proximity to other sources. It could be a counterpart of SGR 1900+14, and we require additional deep, high-resolution infrared imaging to settle this issue. Regardless, it is near our detection limit, so we can place general upper limits to the emission from SGR 1900+14 of m 50CCD 29.0 mag, J 22.8 mag, and K s 20.8 mag, consistent with previously published limits (Eikenberry & Dror 2000;Vrba et al. 2000). We show these limits along with the CXO X-ray spectrum (Fox et al. 2001) in Figure 9. These limits translate into unabsorbed ("U") X-ray-to-infrared flux ratios f U X /(ν J f U ν,J ) 160 and f U X /(ν Ks f U ν,Ks ) 1.1 × 10 3 (assuming A V = 12.8 mag). This last limit is only a factor of ∼ 10 less than the detected X-rayto-K s limit found for the AXP 4U 0142+61 (F. Hulleman 2001, personal communication; see Figure 9) and thus we can exclude the basic disk model for AXP/SGR emission (Perna et al. 2000; see also Kaplan et al. 2001). If a disk does exist, then it must be truncated at R ≈ 10 10 cm. As seen for SGR 0526−66 , the optical/IR limits require that the X-ray power-law break between the X-ray and the optical bands, but allow for a continuation of the X-ray blackbody.

Source G in Figures 1 and 2 and in
We were able to detect but not photometer many luminous members of the massive star cluster that Vrba et al. (2000) proposed as the origin of SGR 1900+14, as we only had a small number of exposures where the central M giant stars were unsaturated. In the remaining exposures these stars were too saturated to allow good psf subtraction, and their scattered light, bleed trails, and readout trails prevent accurate measurement of the stars in the cluster.

Serendipitous Sources
The 332-MHz image ( Figure 5) shows three sources besides the known SNRs that have shell-like morphologies but are in neither the current version of the Galactic SNR catalog (Green 2000) nor any online database. We consider these sources to be candidate SNRs, but with only limited data at other wavelengths we are unable to confirm an SNR classification for any of the three. The sources could be H II regions or other similar sources and require additional confirmation.
Here we present what limited information we have on each.
Without a measurement at another frequency, it is hard to determine the properties of G043.5+00.6. We were not able to find any archival X-ray or radio data at this position that could definitively classify G043.5+00.6. However, there is a marginal detection in the 1.4-GHz NRAO VLA Sky Survey (NVSS; Condon et al. 1998). The source is present in the NVSS with the same general morphology although it is too faint to accurately measure a flux. This faintness, though, suggests that G043.5+00.6 is not a thermal source, as it would have a NVSS flux of ∼ 1 Jy and would then be easily detectable. There are two X-ray sources from the EXOSAT medium-energy slew survey catalog (Reynolds et al. 1999) that lie in or near G043.5+00.6: EXMS B1906+095 and EXMS B1906+094. Both sources had been associated with the X-ray binary 4U 1907+097, which is within their ∼ 10 ′ position uncertainties, but the positions of the EXOSAT sources are closer to G043.5+00.6 than to 4U 1907+097 (Schwartz et al. 1980). It is therefore conceivable that the EXOSAT sources are associated with regions of G043.5+00.6, but we require additional X-ray imaging to confirm this.
We can estimate the distance to G043.5+00.6 very roughly using the updated Σ-D relation (Case & Bhattacharya 1998). Assuming a spectral index α = 0.5, appropriate for many shell-type SNRs (e.g. Case & Bhattacharya 1998), we find Σ 1 GHz = 3.8 × 10 −21 W m −2 Hz −1 ster −1 , which implies a linear diameter D ≈ 40 pc and distance d ≈ 8.1 kpc. We note that these numbers are extremely imprecise (given both intrinsic uncertainties in the Σ-D relation and in estimates of the 1-GHz surface brightness) probably having errors approaching 100%. However, this value is consistent with G043.5+00.6 being near the tangent point of spiral arm 3 of the Taylor & Cordes (1993) electron density model. Also intriguing are the objects in Figure 5 marked G042.0−00.1 and G041.5+00.4 (see Figure 11). These objects have morphologies that are less well-determined than that of G043.5+00.6 and are further out in the VLA primary beam, so we cannot make very useful measurements, but their shapes are suggestive of either SNRs or H II regions. G042.0−00.1 is centered at α 2000 = 19 h 08 m 11 s , δ 2000 = +08 • 00.5 ′ , is roughly circular with radius 4.3 ′ , and has a 332-MHz flux of 1.8 ± 0.4 Jy (this value has been corrected for the attenuation of the primary beam of the VLA antennas and is therefore highly uncertain). The protruding region to the south-east of G042.0−00.1 was identified as an H II region on the basis of radio recombination lines by Lockman, Pisano, & Howard (1996), but their data do not mention the rest of the complex and their angular resolution should be able to distinguish between the center of G042.0−00.1 and the smaller region. G042.0−00.1 is present in the NVSS with similar morphology to the 332-MHz image and flux S NVSS = 0.3 ± 0.1 Jy, suggesting a non-thermal spectrum and leading to a provisional classification as an SNR.
G041.5+00.4 is centered at α 2000 = 19 h 05 m 48 s , δ 2000 = +07 • 46.1 ′ . It has a complicated morphology, with a brightened rim to the north-east and a compact core to the south-west (α 2000 = 19 h 05 m 37 s , δ 2000 = +07 • 44.8 ′ ), suggestive of a pulsar wind nebula. The source has an overall radius of 4.9 ′ and a 332-MHz flux of 1 ± 1 Jy (again, this is a highly uncertain value). Comparison of the overall 332-MHz flux with the NVSS flux of 1.3 ± 0.2 Jy leads to a flat spectrum like that of core-dominated SNR or an H II region. We need more data to conclusively classify all three of these objects; further analysis and observations are in progress.
The wealth of information available in Figure 5, concerning both already known objects and newly discovered ones, is readily apparent. It serves to illustrate how productive lowfrequency wide-field imaging of the Galactic plane with high resolution can be (e.g. LaRosa et al. 2000).
We would like to thank J. Lazio for his assistance with the low-frequency data reduction and D. Fox for useful comments. DLK is supported by the Fannie and John Hertz Foundation, SRK by NSF and NASA, and MHvK by a fellowship from the Royal Netherlands Academy of Arts and Sciences. The NASA/ESA Hubble Space Telescope is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract No. NAS5-26555.The VLA is operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. The W.M. Keck Observatory is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. We have made extensive use of the SIMBAD database, and we would like to express our appreciation of the astronomers who maintain this database.        Table 2 (circles) and 37 field stars (points). Also plotted are model main-sequence colors, assuming V ≈ m 50CCD , for A V = 10 mag (solid line) and A V = 5 mag (dashed line), both at a distance of 8 kpc.  Table 2 (circles) and 37 field stars (points). Also plotted are model main-sequences, assuming for A V = 10 mag (solid line) and A V = 5 mag (dashed line), both at a distance of 8 kpc. J vs. m 50CCD − J (left); K s vs. J − K s (right).   Fig. 9.-Spectral energy distribution for SGR 1900+14, including CXO power-law + blackbody spectrum (Fox et al. 2001; also see Hurley et al. 1999d) and upper limits. Limits are raw values (open triangles), and corrected for extinction with A V = 12.8 mag and A V = 19.2 mag (filled triangles). Solid thick line marked "CXO" is the measured Chandra spectrum; the thick dotted line marked "BB" is the blackbody component of the fit, and the thick dotdashed line marked "PL" is the power-law component, both extrapolated to lower energies. The thick line marked "AXP 4U 0142+61" is the extinction corrected detected K s flux of 4U 0142+61 (F. Hulleman 2001, private communication) scaled by the ratio of X-ray fluxes of 4U 0142+61 and SGR 1900+14. Also plotted are the basic disk model (thin solid line) from Perna et al. (2000), which the current limits rule out, and a disk truncated at 10 10 cm (thin dashed line).