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The Mid-Infrared Colors of the ISM and Extended Sources at the Galactic Center

A mid-infrared (3.6-8 um) survey of the Galactic Center has been carried out with the IRAC instrument on the Spitzer Space Telescope. This survey covers the central 2x1.4 degree (~280x200 pc) of the Galaxy. At 3.6 and 4.5 um the emission is dominated by stellar sources, the fainter ones merging into an unresolved background. At 5.8 and 8 um the stellar sources are fainter, and large-scale diffuse emission from the ISM of the Galaxy's central molecular zone becomes prominent. The survey reveals that the 8 to 5.8 um color of the ISM emission is highly uniform across the surveyed region. This uniform color is consistent with a flat extinction law and emission from polycyclic aromatic hydrocarbons (PAHs). Models indicate that this broadband color should not be expected to change if the incident radiation field heating the dust and PAHs is<10^4 times that of the solar neighborhood. The few regions with unusually red emission are areas where the PAHs are underabundant and the radiation field is locally strong enough to heat large dust grains to produce significant 8 um emission. These red regions include compact H II regions, Sgr B1, and wider regions around the Arches and Quintuplet Clusters. In these regions the radiation field is>10^4 times that of the solar neighborhood. Other regions of very red emission indicate cases where thick dust clouds obscure deeply embedded objects or very early stages of star formation.

µm emission. These red regions include compact H II regions, Sgr B1, and wider regions around the Arches and Quintuplet Clusters. In these regions the radiation field is 10 4 times that of the solar neighborhood. Other regions of very red emission indicate cases where thick dust clouds obscure deeply embedded objects or very early stages of star formation.

Introduction
The Galactic Center (GC) is naturally a region of great interest. The gravitational center of the Galaxy is marked by the supermassive black hole at the radio source Sgr A*. Remarkably, two young, dense stellar disks have been discovered within 0.5 pc orbiting Sgr A* (Paumard et al. 2006). Further out is a molecular ring or the "circumnuclear disk" enclosing an ionized a mini-spiral structure (e.g. Christopher et al. 2005). These features are within a radius of ∼ 7 pc. A concentration of molecular material is observed out to a radius of ∼ 200 pc along the Galactic Plane; this region is known as the Central Molecular Zone (CMZ). Aside from the stellar bulge of the Galaxy, the CMZ encompasses nearly all of the Galactic structures and objects that are characteristic of, or unique to, the GC.
The CMZ is estimated to contain about 10% of the Galaxy's molecular gas content (e.g. Morris & Serabyn 1996;Mezger 1978), or about 5 − 10 × 10 7 M ⊙ (e.g. Armstrong & Barrett 1985;Stark et al. 1989). It has a two-component structure. The inner region consists of a set of low velocity clouds lying on or close to the Galactic plane, including the giant Sagittarius molecular clouds, which together comprise the "disk population" (Morris & Serabyn 1996). The outer set of clouds form a comparatively high velocity (130-200 km s −1 ) quasi-continuous ring structure that surrounds the first one at a radius of ∼ 180 pc (the so-called "180-pc molecular ring"; Morris & Serabyn 1996). Analysis of NH 3 and H + 3 absorption lines in the CMZ has indicated that the molecular gas contains a widespread, relatively hot component (T 175K, Wilson et al. 1982;Oka et al. 2005), which is not found apart from dense molecular cores elsewhere in the Galaxy. Analysis by Hüttemeister et al. (1993) indicates the molecular gas has at least two phases: hot phase (T ∼ 200K, n H 2 10 3 cm −3 ) with a filling factor of ∼ 25%, and a cooler component (T ∼ 25K, n H 2 10 4 cm −3 ) which fills the remaining 75% of the volume of the molecular clouds.
The average dust temperature in the CMZ appears to be much lower than that of the gas, T dust ∼ 25 − 30 K (Odenwald & Fazio 1984;Cox & Laureijs 1989). The much higher gas temperatures therefore indicate the presence of some efficient gas heating mechanisms, and supersonic winds and turbulence had been suspected; the large gas line widths, 15-50 km s −1 , are consistent with such ideas, although the details of the mechanisms are debated (SNRs, outflows, viscous heating effects, etc.). Observations of widespread SiO emission indicates the recent or current presence of strong shocks in the CMZ (Martin-Pintado et al. 1997;Hüttemeister et al. 1998). The discrepancy between dust and gas temperatures has recently been accounted for in terms of the impact of cosmic ray electrons heating molecular clouds directly and raising their temperatures (Yusef-Zadeh et al. 2007b).
In addition to the molecular gas, the CMZ is also infused with neutral and ionized components as revealed by enhanced infrared continuum and line emission (e.g. Rodríguez-Fernández et al. 2004, and by enhanced radio continuum emission (e.g. Altenhoff et al. 1979).
The launch of the Spitzer Space Telescope (Werner et al. 2004) provides a new opportunity to investigate the GC at mid-to far-infrared wavelengths with unprecedented sensitivity and spatial resolution. Spitzer's Infrared Array Camera, IRAC,  provides 5 ′ field of view imaging capabilities with < 2 ′′ spatial resolution at 4 broadband mid-infrared channels nominally at 3.6, 4.5, 5.8, and 8 µm. We used IRAC to carry out a 2 • × 1. • 4 (∼ 280 × 200 pc) survey of the GC region, which includes the entire CMZ. An initial summary of the results of this survey has been presented by Stolovy et al. (2006). A more extensive overview will be provided by S. Stolovy et al. (2008, in preparation). The measurement, cataloging, and characterization of point sources observed in our IRAC GC survey are described by Ramírez et al. (2008). In IRAC's 3.6 and 4.5 µm channels, the GC is dominated by stellar emission, both for resolved sources and in the unresolved background. At 5.8 and 8 µm wavelengths, most point sources become significantly fainter, and a strong, structured background produced by the diffuse ISM is dominant. In this study, we focus on the diffuse ISM as revealed by IRAC. In particular we investigate what the broadband colors indicate about the nature of the dust producing the emission, and the environment in which that dust is found. Section 2 of this paper contains a brief summary of the IRAC observations. Section 3 presents the colors observed for the diffuse emission and for a few, more compact sources of special interest. Mostly this focuses on the 5.8 to 8 µm color of the ISM emission. The shorter wavelength colors are dominated by starlight and extinction (reddening) effects. Section 4 discusses the implications of the uniformity of the color on large scales and the few exceptional regions on small scales. Though limited to the single 5.8 to 8 µm color, our analysis yields several general robust conclusions concerning the relative abundance of PAHs, the strength of the interstellar radiation field, and the shape of the extinction law in the GC. The last section ( §5) contains a summary of this paper.

Observations
We observed the central region of the Galaxy using 12 Astronomical Observation Requests (AORs) which mapped the region in ∼ 0. • 25 wide strips at roughly constant declination (a restriction due to Spitzer's pointing constraints and the low ecliptic latitude of the field). Most strips were ∼ 1. • 75 long, except that the one containing Sgr A was split into two parts (to mitigate the potential effects of latent images, with the brightest regions of the survey scheduled last). Additionally, those strips at the highest and lowest declinations were shorter, in order to limit the area of complete coverage to |l| < 1. • 0 and |b| < 0. • 7. The total area covered in all four channels was ∼ 3.5 square degrees, including the additional irregular edges of the region. The medium scale 5-point Gaussian dither pattern was used. This pattern makes small offsets to the telescope pointing with an average offset of ∼ 30 ′′ . We did not use the smaller scale pattern because it would have provided insufficient dithering for mitigating the stray light artifacts of bright stars, and would have hampered self-calibration of the data (see the Appendix). Frame times of 2 seconds were used throughout the survey. No regions of extended emission were saturated except at Sgr A at 8 µm, but more than 10 4 bright point sources were saturated in at least one channel (Ramírez et al. 2008). Most saturated sources are isolated point sources, but high densities of saturated sources are found within the bright cluster at Sgr A and the Quintuplet Cluster. These areas, and a sample of 12 other saturated sources, were re-observed in subarray mode with 0.02-second frame times to obtain unsaturated data. All analysis in this work was performed on mosaicked images that incorporated the subarray data for these sources.
We constructed our own mosaics from the total set of Basic Calibrated Data (BCD, the standard Spitzer data products intended for detailed scientific analysis) from all AORs. The primary description of our processing is given by Stolovy et al. (2006Stolovy et al. ( , 2008. In the appendix of the present paper, we provide additional detail concerning aspects of the processing that are particularly relevant in studying the brightness and color of extended emission. The 8 µm mosaic of the Galactic Center region is shown in Figure 1. The general structure of the ISM emission in this region appears similar to other regions in the Galactic plane as revealed by the GLIMPSE survey (Benjamin et al. 2003). Both in the GC and other regions, much of the ISM emission is resolved into elongated arcs and filamentary structures (S. Stolovy et al. 2008, in preparation). Embedded among these are very bright compact regions that contain specific local heating sources. Superimposed infrared dark clouds (IRDCs) have a similar structure to the emission, but are not well-correlated with the emission. The ISM emission has a strong unresolved background component, which is often suppressed in depictions meant to highlight the structure of the emission (such as Fig. 1). The most fundamental difference between emission in the GC and that in other regions of the Galactic plane, is that the GC emission is several times brighter. Figure 2 shows a wider view of the Galactic Center region as seen by SPIRIT III on the MSX mission (Price et al. 2001), revealing that the emission within our survey region is brighter than that typically found in the inner galaxy (at |l| > 1. • 5). Figure 3 provides a more quantitative illustration of the relative brightness of our survey region. The intensity enhancement is loosely correlated with the large-scale distribution of molecular emission in the galactic center region (e.g. Martin et al. 2004). This indicates that much of the low-latitude (|b| < 0. • 2) structure in the galactic center mosaic really is in the central ∼ 150 pc radius of the Galaxy, but that the emission does not directly arise from the molecular component of the CMZ.
In contrast, emission at higher latitudes (|b| > 0. • 2) often appears to be more local features superimposed along the line of sight. One indication of this is association with Hα emission, which would be too heavily extincted to be visible at the distance to the GC, with the possible exception of objects seen at the extreme latitude limits of our survey (Stolovy et al. , 2008. Another is the association with very dark clouds, which must lie relatively close to the sun. Otherwise, the general foreground emission along the line of sight would reduce the contrast of these features. The large angular scale and presence of fine detail in the higher latitude features also suggest they are relatively nearby clouds along the lines of sight. Pauls & Mezger (1975) use the velocity and line widths of 6 cm radio recombination lines to distinguish emission from the GC (large velocities and/or broad lines), from more local emission along the line of sight (velocities near zero, and narrow line widths). All of the regions they examined at |b| > 0. • 16 are characterized as being along the line of sight, rather than at the GC. Figure 4 shows the ratio of 8 µm / 5.8 µm emission, i.e. the I(8µm)/I(5.8µm) color. In this ratio virtually all of the diffuse emission features disappear, indicating that they have a very uniform color that is independent of brightness. Dark regions in the ratio map indicate bluer colors associated with numerous individually resolved stars, and with the unresolved stellar background which becomes comparatively bright in the innermost regions of the galaxy (central ∼ 10 ′ ). There are only a few unusually red point sources. Two of these are asteroids (1567 Alikoski and 459 Signe), and two others are PN M1-26 (HD 316248; Sahai & Trauger 1998) and the candidate PN 359. 491+0.316 (IRAS 17399-2910;Becker et al. 1994). Brighter (i.e. redder) extended regions in the ratio map are associated with a few special types of emission features which will be discussed below, and with the darkest of the IRDCs [e.g. the complex near (l, b) = (0. • 24, −0. • 45)]. The apparent redness of the IRDCs may be an artifact of residual errors in the subtraction of the zodiacal light or instrumental offsets (see the Appendix and Table 1). Small offset errors would have the greatest impact on the faintest regions of the images, which are the IRDCs and regions at the extreme latitudes along the edges of the mosaics.

ISM Colors Corrected for the Effect of Unresolved Starlight
To better isolate the ISM emission we have subtracted stellar emission in a manner similar to that applied to external galaxies (Pahre et al. 2004). In effect, we decompose the 8 and 5.8 µm emission into ISM and stellar components using the 4.5 µm emission as strong tracer of the starlight. We assume that there exist fixed spatial templates of each component, such that where I(λ) is the observed emission at λ = 4.5, 5.8 or 8 µm, I ISM and I stellar are fixed (but unknown) templates of the ISM and stellar emission, and α λ and β λ are the scaling factors for the respective templates. The ratios of the α λ or β λ coefficients yield the colors of the ISM and stellar emission at the relevant wavelengths. If the ISM emission were negligible at 4.5 µm and stellar emission were negligible at 8 µm, then α 4.5µm = β 8µm = 0 and a linear least squares fit of the form I(5.8µm) = A I(8µm) + B I(4.5µm) + C could be used to determine: A = α 5.8µm /α 8µm , the mean 5.8/8µm color of the ISM; B = β 5.8µm /β 4.5µm , the mean 5.8/4.5µm color of the starlight; and C, a constant term representing a linear combination of the mean systematic errors (including those of the detector offsets, zodiacal light, and the stellar model) in all three channels. However, because there is no a priori knowledge of the magnitude or sign of the systematic errors in each channel, we attribute this term to only the 5.8 µm channel, which exhibits the greatest instability in detector offset.
In reality, the longer wavelength stellar emission and shorter wavelength ISM emission are not negligible. Therefore the interpretation of the A, B, and C coefficients derived in Eq. 2 is slightly more complicated. To derive both β 5.8µm /β 4.5µm and β 8µm /β 4.5µm from A and B, we must assume the color of the starlight at one pair of wavelengths. We choose to set the same [5.8] − [8.0] = +0.04 mag theoretical color that was adopted by Pahre et al (2004), as these colors should be least affected by extinction. Then we can appropriately scale and subtract I(4.5µm) from the 5.8 and 8 µm emission to remove the mean stellar emission at each of these wavelength. The derived [4.5] − [5.8] and [4.5] − [8.0] colors are about 0.1 mag redder than those used by Pahre et al. (2004). Because I(4.5µm) also contains a weak I ISM component, after this subtraction of the starlight the ISM emission is slightly underestimated. The factor by which the ISM is reduced can be calculated from A and B, if we again assume the color of the ISM emission at one set of wavelengths. We use the "BARE-GR-FG" interstellar dust model of Zubko, Dwek, & Arendt (2004) to calculate that α 4.5µm /α 5.8µm = 0.194. From this we find that using I(4.5µm) as a template of the starlight leads to underestimating the ISM emission by factors of 0.85 and 0.97 at 5.8 and 8 µm respectively. To summarize we construct: and The 8/5.8 µm ISM colors found after subtraction of the starlight (and offset error) are slightly redder and no longer differ at regions of high extinction, but otherwise appear quite uniform over a wide range of brightness. This is depicted in Figure 5, where the color is plotted as a function of the 8 µm brightness. The colors shown here are those of the integrated emission along the line of sight. The shading of this figure is scaled to the log of the number of pixels with a given color -brightness combination. This logarithmic scaling enhances the visibility of the outliers in the color. Some of the unusually red extended regions that stand out in the I(8µm)/I(5.8µm) colors (both before and after subtraction of starlight) are highlighted and are discussed in §3.2. The median color (±1σ standard deviation) is I ISM (8µm)/I ISM (5.8µm) = 2.71 ± 0.19, independent of the 8 µm brightness. The locations of bright stellar sources have been excluded from this plot with an intensity threshold cut at 70 MJy sr −1 at 4.5 µm and by excluding regions with I(8µm)/I(5.8µm) < 2.3 (before the subtraction of starlight), but remaining faint resolved stars and systematic errors in using the 4.5 µm emission as the template cause much of the dispersion toward redder colors at low intensities (<50 MJy sr −1 ). The observed I ISM (8µm)/I ISM (5.8µm) colors in the GC are similar to (though slightly bluer than) the colors observed by Flagey et al. (2006) in 6 regions at longitudes l = 27. • 5, 105. • 6, and 254. • 4. (These latter colors are indicated by the arrows at the left edge of Figure 5.) The difference is established using a two-tailed Kolmogorov-Smirnov test, which yields a maximum deviation between the cumulative distribution function of D = 0.58. For the given number of data points, if the two distributions are intrinsically the same, the probability of D exceeding this value is 0.02. Therefore, this test finds only a 2% chance that the Flagey et al. (2006) results are consistent with a random sampling of the colors that we measure near the GC. However, while this statistic indicates the colors are likely to be different, we cannot rule out that procedural differences or systematic errors are the cause of the apparent differences.

Shorter Wavelength Colors
We are unable to make an analogous comparison to I(8µm)/I(3.6µm) and I(8µm)/I(4.5µm) colors measured by Flagey et al. (2006) because the very bright and confusing unresolved stellar background toward the GC overwhelms nearly all diffuse ISM emission at the shorter wavelengths. Point source-subtracted images (see Ramírez et al. 2008) provide some aid in distinguishing non-stellar 3.6 and 4.5 µm emission, but all such regions appear to be H II regions and PDRs associated with specific heating sources, rather than diffuse ISM heated by the interstellar radiation field (ISRF).

Unusual Emission Regions
Along the Galactic plane, there are a few large ( 4 ′ ) emission features that stand out in the I(8µm)/I(5.8µm) color map with unusually red colors. The largest of these is an oblong shell-like structure centered on the Arches cluster which contains ∼200 stars with masses > 20 M ⊙ (Figer et al. 2002). The shell is much more clearly delineated by the color of its emission, than by the 8 µm surface brightness ( Figure 6, top row). On the north, this shell seems bounded by the innermost of the thermal radio arches surrounding the cluster. The shell is only faintly discernible in the shorter wavelength colors. A similar but smaller region appears near the Quintuplet cluster ( Figure 6, middle row), which is comprised of a cluster of hot O/B and WR stars responsible for ionizing the nearby ISM (Cotera et al. 1996). However, the five cocoon stars (Moneti at al. 2001) that give this cluster its name are so bright at IRAC wavelengths that artifacts from these sources obscure much of this excess.
There are weaker indications of unusually red emission in the Sgr B1 region as well ( Figure 6, bottom row). The redder infrared emission is associated with 5.8 and 8 µm counterparts of the "ionized bar" and "ionized rim" features identified in 3.6, 6 and 20 cm radio observations by Mehringer et al. (1992). Unlike the Arches and Quintuplet regions, however, the most luminous stars here are not grouped in a tight easily-identified cluster, but are apparently dispersed throughout the volume.
All these extended red regions (the vicinities of the Arches and Quintuplet Clusters and Sgr B1) can be distinguished in Figure 5 as the cluster of outlier points at 250 < I(8µm) < 1000 MJy sr −1 and I ISM (8µm)/I ISM (5.8µm) > 3.2. We expect that as with the compact H II regions, the excessively red emission is caused by thermal emission from very warm dust in strong radiation fields. However, in these cases the excess emission appears over a wider region because of the larger number of luminous stars in these clusters and lower column densities around these more evolved systems allowing a strong radiation field to propagate over a larger volume.
The Pistol Nebula on the edge of the Quintuplet Cluster is a different type of red emission region. The 8 µm IRAC data reveal an entire shell or bubble around the Pistol Star, with a diameter of ∼ 22 ′′ (almost 1 pc at the distance of the GC). The brighter portion of this bubble is likely being ionized by the hot stars in Quintuplet cluster . The Pistol Nebula exhibits the most extremely red I ISM (8µm)/I ISM (5.8µm) color in the region, as it is only marginally distinguishable at 5.8 µm. (The 5.8 µm emission is very difficult to quantify because of strong banding artifacts from the extremely bright stars of the Quintuplet Cluster.) Physically, the emission from this region is associated with the ejecta from single star (the Pistol Star) which is an LBV (Figer et al. 1998). In this case, ISO spectroscopy of the nebula already suggests that the 8 µm excess here is due to a relatively red continuum (with deep silicate extinction at 10 µm) and lack of PAH emission (Moneti et al. 1999). The [Ar II] 6.99 µm is the strongest line between 5 and 10 µm, but it it not bright enough to be the dominant cause for the red color. There are no other similar (bright, sharply defined, and very red) circumstellar bubbles evident in our galactic center survey. Two partial shells around massive stars near the GC have been identified by Mauerhan et al. (2007). Neither of these shells exhibit especially red I ISM (8µm)/I ISM (5.8µm) color, although in one case red emission is present interior to the shell.
There are only two extended (∼ 20") emission sources > 0.1 • away from the Galactic plane that show strong 8 µm excesses (Figure 7). Unlike the bright compact H II regions, these sources are not very bright at 8 µm, and thus even fainter or absent at 5.8 µm. However, both brighten rapidly at longer wavelengths and were detected as IRAS point sources at 25 µm (IRAS 17456-2850 and IRAS 17425-2920) and are visible in the MSX band E (21.3 µm) images. Both are located near (in projection) extended star forming regions (RCW 137 in the case of IRAS 17425-2920) with strong cirrus emission and IRDCs. Both are slightly extended and bowed or arced. Neither source contains an obvious point source, nor has a known energizing star in the crowded field. IRAS 17425-2920 does not appear in the 1616 GHz radio maps of the Sgr C region (Liszt & Spiker 1995). The characteristics and the locations of these sources on the outskirts of active star forming regions suggest that these sources represent a very early stage of star formation. Alternately, they could represent local heating by a fully-formed star that is interacting with an otherwise quiescent molecular or neutral cloud. A third option is that these objects indicate some sort of shock phenomenon and the unusual colors are the result of strong line emission. For example, the candidate supernova remnant SSTGFLS J222557+601148 exhibits only line emission with no continuum (Morris et al. 2006). This particular object is bright in the MIPS 24 µm band by virtue of a strong [O IV] 26 µm line, but invisible to IRAC. In the Cas A SNR, [Ar II] at 6.99 µm can contribute 40% of the emission in the IRAC 8 µm band (Ennis et al. 2006). Outside of our survey area at l = 7 • , Rho et al. (2006) detect a similar source in the IRAC and MIPS observations of the Trifid Nebula ). The region they call "Trifid Junior" lies the north of the H II region and the reflection nebula. It is a small arc of emission at 24 µm which can only barely be seen at 8 µm and is absent at 5.8 µm. The general appearance of the local (∼ 5 ′ ) vicinity is very similar to that of IRAS 17425-2920.

Molecular Emission: A Dearth of 4.5 µm Excess Sources
Molecular outflows from young stellar objects are often dramatically highlighted in IRAC images because of CO and H 2 emission lines that are especially strong in the 4.5 µm IRAC band (e.g. Smith et al. 2006). This wavelength is usually depicted as green in multi-color images, and such outflows are therefore recognized by their green colors in such images. While some star forming regions can contain dozens of outflow sources (e.g. Orion, DR21, NGC 1333), there is only a single small source that is strong at 4.5 µm and has a jet-like morphology within our GC survey (Figure 8). There are dozens of other sources that also have 4.5 µm excesses, but these are unresolved or only slightly diffuse, and usually within or near more extended regions of bright 8 µm emission. Many of these may be outflows from young stellar objects that are too distant or intrinsically small to be resolved. Further discussion of these sources can be found in Yusef-Zadeh et al. (2007a).
In several of the large, more local star-forming regions at higher Galactic latitudes there are broad diffuse regions of relatively enhanced 4.5 µm emission [e.g. near (l, b) = (0. • 48, −0. • 67)]. Rather than being an indication of molecular emission, these may be cases where we spatially resolve an H II region from its surrounding PDR. The recent study of M 17 by Povich et al. (2007) shows that the absence of PAH emission within the H II region where a strong EUV radiation field is present is well correlated with an apparent excess at 4.5 µm. Thus the excess is really only a decrease in the strength of the PAH features that dominate the 5.8 and 8 µm emission (and to a lesser degree 3.6 µm emission) in the surrounding PDR.

Unseen Objects
There are 6 known SNRs (Sgr A East, G0.3+0.0, G0.9+0.1, G1.0-0.1, G359.0-0.9 and G359.1-0.5; Green 2006) that lie at least partially within our survey regions, but at 3.6 -8 µm none of these are visible by either their intensity or color. This is not unexpected given the high degree of confusion and the fact that SNRs are often difficult to distinguish in the infrared even with 2 ′′ spatial resolution. A thorough examination of the IRAC GLIMPSE survey (10 • < |l| < 65 • ) by Reach et al. (2006) found IRAC counterparts for only 18 of 95 SNRs. Those that were detected are undergoing strong interactions with dense, often molecular, clouds. The lack of detections in our survey means that none of the 6 SNRs toward the GC is undergoing a strong enough interaction to be visible above the brightest background confusion in the Galaxy. The southern portion of IC 443 is a well-studied example of a SNR / molecular cloud interaction (e.g. Cesarsky et al. 1999). Emission with the same surface brightness as IC 443 (AORID = 4422912 and 4423168) could be distinguished at higher latitudes, but would be lost in the confusion at |b| 0. • 2. Fainter examples of shocked molecular emission such as W44 ) would likely be lost in the confusion at any location within our survey.
The enigmatic infrared source known as AFGL 5376 does not show any distinct emission at any IRAC wavelength, despite being a bright extended source at ∼ 24 µm wavelengths (Uchida et al. 1990(Uchida et al. , 1994Law 2007). We do note that there are several small IRDCs and filaments (in the vicinity of (l, b) = (0. • 7244, −0. • 4759)) that can been seen in extinction against the 24 µm emission of AFGL 5376 as well as the general 3.6 -8 µm background.
There is also no particular correlation between the diffuse infrared emission and the nonthermal radio filament and threads that are present in the Galactic center region (Yusef-Zadeh et al. 2004;Stolovy et al. 2006Stolovy et al. , 2008.

Discussion
Overall, the GC exhibits homogenous IRAC colors that are quite similar to those of other regions of the Galactic plane. Such a uniform I ISM (8µm)/I ISM (5.8µm) color for the diffuse emission indicates that there is little reddening associated with extinction at these wavelengths, despite evidence for highly variable extinction in the form of IRDCs and extinction measurements made at shorter wavelengths (e.g. Gosling et al. 2006). This is expected given the flattening of Galactic extinction curve (Lutz et al. 1996;Indebetouw et al. 2005;Flaherty et al. 2007). Even if the extinction did produce reddening, it would be extremely unlikely for the line of sight reddening to be correlated with variation in the intrinsic color of the emission to produce such an overall uniform color.
The uniform colors are also evidence that the emission spectrum is largely independent of the local and global locations and environments. This strongly suggests that emission from polycyclic aromatic hydrocarbons (PAHs) dominates the diffuse ISM emission (5.8-8 µm), all the way to the Galactic Center. The relative strengths of PAH emission bands that fall within the IRAC 5.8 and 8 µm filter bandpasses are unaffected by the strength of the radiation field heating the grains over a wide range of intensities (Li & Draine 2001). Even as the interstellar radiation field strength increases toward the GC, there is no evident change in the color of the PAH emission. Figure 9 shows model spectra of the ISM emissivity normalized to the strength of the interstellar radiation field (Mathis et al. 1983), as characterized by a scale factor χ. The dust model used here is composed of PAHs and bare graphite and silicate grains (the "BARE-GR-FG" model of Zubko et al. 2004). This model is similar to those of Li & Draine (2001), but is derived via additional constraints provided by elemental abundances.
For the standard ISRF (χ = 1), most emission at λ < 60 µm is from very small grains ( 0.01 µm) and PAHs. These grains have a low heat capacity and are heated to high temperatures by absorption of a single photon. They radiate that energy quickly and cool below equilibrium temperatures before the next energetic photon hits. Thus, for these stochastically heated particles, the grains do not get any warmer as the radiation field increases; they only go through the same temperature spikes more frequently. Therefore, the spectrum normalized by χ remains roughly constant. This is especially evident at λ < 10 µm. Larger grains ( 0.01 µm) have higher heat capacities and absorb photons frequently enough that they radiate at their equilibrium temperatures. For the local ISRF (χ = 1), this thermal emission peaks at ∼ 150 µm. As χ increases, the larger grains get warmer and the thermal emission moves to shorter wavelengths. Eventually, when χ 10 4 , the thermal emission begins to enter the 8 µm IRAC band. At this point the IRAC colors can be affected.
In Figure 10 we plot I ISM (8µm)/I ISM (5.8µm) predicted by the models of a function of radiation field strength. The figure shows the "BARE-GR-FG" model is consistent with the observed I ISM (8µm)/I ISM (5.8µm) mean color as long as χ 10 4 . It is also evident that while the color initially reddens as the thermal emission from the hot dust enters the 8 µm band, the color begins getting bluer when the thermal emission pushes into the 5.8 µm band as well. However, the reddest colors attained by this model (at χ ∼ 10 5 ) are not nearly red enough to explain the unusually red regions that are we identify in the GC region ( §3.2). Alternatively, if the PAH component is removed from the "BARE-GR-FG" model, then at χ ∼ 10 5 the color of the emission can be sufficiently red to match the observations. At χ 10 4 , the colors of the emission would be much bluer than observed if PAHs were generally absent from the ISM in the GC. Qualitatively similar results are obtained from the Draine & Li (2007) dust models, which include a parameter (q P AH ) that characterizes the relative abundance of PAHs.
The red shell around the Arches Cluster provides a nice example of a region with a locally strong radiation field. Figer et al. (2002) have estimated that luminosity of the Arches Cluster as 10 7.8 L ⊙ . The extended red shell has a radius of ∼ 100" or 3.9 pc. Thus the energy density at the outer edge of this shell is ∼ 4.4 × 10 −9 erg cm −3 . Compared to the nominal strength of the local ISRF, 7 × 10 −13 erg cm −3 (Spitzer 1978), we find that χ ∼ 6300 at the edge of the shell. This value of χ indicates that we should indeed expect to see emission with red I ISM (8µm)/I ISM (5.8µm) colors in the vicinity of the Arches Cluster, and that the region of red emission cannot be much larger than what is observed.
The red emission region near the Quintuplet Cluster is related to a similarly enhanced radiation field. This region corresponds to the portion of the "Sickle" that is nearest to the Quintuplet Cluster. Spectroscopy of this region by Simpson et al. (1997Simpson et al. ( , 2007 reveal gas in a state of high excitation, consistent with radiation fields of χ > 10 3 . Calculations by Rodríguez-Fernández et al. (2001) indicate that the Quintuplet Cluster will fill a smaller volume with a strong radiation field than the Arches Cluster. This is in agreement with the smaller extent of the red emission in the vicinity of the Quintuplet Cluster compared to the Arches Cluster (Fig. 6). In particular, the extent of the radiation field strong enough to cause redder I ISM (8µm)/I ISM (5.8µm) colors does not reach to the (Galactic) southern end of the "Radio Arc Bubble". This is consistent with the fact that while 20 − 25µm observations (IRAS, MSX, MIPS) show a complete bubble, the portions more distant from the Quintuplet cluster do not exhibit unusually red I ISM (8µm)/I ISM (5.8µm) colors. Comparison of the spectroscopic results of Simpson et al. (2007) indicates that the I ISM (8µm)/I ISM (5.8µm) color is fairly well correlated (correlation coefficient r > 0.6) with the continuum intensity at λ > 10 µm, and with H 7-6 12.37 µm, [Ne II] 12.81 µm, [P III] 17.89 µm, [S III] 18.71, 33.48 µm, and [S IV] 10.51 µm. The correlation of the I ISM (8µm)/I ISM (5.8µm) color with the λ > 10 µm continuum color, or with line ratios is weaker. An exception is that correlation with [Si II] 34.83 µm / [Fe II] 26 µm (or 17.94 µm) is quite strong, r > 0.75. This is rather surprising as both species have similar ionization potentials and would be similarly affected by depletion into (or destruction of) dust grains.

Summary
Our Spitzer/IRAC survey of the Galactic Center (Stolovy et al. , 2008 in preparation) reveals prominent 5.8 and 8 µm emission from the ISM in the central ∼ 200 pc of the Galaxy in unprecedented detail. The infrared emission is loosely associated with molecular emission from the CMZ, and several times brighter than infrared emission elsewhere in the Galactic plane. In general, the I ISM (8µm)/I ISM (5.8µm) color of the emission is very uniform across the CMZ, and similar to that seen in other parts of the Galactic plane (e.g. Flagey et al. 2006). Clearly we cannot infer too much from a single color, given the large number of parameters that can be adjusted for the dust models used to interpret the emission. However, we can generally conclude the predominant uniformity of the observed I ISM (8µm)/I ISM (5.8µm) color in the GC implies (1) the emission is dominated by PAHs, (2) the general ISRF in the central 200 pc of the Galaxy does not reach extreme values (χ 10 4 ), and (3) the interstellar extinction law at these wavelengths is flat (e.g. Indebetouw et al. 2005;Flaherty et al. 2007) leading to no significant reddening despite obvious cases of locally high extinction.
There are several extended regions that exhibit much redder than average I ISM (8µm)/I ISM (5.8µm) colors despite not being especially bright. Red regions around the Arches Cluster, the Quintuplet Cluster, and Sgr B1 are apparently produced by a combination of locally strong radiation fields (χ 10 4 ), and a lack of PAHs (perhaps a consequence of the strong radiation field). Other very red regions that are more compact and further from the Galactic plane appear to be associated with the onset of star formation, although they do not appear to contain any detectable point sources.
This work is based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. This research has made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We thank the referee for pointing out the non-negligible influence of the 4.5 µm ISM emission in our analysis.

A.1. Post-BCD Processing
The BCD are single frames (exposures) calibrated with the standard calibration data applied by the Spitzer Science Center (SSC) pipeline. While on average this calibration is quite good, in some cases the state of the IRAC detectors may be different from what the standard calibration assumes. Thus the BCD can still contain some types of instrumental artifacts, and the individual frames may not blend smoothly into a large mosaicked image, especially when that image incorporates multiple AORs.
One significant artifact in the 3.6 and 4.5 µm channels is "column pulldown". This is an electronic artifact that induces low intensities in detector columns coincident with bright (usually saturated) sources. A standard correction was applied to individual BCD frames for the 3.6 and 4.5 µm channels. This greatly mitigates the effect, but does not reduce it to undetectable levels. A similar correction for banding was applied to 8 µm BCD frames to eliminate excess intensity in detector rows containing very bright sources. The 5.8 µm data could have benefitted from such corrections in both detector rows and columns. However the character of the artifact is sufficiently different from the behavior at 8 µm, that attempts at correcting 5.8 µm banding were not successful enough to apply them in an automated way. Thus these artifacts remain, and are a source of systematic error in particular regions where saturated sources are clustered.
The BCD frames for each AOR were self-calibrated using the procedure described by Fixsen et al. (2000). This procedure uses the bright infrared background emission (combined zodiacal light and Galactic emission) as a calibration source. The background is not uniform, but it is stable. Thus, when the dithering and mapping place different detector pixels on the same sky location, the resulting values for properly calibrated data should be the same, apart from a random noise variation from gaussian detector noise and Poisson photon shot noise. Not all detector pixels observe a single sky location, but with a large data set such as ours, if two different detector pixels did not observe a common sky location then there is an intermediate pixel (or a chain of several such pixels) that has observed at least one sky location in common with each of the original two pixels. Through such inter-comparisons, the relative calibrations of each detector pixel can be determined. This is solved using a linear least-squares algorithm. In the framework of this procedure the data (D i , i.e. each pixel of each BCD frame) are modeled as the sum of the true sky intensity (S α at sky coordinates α), a fixed detector offset (F p as a function of detector pixel p), and a variable offset (F q a single value for all pixels of frame q): The least-squares self-calibration allows us to determine the offset terms in the model (i.e. the proper detector calibration), and then invert Eq. A1 to derived the true sky intensity, S α from the data, D i . The images of the fixed detector offsets (F p ) are shown in Figure 11 for each IRAC channel and each AOR. For the 3.6 and 4.5 µm data, the obvious feature is the diffuse stray light "butterfly wing" pattern at the top of the detector. This feature represents the integrated stray light of all the sources of the bright unresolved background, which is not effectively removed through the standard BCD calibration. The amplitude of this artifact is small and is not evident directly in the single BCD frames or the mosaics, but if left uncorrected it would induce artifacts in the ratio of 3.6 / 4.5 µm emission as shown in Figure 12. This figure also illustrates that the column pulldown correction is not entirely effective at both 3.6 and 4.5 µm for the very brightest sources, but the pulldown correction usually does improve the data for weaker, more frequent cases. The fixed detector offsets for the 5.8 and 8 µm channels also show (weakly) the diffuse stray light patterns for these channels (Fig. 11). Additionally the 8 µm results clearly show the accumulation of latent images from bright sources, the clearing of these after an anneal, and the further accumulation of new latent images. The strong latent images acquired at 8 µm during AOR 0s, are from the central cluster at Sgr A. The relative strength and extent of the latent image indicates that the Sgr A cluster is clearly the highest density of bright sources in our survey.
The variable detector offset (F q ) is illustrated in Figures 13 and 14. The time sequence of these shows a regular variation each time one of the AOR strips crosses the brightest part of the Galactic plane. For the InSb detectors used at 3.6 and 4.5 µm, |F q | 0.5 MJy/sr and is usually much smaller than the observed sky intensities (which do not include zodiacal light) as shown in Table 1. The variable offsets can be much larger for the 5.8 and 8 µm detectors, as seen in Fig. 13. Here the variable offsets are correcting for any residual errors in the "firstframe" effect (a change in detector offset as a function of time since the preceding frame), and more noticeably for detector droop (where the detector offset changes as a function of the total illumination on the detector). Figure 15 shows the ratio of our mosaic compared to a standard post-BCD pipeline mosaic for one strip of our survey (AORID = 13368320) at 5.8 and 8 µm. This ratio clearly shows artifacts at the frame boundaries, indicating offset errors in one (or both) of the mosaics. These errors are largest at the bright emission through the center of the strip at b = 0, especially for the 8 µm channel. The facts that (a) the frame seams are faintly visible in the post-BCD mosaic and not our self-calibrated versions, and (b) the subsequent science analysis reveals no similar discrepancies in the ratios of 5.8 to 8 µm emission despite independently derived corrections for each channel, indicate that the variable offsets derived from the self-calibration are appropriate. In other words, the selfcalibration procedure does remove artifacts from the images which would otherwise cause systematic errors when comparing intensities at different wavelengths.

A.2. Extended Source Calibration and Zodiacal Light
The absolute calibration of IRAC is based on the photometry of point sources in a fixed size aperture. In the present study, we are interested in the surface brightness of extended emission. Therefore, we have applied extended source corrections to the data used in this analysis. Specifically, the factors used here are: [0.944,0.937,0.772,0.737] at 3.6, 4.5, 5.8, and 8 µm respectively (Reach et al. 2005;Cohen et al. 2007).
Accurate determination of the colors of diffuse emission also relies on having properly subtracted any instrumental offsets. The major difficulty in measuring the correct instrumental offset is that IRAC does not use its shutter to take dark frames. Instead, "skydarks" (observations at a dark location near the north ecliptic pole) are used to in the pipeline to remove instrumental offsets. This works well as long as the instrumental offsets remain stable between the time that the skydarks are measured and the science data are observed. IRAC's 5.8 µm channel show the greatest instabilities in its offsets. The lack of constraints that would be provide by true dark frames also leads to an offset degeneracy in the selfcalibration of these data. Inspection of Eq. 1 reveals that the procedure is insensitive to an arbitrary offset to the derived sky intensity S α as long as the same offset is subtracted from the derived instrumental offset F p (or F q ).
Finally, for measurement of the intensity and color of diffuse Galactic emission the zodiacal light foreground also must be subtracted. The zodiacal light must be estimated via a model. The model used for Spitzer is derived from the Kelsall et al. (1998) model of the zodiacal light emission as observed by the DIRBE instrument on the COBE spacecraft. The model is estimated to be accurate to ∼ 2%, however larger unrecognized systematic errors may be present. Its use for IRAC data requires interpolation to IRAC's wavelengths, and, more significantly, extrapolation from the view in Earth orbit to the view from Spitzer's location several tenths of an astronomical unit away from Earth. Additionally, because of the skydark subtraction, the zodiacal light must be known along the line of sight to the north ecliptic pole as well as toward the Galactic center. Thus the zodiacal light intensity we subtract from the mosaic images is the difference between that estimated for the Galactic Center observations and that of the skydark observations. This is calculated from the mean values of the ZODY EST and SKYDRKZB keywords in the BCD headers, except for the 5.8 µm channel where these keyword values are in error and are inconsistent with the model spectrum; thus the estimates from SPOT 1 are used instead for the 5.8 µm data. We do not account for the small gradient in the zodiacal light intensity. Table 1 lists the median and mean intensities of the final GC mosaics, along with the self-calibration offsets, and the estimated intensities of the zodiacal light. The zodiacal light is essentially negligible for the 3.6 and 4.5 µm data, but could potentially influence 5.8 and 8 µm results where intensities are low. The variable offset per frame (F q ) can reach large values, but as revealed in Figure 14, the largest values are closely confined to Galactic plane Note. -I = the observed intensity in the final image. Z = the difference between the zodiacal light toward the GC and that toward the skydark calibration field, F p = self-calibration detector offset per detector pixel p, F q = self-calibration detector offset per frame q.      Flagey et al. (2006). The cluster of points shaded in green indicate those found in the shell around the Arches Cluster (see Fig. 6), and the cluster of points shaded in magenta indicate those found in the more compact regions at higher latitudes (see Fig. 7). The large dispersion in the color at I ISM (8 µm) < 35 MJy sr −1 is from faint locations at |b| > 0.5 where errors in the subtraction of stellar and zodiacal backgrounds can have large effects on the colors. (In the middle and right columns, lighter shades indicate redder colors.) All these regions contains extended emission that, despite being relatively faint, is very red in I ISM (8µm)/I ISM (5.8µm) colors but not in I(4.5µm)/I(3.6µm) colors. Fig. 8.-The left panel shows 3.6 (blue), 4.5 (green), and 5.8 (red) µm emission centered on 2MASS 17482532-2911590. There is a slightly extended diffuse feature just ∼ 10" south of the point source. This extended emission is relatively strong at 4.5 µm as indicated by the greenish color, and a high I(4.5µm)/I(3.6µm) ratio (shown in the right panel on a linear range of [0,3]). This color is often an indication of molecular emission in outflows or Herbig-Haro objects associated with very young stars. Fig. 9.-The colored curves show the emission spectra for a selected dust grain model ("BARE-GR-FG" from Zubko et al. 2004) subjected to radiation fields of χ = 10 [0,1,2,3,4,4.5,5,5.5,6,7] times the local ISRF. The black lines indicate the filter responses (arbitrarily normalized) for the 4 IRAC channels, and for the MIPS 24 and 70 µm channels. The PAH spectrum is unchanged with radiation field strength, and hence the IRAC flux ratios remain unaltered until the larger, roughly equilibrium temperature grains get so warm that they encroach on the 8 µm IRAC band. Note that the continuum, not PAH emission bands, dominates the 3.6 and 4.5 µm IRAC filters. The 3.3 µm PAH feature contributes < 30% of the total flux in the IRAC 3.6 µm channel for these models.  Figure 9. The dashed line is the median ratio over the whole GC region. (right) This plot shows the corresponding the I(24 µm)/I(8 µm) ratio. Involving longer wavelengths, this color is more sensitive to changes in the cooler, larger dust grains and lower values of the radiation field χ. Fig. 11.-(a-d) Images of the offsets (F p ) derived for each AOR in a time ordered sequence. These show the diffuse stray light artifacts that remained in the BCD, and the accumulation of long-term latent images at 8 µm. The last three AORs (0n, 0s, 1w) were separated from the first by a downlink and anneal, which cleared previous long term latent images. The Sgr A region creates strong latent images during AOR 0s, which then remain throughout AOR 1w. Display ranges are [-2,2], [-2,2], [-5,5], [-4,4] MJy sr −1 for 3.6 through 8 µm respectively. Fig. 12.-Ratios of standard post-BCD mosaics to our self-calibrated BCD mosaics for one AOR strip of the survey (AOR key = 13371648) at 3.6 µm (left) and 4.5 µm (right). Display range = [0.9,1.2]. The light horizontal bands indicate the regions that are affected by the diffuse stray light in the post-BCD mosaics. These artifacts are present in the ratio of 3.6/4.5 µm post-BCD mosaics, but not in ratios of our self-calibrated mosaics. The dark vertical lines illustrate the changes due to the application of the column pull-down correction.   Figure 13. The 5.8 and 8 µm channel offsets are relatively large and well-correlated with the local sky intensity, and thus these maps resemble very low resolution images of the Galactic center region. Fig. 15.-Ratio of the standard post-BCD mosaic to our self-calibrated BCD mosaic for one AOR strip of the survey (AOR key = 13368320) at 8 µm. Display range = [0.8,1.3] MJy sr −1 . This ratio highlights the corrections made by the fixed and variable offsets derived via self-calibration. Some of the sharpest frame boundaries evident here (in the expanded view) are visible in the standard post-BCD mosaics, but not in the self-calibrated mosaics. The bright region around Sgr B1 is at the center of this 0. • 25 × 1. • 75 strip. The lack of banding correction in the post-BCD mosaic causes the horizontal streaks at the locations of bright stars.