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Phe lowest axisvnimetric mode has j=2 and setting we find
The lowest axisymmetric mode has $j=2$ and setting we find
radius of ag=1«109 un we can estimate the total dust mass. AM. of the cloud using Tuscrting the scattering cross-section we nieasure for Schiila’s dust cloud. we tls obtain an approximate dust mass of Ma~3«101 ke.
radius of $a_d=1\times10^{-6}$ m, we can estimate the total dust mass, $M_d$, of the cloud using Inserting the scattering cross-section we measure for Scheila's dust cloud, we thus obtain an approximate dust mass of $M_d\sim3\times10^7$ kg.
This estimate is of course based ol nuuerous assuuptious. particularly average dust erain size and bulk density. as specified above. and is only computed to provide approximate plysical contest to the photometric excess that was measured.
This estimate is of course based on numerous assumptions, particularly average dust grain size and bulk density, as specified above, and is only computed to provide approximate physical context to the photometric excess that was measured.
For reference. cluploving the differeut erain size distribution and deusitv assunptious made bv Jewittetaf(2011) (p=2000 kein Pag=110 9) andBodewitsefal.(2011) (p=2500 ke Frayὃν10tan). we would instead obtain AZ~1«10° ke and AL,~5«10? ke. respectively.
For reference, employing the different grain size distribution and density assumptions made by \citet{jew11} $\rho=2000$ kg $^{-3}$; $a_d=1\times10^{-6}$ m) and\citet{bod11} $\rho=2500$ kg $^{-3}$; $a_d=1\times10^{-4}$m), we would instead obtain $M_d\sim4\times10^7$ kg and $M_d\sim5\times10^9$ kg, respectively.
For comparison. Larson(2010) imeasured 1.24 mag on 2010 December 3. while Bodewitsefal,(2011) measured Aimy=0.66 mag on 2010 December LELS and Jewittetef(2011) measured Any=1.26 mag on 2010 December 28 and Ani=1.00 1uag ou 2011 January 5.
For comparison, \citet{lar10} measured $\Delta m_V=1.24$ mag on 2010 December 3, while \citet{bod11} measured $\Delta m_V=0.66$ mag on 2010 December 14-15 and \citet{jew11} measured $\Delta m_V=1.26$ mag on 2010 December 28 and $\Delta m_V=1.00$ mag on 2011 January 5.
However. given the differeut observational and instrumental circumstances ivolved (especially given the technical difficulty of observiug the extremely bright nucleus aud comparatively extremely faint dust cloud simultancously). with the exception of the decline between the two measurements by Jewittct (2011).. we do not consider auv dust nass fluctuations nuplied by comparing these disparate data sets to he reliable.
However, given the different observational and instrumental circumstances involved (especially given the technical difficulty of observing the extremely bright nucleus and comparatively extremely faint dust cloud simultaneously), with the exception of the decline between the two measurements by \citet{jew11}, we do not consider any dust mass fluctuations implied by comparing these disparate data sets to be reliable.
VoR color measurements (Table 1)) of the dust-contaiuinated nucleus iudicate that it is redder than the Sun. with the entire dust cloud with uucleus flux subtracted ("Dust4 in Table 1)) having au effectively identical red. color. simular to colors measured for other active comoets and active Centaurs (Bodewitsefo£2011:Jewitt 2009)..
$V-R$ color measurements (Table \ref{photom}) ) of the dust-contaminated nucleus indicate that it is redder than the Sun, with the entire dust cloud with nucleus flux subtracted $_{\rm A}$ ” in Table \ref{photom}) ) having an effectively identical red color, similar to colors measured for other active comets and active Centaurs \citep{bod11,jew09b}.
Assunüug that the dust contaminating our uucleus photometry has a scatteriug cross-section larger than that of the nucleus (estimated above) aud the same color as the dust cloud as a whole. we fud a dust-subtracted color for the nucleus of VRocOL mag. consistent with a D-type spectral classification (FornasieretaL2007).
Assuming that the dust contaminating our nucleus photometry has a scattering cross-section larger than that of the nucleus (estimated above) and the same color as the dust cloud as a whole, we find a dust-subtracted color for the nucleus of $V-R\approx0.44$ mag, consistent with a D-type spectral classification \citep{for07}.
. To test whether the northern aud southern phunes of Scheila’s dust cloud exhibit any compositional differences. we mcasiure their colors individually παμε aud "Ότο). but within estimated uncertainties. find no significant color cüffereuces.
To test whether the northern and southern plumes of Scheila's dust cloud exhibit any compositional differences, we measure their colors individually $_{\rm B}$ ” and $_{\rm C}$ ”), but within estimated uncertainties, find no significant color differences.
The iost sensitive probe of sublimating eas in a comet is CN omission at38s0A.
The most sensitive probe of sublimating gas in a comet is CN emission at.
. We show the spectral iuage of Scheila from 3700A to LLOOA in Figure Jaa. Iu this tage. the horizoutal continumu corresponds to reflected light from the nucleus.
We show the spectral image of Scheila from $3700{\rm \AA}$ to $4100{\rm \AA}$ in Figure \ref{keck_obs}a a. In this image, the horizontal continuum corresponds to reflected light from the nucleus.
OIT sky: cmission lines are visible as vertical bands; aud dark Eraunuhofer lines in the Solar spectrum aud promincut Ca II (3933À)) and Ik (3966A)) absorption lines are also visible.
OH sky emission lines are visible as vertical bands, and dark Fraunhofer lines in the Solar spectrum and prominent Ca H ) and K ) absorption lines are also visible.
Iu a two-dimensional spectral image. the intensity. ο. of a cometary cluission lines should be highest near the ceuter of the σοιπμ and eradually decrease with increasing distance frou the coutinuun. moving iu the spatial direction.
In a two-dimensional spectral image, the intensity, $I_e$, of a cometary emission lines should be highest near the center of the continuum and gradually decrease with increasing distance from the continuum, moving in the spatial direction.
No spectral features near 3880À exhibit such behavior. and we therefore conclude that no gas is detected.
No spectral features near ${\rm\AA}$ exhibit such behavior, and we therefore conclude that no gas is detected.
We also search for CN cussion iu a one-dimensional spectrum extracted frou the spectral inaee using a UO.S71 rectangular aperture centered ou the continu.
We also search for CN emission in a one-dimensional spectrum extracted from the spectral image using a $1\farcs0\times8\farcs1$ rectangular aperture centered on the continuum.
Sky backeround was measured aud subtracted using flanking regions 10" to 16" from the nucleus.
Sky background was measured and subtracted using flanking regions $''$ to $''$ from the nucleus.
Calibration was performed using a nearby fux standard star aud a solar analog star (Fig.
Calibration was performed using a nearby flux standard star and a solar analog star (Fig.
το),
\ref{keck_obs}b b).
Shaded regious in Figures ?bb and Tec indicate where CN omission is expected. but we again find no evidence of enission. consistent with work by Bodewitsefaf.(2011).. Howell&Lovell(2011)..aud Jehinetaf.(2011).
Shaded regions in Figures \ref{keck_obs}b b and \ref{keck_obs}c c indicate where CN emission is expected, but we again find no evidence of emission, consistent with work by \citet{bod11}, , \citet{how11},and \citet{jeh11}.
. To estimate Scheila’s CN production rate. we remove the continui using a scaled solar analog spectruui (Fie.
To estimate Scheila's CN production rate, we remove the continuum using a scaled solar analog spectrum (Fig.
του). which should then leave ouly gas cussion.
\ref{keck_obs}c c), which should then leave only gas emission.
Standard errors in three wavelength reeious in the residual spectrum 3830Ar 3900A.. where CN cinission is expected: andΙΟΡΟΑ:: each in width) are 2.7«10.L0, 2.0&«104) and Laslo!orgcn 2s! A!vespectively (Fig.
Standard errors in three wavelength regions in the residual spectrum;, where CN emission is expected; and; each in width) are $2.7\times10^{-17}$, $2.0\times10^{-17}$, and $1.8\times10^{-17}~{\rm erg~cm}^{-2}$ $^{-1}$ $^{-1}$,respectively (Fig.
Tec).
\ref{keck_obs}c c).
We choose 2.7.10 cre ? s| At as a conservative estimate of the wneertainty in the CN baud.
We choose $2.7\times10^{-17}$ erg $^{-2}$ $^{-1}$ $^{-1}$ as a conservative estimate of the uncertainty in the CN band.
Since 9P/Tempel 1 observations (Aleechetad.WwH1) utilized the same iustruuicutal settings used here. we assume that any CN band iu Scheila’s spectu will have the same profile as in 9P's spectrmm.
Since 9P/Tempel 1 observations \citep{mee11} utilized the same instrumental settings used here, we assume that any CN band in Scheila's spectrum will have the same profile as in 9P's spectrum.
We therefore fud a peak value of any CN baud of (2.7s10PL ee en2 1A 1) or oe 2s PAL,
We therefore find a peak value of any CN band of $3\times(2.7\times10^{-17}$ erg $^{-2}$ $^{-1}$ $^{-1}$ ), or $\sim8\times10^{-17}$ erg $^{-2}$$^{-1}$ $^{-1}$.
We calculate the inteerated CN baud fux. fox. by sunning the cmission flux iu the shaded region. obtaining foxSS1019 erg 72 1.
We calculate the integrated CN band flux, $f_{\rm CN}$ by summing the emission flux in the shaded region, obtaining $f_{\rm CN}=8.8\times10^{-16}$ erg $^{-2}$ $^{-1}$.
We then convert fex to the total παπανο of CN molecules. Ver. usine where A and ry, are in cm aud AU. respectively. and g(7) is the resonance fluorescence efficiency. which describes the wmmber of photous scattered per second per radical. in cre s +.
We then convert $f_{\rm CN}$ to the total number of CN molecules, $N_{\rm CN}$, using where $\Delta$ and $r_h$ are in cm and AU, respectively, and $g(l)$ is the resonance fluorescence efficiency, which describes the number of photons scattered per second per radical, in erg $^{-1}$ $^{-1}$.
Diving our observations. Scheila had a radial velocity of ὃν=2.5 km lofor which g(lrj)=2.7«101 eve st + when the Swings effect is taken iuto account(Schleicher 2010)..
During our observations, Scheila had a radial velocity of $\dot{r}_h=2.5$ km $^{-1}$ , for which $g(l,r_h)=2.7\times10^{-13}$ erg $^{-1}$ $^{-1}$ when the Swings effect is taken into account\citep{sch10}. .
Substituting fex=s.8«1016 ere 2s ο, we obtain Nox=5.33s1025.
Substituting $f_{\rm CN}=8.8\times10^{-16}$ erg $^{-2}$ $^{-1}$, we obtain $N_{\rm CN}=5.33\times10^{26}$.
A simple Haser(1957) model is used to derive the CN production rate. Qox. from New. assuming isotropic outgassing.constant radial expansion of the gas coma. aud a 2-step exponeutial decay process.
A simple \citet{has57} model is used to derive the CN production rate, $Q_{\rm CN}$, from $N_{\rm CN}$ , assuming isotropic outgassing,constant radial expansion of the gas coma, and a 2-step exponential decay process.
We use 7,=τον and p,-223«4102 as the effective Haser scale leugthbs af rj=1 AU (AIlemetαἱ. 1995).. and adopt a eas volocitv of (Diverctal. 1997)..
We use $l_{p}=1.3\times10^{4}$ and $l_{p}=2.2\times10^{5}$ as the effective Haser scale lengths at $r_{h}=1$ AU \citep{ahe95}, , and adopt a gas velocity of \citep{biv97}. .
Integratiug the computed spatial cohunn density model over a rectangle 1800 Ίαν «
Integrating the computed spatial column density model over a rectangle 1800 km $\times$
ddo not differ too ereatly from those obtained with he red continu measured from the Fabry-Perot data cubes.
do not differ too greatly from those obtained with the red continuum measured from the Fabry-Perot data cubes.
The typical difference is 3 wwhich fs within the typical uncertainties for our uecasureinenuts.
The typical difference is 3 which fits within the typical uncertainties for our measurements.
This sugecsts that star formation rigecring effects. although locally nuportaut. do mot change the overall weights applied to the velocities. (
This suggests that star formation triggering effects, although locally important, do not change the overall weights applied to the velocities. (
IT) We consider velocity deviations due to density waves by naking sure that we derive consistent values whem we nove along a given slit or across the disk.
II) We consider velocity deviations due to density waves by making sure that we derive consistent values when we move along a given slit or across the disk.
This strongly sugeests that such velocity deviations do cancel to first order.
This strongly suggests that such velocity deviations do cancel to first order.
The ideal form of the Tremaine-Weinberg method uses he luninosity-weighted velocities and the distances in he full radial range of the galaxy disk. ic. the iutegrals ranee from κ. to ox.
The ideal form of the Tremaine-Weinberg method uses the luminosity-weighted velocities and the distances in the full radial range of the galaxy disk, i.e. the integrals range from $-\infty$ to $\infty$.
Observations. ou the other haud. cover oulv a limited radial range. aud ours reach typically he ος radius of the galaxy disk.
Observations, on the other hand, cover only a limited radial range, and ours reach typically the $r_{25}$ radius of the galaxy disk.
It is thus important to see whether our observations are deep enough so that hey do not vield correct
It is thus important to see whether our observations are deep enough so that they do not yield incorrect.
We uote that the method gives pattern speeds strictly when the nuits of the integrals reach well bevoud the bar.
We note that the method gives pattern speeds strictly when the limits of the integrals reach well beyond the bar.
However. iu lis test. we restrict the radial range of integration iu xogressive steps. starting from au inucrimost radius of LO jxxels and increasing in steps of LO pixels uutil the eutire observed disk is covered (see Fig. 62).
However, in this test, we restrict the radial range of integration in progressive steps, starting from an innermost radius of 10 pixels and increasing in steps of 10 pixels until the entire observed disk is covered (see Fig. \ref{fig:radialdependency}) ).
Each derivation vields a differeut. (V2/6X5 value. treated as if it were a n for ;(see Fig. ον 7)).
Each derivation yields a different $\langle V \rangle$ $\langle X \rangle$ value, treated as if it were a value for (see Fig. \ref{fig:radialOpvalues}) ).
As shown in the detailed analvsis by these values decrease asviuptotically as the disk coverage increases.
As shown in the detailed analysis by \citet{Zimmeretal2004} these values decrease asymptotically as the disk coverage increases.
In particular the differences between he outermost values of these trials and the values obtained using the complete disks are much simaller than hose between the values obtained linitine the feld to a few tens of pixels aud those fouud when using the full disk region.
In particular the differences between the outermost values of these trials and the values obtained using the complete disks are much smaller than those between the values obtained limiting the field to a few tens of pixels and those found when using the full disk region.
Furthermore. we confini that by using all the jxxels inside the slit G.c.. without the radial truication as shown in Fig. 6)).
Furthermore, we confirm that by using all the pixels inside the slit (i.e., without the radial truncation as shown in Fig. \ref{fig:radialdependency}) ),
the derived. Q,s are unich inore stable.
the derived s are much more stable.
This is in support of the result by ? who found that uost of the variation in the derived patteri speed occurs within au inner radius which they set to Tor their observations.
This is in support of the result by \citet{Zimmeretal2004} who found that most of the variation in the derived pattern speed occurs within an inner radius which they set to for their observations.
A pattern speed derived over a radial. rauge sienificautly shorter than the bar leneth will not eive a ucaningful value since the integrals do not converge. so we do not wish to justify here any variations (or indeed absence of variations) calculated within these ranges.
A pattern speed derived over a radial range significantly shorter than the bar length will not give a meaningful value since the integrals do not converge, so we do not wish to justify here any variations (or indeed absence of variations) calculated within these ranges.
Moreover. as foundbv ? for 551. quasi-steady state continuity is uot satisfiedby a single valued pattem speed. so it is nuaportaut to analyse the derived vvariatious to male sure they converge on a unique ονvalue, which cau then be taken as the main value for a single and dominant i)=2 perturbation (2)..
Moreover, as found by \citet{Shettyetal2007} for 51, quasi-steady state continuity is not satisfied by a single valued pattern speed, so it is important to analyse the derived variations to make sure they converge on a unique value, which can then be taken as the main value for a single and dominant $m=2$ perturbation \citep{Henryetal2003}.
In our work we have used integrations over the full observed field for all ealaxies to derive the pattern speeds preseuted in Table and Fig. 1
In our work we have used integrations over the full observed field for all galaxies to derive the pattern speeds presented in Table \ref{tab:patternspeeds} and Fig. \ref{fig:allmaps}.
Iu an ongoing study of the uuuethod based ou N-body SPI uunerical simmlatious carried out by Isabel Pérez |(Pérezctal.inpreparation.seealso7). we find that in the case of oue constant pattern speed ina disk. scans of uutruucated pseudo-slits in the Tuner regions as well as outer regions. all recover the input oofthe simmlatious.
In an ongoing study of the method based on N-body + SPH numerical simulations carried out by Isabel Pérrez \citep[P\'erez et al. in preparation, see also ][]{Beckmanetal2008} we find that in the case of one constant pattern speed in a disk, scans of untruncated pseudo-slits in the inner regions as well as outer regions, all recover the input of the simulations.
©, When truncating the slits to cut the outer sections of the disk (compare with Fie. 6)).
When truncating the slits to cut the outer sections of the disk (compare with Fig. \ref{fig:radialdependency}) ),
the inteerals applied to immer truncated slits result in different (V2d /C6X5 values when compared with using the full ίατο ealaxy.
the integrals applied to inner truncated slits result in different $\langle V \rangle$ $\langle X \rangle$ values when compared with using the full simulated galaxy.
This i$ a clear demonstration of the necessity of using full slit coverage aud deep observations (even in the iuner regious) due to the non-convereius integrals.
This is a clear demonstration of the necessity of using full slit coverage and deep observations (even in the inner regions) due to the non-converging integrals.
However. as expected. we have found. that as soon as the radial disk section is comparable with or larger than the size of the bar. the input iuto the simulations can be recovered.
However, as expected, we have found that as soon as the radial disk section is comparable with or larger than the size of the bar, the input into the simulations can be recovered.
Ὃν This feature is shared by real data as well as by umuerical simulations.
This feature is shared by real data as well as by numerical simulations.
Furthermore. we have applied the nuuethod using the simulated velocity ficld weighted by the simulated density maps as well as the simulated star-formation maps and found that although using the density iaps provide more accurate results. the star-formation maps also deliver comparable (to within the errors) pattern speeds. Πρίντις that the vvolocities are reliable when applying the uuuethod.
Furthermore, we have applied the method using the simulated velocity field weighted by the simulated density maps as well as the simulated star-formation maps and found that although using the density maps provide more accurate results, the star-formation maps also deliver comparable (to within the errors) pattern speeds, implying that the velocities are reliable when applying the method.
For the galaxies preseuted here. Table 1 shows that the aad coutiuuun deliver comparable (note that the continuum maps used for this exercise also conie from our data cubes).
For the galaxies presented here, Table \ref{tab:patternspeeds} shows that the and continuum deliver comparable (note that the continuum maps used for this exercise also come from our data cubes).
The exception is NGC L519. where the contimuuu nuage assigns higher weights to the velocities imside the bar. and thereby iutroduces significaut difference between the dderived using the contiuuuuau and2, nunaps.
The exception is NGC 4519, where the continuum image assigns higher weights to the velocities inside the bar, and thereby introduces significant difference between the derived using the continuum and maps.
The higher pattern speed of the bar. using the οὐΠτα image delivers a higher total£,.
The higher pattern speed of the bar, using the continuum image delivers a higher total.
. Tn the presence of multiple structures with different patteru speeds. when deriving the VW) /CN) for au inner component. such as an inner bar inside a large bar. by applving the notional slits on the iuner parts of a two-dimensional map. the majority of the pixels will belong to he disk region outside the bar.
In the presence of multiple structures with different pattern speeds, when deriving the $\langle V \rangle$ $\langle X \rangle$ for an inner component, such as an inner bar inside a large bar, by applying the notional slits on the inner parts of a two-dimensional map, the majority of the pixels will belong to the disk region outside the bar.
These pixels will therefore affect the derived: pattern speed of the immer structure.
These pixels will therefore affect the derived pattern speed of the inner structure.
This effect was tested by c.g. ?.theirFie.7 who fouud hat once the slit coverage is comparable with the size of he main bar. the inethod delivers reliable. Q,s. Moreover. cutting the disk in radial bius secuis to reveal the presence of multiple ottern speeds when the disk section is comparable to enu with a correspondingly distinct pattern speed.
This effect was tested by e.g., \citet[][ their Fig.~7 ]{Zimmeretal2004} who found that once the slit coverage is comparable with the size of the main bar, the method delivers reliable s. Moreover, cutting the disk in radial bins seems to reveal the presence of multiple pattern speeds when the disk section is comparable to structure with a correspondingly distinct pattern speed.
Iu Fig.
In Fig. \ref{fig:allmaps}, ,
we illustrate the (V5. CX poiuts from the Inner regions of two galaxies where this 5effect can be seen.
we illustrate the $\langle V \rangle$, $\langle X \rangle$ points from the inner regions of two galaxies where this effect can be seen.
ILowever. as the derivation of secondary pattern speeds Is nore complicated (seee.g...2).. this test is to be used with caution iu those cases.
However, as the derivation of secondary pattern speeds is more complicated \citep[see e.g., ][]{Meidtetal2008}, this test is to be used with caution in those cases.
Finally. we contirui that the dadiffereunces using the continu©, maages instead of the ssurface briehtuess maps. does not significantly chauge the location of the resonances throughout our analysis.
Finally, we confirm that the differences using the continuum images instead of the surface brightness maps, does not significantly change the location of the resonances throughout our analysis.
Morphologically. the bar radius r(bar)is commonly determined by analysing cllipse fitting profiles
Morphologically, the bar radius r(bar)is commonly determined by analysing ellipse fitting profiles
may contribute to future growth in massive galaxies.
may contribute to future growth in massive galaxies.
We are also pursuing a companion project on the IRAM 30m telescope, COLDGASS?,, which will obtain accurate and homogeneous molecular gas masses for a subset of ~300 galaxies from the GASS sample.
We are also pursuing a companion project on the IRAM 30m telescope, COLD, which will obtain accurate and homogeneous molecular gas masses for a subset of $\sim 300$ galaxies from the GASS sample.
These data will allow us to characterize the balance between atomic and molecular gas in the galaxies in our sample, and understand the physical processes that determine how the condensed baryons are partitioned into stars, HI and Hg in the local Universe.
These data will allow us to characterize the balance between atomic and molecular gas in the galaxies in our sample, and understand the physical processes that determine how the condensed baryons are partitioned into stars, HI and $_2$ in the local Universe.
In this paper, we report on UGC8802, an extraordinary galaxy blindly selected for inclusion in the GASS sample (under the catalog name GASS35981, used hereafter), which contains a reservoir of HI >1010 Mo, at least equal in mass to the galaxy's entire stellar content.
In this paper, we report on UGC8802, an extraordinary galaxy blindly selected for inclusion in the GASS sample (under the catalog name GASS35981, used hereafter), which contains a reservoir of HI $>10^{10}$ $_\odot$, at least equal in mass to the galaxy's entire stellar content.
This galaxy contains less than one tenth this mass in H5 and also has a rather modest star formation rate (SFR).
This galaxy contains less than one tenth this mass in $_2$ and also has a rather modest star formation rate (SFR).
We describe the spectroscopic follow-up that enables us to conclude that the outer disk of this galaxy is currently forming from gas that has likely accreted from the external environment.
We describe the spectroscopic follow-up that enables us to conclude that the outer disk of this galaxy is currently forming from gas that has likely accreted from the external environment.
In the following, we adopt a standard ACDM cosmology with Hg= 70km s! Mpc-!, Qm=0.3 and Q4=0.7.
In the following, we adopt a standard $\Lambda$ CDM cosmology with $_0=70$ km $^{-1}$ $^{-1}$, $\Omega_m=0.3$ and $\Omega_\Lambda=0.7$.
GASS35981 (also SDSS J135308.36+354250.5, in addition to UGC8802), was selected for inclusion in the GASS parent sample because it has photometry from both SDSS and GALEX, is located ina region of sky accessible to Arecibo, and has stellar mass of M,=2x10!°Mo and redshift of z=0.0411 that fit into our targeted range.
GASS35981 (also SDSS $+$ 354250.5, in addition to UGC8802), was selected for inclusion in the GASS parent sample because it has photometry from both SDSS and GALEX, is located ina region of sky accessible to Arecibo, and has stellar mass of $_*=2 \times 10^{10} {\rm M}_\odot$ and redshift of $z=0.0411$ that fit into our targeted range.
GASS35981 has pre-existing HI observations available from the Cornell HI Digital Archive (Springob et al.
GASS35981 has pre-existing HI observations available from the Cornell HI Digital Archive (Springob et al.
2005).
2005).
The mass of HI in GASS35981 is estimated from the line flux tobe 2.1x 10!°Mo, and the rest-frame velocity width of the HI line is Wso=360+25 km s! (Figure 1)).
The mass of HI in GASS35981 is estimated from the line flux tobe $2.1\times 10^{10}$ $_\odot$, and the rest-frame velocity width of the HI line is $W_{50}=360\pm25$ km $^{-1}$ (Figure \ref{gas_prof}) ).
We display the HI archive spectrum in Figure 1..
We display the HI archive spectrum in Figure \ref{gas_prof}.
We note that there is no evidence for contamination from possible companion galaxies within the 3.5’ Arecibo beam: one nearby galaxy has a spectroscopic redshift of 0.14, and two additional faint companions have SDSS photometric redshifts consistent with z=0.14.
We note that there is no evidence for contamination from possible companion galaxies within the $\arcmin$ Arecibo beam: one nearby galaxy has a spectroscopic redshift of 0.14, and two additional faint companions have SDSS photometric redshifts consistent with $z=0.14$.
The lower panel of Figure 1 shows that GASS35981 (blue star) lies near the extreme end of HI fractions observed by GASS.
The lower panel of Figure \ref{gas_prof} shows that GASS35981 (blue star) lies near the extreme end of HI fractions observed by GASS.
Indeed, its gas fraction is comparable to the highest values measured for all galaxies in this stellar mass range (e.g., Giovanelli et al.
Indeed, its gas fraction is comparable to the highest values measured for all galaxies in this stellar mass range (e.g., Giovanelli et al.
2007).
2007).
Indeed, even in Hl-selected samples, which are biased towards objects like GASS35981, galaxies in this stellar mass range >1010 with such high gas fractions are quite (M,uncommon Mo)(e.g., the “HI Giants” described by Garcia-Appadoo et al.
Indeed, even in HI-selected samples, which are biased towards objects like GASS35981, galaxies in this stellar mass range $_*>10^{10}$ $_\odot$ ) with such high gas fractions are quite uncommon (e.g., the “HI Giants” described by Garcia-Appadoo et al.
2009).
2009).
It is also clear that GASS35981 lies significantly above the best-fit ‘gas fundamental plane’ (dotted line) relating stellar mass surface density,NUV-r color, and gas fraction, as described in Catinella et al. (
It is also clear that GASS35981 lies significantly above the best-fit `gas fundamental plane' (dotted line) relating stellar mass surface density, color, and gas fraction, as described in Catinella et al. (
2010).
2010).
Because this galaxy was an interesting outlier and a presumed easy target, GASS35981 was selected for inclusion in our initial COLD GASS pilot program to obtain molecular gas measurements with the IRAM 30m telescope.
Because this galaxy was an interesting outlier and a presumed easy target, GASS35981 was selected for inclusion in our initial COLD GASS pilot program to obtain molecular gas measurements with the IRAM 30m telescope.
Observations of GASS35981 in the J—1-0 rotational transition of CO were made at 3mm with the IRAM 30m telescope, in three different pointings: one at the galaxy center, and one each to the north and south, one beam-width (22") away along the galaxy major axis.
Observations of GASS35981 in the J=1–0 rotational transition of CO were made at 3mm with the IRAM 30m telescope, in three different pointings: one at the galaxy center, and one each to the north and south, one beam-width $\arcsec$ ) away along the galaxy major axis.
Data were taken in June and August 2009, using the WILMA and 4MHz backends simultaneously to record the data, and the software to process them.
Data were taken in June and August 2009, using the WILMA and 4MHz backends simultaneously to record the data, and the software to process them.
Individual scans were examined, and a linear baseline subtracted from each of them.
Individual scans were examined, and a linear baseline subtracted from each of them.
After rejection of scans with unstable baselines due to, e.g., poor atmospheric conditions, the data were combined and binned to a spectral resolution of 21 km s-!.
After rejection of scans with unstable baselines due to, e.g., poor atmospheric conditions, the data were combined and binned to a spectral resolution of 21 km $^{-1}$ .
'The CO line is detected in the central pointing with S/N= 5.8, for an integrated line flux of
The CO line is detected in the central pointing with $=5.8$ , for an integrated line flux of
(AMaceroni BRuciuski 1997: Maceroni \Toutalban 2001).
(Maceroni Rucinski 1997; Maceroni Montalban 2004).
This was modelled: as cousisting of two ALS dwarfs which are aluost. but not quite in contact. with masses of WLAL. aud 0.11AZ...
This was modelled as consisting of two M3 dwarfs which are almost, but not quite in contact, with masses of $0.44~M_\odot$ and $0.41~M_\odot$.
This object is not wesent in the SuperWASP database owing to its füntuess.
This object is not present in the SuperWASP database owing to its faintness.
Iun preseutius the short-period end of the coutact dnarv period distribution based ou the Al-Sky Automated Survey (ASAS). Rucinski (2007) found only τος systems in the period range 0.200 d <P« 1225 d (and none at shorter periods).
In presenting the short-period end of the contact binary period distribution based on the All-Sky Automated Survey (ASAS), Rucinski (2007) found only three systems in the period range 0.200 d $< P <$ 0.225 d (and none at shorter periods).