ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

If accretion has been stochastic. then galaxies that have recently acquired gas from the external environment might be expected to have unusually high HI mass fractions.

If accretion has been stochastic, then galaxies that have recently acquired gas from the external environment might be expected to have unusually high HI mass fractions.

The neutral gas content of a galaxy has long been known to correlate with its other physical properties.

HI mass fractions increase smoothly along the Hubble sequence from the early-type (S0) to the late-type (Im) end (Roberts&Haynes1994).

HI mass fractions increase smoothly along the Hubble sequence from the early-type (S0) to the late-type (Im) end \citep{RH94}.

. There are also correlations between HI mass fraction and galaxy properties such as stellar mass (AZ,.). stellar mass surface density (μ.) and broadband colour.

There are also correlations between HI mass fraction and galaxy properties such as stellar mass $M_*$ ), stellar mass surface density $\mu_*$ ) and broadband colour.

These correlations form the basis for the so-called “photometric gas fraction” technique for predicting the HI gas content galaxies (Kannappan2004:Zhangetal.2009).

These correlations form the basis for the so-called “photometric gas fraction” technique for predicting the HI gas content galaxies \citep{Kannappan04, Zhang09}.

. Rather little attention has been given to thescaffer in gas fraction at a given value of Μ.. µε or colour.

Rather little attention has been given to the in gas fraction at a given value of $M_*$, $\mu_*$ or colour.

Recently. Zhang et al (2009) found that more gas-rich galaxies were systematically deficient in metals at fixed stellar mass and they interpreted this as evidence for recent cosmological infall of gas in these systems.

Recently, Zhang et al (2009) found that more gas-rich galaxies were systematically deficient in metals at fixed stellar mass and they interpreted this as evidence for recent cosmological infall of gas in these systems.

However. up to now there has been ho systematic study of how galaxy properties vary as a function of HI over- or under-abundance.

However, up to now there has been no systematic study of how galaxy properties vary as a function of HI over- or under-abundance.

One problem that has hampered progress on this front has been the lack of suitable data sets.

Blind HI surveys. such as the HI Parkes All Sky Survey (HIPASS. Barnes et al 2001) or the Arecibo Legacy Fast ALFA survey (ALFALFA. Giovanelli et al 2005) have produced large. unbiased samples of galaxies selected by HI mass.

Blind HI surveys, such as the HI Parkes All Sky Survey (HIPASS, Barnes et al 2001) or the Arecibo Legacy Fast ALFA survey (ALFALFA, Giovanelli et al 2005) have produced large, unbiased samples of galaxies selected by HI mass.

However. these surveys are still relatively shallow. so the fraction of HI-poor galaxies is small.

However, these surveys are still relatively shallow, so the fraction of HI-poor galaxies is small.

Recently. Catinella et al (2010. hereafter C10) explored how HI gas fraction scales as a function of galaxy stellar mass. galaxy structural parameters and NUV-r colour. using data from the GALEX Arecibo SDSS Survey (GASS).

Recently, Catinella et al (2010, hereafter C10) explored how HI gas fraction scales as a function of galaxy stellar mass, galaxy structural parameters and NUV-r colour, using data from the GALEX Arecibo SDSS Survey (GASS).

The survey is an ongoing large programme at the Arecibo radio telescope that is gathering high quality HI-line spectra for an unbiased sample of ~1000 galaxies with stellar masses greater than 107A. and. redshifts in the range 0.025<z< 0.05. selected from the Sloan Digital Sky Survey (SDSS. Yorketal. (2000))) spectroscopic and Galaxy Evolution Explorer (GALEX. Martinetal. (2005))) imaging surveys.

The survey is an ongoing large programme at the Arecibo radio telescope that is gathering high quality HI-line spectra for an unbiased sample of $\sim 1000$ galaxies with stellar masses greater than $10^{10} M_{\odot}$ and redshifts in the range $<z<$ 0.05, selected from the Sloan Digital Sky Survey (SDSS, \citet{York00}) ) spectroscopic and Galaxy Evolution Explorer (GALEX, \citet{Martin05}) ) imaging surveys.

The galaxies are observed until detected or until a low gas mass fraction limit €1.5-5 ) is reached.

The galaxies are observed until detected or until a low gas mass fraction limit (1.5-5 ) is reached.

CIO quantify in detail how the atomic gas fraction of the galaxies in their sample decreases as a function of stellar mass. stellar mass surface density. galaxy bulge-to-dise ratio (as measured by the concentration index of the r-band light). and global NUV-r colour.

C10 quantify in detail how the atomic gas fraction of the galaxies in their sample decreases as a function of stellar mass, stellar mass surface density, galaxy bulge-to-disc ratio (as measured by the concentration index of the r-band light), and global NUV-r colour.

The fraction. of galaxies with significant (more than a few percent) HI decreases sharply above a characteristic stellar surface mass density of 107? M. 7.

The fraction of galaxies with significant (more than a few percent) HI decreases sharply above a characteristic stellar surface mass density of $10^{8.5}$ $_{\odot}$ $^{-2}$.

The fraction of gas-rich galaxies decreases much more smoothly with stellar mass.

In this paper. we extend the study of CIO to study how theresolved UV/optical photometric properties of galaxies depend on atomic gas fraction.

In this paper, we extend the study of C10 to study how the UV/optical photometric properties of galaxies depend on atomic gas fraction.

We have developed a photometric pipeline that ransforms the SDSS and GALEX images to the same geometry and etfective resolution. so that our photometric measurements race the same part of each galaxy at different wavelengths (Wang et al 2009).

We have developed a photometric pipeline that transforms the SDSS and GALEX images to the same geometry and effective resolution, so that our photometric measurements trace the same part of each galaxy at different wavelengths (Wang et al 2009).

We use these images to study UV-optical colours in jeir inner and outer regions of the galaxies in our sample and to Paltudy how colour cand hence specific star formation rate) "pend on HI content.

We use these images to study UV-optical colours in their inner and outer regions of the galaxies in our sample and to study how colour (and hence specific star formation rate) depend on HI content.

In order to extend the range in MCHD/AM jit we are able to probe. we combine the sample analyzed in CIO with galaxies from the GASS "parent sample" (described in Section 2)) which have HI detections in the ALFALFA survey.

In order to extend the range in $M_*$ that we are able to probe, we combine the sample analyzed in C10 with galaxies from the GASS “parent sample” (described in Section \ref{sec:data}) ) which have HI detections in the ALFALFA survey.

This considerably improves the statistics for galaxies with large values of MCHD/M,. which are relatively rare.

This considerably improves the statistics for galaxies with large values of $M_*$, which are relatively rare.

In order to isolate the effect of the presence or absence of gas on the photometric properties of galaxies. we compare our results with three control samples. whieh have similar optical properties to the galaxies in our sample with HI mass measurements. but are selected without regard to HI content.

In order to isolate the effect of the presence or absence of gas on the photometric properties of galaxies, we compare our results with three control samples, which have similar optical properties to the galaxies in our sample with HI mass measurements, but are selected without regard to HI content.

These control samples are described in detail in Section 2.3..

These control samples are described in detail in Section \ref{subsec:controlsample}.

Sections 3.1 and 3.3 present the methods used to measure our photometric parameters. and in Section + we discuss the trends in galaxy size. NUV—r and specific star formation rate gradients. and higher order morphological parameters such as asymmetry and smoothness as a function of HI mass fraction.

Sections \ref{subsec:phot_tech} and \ref{subsec:SF_tech} present the methods used to measure our photometric parameters, and in Section 4 we discuss the trends in galaxy size, $NUV-r$ and specific star formation rate gradients, and higher order morphological parameters such as asymmetry and smoothness as a function of HI mass fraction.

Finally Section 5. summarizes our results.

Finally Section \ref{sec:summary} summarizes our results.

Throughout this paper. we assume a cosmology with 5270 km s! !. Q,,=0.3. and Q,=0.7 (Tegmarketal.2004).

Throughout this paper, we assume a cosmology with $H_0$ =70 km $^{-1}$ $^{-1}$ , $\Omega_m$ =0.3, and $\Omega_\Lambda$ =0.7 \citep{Tegmark04}.

.. A Chabrier initial mass function is assumed in the stellar population synthesis analysis (Chabrier 2003).

A Chabrier initial mass function is assumed in the stellar population synthesis analysis \citep{Chabrier03}.

We select galaxies with stellar masses M,>10!A7; in the redshift range 0.025<z«0.05 from the sixth data release (DR6) of the SDSS Survey. which lie within the maximal ALFALFA footprint.

We select galaxies with stellar masses $M_*>10^{10}M_{\odot}$ in the redshift range $0.025<z<0.05$ from the sixth data release (DR6) of the SDSS Survey, which lie within the maximal ALFALFA footprint.

We match this catalogue with the fourth data release (GR4) of the GALEX survey and this yields a sample of 10468 galaxies.

We further limit the sample to face-on galaxies with b/a>0.4 tor an inclination of less than 66.4 deg) where 5 and à are the minor and major axes of an ellipsoidal tit (from SDSS r band) to each galaxy.

We further limit the sample to face-on galaxies with $b/a>0.4$ (or an inclination of less than 66.4 deg) where $b$ and $a$ are the minor and major axes of an ellipsoidal fit (from SDSS $r$ band) to each galaxy.

This yields a final sample of 8429 galaxies. which we refer to as thesample from now on.

This yields a final sample of 8429 galaxies, which we refer to as the from now on.

We restrict the sample to face-on systems. so that our estimates of star formation rate. which are derived from the UV/optical photometry. will be less attected by dust attenuation effects (Section 3.3)).

We restrict the sample to face-on systems, so that our estimates of star formation rate, which are derived from the UV/optical photometry, will be less affected by dust attenuation effects (Section \ref{subsec:SF_tech}) ).

This sample consists of galaxies in the parent sample for which we have catalogued HI detections from the 40% ALFALFA survey (a 40. which will be discussed in Martinetal.(2010) and Haynesetal. 01000). as well as galaxies that are included in the first data release (DRI) of the GASS survey (see CIO for a detailed description).Our HI sample consists of 458 galaxies with amedian HI mass fraction of ~ 30%.

This sample consists of galaxies in the parent sample for which we have catalogued HI detections from the $40\%$ ALFALFA survey $\alpha.40$ , which will be discussed in \citet{Martin10} and \citet{Haynes10}) ), as well as galaxies that are included in the first data release (DR1) of the GASS survey (see C10 for a detailed description).Our HI sample consists of 458 galaxies with amedian HI mass fraction $_*$ of $\sim$ $\%$ .

Figure |. shows the distribution of for these galaxies.

Figure \ref{fig:HI_frac} shows the distribution of $_*$ for these galaxies.

The left-top panel of

model.

This arises because the old bright spot model could not describe the complex bright spot profile present and an innacurate value of the mass ratio is found as a result.

Our new bright spot model is much better in this respect, and is able to take into account a wider variety of geometric effects and orientations.

Given that our white dwarf radius is consistent with that of Feline et al. (

2004b), this seems the most likely cause of such a large change.

It is worth noting that our new donor masses for both DV UMa and XZ Eri, are both consistent with the masses obtained by Feline et al. (

2004b) using the derivative method, which, unlike our parameterised model, does not make any attempt to recreate the bright spot eclipse profile (e.g. Wood et al.

1986; Horne et al.

1994; Feline et al.

2004a; Feline et al.

2004b).

Our new fits to SDSS 1502 decrease the donor mass by 2.90 (AM,= 0.012Μᾳ9) from that of ?..

\nocite{wood1986, horne1994, feline2004a, feline2004b} Our new fits to SDSS 1502 decrease the donor mass by $2.9\sigma$ $\Delta$$M_{r}=0.012M_{\odot}$ ) from that of \citet{littlefair2008}.

Our mass ratio and inclination are consistent with those of ?,, however our white dwarf radius, Πω, has increased by 13 percent (3.40).

Our mass ratio and inclination are consistent with those of \citet{littlefair2008}, however our white dwarf radius, $R_{w}$, has increased by $13$ percent $3.4\sigma$ ).

We believe the primary reason for this change was that the original fit was heavily binned, and thus more susceptible to the bug outlined in section 3.3..

We believe the primary reason for this change was that the original fit was heavily binned, and thus more susceptible to the bug outlined in section \ref{sec:pme}.

The most important change of all of our re-modelled systems is for that of SDSS 1501.

Whilst our donor mass has only increased by 1.90 from that of ?,, we note that our uncertainties are large (cM,= 0.010Μ9) and the mass difference is large enough to take this system from being a post-period-bounce system, to a pre-period-bounce system.

Whilst our donor mass has only increased by $1.9\sigma$ from that of \citet{littlefair2008}, we note that our uncertainties are large $\sigma$$M_{r}=0.010M_{\odot}$ ) and the mass difference is large enough to take this system from being a post-period-bounce system, to a pre-period-bounce system.

Lt is interesting to note that on the opposite side of the nucleus (region E) neither the Ho. nor the kinematic maps show anv remarkable features.

It is interesting to note that on the opposite side of the nucleus (region E) neither the $\alpha$ nor the kinematic maps show any remarkable features.

The ionization state. however. is similar to region D. “Phe Ha dispersion map doces show a large perturbation in region E. Perhaps this is the signature of the counter jet although it could also be due to interaction with the companion galaxy.

The ionization state, however, is similar to region D. The $\alpha$ dispersion map does show a large perturbation in region F. Perhaps this is the signature of the counter jet although it could also be due to interaction with the companion galaxy.

The total OLLI] luminosity that we derive. from the lux distribution shown in panel (g) isM 4.2.«107I erg s (ignoring the contribution from region C. which is due to star formation).

The total [OIII] luminosity that we derive from the flux distribution shown in panel (g) is $4.2 \times 10^{40}$ erg $^{-1}$ (ignoring the contribution from region C, which is due to star formation).

Thisis equivalent to the Ho. luminosity of he ENLI. Lg,=4. 107"«erg 7.

This is equivalent to the $\alpha$ luminosity of the ENLR, $L_{\rm H\alpha} = 4\times10^{40}$ erg $^{-1}$.

Note that this will include the contribution from the eas disk. which cannot »e properly separated from the highly ionized extra-planar eas.

Note that this will include the contribution from the gas disk, which cannot be properly separated from the highly ionized extra-planar gas.

In comparison. the mechanical power of the jet. itself can be estimated. using an empirical conversion from racio uminosity. (equation (1) of Best et al.

In comparison, the mechanical power of the jet itself can be estimated using an empirical conversion from radio luminosity (equation (1) of Best et al.

2007).

The observed racio Lux (table 1) ) implies' Li,=(4332).«10712 erg roughly an order of magnitude more than the energy which is reracliated as emission lines.

The observed radio flux (table \ref{t:properties}) ) implies $L_{\rm mech} = (4 \pm 2)\times10^{42}$ erg $^{-1}$, roughly an order of magnitude more than the energy which is reradiated as emission lines.

Phe X-ray luminosity (table 1)) is larger than both £i, and the emission line edies (Hla and ΟΠ) as expected(e.g. Heckman et al.

The X-ray luminosity (table \ref{t:properties}) ) is larger than both $L_{\rm mech}$ and the emission line luminosities $\alpha$ and [OIII]) as expected (e.g. Heckman et al.

e2004).

2004).

If we assume that most of the mechanical energv of the jet is converted. to kinetic energy in the extra-planar eas then we can compute an upper limit on the mass of ionized hydrogen in the IENLIt.

If we assume that most of the mechanical energy of the jet is converted to kinetic energy in the extra-planar gas then we can compute an upper limit on the mass of ionized hydrogen in the ENLR.

Over the lifetime of availablethejet. I0? vr (e.g. Sanders 1984). the upper limit on the encrey in the EENLIU is: s4107 cre +10° ves1.3.107" erg.

Over the lifetime of the jet, $\lesssim 10^6$ yr (e.g. Sanders 1984), the upper limit on the available energy in the ENLR is: $\lesssim 4 \times 10^{42}$ erg $^{-1} \times \lesssim 10^6$ yr $ \simeq 1.3 \times 10^{56}$ erg.

The jet lifetime is also consistent with its small size of =5 kpe. assuming a canonical jet velocity of c.

The jet lifetime is also consistent with its small size of $\lesssim 5$ kpc, assuming a canonical jet velocity of $\gtrsim 0.1c$.

With the typical gas velocities observed in our data of Vins=MAS1oF~300 km this kinetic energy would. correspond to an upper limit on the mass in ionized ivdrogen of ~1.4107M..

With the typical gas velocities observed in our data of $V^2_{\rm RMS} = V^2_{\rm rot} + \sigma^2 \sim 300$ km $^{-1}$, this kinetic energy would correspond to an upper limit on the mass in ionized hydrogen of $\sim1.4 \times10^{8} M_\odot$.

In our interpretation. we associate the structure observed in region LD with jet driven mass outllow.

In our interpretation, we associate the structure observed in region D with jet driven mass outflow.

The raction of the total kinetic energy. needed. to. power this outflow is simply the fraction of mass in region D multiplied w the ratio of the bulk velocity (150 km 3) to Vp~ km s+ squared.

The fraction of the total kinetic energy needed to power this outflow is simply the fraction of mass in region D multiplied by the ratio of the bulk velocity $\sim 150$ km $^{-1}$ ) to $V_{\rm RMS} \sim 300$ km $^{-1}$ squared.

Under the assumption of constant eas densitv. the fraction of mass can be estimated. as the raction of OLLI] luminosity in region D. ~0.07.

Under the assumption of constant gas density, the fraction of mass can be estimated as the fraction of [OIII] luminosity in region D, $\sim 0.07$.

Therefore he fraction of ENLI kinetic energy in this bulk outllow is M13.10.0007«(150/300)?~2.3.1073 erg.

Therefore the fraction of ENLR kinetic energy in this bulk outflow is $\sim 1.3 \times 10^{56} \times 0.07 \times (150/300)^2 \simeq 2.3 \times 10^{54}$ erg.

Phat is. only about 2 percent of the mechanical energy is requirecl to »ower the outflow.

That is, only about 2 percent of the mechanical energy is required to power the outflow.

The derived energies are order of magnitude estimates only but are all internally consistent.

The low mass loading ancl velocity associated with the outflow makes it unlikely that this process has a profound. impact on the cold. gas content of this galaxy.

The low mass loading and velocity associated with the outflow makes it unlikely that this process has a profound impact on the cold gas content of this galaxy.

However. the implied: mechanical energv of the jet is 50r times greater on this basis only a small fraction of the jet energy is used to power the outflow.

However, the implied mechanical energy of the jet is 50 times greater — on this basis only a small fraction of the jet energy is used to power the outflow.

A much larger fraction is available to heat the gas which we observe as the highly ionized. large ENNLI in this galaxy.

A much larger fraction is available to heat the gas which we observe as the highly ionized, large ENLR in this galaxy.

Ht djs notable that the jet energy. is comparable to the cooling luminosity of a 1 keV ealaxy(LOMOAL.) "m"halo.

It is notable that the jet energy is comparable to the cooling luminosity of a 1 keV $\sim 10^{13.5} {\rm M_{\odot}}$ ) halo.

This is an important. point in this [rom the AGN seems to have little direct. effect. on the galaxy: any inlluence it can have occurs through the heating of eas in the galaxy’s halo.

This is an important point — in this galaxy feedback from the AGN seems to have little direct effect on the galaxy: any influence it can have occurs through the heating of gas in the galaxy's halo.

This scenario is very much wih current galaxy formation models(ee..

This scenario is very much consistent with current galaxy formation models (eg.,

. Croton et consistent22006. Bower et 22006nP IS).

Croton et 2006, Bower et 2006, 2008).

Compared with powerful. QSOs (eg. Nesvadha et

Compared with powerful QSOs (e.g. Nesvadba et al.

=Yr) and radio galaxies (eg..

2007) and radio galaxies (eg.,

. Jost οἱ "200ACIN) the "backjet energy is

Best et 2007) the jet energy is small.

Nevertheless. it is the impact [oc in 1077.-- 1077AL. haloes that is responsible for shaping the galaxy function.

Nevertheless, it is the impact of AGN feedback in $10^{12}$ $10^{13} {\rm M_{\odot}}$ haloes that is responsible for shaping the galaxy luminosity function.

The jet of this low mass AGN imparts more of its kinetic energy into the cold gas by means of kinetic heating than by directed outflow.

LEU observations of galaxies hosting raclio AGN. such as presented in this letter. provide key insight into the coupling between the jet and the gas.

IFU observations of galaxies hosting radio AGN, such as presented in this letter, provide key insight into the coupling between the jet and the gas.

We thank the referee. Montserrat. Villar-Martin. for the constructive comments and suggestions.

We thank the referee, Montserrat Villar-Martin, for the constructive comments and suggestions.

We also like to thank Chris Done. Isabelle Gavignaud. Martin Krause and Alare Schartmann for helpful discussions.

We also like to thank Chris Done, Isabelle Gavignaud, Martin Krause and Marc Schartmann for helpful discussions.

We have assumed that a PPA source would appear as aminous AGN. with the gaseous ΠΙΟ perhaps being supplied » the preceding merger of the binarv's host galaxies.

We have assumed that a PTA source would appear as luminous AGN, with the gaseous fuel perhaps being supplied by the preceding merger of the binary's host galaxies.

The degree to which galaxy mergers dictate ACN activity remains an open question. and the link between SALBLI παν and GN activity is even less certain.

The degree to which galaxy mergers dictate AGN activity remains an open question, and the link between SMBH binarity and AGN activity is even less certain.

Lt is possible hat many SMDLILI binaries that are incliviclually detected by *PAs will have no EAL counterpart at all.

It is possible that many SMBH binaries that are individually detected by PTAs will have no EM counterpart at all.

Our results were calculated using a simple semi-analvtic accretion disc model. a central assumption of which is that he binarys tidal torques are able to open a central cavity in the disc.

Our results were calculated using a simple semi-analytic accretion disc model, a central assumption of which is that the binary's tidal torques are able to open a central cavity in the disc.

In the raciation-dominatec regions of interest. strong horizontal advective Duxes or vertical thickening of he disc may act to close such à cavity ancl wipe out the eatures we predict.

In the radiation-dominated regions of interest, strong horizontal advective fluxes or vertical thickening of the disc may act to close such a cavity and wipe out the features we predict.

Phe features would also not be present if the disc and binary orbits do not lie on the same plane. as in the binary model sof the variable BL Lac object O.J 287 (Lehto&Valtonen1996).

The features would also not be present if the disc and binary's orbits do not lie on the same plane, as in the binary model of the variable BL Lac object OJ 287 \citep{LV96}.

. Absorption and reprocessing by he binary’s host galaxy may also act to mask or mimic the intrinsic thermal AGN emission we have moclelect.

Absorption and reprocessing by the binary's host galaxy may also act to mask or mimic the intrinsic thermal AGN emission we have modeled.

As we were completing this work. we became aware of a concurrent independent. study by Sesanaetal.(2011). addressing similar questions.

As we were completing this work, we became aware of a concurrent independent study by \cite{Sesana+11}, addressing similar questions.

“PT acknowledges fruitful discussions. with Alberto Sesana. Massimo. Dotti. and Constanze Roelelig.

TT acknowledges fruitful discussions with Alberto Sesana, Massimo Dotti, and Constanze Röddig.

Phe authors thank Jules Halpern aud Jeremy Goodman. for insightful. conversations. and are erateful to the anonymous referee. for suggestions that improved the clarity of the manuscript.

The authors thank Jules Halpern and Jeremy Goodman for insightful conversations, and are grateful to the anonymous referee for suggestions that improved the clarity of the manuscript.

This work was supported by NASA NEEP. grants NNNOSAII35G. (to IXM) and NNITLOZDAOOIN (to ZHI) and by the Polánnyi Program of the Hungarian National Ollice for Rescarch and Technology. (ΝΑΤ to ZL).

This work was supported by NASA ATFP grants NNXO8AH35G (to KM) and NNH10ZDA001N (to ZH) and by the Polánnyi Program of the Hungarian National Office for Research and Technology (NKTH; to ZH).

This research was supported in part by the Perimeter Institute for Theoretical Physics.

quasars without detected AALs, matched in redshift and i-band magnitude to the AAL quasars.

quasars without detected AALs, matched in redshift and $i$ -band magnitude to the AAL quasars.