Abstract
We present the results of near-infrared [Fe II] and H 2 line imaging and spectroscopic observations of the supernova remnant 3C 396 using the Palomar 5 m Hale telescope. We detect long, filamentary [Fe II] emission delineating the inner edge of the radio emission in the western boundary of the remnant in imaging observations, together with a bright [Fe II] emission clump close to the remnant center. There appears to be faint, diffuse [Fe II] emission between the central clump and the western filamentary emission. The spectroscopic observations determine the expansion velocity of the central clump to be ∼ 56 km s −1 . This is far smaller than the expansion velocity of 3C 396 obtained from X-ray observations, implying the inhomogeneity of the ambient medium. The electron number density of the [Fe II] emission gas is ≤ 2,000 cm −3 . The H 2 line emission, on the other hand, lies slightly outside the filamentary [Fe II] emission in the western boundary, and forms a rather straight filament. We suggest that the [Fe II] emission represents dense clumps in the wind material from the red supergiant phase of a Type IIL/b progenitor of 3C 396 which have been swept up by the

supernova remnant shocks. The H 2 emission may represent either the boundary of a wind bubble produced during the main-sequence phase of the progenitor or molecular clumps left over inside the bubble. We propose that the near-infrared [Fe II] and H 2 emission observed in several supernova remnants of Type IIL/b SNe likely has the same origin.
Subject headings: circumstellar matter -infrared: ISM -ISM: individual(3C 396) -shock waves -supernova remnant
Introduction
The near-infrared (IR) [Fe II] lines are useful tools for studying the radiative shocks of supernova remnants (SNRs). The increased abundance of the gas phase iron by shockinduced sputtering of the dust grains and/or the creation of an extensive partially ionized zone by shock heating can substantially enhance the [Fe II] emission (e.g., Greenhouse 1991; Mouri et al. 2000) . This, together with less severe extinction effects than the optical and the recent advents of near-IR imaging cameras covering sufficiently wide fields (e.g., ≥ 5 ′ ), can potentially make the observations of the near-IR [Fe II] lines a very competitive and efficient way to study the radiative shocks of SNRs, especially in the Galactic plane where the extinction is usually large.
A good example is the young core-collapse SNR G11.2−0.3 where we recently found the brightest near-IR [Fe II] emission of all known SNRs using Wide-field Infrared Camera aboard the Palomar 5-m telescope (Koo et al. 2007 ). Besides G11.2−0.3, a few other SNRs -such as RCW 103 (Oliva et al. 1999 ), 3C 391 (Reach et al. 2002) , W44 (Reach et al. 2005) , and W49B (Keohane et al. 2007 ) -have been detected in the near-IR [Fe II] line emission.
1 Note that all these SNRs have a core-collapse nature and, except for W44, all of them are relatively young SNRs. Interestingly, in addition to the [Fe II] lines, near-IR H 2 emission of shocked molecular gas has been detected from all these SNRs. The H 2 emission is distributed close to, but slightly outside, the [Fe II] emission from the center of the SNRs (e.g., Oliva et al. 1990 ; Keohane et al. 2007; Koo et al. 2007) , which is opposite to the standard picture of the molecular shocks where molecules form in the downstream behind a recombination region. This apparent pattern of the positional 'reversal' between the [Fe II] and H 2 emission may have important clues to understanding the pre-supernova environment of SNRs, calling more systematic investigation of the [Fe II] and H 2 emission together in (young) core-collapse SNRs.
In this paper we present the near-IR imaging and spectroscopic observations of [Fe II] and H 2 lines from another core-collapse SNR 3C 396 (also known as G39.2-0.3). In radio 3C 396 is a shell-type SNR of ∼ 8 ′ ×6 ′ size (Anderson & Rudnick 1993) composed of multiple elliptical arcs and filaments (Figure 1) , with a thick, half ring-like feature in the west. The ASCA X-ray observation revealed the existence of a pulsar wind nebula within a thermal Xray shell (Harrus & Slane 1999) , and the former was later resolved by Chandra, although no X-ray pulsation was detected (Olbert et al. 2003) . Recently Spitzer detected bright mid-IR emission from this SNR along its western boundary in the radio (Lee 2005; Reach et al. 2006) . The mid-IR emission is filamentary, and is prominent in the 4.5 and 5.8 µm bands. If the origin of the mid-IR emission is ionic and molecular lines as suggested (Lee 2005; Reach et al. 2006) , we expect the existence of the corresponding near-IR emission of the [Fe II] lisarodriguez@example.com µm and H 2 lisarodriguez@example.com µm from this SNR, which motivated our observations presented in this paper. The distance to 3C 396 is rather uncertain; however, most of the previous observations of H and OH indicate that it is beyond the tangential lisarodriguez@example.com kpc, and lisarodriguez@example.com kpc (Caswell et al. 1975; Becker & Helfand 1987; Green 1989) . In this paper, we adopt the distance of 8.5 kpc determined based on the velocity of a nearby molecular cloud in 3C 396 (see § 4). We organize this paper as follows. We first give a description of our near-IR imaging and spectroscopic observations of 3C 396 in § 2, together with the basic data reduction processes. In § 3 and § 4, we present the results of the [Fe II] and H 2 line observations, respectively. We discuss the origin of the [Fe II] and H 2 emission in § 5 and summarize the paper in § 6.
Observations and Data Reduction
Near-IR narrow-band imaging observations of the SNR 3C 396 were carried out to search for [Fe II] and H 2 emission using Wide-field Infrared Camera (WIRC) aboard the Palomar 5-m Hale telescope on 2005 July 15 and August 28 and 29. WIRC is equipped with a Rockwell Scientific now Teledyne Hawaii II HgCdTe 2K IR focal plane array, covering a ∼ 8.
′ 5 × 8. ′ 5 field of view with a 0. ′′ 25 pixel scale. Our WIRC field covers the entire area of 3C 396 seen in radio except for a tiny portion of the faint tail in the east (see Figure 1) . The narrow-band filters for the 1.64 µm [Fe II] line and 2.12 µm H 2 line were used, together with H-cont and K-cont filters to subtract out continuum emission. The detailed observing parameters are given in Table 1 . For the basic data reduction, the dark and sky background were subtracted out from each dithered frame using the median filtering of stacked dithered frames, and then the dithered frames were combined to a final image after flat fielding. The photometric solutions of the [Fe II] and H 2 images were obtained using the aperture photometry of 14 isolated stars from the 2MASS point source catalog (Skrutskie et al. 2006) . The obtained instrumental magnitudes of the [Fe II] and H 2 images were found to have a very good correlation (r ≃ 1) with the 2MASS H and K s magnitudes. Astrometric solutions were also obtained using the 2MASS stars. The 1-σ uncertainties of both R.A. and decl. directions in the [Fe II] and H 2 images were smaller than 0.
′′ 05 which is within the systematic accuracy of the 2MASS astrometry.
The left panels in Figures 2 and 3 present the WIRC images obtained with the [Fe II] 1.64 µm and H 2 2.12 µm narrow-band filters respectively, where we see filamentary nonstellar emission in both the [Fe II] and H 2 images which are highly contaminated by the emission from numerous stars in the field. In order to see the non-stellar emission more clearly we subtracted out the stellar emission using the continuum emission images obtained with the H-cont and K-cont filters. For this we first performed PSF photometry of the H-cont image to obtain a list of stars for continuum subtraction, and then subtracted out the corresponding stars in the [Fe II] image. To enhance the image quality we applied a median filtering with a square box of 1.
′′ 75 × 1. ′′ 75 and then smoothed the resulting image using a Gaussian with an 1 ′′ FWHM. It was difficult to remove several very bright stars with the above method, so we simply masked them out. We followed the same procedure with the H 2 and K-cont images. The right panels in Figures 2 and 3 After the aforementioned imaging observations, follow-up spectroscopic observations were carried out using the Long-slit Near-IR Spectrograph aboard the same Hale Telescope (Larkin et al. 1996) . The spectrograph is equipped with a 256 × 256 HgCdTe NICMOS detector from Rockwell International now Teledyne, providing a 38 ′′ slit length and 1 ′′ (= 6 detector pixels) slit width. Both low-resolution (R ≃ 700) and high-resolution (R ≃ 5,000) modes of the spectrograph were used for the observations. The spectral coverages of the low-resolution mode were ∼ 0.06 and 0.12 µm for the [Fe II] (1.64 µm) and H 2 (2.12 µm) lines respectively, while that of the high-resolution mode was ∼ 0.02 µm for the [Fe II] Table 2 . Both flat fielding and atmospheric opacity correction were performed by dividing the spectrum of a standard G-type star obtained at a similar airmass just after the target observations. Then the target spectrum was multiplied by a blackbody curve with an effective temperature of the standard G-type star. The OH air glow lines (Rousselot et al. 2000) were used for wavelength calibration and the Earth motion was corrected. s −1 sr −1 , respectively. Compared to the radio emission, this filamentary [Fe II] emission is confined by radio contours delineates the inner edge of a large, bright half-shell feature prominent in radio in the western part of the remnant (Figures 1 and 2 ). The radio halfshell feature is relatively thick (∼ 0.
′′ 5 or 1 pc), and the brightnesses of the both [Fe II] and radio emission drop sharply across the western boundary of this feature. The central [Fe II] emission and the emission in the west appear to be connected though low-surface brightness [ (Quinet et al. 1996 ) to 1.49 (Smith & Hartigan 2006) . From [Fe II]-pk1, we obtain [Fe II] 1.26 to 1.64 µm line ratio of 0.232 ± 0.026, which gives the visual extinction in the range of A V = 17 -21 mag based on the extinction cross section per hydrogen nuclei for the carbonaceous-silicate model of interstellar dust with R V = 3.1 (Draine 2003) . The estimated column density of hydrogen nuclei is N H = 3.4 -4.2 × 10 22 cm −2 , which is between the values obtained by X-ray and radio H I observations (Becker & Helfand 1987; Harrus & Slane 1999; Olbert et al. 2003) . We adopt N H = 3.4 × 10 22 cm −2 , which is approximately the middle value of X-ray and H I observations, in this paper. On the other hand, the intensity ratios of the two undetected [Fe II] lisarodriguez@example.com and 1.66 µm to the 1.64 µm transition is useful to estimate the upper limit of electron number density (e.g., Koo et al. 2007 ). We solve the rate equation of 16 levels of the ionized iron using the atomic parameters assembled by the CLOUDY program (ver. C05.05; Ferland 1998), and derived an upper limit of the electron number density to be 2,000 cm −3 for both [Fe II]-pk1 and [Fe II]-pk2 at assumed temperature of 5,000 K. .) The large linewidth of the averaged spectrum is partly due to a velocity shift which is apparent in the position-velocity diagram in Figure 5 (bottom) where we can identify a velocity shift from 93 to 145 km s −1 across the peak position from the northeast to the southwest direction. The measured velocity +117 km s s −1 sr −1 ) can be explained by a shock propagating into the gas of n H = 10 3 -10 4 cm −3 with a velocity of 30 -50 km s −1 (Hartigan et al. 2004 ).
High-Resolution Spectroscopy

Results of H 2 Observations

Distribution of H 2 Emission and Molecular Gas Toward 3C 396
There is only one conspicuous feature in the H 2 2.12 µm emission which is a long (∼ 3 ′ , or 7 pc) filamentary emission along the southwestern boundary of 3C 396 (Figure 3 ). The peak position (H 2 -pk1) is slightly south of the position of [Fe II]-pk2, and its surface brightness is 2.4 ± 0.2 × 10 −5 ergs cm −2 s −1 sr −1 . The total integrated flux of the H 2 2.12 µm emission is 4.1 ± 1.3 × 10 −13 erg cm −2 s −1 . Figure 6 compares the relative locations of the [Fe II] and H 2 emission, superimposed on the radio contours. It is apparent that the H 2 emission is located slightly farther than the [Fe II] emission from the SNR center, and the former is linear as opposed to the latter which is curved along the bright radio emission. The [Fe II] and H 2 filaments are adjacent near the position of [Fe II]-pk1, but are apparently separated at the edges.
Across the southwest boundary of 3C 396, the radio emission drops sharply (Figure 1) . This, together with our detection of the H 2 emission, strongly suggests that 3C 396 has encountered dense medium in this region, most likely a molecular cloud. In order to investigate this more we examined the distribution of molecular gas around 3C 396 using the Galactic Ring 13 CO J=1-0 line Survey (Jackson et al. 2006 ) and found molecular gas at v LSR = 68-70 km s −1 surrounding 3C 396 (Figure 7 ). These velocities are consistent with the previously proposed systematic velocity of 3C 396 (Caswell et al. 1975) , and the location of the H 2 filament that we found in this study matches the inner boundary of the molecular gas that forms a ring-like structure. Therefore, it is likely that 3C 396 is indeed interacting with some of these molecular clouds. If we adopt the velocity (69 km s −1 ) of the molecular gas as the systemic velocity of 3C 396, its distance is 8.5 kpc which we use in this paper as the distance to 3C 396.
Spectroscopic Observations and Molecular Shocks
Our H 2 spectrum of H 2 -pk1 averaged over a 24 ′′ area around the peak is shown in Figure 8 . The H 2 2.12 µm 1-0 S(1) line stands out clearly, while the detection of H 2 2.03 µm 1-0 S(2) line is marginal. The H 2 2.07 µm 2-1 S(3) line is not detected, which could be partly due to the incomplete OH airglow subtraction. The ratio of the H 2 2.03 µm 1-0 S(2) line to the 2.12 µm 1-0 S(1) line is 0.25 ± 0.08, giving an excitation temperature of T ex = 820 ± 350 K based on the transition probabilities of Wolniewicz et al. (1998) and an ortho to para hydrogen ratio of 3. This is lower than the general value of ∼ 2, 000 K derived from the ro-vibrational H 2 1-0 S(1) and 2-1 S(1) line ratios in thermalized molecular shock regions (Burton et al. 1989) . It may be that in the case of 3C 396 the density of the molecular gas is smaller than the critical density for thermalization, so that the excitation temperature is lower.
The dereddened surface brightness of H 2 -pk1 is 1.4 × 10 −4 ergs cm −2 s −1 sr −1 . This surface brightness can be produced by a C-shock penetrating into a molecular cloud of n H = 10 4 cm −3 and B = 50 -100 µG with a shock velocity of 20-30 km s −1 (lisarodriguez@example.com. 1983 ). Since we are observing tangentially to the shock front in 3C 396, the real density could be lower than the model value which was calculated for a normal surface brightness.
Origin of the [Fe II] and H 2 Filaments
We detected the [Fe II] 1.64 µm and H 2 2.12 µm filaments which are close to each other in the western boundary of 3C 396. While the [Fe II] filament is curved and distributed along the bright radio emission, the H 2 filament appears to be straight and is located just outside the bright radio emission. The [Fe II] and H 2 filaments are adjacent in the middle, but are apparently separated at the edges. The former is brighter than the latter.
What would be the origin of these filaments? In principle, one can observe the emission from both Fe + ions and H 2 molecules in a single SNR shock if the shock is radiative and has swept up enough column density to have reformation of molecules behind the shock (e.g., Hollenbach & McKee 1989b) . However, in such cases the H 2 emission originates in the downstream, which is opposite to the 'reversal' of the relative positions between the [Fe II] and H 2 emission observed in 3C 396 ( Figure 6 ). Interestingly 3C 396 is not the first SNR where the 'reversal' is detected. It has been detected previously in several SNRs, including RCW 103 (Oliva et al. 1989 (Oliva et al. , 1990 , G11.2−0.3 (Koo et al. 2007) , and W49B (Keohane et al. 2007 ). (The 'reversal' was also observed in W44 (Reach et al. 2005 ), but it is excluded in our discussion because it is an old SNR and also [Fe II] emission is fainter than H 2 emission, on the contrary to the other SNRs.) Table 4 summarizes the properties of these four SNRs (including 3C 396) where we list their distances, radii, and ages, along with their interactions with ambient molecular clouds. For the origin of the 'reversal' between the near-IR [Fe II] and H 2 emission, Oliva et al. (1990) considered the excitation of H 2 by ultraviolet or Xrays from a SNR, although it appears unlikely because of the difficulty in exciting the cold gas. For W49B, Keohane et al. (2007) proposed an explosion inside a wind bubble within a molecular cloud and H 2 emission being produced by a 'jet' propagating beyond the bubble; however, it is inadequate to explain all the four SNRs. Table 4 are core-collapse SNRs with a relatively small radius of 10 pc. According to Chevalier (2005) , G11.2−0.3 and RCW 103 are remnants of Type IIL/b SNe that had intermediate-mass (25-35 M ⊙ ) progenitors with strong red supergiant (RSG) winds, and the SNR shocks are interacting with the wind material extending to 5-7 pc. The radio morphology of 3C 396, together with the detection of the [Fe II] and H 2 filaments, indicates that it may also be a remnant of a Type IIL/b SN. In the radio emission (Figure 1 ), there exists a thick, half ring-like bright radio shell (BRS) in the western part of the remnant. Beyond this BRS of ∼ 2.
All the SNRs in
′ 5 (or 6 pc) radius, there appears faint plateau-like emission especially in the north and south of the remnant. We interpret this BRS as the RSG wind material confined by the ambient pressure and the faint extended plateau-like emission as originating from the SNR shock that has crossed the boundary of the wind material. This interpretation is consistent with the size of the BRS which is comparable to to the typical size of the RSG wind material confined by the ambient pressure (Chevalier 2005) . In this scenario, the [Fe II] emission may trace the shocked material of the RSG wind, while the H 2 emission may represent interactions between the SNR shock and a wind blown bubble produced during the main sequence (MS) phase of the progenitor star by its fast stellar winds (e.g., lisarodriguez@example.com. 1984b) . This also provides a natural explanation of the positional 'reversal' between the [Fe II] and H 2 emission. The H 2 filament could be either the inner boundary of the MS wind blown bubble or molecular clumps left over inside the bubble. We note that the 4.5 and 5. We suggest that the [Fe II] emission in 3C 396 represent dense clumps in the RSG wind of the progenitor star as follows. The expansion velocity of [Fe II]-pk1 is very small compared to what we expect from a young SNR. If we use the hot gas temperature of 0.62 keV derived from X-ray observations (Harrus & Slane 1999) , the implied velocity of the SNR shock in 3C 396 is in fact 600 km s −1 . This discrepancy indicates that the ambient medium may be inhomogeneous, so that the SNR shocks propagating into dense clumps are slow and radiative while the shocks propagating through the diffuse interclump medium are fast and non-radiative. Generally in a shocked gas the [Fe II] lines are emitted in the far downstream where hydrogen atoms are mostly neutral at T = 10 3 -10 4 K (lisarodriguez@example.com. 1984a; Hollenbach et al. 1989a; Oliva et al. 1989) . This is because the ionization potential of the Fe atom is low (7.9 eV) and the far-ultraviolet photons from hot shocked gas can penetrate far downstream to maintain the ionization state of Fe + . The upper limit of the electron density that we derived in § 3.2.1 gives an upper limit of the density of hydrogen nuclei n H ≈ n e /0.11 2×10 4 cm −3 if the mean ionization fraction is 0.11 (Oliva et al. 1989 ). Assuming 5,000 K as the temperature of the [Fe II] line emitting region, the thermal pressure is p/k B 1 × 10 8 cm −3 K where k B is Boltzmann constant, which is comparable to the thermal pressure of the X-ray emitting hot gas (Harrus & Slane 1999) . This indicates that the shock in the [Fe II] gas can be driven by the thermal pressure of the hot gas. The preshock density may be estimated from n 0 ≈ p/µ H v 2 s ∼ 50 cm −3 where µ H (= 2.34 × 10 −24 g) is the mean mass per hydrogen nuclei including the cosmic abundance of He and v s ≃ 56 km s −1 is the shock speed. The density of the RSG wind at r = 6 pc is n w =Ṁ/4πr 2 µ H v w ∼ 1 cm (e.g., Chevalier 2005) . Therefore, the preshock density of the [Fe II] emitting gas is an order of magnitude greater than the mean density of the RSG wind, suggesting that the former represents dense clumps in the RSG wind material.
Conclusion
Stars of intermediate-mass (25-35 M ⊙ ) have fast winds in the MS phase and slow dense winds during the RSG phase. Just before a core-collapse SN explosion, therefore, they are surrounded by a dense circumsteller region of 5-7 pc radius created by the RSG winds embedded within a low-density bubble produced by the MS winds (Chevalier 2005) . Once explodes, the SN shock encounters these pre-existing structures to generate in principle a multi-shell structure (e.g. Dwarkadas 2005) , although the details can differ from case to case depending on the distribution of the interstellar medium. In this paper, based on the detection of the near-IR [Fe II] and H 2 filaments, we propose that the SNR 3C 396 is a remnant of a Type IIL/b SN that shows such a multi-shell structure produced by the interactions of the SNR shocks with the RSG wind material (i.e., [Fe II] filament) and also with the MS wind blown bubble (i.e., H 2 filament). This scenario can provide a natural explanation of the relative locations between the [Fe II] and H 2 emission, where the latter is placed farther from the center of the remnant than the former, and is also consistent with the Spitzer 4.5 and 5.8 µm emission of 3C 396. The [Fe II] emission likely represents dense clumps driven by the thermal pressure of hot X-ray emitting gas in the remnant. We identify that there exists a molecular cloud that appears to surround 3C 396. The velocities of the molecular gas are very similar to the systematic velocity of 3C 396, indicating that this molecular gas may be indeed interacting with 3C 396 to produce the H 2 emission. Further sensitive molecular line observations are needed to investigate this possibility more. Finally, we note that there are three other SNRs having [Fe II] and H 2 emission somewhat similar to that of 3C 396. We expect that these SNRs are also remnants of Type IIL/b SNe (e.g., Chevalier 2005) that have comparable environments produced by the stellar winds of similar progenitors.
We would like to thank Kristy Dyer for providing the VLA image of 3C 396. This 

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