------------------------------------------------------------------------ gc_sio.tex submitted to A&A Main Journal Content-Type: text/plain; charset=us-ascii; name="gc_sio.tex" Content-Transfer-Encoding: 7bit Content-Disposition: inline; filename="gc_sio.tex" \documentstyle[psfig_laa]{laa} %\documentstyle[referee,psfig_laa]{l-aa} %\renewcommand{\baselinestretch}{1.3} % \begin{document} % \thesaurus{09.03.1,09.13.2,09.19.1,10.03.1,13.19.3} % % Abbreviations % \def\etal {et al.} \def\ie {i.\,e.} \def\eg {e.\,g.} \def\etc {etc.} \def\ccpers {\hbox{${\rm cm}^3{\rm s}^{-1}$}} \def\DEGR {\hbox{$^{\circ }$}} \def\vlsr {\hbox{${v_{\rm LSR}}$}} \def\vel {\hbox{${v_{\rm LSR}}$}} \def\vhel {\hbox{${v_{\rm HEL}}$}} \def\delv {\hbox{$\Delta v_{1/2}$}} \def\dvel {\hbox{$\Delta v_{1/2}$}} \def\TL {$T_{\rm L}$} \def\TC {$T_{\rm c}$} \def\TEX {$T_{\rm ex}$} \def\TMB {$T_{\rm MB}$} \def\TKIN {$T_{\rm kin}$} \def\TREC {$T_{\rm rec}$} \def\TSYS {$T_{\rm sys}$} \def\TVIB {$T_{\rm vib}$} \def\TROT {$T_{\rm rot}$} \def\TDUST {$T_{\rm d}$} \def\TASTAR {$T_{\rm A}^{*}$} \def\TVIBST {$T_{\rm vib}^*$} \def\TB {$T_{\rm B}$} % \def\RZWCO {${\cal R}_{2/1}^{\rm C^{18}O}$} \def\RDRCO {${\cal R}_{3/2}^{\rm C^{18}O}$} \def\RMSIO {${\cal R}_{5/2}^{\rm ^{28}SiO}$} \def\RISOSIO {${ r }_{28/29}^{\rm SiO}$} \def\RISOZSIO {${ r}_{29/30}^{\rm SiO}$} % \def\solmass {\hbox{M$_{\odot}$}} \def\solum {\hbox{L$_{\odot}$}} \def\irlum {\hbox{$L_{\rm IR}$}} \def\ohlum {\hbox{$L_{\rm OH}$}} \def\blum {\hbox{$L_{\rm B}$}} \def\numd {\hbox{$n\,({\rm H}_2$)}} \def\kms {\hbox{${\rm km\,s}^{-1}$}} \def\kmsyr {\hbox{${\rm km\,s}^{-1}\,{\rm yr}^{-1}$}} \def\kmsmpc {\hbox{${\rm km\,s}^{-1}\,{\rm Mpc}^{-1}$}} \def\Kkms {\hbox{${\rm K\,km\,s}^{-1}$}} \def\percc {$\hbox{{\rm cm}}^{-3}$} %cm-3 \def\cmsq {$\hbox{{\rm cm}}^{-2}$} %cm-2 \def\cmsix {$\hbox{{\rm cm}}^{-6}$} %cm-6 \def\arcsec {\hbox{$^{\prime\prime}$}} \def\arcmin {\hbox{$^{\prime}$}} \def\ffam {\hbox{$\,.\!\!^{\prime}$}} \def\ffas {\hbox{$\,.\!\!^{\prime\prime}$}} \def\ffM {\hbox{$\,.\!\!\!^{\rm M}$}} \def\ffm {\hbox{$\,.\!\!\!^{\rm m}$}} \def\ffd {\hbox{$\,.\!\!^{\circ}$}} \def\HI {\hbox{HI}} \def\HII {\hbox{HII}} % % molecules % \def\MOLH {\hbox{${\rm H}_2$}} %H2 \def\AMM {\hbox{${\rm NH}_{3}$}} %NH3 \def\TCO {\hbox{${\rm ^{12}CO}$}} %12CO \def\CEIO {\hbox{${\rm C}^{18}{\rm O}$}} %C18O \def\CSEO {\hbox{${\rm C}^{17}{\rm O}$}} %C17O \def\CTHFOS {\hbox{${\rm C}^{34}{\rm S}$}} %C34S \def\THCO {\hbox{$^{13}{\rm CO}$}} %13CO % % \title{Molecular gas in the Galactic center region} \subtitle{III. Probing shocks in molecular cores\thanks{Based on observations obtained at the Swedish-ESO Submillimeter Telescope (SEST, Project C-0518, 1996), the James Clark Maxwell Telescope (JCMT, operated by the Joint Astronomy Centre on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada) and the Heinrich-Hertz-Telescope (HHT, operated by the the Submillimeter Telescope Observatory).}} % \author{S.~H\"{u}ttemeister\inst{1,2}, G.~Dahmen\inst{3,4}, R.~Mauersberger\inst{5}, C.~Henkel\inst{4}, T.L.~Wilson\inst{4,6}, J.~Mart\'{\i}n-Pintado\inst{7} } % \offprints{S. H{\"u}ttemeister, RAIUB} % \institute{ Radioastronomisches Institut der Universit\"{a}t Bonn, Auf dem H\"{u}gel 71, D - 53121 Bonn, Germany \and Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, U.S.A. \and Physics Department, Queen Mary \& Westfield College, University of London, Mile End Road, London E1 4NS, England \and Max-Planck-Institut f{\"u}r Radioastronomie, Auf dem H{\"u}gel 69, D - 53121 Bonn, Germany \and Steward Observatory, The University of Arizona, Tucson, AZ 85721, U.S.A. \and Submillimeter Telescope Observatory, The University of Arizona, Tucson, AZ 85721, U.S.A. \and Centro Astron{\'o}mico de Yebes, Apartado 148, 19080 Guadalajara, Spain } % \date{Received / Accepted} % \maketitle % \markboth{S.~H\"uttemeister et al.: Molecular Gas in the GC Region III.~Probing shocked molecular gas}{} % \begin{abstract} Multiline observations of C$^{18}$O and SiO isotopomers toward 33 molecular peaks in the Galactic center region, taken at the SEST, JCMT and HHT telescopes, are presented. The C$^{18}$O presumably traces the total H$_2$ column density, while the SiO traces gas affected by shocks and high temperature chemistry. The $J =2\to 1$ line of SiO is seen only in few regions of the Galactic disk. This line is easily detected in all Galactic center sources observed. A comparison of the strength of the rare isotopomers $^{29}$SiO and $^{30}$SiO to the strength of the main isotopomer $^{28}$SiO implies that the $J = 2 \to 1$ transition of $^{28}$SiO, is optically thick. The isotope ratio of $^{29}$Si$/^{30}$Si of 1.6 in the Galactic center clouds is consistent with the terrestrial value. Large Velocity Gradient models show that the dense component ($n_{\rm H_2} \geq 10^4$\,\percc) in typical molecular cores in the Galactic center is cool (\TKIN{} $\approx 25$\,K), contrary to what is usually found in Giant Molecular Clouds in the disk, where the densest cores are the hottest. High kinetic temperatures, $> 100$\,K, known to exist from NH$_3$ studies, are only present at lower gas densities of a few 10$^3$\,\percc , where SiO is highly subthermally excited. Assuming that \CEIO\ traces all of the molecular gas, it is found that in all cases but one, SiO emission is compatible with arising in gas at higher density that is (presently) relatively cool. The relative abundance of SiO is typically $10^{-9}$, but differs significantly between individual sources. It shows a dependence on the position of the source within the Galactic center region. High abundances are found for the locations bar potential models predict for preferred collision regions. These results can be used to relate the amount of gas that has encountered shocks within the last $\sim 10^6$ years to the large scale kinematics in the inner $\sim$\,500\,pc of the Galaxy. \end{abstract} % \keywords{ISM: clouds, molecules, structure -- Galaxy: center -- Radio Lines: ISM } % \section{Introduction} The molecular environment in inner $\sim 8^{\circ}$ of the Galaxy differs drastically from that in the Galactic disk. Large scale surveys of CO isotopomers (e.g.\ Bally et al.\ 1987, Heiligman 1987, Jackson et al.\ 1996, Dahmen et al.\ 1997a (Paper\,I), Bitran et al.\ 1997) show that the cloud and the intercloud medium in the gaseous bulge of the Milky Way are molecular. The gas is characterized by large linewidths, indicating a high degree of turbulence. As shown by (Dahmen et al.\ 1998 (Paper\,II) ), there is no simple relationship between CO line intensities and H$_2$ column densities densities and thus no global $N({\rm H}_2)/I$(CO) conversion factor. Except for a few exceptional regions such as Sgr\,A and Sgr\,B2, little evidence for ongoing massive star formation in Galactic center Giant Molecular Clouds (GMCs) is found, as is demonstrated by a general lack of strong FIR or radio continuum point sources associated with these clouds (Odenwald \& Fazio 1984, G\"usten 1989). Ambient dust temperatures are fairly low at \TDUST\ $\sim 25 - 30$\,K (e.g.\ Cox \& Laurijs 1989). Based on ISO data, this is also a typical \TDUST\ in the sources studied in this paper. In the Galactic disk, quasithermal SiO emission is tightly correlated with high temperature regions (e.g.\ Ziurys et al.\ 1989). A close association with outflows strongly suggests that grain disruption by shocks is the major mechanism for releasing SiO into the gas phase (Mart\'{\i}n-Pintado et al.\ 1992), although high temperature gas phase chemistry (Langer \& Glassgold 1990) may play a minor role. In the Galactic center region, SiO is much more widespread (see e.g.\ the survey in the $J=1\to 0$ transition by Mart\'{\i}n-Pintado et al.\ 1997), which can was interpreted as evidence for large scale or ubiquitous (fast) shocks. In this paper, \CEIO\ and SiO data are presented for molecular cloud cores of the Bally et al.\ (1987) CS survey. Our measurements allow us to (1) trace the H$_2$ column density, (2) determine or constrain density and temperature structure, (3) estimate SiO abundances and (4) obtain information about silicon isotope ratios. Since our sources are selected on the basis of their intensity in CS, a general high density tracer, the sample is not a priori biased toward strong SiO emission. NH$_3$ data (H\"uttemeister et al.\ 1993b) show that at least two phases of different kinetic temperatures are present within all cloud cores without massive star formation: A cool component with \TKIN\ of 20\,--\,30\,K, close to the temperature of the dust in the Galactic center and a warm component with \TKIN\,$\geq 120$\,K. \TKIN\,$\sim 75$\,K is considered by many authors, (e.g.\ the reviews by Morris \& Serabyn 1996, Mezger et al.\ 1996) as `typical'. This is, however, just the average over the hot and cool component and has no meaning as a distinct physical component. A major aim of this study is to determine the physical parameters, distribution and origin of these phases and to decide which of them, cool or warm, is associated with the bulk of the SiO emission. \section{Observations} The observations of the $J = 1 \to 0$ and $2\to 1$ transitions of \CEIO\ and the $J = 2\to 1$ transition of $^{28}$SiO, $^{29}$SiO and $^{30}$SiO were carried out at the Swedish-ESO Submillimetre Telescope (SEST) at La Silla (Chile). All C$^{18}$O($3 \to 2$) data were taken at the James Clark Maxwell Telescope (JCMT) on Mauna Kea, (Hawaii), where we also obtained some $^{28}$SiO($5 \to 4$) and \CEIO ($2 \to 1$) measurements for calibration purposes. Finally, the $5 \to 4$ transition of $^{28}$SiO was measured at the Heinrich Hertz Telescope (HHT) on Mt.\ Graham (Arizona). Table \ref{log} summarizes the observations, including line frequencies, telescope parameters, receivers and spectrometers used and typical system temperatures. A telescope beamsize of 46$''$ corresponds to 1.9\,pc at a distance of the Galactic center of 8.5\,kpc. \begin{table*} \begin{flushleft} \caption{\label{log} Summary of the observations } \begin{tabular}{lrrrrrr|rrr|r} Line & \multicolumn{1}{c}{Frequency} & \multicolumn{1}{c}{Telescope} & \multicolumn{1}{c}{Receiver$^{\rm a)}$} & \multicolumn{1}{c}{\TSYS $^{\rm b)\multicolumn{1}{c}{$\theta_{\rm B}$} & \multicolumn{1}{c}{$\eta_{\rm MB}$} & \multicolumn{3}{c}{Spectrometer} & \multicolumn{1}{c}{Date} \\ & \multicolumn{1}{c}{} & & & \multicolumn{1}{c}{} & \multicolumn{1}{c}{} & & \multicolumn{1}{c}{Type$^{\rm c)}$} & \multicolumn{1}{c}{Bandwidth} & \multicolumn{1}{c}{Resolution} & \\ & \multicolumn{1}{c}{GHz} & & & \multicolumn{1}{c}{K} & \multicolumn{1}{c}{$''$} & & \multicolumn{1}{c}{} & \multicolumn{1}{c}{MHz} & \multicolumn{1}{c}{kHz} & \\ \hline C$^{18}$O($1\to 0$) & 109.782160 & SEST & Schottky SSB & 400 & 46 & 0.72 & AOS & 86 & 80 & 9/90,7/92 \\ C$^{18}$O($2\to 1$) & 219.560319 & SEST & SIS SSB & 700--1100 & 24 & 0.55 & AOS & 995 & 1400 & 7\&9/92 \\ & & JCMT$^{\rm d)}$ & SIS DSB & 600 & 20 & 0.69 & DAS & 250 & 189 & 4/96 \\ C$^{18}$O($3\to 2$) & 329.350500 & JCMT & SIS DSB & 1600--1800 & 14 & 0.58 & DAS & 500 & 378 & 4/96 \\ $^{28}$SiO($2\to 1$) & 86.846998 & SEST & Schottky SSB & 350--400 & 57 & 0.75 & AOS & 86 & 80 & 10/94 \\ $^{28}$SiO($5\to 4$) & 217.104935 & HHT & SIS DSB & 700 & 33 & 0.83 & AOS & 990 & 1000 & 4/97 \\ & & JCMT$^{\rm e)}$ & SIS DSB & 600 & 20 & 0.69 & DAS & 250 & 189 & 4/96 \\ $^{29}$SiO($2\to 1$) & 85.759132 & SEST & SIS SSB & 140 & 57 & 0.75 & AOS & 86/995 & 80/1400 & 3/96 \\ $^{30}$SiO($2\to 1$) & 84.746036 & SEST & SIS SSB & 140 & 57 & 0.75 & AOS & 86/995 & 80/1400 & 3/96 \\ \hline \\ \end{tabular} \\ a): SSB: Single sideband receiver, with a sideband rejection of $\sim 20$\,dB ; DSB: Double sideband receiver \\ b): on a \TMB\ scale \\ c): AOS: Acousto-Optical Spectrometer; DAS: Dutch Autocorrelation Spectrometer \\ d): Only three sources were observed for cross-telescope calibration with SEST \\ e): Only two sources were observed for cross-telescope calibration with the HHT \end{flushleft} \end{table*} Typically, the telescope pointing was checked every three hours. At SEST, two strong SiO masers close to the Galactic center, AH\,Sco and VX\,Sgr, were used. At the JCMT, the pointing was done using five-point continuum maps on NGC\,6334\,I. At the HHT, continuum cross scans on Sgr\,B2 were used. At SEST and JCMT, the pointing accuracy was always better than $5''$. At the HHT, the accuracy was always better than $10''$ and usually $5''$. At all telescopes, calibration was done with the standard chopper wheel method, giving temperatures on the $T_A^*$ scale. These were converted to \TMB , the main beam brightness temperature (see, e.g., Downes 1989 and Rohlfs \& Wilson 1996 for definitions), using the main beam efficiencies, $\eta_{\rm MB}$, listed in Table\,\ref{log}. To account for structure in our extended sources, we mapped the beam used to measure the \CEIO ($1\to 0$) line in the ($2\to 1$) and ($3\to 2$) transitions. Due to time restrictions, this was not possible for the SiO($5\to 4$) line observed at the HHT. From maps obtained at the JCMT for two sources in SiO($5\to 4$), the intensity in a 33$''$ beam is higher by 8\% and 15\% than in 46$''$ and 57$''$ beams, respectively. This is caused by the structure of the individual sources, and the numbers are, therefore, uncertain. We will, however, use these as correction factors when comparing HHT SiO($5\to 4$) data to measurements obtained with larger beamsizes. To adjust the calibration scales of the three telescopes, we have observed a number of sources in the 1.3\,mm band which is accessible to all instruments (see Table\,\ref{log}). For \CEIO ($2\to 1$), the integrated intensities at the 15\,m-telescopes SEST and JCMT agree to better than 10\%, i.e.\ to within the uncertainty of the calibration of the individual measurement. The SiO($5\to 4$) JCMT maps, convolved to the larger size of the HHT beam, also give intensities that are consistent to within 10\% with the HHT data. Thus, the \TMB -scales are in excellent agreement, and we are confident that our complete data set has been placed on a compatible intensity scale. We carried out all observations employing position switching. This was necessary given the large angular extent of the sources, especially in \CEIO . Offset positions were typically chosen $15'$ (SiO) or $30'$ (\CEIO ) away from the source, perpendicular to the Galactic plane. This observing mode causes baselines that are not always perfectly flat. Thus, polynomial baselines of order $\leq 3$ were subtracted from the data. Baselines were determined over a region of $\pm 150 - 200$\,\kms\ from the line center. For the JCMT data, which were observed in exceptionally good weather ($\tau_{225} < 0.05$ throughout), linear baselines were usually sufficient. For the final analysis, all spectra were smoothed to a velocity resolution of $\sim 2$\,\kms . \section{Results} The 33 sources we study in detail are selected as distinct peaks in the CS survey of Bally et al.\ (1987). They are, therefore, cores that should have large amounts of dense gas, which are embedded in a smooth, lower intensity molecular intercloud medium. Since we wish to derive the properties of typical clouds in the Galactic center environment, the non-typical molecular peaks associated with Sgr\,A and Sgr\,B2 were not included in the sample. \begin{figure*} \psfig{file=spectra.ps,bbllx=0pt,bblly=55pt,bburx=490pt,bbury=700pt,height=21.0cm} \caption{\label{spectra} Sample spectra of eight sources with very complete observations. In the top panels, the \CEIO\ lines are displayed: The ($1\to 0$) transition is drawn as a dotted line, the ($2\to 1$) transition is drawn as a thin solid line and the ($3\to 2$) transition is displayed as a thick solid line. The bottom panels show the SiO spectra. $^{28}$SiO($2\to 1$) is denoted as a thin solid line, $^{28}$SiO($5\to 4$) is given by a thick solid line; the $^{29}$SiO($2\to 1$) spectrum is drawn as a thick dotted line and the $^{30}$SiO($2\to 1$) is shown as a thin dotted line. The intensities of both rare isotopomers were multiplied by 5 to allow clearer comparisons. Black squares overlayed on the map of the overall distribution of the $J=1\to 0$ line of \CEIO\ (taken from Paper\,I) mark all 33 sources we observed except Clump\,1 at $l \sim -5^{\circ}$. } \end{figure*} The ($1\to 0$) and ($2\to 1$) transitions of \CEIO\ and the ($2\to 1$) transition of $^{28}$SiO were observed toward all sources of our sample. Both the \CEIO\ and (more notably) the SiO lines were easily seen in all sources. We observed the 10 strongest \CEIO\ sources in the ($3\to 2$) submillimeter transition of \CEIO , again detecting all. The rare isotopomers $^{29}$SiO and $^{30}$SiO were observed and detected toward 12 clouds showing strong emission in $^{28}$SiO. Of 8 sources measured in the $^{28}$SiO($5\to 4$) line, 7 showed emission. In Fig.\,\ref{spectra} and Table\,\ref{rat}, we give an overview of our results. Sample spectra for eight sources are displayed in Fig.\,\ref{spectra}. We mark the position of all sources observed on a survey map of the ($1\to 0$) transition of \CEIO\ (Papers I and II). Galactic coordinates are used for source names. For the corresponding equatorial coordinates, refer to H\"uttemeister et al.\ (1993b). Most of the clouds we have studied are located within the molecular bulge region, while three lie within the `Clump\,2' complex (Bania 1977, Stark \& Bania 1986) at $l \sim 3^{\circ}$. `Clump\,1' is located $\sim 5^{\circ}$ south of the Galactic center. Integrated line intensities for the \CEIO ($2\to 1$) and $^{28}$SiO($2\to 1$) transitions and line intensity ratios derived from integrated intensities are presented in Table\,\ref{rat}. The errors are determined from the (usually small) formal error of a Gaussian fit to the lines or, in case of non-Gaussian lineshapes, from the rms noise in the spectrum. In addition, a calibration uncertainty of 10\% for data taken with the same telescope or 20\% for data obtained with different instruments was assumed. For Gaussian lines, the formal error of the fit agrees closely with the error obtained from the rms noise in the spectrum. \begin{table*} \caption{ \label{rat} Integrated intensities of the \CEIO ($2\to 1$) and SiO($2\to 1$) transitions and line intensity ratios of \CEIO\ and SiO.} \begin{tabular}{l|rrll|rrlll} Source & $v_{\rm LSR}^{\rm C^{18}O}$ & $I_{(2-1)}^{\rm C^{18}O}$ & ${\cal R}_{2/1}^{\rm C^{18}O}$ & ${\cal R}_{3/2}^{\rm C^{18}O}$ & $v_{\rm LSR}^{\rm SiO}$ & $I_{(2-1)}^{\rm ^{28}SiO}$ & ${\cal R}_{5/2}^{\rm ^{28}SiO}$ & $r_{28/29}^{\rm SiO}$ & $r_{29/30}^{\rm SiO}$ \\ & & & & 10$^{15}$ & & & & & 10$^{13}$ \\ & km/s & K\,km/s & & & km/s & K\,km/s & & & \\ \hline Clump\,1 & 97 & 4.5(0.2) & 1.46(0.27) & & 97 & 0.6(0.2) & & \\ M--0.96+0.13 & 135 & 5.1(0.5) & 1.31(0.24)$^{\rm a)}$ & & 135 & 4.6(0.3) & & & \\ M--0.50--0.03 & --100 & 15.3(0.4) & 1.44(0.16) & 0.49(0.11) & --98 & 4.0(0.3) & & & \\ & --2 & 1.6(0.2) & 1.12(0.31)$^{\rm c)}$ & 0.38(0.11)$^{\rm c)}$ & --2 & $\leq 0.1$ & & & \\ M--0.42--0.01 & --84 & 10.5(0.4) & 1.05(0.12) & & --80 & 3.4(0.2) & & \\ & --4 & 2.2(0.3) & 1.35(0.34) & & --4 & $\leq 0.2$ & & & \\ M--0.39+0.02 & --80 & 13.8(1.2) & 1.89(0.27) & & --79 & 3.3(0.2) & & & \\ M--0.33--0.01 & --47 & 12.8(0.5) & 1.19(0.18) & & --51 & 4.4(0.2) & & & \\ M--0.45--0.09 & --61 & 4.0(0.4) & 0.77(0.25)$^{\rm a)}$ & & --61 & $\leq 0.3$ & -- & -- & \\ & --10 & 11.4(0.5) & 2.09(0.39) & & --12 & 16.7(1.5) & 0.23(0.05) & 15.9(3.1) & \\ M--0.30+0.00 & --33 & 12.1(0.7) & 1.23(0.25)$^{\rm b)}$ & & --38 & 5.0(0.3) & & & \\ M--0.44--0.10 & --15 & 6.1(0.6) & 0.60(0.31)$^{\rm a),b)}$ & & --20 & 19.8(2.0) & & 12.3(1.8) & 1.7(0.3) \\ M--0.51--0.16 & --61 & 27.2(0.8) & 2.67(0.14) & 0.51(0.08) & --57 & 30.6(0.5) & 0.09(0.03) & 11.0(1.4) & 1.6(0.3) \\ & --2 & 1.7(0.2) & 1.33(0.75)$^{\rm c)}$ & 0.41(0.14)$^{\rm c)}$ & --2 & $\leq 0.1$ & -- & -- & -- \\ M--0.50--0.18 & --57 & 14.3(0.7) & 2.40(0.67) & & --55 & 9.3(0.4) & & & \\ M--0.48--0.19 & --50 & 6.9(0.4) & 1.47(0.32)$^{\rm a)}$ & & --53 & 4.0(0.2) & & & \\ M--0.38--0.13 & --33 & 10.4(0.7) & 1.11(0.26) & 0.57(0.12) & --32 & 1.7(0.2) & & & \\ M--0.45--0.19 & 21 & 4.5(0.2) & 1.10(0.19)$^{\rm c)}$ & & 21 & $\leq 0.2$ & & & \\ & 58 & 3.9(0.3) & 0.50(0.31)$^{\rm a)}$ & & 58 & 5.8(0.2) & & & \\ M--0.43--0.21 & --34 & 22.4(0.9) & 2.07(0.27) & & --27 & 6.5(0.3) & & & \\ & 19 & 3.5(0.3) & 1.13(0.21)$^{\rm c)}$ & & 19 & $\leq 0.2$ & & & \\ & 53 & 2.7(0.6) & 0.71(0.22)$^{\rm a)}$ & & 54 & 5.7(0.4) & & & \\ M--0.15--0.07 & 5 & 37.5 (0.4) & 1.40(0.15)$^{\rm b)}$ & 0.75(0.14)$^{\rm b)}$ & 4 & 31.5(1.1) & 0.19(0.04)$^{\rm b)}$ & 11.3(1.3) & 1.6(0.2) \\ M--0.32--0.19 & --32 & 7.3(0.7) & 1.22(0.28)$^{\rm a)}$ & & & & & \\ & 25 & $\leq$3.0 & & & 25 & 2.9(0.2) & & & \\ M+0.24+0.02 & 35 & 41.7(3.7) & 1.06(0.16)$^{\rm b)}$ & 0.51(0.11)$^{\rm b)}$ & 36 & 24.4(0.7) & 0.07(0.02) & 11.4(1.3) & 1.6(0.3) \\ M+0.10--0.12 & 22 & 22.1(0.9) & 1.46(0.21) & & 29 & 11.5(1.2) & & & \\ M+0.43--0.05 & --2 & 4.5(0.3) & 0.95(0.28)$^{\rm a)}$ & & 12 & 11.0(0.8) & & 13.3(2.2) & \\ & 38 & 4.9(0.3) & 0.69(0.15)$^{\rm a)}$ & & & & & \\ M+0.41+0.03 & 30 & 10.5(0.4) & 0.64(0.10)$^{\rm b)}$ & & 21 & 10.4(1.1) & & & \\ M+0.50+0.00 & 30 & 22.9(2.4) & 1.06(0.21)$^{\rm b)}$ & 0.76(0.15)$^{\rm b}$ & 26 & 20.9(0.3) & & 13.1(2.1) & 1.7(0.3) \\ M+0.68--0.12 & 30 & 35.4(0.3) & 1.48(0.17)$^{\rm b)}$ & 0.70(0.14)$^{\rm b)}$ & 20 & 39.2(2.3) & 0.18(0.05)$^{\rm b)}$ & 11.9(1.4) & 1.2(0.2) \\ M+0.83--0.10 & 24 & 13.4(0.8) & 0.81(0.18)$^{\rm a),b)}$ & & 24 & 10.4(0.4) & & & \\ & 99 & 8.7(0.5) & 1.66(0.29) & & 96 & 2.5(0.2) & & & \\ M+1.46+0.27 & 80 & 6.6(0.4) & 0.95(0.32)$^{\rm a)}$ & & 81 & 20.4(0.3) & & 13.8(1.6))$^{\rm b)}$ & 1.5(0.3) \\ M+0.83--0.18 & 41 & 10.6(1.0) & 1.03(0.13) & 0.53(0.12) & 43 & 26.0(0.4) & 0.19(0.04) & \enspace7.4(0.8) & 2.0(0.2) \\ M+1.15--0.09 & --18 & 1.9(0.2) & 1.05(0.20)$^{\rm c)}$ & & --18 & $\leq 0.2$ & & & \\ & 60 & 7.8(1.2) & 0.68(0.23)$^{\rm a),b)}$ & & 67 & 21.0(1.1) & & & \\ M+1.13--0.12 & --15 & 12.9(0.9) & 1.48(0.18)$^{\rm c)}$ & & --15 & 6.4(0.7) & & & \\ & 80 & 1.8(0.3) & 0.47(0.22)$^{\rm a)}$ & & 70 & 17.5(1.2) & & & \\ M+1.31--0.13 & 30 & 11.6(0.6) & 1.95(0.39)$^{\rm a)}$ & 0.25(0.13)$^{\rm a)}$ & 35 & 11.5(0.5) & $\leq 0.04$$^{\rm d)}$ & 20.9(2.3) & 1.5(0.3) \\ & 80 & 7.8(0.6) & 1.58(0.35)$^{\rm a)}$ & 0.78(0.26)$^{\rm a)}$ & 87 & 16.8(0.5) & $\leq 0.04$$^{\rm d)}$ & 18.0(2.7) & 1.4(0.2) \\ M+0.94--0.36 & --49 & 18.5(1.3) & 2.05(0.33)$^{\rm b)}$ & & --41 & 3.4(0.5) & & & \\ M+3.06+0.40 & 12 & 1.6(0.1) & 0.82(0.15)$^{\rm c)}$ & & 12 & $\leq 0.1$ & & \\ & 22 & 2.5(0.3) & 0.61(0.16)$^{\rm a)}$ & & 30 & 9.3(2.1) & & & \\ M+3.06+0.34 & 27 & 12.1(3.6) & 1.49(0.48) & 0.24(0.09) & 27 & 15.2(0.3) & 0.16(0.06) & 10.1(1.2) & 1.9(0.4) \\ M+2.99-0.06 & --13 & 2.6(0.3) & 1.21(0.27) & & --12 & 8.5(0.3) & & & \\ \hline \\ \end{tabular} \\ Abbreviations used throughout the paper: \RZWCO\,$=\,I_{(2-1)}^{\rm C^{18}O}/I_{(1-0)}^{\rm C^{18}O}$, \RDRCO\,$=\,I_{(3-2)}^{\rm C^{18}O}/I_{(2-1)}^{\rm C^{18}O}$, \RMSIO\,$=\,I_{(5-4)}^{\rm ^{28}SiO}/I_{(2-1)}^{\rm ^{28}SiO}$, \RISOSIO\,$=\,I_{(2-1)}^{\rm ^{28}SiO}/I_{(2-1)}^{\rm ^{29}SiO}$, $r_{29/30}^{\rm SiO}\,=\,I_{(2-1)}^{\rm ^{29}SiO}/I_{(2-1)}^{\rm ^{30}SiO}$ \\ a): The result is uncertain, due to either baseline uncertainties or low signal-to-noise ratios \\ b): Non-gaussian lineshape or several blended velocity components; ratios have been calculated by adding channels with emission \\ c): very narrow line, possibly from local ISM \\ d): 3$\sigma$ \\ \end{table*} For a given species, the center velocities and lineshapes of the different transitions always agree, to within the noise. Between SiO and \CEIO , however, there can be significant differences. The sources M+0.50+0.00 and M+0.24+0.02 are good examples: In both cases, the \CEIO\ lines show two distinct peaks, while the SiO transitions are single-peaked. In M+0.50+0.00, the central velocity of the SiO is close to the weaker \CEIO\ line; in M+0.24+0.02 SiO and the stronger \CEIO\ peak agree. It is also noteworthy that narrow \CEIO\ lines close to a \vlsr\ of 0\,\kms , likely to be of local origin, never have a counterpart in SiO. This is illustrated by the source M--0.51--0.16 in Fig.\,\ref{spectra} and demonstrates the unusual nature of the Galactic center sources as compared to Galactic disk clouds. Since the lines of $^{29}$SiO and $^{30}$SiO are very likely optically thin, we can directly check whether the ratio of $^{29}$Si/$^{30}$Si in the Galactic center region agrees with the terrestrial value. We find a line intensity ratio of $1.6 \pm 0.2$ in our sample, in excellent agreement with the terrestrial isotope ratio of 1.5. This indicates that this ratio does not depend on the galactocentric distance (Wilson \& Rood 1994, Penzias 1981). For $^{28}$SiO/$^{29}$SiO, we take the ratio to be the terrestrial value of 20, as suggested by Penzias (1981). Since the $^{28}$SiO($2\to 1$) transition is optically thick, we cannot check this assumption, but it is supported by the largest line ratio \RISOSIO\ (toward the exceptional source M+1.31--0.13, see Section 4.3), which is indeed close to 20. \section{The physical conditions of the gas} \subsection{Determinination of densities and temperatures} A Large Velocity Gradient (hereafter LVG) model was used to estimate physical conditions of the gas, namely density (\numd ), kinetic temperature (\TKIN ) and -- in the case of SiO -- molecular abundance ($X$(SiO)) (see Scoville \& Solomon 1974 and De Jong et al.\ 1975 for the general properties). We used a modified version of the code of Henkel et al.\ (1980). The basic LVG assumption is that a of systematic velocity gradient large compared to random motions and monotonic. This allows us to treat the molecular excitation as a local problem. This is certainly an idealisation for Galactic center clouds. However, an application of an LVG code requires no detailed knowledge of the velocity field. In the case of CO lines, H$_2$ densities determined under the assumption of LVG and microturbulence (the opposite extreme) do not differ by more than a factor of three (White 1977). Our analysis is also based on the assumption that, for a given molecule, all transitions observed arise in the same volume. Since the excitation depends on critical density and optical depth, and thus is not identical for different transitions or different isotopomers, this assumption may not be strictly correct, especially if the medium is very clumpy (see, e.g., the discussion in Oka et al.\ 1997 for $^{12}$CO). For \CEIO , this is not a critical problem since the LVG calculations show that all transitions measured by us are optically thin, or, in the low temperature/high density scenario described below, reach, at most, $\tau \approx 1$. Therefore, \CEIO\ line intensity ratios are not strongly affected by the possibly different locations of cloud `photospheres' from which the bulk of the photons are emitted. All transitions of \CEIO\ should trace all molecular gas at densities $\geq 10^3$\,\percc, as long as \CEIO\ is not selectively dissociated by UV radiation. Away from the star forming regions Sgr\,A and Sgr\,B2, the UV radiation field in the Galactic center region is not likely to be strong (e.g.\ Nagakawa et al.\ 1995). In addition, because we are considering density peaks of the molecular gas (and dust), the destruction of C$^{18}$O should not be a problem; from the H$_2$ column densities derived below, dust extinction is likely to provide sufficient shielding. \begin{figure} \psfig{file=lvg1.ps,bbllx=0pt,bblly=40pt,bburx=290pt,bbury=500pt,height=10.5cm} \caption{ \label{lvg1} Line intensity ratios from our LVG model calculations for \CEIO\ (upper panel) and $^{28}$SiO and $^{29}$SiO (lower panel). The contours are at: \protect\\ \RZWCO (solid): 0.3, 0.45, 0.6, 1.0, 1.5, 2.0, 3.0, 3.5, 3.75 \protect\\ \RDRCO (dashed): 0.1, 0.2, 0.35, 0.5, 0.75, 1.0, 1.5, 2.0 \protect\\ \RMSIO (solid): 0.04, 0.05, 0.1, 0.2, 0.5, 0.8, 0.9, 1.0 \protect\\ \RISOSIO (dashed): 1.2, 1.5, 2.0, 5.0, 10.0, 15.0, 19.0, 19.75, 20.0} \end{figure} In our model calculations, the population $n(J)$, excitation temperature \TEX , brightness temperature \TB , and optical depth $\tau$, are determined for each rotational level up to a maximum $J= 20$ for \CEIO\ and $J= 13$ for SiO. Collision rates are taken from Green \& Chapman (1978). We have chosen a constant velocity gradient of 5\,\kms\,pc$^{-1}$, a value estimated from the velocity extent and diameter of Galactic center clouds as mapped by Bally et al.\ (1987) or Lindquist et al.\ (1995). The background continuum temperature was assumed to be 2.7\,K. \subsubsection{C$^{18}$O} We assume that the abundance of \CEIO\ relative to H$_2$ is $4 \cdot 10^{-7}$. This ratio is based on the standard Galactic center $^{16}$O/$^{18}$O isotopic ratio of 250 (Wilson \& Matteucci 1992, Wilson \& Rood 1994) and a CO/H$_2$ ratio of 10$^{-4}$ (e.g.\ Frerking et al.\ 1982, Duvert et al.\ 1986 (dark clouds), Blake et al.\ 1987 (OMC-1)). According to the calculations of Farquhar et al.\ (1994), this ratio is stable against a possibly enhanced cosmic ray flux close to the Galactic center. For \CEIO , we then calculate a grid by varying the H$_2$ density, \numd , from $10^{2}$\,\percc\ to $10^{6}$\,\percc\ and the kinetic temperature, \TKIN , from 5 to 200\,K. We found the range of \numd\ and \TKIN\ by comparing the observed line intensity ratios (Table\,\ref{rat}) to the ratios predicted by the model (Fig.\,\ref{lvg1}). It is important to use line intensity ratios, not just intensities, since intensities are affected by an a priori unknown beam filling factor, $f_{\rm b}$. For the intensity ratios, we took $f_{\rm b}$ to be identical for all transitions of a given molecule. The integrated intensity ratio of the \CEIO ($2\to 1$) to ($1\to 0$) transition, \RZWCO , from all sources is $1.3 \pm 0.5$, with extreme \ values ranging from 0.5 to $\sim 2$. The ratio of the ($3\to 2$) to the ($2\to 1$) intensity, \RDRCO , is $0.5 \pm 0.2$. Averaged over all sources, the LVG model fits yield \TEX\,$= 11( \pm 4)$\,K for the \CEIO ($2\to 1$) transition, that is, \CEIO\ is subthermally excited. The beam averaged peak optical depth $\tau_{21}$ is always $< 0.2$. The measured peak temperatures \TMB\ indicate a beam filling factor $f_{\rm b}^{\rm C^{18}O}$ of $\sim 1$ for \numd\ $(2 - 4)\,10^3$\,\percc . $f_{\rm b}^{\rm C^{18}O}$ is lower if \numd\ is higher. From the structure seen in our maps of the \CEIO\ 1$\to$0 beam in higher transitions, we estimate $f_{\rm b}^{\rm C^{18}O}$, averaged over all gas components, to be $\sim 0.5$. These maps, however, have a low signal-to-noise ratio. From the model calculations (Fig.\,\ref{lvg1}, upper panel), maximum values up to 4 in \RZWCO\ and up to 2 in \RDRCO\ are expected for hot (\TKIN $> 100$\,K) and dense (\numd\,$\sim 10^{4} - 10^{5}$\,\percc) gas. Such high ratios are reached in none of our sources. Possible combinations of \numd\ and \TKIN are {\em either} a high \TKIN\ of ($\geq 100$\,K ) and \numd\ of (1 -- 4)\,10$^3$\,\percc\ {\em or} a lower \TKIN\ ($\leq 20$\,K -- 30\,K) at a higher \numd\,$\geq 10^4$\,\percc. We have carried out this analysis individually for all clouds. For all sources our \CEIO\ data {\em unambiguously} require gas with densities of $\geq 10^4$\,\percc\ in these Galactic center clouds to be {\em cool}. From ammonia measurements (H\"uttemeister et al.\ 1993b) it is known that both hot and cool gas is present within the area covered by our beam. Thus, it is likely that both scenarios possible from the \CEIO\ data are realized. A continuum of temperatures and densities is compatible with our data. Then \TKIN\ rises steadily as \numd\ decreases. Therefore, density peaks (indicated by strong CS lines) in Galactic center region GMCs not presently undergoing massive star formation can be regarded as {\em `cool dense cores'} in contrast to the hot dense cores generally encountered within GMCs in the disk. \subsubsection{SiO} It has been shown that the fractional abundance of SiO, $X$(SiO), can change by more than six orders of magnitude between quiescent cold material and hot, shocked dense gas (Ziurys et al.\ 1989, Mart\'{\i}n-Pintado et al.\ 1992). Since SiO is readily detected in all our sources, $X$(SiO) {\em cannot} be low. However, we cannot assume any specific value. Therefore, we have used $X$(SiO) as a free parameter in the models, varying it from $10^{-12}$ to $10^{-6}$. Grids have been calculated as a function of \numd\ and $X$(SiO) for \TKIN $=$ 25\,K, 50\,K, 75\,K, 100\,K, 150\,K. The average line intensity ratio between the SiO\,$(5\to 4$) and ($2\to 1$) transition is \RMSIO\,$=\,0.16 \pm 0.05$, for those sources where we detected SiO($5\to 4$). M+1.31--0.13 is exceptional and will be discussed further in Section\,4.3. Since the signal-to-noise (S/N) ratio in the ($5\to 4$) line is not always high and the ($2\to 1$) SEST beam was not mapped in this transition, it is not clear whether source to source differences (apart from M+1.31--013) are significant. For the line ratios between the ($2\to 1$) transitions of $^{28}$SiO and $^{29}$SiO (excluding M+1.31--0.13), we find a mean ratio of \RISOSIO$\,=\,12.2\pm 2.2$ with extreme values ranging from 7.4 to 15.9. Since these data, even in the $^{29}$SiO transition, have very good S/N ratios and were measured with the same telescope and beamsize, there is no doubt that the source-to-source differences are significant. Fig.\,\ref{lvg1} (lower panel) shows the line ratios \RMSIO\ and \RISOSIO\ for \TKIN = 75\,K from our LVG calculations. Lowering \TKIN\ to 25\,K requires a slight increase of the H$_2$ densities for a given line ratio, by a factor of $\sim 1.25$. Raising \TKIN\ to 150\,K corresponds to slightly lower densities. In general, the effects of changing \TKIN\ on the line ratios are negligible: SiO line ratios are almost insensitive to cloud temperature. To derive \TEX\ and \numd , we have analysed all sources individually. Excluding M+1.31--0.13, we find \TEX\ ranging from 5\,K to 10\,K. \numd\ changes with the assumed SiO abundance: A higher $X$(SiO) requires a lower H$_2$ density (Fig.\,\ref{lvg1}). Curves of constant \RMSIO\ and \RISOSIO\ run parallel for a range of densities and SiO abundances. Typical ranges are $X$(SiO)\,$\approx 10^{-6}$ for \numd\,$\approx (1 - 2.5)\,10^3$\,\percc\ to $X$(SiO)\,$\approx 10^{-9}$ for \numd\,$\approx 4\,10^4 - 10^5$\,\percc . Lower SiO abundances or higher densities are not possible, since then the curves diverge. Plotting the measured \TMB 's for the $^{28}$SiO($2\to 1$) transition in the LVG plots (e.g.\ Fig.\,\ref{lvg}, lower left panel), we find that a considerably higher \TB\ is needed to fit the observed line ratios. Thus we estimate that the beam filling factor $f_{\rm b}^{\rm SiO}$ (\TMB\,=\,$f_{\rm b}^{\rm SiO}$\,\TB ) is $\sim 0.2$ for the `typical' sources. This analysis is biased toward more intense SiO lines, since only those were observed in transitions other than $^{28}$SiO$(2\to 1)$. The strength in SiO toward these positions may be either due to a large amount of molecular gas, or to an abundance of SiO that is above average. This question will be addressed in the next section. \subsection{A joint view of C$^{18}$O and SiO} \subsubsection{Column densities and SiO abundances} Further insight in the structure of the clouds is gained by combining what can be learned from \CEIO\ and SiO. \begin{table} \caption{ \label{col} Total beam averaged column densities of \CEIO\ ($\bar N^{\rm C^{18}O}$) and SiO ($\bar N^{\rm SiO}$) and beam averaged abundances of SiO, $X$(SiO). ($\bar N^{\rm H_2}$ follows from ($\bar N^{\rm C^{18}O}$, assuming $X$(\CEIO ) to be $4 \cdot 10^{-7}$. Errors are given in parentheses.} \begin{tabular}{l|rr|rrr} Source & $v_{\rm LSR}^{\rm C^{18}O}$ & $\bar N^{\rm C^{18}O}$ & $v_{\rm LSR}^{\rm SiO}$ & $\bar N^{\rm SiO}$ & $X^{\rm SiO}$ \\ & & 10$^{15}$ & & 10$^{13}$ & 10$^{-9}$ \\ & km/s & cm$^{-2}$ & km/s & cm$^{-2}$ & \\ \hline Clump\,1 & 97 & 2.6(0.9) & 97 & $\sim 0.1$ & 0.1\\ M--0.96+0.13 & 135 & 2.9(1.1) & 135 & $\sim 0.9$ & 1.7 \\ M--0.50--0.03 & --100 & 8.9(1.4) & --98 & $\sim 0.8$ & 0.4 \\ & --2 & 1.1(0.2) & --2 & $< 0.02$ & $<0.07$\\ M--0.42--0.01 & --84 & 7.5(0.8) & --80 & $\sim 0.6$ & 0.3 \\ & --4 & 1.2(0.1) & --4 & $< 0.1$ & $<0.03$ \\ M--0.39+0.02 & --80 & 6.6(0.6) & --79 & $\sim 0.6$ & 0.4 \\ M--0.33--0.01 & --47 & 7.7(2.5) & --51 & $\sim 0.8$ & 0.4 \\ M--0.45--0.09 & --61 & 3.6(0.9) & --61 & $< 0.09$ & $<0.1$ \\ & --10 & 5.5(0.5) & --12 & 3.2(0.4) & 2.3 \\ M--0.30+0.00 & --33 & 7.2(1.5) & --38 & $\sim 1.0$ & 0.6 \\ M--0.44--0.10 & --15 & 7.4(2.9) & --20 & 4.4(1.0) & 2.4 \\ M--0.51--0.16 & --61 & 13.1(3.5) & --57 & 4.0(1.5) & 1.2 \\ & --2 & 1.0(0.3) & --2 & $< 0.03$ & $<0.1$ \\ M--0.50--0.18 & --57 & 7.6(1.5) & --55 & $\sim 1.8$ & 1.0 \\ M--0.48--0.19 & --50 & 3.6(0.8) & --53 & $\sim 0.7$ & 0.7 \\ M--0.38--0.13 & --33 & 6.6(0.9) & --32 & $\sim 0.3$ & 0.2 \\ M--0.45--0.19 & 21 & 3.0(0.5) & 21 & $< 0.04$ & $<0.05$ \\ & 58 & 4.4(1.0) & 58 & $\sim 1.1$ & 1.0 \\ M--0.43--0.21 & --34 & 10.4(0.4) & --27 & $\sim 1.2$ & 0.5 \\ & 19 & 2.2(0.3) & 19 & $< 0.04$ & $<0.07$ \\ & 53 & 2.8(0.6) & 54 & $\sim 1.1$ & 0.6 \\ M--0.15--0.07 & 5 & 19.1(1.0) & 4 & 6.2(1.0) & 1.3 \\ M--0.32--0.19 & --32 & 4.3(0.7) & & & \\ & 25 & $\leq 1.6$ & 25 & $\sim 0.6$ & $>1.5$ \\ M+0.24+0.02 & 35 & 26.1(5.0) & 36 & 5.6(1.2) & 0.9 \\ M+0.10--0.12 & 22 & 11.5(1.9) & 29 & $ \sim 2.2 $ & 0.8 \\ M+0.43--0.05 & --2 & 3.4(0.8) & 12 & 2.4(0.6) & 2.8 \\ & 38 & 4.9(0.9) & & & \\ M+0.41+0.03 & 30 & 14.1(1.0) & 21 & $\sim 2.0$ & 0.6 \\ M+0.50+0.00 & 30 & 13.2(2.9) & 26 & 5.0(1.6) & 1.5 \\ M+0.68--0.12 & 30 & 17.2(2.0) & 20 & 7.6(0.7) & 1.8 \\ M+0.83--0.10 & 24 & 12.1(3.5) & 24 & $\sim 1.9 $ & 0.6 \\ & 99 & 4.4(0.4) & 96 & $\sim 0.5$ & 0.5 \\ M+1.46+0.27 & 80 & 4.9(1.3) & 81 & 4.9(1.0) & 4.0 \\ M+0.83--0.18 & 41 & 6.8(1.6) & 43 & 4.9(0.2) & 2.9 \\ M+1.15--0.09 & --18 & 1.2(0.2) & --18 & $<0.06$ & $<0.02$ \\ & 60 & 8.9(1.9) & 67 & $\sim 4.0$ & 1.8 \\ M+1.13--0.12 & --15 & 7.4(0.8) & --15 & $\sim 0.5$ & 0.3 \\ & 80 & 3.3(1.3) & 70 & $\sim 3.3$ & 4.0 \\ M+1.31--0.13 & 30 & 6.9(3.5) & 35 & $\sim 10 $ & 5.8 \\ & 80 & 4.0(0.4) & 87 & $\sim 10 $ & 10.0 \\ M+0.94--0.36 & --49 & 8.8(0.3) & --41 & $\sim 0.6$ & 0.3 \\ M+3.06+0.40 & 12 & 1.5(0.3) & 12 & $< 0.02$ & $<0.05$ \\ & 22 & 3.7(1.0) & 30 & $\sim 1.7$ & 1.8 \\ M+3.06+0.34 & 27 & 10.0(3.0) & 27 & 2.9(0.3) & 1.2 \\ M+2.99-0.06 & --13 & 3.3(0.7) & --12 & $\sim 1.6$ & 1.9 \\ \hline \\ \end{tabular} \\ \end{table} We calculate the total (beam averaged) column densities for \CEIO\ and $^{28}$SiO, applying the same procedure to both molecules: First, we use the observed \TMB\ and the best fit \TEX\ from the LVG model for the ($2\to 1$) transitions (the lines with the best S/N ratio) to calculate the (beam averaged) optical depth $\bar \tau_0$. Taking $ \int\,\bar \tau(v)\,{\rm d}v\,= 1.06\,\bar \tau_0\, \Delta v_{1/2}$, we then derive the column density in the $J=1$ level, $\bar N_1$. The total beam averaged column density $\bar N$ is determined by dividing $\bar N_1$ by the fraction of the population residing in the $J=1$\ level, derived from the LVG model. (See Rohlfs \& Wilson (1996) for a collection of the (standard) formulae we used.) For non-gaussian lines, the second moment of the line has been used as a measure of the line width instead of $\Delta v_{1/2}$. This has been combined with a \TMB\ which reproduces the observed integrated intensity. This is correct as long as \TEX\ does not change across the line profile. Since the line shapes in all transitions observed for a given molecule agree closely, this is a reasonable assumption. For the sources where we observed only the SiO($2\to 1$) line, we have used average values for \TEX\ and the fraction of the population in the $J = 1$ level. These are 7\,K and 0.36, respectively. Total beam averaged column densities are given in Table\,\ref{col}. The usually optically thin, easily excited transitions of \CEIO\ are known to be excellent tracers of H$_2$ column density, $\bar N_{\rm H_2}$, over a wide variety of H$_2$ densities and kinetic temperatures. In particular, the integrated intensity of the $J = 2 \to 1$ transition can be directly related to $\bar N_{\rm H_2}$ (Genzel 1992). Our LVG modelling confirms that for a large range of temperatures the fraction of the total population in the \CEIO\ $J = 1$ level remains almost constant between \numd\,$=10^2$\,\percc\ and $10^4$\,\percc. In most cases, we find that column densities agree with the formula given Mauersberger et al.\ (1992) and Genzel (1992) ($\bar N_{\rm H_2} = 1.12\ 10^{21} \int T_{\rm MB}^{\rm C^{18}O} (2\to 1)\,{\rm d}v $ for the Galactic center isotopic ratios) to within 10\%. Differences up to a factor of 2 are found if the excitation temperature is exceptionally low. The knowledge of $\bar N_{\rm H_2}$ and $\bar N^{\rm SiO}$ allows us to derive a beam averaged SiO relative abundance, assuming that the \CEIO\ and SiO emission arises from the same gas. This is given in Table\,\ref{col}, and typically ranges from 0.5\,10$^{-9}$ to 5\,10$^{-9}$. Since \CEIO\ presumably traces {\em all} H$_2$ with \numd\,$\geq 10^3$\,\percc, the true abundance of SiO cannot be {\em lower} than the value thus determined. It can be higher, if only part of the gas within the beam is abundant in SiO, which, considering the differences in line centers and line profiles between \CEIO\ and SiO (e.g.\ M+0.24+0.02, see Fig.\,\ref{spectra}), is likely. \begin{figure*} \psfig{file=lvg.ps,bbllx=25pt,bblly=40pt,bburx=560pt,bbury=680pt,% height=12.2cm,angle=-90} \caption{ \label{lvg} LVG Model calculations for two individual sources, demonstrating the range of physical conditions encountered. The observed line intensity ratios are plotted including their 1\,$\sigma$ uncertainties. The \TB - curves are the observed \TMB\ and the best fit temperature derived from the line ratios. The coresponding beam filling factor $f_{\rm b}$ is also given. For M+1.31--0.13, the observed and the best fit temperature are the same, thus $f_{\rm b} = 1$. $\tau^{21}$ is the peak optical depth of the $^{28}$SiO line (not beam averaged). } \end{figure*} For those sources where the SiO($2\to 1$) transition is weak and which were only observed in this SiO transition, the H$_2$ column density obtained from \CEIO\ is not systematically lower than for the more intense sources. On average, in these clouds $X$(SiO) is lower by a factor of $\sim$\,3. Source-to-source variations, are significant, even though the properties in dense gas as traced by CS and total gas as traced by \CEIO\ do not differ. This agrees with the variation found in the survey of Mart\'{\i}n-Pintado et al.\ (1997). Note, however, that we find SiO with an abundance that is high when compared to typical disk GMCs in all sources. \subsubsection{Density and temperature} We now combine the LVG models for \CEIO\ and SiO to derive a consistent solution for \numd\ and \TKIN . From the NH$_3$ studies of H\"uttemeister et al.\ (1993b), we know that, for a typical cloud, roughly 25\% of the neutral gas mass is contained in the hot ($> 100$\,K) component, which, from our \CEIO\ analysis, must have a low density. Thus, the remaining $\sim 75$\% of the total gas mass is at a density of $\geq 10^4$\,\percc. Let us assume that the SiO emission arises in only this component. The corresponding relative abundances of SiO, $X$(SiO), are 25\% above the beam averaged values given in Table\,\ref{col}. $X$(SiO) lies within the range allowed by the SiO line ratios (Fig.\,\ref{lvg1}) and is consistent with an H$_2$ density of typically (1~--~4)\,10$^4$\,\percc. The beam filling factor $f_{\rm b}$ in this dense component must be the same for both the SiO and the part of the \CEIO\ emission that arises in this component. We can estimate the main beam brightness temperature of the dense, \CEIO\ emitting gas to be 0.75 of the total \CEIO\ intensity. Using the beam filling factor determined for the dense gas from SiO lines, the brightness temperature of the \CEIO $(2 \to 1)$ line is $T_{\rm B}^{\rm C^{18}O}$(75\%)\,$\approx$\,0.75\,\TMB\,$/f_{\rm b}$. This simple relation can be used since the \CEIO ($2\to 1$) line has an optical depth $\leq 1$. If the solution is self-consistent, the H$_2$ density indicated by this value of \TB\ and the observed line ratios must agree with the density derived for SiO. To within a factor of 2, we find that this is indeed the case. Thus, it is possible that the SiO emission originates in gas having \numd\,$\sim 10^4$\,\percc $- 4\,10^4$\,\percc. This gas must be relatively cool, which is a surprising result in the light of the clear association of the SiO emission with hot shocked gas in the Galactic disk. The intrinsic optical depth $\tau_0$ of the SiO($2\to 1$) then ranges from 4 - 8. Can the SiO emission be associated with the hot, thin gas component? To check this, we used the same procedure as before: From the assumption that the SiO emission arises in $\sim 25$\% of the gas, we get a value for $X$(SiO) and a corresponding H$_2$ density. The beam filling factor $f_{\rm b}$ for SiO and, consequently, for the part of \CEIO\ originating in the same gas does not change. Now, $T_{\rm B}^{\rm C^{18}O}$(25\%)\,$\approx 0.25\,$\TMB\,$f_{\rm b}$. The additional requirement that this gas should be hot yields a H$_2$ density that is, for many clouds, lower by an order of magnitude than what is derived for the same component from SiO. Thus, this is not a consistent solution. If the SiO arises in a still smaller portion of the total gas, the discrepancy becomes even worse. Note that we do not claim that the hot, thin component is devoid of SiO. We consider it likely that the SiO abundance in this component is similar to what we determine for the dense component. However, the SiO excitation in this component is extremely subthermal, leading to (1) lower line intensities and (2) different line intensity ratio than those we observe (specifically, no emission from the (5$\to$4) transition). Finally, we can rule out the presence of a hot, high density (\numd\,$\geq 10^4$\,\percc) component as the origin of the SiO emission: We know from the line ratios \RMSIO\ and \RISOSIO\ that $X$(SiO) must decrease with increasing density, independent of \TKIN . Thus, a hot, high density component containing the bulk of the SiO would have to contain a significant fraction of the total H$_2$ column density and mass. Such a component should be visible in \CEIO . Thus we conclude that it does not exist in these clouds. Therefore the only consistent scenario that accounts for both the SiO and \CEIO\ results requires that most of the SiO emission arises from a {\em cool} (20\,K -- 30\,K) gas component with \numd\ $\sim 10^4$\,\percc . So far, we have {\em assumed that the \CEIO\ abundance in all components considered is known and constant.} If this is not the case, i.e.\ if \CEIO\ is {\em not} a reliable H$_2$ column density tracer, the line of argumentation given above cannot be maintained. In this case, it becomes possible to claim that the SiO emission arises in a hot, dense component not seen in \CEIO . This is, however, only possible if \CEIO\ is selectively underabundant in the component where the SiO abundance is high. Since selective \CEIO\ dissociation due to lack of shielding is efficient only in diffuse, low density gas (if the UV field is strong enough), it is difficult to find a scenario causing \CEIO\ to be underabundant in dense, warm gas. \subsection{A detailed investigation of two sources} \paragraph{M+0.83--0.18} In Fig.\,\ref{lvg} (left panels), we illustrate the procedure of deriving the physical cloud parameters described above for the source M+0.83--0.18. This is representative of the type of cloud where SiO emission arises in a cool, moderately dense component. The SiO line ratios in this source (lower left panel in Fig.\,\ref{lvg}) are well fit by \TEX\,$= 10$\,K. Then, $X$(SiO) is 3.9\,10$^{-9}$ if SiO arises from 75\% of the total gas. The observed \TMB\ is 0.48\,K, while a \TB\ of 7.5\,K is required to fit the range of \numd\ and $X$(SiO). This results in the lowest value for $f_{\rm b}$ of all sources, 0.06, and the highest optical depth $\tau_0$, 8. Applying the same $f_{\rm b}$ to the \CEIO\ data (upper left panel), we have \TB (\CEIO)\,$=$\,0.75\,\TMB /0.06\,=\,5.1\,K. If we take the resulting \numd\ from the SiO analysis, 4\,10$^4$\,\percc , and transfer this to \CEIO , we find that \TKIN\ is only $\sim 12\,$K. If we require \TKIN\ to be identical to the value we have measured for the cool component from NH$_3$, 21\,K (H\"uttemeister et al.\ 1993b), the density obtained for \CEIO\ is 2\,10$^4$\,\percc . Considering all the assumptions, there is satisfactory agreement. Could the results be consistent with the `high temperature scenario' described in the last section? We find that \numd\ required from an analysis of the SiO data is $\sim 3\,10^4$\,\percc, while the \CEIO\ data at the appropriate \TB\ yield $\sim 3\,10^3$\,\percc , clearly a far less consistent solution. \paragraph{M+1.31--0.13} Now we analyze the one source which is very different from all the others in our sample: M+1.31--0.13 (right panels of Fig.\,\ref{lvg}). This is the only source where the ($5\to 4$) transition of $^{28}$SiO was not detected, even though the ($2\to 1$) line is strong. It also is the one source where \RISOSIO\ approaches the terrestrial isotopic ratio of 20, indicating that $\tau_0$ for $^{28}$SiO($2\to 1$) cannot be large. M+1.31--0.13 has two velocity components. We analyze the (slightly) more intense \vlsr\,$\sim 85$\,\kms\ component, but the \vlsr\,$\sim 35$\,\kms\ component shows the same characteristics. From the SiO line ratios, we derive an extremely low excitation temperature, 3.2\,K. The large column density for SiO depends strongly on the exact value of \TEX . Therefore, we can give only an estimate of $\bar N^{\rm SiO}$ in Table\,\ref{col}. The corresponding $X$(SiO) is $\sim 10^{-8}$, by far the highest in the sample. In addition, for this source we find \TMB\,$=$\,\TB , i.e.\ $f_{\rm b} = 1$. As expected, this results in a lower $\tau_0$, $\sim 2$, in the $^{28}$SiO($2\to 1$) line. For the H$_2$ density, we find \numd\,$\approx 2\,10^3$\,\percc . This is a factor of $\sim 10$ lower than for all other sources. Since our calculations give $\tau_0 \sim 2$ and $f_{\rm b} = 1$ for SiO, we conclude that in this source SiO traces {\em all} the gas, as does \CEIO . Hence, $f_{\rm b} = 1$ and \TB\,$=$\TMB\ is also valid for \CEIO . The resulting \numd\ from \CEIO\ is $2\,10^3$\,\percc , identical to the result obtained from SiO. As expected, \TKIN\ is high, $>\,100$\,K. Thus, for this one source, we have found a fully self-consistent solution of an entirely different type: The SiO arises in thin, hot gas and has an abundance that is far above average, even for the clouds that are rich in SiO. Is there evidence that M+1.31--0.13 lacks the dense and cool gas component? Remarkably, our NH$_3$ data show that this cloud is {\em the only one in the entire sample} in which the $(J,K)=(4,4)$ inversion transition at 200\,K above ground is more intense than the (1,1) transition at 23\,K above ground. Thus, in this cloud a cool, dense gas component containing the bulk of the gas is not present. Most, perhaps all of the gas is hot and thin. Relating this to the high SiO abundance, it is likely that SiO does indeed form at high temperatures in the Galactic center region. In M+1.31--0.13, the SiO formation process is either still ongoing or has occured very recently. The SiO has not yet had time to recondense onto dust grains, and the cloud itself has not had time to form dense cool cores. \section{Discussion} \subsection{Cloud properties as a tracer of large scale dynamics} We will now explore how the relative SiO abundances may be related to the position of the cloud in the large scale Galactic center environment. In Fig.\,\ref{nl} we plot the the beam averaged SiO abundance, $X^{\rm SiO}$, on a longitude-velocity diagram of the large scale distribution of \CEIO . While any estimates of line-of-sight locations in the Galactic center are rather uncertain, this plot gives a measure of location. \begin{figure} % \psfig{file=nl_xfig.ps,height=6.5cm,angle=-90} \psfig{file=nl_xfig.ps,bbllx=0pt,bblly=55pt,bburx=530pt,bbury=700pt,height=6.7cm,angle=-90} \caption{\label{nl} The beam averaged SiO abundance plotted on a longitude- velocity diagram of the large scale distribution of \CEIO , adapted from Paper\,I (spatial resolution 9$'$). } \end{figure} There is a clear trend in the large-scale distribution of SiO abundances. For $-0\ffd 35 \leq l \leq 0\ffd 6$, between the location of the Sgr\,C and Sgr\,B2 regions, we find a pronounced lack of very large $X$(SiO). Most clouds with very high abundances, including the exceptional source M+1.31--0.13, are located at $l > 0\ffd 8$. The `Clump\,2' Region is not included in this plot, since, at $l \approx 3^{\circ}$, this is not really part of the continuous molecular bulge or, in the notation of Morris \& Serabyn (1996), the Central Molecular Zone (CMZ), although it shares many characteristics of the Galactic center gas. This detection of an `SiO hole' at $-0\ffd 35 \leq l \leq 0\ffd 6$ is confirmed by an extension of the survey of Mart\'{\i}n-Pintado et al.\ (1997) (unpublished data). In recent years, the large scale dynamics of the gas in the CMZ, characterized by large non-circular motions and a distinct `parallelogram' shape of $lv$-diagrams based on $^{12}$CO and $^{13}$CO (e.g.\ Bally et al.\ 1988), have been explained in terms of a model involving a rotating (stellar) bar with corotation at 2.4\,kpc and oriented at an angle of $\sim 20^{\circ}$ with the line of sight to the Galactic center (see Morris \& Serabyn 1996 and references therein). Within such a potential, there exist two categories of closed elliptical orbits, called $x_1$ and $x_2$-type orbits. Inside a cusped orbit, the $x_1$-orbits, which are elongated along the bar axis, become self-intersecting. Clouds on these orbits encounter a shock and within a dynamical time plunge into $x_2$-orbits that lie considerably deeper within the potential and mimic circular orbits (Mulder \& Liem 1986, Athanassoula 1988, 1992). This encounter breaks the flow along the cusped orbit into a spray that spreads out into the interior of the orbit (Binney et al.\ 1991, Athanassoula 1992, Jenkins \& Binney 1994). According to this model, Sgr\,B2 and Sgr\,C can be interpreted to be at the locations of the intersections of the $x_1$- and the $x_2$-orbits. The dense, virialized clouds on the $x_2$-orbits, inside of Sgr\,B2 and Sgr\,C, are more likely to form stars, while the less dense clouds on the $x_1$ orbits are too disrupted by shocks for efficient star formation. Generally, this is supported by the fact that stars of young and intermediate age are restricted to the area inside of Sgr\,B2 and Sgr\,C, suggesting sustained star formation in this region (Serabyn \& Morris 1996). In Sgr\,B2, the ongoing star formation, the large amount of molecular gas, the existence of a very hot molecular component (H\"uttemeister et al.\ 1993a, 1995, Flower et al.\ 1995) and the high optical depth of the $^{12}$CO emission (Paper\,II) also provide ample evidence for this picture. The Sgr\,C complex is a much more quiescent region, but it contains a large H\,{\sc II} region and GMCs. Because of the the sprayed gas crashing into material that is still on $x_1$-orbits, collision regions are expected along the acceleration part of the cusped orbits. Due to perspective, these main collision areas should be at higher positive longitudes than Sgr\,B2 and between Sgr\,A and Sgr\,C. While the situation at negative $l$ is confused (the gas on $x_2$ orbits and the collision regions are expected at similar Galactic longitudes), at positive $l$, especially the `$1\ffd 5$--complex' (Bally et al.\ 1988) shows all characteristics expected from a collision region (see the discussion in Paper\,II). Inspecting Fig.\,\ref{nl} with this picture in mind, we find remarkable support from our data. The exceptional cloud M+1.31--0.13 and three other sources with high SiO abundances are located within the 1.5$^{\circ}$--complex, exactly where strong shocks are expected. Generally, the high SiO abundances at $l > 0\ffd 8$ can be identified with a region in which the probability of gas encountering shocks is high. Thus, for the first time, we can tie the physical and chemical parameters of Galactic center molecular peaks to the large scale dynamics of the region. Since CS traces all dense gas, not just the part that has been subjected to shocks, we do not expect a similar `zone of avoidance'. Indeed, a plot similar to Fig\,\ref{nl} comparing the CS abundance in these clouds (data presented in H\"uttemeister 1993), to the total H$_2$ column density does not show such an effect. Of course, the correspondence is not absolute. Not every cloud at $l > 0\ffd 8$ {\em must} experience strong shocks. Note, however, that the clouds with very low SiO abundances in this range of $l$ have negative velocities and might not be located in the collision region. Also, there are certainly strong shocks occuring in the Galactic center that are unrelated to large scale dynamics. Mart\'{\i}n-Pintado et al.\ (1997) explain the `SiO clouds' they find by shocks of a variety of origins: Interaction with SNRs close to Sgr\,A, interaction with non-thermal filaments in the radio arc and cloud-cloud collisions (as expected from large scale dynamics) or expanding bubbles in the vicinity of Sgr\,B2. Their map extends from $-0\ffd 2$ to $+0\ffd 8$. Therefore, they have missed the systematic signature of large scale effects present at higher positive longitudes. Note that this `bar model' is not the only possibility to explain the large scale dynamics of the Galactic center region. von Linden et al.\ (1993) suggest that an accretion disk can also reproduce the basic structure seen in position-velocity plots. While this is certainly true for the inner region ($x_2$-orbits are almost indistinguishable from circular orbits) it is not clear whether such a scenario can also explain the gas distribution and the different chemical properties in the entire CMZ. Based on the data presented in Papers I and II and this work, we will examine this question closely in a forthcoming paper (von Linden et al.\ 1998). \subsection{The origin of SiO} Even though the SiO emission we observe is mostly associated with cool gas, we consider it likely that it forms under conditions of high kinetic temperature and grain erosion by shocks (Mart\'{\i}n-Pintado et al.\ 1992). Both the distribution discussed in the previous section and the existence of the hot cloud core M+1.31--0.13 can be taken as evidence for this. From the fact that only one cloud in our sample, namely M+1.31--0.13, shows hot gas emitting in SiO lines, we can conclude that shocked gas forms dense, cooler cores fairly fast. Extrapolating from the analysis in Hollenbach (1988), we find a very short cooling time scale for gas at a density of 10$^4$\,\percc . A competing process is the condensation of gas phase SiO onto grain mantles. We estimate this `freeze out' time scale to be of order $10^6$\,yr or slightly less (Rohlfs \& Wilson 1996). Thus, gas with a high gas phase relative abundance of SiO has been cool for most of its lifetime. This is supported by the fact that we find only one example of a cloud where the gas phase SiO is arising from gas that is still at a high temperature. \begin{figure} \psfig{file=complex_spe.ps,bbllx=170pt,bblly=55pt,bburx=520pt,bbury=600pt,height=4.2cm,angle=-90} \caption{\label{complex} The chemical complexity of sources strong in SiO is demonstrated by our spectrum of M+0.83--0.18, covering $\sim 1.8$\,GHz and including both $^{29}$SiO and $^{30}$SiO. The spectrum shown is a composite of two single SEST low resolution spectra.} \end{figure} After more than 10$^6$ years, the gas slowly loses the chemical memory of having been shocked. From the typical radial velocities of the clouds, we estimate that, during this time, these can move $\sim 50$\,pc. If the very SiO-rich clouds are preferentially located on $x_1$-orbits along the bar, these will remain sufficiently close in longitude to the region where they encountered the shock to be recognised as a distinct population. Since we find some SiO in all cloud cores, additional formation of gas phase SiO is required, probably by local turbulence and/or cloud-cloud collisions causing some shocks. This is a general characteristic of the cores of Galactic center GMCs. Ongoing local shock activity is likely to be also necessary to maintain the hot gas component seen in all clouds (Flower et al.\ 1995). It is known that gas in post-shock regions, away from chemical equilibrium, is characterized by abundant complex molecules (e.g.\ Brown et al.\ 1988). This is exactly what can be observed toward the positions within the CMZ that are strong in SiO. We show just one example in Fig.\,\ref{complex}. In the Galactic disk, such complex spectra are typical for very confined regions, while in the CMZ they seem to be ubiquitous. This chemical complexity can only be maintained for $\sim 10^6$ years. On larger timescales, an equilibrium state is reached. This consists of a gas phase component of mostly diatomic molecules (Herbst \& Leung 1989). This is another indication of the high frequency of shocks in the CMZ. An analogy to the CMZ may be found in the prototypical starburst galaxy NGC\,253, where Mauersberger \& Henkel (1991) detected extended SiO emission, and where complex molecules such as CH$_3$OH, CH$_3$CCH and CH$_3$CN are observed (Henkel et al.\ 1987, Mauersberger et al.\ 1991, H\"uttemeister et al.\ 1997). \section{Conclusions} Based on our measurements and analysis of 33 molecular density peaks in the Galactic center in the $J = 1\to 0, 2\to 1$ and $3\to 2$ transition of \CEIO , \ the $J = 2\to 1$ and $5\to 4$ transition of $^{28}$SiO and the $J = 2\to 1$ transition of $^{29}$SiO and $^{30}$SiO, we find: \begin{enumerate} \item All sources are easily detected in all transitions of \CEIO\ searched for and in $^{28}$SiO($2\to 1$), demonstrating that the properties of molecular peaks in the Galactic center region are markedly different from the Galactic disk, where thermal SiO emission is confined to very small regions in the vicinity of outflows associated with star formation. Local gas on the line-of-sight, distinguished from Galactic center gas by its narrow lines, is not detected in SiO. \item The rare isotopomers $^{29}$SiO and $^{30}$SiO are detected in all 12 studied sources. The $J=5\to 4$ transition of the main isotopomers is seen in 7 out of 8 sources. The line intensity ratio of $^{29}$SiO and $^{30}$SiO shows that the terrestrial isotope ratio, $^{29}$Si/$^{30}$Si\,$=$\,1.5, holds for the Galactic center region. \item From LVG model calculations applied to the \CEIO\ line intensity ratios \RZWCO\ and \RDRCO , cool (\TKIN\,$\sim 20 -30$\,K) gas toward the molecular peaks has moderately high densities (\numd\,$\sim (1 - 4)\,10^4$\,\percc), while high kinetic temperatures of $> 100$\,K correspond to H$_2$ densities that are an order of magnitude lower. This is contrary to what is found in the disk, were the cores of GMCs are usually hot, and is an indication that the `cool cores' in `typical' Galactic center GMCs, away from Sgr\,A and Sgr\,B2, are not presently forming high mass stars. \item Combining the results of LVG models for \CEIO\ and SiO, and using \CEIO\ as a tracer of total H$_2$ column density, a beam averaged SiO abundance is derived for all clouds. This varies significantly from source to source, ranging from $\sim 0.3\,10^{-9}$ to $> 5\,10^{-9}$. The $^{28}$SiO($2\to 1$) transitions are optically thick, with $\tau_0$ ranging from $4 - 8$. \item Including information on the temperature structure of the clouds from NH$_3$, it is shown that for most clouds a self-consistent solution accounting for the properties of \CEIO\ and SiO is only possible if the bulk of the SiO emission arises in the cool, dense gas component. \item One source, M+1.31$-$0.13, is different from all the others: SiO arises in the hot, thin gas component. Since this is also the cloud with the highest SiO abundance in the entire sample, the SiO formation process, probably grain erosion by shocks, is likely to be still ongoing in this source. \item The SiO abundance in individual clouds is related to the large scale gas dynamics in the Galactic center region. The highest abundances are found at Galactic longitudes of $l > 0.8^{\circ}$, which can be identified with `collision regions' likely to encounter shocks in terms of the bar model of gas dynamics. \item As in the disk, SiO in the Galactic center region is likely to originate in shocks. Since the hot, thin post-shock gas forms dense, cool cores faster than SiO recondenses to dust grains, the SiO rich gas is cool for most of its lifetime. On a timescale of $\sim 10^6$\,yr, the SiO molecules freeze out on grain mantles and the clouds lose their chemical memory of the shock. \end{enumerate} \acknowledgements{We thank the staff of the telescopes, especially L.\ Knee and L.-\AA.\ Nyman (SEST), and the SMTO/HHT staff for their help in taking the observations. GD gratefully acknowledges fiancial support from the Particle Physics and Astronomy Research Council (PPARC) (UK), for the observing run at JCMT. RM was supported by a Heisenberg fellowship by the Deutsche Forschungsgemeinschaft. 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