THE BOLOCAM GALACTIC PLANE SURVEY: λ = 1.1 AND 0.35 mm DUST CONTINUUM EMISSION IN THE GALACTIC CENTER REGION

Estimation of the dust column densities and masses from the 1.1 mm emission requires knowledge of the dust properties such as emissivity, spectral index, and temperature, which requires input from additional data. Column densities, densities, and masses are estimated for a subset of the brightest clumps based on dust parameters appropriate for the solar neighborhood (Table 2). Masses, column densities, and densities estimated for all Bolocat entries are tabulated in Table 3 in the electronic version of this paper. For all mass estimation given in this paper, the images and fluxes reported in the V1.0 data release have been multiplied by an empirically determined scaling factor of 1.5 as discussed above.

The mass within the 33'' BGPS effective beam (or equivalently, in a 40'' diameter circular aperture) is given by

$$\begin{equation*} M = {{1.0 \times 10^{-23} { S_{1.1} ({\rm Jy}) } D^2} \over {\kappa _{1.1} B_{1.1}(T) }}, \end{equation*}$$

where B_1.1(T) is the Planck function for dust temperature T, and κ_1.1 is the dust opacity at λ = 1.1 mm in units of cm² g⁻¹. As a function of wavelength, κ_λ = κ₀(λ₀/λ)^β where β is the dust emissivity. Hildebrand (1983) calibrated κ at λ = 0.25 mm, giving a value κ_0.25 = 0.1 cm² g⁻¹ for interstellar dust and a gas-to-dust ratio of 100. For β = 1.5, the Hildebrand opacity implies κ_1.1 = 1.08 × 10⁻² cm² g⁻¹. Ossenkopf & Henning (1994) (OH94) calculated the opacity of grains with ice mantles that have coagulated for 10⁵ years at a gas density of 10⁶ cm⁻³. Enoch et al. (2006) interpolated logarithmically the OH94 opacities to obtain κ_1.1 = 1.14 × 10⁻² cm² g⁻¹, close to the estimate based on the Hildebrand opacity. As discussed in Aguirre et al. (2010), the effective central frequency of the Bolocam 1.1 mm filter is 271.1 GHz. Log-interpolating the OH94 opacity at 1.00 mm to this frequency implies κ_1.1 = 1.14 × 10⁻² cm² g⁻¹. Adopting this value gives

$$\begin{equation*} M = 14.26 (e^{13.01/T(K)} - 1) S({{\rm Jy}}) D^2_{\rm kpc} (M_{\odot }), \end{equation*}$$

where D²_kpc is in units of 1.0 kpc and T is in Kelvin.

The greater metallicity of the Galactic center may result in a larger dust opacity and possibly a different dust emissivity law. This effect has not been taken into account in the present analysis. Due to the spatial filtering inherent in BGPS, and large uncertainties inherent in the estimation of masses using spectral lines, it is not at present possible to determine clump masses with sufficient precision to detect a difference in the dust properties. Use of a larger dust opacity would decrease the estimated masses and column densities given here.

The dust temperature in the CMZ has been studied by Sodroski et al. (1994) who used COBE/DIRBE 140–240 μm data to conclude that 15%–30% of the emission arises from molecular clouds with temperatures of about 19 K, 70%–75% arises from the H i phase with a dust temperature of 17–22 K, and less than 10% comes from ∼29 K dust in the extended H ii phase. Rodríguez-Fernández et al. (2004) used the Infrared Space Observatory (ISO) to study a representative sample of 18 molecular clouds not associated with prominent thermal radio continuum emission. They found typical dust temperatures, measured from λ = 40–190 μm, required a combination of a cold, 15 K component and a warm component with a temperature ranging from 27 to 42 K. Thus, in the analysis presented here, the dust column densities and masses are parameterized in terms of a fiducial dust temperature of 20 K. For T = 20 K,

$$\begin{equation*} M = 13.07 S({{\rm Jy}}) D^2_{\rm kpc} (M_{\odot }) \end{equation*}$$

which at the 8.5 kpc distance to the Galactic center gives M = 944S(Jy)M_☉. This mass can be converted to a column density averaged over the beam,

$$\begin{equation*} N({\rm H_2}) = {{M} \over { \pi \mu _{{\rm H_2}} m_H r_B^2}} ({\rm cm^{-2}}), \end{equation*}$$

where M is in grams, $\mu _{{\rm H_2}}$ = 2.8 is the mean molecular weight per hydrogen molecule (Kauffmann et al. 2008), $r_{B} = D_{\rm GC} \times {\rm FWHM} /(205265 \times \sqrt{4 {\rm ln}{2}})$ is the effective beam radius in centimeters at the Galactic center, and D_GC = 8.5 kpc. Following Kauffmann et al. (2008), the mean molecular weight per particle, μ_p ≈ 2.37 should be used when estimating quantities such as collision rate or pressure. For BGPS, FWHM ≈33'', giving effective beam radius or R_B = 198, which at a distance of 1 kpc corresponds to r_B(kpc) = 2.97 × 10¹⁷ cm and at the distance to the Galactic center is r_B = 2.52 × 10¹⁸ cm. The resulting column density per beam is

$$\begin{eqnarray*} N({\rm H_2}) &=& 2.19 \times 10^{22} [e^{13.0 / T(K)} - 1] S({\rm Jy})\nonumber\\ &\approx & 2.0 \times 10^{22} S({\rm Jy}) (\rm cm^{-2}), \end{eqnarray*}$$

where the right side was evaluated for T = 20 K. This column density corresponds to a visual extinction of A_V ≈ 10.7 mag assuming the usual Bohlin & Savage conversion of N(H) = 1.9 × 10²⁰ cm⁻² per visual magnitude in the standard Johnson/Cousins V filter centered at 0.55 μm. These column densities can be converted into surface densities; Σ(g cm^-2) = 4.65 × 10⁻²⁴N(H₂) or Σ(M_☉ pc⁻²) = 0.048A_V.

For extended sources, masses are given by M = μm_HN(H₂)A_S where A_S is the area subtended by the aperture used to estimate the flux. At the Galactic center, the mass is given by M_tot = 944〈S_1.1(Jy)〉[Ω_S/Ω_B] M_☉ where <S_1.1(Jy)> is the flux averaged over the solid angle Ω_S subtended by the measurement aperture and Ω_B = πr²_B ≈ 2.9 × 10⁻⁸ is the BGPS beam solid angle in steradians. A temperature of T = 20 K is assumed to evaluate the numerical coefficient.

An estimate of the volume density can be made by assuming that each clump is spherical and uniformly filled with emitting matter. Table 2, Columns 9 and 10 give the densities estimated from the fluxes measured in 40'' and 300'' diameter apertures under the assumption that each clump is either 40'' or 300'' in diameter. Interferometric measurement of clumps in the solar vicinity show that much of the millimeter continuum emission arises from regions having diameters of order 10³ AU or less. Thus, BGPS clumps are likely to consist of clusters of unresolved objects much smaller than the beam. The BGPS-based density estimates are volume-averaged quantities. Inside cores, the densities are likely to be much higher while in-between cores, densities may be lower.

Table 2 lists the locations, fluxes, masses and densities for a sample of regions, many of which are marked in Figure 9. S₄₀ gives the flux measured with the FWHM = 33'' Bolocam effective resolution at the location of the brightest emission in each source. The uncertainties in the listed quantity range from about 20% to over 30% for bright sources due to a combination of calibration errors, and incomplete flux recovery in the iterative mapping phase of data reduction. The 1.1 mm images have also been smoothed with a "top-hat" function having a 20 pixel radius to provide an estimate of the total flux produced by extended objects in a 300'' diameter, circular region. At the distance of the Galactic center (8.5 kpc), these masses are given by M(300) = 5.0 × 10⁴S₃₀₀(Jy) M_☉ for 20 K dust where S₃₀₀(Jy) is the average flux density in the 300'' aperture. (Note that for several high-latitude regions suspected to be in the foreground, a distance of 3.9 kpc is assumed in estimating masses as indicated in the comments column.) These mass estimates are lower bounds because they only refer to the flux in small-scale structure to which Bolocam is most sensitive. Because BGPS data start to lose sensitivity to structure on scales ranging from 3' to 6', 300'' is the largest aperture for which reliable photometry can be obtained. As discussed below, comparison of mass determinations in the 300'' aperture to previously published measurements for regions such as Sgr A and Sgr B2 agree to within the estimated uncertainties.

Two fields around Sgr A and Sgr B2 were observed at 350 μm. The 350 μm to 1100 μm flux ratio can be used to constrain dust temperatures or emissivities. Three corrections must be applied to the SHARC-II data prior to comparisons with BGPS; matching of angular resolution, correction for signals picked up in the sidelobes, and matching of the spatial frequency response functions. No correction has been made for contamination by spectral lines.

The SHARC-II 350 μm images were obtained with a 9'' FWHM beam. Thus, they have to be convolved with a Gaussian function with σ = 135 to match the 33'' effective BGPS beam. Because the data values in both the 350 and 1100 μm images are in units of Janskys per beam, the pixel values in the convolved SHARC-II images must be multiplied by the ratio of beam areas (13.4 when scaling from 9'' to a 33'' beam).

Inspection of the 350 μm beam delivered by the CSO on bright point sources such as planets shows that in addition to the main 9'' beam, there is a low-level error pattern several arcminutes diameter with a roughly hexagonal shape. The main-beam efficiency of the CSO at 350 μm is only about 0.4 ± 0.15. In a complex field such as the Galactic center with an abundance of bright, extended emission, fluxes will be overestimated when compared with a point source calibrator in the main beam due to the signal entering the sidelobes. Because bright emission near both Sgr A and Sgr B2 is extended on scales of many arcminutes, we assume that about 30% of the signal at a typical location in the maps arises from signal picked up in the sidelobes. As an approximate correction for this effect, we scale the convolved 350 μm images by a factor 0.7 to correct for clumps contributing emission through the sidelobes.

Comparison of the two 350 μm maps of the Sgr A region discussed above shows that the rapidly scanned SHARC-II images recover flux on larger angular scales than the 1.1 mm BGPS images. Thus, the SHARC-II images are spatially filtered with an "unsharp mask." The mask is constructed by convolving the 350 μm with a Gaussian function having a kernel σ_k, and subtracting a fraction f of the mask from the original image. The optimal values of σ_k and f were chosen by finding the closest match between the spatial structure of the resulting "unsharp-masked" SHARC-II image with the corresponding BGPS image. For the Sgr A field, the best choice is σ_k = 90'' (FWHM = 212'') and f = 0.8; for Sgr B2, σ_k = 68'' (FWHM = 160'') and f = 0.5. The beam-matched, spatially filtered 350 μm images were divided by the corresponding 1.1 mm BGPS images to form flux-ratio maps.

The observed intensity of continuum emission at each location in a map is given by $S_{\nu } = \Omega _{\rm beam} [1 - e^{- \tau _{\nu d}}] B_{\nu }(T) ({{}\rm Jy })$ where Ω_beam ≈ 2.9 × 10⁻⁸ sr is the beam solid angle corresponding to the 33'' effective resolution of the 1.1 mm data and the convolved, beam-matched, 350 μm observations, τ_ν is the optical depth of the emitting dust, and B_ν(T) is the Planck function at each frequency. The optical depth in the submillimeter to millimeter regime scales with frequency as ν^β where β is the emissivity power-law index. For small grains, β is close to 2 in the millimeter to submillimeter regime but can be smaller if the grains have experienced substantial growth or are coated by ice mantles (Ossenkopf & Henning 1994). In the optically thin limit, the flux ratio at each point in a ratio map is given by

$$\begin{equation*} R = S_{350\,\mu {\rm m}} / S_{1.1\,{\rm mm}} = { { \nu _{350}^{3 + \beta } [{\rm exp} (h \nu _{1.1} / k T_d) -1]} \over { \nu _{1.1}^{3 + \beta } [{\rm exp} (h \nu _{350} / k T_d) -1]} } \end{equation*}$$

which cannot be solved analytically for T_d. A look-up table is generated for each choice of β and T_d and used to analyze points in the ratio maps. For hot β = 2.0 dust where both 350 μm and 1100 μm are in the Rayleigh–Jeans limit, the observed intensity ratio should be around R = 100. Larger ratios require β>2. Where the pixels have good signal at both 350 μm and 1.1 mm, flux ratios have values R< 70, corresponding to dust temperatures less than 70 K for β = 2 dust. The interpretation of these results will be discussed below.

The Bolocat clump catalog contains the highest concentration of clumps in the sky toward the inner few degrees of the Galaxy, and the 1.1 mm emission is about five times brighter than adjacent parts of Galactic plane. The region between l = 0° and 15 contains 542 Bolocat entries, or 361 objects per square degree. In contrast, the Galactic disk away from the Galactic center between l = 55 and 115 contains 511 entries in 6 deg² or 85 objects per square degree on average. This field is chosen for comparison since it likely has a similar surface density of foreground BGPS clumps as the Galactic center region. The symmetrically located field in the fourth quadrant was not observed to comparable depth due to its lower elevation as seen from Mauna Kea. The surface density of clumps is somewhat larger around l = 20°–35° where the line of sight passes through the tangent points in the molecular ring. Thus, the surface density of clumps in the CMZ is about 4.3 times higher than in the plane. When Bolocat entries located at higher latitudes above and below the CMZ are excluded, the surface density of Bolocat clumps located close to the Galactic center is more than a factor of 5 higher. This is a lower limit since confusion is likely to make Bolocat have lower fidelity in the Galactic center. Thus, at least 80% of Bolocat entries in the inner few degrees of the Galaxy are likely to be within 500 pc of the nucleus, located at approximately 8.5 kpc from the Sun. Less than 20% are likely to be located in the molecular ring at distances of 2 to 6 or 11 to 14 kpc from the Sun.

The concentration of clumps in the CMZ permits the generation of a clump mass spectrum for a large sample of objects located at a common distance between about 8 and 9 kpc. The formulae in Section 2.1 are used to estimate individual clump masses from the peak flux per beam in a 40'' diameter (1.64 pc at D = 8.5 kpc) aperture tabulated in Column 13 in Bolocat. Figure 5 shows the mass spectrum using the assumption that all clumps have the same temperature of T_dust = 20 K. For masses above about 100 M_☉, the mass spectrum can be fit by a power-law, dN/dM = kM^α with an index α≈ –2.14 ± 0.1. On this scale, the Salpeter initial mass function (IMF) of stars has a slope of –2.35. Between about 400 and 10⁴ M_☉, the slope is slightly steeper with α≈ –2.58 ± 0.1, but there are several very massive clumps such as Sgr B2 which lie well above the extrapolated curve. Figures 6–8 show mass spectra under different assumptions about the dust temperature and its variation with either Galactocentric distance or the peak 1.1 mm flux in each clump. These spectra are based on the fluxes reported in the V1.0 data release and have not been scaled by the factor of 1.5.

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Great caution must be exercised in generating mass spectra for diffuse objects and in the comparison of such spectra with those obtained with different telescopes, at different wavelengths, and using different tracers. Mass spectra will be biased by the properties of the data acquisition strategy, the properties of reduction pipelines, the definitions and algorithms used to define "objects." Nevertheless, mass spectra are frequently presented in the literature on molecular clouds and their condensations. Despite their potential pitfalls, mass spectra were computed because BGPS contains one of the largest samples of dense clumps ever obtained.

Miyazaki & Tsuboi (2000) present a mass spectrum for molecular clouds in the Galactic center, finding that for masses greater than 10⁴ M_☉, the slope of the mass function is α = –1.59 ± 0.07, considerably shallower than the slope of BGPS clumps. Inspection of their J = 1–0 CS images shows that the emission is continuous, area filling, and dominated by extended structure in both the l − b and l−V maps; it is far less "lumpy" than the BGPS maps which are dominated by small-scale structure. Thus, it is not surprising that the overlap in mass between the CS and BGPS mass spectra is small or that the majority of CS features are considerably more massive than the BGPS clumps. The slope of the 1.1 mm BGPS mass spectrum above about 100 M_☉ is similar to the mass spectrum of low-mass cloud cores in the solar vicinity which has a slope of around –2.3, similar to the stellar IMF (Motte et al. 1998, 2001; André et al. 2007; Enoch et al. 2008); it is considerably steeper than found for CO-emitting GMCs in the Galactic plane.

Because at the distance of the Galactic center, most BGPS clumps will form clusters rather than individual (or multiple) stars, it may be more appropriate to compare the BGPS mass spectra with the mass spectra of star clusters. Observations of Galactic star clusters (Battinelli et al. 1994) and OB associations (McKee & Williams 1997) reveal steep mass spectra scaling as M⁻². Massive young star clusters in the merging galaxy pair known as the "antennae" also have dN/dM = N₀M⁻² (Zhang & Fall 1999) in the mass range 10^3.5–10^6.5 M_☉. Similar results have been found for clusters in M51 (Bik et al. 2003) and older clusters in the Large Magellanic Cloud (de Grijs & Anders 2006). Nearby low-mass star-forming cores, more distant and more massive BGPS clumps in the Galactic center, and Galactic and extra-galactic star clusters share the common trait of having steep mass spectra. That all three classes of object are gravitationally bound may have something to do with this trait.

The BGPS data acquisition and reduction pipeline results in images with highly attenuated response to extended structure, effectively acting as a high-pass filter that attenuates low spatial frequencies (Aguirre et al. 2010). Furthermore, the Bolocat watershed algorithm (Rosolowsky et al. 2010) tends to subdivide sources with complex structure into multiple clumps. The Sgr B2 complex provides a good example. While single-dish maps in tracers such as CO or other high-dipole moment molecules such as CS show that this region consists of a single cloud complex about 02 × 03 in extent (e.g., Figure 3), the BGPS images such as Figure 9 highlight local maxima and interconnecting ridges. Bolocat subdivides the Sgr B2 complex into several dozen clumps. Attenuation of low spatial frequencies combined with the watershed algorithm of Bolocat may remove objects from the high-mass, large-extent portion of the mass spectrum and re-distribute the flux into a larger number of lower-mass entries. This has the effect of steepening the power-law index of the derived mass spectrum. Detailed simulations to quantify this effect are needed and will be the subject of a future paper. In complex fields such as the Galactic plane, line-of-sight blending of physically unrelated sources may lead to further confusion which can alter the slope of the mass function. We urge the reader to exercise great caution in the interpretation of mass spectra of diffuse objects.

Bolocat (Rosolowsky et al. 2010) tabulates fluxes in four different ways. In addition to the flux in a 40'' aperture, fluxes are also measured in 80'' and 120'' apertures. The catalog also tabulates the total flux in the area assigned to each clump along with a beam-deconvolved "effective" radius of the clump (the radius of a circle that has the same area as the clump). Mass spectra based on fluxes measured in larger apertures have shallower slopes by about 0.1–0.2.

The mass spectrum derived from the Bolocat fluxes measured in a 120'' diameter aperture using a constant dust temperature of 20 K has a power-law index α = –2.04 (Figure 6), slightly shallower than the index based on measurements in a 40'' aperture. This difference may be due to increased source confusion within the larger aperture. The relatively small slope change may suggest that source blending and confusion already affects the slope based on fluxes in the smaller aperture. The slope also depends slightly on the widths of the bins used to generate the histograms.

A curious puzzle about the Galactic center is that the dust temperature appears to be lower than the gas temperature despite the high average density. Both the COBE and ISO satellites determined that the average dust temperature in the CMZ ranges from 15 to 22 K (Reach et al. 1995; Lis et al. 2001). Thus, the use of T_dust = 20 K in mass estimation is justified. However, the dense gas associated with this dust is considerably warmer with 80% of the ammonia having temperatures between 20 and 80 K and 18% warmer than 80 K (Nagayama et al. 2007).

Is the slope of the derived mass spectrum sensitive to variations in T_dust? Two types of dust temperature variations are considered to answer this question: first, the dust temperature of a clump is assumed to decrease with increasing projected distance from Sgr A as a power law. Second, the dust temperature of each clump is assigned a value depending on the observed flux; a 100 Jy source such as Sgr B2 is assumed to have dust with T = 50 K while a 1 Jy source is assumed to have T = 20 K. This corresponds to a dust temperature varying with flux, F, as F^0.2.

Given a source of luminosity L located at a distance r from a grain, the mean grain temperature depends on the luminosity L and distance r as

$$\begin{equation*} T(r) = T_0 L^{1 / \gamma _g} r^{- 2 / \gamma _g} \end{equation*}$$

(Scoville & Kwan 1976) where T₀ is a normalization constant. For blackbodies larger than the emitted and absorbed wavelengths, γ_g = 4. For interstellar grains, γ_g ranges between 5 and 6 for grain emissivity power-law indices ranging from 1 to 2. Assume that the heating luminosity L is generated by stars in Galactic disk and bulge, that the luminosity-to-mass ratio is constant, and grains at a given distance r from the Galactic center are heated mostly by the stars within a sphere of radius r centered on the Galactic center. Thus, the enclosed luminosity, L(r), is a function of r which can be estimated by assuming that the density of stars decreases with increasing distance from the Galactic center as ρ(r') = ρ₀r'^−δ. For δ = 2 (the singular isothermal sphere), integrating the mass distribution from r' = 0 to r implies that the mass enclosed within a radius r is proportional to r. Thus, the luminosity L(r) is also proportional to r. Inserting this in the above equation for T(r) gives a simple approximation for the radial variation of the grain temperature

$$\begin{equation*} T(r) = T_0 r^{- 1 / \gamma _g}, \end{equation*}$$

where T₀ is normalized to a temperature at some radius. This estimate applies to grains which are located mostly at cloud surfaces or otherwise are exposed to the visual and near-IR component of starlight. As a concrete example, choose β = 1 and γ_g = 5 so that T(d) = T₀r^−0.2 because it yields reasonable cloud surface temperatures over a wide range of Galactocentric distances. Taking T₀ = 50 K at r = 10 pc implies T = 27.5 K at r = 200 pc, 15.5 K at r = 3 kpc, and 12.5 K at 8.5 kpc. This parameterization provides reasonable base temperatures at the surfaces of clouds lacking internal or close-by heating sources. It is used to test the sensitivity of the mass spectrum to a systematic decrease of grain temperature with increasing galactocentric distance.

Systematic temperature variations do not have a large impact on the derived slope of the mass spectrum. However, the elevated temperature does decrease the estimates of total mass. The mass spectrum for γ_g = 5 corresponding to a radial temperature gradient expected for a grain emissivity that decreases with wavelength as a power law with index β = 1 has a slope of α = –2.44 (Figure 7).

Finally, it is possible that internal heating by massive stars results in an average dust temperature that is a function of 1.1 mm flux. To determine the consequences of such a model on the derived mass spectrum, we assume that the dust temperature varies as T = T₀F^0.2_1.1 where the normalization constant T₀ is chosen to give a dust temperature of 20 K for a 1.1 mm flux density per beam of 1 Jy. This model results in a steeper mass spectrum shown in Figure 8 with a slope of α = −2.64.

These results show that the mass spectrum of millimeter clumps having masses above about 70 M_☉ is represented by a steep power law comparable to low-mass star-forming cores with masses below about 10 M_☉ (Motte et al. 1998, 2001; André et al. 2007). The derived power-law index is relatively insensitive to assumptions about the dust temperature. In summary, the power-law index of Bolocat clumps in the Galactic center is 2.14 (+0.4,–0.1). The mass spectrum in the Galactic center is steeper than GMC mass spectra in the inner galaxy and LMC and comparable to the outer Galaxy and M33 (Rosolowsky 2005). However, great caution must be exercised in the interpretation of these mass spectra owing to the data acquisition and reduction methodology.

The total H₂ mass of the 1428 Bolocat clumps between l = –15 and 45, measured in a 40'' aperture, assuming that all are located at a distance of 8.5 kpc, have a dust temperature of 20 K, and multiplied by the empirically determined scaling factor of 1.5 needed to bring BGPS fluxes into agreement with MAMBO (Motte et al. 2003, 2007) and SIMBA (Matthews et al. 2009) gives M(H₂) ∼ 6.6 ± 1 × 10⁵ M_☉. Assuming a temperature gradient that decreases with distance from Sgr A as proposed above gives the lowest total mass of order ∼5.7 ± 1.2 × 10⁵ M_☉. Assuming a dust temperature that scales with detected 1.1 mm flux as described above gives a total mass of ∼1.1 ± 2.3 × 10⁶ M_☉ in a 40'' aperture. In the 120'' aperture (assuming a constant dust temperature of 20 K), this mass grows to ∼5.5 ± 1.2 × 10⁶ M_☉. Using the total integrated flux (Bolocat Column 19) within the effective radius (Bolocat column 12) results in a total mass ∼6.2 ± 1.8 × 10⁶ M_☉.

Summing all BGPS flux in this field above 100 mJy gives a mass of about ∼6.3 ± 1.8 × 10⁶ M_☉. The total mass of molecular gas within 500 pc of the center has been estimated to be between 3 and 8 × 10⁷ M_☉ (Sodroski et al. 1994; Dahmen et al. 1998; Tsuboi et al. 1999). Thus, the total 1.1 mm emission from clumps in a 40'' aperture represents about 0.8%–2.2% of the total mass in molecular gas while the mass in a 120'' aperture is about 7%–18% of the total. Summing the flux above 100 mJy corresponds to 8%–21% of the total molecular gas mass. This is about a factor of 2 larger than what was found in the Gem OB1 association toward the Galactic anticenter (M. K. Dunham et al. 2009, in preparation), indicating that a larger fraction of the molecular gas near the Galactic center is in compact clumps.

It is important to recall that BGPS filters out extended emission on scales larger than about 5' and that the 1.1 mm data are only sensitive to dust in compact structure. Such spatial- filtering is not present in the single-dish spectral line data which therefore can trace gas in extended regions. Thus, the small mass fraction detected by BGPS may be a consequence of the spatial filtering or diffuse, extended emission. Spatial filtering of the spectral line data to match BGPS could be used to constrain the relative abundances of dust and the tracers producing the lines.

Figure 1 shows an image of the 1.1 mm emission from the Galactic center region. Figure 2 shows this image as contours superimposed on the Spitzer Space Telescope IRAC 8 μm image. As at most long IR, submillimeter, and radio wavelengths, the area-integrated 1.1 mm emission is about five times brighter than in adjacent portions of the inner Galactic plane. While most of the Galactic plane millimeter-wave continuum emission is dominated by discrete clumps (Aguirre et al. 2010; Rosolowsky et al. 2010; M. K. Dunham et al. 2009, in preparation), the CMZ contains bright arcminute-scale extended structures interconnected by a lacy network of filaments. Figure 9 shows the inner 2 deg² of the Galactic plane at 1.1 mm. Figure 10 shows the same field at 8.0 μm as observed by the Spitzer Space Telescope (Arendt et al. 2008) with contours of 1.1 mm emission superimposed.

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The chain of bright 1.1 mm clumps extending from Sgr B2 (G0.68,–0.03) to G0.26+0.03 and the surrounding lower level diffuse emission is seen in silhouette against the bright background in the infrared image. The 20' halo of 1.1 mm emission surrounding Sgr B2 can also be seen in silhouette against the infrared background. Thus, most of the BGPS emission between l = 10 and 025 must originate in front of most of the stars and mid-infrared light in the central bulge of the Galaxy. This is consistent with the suspected orientation of the central bar whose positive longitude part is closer.

Several features of the dust continuum emission from the CMZ are apparent. First, most of the 1.1 mm emission with a flux density greater than about 1 Jy beam⁻¹ originates from a few dozen clumps (in the Bolocat clump catalog, many of these objects are subdivided further). The two brightest peaks correspond to Sgr B2 and the Sgr A*. At lower levels, the CMZ can be decomposed into hundreds of individual clumps and filaments. Bolocat contains 1428 entries in the field covered in Figure 1. Second, most of the 1.1 mm emission in the CMZ closely follows the distribution of dense molecular gas traced by high-dipole moment molecules and the large line-width CO emission (Bally et al. 1987, 1988; Oka et al. 1998, 2001b). In the inner square degree, both the 1.1 mm dust continuum and the molecular tracers of the CMZ exhibit a smaller scale height than emission from the Galactic disk measured beyond |l|>1°. Near Sgr A (l≈ 0°), the vertical extent (orthogonal to the Galactic plane) of the brightest BGPS emission and broad-line molecular gas is only about 2'–3' (5–8 pc); near Sgr B2 around l≈ 07 this layer has a thickness of about 10'–15' (25–40 pc). Beyond |l|∼ 1°, the scale height of 1.1 mm clumps associated with the CMZ increases dramatically. Around l = 13 and 32 the gas layer is nearly 60' (150 pc) thick. Away from the Galactic center, between l = 5° and 35° the dust layer has an average thickness of about 08 or 125 pc. The spatially averaged 1.1 mm flux from the Galactic center features is about three to five times higher than the average flux from the roughly ± 04 scale-height clump population of the molecular ring which is at least 3 kpc in front of (or behind) the Galactic center. Third, many 1.1 mm features can be seen in silhouette in the 8 μm IRAC and 24 μm MIPS images. Fourth, at levels up to several hundred mJy beam⁻¹, a complex network of filaments forms a frothy background of emission which surrounds voids. The filaments may either trace the walls of stellar-wind bubbles or supernova remnant (SNR), the warm edges of molecular clouds, cavities produced by H ii regions, or "fossil" cavities whose energy sources are extinct.

Comparison with the 20 cm radio maps (Figure 11) shows that only a small subset of these 1.1 mm cavities is associated with radio continuum emission. Figure 11 shows the 20 cm emission imaged with 30'' resolution (Yusef-Zadeh et al. 2004) superimposed on the 1.1 mm image. The Sgr A, B1, B2, C, and D regions that contain young massive stars and star clusters stand out in both the radio continuum and at 1.1 mm. The non-thermal radio continuum filaments tend to be located adjacent to 1.1 mm clumps. For example, the largest and brightest group of filaments near l = 018 crosses the Galactic plane at the positive latitude end of the ridge of clumps associated with the 50 km s⁻¹ cloud that stretches from l = 0°–017. The fainter non-thermal filament located directly above (in latitude) Sgr A known as N5 (Yusef-Zadeh et al. 2004) passes through the cluster of clumps associated with the H ii region G359.96,+0.17. The brightest part of filament N8 abuts a diffuse 200 mJy BGPS clump at [359.78,0.17] (not marked).

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Some of the diffuse 1.1 mm emission between l = 024 and –3596 located at positive latitudes is visible in emission at 8 μm. At positive longitudes, this region contains the thermal "arched filaments" near the Arches cluster. The 1.1 mm emission may either trace warm dust or free–free emission from dense, photoionized plasma illuminated by the Arches cluster. At negative longitudes, 1.1 mm clumps associated with the high-latitude, positive velocity portion of the "EMR" (b = 00 to about 02; V_LSR = 100–200 km s⁻¹) are also associated with 8 μm emission. At negative longitudes, the 20 km s⁻¹ cloud located south of Sgr A* is visible in silhouette and must therefore be in front of the central cluster. A pair of oppositely facing comet-shaped clouds near the right edge of Figures 9–11 (associated with Sgr C and G359.62–0.24) are both seen prominently in extinction. Except for these clouds, few of the 1.1 mm sources at negative longitudes are visible in extinction in the IRAC data.

The fainter 1.1 mm continuum is structured and organized into a filamentary network of arcs and circular features. This is most evident in the left half of Figure 9. Some rings and partial rings surround radio continuum features and therefore probably mark dust heated in photon-dominated regions (PDRs) and the edges of cavities created by massive stars and star clusters. However, most rings, arcs, and filaments are not associated with known H ii regions or SNRs. Massive stars that have shed their natal cocoons are hard to separate from the multitude of intrinsically fainter foreground stars, because in the IR distant massive stars and closer lower-mass stars have similar colors and spectral shapes. These cavities may also be the remnants of older generations of stars and clusters whose most massive members have either exploded as supernovae more than a million years ago so that all radio continuum has long since faded, or whose stars were ejected by dynamical processes that create runaway, high-velocity stars. Runaway stars are most common among spectral type O (Gies & Bolton 1986; Gies 1987); many of these stars were ejected by dynamic interactions in dense clusters (Gualandris et al. 2004). The shells and filaments may be tracers of fossil star and cluster formation in the CMZ. However, it is also possible that randomly superimposed filaments create the impression of cavities.

The large kidney-bean-shaped clump at [0.26,0.03] (see Figures 9 and 12) is remarkable for being one of the brightest features at 1.1 mm and for being the most prominent infrared dark cloud (IRDC) in the entire Galactic center region (Liszt 2009). This feature is warm (Figure 13) and seen in absorption at 24 μm (Figure 14). Although Bolocat lists two entries for this object (nos. 96 and 99), dense-gas tracers such as CO, CS, and HCN indicate that the entire structure is at approximately the same radial velocity (V_LSR≈ 40 km s⁻¹ with a line width extending from 20 to 60 km s⁻¹). Liszt (2009) use the Midcourse Space Experiment (MSX) data to set a lower bound on the column density of N(H₂) >1.0 × 10²³ cm⁻² (A_V> 53 mag), consistent with the 1.1 mm based column density estimate in a 40'' aperture. The higher resolution 350 μm observations (Figure 12) indicate that it has considerable substructure with consequently larger column densities.

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Figure 13 shows the ratio map formed from the 350 μm and 1.1 mm data in the vicinity of Sgr A. In regions with significant emission at both wavelengths, there is a general trend that the clump centers where the fluxes peak have the smallest flux ratios. The flux ratios tend to increase toward the clump edges. Ratios at a number of locations in Figure 13 are indicated by the blue arrows. In Figure 13, the flux in the 1.1 mm BGPS maps have been scaled up by a factor of 1.5. Without this scaling, a large fraction of the pixels have ratios larger than 100. Figure 13 shows that clumps in the CMZ have warmer dust than either foreground or background clumps in the molecular ring.

Sgr B2 is the brightest clump in the entire BGPS survey. Figure 15 shows the 350 μm image. As shown in Table 2, it has the largest total mass (>5 × 10⁵ M_☉), highest volume-averaged density (n(H₂)>10⁴ cm⁻³), and greatest column density (N(H₂)>2 × 10²⁴ cm⁻² in 33'' beams centered on Sgr B2 Main and North) of any BGPS clump. Thus, the average extinction toward Sgr B2 M and N corresponds to at least 1000 mag at visual wavelengths. Molecular line measurements indicate that the Sgr B2 complex has a mass of at least 5 × 10⁶ M_☉, at least a factor of 10 greater than implied by the total 1.1 mm flux detected by BGPS and listed in Table 2 (Jones et al. 2008). Over 50 compact H ii regions and powerful masers indicate that a major burst of star formation is occurring here (Gaume et al. 1995; de Pree et al. 1998; De Pree et al. 2005). The molecular line data provide evidence that the intense burst of star formation occurring in Sgr B2 may have been caused by the supersonic collisions between two giant Galactic center molecular clouds (Jones et al. 2008; Liszt 2009). Sgr B2 may be located near the high-longitude end of "x2" family of orbits in a bared potential (Contopoulos & Papayannopoulos 1980) where it may encounter gas moving along the "x1" family of orbits. Sgr B2 is located at the high-longitude end of a chain of 1.1 mm clouds, mid-IR extinction features, embedded 24 μm sources, and H ii regions starting near the less-extincted and older Sgr B1 complex (Yusef-Zadeh et al. 2009). Sgr B2 may represent the most recent event in a chain of sequential star formation that started near Sgr B1. The Sgr B2 complex may be giving birth to a massive star cluster or even a super star cluster comparable to the somewhat older Arches and Quintuplet clusters located in the CMZ.

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Figure 16 shows the 350 μm/1.1 mm flux-ratio map derived from the 350 μm image matched in spatial frequency response to BGPS divided by a 1.1 mm map scaled-up by a factor of 1.5. No correction for line contamination has been applied. Sgr B2 is likely to be most affected by such contamination; Lis & Goldsmith (1991) found about 30% of the 350 GHz flux is due to lines and contamination. The two brightest peaks, corresponding to Sgr B2 North and Sgr B2 Middle (or Main) have the smallest flux ratios (25 and 34, respectively) which for β = 2 dust, implies temperatures of 14–17 K. It is likely that the dust is internally heated by forming massive stars to at least 20 K, and possibly 40 K (Lis et al. 1991). As these authors proposed, it is possible that the low flux ratio is the result of an unusually low β; for β = 1.5, ratios of 25 and 34 imply 21 K and 33 K.

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Figures 17 and 18 show the BGPS 1.1 mm contours superimposed on the Spitzer 8 and 24 μm images, respectively. These images clearly show that most of the dust associated with the Sgr B2 region is in front of the bulk of the 8 and 24 μm emission. Furthermore, several of the roughly circular cavities in the BGPS images are translucent in the IR data. The large Sgr B1 complex of H ii regions and infrared sources has created a giant cavity on the low-longitude and low-latitude side of Sgr B2. Highly obscured nebulosity and IR sources abut the Sgr B2 submillimeter/millimeter peaks on the low-latitude side. However, no IR sources can be seen at the locations of Sgr B2N and Sgr B2M. More discussion of the Sgr B2 region is given in the Appendix.

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The second most prominent peak in the 1.1 mm data (Sgr B2 is the first) is centered on Sgr A* which is thought to mark the location of the 4.5 ± 0.4 × 10⁶ M_☉ black hole at the Galactic center (Genzel et al. 2003; Ghez et al. 2008). The extended radio source known as Sgr A consists of Sgr A East, a several arcminute diameter oval of non-thermal radio continuum emission thought to trace an SNR, and Sgr A West, an H ii region which consists of a "mini-spiral" of dense, thermal plasma with the central black hole, Sgr A* at its center. The "mini-spiral" is surrounded by the ∼2 pc radius "circumnuclear ring" (also referred to as the "circumnuclear disk"—CND) of dense molecular gas and dust (Morris & Serabyn 1996). The CND is clearly seen in the 350 μm image shown in Figure 12.

Emission from Sgr A* peaks near a wavelength of 1 mm (Serabyn et al. 1997). Between 1993 and 1996, the peak flux near this wavelength was around 3.5 Jy (Figure 2 in Serabyn et al. 1997). In 2005 July, Sgr A* was about 1.6 times brighter at 1.1 mm with a flux density of 5.7 Jy. This may be an underestimate because the BGPS pipeline may not fully recover the flux of bright sources or those surrounded by an extended envelope (Aguirre et al. 2010; Rosolowsky et al. 2010). Diffuse emission produced by dust or hot plasma along the line of sight contributes at most 1 Jy per beam. Thus, the increased brightness of Sgr A* in the 1.1 mm maps is most likely due to variability which is consistent with a non-thermal origin for most of the 1.1 mm continuum emission from that source.

Figure 12 shows the 350 μm SHARC-II image of the 30' field surrounding Sgr A. Figure 14 shows contours of 350 μm emission superimposed on the Spitzer 24 μm image. In contrast to 1.1 mm, there is very little 350 μm emission from Sgr A* itself. At 350 μm, Sgr A* has a flux density between 0.5 and 1.5 Jy depending on how the large background flux is subtracted. While emission from the CND and Sgr A* are blended at 1.1 mm by the 33'' effective resolution, in the 350 μm image, the CND is resolved (Figure 12) and appears as a nearly complete oval. Several dozen discrete clumps and filaments can be identified in the region surrounding the ring in addition to the prominent 20 and 50 km s⁻¹ clouds. Interferometric observations of molecules such as HCN resolve the CND into a collection of dense clumps (Christopher et al. 2005). The detection of maser emission and hyper-compact sources of radio continuum emission indicates ongoing massive star formation in the CND (Yusef-Zadeh et al. 2008).

The central 30' region contains several chains of 1.1 mm clumps and filaments elongated toward Sgr A (turquoise lines in Figures 9–11 red lines in Figures 12 and 14). These chains are narrower at the end facing Sgr A and may therefore be cometary clouds sculpted by this source. The two most prominent chains are located about 1' south of the Arches cluster and 2' south of the Quintuplet cluster and Pistol Star; these chains are each about 6'–8' long. A shorter, fainter chain is located between the Arches and Quintuplet clusters, and a nearly 10' long chain is located about 3' north of the Arches cluster. These features are also seen in the 350 μm image (Figures 12 and 14) where they break up into complicated substructures with an overall elongation toward or away from Sgr A.

About a half dozen compact clumps located between the IRDC at [0.26,0.03] and the Quintuplet cluster near l∼022 are also elongated along the Galactic plane and may be smaller cometary clouds less than 1'–2' in length. These are best seen in Figure 9. Figure 11 shows that these clumps are located at the high-longitude side of the prominent non-thermal filaments that cross the Galactic plane at l≈ 018.

Mid-infrared images from the MSX and ISO satellites reveal a very bright, limb-brightened bubble of infrared emission surrounding the Quintuplet and Arches clusters (Moneti et al. 2001). Inspection of the 24 μm images obtained by the Spitzer Space Telescope (program P20414; PI: F. Yusef-Zadeh) also shows the prominent elliptical bubble with a major axis bounded toward positive latitudes by the Arched Filaments around [l, b] ≈ [+011, +010] and toward negative latitudes by a rim at [l, b] ≈ [+012, −018] and extending from l≈ 0.0 to 0.18. The oval infrared shell is clearly seen in Figure 14. The location of the 24 μm ovoid is indicated by the large ellipse in Figures 9–14.

The chains facing Sgr A are in the projected interior of the infrared "bubble" surrounding the Arches and Quintuplet clusters. It is likely that this bubble traces the walls of a cavity that extends from the non-thermal filaments near l≈ 018 to at least l ≈ 00 near the 50 km s⁻¹ cloud and contains both the Arches and Quintuplet clusters and possibly the Sgr A complex. Figure 11 shows that this region is filled with diffuse 20 cm continuum emission. Figure 12 shows this region at 350 μm.

The dashed line in Figure 2 outlines the boundary of a large region relatively devoid of clouds as traced by 8 or 24 μm emission. This feature may trace an extension of the bubble associated with the Arches and Quintuplet clusters that opens up above and below the Galactic plane. While this feature can be traced to the edge of the observed field above Sgr A, below the Galactic plane it becomes confused with the foreground bubble centered on [l, b] = [+041, −050]. An extended bubble blowing out of the Galactic plane may have been created by energy release from Sgr A, Arches, and the Quintuplet regions.

The Galactic center contains several remarkable cloud complexes which exhibit line widths ΔV> 100 km s⁻¹ over a region typically about 05 in diameter (Bania et al. 1986; Stark & Bania 1986; Oka et al. 1998; Liszt 2006, 2008). Although Bania's Clump 1 at [l, b] ≈ [–355°, +04] and Clump 2 [l, b] ≈ [+32, +02] were the first to be noted, similar but less distinct, large line-width complexes exist near [+13,+02] and near [+50,+02] (Oka et al. 1998, 2001a; Dame et al. 2001; Bally et al. 1987, 1988). Bania's Clump 2, located at a projected distance of about 450 pc from the Galactic center, is the most prominent of these features with a line width of ΔV≈ 200 km s⁻¹. The feature extends from the mid-plane to about b = +08 and consists of about 16 CS-emitting clumps each with masses of order 5 × 10⁵ M_☉ and H₂ densities in excess of 2 × 10⁴ cm⁻³ (Stark & Bania 1986) surrounded by an envelope of CO-emitting material. Such dense and massive clumps are usually associated with ongoing massive star formation.

Figure 19 shows the BGPS 1.1 mm map of the region around l = 3°. Some BGPS data were obtained up to b = ± 15 that shows several additional features associated with Bania's Clump 2. These data are noisier than the BGPS data along the plane and are not shown here. The 1.1 mm clumps coincide with the CS and CO clouds in Bania's Clump 2. Figure 20 shows the Spitzer Space Telescope 8 μm image from GLIMPSE (Benjamin et al. 2005) with superimposed 1.1 mm contours. The most striking result is the absence of extended infrared sources or bubbles associated with most Clump 2 features. In contrast to Clump 2, 1.1 mm clumps located along the Galactic plane tend to contain infrared sources at 4.5, and 5.8 μm, extended 8 μm emission, and prominent bubbles, indicating massive star and star cluster formation. Inspection of the 24 μm MIPS scan maps available from the Spitzer Science Center archives also shows a general absence of mid-IR sources in Bania's Clump 2. This result is consistent with the lack of 20 cm radio continuum emission (Yusef-Zadeh et al. 2004).

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Comparison of the 1.1 mm and Spitzer images shows that there is a slight depression in the spatially averaged 8 and 24 μm emission from the bulge of the Galaxy at the location of the 1.1 mm features in Clump 2. The Clump 2 dust features are weak IRDCs indicating that most of the dust in this region must be somewhat in the foreground of the bulge. This is consistent with the suspected location of the Clump 2 material in the near-side of the central bar.

Two clumps located in Bania's Clumps appear to be foreground objects. The 1.1 mm clump located at [l, b] = [–3.405, +0.877] is associated with IRAS 17470-2533 at the high-latitude end of Clump 2. Located several arcminutes southwest of the clump, there is an arcminute diameter bubble of 8 μm emission which contains a small cluster. The BGPS clump is seen as an IRDC projected in front of the background star field and diffuse emission. Its high latitude and appearance as an IRDC is consistent with this clump and adjacent H ii region being located in the foreground in the Galactic disk. The bright BGPS clump at [l, b] = [+3.34, +0.44] is associated with IRAS 17484-2550 and the clump CO003.34+00.44 in the Galactic center survey of Oka et al. (1998, 2001b). The radial velocity ranges from 0 to 40 km s⁻¹. A low-contrast IRDC is associated with this feature.

Several prominent 1.1 mm clumps are located at relatively high Galactic latitudes |b|> 025. Some of these clumps are associated with prominent H ii regions or bubbles visible in the IRAC images. The Spitzer 8.0 μm images (Figure 10) reveal a prominent bubble centered on the 1.1 mm clump located at [l, b] = [041, −050] (Figure 9). In the 8.0 μm Spitzer image, the bubble appears as a nearly circular ring of emission with a radius of about 7' (∼17 pc) with a dark region centered on the 1.1 mm clump. The 1.1 mm clump is located at the apex of a cometary cloud seen at 8 μm that projects from a more negative latitude toward the center of the bubble (Figures 9 and 10). The bubble is rimmed by several 1.1 mm clumps, the brightest of which is located at [l, b] = [028, –048]. Several additional 1.1 mm clumps are located at larger projected radii from the center of the bubble along its rim. Figure 11 shows that the bubble interior is lit up with faint 20 cm radio continuum emission. This region is reminiscent of the IC 1396 bubble in Cepheus with its intruding tongue of molecular gas.

Several other relatively low-latitude 1.1 mm clumps such as [l, b] = [359.71, −0.37], [l, b] = [359.91, −0.31], and [l, b] = [359.97, −0.46] form a chain below the bulk of dust associated with the Galactic center. This chain of clouds and those associated with the Spitzer ring have been recently noted by Nagayama et al. (2009) who identified them in near-IR extinction maps as well as in CO data. These authors estimate distances using the cumulative numbers of stars and reddening to determine distances and column densities. The distance to this chain of clouds located about 02–04 below the Galactic plane ranges from 3.6 to 4.2 kpc. K-band extinctions range from 0.4 to 1.0 mag. In Table 2, we assume a uniform distance of 3.9 kpc for mass estimation.

Another high-latitude region is associated with the compact cluster of 1.1 mm clumps listed in Table 1 and in the figures as [l, b] = [359.96, +0.17]. The Spitzer image (Figure 10) shows bright 8.0 μm emission at the high-longitude (northeast) side of the cluster of BGPS clumps. This region may be a foreground H ii region complex exhibiting a "champagne flow" toward high Galactic longitudes from a molecular cloud containing several dense clumps.

The clumps at [l, b] = [0.53, +0.18] and [l, b] = [0.84, +0.18] are located well away from the dense-gas layer of the CMZ which suggests that they may be foreground clumps in the Galactic plane far from the Galactic center region. Thus, for these objects, as well as the cluster of clumps near [l, b] = [359.96, +0.17], we assume a distance of 3.9 kpc.

Millimeter and submillimeter imaging of the CMZ reveals a network of dust filaments, shells, and dense clumps. The ∼30 pc diameter infrared bubble between Sgr A and l≈ 02 is associated with a 1.1 mm and 350 μm cavity that contains the Quintuplet and Arches clusters and Sgr A. Several chains of clumps in the bubble interior point toward Sgr A. The high combined luminosity of the Arches and Quintuplet clusters and of the circumnuclear cluster surrounding Sgr A* can plausibly produce the cavity. Below, the physical parameters of this feature are estimated.

The central few hundred parsecs of the Galaxy contain at least four different types of gas reservoir: cold (T <10² K) and dense molecular clouds (n(H₂) > 10⁴ cm⁻³) that comprise the CMZ (Morris & Serabyn 1996; Oka et al. 1998), a warm inter-cloud medium (T ≈ 10²–10³ K) traced by species such as H⁺₃ in absorption (Oka et al. 1998, 2005; Goto et al. 2008) and diffuse ¹²CO and CI in emission (Martin et al. 2004), a warm ionized component (T ≈ 10⁴ K) traced by recombination and fine structure lines, radio scattering (Lazio & Cordes 1998), and free–free emission and absorption, and an ultra-hot component (T >10⁶ K) traced by X-rays (Wang et al. 2002).

The volume filling factor of the dense, gravitationally bound cold molecular phase is in the range 0.01 (Cotera et al. 2000) to 0.1(Morris & Serabyn 1996). The remaining volume is either filled with gravitationally unbound (to individual clouds or star clusters) warm neutral atomic and molecular, photoionized, or X-ray-emitting gas (Martin et al. 2004; Oka et al. 2005; Goto et al. 2008). Abundant 3–4 μm H⁺₃ absorption toward massive stars in the central 02 of the Galaxy indicates that warm, but mostly neutral gas unbound to any particular cloud may occupy a large fraction of the volume in the CMZ. However, the large angular diameters of background extra-galactic radio sources produced by scattering of their radio continuum emission by a foreground screen of electrons (Lazio & Cordes 1998) observed toward the CMZ indicates the presence of relatively dense H ii regions toward most lines of sight through the CMZ. Yusef-Zadeh et al. (2007) show that 6.4 keV Fe Kα emission fills the projected area of the cavity delineated by the Spitzer 24 μm bubble and the 1.1 mm cometary clouds (see their Figure 13). However, the strongest 6.4 keV emission appears to be associated with the molecular clouds traced by the 1.1 mm continuum. These authors suggest that the X-ray emission is produced by the interaction of low-energy cosmic rays with dense clouds.

We consider a scenario in which UV radiation and winds emerging from the Arches, Quintuplet, and Sgr A* clusters and the central black hole are responsible for carving out the cavity. The massive Quintuplet and Arches clusters may have contributed to formation of the individual cells in which these clusters are embedded. Toward positive longitudes the outer boundary of the cavity is near l≈ 020 at a projected distance of about 30 pc from Sgr A* (Figure 9). As shown in Figure 12, the high longitude end of the cavity is marked by a chain of small dust clouds located around l = 022 (northeast of the Quintuplet cluster marked with a "Q" and below the bright 1.1 mm cloud at [026,+003]). The cell associated with the Quintuplet cluster may extend to l = 028, or about 40 pc from Sgr A. Below (to the right of) Sgr A, the intense emission from the 20 and 50 km s⁻¹ clouds masks the cavity if it extends this far. The 20 and 50 km s⁻¹ clouds are physically close to Sgr A; they may block the penetration of UV irradiation and winds. The Spitzer 24 μm image shows that the low-longitude end of the cavity is near l = 0°.

An order-of-magnitude estimate of the parameters of an ionized cavity can be obtained from its size and the 20 cm radio continuum flux, which imply that the average electron density in the cavity cannot be much above 100 cm⁻³. Assuming that the cavity interior is ionized and has a uniform hydrogen density, the Lyman continuum luminosity (in units of the number of ionizing photons emitted per second) required to ionize the cavity is L(LyC) = 3 × 10⁵⁰n²₁₀₀r³₁₀ where n₁₀₀ is the average ionized hydrogen density in units of 100 cm⁻³ and r₁₀ is the mean radius of the cavity in units of 10 pc. The above parameters imply an ionized mass of about $M_{{\rm H}\,{\mathsc{ii}}} \approx 1.4 \times 10^4 n_{100} r^3_{10}$ solar masses around the Arches, Quintuplet, and Sgr A.

The Lyman continuum luminosity of the Sgr A region has been estimated from the flux of free–free radio continuum to be around L(LyC) ∼ few × 10⁵¹ ionizing photons per second. The entire CMZ has a Lyman continuum luminosity of around (1–3) × 10⁵² ionizing photons per second (Morris & Serabyn 1996). This radiation field is produced mostly by young massive stars with a possible contribution from the central black hole in the immediate vicinity of Sgr A. Dust within the cavity and the CMZ may absorb a substantial fraction of the Lyman continuum radiation. However, even if nearly 90% is absorbed, the UV radiation field is sufficient to maintain the central cavity in an ionized state for a mean density of order 10²–10³ cm⁻³.

The orbital time scale in the center of the Galaxy sets a constraint on the age of the cavity and its cometary clouds. Using a typical orbit velocity V_o = 120 km s⁻¹, a test particle in a circular orbit would have an orbital period of t_o = 2πr/V_o∼ 0.5 r₁₀V⁻¹₁₂₀ Myr where V₁₂₀ is the orbit velocity in units of 120 km s⁻¹ and r₁₀ is the distance from the Galactic center in units of 10 pc. Portions of clouds at different radii that are gravitationally unbound from each other would be sheared on a time scale comparable to the orbit time. The presence of recognizable cometary clouds implies that they have been ionized on a time scale shorter than the orbit time. On the other hand, dense clouds are expected to orbit on largely radial trajectories. Thus, the observed cometary features may be objects that have only recently fallen into the interior of the cavity with a speed comparable to the orbit speed. In this case, the elongation of these features away from Sgr A may be a result of the survival of dense gas and dust in the region shielded by the dense head facing Sgr A. Alternatively, this cavity may be only recently ionized by the Arches and Quintuplet clusters. If these clusters have high space velocities, they may have completed several orbits around the nucleus during their lifetimes (Stolte et al. 2008) and the cometary clouds in their vicinity may be a consequence of this motion.

The formation mechanism of massive stars within the central parsec of the Galaxy has been a long-standing mystery. Wardle & Yusef-Zadeh (2008) present a model in which a GMC passes through the center. Gravitational focusing of trajectories by the central black hole and the central stellar cluster of older stars causes gas to collide and dissipate angular momentum and kinetic energy in the wake. This results in the trapping of some gas in eccentric, circumnuclear orbits, forming a disk of compressed material. Gravitational instabilities in the disk can result in the formation of massive stars.

It has been long suspected that the two most prominent molecular clouds near Sgr A, the 20 and 50 km s⁻¹ clouds, are located within the central 30 pc (Morris & Serabyn 1996; Coil & Ho 1999; Wright et al. 2001; McGary et al. 2001; Christopher et al. 2005; Lee et al. 2008). High-dipole moment molecules such as HCN and CS show that they are among the densest and warmest clouds in the CMZ (Miyazaki & Tsuboi 2000). Together, these two clouds give the CMZ its apparently small (few arcminute) scale height near Sgr A.

The 50 km s⁻¹ cloud is located about 10 pc in projection from the CND toward higher longitudes. Several 350 μm filaments connect the 50 km s⁻¹ cloud to the CND (Figure 12); these features correspond to ammonia streamers seen by McGary et al. (2001). Lee et al. (2008) present a three-dimensional model based on near-infrared observations of H₂ and argue that the 2 pc radius CND was recently fueled by the passage of the 50 km s⁻¹ cloud within a few pc of the nucleus. Assuming that the 50 km s⁻¹ cloud is on an x2 orbit, the projected separation of about 10 pc from Sgr A* implies that for a proper motion of 100 km s⁻¹, it passed by the nucleus about 10⁵ years ago. The passage of the 50 km s⁻¹ cloud may have formed the CND, leading to the current episode of massive star formation activity within its clumps.

The 20 km s⁻¹ cloud is the most prominent object at 350 μm in the inner square degree of the Galactic center (Figures 12 and 14); it is located about 7–20 pc from the nucleus in projection. It is elongated with a length-to-width ratio of about 5–10:1 with a major axis that points about to 2' below the CND (Figure 12). Table 2 shows that for an assumed grain temperature of 20 K, the 20 km s⁻¹ cloud is the second most massive and densest cloud in the CMZ (Sgr B2 is the most massive and dense). A bright filament of 350 μm and 1.1 mm dust emission connects the 20 km s⁻¹ cloud to the CND; this feature is likely to be the same as the one detected in ammonia by Coil & Ho (1999). This structure provides support for the model in which CND continues to be actively fueled, currently by flow from the 20 km s⁻¹ cloud (Coil & Ho 1999).

If the center of the 20 km s⁻¹ cloud is currently located about 20 pc from the nucleus, and on a near collision course with the nuclear cluster, it may provide fuel for further star formation in the nucleus in the near future. Assuming that this cloud is on an x2 orbit (Binney et al. 1991; Marshall et al. 2008; Rodriguez-Fernandez & Combes 2008) whose velocity is mostly orthogonal to our line of sight, and has speed of about 100 km s⁻¹, it will come closest to the nuclear region in a few hundred thousand years.

The IRS 16 cluster of massive stars surrounding Sgr A* could not have formed from any currently visible cloud. The IRS 16 cluster has an age of at least several Myr; Tamblyn & Rieke (1993) give an age of 7–8 Myr. The formation of a previous circumnuclear ring would have required the passage of a cloud about 5–10 Myr ago. Any such cloud would have moved too far from the nucleus to have a clear connection to it and may have been disrupted by a combination of tidal forces, UV radiation, and winds.

Figure 19 shows the 1.1 mm image of the region containing Bania's Clump 2. Figure 20 shows this image superimposed on the Spitzer Space Telescope IRAC image. These images show that despite the presence of dozens of clumps emitting brightly at 1.1 mm, Clump 2 lacks significant mid-infrared emission from either extended or point sources. It is not forming stars at a rate comparable to similar clumps elsewhere in the Galaxy. In the Galactic plane, H ii regions and nebulae excited by young, massive stars are frequently seen at 3.6–5.8 μm. The 8 μm Spitzer images trace O, B, and A stars by the presence of bubbles rimmed by bright emission from UV-excited polycyclic aromatic hydrocarbon (PAH) molecules and nano-particles. Few of these signatures of massive stars are observed toward Bania's Clump 2. The longer wavelength Spitzer images (e.g., 24 μm) trace highly embedded young stars and clusters. Inspection of the 24 μm Spitzer images shows that Bania's Clump 2 also lacks a substantial population of embedded sources visible at 24 μm. Perhaps the clumps are embedded in a medium having an unusually high pressure or large non-thermal motions which suppress gravitational instabilities on the mass scale of stars or even clusters. Many of the Clump 2 features are seen in silhouette against background star light and diffuse emission from the ISM in the 3.6–8 μm Spitzer images, indicating that they are in front of much of the central bulge.

The absence of emission in the Spitzer images is not the result of high extinction. Regions having similar amounts of 1.1 mm emission closer to the mid-plane of the Galaxy show a factor of two to four times more emission between 6 and 24 μm. Furthermore, the brightest 1.1 mm emission maxima toward Bania's Clumps 2 imply localized extinctions of A_V< 10 mag averaged over a 40'' aperture. However, the low spatial frequency emission resolved out by BPGS may amount to as much as A_V = 30 mag based on the line-of-sight column density of CO. But, this would be insufficient to completely hide all Spitzer mid-IR emission.

The exceptionally large line widths and velocity extent of this complex may be caused by three processes. First, the abrupt change in the orientation of the orbital velocity where a spur or inner dust lane encounters the ridge of dense gas at the leading edge of the Galactic bar. Second, at the ends of the so-called x1 orbits in a tri-axial potential where they become self-intersecting (Contopoulos & Papayannopoulos 1980; Binney et al. 1991; Marshall et al. 2008; Rodriguez-Fernandez & Combes 2008). In either mechanism cloud–cloud and cloud–ISM collisions result in high ram pressures, shocks, and turbulence that can produce large line widths and non-thermal motions. Third, broad-line features such as Bania's Clump 2 may consist of unrelated clouds along an x1 orbit which has a leg aligned with our line of sight. In this scenario, Bania's Clump 2 traces gas in a dust lane at the leading edge of the bar which is elongated along our line of sight.

Figure 21 shows a cartoon of the possible configuration of the central regions of the Milky Way as seen from the north Galactic pole. The x1 orbits are elongated along the major axis of the Galactic bar in a frame of reference rotating with the bar. As a family of orbits with decreasing semimajor axes (corresponding to lower angular momenta about the nucleus) are considered, the ends first become more "pointed," then they become self-intersecting. Clouds whose orbits decay eventually enter the regions of self-intersection where they may encounter other clouds or the lower density ISM on the same trajectory with supersonic velocities. The resulting shock waves and consequent dissipation of orbital energy causes the clouds to rapidly migrate onto the smaller "x2" orbits whose major axes are orthogonal to the bar. Most of the observed dense gas in the CMZ is thought to occupy the near side of the X2 orbits at positive longitudes where they are seen as IRDCs.

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Models of gas flows in a tri-axial potential can reproduce the major features observed in the longitude–velocity diagrams of molecular tracers of the inner Galaxy (Bissantz et al. 2003; Rodriguez-Fernandez et al. 2006; Rodriguez-Fernandez & Combes 2008; Rodriguez-Fernandez 2009; Pohl et al. 2008; Liszt 2009). It is suspected that the major axis of the central bar is currently oriented within 15°–45° of our line of sight. The conventional interpretation of the Galactic center molecular line emission places the high-velocity rhombus evident in l–V diagrams in the "180 pc molecular ring" (sometimes called the EMR—see Figure 4 in Morris and Serabyn 1996). This feature is tilted with respect to the Galactic plane so that the positive latitude part is seen at positive velocities. In Figure 4, this feature is labeled as the "leading edge of the bar" since in most models, the gas that is transiting from x1 to x2 orbits is found in such a dust lane (see Figure 5 in Morris and Serabyn 1996).

Bania's Clump 2 is located at higher longitudes than the 180 pc molecular ring. Thus, it is either located where the x1 orbits become self-intersecting, or where a spur encounters the dust lane at the leading edge of the bar, or traces gas in the dust lane which is oriented nearly parallel to our line of sight. Rodriguez-Fernandez et al. (2006) show the likely location of Bania's Clump 2 in a hypothetical face-on view of the Galaxy (see their Figure 2). The superposition along the line of sight of clouds moving near the apex of the innermost, self-intersecting x1 orbits can exhibit a large range of radial velocities within a very restricted spatial region. This phenomenon is caused by the near co-location of clouds moving toward and away from the Galactic center. Some of these clouds will be colliding, resulting in compression, heating, and turbulence generation. While low-velocity cloud–cloud collisions can trigger star formation, high-velocity collisions are likely to be disruptive; sufficiently fast collisions dissociate molecules and raise post-shock temperatures.

The presence of 1.1 mm clumps and associated gas seen in high-dipole moment molecules such as CS and HCN suggests that the smallest mass that can become gravitationally unstable is comparable to the observed clump mass of 10³–10⁴ M_☉. The low rate of star formation in these clumps implies that fragmentation is suppressed in this environment, possibly by shear or large fluxes of energy and momentum from large scales to small.

If cloud collisions and supersonic internal motions suppress star formation, the dissipation of even a small fraction of this energy must emerge as radiation, resulting in the excitation of fine structure cooling lines such as the 63 μm [O i], 157 μm [C ii], and 205 μm [N ii], or high-J lines of CO and OH. A substantial portion of the internal energy may also go into heating grains which would emerge as submillimeter continuum radiation. The Herschel Space Observatory may be able detect the resulting fluxes.

While such shocks may promote star formation in other environments such as in Sgr B2, they have failed to do so here. Perhaps the millimeter sources in Bania's Clump 2 are in an earlier, pre-stellar stage of evolution. Clumps could evolve to denser configurations capable of star formation in a few crossing times. Using a typical Galactic center cloud line width of ΔV = 20 km s⁻¹ for each clump, a clump with a radius of r = 1' (∼ 2.5 pc) will evolve on a time scale t ∼ 2R/ΔVf ∼ 2.5 × 10⁵f⁻¹ years where f (= 0.1–1) is the area-filling factor of gas in the clump. Thus, it is possible that we are seeing the Clump 2 objects in an early state of evolution and that they will become active star formers in the next few million years.

The Clump 2 complex has a diameter of about 0.5° (70 pc), implying R₂ ∼ 2.3 × 10²⁰ cm (25 pc), and a velocity extent of about 2V₂ = 200 km s⁻¹. The time scale on which the entire Clump 2 region should evolve is about t₂ ∼ 2R₂/ΔV₂f₂ ∼ 3.6 × 10⁵f⁻¹ years. Figure 19 implies an area-filling factor of about f₂∼ 0.1 for the 1.1 mm clumps, so the evolutionary time scale for the entire complex is a few Myr, comparable to evolutionary time scales of individual clumps. Both of these time scales are short compared to the ∼ 15 Myr orbital time scale at the projected separation between Sgr A* and Clump 2.