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It is in this zegion where differences inthe radio structures of the two families of radio galaxies become evident and where the emission lines in the optical spectra are produced. | It is in this region where differences in the radio structures of the two families of radio galaxies become evident and where the emission lines in the optical spectra are produced. |
Moreover. the propertiesof this region mist be closely Luked with the conditious in the region of accretion. which define the source of ionization or the different core power. | Moreover, the properties of this region must be closely linked with the conditions in the region of accretion, which define the source of ionization or the different core power. |
Tn a sinall fraction of sources. the external medii at larger scales (possibly at the ealactic halo) nueght determine key propertics of the large scale radio structure (FR I/II sources. or cuviromment induced asviunetrv). | In a small fraction of sources, the external medium at larger scales (possibly at the galactic halo) might determine key properties of the large scale radio structure (FR I/II sources, or environment induced asymmetry). |
The relation between the intrinsic size aud the radio power of radio galaxies is a powerful tool for studying the evolution of radio galaxies (Shklovskii 1963.. Scheucr 1971.. Neeser et al. 1995.. | The relation between the intrinsic size and the radio power of radio galaxies is a powerful tool for studying the evolution of radio galaxies (Shklovskii \cite{shklovskii}, , Scheuer \cite{scheuer}, Neeser et al. \cite{neeser}, , |
aiser et al. 1997)). | Kaiser et al. \cite{kaiser}) ). |
We display in Fig. | We display in Fig. |
14 the power - size diagrain for a large compilation of radio galaxies taken from the B2 sample. the Peacock Wall (1981)) sample. our sample of large angular size radio galaxies (Paper D). aud a compilation of GRCs frou: Ishwara-Chaudra Saikia (1999)). Schocumakers et al. (200043) | \ref{pd} the power - size diagram for a large compilation of radio galaxies taken from the B2 sample, the Peacock Wall \cite{peacock}) ) sample, our sample of large angular size radio galaxies (Paper I), and a compilation of GRGs from Ishwara-Chandra Saikia \cite{ishwara}) ), Schoenmakers et al. \cite{arno1}) ) |
and Machalski et al. (2001)). | and Machalski et al. \cite{machalski}) ). |
Both FR I and FR IH radio ealaxies are represented in this plot. | Both FR I and FR II radio galaxies are represented in this plot. |
In order to separate the influence of the sample sclection criteria and of the evolution of radio galaxies we have plotted also a line representing the seusitivitv limits affecting our sample. as discussed in Paper I. This sensitivity limit mostly affects the selection of FR I radio galaxies. and in principle of very huge FR Is. | In order to separate the influence of the sample selection criteria and of the evolution of radio galaxies we have plotted also a line representing the sensitivity limits affecting our sample, as discussed in Paper I. This sensitivity limit mostly affects the selection of FR I radio galaxies, and in principle of very large FR IIs. |
We also plot the evolutionary tracks of FR II radio galaxies determined by Naiser et al. (1997)) | We also plot the evolutionary tracks of FR II radio galaxies determined by Kaiser et al. \cite{kaiser}) ) |
from a amodel of the cocoon of FR II radio galaxies which takes into account the energv loss processes for the relativistic electrons. | from a model of the cocoon of FR II radio galaxies which takes into account the energy loss processes for the relativistic electrons. |
The lines plotted correspond to their Fig. | The lines plotted correspond to their Fig. |
1. for three differeut jet powers. aud to the case when the euerev of the magnetic field and of the particles is in equipartition (their case 3). | 1, for three different jet powers, and to the case when the energy of the magnetic field and of the particles is in equipartition (their case 3). |
This model explains the dearth of high Iuuinosity GRGs. even if recent systematic searches of GRGs are reaching lower aud lower sensitivity limits (Paper I: Schocumalsers et al. 20002:: | This model explains the dearth of high luminosity GRGs, even if recent systematic searches of GRGs are reaching lower and lower sensitivity limits (Paper I; Schoenmakers et al. \cite{arno1}; |
Machalski et al. 2001)). | Machalski et al. \cite{machalski}) ). |
Moreover. the energy losses in the radio lobes of extended radio galaxies wight also explain the relatively Heh core prominence in these objects (Sect. 2.3)). | Moreover, the energy losses in the radio lobes of extended radio galaxies might also explain the relatively high core prominence in these objects (Sect. \ref{pcore}) ). |
Note that the apparent correlation between the radio power and the source size in our sample (filled circles in Fig. 11)) | Note that the apparent correlation between the radio power and the source size in our sample (filled circles in Fig. \ref{pd}) ) |
is well reproduced taking into account the biases iutroduced by the sample selection criteria in total flux ceusity aud angular size. thus the need of considering several samples with different sclection criteria to derive valuable conclusions from power-size relations. | is well reproduced taking into account the biases introduced by the sample selection criteria in total flux density and angular size, thus the need of considering several samples with different selection criteria to derive valuable conclusions from power-size relations. |
The lack of high power GRGs is well visible in Fig. | The lack of high power GRGs is well visible in Fig. |
11 aud coufiriued by the laree sample of radio sources: iu the linear size range <1 Mpe we have iiiv sources with radio power between 10°? and 1075 Πε. while very fev GRGs have radio power >107 W/IIz. | \ref{pd} and confirmed by the large sample of radio sources: in the linear size range $<$ 1 Mpc we have many sources with radio power between $10^{29}$ and $10^{28}$ W/Hz, while very few GRGs have radio power $>10^{27}$ W/Hz. |
It cannot be justified by selection effects. | It cannot be justified by selection effects. |
A lack of giant radio sources is expected by evolutionary models of radio sources (Ixaiscer οἳ al. 1997: | A lack of giant radio sources is expected by evolutionary models of radio sources (Kaiser et al. \cite{kaiser}; |
Dluudell et al. 19993). | Blundell et al. \cite{blundell}) ). |
We estimate that a radio power loss of a factor 10 to 100 is required from the sizepower relation. | We estimate that a radio power loss of a factor 10 to 100 is required from the size–power relation. |
This loss is in remarkably agrecuent the total radio power loss expected from the P. vs [5 correlation (see Sect. 2.3]). | This loss is in remarkably agreement the total radio power loss expected from the $P_c$ vs $P_t$ correlation (see Sect. \ref{pcore}) ). |
We sunununarize here the results preseuted iu previous sections, but from the poiut of view of giant radio galaxies, | We summarize here the results presented in previous sections, but from the point of view of giant radio galaxies. |
To improve the statistics of this type of racio galaxies. we consider when possible a compilation of 115 GRGs taken fron Ishwara-Chaudra Sailia (1990)). Schoemmakers ct al. (2000a)). | To improve the statistics of this type of radio galaxies, we consider when possible a compilation of 115 GRGs taken from Ishwara-Chandra Saikia \cite{ishwara}) ), Schoenmakers et al. \cite{arno1}) ), |
Alachalski et al. (2001)). | Machalski et al. \cite{machalski}) ), |
aud our sample (Paper I) (the same compilation shown in Fie. 11)). | and our sample (Paper I) (the same compilation shown in Fig. \ref{pd}) ). |
From Fie. | From Fig. |
11. we fiud that most CRCs have sizes below 3 Mpc. with onlv a limited umber of objects surpassing this size. | \ref{pd} we find that most GRGs have sizes below 3 Mpc, with only a limited number of objects surpassing this size. |
A cutoff in the linear size of GCRGs was already noted bv Ishlwara-Chandra Sailkia (1999)) from the analvsis of a sample of 50 (Πας, aud by Schocumalers et al. (2001)) | A cutoff in the linear size of GRGs was already noted by Ishwara-Chandra Saikia \cite{ishwara}) ) from the analysis of a sample of 50 GRGs, and by Schoenmakers et al. \cite{arno2}) ) |
from a sample of LF CRs. | from a sample of 47 GRGs. |
Our sample helps to. confirm: this cutoff between 2 aud 3 Alpe. althougaay[um we note it nüeht be a combined effect of the decrease in huniuositv with source size (Naiser et al. 1997)). | Our sample helps to confirm this cutoff between 2 and 3 Mpc, although we note it might be a combined effect of the decrease in luminosity with source size (Kaiser et al. \cite{kaiser}) ), |
the scusitivity limitatiou of the NWSS aud the bias introduca by our selection criteria (see Fig. 11)). | the sensitivity limitation of the NVSS and the bias introduced by our selection criteria (see Fig. \ref{pd}) ). |
Tn Fig. | In Fig. |
12 we plot the distribution with redshift of the GBCis in our sample. together with the same distribution for the compilation of CRCs. | \ref{lumsize} we plot the distribution with redshift of the GRGs in our sample, together with the same distribution for the compilation of GRGs. |
We find that most kuown GC have a redshift 2<0.25. | We find that most known GRGs have a redshift $z \le 0.25$. |
In our sample we know it is mmostly due to our selection criteria: we do not select sources smaller than 1 Mpc at 2>0.2. xd the αππα selectable size increases rapidly with : (see Paper I). | In our sample we know it is mostly due to our selection criteria: we do not select sources smaller than 1 Mpc at $z>0.2$, and the minimum selectable size increases rapidly with $z$ (see Paper I). |
That means that at 2>0.2 many CRCs are Όσιος missed and oulythose withsizes of about 2 \Ipe or larger cau be selected. but these are rare objects. | That means that at $z > 0.2$ many GRGs are being missed and onlythose withsizes of about 2 Mpc or larger can be selected, but these are rare objects. |
Moreover. most | Moreover, most |
of the galaxy core. fouud evidence for iutermediate-aged. aud old stellar populations. along with a much vounecr componcut that fades very quickly with radius. | of the galaxy core, found evidence for intermediate-aged, and old stellar populations, along with a much younger component that fades very quickly with radius. |
The nucleus also has a dominant 1 Cr old population 2010). | The nucleus also has a dominant 1 Gyr old population . |
. UV spectroscopy also imdicates some very voung (8 10 My) stars in the galaxw ceuter1998). | UV spectroscopy also indicates some very young $\leq$ 10 Myr) stars in the galaxy center. |
. Farther out. NGC LOL las a very huge (720 kpc) stellar disk that is known to be dominated by red giants2003). | Farther out, NGC 404 has a very large $>$ 20 kpc) stellar disk that is known to be dominated by red giants. |
. This disk also coutaius a large amount of eas in a well-defined inner disk (<5 kpc) aud a warped outer disk£).. although the density of gas is low (2 41029 AL. ? throughout) aud ouly a faint rius of star formation is detected in the far violet 1990). | This disk also contains a large amount of gas in a well-defined inner disk $<$ 5 kpc) and a warped outer disk, although the density of gas is low $<$ $\times$ $^{20}$ $_{\odot}$ $^{-2}$ throughout) and only a faint ring of star formation is detected in the far ultra-violet . |
. The origin of the gas disk of NGC LOL is somewhat nivsterious. | The origin of the gas disk of NGC 404 is somewhat mysterious. |
sugeest that the disk is likely due to a recent(0.51 Car ago) merecr event with a ο galaxy. motivated by their detection of a possible warp in the disk. | suggest that the disk is likely due to a recent (0.5–1 Gyr ago) merger event with a dIrr galaxy, motivated by their detection of a possible warp in the disk. |
However. NGC LOL is very isolated for a ealaxy of its massdetails). | However, NGC 404 is very isolated for a galaxy of its mass. |
. With no other galaxies of auv kiud detected within 1 \Ipe of NGC LOL stich a merger is not obviously consistent with the ealaxw’s cuviromment. | With no other galaxies of any kind detected within 1 Mpc of NGC 404, such a merger is not obviously consistent with the galaxy's environment. |
The eas dis- could therefore simply be due to late time gas intall frou filameuts rather than accretion of bound. couldobjects. | The gas disk could therefore simply be due to late time gas infall from filaments rather than accretion of bound objects. |
Firthermore. a gas disk of such low density have a very long lifetime eiven the low star formation rate. sugecsting that it could have been accreted far earlier. | Furthermore, a gas disk of such low density could have a very long lifetime given the low star formation rate, suggesting that it could have been accreted far earlier. |
By comparing the stellar populations aud eas densities. we can constrain the origiu of this uuusual disk. | By comparing the stellar populations and gas densities, we can constrain the origin of this unusual disk. |
We examine the stellar populations of NGC 101 iu detail using deep observations from UST. | We examine the stellar populations of NGC 404 in detail using deep observations from HST. |
We determine the star formation history (SEII) of several portions of the galaxy by fittine the distribution of stars iu color-magnitude diagrams (CMDs) with model distributious determined from stellar evolution isochrones2002). | We determine the star formation history (SFH) of several portions of the galaxy by fitting the distribution of stars in color-magnitude diagrams (CMDs) with model distributions determined from stellar evolution isochrones. |
. Section 2 describes our data set aud analysis procedures. | Section 2 describes our data set and analysis procedures. |
Section 3 preseuts the results of our measurements. | Section 3 presents the results of our measurements. |
Section [| interprets the ueasurements in the coutext of the eas properties and our understanding of SO ealaxies. aud Section 5 sunnuurlizes our conclusions. | Section 4 interprets the measurements in the context of the gas properties and our understanding of S0 galaxies, and Section 5 summarizes our conclusions. |
We assume a distance of 3.05 Mpc for conversions of aneular measurements to physical distances aud adopt au inclination angle /—117 for surface deusitv measurements. | We assume a distance of 3.05 Mpc for conversions of angular measurements to physical distances and adopt an inclination angle $i$ $^{\circ}$ for surface density measurements. |
We adopt a five-vear WALAP cosinoloey for all conversous between time aud redshift. | We adopt a five-year WMAP cosmology for all conversions between time and redshift. |
From 2007-Aug-08 to 2007-Sep-20. we observed a field in the NGC 101 disk located at R.ÁÀ. (2000) = 17.32325 (01:09:17.6). decl. ( | From 2007-Aug-08 to 2007-Sep-20, we observed a field in the NGC 404 disk located at R.A. (2000) = 17.32325 (01:09:17.6), decl. ( |
2000)= 35.71856 (|35:11:55) mna rotation angele DÀ.V3250.0 degrees. | 2000) = 35.74856 (+35:44:55) with a rotation angle V3=50.0 degrees. |
Frou, 2009-Febn to 2009-Feb-20. we Vperformed shallower observations 2 fields located at R.À. (2000) =17.36697 (01:09:28.1). decl. ( | From 2009-Feb-16 to 2009-Feb-20, we performed shallower observations for 2 fields located at R.A. (2000) =17.36697 (01:09:28.1), decl. ( |
2000) = 35.761172m (1235:15:10). with a rotation angle PA_VV3= and R.A. (2000) 17.333368 (01:09:20.0). decl. ( | 2000) = 35.76117 (+35:45:40) with a rotation angle V3=230.0 and R.A. (2000) =17.333368 (01:09:20.0), decl. ( |
2000) = 35.270205 (135:12:07) with a rotation angle V32230.0. | 2000) = 35.70205 (+35:42:07) with a rotation angle V3=230.0. |
Figure d shows outlines of the fields” locations. | Figure \ref{field_loc}
shows outlines of the fields' locations. |
Our field locations were chosen to maximize the number of disk stars aud avoid crowding. | Our field locations were chosen to maximize the number of disk stars and avoid crowding. |
Iu the deep field. we obtained 15 full-orbit exposures with the WEPC2 through the E606. (wide V) filter. and 29 full-orbit exposures through the Faliw (7 equivalent) filter. | In the deep field, we obtained 15 full-orbit exposures with the WFPC2 through the F606W (wide $V$ ) filter, and 29 full-orbit exposures through the F814W $I$ equivalent) filter. |
These data totaled 39000 s and 75100 s of exposure time in EFGOGW and Feliw. respectively. | These data totaled 39000 s and 75400 s of exposure time in F606W and F814W, respectively. |
Iu the other two fields. we obtained 2 orbits through F606W. totaling 1800 s. and Lb orbits through FSLIW. totaling 9600 s. All images were calibrated in the IST pipeline with CALWDP2 using OPUS version 2006_66a for the 2007 data and 2008_55¢ for the 2009 data. | In the other two fields, we obtained 2 orbits through F606W, totaling 4800 s, and 4 orbits through F814W, totaling 9600 s. All images were calibrated in the HST pipeline with CALWP2 using OPUS version 6a for the 2007 data and 5c for the 2009 data. |
To expand our radial coverage. we also reduced 2 fields in the outer disk. previously studied aud named S2 aud S3details). | To expand our radial coverage, we also reduced 2 fields in the outer disk, previously studied and named S2 and S3. |
.. These fields lie ~s’ (7 kpe) SW of the uucleus aud were taken ax part of 60-5369. | These fields lie $\sim$ $'$ (7 kpc) SW of the nucleus and were taken as part of GO-5369. |
They coutain 1200 sec of exposure in F606W cach. | They contain 1200 sec of exposure in F606W each. |
Field $3 contains L200 sec of exposure πι Fall. while S2 contains oulv 2100 sec in FaliW. The data reduction and photometry for the ANGST survey are fully described i(2009). | Field S3 contains 4200 sec of exposure in F814W, while S2 contains only 2100 sec in F814W. The data reduction and photometry for the ANGST survey are fully described in. |
. For couvenience we provide a brief summary of the techniques here. | For convenience we provide a brief summary of the techniques here. |
photometry was mieasured simultancously for all of the paclaceIwobjects in the uncoubined inages using the software HSTPIIOT2000).. | The photometry was measured simultaneously for all of the objects in the uncombined images using the software package HSTPHOT. |
This package is optimized for mieasuriug photometry of stars ou WEPC2 nuages using the wellcharacterized and stable point spread function (PSF) calculated with The software fits the PSF to all of the stars in cach individual fune to find PSF amaenitudes. | This package is optimized for measuring photometry of stars on WFPC2 images using the well-characterized and stable point spread function (PSF) calculated with The software fits the PSF to all of the stars in each individual frame to find PSF magnitudes. |
It then determines and applies the aperture correction for cach image usiug the most isolated stars. corrects for the charge transfer cficiency of the WEPC2detector’... combines the results frou the individual exposures. aud converts the measured count rates to the VECΔίας svstenm. | It then determines and applies the aperture correction for each image using the most isolated stars, corrects for the charge transfer efficiency of the WFPC2, combines the results from the individual exposures, and converts the measured count rates to the VEGAmag system. |
The HSTPIIOT output was then filtered to only allow objects classified as stars with signal-to-noise (total counts from the star to total noise) 2 bin both filters. | The HSTPHOT output was then filtered to only allow objects classified as stars with signal-to-noise (total counts from the star to total noise) $>$ 4 in both filters. |
The list was firther eulled using sharpucss (F606raryFSLIWνα< 027) and crowding (F606uuu|FSl1W.4< 0.7). | The list was further culled using sharpness $|F606W_{sharp} + F814W_{sharp}| < 0.27$ ) and crowding $F606W_{crowd} + F814W_{crowd} < 0.7$ ). |
The sharpness eut was chosen based ou the distribution of values in the original catalog. | The sharpness cut was chosen based on the distribution of values in the original catalog. |
The crowding paraucter gives the difference between the maguitude of a star measured before aud after subtracting the neighboring stars in the nuage. | The crowding parameter gives the difference between the magnitude of a star measured before and after subtracting the neighboring stars in the image. |
When this valueis lavec. it 3ugeests that the stars photometry was siguificautly affectedby crowding. aud we therefore exclude it from our catalog. | When this value is large, it suggests that the star's photometry was significantly affected by crowding, and we therefore exclude it from our catalog. |
Quality cuts based on the \ values were also considered. but they were rejected when a correlation was found between y aud the local background. | Quality cuts based on the $\chi$ values were also considered, but they were rejected when a correlation was found between $\chi$ and the local background. |
Our final star catalogs contained 10793. 22:332. and 33365 stars for the deep. NE. and SW Ποια». respectively. | Our final star catalogs contained 40793, 22332, and 33365 stars for the deep, NE, and SW fields, respectively. |
The archival outer disk fields 82 aud $3 | The archival outer disk fields S2 and S3 |
fraction the projectile, the angle-adjusted reduced mass 1s The disruption criteria and critical impact velocity for a specific mass ratio and impact angle are calculated by adjustments to the head-on equal-mass principal disruption curve. | fraction the projectile, the angle-adjusted reduced mass is The disruption criteria and critical impact velocity for a specific mass ratio and impact angle are calculated by adjustments to the head-on equal-mass principal disruption curve. |
To achieve disruption with a smaller projectile, the impact velocity rises to and, from Equation 2, the catastrophic disruption criteria increases by Next, accounting for the projectile interacting mass fraction for oblique impacts, where the prime notation indicates the value for an oblique impact. | To achieve disruption with a smaller projectile, the impact velocity rises to and, from Equation \ref{eqn:qstarred}, the catastrophic disruption criteria increases by Next, accounting for the projectile interacting mass fraction for oblique impacts, where the prime notation indicates the value for an oblique impact. |
Finally, from Equation 1,, the critical impact velocity adjusted for both the mass ratio and impact angle is The for and V'* are curvesas a function of Βσι, impact equationsangle Q5;(within μα), mass ratio (within jj), and two material parameters, c* and jg. | Finally, from Equation \ref{eqn:qr}, the critical impact velocity adjusted for both the mass ratio and impact angle is The equations for $Q^{\prime *}_{RD}$ and $V^{\prime *}$ are curvesas a function of $R_{C1}$, impact angle (within $\mu_{\alpha}$ ), mass ratio (within $\mu$ ), and two material parameters, $c^*$ and $\bar
\mu$. |
The specific collision outcome regime is determined using the mutual escape velocity, disruption criteria, and impact parameter. | The specific collision outcome regime is determined using the mutual escape velocity, disruption criteria, and impact parameter. |
When the impact velocity is below the mutual escape velocity of the interacting mass (Mt,=+ Με), the two bodies are assumed to merge completelyaM, (perfect merging). | When the impact velocity is below the mutual escape velocity of the interacting mass $M^{\prime}_{\rm tot}=\alpha M_{\rm p}+M_{\rm t}$ ), the two bodies are assumed to merge completely (perfect merging). |
Above VJ,esc=V/2GM,+ the collisions outcomes are divided into two groups:/(Jt grazingRt), and non-grazing. | Above $V^{\prime}_{\rm esc}=\sqrt{2 G M^{\prime}_{\rm tot} / (R_{\rm
p}+R_{\rm t})}$ , the collisions outcomes are divided into two groups: grazing and non-grazing. |
The transition between non-grazing and grazing outcomes is demarcated by a critical impact parameter, be. | The transition between non-grazing and grazing outcomes is demarcated by a critical impact parameter, $b_{\rm crit}$. |
Following(2010), the center of the projectile is tangent to the target at the critical impact parameter, Non-grazing (b« collisions transition from perfect merging to the disruption berit) regime with increasing impact velocity. | Following, the center of the projectile is tangent to the target at the critical impact parameter, Non-grazing $b<b_{\rm crit}$ ) collisions transition from perfect merging to the disruption regime with increasing impact velocity. |
In the disruption regime, the outcome maybe partial accretion or partial erosion depending on the mass of the largest remnant, My. | In the disruption regime, the outcome maybe partial accretion or partial erosion depending on the mass of the largest remnant, $M_{\rm
lr}$. |
The largest remnant mass is calculated using the impact energy and the catastrophic disruption criteria. | The largest remnant mass is calculated using the impact energy and the catastrophic disruption criteria. |
For collisions with 0«Qr/Qkp1.8, My is proportional to impact energy: Because single line fit a wide variety of simulation results, we anamed this linear relationship the “universal law for the mass of the largest remnant" 2009). | For collisions with $0<Q_R/Q^{\prime *}_{RD}<1.8$, $M_{\rm lr}$ is proportional to impact energy: Because a single line fit a wide variety of simulation results, we named this linear relationship the “universal law for the mass of the largest remnant” . |
. In the case of super-catastrophic(Stewart Qn/Q'5p> 1.8, the largest remnant follows a power disruption,law, | In the case of super-catastrophic disruption, $Q_R/Q^{\prime *}_{RD} \ge 1.8$ , the largest remnant follows a power law, |
davs and 1.33 days. again. wilh no known instrumental effect (hat could be responsible. | days and 1.33 days, again, with no known instrumental effect that could be responsible. |
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