source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
The latter is very schematic and assumes a disk radius of ~ 1000 AU (Cesaroni et al. 2007)) | The latter is very schematic and assumes a disk radius of $\sim$ 1000 AU (Cesaroni et al. \cite{ppv}) ) |
and the same center and inclination as the mmaser jet (see Sect. 6.1)). | and the same center and inclination as the maser jet (see Sect. \ref{scone}) ). |
If the masers lie in the disk and co-rotate with it. the corresponding velocity vectors should also He in the plane of the disk. | If the masers lie in the disk and co-rotate with it, the corresponding velocity vectors should also lie in the plane of the disk. |
Since the latter is basically edge-on and the masers cluster Is at a (projected) distance of ~072 or 340 AU from the star. one would expect the proper motions to be relatively small (<4 s!.. assuming Keplerian rotation about a 7 M. star) and directed parallel to the disk. ie. NE-SW. | Since the latter is basically edge-on and the masers cluster is at a (projected) distance of $\sim0\farcs2$ or 340 AU from the star, one would expect the proper motions to be relatively small $\la4$ , assuming Keplerian rotation about a 7 $M_\odot$ star) and directed parallel to the disk, i.e. NE–SW. |
In contrast. the absolute proper motions measured by us have speeds of ο... aand are all roughly perpendicular to the disk. as shown in Fig. 2.. | In contrast, the absolute proper motions measured by us have speeds of $\sim$ 18 and are all roughly perpendicular to the disk, as shown in Fig. \ref{fmpm}. |
It is also worth noting that the dispersions of the vvelocities both along the line of sight and in the plane of the sky are small compared to the mean speed of the features. unlike mmasers whose velocities span an order of magnitude. | It is also worth noting that the dispersions of the velocities both along the line of sight and in the plane of the sky are small compared to the mean speed of the features, unlike masers whose velocities span an order of magnitude. |
Thus. we note that the mmasers are moving much slower than the mmasers. | Thus, we note that the masers are moving much slower than the masers. |
Our explanation of all these facts is thatorbit. | Our explanation of all these facts is that. |
Whether such a peculiar motion is due to a wrong estimate of the Solar motion with respect to the LSR. to an inadequate Galaxy rotation curve. or to a deviation from it. is irrelevant for our purposes. | Whether such a peculiar motion is due to a wrong estimate of the Solar motion with respect to the LSR, to an inadequate Galaxy rotation curve, or to a deviation from it, is irrelevant for our purposes. |
The important point is that the absolute proper motion of the mmasers can be used to correct any absolute proper motion measurement and determine the true motion with respect to the disk+star system to within about +4 ((the expected rotational velocity). | The important point is that the absolute proper motion of the masers can be used to correct any absolute proper motion measurement and determine the true motion with respect to the disk+star system to within about $\pm4$ (the expected rotational velocity). |
Incidentally. we note that deviations from the standard Galaxy rotation curve are quite common. as demonstrated by Reid et al. (2009b)) | Incidentally, we note that deviations from the standard Galaxy rotation curve are quite common, as demonstrated by Reid et al. \cite{reid09b}) ) |
in their study of the Galactic velocity field. | in their study of the Galactic velocity field. |
Note also that the existence of this “spurious” velocity component of - 18 iin the plane of the sky has less impact on mmasers. due to the much larger velocities of the features (up to - 100 y. | Note also that the existence of this “spurious” velocity component of $\sim$ 18 in the plane of the sky has less impact on masers, due to the much larger velocities of the features (up to $\sim$ 100 ). |
In conclusion. our mmeasurements can be used (combined with the corresponding Doppler shifts) to determine the 3D velocity of the start+disk system. | In conclusion, our measurements can be used (combined with the corresponding Doppler shifts) to determine the 3D velocity of the star+disk system. |
In practice. in order to obtain the velocity in the plane of the skystar. any absolute proper motion must be corrected by subtractingmasers. ie. -16 aalong the right-ascension axis and +7.6 aalong the declination. | In practice, in order to obtain the velocity in the plane of the sky, any absolute proper motion must be corrected by subtracting, i.e. $-16$ along the right-ascension axis and +7.6 along the declination. |
Better knowledge of the distance and. especially. of the systemic 3D velocity of aallows us to improve on the model presented by Moscadelli et al. (2005)). | Better knowledge of the distance and, especially, of the systemic 3D velocity of allows us to improve on the model presented by Moscadelli et al. \cite{mosca05}) ). |
Also. one can accurately compare the distributions and velocities of the mmasers with those of the mmasers. and thus draw a better picture of the disk and Jet structure and of their relationship. | Also, one can accurately compare the distributions and velocities of the masers with those of the masers, and thus draw a better picture of the disk and jet structure and of their relationship. |
Before proceeding to this comparison. it is worth noting that the two types of masers have been observed at very different epochs. namely November 9 2000 for the mmasers (Moscadell et al. 2005)) | Before proceeding to this comparison, it is worth noting that the two types of masers have been observed at very different epochs, namely November 9 2000 for the masers (Moscadelli et al. \cite{mosca05}) ) |
andNovember 6 2004 for the mmasers (this paper). | andNovember 6 2004 for the masers (this paper). |
The maserspotscould have moved over this ~4+ yr period. thus questioning a positional comparison between the two maser types. | The maserspotscould have moved over this $\sim$ 4 yr period, thus questioning a positional comparison between the two maser types. |
However. from the known proper | However, from the known proper |
name these catalogs FOFZ and FOF! with (n=0.1.2.3) for the DM and gas components respectively. | name these catalogs $^{DM}_n$ and $^{gas}_n$ with (n=0,1,2,3) for the DM and gas components respectively. |
For more details on (he simulation. the FOF identification and general cluster properties we refer the reader to Gottlóber&Yepes(2007). | For more details on the simulation, the FOF identification and general cluster properties we refer the reader to \citet{2007ApJ...664..117G}. |
For our analvsis we have selected all halos with masses larger than LottM... at [our different redshifts. z=0. 0.3 and 0.5. | For our analysis we have selected all halos with masses larger than $10^{14}$, at four different redshifts, $=0$, $0.3$ and $0.5$. |
The number of selected halos for every redshilt. and the hiehest halo mass are quoted in Table 2.. | The number of selected halos for every redshift, and the highest halo mass are quoted in Table \ref{table:lista}. |
Our objective is to measure the physical separation between the dominant gas clump in the cluster. with respect to its predominant DM structure. | Our objective is to measure the physical separation between the dominant gas clump in the cluster, with respect to its predominant DM structure. |
We find the separation based on the FOF catalogs described above. | We find the separation based on the FOF catalogs described above. |
The main advantage of the FOF algorithm is that it can be easily applied to the subsets of DM particles and gas particles and identifies gas chimps even 1f μον don have a pronounced density peak. | The main advantage of the FOF algorithm is that it can be easily applied to the subsets of DM particles and gas particles and identifies gas clumps even if they don't have a pronounced density peak. |
Moreover. the set. of linking leneths in the hierarchical FOF allows to find objects at. different. overdensities. | Moreover, the set of linking lengths in the hierarchical FOF allows to find objects at different overdensities. |
In order to test the robustuess of the results from the algorithm based on FOF catalogs we have compared them with those obtained using a spherical overclensity algorithm. the Amiga Lalo Finder. described in the next section. | In order to test the robustness of the results from the algorithm based on FOF catalogs we have compared them with those obtained using a spherical overdensity algorithm, the Amiga Halo Finder, described in the next section. |
We define the clusters as the objects detected in the DM FOF, catalog. | We define the clusters as the objects detected in the DM $_{0}$ catalog. |
FOF) provides objects with about. virial overdensitv. whereas FOF; provides substructures with about S higher densities. | $_0$ provides objects with about virial overdensity, whereas $_i$ provides substructures with about $8^i$ higher densities. |
In the MN-1024 simulation a tvpical object of 10!!2.IM. at the low mass end of our sample consists of 10000 DM particles and about the same number of gas particles. | In the MN-1024 simulation a typical object of $10^{14}
\hMsun$ at the low mass end of our sample consists of $\sim 10000$ DM particles and about the same number of gas particles. |
In the series of low resolution simulations we consider only clusters with masses larger than 3.2xLOMA!M. which consists of ~4000 DM particles and about the same number of gas particles. | In the series of low resolution simulations we consider only clusters with masses larger than $3.2 \times 10^{14} \hMsun$ which consists of $\sim 4000$ DM particles and about the same number of gas particles. |
In all cases we find (hat the clusters exhibit a massive dominant gas clump and various DM champs. | In all cases we find that the clusters exhibit a massive dominant gas clump and various DM clumps. |
In both DM auc gas catalogs we define the dominant clump position in the following wav. | In both DM and gas catalogs we define the dominant clump position in the following way. |
We identifv first the center and radius of the cluster as given by the DM FOF, catalog. | We identify first the center and radius of the cluster as given by the DM $_0$ catalog. |
Next we find (he most massive group in the FOF, catalog with its center inside (he radius defined previously. | Next we find the most massive group in the $_1$ catalog with its center inside the radius defined previously. |
The clump defines a new center and radius. | The clump defines a new center and radius. |
We iterate (his procedure until we cannot find a new eroup in the Following FOF; catalog. | We iterate this procedure until we cannot find a new group in the following $_i$ catalog. |
The outputs ol the algorithm are the mass and radius of the cluster defined from the DM FOF, catalog and the centers of mass of the dominant DM and gas clamps as well as their separation. | The outputs of the algorithm are the mass and radius of the cluster defined from the DM $_0$ catalog and the centers of mass of the dominant DM and gas clumps as well as their separation. |
lt is well known that the FOF algorithm at the level FOF) will connect two clusters that are starting to merge as one FOF object. | It is well known that the FOF algorithm at the level $_{0}$ will connect two clusters that are starting to merge as one FOF object. |
Since such cases could lead (o spurious large separations between the DAI ancl eas centers (μον have been excluded [rom the analvsis. | Since such cases could lead to spurious large separations between the DM and gas centers they have been excluded from the analysis. |
These cases represent always less than 3% of the total number of clusters. | These cases represent always less than $3\%$ of the total number of clusters. |
theLo clomain using the rotation curve of ?.. | the domain using the rotation curve of \citet{brand93}. |
We adopt tle solar rotation and radius as mettioned in the previous paragraph aud use these throughou the discussion. | We adopt the solar rotation and radius as mentioned in the previous paragraph and use these throughout the discussion. |
The choice of rotation curve. solar parameters aud spiral arm shape have minimal effect. but these are acIclressecl iu detai int appeudix. | The choice of rotation curve, solar parameters and spiral arm shape have minimal effect, but these are addressed in detail in the appendix. |
Sources already. associated with star [oriiug 'eglons towards the Galactic cent'e (S B2. the Galactic Centre Zoje and sites within 3.97spe of the Galactic centre. see ? for cletai ALL tlie parallel sections of he 3-kpc aruis between lougitudes 15° are clistineuisied by different symbols. | Sources already associated with star forming regions towards the Galactic centre (Sgr B2, the Galactic Centre Zone and sites within 3.5–kpc of the Galactic centre, see \citet{caswell10mmb1} for details) and the parallel sections of the 3–kpc arms between longitudes $\pm$ $^{\circ}$ are distinguished by different symbols. |
With these excluced from consideration. {1en iucorporaine an arm thickess of 1 kp wit La velocity tolerance of τι 1. we fiud the spiral arms account for of jase | With these excluded from consideration, then incorporating an arm thickness of 1 kpc with a velocity tolerance of $\pm$ $^{-1}$, we find the spiral arms account for of the masers. |
OT Moite Carlo simulation (frou 100 realisatious) gives a1r average (both mean aud medi: liuaber ¢of assxciated sources of wi tha staudard deviation of ---uplyit& the association ‘eal inase"s Insatistically siguilicant (noue of the Monte Carlo simulations exceed. associat. | The Monte Carlo simulation (from 100 realisations) gives an average (both mean and median) number of associated sources of with a standard deviation of, implying the association of real masers is statistically significant (none of the Monte Carlo simulations exceed association). |
We chose an arin thickuess of 1. kpe based on estimates of the i1ler-arni sepa‘ation aud typica idtlis of t1e aris iu models (seeforexample22??).. | We chose an arm thickness of 1 kpc based on estimates of the inter-arm separation and typical widths of the arms in models \citep[see for example][]{gomez04, sewilo04,mcclure04,levine06}. |
Varyiug the arm hickuess between 0.5 aud 1.5 kpc resulS in the level of association varying between LO and7S8%. | Varying the arm thickness between 0.5 and 1.5 kpc results in the level of association varying between 49 and. |
. Qir velocity tolerance of ct 1 Cds Chosen based ou the kinematic uucertainty of the masers iu ‘elation to the high-mass star formine 'eelon they are traciug (?22).. | Our velocity tolerance of $\pm$ $^{-1}$ is chosen based on the kinematic uncertainty of the masers in relation to the high-mass star forming region they are tracing \citep{reid09, pandian09}. |
Varyiug the velocity toleraice between land 'esults in an association of between 258 and8056. | Varying the velocity tolerance between $^{-1}$ and $^{-1}$ results in an association of between 58 and. |
.. The unissociated sources (approximately 205€)) ie almost exclusively iu the longitude regions 187 to 287 with velocities >LOC + and between —]9* aud —21° with velocities between ϐ and —130kkmss |. | The unassociated sources (approximately ) lie almost exclusively in the longitude regions $^{\circ}$ to $^{\circ}$ with velocities $>$ $^{-1}$ and between $-$ $^{\circ}$ and $-$ $^{\circ}$ with velocities between 0 and $-$ $^{-1}$. |
These sources will be «]scussed iu he following sectious. | These sources will be discussed in the following sections. |
TheLe ¢istribujon ralses the cuestion: what are the causes of the statistically sigtiicant over-densities — ure hey related to Calactic structure or raucdom fluctuatious in star formation? | The distribution raises the question: what are the causes of the statistically significant over-densities $-$ are they related to Galactic structure or random fluctuations in star formation? |
Reeious of star formaion are coitained within Ciaut Molecular Clouds (CINICs) which have pvsical sCales of I0s os parsecs. | Regions of star formation are contained within Giant Molecular Clouds (GMCs) which have physical scales of 10s of parsecs. |
With ie hiel-lnass star formatiou traced by the 6.7—CHz masers. we are loosing at reelOlls al ypleal distauces ol >2 kpc (he Lacistributiou within E28" suggest! iostly ES kpc). | With the high-mass star formation traced by the 6.7–GHz masers, we are looking at regions at typical distances of $\ge$ 2 kpc (the distribution within $\pm$ $^{\circ}$ suggesting mostly $\ge$ 4 kpc). |
At jese cliSances most CGICs subtend auguar scales stualler tinn 1. | At these distances most GMCs subtend angular scales smaller than $^{\circ}$. |
The coherent deise l'egions «X sources seen i Fiewe 1.. aud listed in Table 1.. are on scaes of »1* (hiiehitecd yy their preseice pos-sIn0olug in Figure 3 ) | The coherent dense regions of sources seen in Figure \ref{lvdistribution}, and listed in Table \ref{resotable}, are on scales of $>$ $^{\circ}$ (highlighted by their presence post-smoothing in Figure \ref{lvdensity}) ). |
This lases random. fluctuations. unred to aalactie scale processes. an unlikely cause of tje ]ncreased star formatior. | This makes random fluctuations, unrelated to Galactic scale processes, an unlikely cause of the increased star formation. |
This section analyses he 6.7-GHz 1jetlano inaserLe distribution i1 relatio ito the major st"uctural features of the inner Calaxy. | This section analyses the 6.7–GHz methanol maser distribution in relation to the major structural features of the inner Galaxy. |
We begi1 with the Galactic bar. hen the 3-kpe arms. aud ollow with tle starting volts of the spiral arius. or tLe spiral arn orieins. | We begin with the Galactic bar, then the 3–kpc arms, and follow with the starting points of the spiral arms, or the spiral arm `origins'. |
Excluded from the ¢ISCUSSIOLL are liisers al eh latitucles. as these are uulisely to lie withit Lkkpe of tje Galactic cenre. | Excluded from the discussion are masers at high latitudes, as these are unlikely to lie within kpc of the Galactic centre. |
With a scale helelit ol star formation estimated to be ~30— ppc (2).. ally Inasers wliose Latitudes tinply a Calactic eight. z. greater than £60—80ppe at a jeliocenric distance of Ekkpe cai be excluded from inner Galaxy cousklerations. | With a scale height of star formation estimated to be $\sim$ $-$ pc \citep{fish03}, any masers whose latitudes imply a Galactic height, z, greater than $\pm$ $-$ pc at a heliocentric distance of kpc can be excluded from inner Galaxy considerations. |
This correspoxls to latitudes 0.587 to 1.0° aud we tierefore adopt 0.97 as a cut-off. | This corresponds to latitudes $^{\circ}$ to $^{\circ}$ and we therefore adopt $^{\circ}$ as a cut-off. |
Within the lougitude region d EPS” there are 29 masers with latitudes in excess of 0.97 [Οι | Within the longitude region $\pm$ $^{\circ}$ there are 29 masers with latitudes in excess of $^{\circ}$ from |
snapshot duration. | snapshot duration. |
The observing log is reported in Table 1. where we list the targets. the observing date. the data set uae. the spectral clement used in the observations. and the exposure time aud ummber of exposures obtained. | The observing log is reported in Table 1, where we list the targets, the observing date, the data set name, the spectral element used in the observations, and the exposure time and number of exposures obtained. |
Aliases ave identified in the table notes. | Aliases are identified in the table notes. |
The STIS data were calibrated using the standard pipeline svstem. as in the LMC data (Shawctal.2001). | The STIS data were calibrated using the standard pipeline system, as in the LMC data \citep{sha01}. |
. Figures 1 through 6 show the observed SAIC PNs in the three observing modes: left panels show the broad band CCD images: central paucls show the aud images: right paucls show. when available. the 5007 images. | Figures 1 through 6 show the observed SMC PNs in the three observing modes: left panels show the broad band CCD images; central panels show the and images; right panels show, when available, the 5007 images. |
Spectral analysis of the SAIC PNs have been performed similarly to that of LAC PNs 2002). | Spectral analysis of the SMC PNs have been performed similarly to that of LMC PNs \citep{sta02}. |
. For most PNs observed. the combination of dispersion aud spatial scale allows a clear separation of the nonochromatic dmaees for all cussion lines. | For most PNs observed, the combination of dispersion and spatial scale allows a clear separation of the monochromatic images for all emission lines. |
Exceptions are J 27. where broad aud monochromatic inages are at the limit of detectability. and MÀ 1682. where the aud images may have partial overlap. | Exceptions are J 27, where broad and monochromatic images are at the limit of detectability, and MA 1682, where the and images may have partial overlap. |
No innages are severely overlapped. thus the one dimensional spectral extraction was adequate for nebula line Hux and ratio analysis. | No images are severely overlapped, thus the one dimensional spectral extraction was adequate for nebular line flux and ratio analysis. |
We extract the ouc-dimeusional spectra aud applied a photometric calibration using he standard STIS calibration pipelineLB module (MeCGrath.Busko.&IIodge1999). | We extract the one-dimensional spectra and applied a photometric calibration using the standard STIS calibration pipeline module \citep{McGrath_etal99}. |
. We used extraction )Xxes for the nebulae laree enough to encompass all the nebular features. but suug enough as to exclude uost of the sky background from the extraction. | We used extraction boxes for the nebulae large enough to encompass all the nebular features, but snug enough as to exclude most of the sky background from the extraction. |
Sky background regious were selected for each object to avoid stray stellar photons from field stars. | Sky background regions were selected for each object to avoid stray stellar photons from field stars. |
The backeround was then averaged aud subtracted. | The background was then averaged and subtracted. |
We measured emission line intensities with task. fitting eaussiaus to individual lines. while estimating the continu level. | We measured emission line intensities with task, fitting gaussians to individual lines, while estimating the continuum level. |
In the cases in which the enüssion ues were notably uou eaussian. we estimated the liue flix as measured from the area above the coutiuuuu. | In the cases in which the emission lines were notably non gaussian, we estimated the line flux as measured from the area above the continuum. |
Iu Table 2 we report the measured line intensities for the SAIC PNs in our sample. | In Table 2 we report the measured line intensities for the SMC PNs in our sample. |
Column (1) gives the conunon names: column (2) gives the logarithinic intensities. not corrected for oxtinction. in erg ci7 +: column (3) lists the logarithmic optical extinction at COsterbrock1989): colunums (1) to (11) give the line intensitics for cach nebula. relative to —1100. not corrected for extinction. | Column (1) gives the common names; column (2) gives the logarithmic intensities, not corrected for extinction, in erg $^{-2}$ $^{-1}$; column (3) lists the logarithmic optical extinction at \citep{ost89}; columns (4) to (14) give the line intensities for each nebula, relative to 100, not corrected for extinction. |
Line ideutification are eiven at the heads of the colimus. | Line identification are given at the heads of the columns. |
From the aualvsis of our spectral line measurements. aud eiven the similarities of this data set to the ones analyzed in Staughelliuietal.(20023.. we confini that the errors in the line iuteusities of Table 2 (OlogF) ive Gin dex): ólosF«(05 ifloeP>12.25: dloeP< .15if 12.25>logF12.75: dloeF«.2 if 12.75>loeF 135:0logF< 25if (3.5>loeFILLS: and dloeF «.55iflogF<11.5. | From the analysis of our spectral line measurements, and given the similarities of this data set to the ones analyzed in \citet{sta02}, we confirm that the errors in the line intensities of Table 2 $\delta {\rm log} F$ ) are (in dex): $\delta {\rm log} F < .05$ if ${\rm log} F > -12.25$; $\delta {\rm log} F<.15$ if $-12.25 > {\rm log} F > -12.75$; $\delta {\rm log} F < .2$ if $-12.75 > {\rm log} F >-13.5$; $\delta {\rm log} F< .25$ if $-13.5 > {\rm log F} > -14.5$; and $\delta {\rm log} F<.55$ if ${\rm log} F < -14.5$. |
Iu Figure 7 we compare the measured line mteusities with those available iu the literature. | In Figure 7 we compare the measured line intensities with those available in the literature. |
We used the data frou selected references (Dopita&Meatheringhaim1990.1991a.b:Vassiliadisetal. | We used the data from selected references \citep{dop90,dop91a,dop91b,vdm}. |
1992).. Some authors eive the intensity corrected for extinction. so before comparison to the prescut results we uncorrect the intensity ratios by using the extinction coustaut eiven iu the same reference. | Some authors give the intensity corrected for extinction, so before comparison to the present results we un–correct the intensity ratios by using the extinction constant given in the same reference. |
We have comparisous for several emission lines in MG δν MG. 13. SMP 1. SMP 6. SMP 11. SMP 13. SMP L1. SMP 17. SMP 25. aud SAMIP 26. | We have comparisons for several emission lines in MG 8, MG 13, SMP 1, SMP 6, SMP 11, SMP 13, SMP 14, SMP 17, SMP 25, and SMP 26. |
A comparison of our fluxes to with previous values (Fig. | A comparison of our fluxes to with previous values (Fig. |
o 7) shows a generally οσους agreement. | 7) shows a generally good agreement. |
The | The |
?3) to break the near-far kinematic distance ambiguity towards 730 molecular clouds (~90 per cent of all clouds identitied). | \citealt{stil2006}) ) to break the near-far kinematic distance ambiguity towards 750 molecular clouds $\sim$ 90 per cent of all clouds identified). |
The combination of a complete set of molecular clouds and kinematic distances provides an enormously important catalogue that can be used for determining distances to star-forming regions. allowing their distribution. luminosity functions. ete.. | The combination of a complete set of molecular clouds and kinematic distances provides an enormously important catalogue that can be used for determining distances to star-forming regions, allowing their distribution, luminosity functions, etc., |
to be mapped with respect to the large-scale Galactic structure. | to be mapped with respect to the large-scale Galactic structure. |
Obtaining heliocentric distances to our sources is a crucial element of our campaign as without these we cannot accurately calculate source luminosities. and thus. distinguish between nearby low and intermediate-mass YSOs from the more generally more distance MYSOs. | Obtaining heliocentric distances to our sources is a crucial element of our campaign as without these we cannot accurately calculate source luminosities, and thus, distinguish between nearby low and intermediate-mass YSOs from the more generally more distance MYSOs. |
As mentioned in the introduction kinematic distance ambiguities attect a large proportion of our sample («900 sources in total) and these need to be resolved before we obtain a robus sample of voung massive stars. | As mentioned in the introduction kinematic distance ambiguities affect a large proportion of our sample $\sim$ 900 sources in total) and these need to be resolved before we obtain a robust sample of young massive stars. |
There have been a number of studies devoted to resolving distances ambiguities using a variety of methods. | There have been a number of studies devoted to resolving distances ambiguities using a variety of methods. |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.