THE SPITZER LOCAL VOLUME LEGACY: SURVEY DESCRIPTION AND INFRARED PHOTOMETRY

New Spitzer IRAC (Fazio et al. 2004) observations were obtained for 180 LVL galaxies. The IRAC observing strategy follows that of SINGS, which shows that stellar and small grain dust emission is typically detected out to the optical radius at a surface brightness level of ∼0.01–0.1 MJy sr⁻¹ (Regan et al. 2006; Dale et al. 2000). For galaxies smaller than the IRAC field of view (D₂₅ ⩽ 300''), the Astronomical Observing Requests (AORs) were constructed using four dithered 30 s integrations. For larger galaxies, a mosaicking strategy with ∼half-array spatial offsets was used, with the sizes of the mosaic "cores" tailored to the optical size of each galaxy. Two sets of IRAC maps were obtained for each source to enable asteroid removal and enhance map sensitivity and redundancy. Combining all eight 30 s frames thus results in a net integration per pixel of 240 s (and 120 s around the ∼25 wide mosaic peripheries). Since each source was observed in all IRAC channels, ample sky coverage is automatically provided by the nonoverlapping nature of the two IRAC fields of view.

The basic calibrated data (BCD) used for post-pipeline processing are from the S18.0 and S18.5 versions of the IRAC pipeline. These versions differ from their predecessors by including improved corrections for muxbleed and the first-frame effect, among other corrections. The multiepoch, multiple-pointing IRAC observations for each galaxy are combined into one single mosaic for each band using the MOPEX mosaicking software. Additional post-BCD processing includes: distortion corrections, rotation of the individual frames (for multiepoch observations), bias structure and bias drift corrections, image offset determinations via pointing refinements from the SSC pipeline (MOPEX's default), detector artifact removal, constant-level background subtraction, and image resampling to 075 pixels using drizzling techniques. The drizzling slightly improves the final point-spread function (PSF) over the native one; the full-width half maxima are ∼16 in the shorter wavelength channels and ∼19 at 8 μm. The final images are in units of MJy sr⁻¹ and have the average sky level removed; sky values are estimated via several "blank" regions located near but beyond the target galaxy emission. Though the LVL IRAC data processing is built upon MOPEX while the SINGS project developed its own IRAC data processing package, the nature of the final data products in the two surveys is essentially the same.

In cases where exceptionally bright target sources saturated or entered the nonlinear regime of the detector during the 30 s exposure, additional 1.2 s images are used to allow for recovery of this information. Pixels affected by these issues, typically in the 5.8 and 8.0 μm frames, are flagged during processing. The correction begins by creating a mosaic of the 1.2 s exposures interpolated onto the same pixel grid as the original mosaic. A difference image is then created from the two mosaics and any residual, systematic difference in the background sky levels is removed. Pixels in the difference image valued at 1 MJy sr⁻¹ or higher are flagged (routinely regions of ∼400 contiguous pixels) and these pixels in the long integration mosaic are replaced by their short integration counterparts. Immediately outside of these saturated areas, the photometry of the 1.2 s-based mosaics is consistent with that from the far-deeper mosaics described above, and farther away in the fainter surface brightness regions the deeper mosaics obviously more effectively detect emission. The nuclear regions for the following galaxies were corrected for saturation: NGC 0253, NGC 2903, NGC 3031, NGC 3034 (at all IRAC wavelengths), NGC 3351, NGC 3593, NGC 3627, NGC 4258, NGC 5195, and NGC 5253.

New Spitzer MIPS (Rieke et al. 2004) observations were obtained for 201 LVL galaxies. Galaxies were imaged in all three MIPS bands centered at 24, 70, and 160 μm, using the highly successful scan mapping strategy employed in the SINGS project. The scan mode was used even on galaxies small enough to fit within the array field of view, because achieving adequate background measurements for extended targets in the photometry mode is less efficient than in the scan mode. Each map was executed at the medium scan rate and includes multiple scan legs tailored to the size of the galaxy and half-array offsets between scan legs. Each galaxy was mapped twice, with the maps separated by 10–40 days to allow time for the field of view to rotate and for asteroids to move out of the field. This second map was performed in the reverse direction (the "backward mapping" mode), with offsets in the cross-scan and in-scan directions. Taken together, these mapping strategies ensure that each point on the galaxy is scanned over in two different directions, which aids reduction of array artifacts on both Si:As and Ge:Ga arrays. The in-scan offset ensures that Ge:Ga stimflashes do not occur at the same point in both maps and thereby improves the calibration. The integration time per point is 160, 80, and 16 s at 24, 70, and 160 μm, respectively.

The MIPS images are processed with the MIPS Data Analysis Tool (DAT; Gordon et al. 2005), supplemented by custom scripts for the specific data reduction and mosaicking of extended sources. The latter include at 24 μm: readout offset correction, array-averaged background subtraction, and exclusion of the first five images in each scan leg due to boost frame transients. At 70 and 160 μm, the custom scripts include a pixel-dependent background subtraction for each map to remove residual detector drifts and background cirrus and zodiacal emission. This method of reduction was used for all the SINGS galaxies as well as very large galaxies (M31, M33, M101, SMC, LMC, etc.). The resulting PSFs have full-width half maxima of ∼6'', 18'', and 40'' at 24, 70, 160 μm, respectively. The pixel scales of the MIPS mosaics are 15, 45, and 90 at 24, 70, and 160 μm, respectively.

Finally, a correction for 70 μm nonlinearity effects is included in the data processing. A correction of the form

$$\begin{equation} f^{70\;{\mbox{$\mu $m}}}_{\rm true} = 0.502 \big(f^{\rm 70\;{\mbox{$\mu $m}}}_{\rm measured}\big)^{1.182}, \end{equation} \tag{ 1 }$$

derived from data presented by K. D. Gordon et al. (2009, in preparation) and slightly different from the form first presented by Dale et al. (2007) for SINGS galaxies, is applied to pixel values above a threshold of ∼44 MJy sr⁻¹. The uncertainties on the parameters in Equation (1) are ∼10%. The correction to the global 70 μm flux density is ⩽1.01 for 83% of the sample, ⩽1.05 for 90% of the sample, and ⩽1.29 for all but two sources. The correction for NGC 0253, a galaxy with a well known super star cluster, is 1.59. The correction for the starburst galaxy NGC 3034 (M82) is 1.83.

Archival IRAC and MIPS data, with spatial coverage and sensitivity similar to or greater than that described in Sections 3.1 and 3.2, are utilized for 78 (IRAC) and 57 (MIPS) galaxies. No new IRAC or MIPS observations were obtained for these subsets of the LVL sample. The data processing procedures for the archival data are the same as those followed for the new observations described above (including the use of the S18 IRAC data processing pipeline), except for the asteroid rejection in the few cases where only one epoch was measured. Table 2 indicates for which galaxies we use the archival Spitzer data.

The 2MASS obtained data for the entire sky at 1.25, 1.65, and 2.17 μm using two automated, ground-based 1.3 m telescopes (Skrutskie et al. 2006). Galaxy photometry is available from the 2MASS Extended Source Catalog for over a million galaxies and the 2MASS Large Galaxy Atlas for several hundred galaxies larger than 1' (Jarrett et al. 2003). Integrated fluxes for several LVL galaxies were adopted from the Large Galaxy Atlas, and these are generally consistent with expectations based on IRAC 3.6 and 4.5 μm fluxes and simple stellar model extrapolations of 2MASS wavelengths. However, most LVL galaxies do not appear in the Large Galaxy Atlas, and for these relatively faint systems, many of the fluxes from the Extended Source Catalog are 0.5–2 mag low based on similar extrapolations from IRAC 3.6 and 4.5 μm data. We find that when Extended Source Catalog fluxes appear unexpectedly faint, it is typically due to the comparatively small apertures used in the automated 2MASS extraction (see, for example, the fairly extreme case of UGC 08245 in Figure 4). Hence, we have independently extracted 2MASS fluxes for the vast majority of the LVL samples using the same apertures and foreground star removals used to determine IRAC and MIPS fluxes, as discussed in the following section. Figure 5 displays the ratios of our near-infrared extractions with those provided in the 2MASS Extended Source Catalog. Included in the figure are results of Kirby et al. (2008) based on deep H band imaging of nearby galaxies with the 3.9 m Anglo–Australian Telescope; Kirby et al. (2008) likewise find that the fainter sources in the Extended Source Catalog have their global fluxes underestimated. The correction factors in Figure 5 rise steeply with decreasing flux densities below 0.1 Jy (∼10 mag). The secureness of the detections below this level also drops quickly, down to the 2σ–3σ level for f_ν ≲ 0.01 Jy.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

4.2.1. Foreground Star and Background Galaxy Removal

4.2.2. Aperture Corrections

Since the IRAC flux calibration is based on point-source photometry for a 12'' radius aperture, the fluxes for all extended sources and aperture radii ≠ 12'' need to have an additional correction applied. These corrections account for the "extended" emission due to the wings of the PSF and also for the scattering of the diffuse emission across the IRAC focal plane. This photometric correction is different from merely subtracting off the sky value (Section 3.1). As described by Dale et al. (2007), the IRAC extended source correction has been derived for a variety of source morphologies and extents. For an effective aperture radius $r=\sqrt{ab}$ in arcseconds derived from the semimajor a and semiminor b ellipse axes provided in Table 1, the IRAC extended source aperture correction is

$$\begin{equation} f^{\rm IRAC}_{\rm true} \big/ f^{\rm IRAC}_{\rm measured} = A {\rm e}^{-r^B} + C, \end{equation} \tag{ 2 }$$

where A, B, and C are listed in Table 3.²⁰ The median IRAC extended source aperture corrections are [0.914, 0.941, 0.826, 0.756] at [3.6, 4.5, 5.8, 8.0] μm, respectively.

Table 3. IRAC Aperture Correction Parameters

λ	A	B	C
3.6 μm	0.82	0.370	0.910
4.5 μm	1.16	0.443	0.940
5.8 μm	1.49	0.207	0.710
8.0 μm	1.37	0.330	0.740

Notes. See Equation (2) and http://ssc.spitzer.caltech.edu/irac/calib/extcal/.

Download table as:  ASCIITypeset image

In contrast to the IRAC aperture corrections, the main reason MIPS aperture corrections are needed is the smearing of light according to the PSF profile; the measured MIPS fluxes need to be slightly boosted to account for light diffracted beyond the extent of the chosen apertures. MIPS aperture corrections are empirically determined from a comparison of fluxes from smoothed and unsmoothed 3.6 μm imaging, an approximate proxy for tracing the MIPS galaxy morphologies. The aperture correction for a given MIPS flux is the ratio of the fluxes from the unsmoothed 3.6 μm image to the flux from the 3.6 μm image smoothed to the same PSF as the MIPS image in question. The median MIPS aperture corrections are [1.01, 1.01, 1.03] at [24, 70, 160] μm, respectively, and the most significant corrections are [1.07, 1.20, 1.68] for UGC 05923.

Many of the optically faint galaxies in the sample are frequently undetected in the infrared, particularly at wavelengths of 5.8 μm and longer. Upper limits are included in Table 2 for sources undetected by infrared imaging. In all cases, "undetected" implies that the measured flux density is below the 5σ upper limit. The 5σ upper limits for Spitzer imaging are derived assuming a galaxy spans all N_pix pixels in the aperture,

$$\begin{eqnarray} f_\nu (5\sigma {\rm \;upper \;limit})_{\rm Spitzer} &=& 5 \sigma _{\rm sky} \Omega _{\rm pix} \sqrt{N_{\rm pix}+{N_{\rm pix}}^2/N_{\rm sky}}\nonumber\\ &&\hspace{-1pc}\approx 5 \sigma _{\rm sky} \Omega _{\rm pix} \sqrt{2 N_{\rm pix}}, \end{eqnarray} \tag{ 3 }$$

where σ_sky is the sky surface brightness fluctuation per pixel (MJy sr⁻¹), Ω_pix the solid angle subtended per pixel, and N_sky (≈N_pix) is the total number of pixels in the sky apertures. The parameter σ_sky is approximately 0.02, 0.03, 0.11, 0.12, 0.2, 0.9, and 1.7 MJy sr⁻¹ at 3.6, 4.5, 5.8, 8.0, 24, 70, and 160 μm, respectively, though somewhat larger values are employed for situations where the sky fluctuations are notably larger due to flatfielding errors, scattered light, cirrus, etc. A similar computation for 2MASS near-infrared upper limits is carried out after converting that survey's mean 10σ point-source sensitivities (∼16.4, 15.5, and 14.8 mag for J, H, and K_s, respectively; Skrutskie et al. 2006) to 5σ values and accounting for the difference in the sizes of the 2MASS point-source aperture (πr²_2MASS; r_2MASS = 4'') and the LVL apertures (πab). In other words,

$$\begin{eqnarray} f_\nu (5\sigma {\rm \;upper \;limit})_{\rm 2MASS} &=& {5 \over 10} f_\nu ({\rm 2MASS} \;10\sigma {\rm \;point}\nonumber\\ &&-{\rm source \;limit}) \sqrt{{\pi a b \over \pi r_{\rm 2MASS}^2}}. \end{eqnarray} \tag{ 4 }$$

The lower panels of Figure 6 display the detection rates for the different Spitzer imaging channels as a function of B band apparent and absolute magnitudes. Nearly all galaxies are detected at all Spitzer wavelengths down to m_B ≈ 14 mag and M_B ≈ −13 mag. Consistent with our presurvey expectations, the m_B ∼ 15.5 mag cutoff for the outer tier of the sample that extends to 11 Mpc (see Section 2) proved to be a useful sample selection criterion, because very few galaxies fainter than m_B ∼ 15.5 mag were detected in MIPS. The inner tier/ANGST portion of the sample extends the sample to much fainter levels, as faint as m_B ≈ 19 mag in the cases of BK03N and M81 Dwarf A. As expected, for the optically faint galaxies, the highest detection rates are found for the stellar-dominated 3.6 and 4.5 μm channels, while the 70 and 160 μm imaging proved to be far more challenging to convincingly detect cold dust emission. A stacking analysis (e.g., Dole et al. 2006) will be employed to obtain a better statistical understanding of the fainter galaxy population at long wavelengths, in particular with respect to the H i emission.

Zoom In Zoom Out Reset image size

Secure flux measurements are available at all IRAS and MIPS wavelengths for a subset of 70 LVL galaxies. The IRAS data are compiled from Rice et al. (1988), Moshir et al. (1990), Sanders et al. (2003), Lisenfeld et al. (2007), and our own archival extractions. Figure 7 provides a comparison of MIPS 24 μm and IRAS 25 μm data. The agreement between 24 and 25 μm fluxes is excellent: νf_ν(24 μm)/νf_ν(25 μm) = 1.01 with a dispersion of 25%.

Zoom In Zoom Out Reset image size

The aggregate emission from all dust grains is a fundamental metric of any galaxy. Figure 7 provides a comparison of the 3–1100 μm total infrared (TIR) for the LVL sample as measured by MIPS and IRAS. The MIPS-based total infrared is estimated from a linear combination of 24, 70, and 160 μm fluxes:

$$\begin{eqnarray} f(T\!I\!R)_{\rm MIPS}&=&1.559 \nu f_\nu (24\;\mu {\rm m}) + 0.7686 \nu f_\nu (70\;\mu {\rm m})\nonumber\\ &&+ 1.347 \nu f_\nu (160\;\mu {\rm m}), \end{eqnarray} \tag{ 5 }$$

and the IRAS-based total infrared is similarly computed from a linear combination of the 25, 60, and 100 μm fluxes:

$$\begin{eqnarray} f(T\!I\!R)_{IRAS}&=&2.403 \nu f_\nu (25\;\mu {\rm m}) - 0.2454 \nu f_\nu (60\;\mu {\rm m})\nonumber\\ && + 1.6381 \nu f_\nu (100\;\mu {\rm m}), \end{eqnarray} \tag{ 6 }$$

which are Equations (4) and (5), respectively, in Dale & Helou (2002; see Equation (22) of Draine & Li 2007 for a variation of Equation (5) above that includes the IRAC 8.0 μm flux). The coefficients in the above two equations stem from fits to a suite of spectral templates applicable to a wide range of normal star-forming galaxies at redshift zero, where "normal" implies the exclusion of AGN and ultraluminous infrared galaxies (see Section 5.3 and Figure 5 of Dale & Helou 2002 for a representative sampling of the suite of templates). The uncertainty in using these prescriptions to compute the total infrared for normal star-forming galaxies is estimated to be of order 25% (Draine & Li 2007).

The MIPS-based version should be more accurate since the infrared wavelength baseline spanned by MIPS is longer than the baseline covered by IRAS, and more importantly, the IRAS detectors do not sample the bulk of the dust in the coldest, most quiescent galaxies. To determine if these differences in wavelength coverage between IRAS and MIPS result in different estimates of the total infrared, the two right-hand-side panels in Figure 7 compare f(TIR)_MIPS and f(TIR)_IRAS. The ratio of MIPS- and IRAS-based total infrared measures has a scatter (21%) similar to that in the 24-to-25 μm comparison, but the average ratio is 1.15. These findings are similar to those of Kennicutt et al. (2009) for a sample of 205 nearby galaxies with both IRAS and MIPS data. The right-hand-side panel in Figure 7 includes semiempirical predictions from models of infrared spectral energy distributions. As alluded to above, part of the systematic offset in TIR_MIPS/TIR_IRAS can be attributed to the relative inability of IRAS to accurately measure the total infrared for cold galaxies. The infrared emission for the coldest galaxies, galaxies with the lowest f_ν(60 μm)/f_ν(100 μm) ratios, peaks beyond IRAS's 100 μm detector, and thus the total infrared as measured by IRAS is systematically low for the coldest galaxies.

Figures 8 and 9 show ultraviolet–Hα-infrared mosaics of NGC 5236 and UGC 05829, spanning wavelengths where the emission is dominated by young stars (0.15 μm), H ii regions (Hα), old stars (3.6 μm), PAHs (8.0 μm), very small grains (24 μm), and large grains (70 μm). The galaxies and wavelengths displayed in these two figures highlight the broad range of environments and galaxies sampled by the LVL survey (see Section 2). Figure 10 provides the panchromatic ultraviolet–infrared broadband spectral energy distributions for all 258 galaxies.²¹ The solid curve is the sum of a dust (dashed) and a stellar (dotted) model. The dust curve is a Dale & Helou (2002) model (least squares) fitted to ratios of the observed 24, 70, and 160 μm fluxes, and then scaled to match the overall infrared brightness. The α_SED listed within each panel parameterizes the distribution of dust mass as a function of heating intensity U in units of the local ultraviolet interstellar radiation field, as described by Dale & Helou (2002):

$$\begin{eqnarray} dM_{\rm dust}(U) &\propto & U^{-\alpha _{\rm SED}} \; dU, \; \; \; \; \; \; 0.3 \le U \le 10^5. \end{eqnarray} \tag{ 7 }$$

To quantify the uncertainty on α_SED displayed within each panel of Figure 10, 1000 Monte Carlo simulations of the fit to each galaxy's far-infrared fluxes were performed, utilizing the tabulated flux uncertainties to add a random (Gaussian deviate) flux offset at each MIPS wavelength. The α_SED uncertainties reflect the standard deviations in the simulations. The stellar curve is a 1 Gyr continuous star formation, solar metallicity curve from Vazquez & Leitherer (2005) fitted to the 2MASS data. The initial mass function for this curve utilizes a double power-law form, with α_1,IMF = 1.3 for 0.1 < m/M_☉ < 0.5 and α_2,IMF = 2.3 for 0.5 < m/M_☉ < 100 (e.g., Kroupa 2002). Though this stellar curve is not adjusted for internal extinction and may not be applicable to many galaxies in the sample, it is included as a fiducial reference against which deviations in the ultraviolet can be compared from galaxy to galaxy.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

The spectral energy distributions for the LVL sample range widely. There are stellar-dominated (NGC 0404, UGC 05373, UGCA 0193) and comparatively dusty (IC 5256, NGC 6503) systems; for sources detected by MIPS, the infrared-to-far-ultraviolet ratio in the sample spans more than three orders of magnitude, from ≲0.1 to over 100 (Section 5.5). There are galaxies with far-infrared spectral energy distributions indicative of warm (UGCA 0281) and cold dust grains (NGC 5055). Compared with what would be expected based on their stellar and far-infrared emission, many galaxies show a dearth of emission from PAHs in the 8.0 μm band (e.g., ESO 245-G005, UGC 01249, UGC 05272). The variations in global spectral energy distributions are discussed in more detail below.

The IRAC–MIPS infrared colors for the LVL sample are displayed in Panel (a) of Figure 11. The f_ν(70 μm)/f_ν(160 μm) ratio typically traces the temperature of large interstellar grains, while the f_ν(8.0 μm)/f_ν(24 μm) ratio has several influences. The flux at 24 μm mostly represents emission from very small grains (grains with effective radii of 15–40 Å; Draine & Li 2007), and the flux at 8.0 μm can have contributions from stars, hot dust, PAHs, and AGNs. Perhaps, due to the diversity of emission mechanisms responsible for 8.0 and 24 μm flux levels, the f_ν(8.0 μm)/f_ν(24 μm) ratio spans nearly two orders of magnitude compared with the single factor of ∼10 stretch in the f_ν(70 μm)/f_ν(160 μm) ratio. Since the local volume lacks "strong" AGN, loosely defined here as AGN that dominates a galaxy's emission over substantial portions of the electromagnetic spectrum, it is unlikely that AGN contribute much to the scatter in Figure 11.