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

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
llowever. in svmbiotic novae like V407 Cvgni. there are some physical conditions which are greatly different from jose in σας: (i) Phe magnetic field of the CSAL is the magnetic field of the stellar wind from the RC. and it is given by Bodeοἱ(1985) where & is: Doltzmann's: constant. m=1021 eis the mean particle mass. and 7; is the temperature of the stellar wind.	However, in symbiotic novae like V407 Cygni, there are some physical conditions which are greatly different from those in SNRs: (i) The magnetic field of the CSM is the magnetic field of the stellar wind from the RG, and it is given by \cite{Bode1985} where $k$ is Boltzmann's constant, $\bar{m} = 10^{-24}$ g is the mean particle mass, and $T_{\rm g}$ is the temperature of the stellar wind.
W407 ονου is a D-type 88 in which the BG has dust shells.	V407 Cygni is a D-type SS in which the RG has dust shells.
Phe temperature for the dust condensation zone is 1000 Ix Cail&Sedlmayr(1999).	The temperature for the dust condensation zone is $\sim$ 1000 K \cite{Gail1999}.
. Llere. we take 7;= 10001. The density p is given by where Ah, is the mass-loss rate of the RG. V is the stellar wind velocity. e is the binary separation. a distance Rand polar angle @ is from the WD center towards the tC.	Here, we take $T_{\rm g}=1000$ K. The density $\rho$ is given by where $\dot{M}_{\rm L}$ is the mass-loss rate of the RG, $V_{\rm w}$ is the stellar wind velocity, $a$ is the binary separation, a distance $R$ and polar angle $\theta$ is from the WD center towards the RG.
For simplicity. we only consider 6=0°.	For simplicity, we only consider $\theta=0^{\rm o}$.
As shown by Bodeetal.(L985)... Bis ~107 Gauss. which is 10° times ueher than that in the CSM of SN~ (	As shown by \cite{Bode1985}, $B$ is $\sim 10^{-2}$ Gauss, which is $10^{3}$ times higher than that in the CSM of SNR. (
i) Por a typical nova. the mass of ejecta is ~10."AL. (Yaronetal.2005).. which is much less than that of the ejecta in a typical supernova.	ii) For a typical nova, the mass of ejecta is $\sim 10^{-6} M_\odot$ \citep{Yaron2005}, which is much less than that of the ejecta in a typical supernova.
Furthermore. we note that he duration of matter ejecting in a tvpical nova is several davs or tens of davs. and the matter ejected has an average expansion velocity Vis over the whole ejecting matter phase and a maximal expansion velocity Vias.	Furthermore, we note that the duration of matter ejecting in a typical nova is $\sim$ several days or tens of days, and the matter ejected has an average expansion velocity $V_{\rm av}$ over the whole ejecting matter phase and a maximal expansion velocity $V_{\rm max}$.
Fhis means that a part of the matter ejected has the high expansion. velocity Vas	This means that a part of the matter ejected has the high expansion velocity $V_{\rm max}$.
Phe shock in à nova is mainly produced by the matter ejected with high velocity.	The shock in a nova is mainly produced by the matter ejected with high velocity.
We assume that Via=Vis and use a parameter η to define a ratio of the mass ejected with a velocity of Vias to the whole ejecta.	We assume that $V_{\rm sh}=V_{\rm max}$, and use a parameter $\eta$ to define a ratio of the mass ejected with a velocity of $V_{\rm max}$ to the whole ejecta.
Ler can be given by the following equation: where {νο is the racius of WD.	$R_{\rm ST}$ can be given by the following equation: where $R_{\rm WD}$ is the radius of WD.
Abdoetal. found that the peak [lux in 5-ravs was observed. after 3-4 davs of a nova outburst from V407 Cvgni on 10 March 2010.	\cite{Abdo2010} found that the peak flux in $\gamma$ -rays was observed after 3-4 days of a nova outburst from V407 Cygni on 10 March 2010.
0 This implies that /p in V407 Cveni should shorter than 3 days.	This implies that $t_{\rm ST}$ in V407 Cygni should shorter than 3 days.
In our model fp depends on the parameter η.	In our model $t_{\rm ST}$ depends on the parameter $\eta$.
We find that fs¢~ days when qg 0.01.	We find that $t_{\rm ST} \sim$ days when $\eta \sim$ 0.01.
‘The novae occurring in SSs are surrounded by the dense stellar winds from the Bs.	The novae occurring in SSs are surrounded by the dense stellar winds from the RGs.
They olfers an environment for high efficient. particle acceleration.	They offers an environment for high efficient particle acceleration.
Pherefore. they may be an important source of the high-energy 5-ravs in the Galaxy.	Therefore, they may be an important source of the high-energy $\gamma$ -rays in the Galaxy.
In general. SSs are the detached interacting binaries in which the WDs accrete the matter of the RCs via stellar winds.	In general, SSs are the detached interacting binaries in which the WDs accrete the matter of the RGs via stellar winds.
ὃν a population synthesis method. Lüetal.(2006). czuried out a detailed investigation of SSs.	By a population synthesis method, \cite{Lu2006} carried out a detailed investigation of SSs.
They found. that. the occurrence rate of the novae in ος is greatly. alleeteck by common-envelope evolution and the stellar wind velocity Vi of theBC.	They found that the occurrence rate of the novae in SSs is greatly affected by common-envelope evolution and the stellar wind velocity $V_{\rm w}$ of theRG.
Following Lüetal.(2006) and Zhuetal.(2010).. [or common-envelope evolution in dilflerent. simulations we use an GQeeAce=0.5 in o-algorithm and 5=1.75 in a 5- algorithm. respectively: for the stellar wind. Vio=ieu where rece is the escape velocity and Vy is determined. by he relation between the mass-loss rates and the terminal wind velocities fitted by Wintersetal.(2003) às: In this work we consider three cases with different input xwameters: ((1) iin case 1. a4.=0.5 and V=NE (ii) lin case 2.5=1.75 and V—S6 ((iii) in case 3. eX.=0.5 and Vi taken as(5).	Following \cite{Lu2006} and \cite{Zhu2010}, , for common-envelope evolution in different simulations we use an $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ in $\alpha$ -algorithm and $\gamma=1.75$ in a $\gamma$ -algorithm, respectively; for the stellar wind, $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ where $v_{\rm esc}$ is the escape velocity and $V_{\rm w}$ is determined by the relation between the mass-loss rates and the terminal wind velocities fitted by \cite{Winters2003} as: In this work we consider three cases with different input parameters: (i) in case 1, $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ and $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ (ii) in case 2, $\gamma=1.75$ and $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ (iii) in case 3, $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ and $V_{\rm w}$ taken as.
Using the model of SSs and erid for novae in Yaronetal. (2005).. we can estimate £7)luax in which 6=0" and η=0.01 or every nova.	Using the model of SSs and grid for novae in \cite{Yaron2005}, we can estimate $E^{\rm p}_{\rm max}$ in which $\theta=0^{\rm o}$ and $\eta=0.01$ for every nova.
According to Kamaeetal. (2006).. p.p interaction can occur when the energv. of proton is higher han 100 eV.κε which. results in. -rav emission.	According to \cite{Kamae2006}, , $p-p$ interaction can occur when the energy of proton is higher than $10^9$ eV, which results in $\gamma$ -ray emission.
.. sTherefore.. we assume that the novae in SSs are 5-rav sources if the £jULES.p in Eq. (1))	Therefore, we assume that the novae in SSs are $\gamma$ -ray sources if the $E^{\rm p}_{\rm max}$ in Eq. \ref{eq:epmax}) )
is higher than 10" eV. Using a population svnthesis method described in Luetal. 2008).. we model10" binary svstems which gives a statistical error for our Monte Carlo simulation lower than 5 percent for the svmbiotic novae.	is higher than $10^9$ eV. Using a population synthesis method described in \cite{Lu2006, Lu2008}, , we model$10^6$ binary systems which gives a statistical error for our Monte Carlo simulation lower than 5 percent for the symbiotic novae.
In order to estimate the occurrence rate of the x-ray sources like V407 Cvgni. we	In order to estimate the occurrence rate of the $\gamma$ -ray sources like V407 Cygni, we
iaenetosphere above the polar cap.	magnetosphere above the polar cap.
Low-level accretion. possibly of supernova fallback material. has been mvoked to ease the problem iu both cases.	Low-level accretion, possibly of supernova fallback material, has been invoked to ease the problem in both cases.
This would point to the existence of a debris disk surrounding the INS (lich could be detected in the optical-intrared range). a long sought astrophysical object. so far possibly observed ouly in the case of the Anomalous N-rav Pulsar (ANP) 0112\|61 (?).	This would point to the existence of a debris disk surrounding the INS (which could be detected in the optical-infrared range), a long sought astrophysical object, so far possibly observed only in the case of the Anomalous X-ray Pulsar (AXP) 0142+61 .
. Very deep imaging of the field of hhave been performed both from the grouud with the ESO (VLT)) and with the (IEST)).	Very deep imaging of the field of have been performed both from the ground with the ESO ) and with the ).
Optical oobservatious (De Luca et al.	Optical observations (De Luca et al.
2001) did not reveal any potential counterpart down to R~27.1 aud V~27.3 while observations iu the optical with the and in the near infrares (NIR) with the showed the preseuce of a faint source (hereafter "source Z) close to the Nav position. with maguitudes Wess026.14 and Ks~20.7(77).	2004) did not reveal any potential counterpart down to $\sim$ 27.1 and $\sim$ 27.3 while observations in the optical with the and in the near infrared (NIR) with the showed the presence of a faint source (hereafter “source Z”) close to the X-ray position, with magnitudes $_{F555W}\sim26.4$ and $\sim$ 20.7.
Ilowever. the association with wwas soon after questioned by ou the basis of precise absolute astrometry of he Ππμασος, which showed a positional offset of source Z with respect to the coordinates.	However, the association with was soon after questioned by on the basis of precise absolute astrometry of the images, which showed a positional offset of source Z with respect to the coordinates.
The same source was observed in the NIR by?.. who reported very red colours. consistent with au AL dwarf. aud also questioned its possible association to bbecause of the inconsistency with the position.	The same source was observed in the NIR by, who reported very red colours, consistent with an M dwarf, and also questioned its possible association to because of the inconsistency with the position.
also observed the field withSpitzerat L5jan aud at δρα, but did not detect auv source at the target position.	also observed the field withat $\mu$ m and at $\mu$ m, but did not detect any source at the target position.
Tere we repor onu a different. iudependent test to assess the association of source Z to5209.. using multi-epoch data collected with the ((Sect. 2)).	Here we report on a different, independent test to assess the association of source Z to, using multi-epoch data collected with the (Sect. \ref{hst}) ).
The same ddataset. completed by erouncd based data collected with theVLT.. is also used to derive strinecut constraints ou the optical/imtrared cussion from ((Sect. 3)).	The same dataset, completed by ground based data collected with the, is also used to derive stringent constraints on the optical/infrared emission from (Sect. \ref{phot}) ).
Results ave discussed in Sect. L..	Results are discussed in Sect. \ref{disc}. .
1296.5\|10.0 has a remarkable. well defined bilatera sviunetre(27).	G296.5+10.0 has a remarkable, well defined bilateral symmetry.
Very likely. the explosion site lies on the svinmnetry axis of the SNR. but the current position of 1s sigmificautly offset from the apparent ceuter of the host SNR.	Very likely, the explosion site lies on the symmetry axis of the SNR, but the current position of is significantly offset from the apparent center of the host SNR.
Iudeed. the geometrical ceuter position evaliatec by is ~δ to the south west of the N-rav source(2)..	Indeed, the geometrical center position evaluated by is $\sim8\arcmin$ to the south west of the X-ray source.
Assuming for the system au age of TOOOO years. such a displacement would tuaply a proper motion of ~TO amas ol. corresponding to a projected velocity of ~GIO kan 4. consistent with the observed. velocity distribution for radio pulsus(7).	Assuming for the system an age of 000 years, such a displacement would imply a proper motion of $\sim70$ mas $^{-1}$, corresponding to a projected velocity of $\sim640$ km $^{-1}$, consistent with the observed velocity distribution for radio pulsars.
. This offers a natura wav to test the putative identification: if iudeed associated with5209.. source Z should show a significant proper motion.	This offers a natural way to test the putative identification: if indeed associated with, source Z should show a significant proper motion.
Thus. we used iuulti-epochi oobservations to search for an angular displacement of source Z. We observed the field of with the on May sth 2007 (Proeranune 10791).	Thus, we used multi-epoch observations to search for an angular displacement of source Z. We observed the field of with the on May 8th 2007 (Programme 10791).
Our observations were originally scheduled for execution with the (WEC) of the (CACS))(22).	Our observations were originally scheduled for execution with the ) of the ).
. Unfortunately. theACS/WEC was put in idle state on January 2007 due to a failure of the on-board electronics.	Unfortunately, the was put in idle state on January 2007 due to a failure of the on-board electronics.
Our observations were re-scheduled aud executed with the C(WEPC'2)).	Our observations were re-scheduled and executed with the ).
A set of four 500s exposures were obtained duriug one spacecraft orbit. through the sliW filter (A=8OL2AÀ: AA=1539 Aj).	A set of four 500s exposures were obtained during one spacecraft orbit, through the 814W filter $\lambda= 8012$ ; $\Delta \lambda=1539$ ).
In order to exploit the maxinuun spatial resolution for the proper motion measurement. was placed at the ceutre of the (PC) Tip (00015/pixel).	In order to exploit the maximum spatial resolution for the proper motion measurement, was placed at the centre of the ) chip 045/pixel).
Our new data add up to observations collected with the oon July 28th and August 7th 2003 (Proeranuue 9872) and available in the public aarchives.	Our new data add up to observations collected with the on July 28th and August 7th 2003 (Programme 9872) and available in the public archives.
This first-epoch ddataset allowed to pick up Source Z as a possible couuterpart to5209.	This first-epoch dataset allowed to pick up Source Z as a possible counterpart to.
. TheWEC (070050/pixel) was used in both visits.	The 050/pixel) was used in both visits.
Two SOGQUOLCOR of L aud 5 exposures were obtained through the broad-banud filters 555WN (A=531\|GÀ:: AA=1193 Aj) and SIIW (A= S333A: AA=2511 ÀJ) for a total iuteeration time of 12800 s and 10200 x. respectively.	Two sequences of 4 and 5 exposures were obtained through the broad-band filters 555W $\lambda= 5346$; $\Delta \lambda=1193$ ) and 814W $\lambda= 8333$ ; $\Delta \lambda=2511$ ), for a total integration time of 12800 s and 10200 s, respectively.
The complete dataset spans a time baselime of ~3.75 vears.	The complete dataset spans a time baseline of $\sim3.75$ years.
We downloaded the data from the Space Telescope European Coordinating Facility (ST-ECF) Scieuce DataArchivel.	We downloaded the data from the Space Telescope European Coordinating Facility (ST-ECF) Science Data.
Ou-the-fly data reduction (bias and flat- correction) and flux. calibration were applied usine theSoftware (STSDAS) through the ST-EC'F Data Archive pipeline.	On-the-fly data reduction (bias and flat-field correction) and flux calibration were applied using the ) through the ST-ECF Data Archive pipeline.
To filter cosimic rav hits. sinele eexposures were combined and averaged using the taskcombine. while single exposures were combined using which also produces a mosaic mage of the two cchips aud applies the correction for the geometric distortions of the camera.	To filter cosmic ray hits, single exposures were combined and averaged using the task, while single exposures were combined using which also produces a mosaic image of the two chips and applies the correction for the geometric distortions of the camera.
the surface density of galaxies at the inner edge of the extended radio relic.	the surface density of galaxies at the inner edge of the extended radio relic.
Remarkably, the transverse extent of the wall is also comparable to the relic, ~2 Mpc.	Remarkably, the transverse extent of the wall is also comparable to the relic, $\sim$ 2 Mpc.
Fig.	Fig.
7 shows a 1-D slice across the Coma halo and relic in radio (1.4 GHz) and brightnesses.	\ref{slice} shows a 1-D slice across the Coma halo and relic in radio (1.4 GHz) and brightnesses.
There is no significant X-ray emission beyond the relic.	There is no significant X-ray emission beyond the relic.
We also plot the surface density of SDSS galaxies in three velocity bins along the same slice, with a width of 2°.	We also plot the surface density of SDSS galaxies in three velocity bins along the same slice, with a width of $^{\circ}$.
Moving away from the cluster, we find a sharp drop in the number of galaxies with 6600 < v < 8200 just as the relic radio emission increases.	Moving away from the cluster, we find a sharp drop in the number of galaxies with 6600 $<$ v $<$ 8200 just as the relic radio emission increases.
The other velocity bins, as expected, are dominated by the cluster itself and drop off gradually without any special behaviour at the relic position.	The other velocity bins, as expected, are dominated by the cluster itself and drop off gradually without any special behaviour at the relic position.
Several infalling groups of galaxies have been identified near this region (e.g., Adami et al.	Several infalling groups of galaxies have been identified near this region (e.g., Adami et al.
2005), but this is the first identification of a very broad transverse feature in the infall pattern into Coma.	2005), but this is the first identification of a very broad transverse feature in the infall pattern into Coma.
Rines et al.	Rines et al.
2003 estimate a virial radius of 2.8 Mpc (using ~1.3rag9, Eke, Cole Frenk 1996, and converting to Hp=70), so the galaxy wall is already within the virialized region but can still be isolated in space and velocity.	2003 estimate a virial radius of 2.8 Mpc (using $\sim$ $_{200}$, Eke, Cole Frenk 1996, and converting to $_0$ =70), so the galaxy wall is already within the virialized region but can still be isolated in space and velocity.
The correspondence between the wall of galaxies and the inner edge of the radio relic suggests a causal link.	The correspondence between the wall of galaxies and the inner edge of the radio relic suggests a causal link.
This poses a problem for the association between the radio relic and the NAT radio source NGC 4789, which is a possible source of seed relativistic electrons (EnBlin et al.	This poses a problem for the association between the radio relic and the NAT radio source NGC 4789, which is a possible source of seed relativistic electrons lin et al.
1998).	1998).
Given its radial velocity of 8365 km/s and the morphology of the bent jet, NGC 4789 is apparently on the back side of the cluster (moving away), while the wall of galaxies, if	Given its radial velocity of 8365 km/s and the morphology of the bent jet, NGC 4789 is apparently on the back side of the cluster (moving away), while the wall of galaxies, if
finds a maximum of 770+110rad m-?.	finds a maximum of $770\pm110$rad $^{-2}$.
The maps are similar moving north across where the RM switches from positive to negative (0?15,0?4),values; Tsuboietal.(1986) measure RM—250rad m? while the present survey finds —220+130 rad m~?.	The maps are similar moving north across $(0\ddeg15,0\ddeg4)$, where the RM switches from positive to negative values; \citet{t86} measure $\approx-250$rad $^{-2}$ while the present survey finds $-220\pm130$ rad $^{-2}$.
There is some agreement at the northwestern edge of the polarized emission of the Radio Arc, shown in Figure 5,, where the RM switches back to positive values.	There is some agreement at the northwestern edge of the polarized emission of the Radio Arc, shown in Figure \ref{rmcomp}, where the RM switches back to positive values.
The exact location of this second RM sign change is slightly different and may reflect the different RM depths each survey is sensitive to.	The exact location of this second RM sign change is slightly different and may reflect the different RM depths each survey is sensitive to.
Yusef-Zadehetal.(1997) presented a detailed study of the polarization properties of the nonthermal filament G359.544-0.18 (RF-C3) at 6 and 3.6 cm.	\citet{y97} presented a detailed study of the polarization properties of the nonthermal filament G359.54+0.18 (RF-C3) at 6 and 3.6 cm.
Figure 3 of that work has a similar 6 cm brightness and RM distribution as the present work, both presented here and in Lawetal.	Figure 3 of that work has a similar 6 cm brightness and RM distribution as the present work, both presented here and in \citet{gcl_vla}.
The RM map of the filament shows three distinct, (2008a)..bright clumps each having relatively uniform values.	The RM map of the filament shows three distinct, bright clumps each having relatively uniform values.
The morphology seen in the present survey is similar to that of Yusef-Zadehetal.(1997),, although it had roughly three times better resolution ccompared to iin the present work).	The morphology seen in the present survey is similar to that of \citet{y97}, although it had roughly three times better resolution compared to in the present work).
The first clump, at RA, Dec = -29:12:30) has RM&—2700 rad m-?, (B1950)compared to (17:40:41,—3960+1100 rad m~? in the present survey.	The first clump, at RA, Dec (B1950) = (17:40:41, –29:12:30) has $\approx-2700$ rad $^{-2}$ , compared to $-3960\pm1100$ rad $^{-2}$ in the present survey.
The second clump, at (17:40:43, —29:12:40), has RM—2000 rad m~?, compared to —2200+440 rad m? in the present survey.	The second clump, at (17:40:43, –29:12:40), has $\approx-2000$ rad $^{-2}$, compared to $-2200\pm440$ rad $^{-2}$ in the present survey.
The third clump, at (17:40:44,29:12:45), has RM—1500 rad m?, compared to —15404:660 rad m~? in the present survey.	The third clump, at (17:40:44,--29:12:45), has $\approx-1500$ rad $^{-2}$, compared to $-1540\pm660$ rad $^{-2}$ in the present survey.
We conclude that, in general, there is good agreement between the RM of the present survey and that of Yusef-Zadehetal. (1997).	We conclude that, in general, there is good agreement between the RM of the present survey and that of \citet{y97}.
. In summary, the polarized intensity and RM of the present 6 cm survey shows good agreement with those of other surveys.	In summary, the polarized intensity and RM of the present 6 cm survey shows good agreement with those of other surveys.
This is consistent with the fact the polarimetric leakgage is expected to have relatively little frequency structure for the VLA feed design (Cotton1994,1999);; any systematic errors in the polarization angle are subtracted when forming the iimage.	This is consistent with the fact the polarimetric leakgage is expected to have relatively little frequency structure for the VLA feed design \citep{c94,c99}; any systematic errors in the polarization angle are subtracted when forming the image.
It also shows that histogram fitting of the vvalues is a reasonable estimate of the RM and its uncertainty at 6 cm in this region.	It also shows that histogram fitting of the values is a reasonable estimate of the RM and its uncertainty at 6 cm in this region.
'The 6 cm polarized continuum intensity of the northern extension of the Radio Arc is several mJy beam~! and spans the entire eastern edge of the survey up to a latitude of b~0?8.	The 6 cm polarized continuum intensity of the northern extension of the Radio Arc is several mJy $^{-1}$ and spans the entire eastern edge of the survey up to a latitude of $b\sim0\ddeg8$.
To test for frequency structure in the polarized intensity, the polarized intensity maps in the two bands were differenced.	To test for frequency structure in the polarized intensity, the polarized intensity maps in the two bands were differenced.
The lack of diffuse emission in the difference map shows that the two maps have similar diffuse emission within roughly 1 mJy.	The lack of diffuse emission in the difference map shows that the two maps have similar diffuse emission within roughly 1 mJy.
The comparable 20 cm mosaic of polarized continuum shows no extended emission down to a level of about 0.1 mJy beam! detailinLawetal.2008a).	The comparable 20 cm mosaic of polarized continuum shows no extended emission down to a level of about 0.1 mJy $^{-1}$ \citep[more detail in][]{gcl_vla}.
. For latitudes(more up to b=0°3, the polarized continuum emission seen in the 6 cm interferometric maps (Fig. 1))	For latitudes up to $b=0\ddeg3$, the polarized continuum emission seen in the 6 cm interferometric maps (Fig. \ref{poln_polc}) )
has a total intensity counterpart in the same data.	has a total intensity counterpart in the same data.
However, north of b=0?3, the total intensity counterpart is too extended to be detected by the VLA 6 cm observations.	However, north of $b=0\ddeg3$, the total intensity counterpart is too extended to be detected by the VLA 6 cm observations.
Since the polarized emission is broken into small spatial scales (as shown in Figure 2)), it is detected throughout the region and the apparent polarization fraction often exceeds100%.	Since the polarized emission is broken into small spatial scales (as shown in Figure \ref{canal}) ), it is detected throughout the region and the apparent polarization fraction often exceeds.
. 'To estimate the polarization fraction without the effect of missing flux, we compare the VLA polarized-intensity maps to continuum maps from the Green Bank Telescope (Lawetal.2008b)..	To estimate the polarization fraction without the effect of missing flux, we compare the VLA polarized-intensity maps to continuum maps from the Green Bank Telescope \citep{gcsurvey_gbt}.
We convolve the VLA maps to the GBT resolution to estimate the polarization fraction; this will be a lower limit, since the VLA emission is laced with depolarized canals.	We convolve the VLA maps to the GBT resolution to estimate the polarization fraction; this will be a lower limit, since the VLA emission is laced with depolarized canals.

llowever. in svmbiotic novae like V407 Cvgni. there are some physical conditions which are greatly different from jose in σας: (i) Phe magnetic field of the CSAL is the magnetic field of the stellar wind from the RC. and it is given by Bodeοἱ(1985) where & is: Doltzmann's: constant. m=1021 eis the mean particle mass. and 7; is the temperature of the stellar wind.

However, in symbiotic novae like V407 Cygni, there are some physical conditions which are greatly different from those in SNRs: (i) The magnetic field of the CSM is the magnetic field of the stellar wind from the RG, and it is given by \cite{Bode1985} where $k$ is Boltzmann's constant, $\bar{m} = 10^{-24}$ g is the mean particle mass, and $T_{\rm g}$ is the temperature of the stellar wind.

W407 ονου is a D-type 88 in which the BG has dust shells.

V407 Cygni is a D-type SS in which the RG has dust shells.

Phe temperature for the dust condensation zone is 1000 Ix Cail&Sedlmayr(1999).

The temperature for the dust condensation zone is $\sim$ 1000 K \cite{Gail1999}.

. Llere. we take 7;= 10001. The density p is given by where Ah, is the mass-loss rate of the RG. V is the stellar wind velocity. e is the binary separation. a distance Rand polar angle @ is from the WD center towards the tC.

Here, we take $T_{\rm g}=1000$ K. The density $\rho$ is given by where $\dot{M}_{\rm L}$ is the mass-loss rate of the RG, $V_{\rm w}$ is the stellar wind velocity, $a$ is the binary separation, a distance $R$ and polar angle $\theta$ is from the WD center towards the RG.

For simplicity. we only consider 6=0°.

For simplicity, we only consider $\theta=0^{\rm o}$.

As shown by Bodeetal.(L985)... Bis ~107 Gauss. which is 10° times ueher than that in the CSM of SN~ (

As shown by \cite{Bode1985}, $B$ is $\sim 10^{-2}$ Gauss, which is $10^{3}$ times higher than that in the CSM of SNR. (

i) Por a typical nova. the mass of ejecta is ~10."AL. (Yaronetal.2005).. which is much less than that of the ejecta in a typical supernova.

ii) For a typical nova, the mass of ejecta is $\sim 10^{-6} M_\odot$ \citep{Yaron2005}, which is much less than that of the ejecta in a typical supernova.

Furthermore. we note that he duration of matter ejecting in a tvpical nova is several davs or tens of davs. and the matter ejected has an average expansion velocity Vis over the whole ejecting matter phase and a maximal expansion velocity Vias.

Furthermore, we note that the duration of matter ejecting in a typical nova is $\sim$ several days or tens of days, and the matter ejected has an average expansion velocity $V_{\rm av}$ over the whole ejecting matter phase and a maximal expansion velocity $V_{\rm max}$.

Fhis means that a part of the matter ejected has the high expansion. velocity Vas

This means that a part of the matter ejected has the high expansion velocity $V_{\rm max}$.

Phe shock in à nova is mainly produced by the matter ejected with high velocity.

The shock in a nova is mainly produced by the matter ejected with high velocity.

We assume that Via=Vis and use a parameter η to define a ratio of the mass ejected with a velocity of Vias to the whole ejecta.

We assume that $V_{\rm sh}=V_{\rm max}$, and use a parameter $\eta$ to define a ratio of the mass ejected with a velocity of $V_{\rm max}$ to the whole ejecta.

Ler can be given by the following equation: where {νο is the racius of WD.

$R_{\rm ST}$ can be given by the following equation: where $R_{\rm WD}$ is the radius of WD.

Abdoetal. found that the peak [lux in 5-ravs was observed. after 3-4 davs of a nova outburst from V407 Cvgni on 10 March 2010.

\cite{Abdo2010} found that the peak flux in $\gamma$ -rays was observed after 3-4 days of a nova outburst from V407 Cygni on 10 March 2010.

0 This implies that /p in V407 Cveni should shorter than 3 days.

This implies that $t_{\rm ST}$ in V407 Cygni should shorter than 3 days.

In our model fp depends on the parameter η.

In our model $t_{\rm ST}$ depends on the parameter $\eta$.

We find that fs¢~ days when qg 0.01.

We find that $t_{\rm ST} \sim$ days when $\eta \sim$ 0.01.

‘The novae occurring in SSs are surrounded by the dense stellar winds from the Bs.

The novae occurring in SSs are surrounded by the dense stellar winds from the RGs.

They olfers an environment for high efficient. particle acceleration.

They offers an environment for high efficient particle acceleration.

Pherefore. they may be an important source of the high-energy 5-ravs in the Galaxy.

Therefore, they may be an important source of the high-energy $\gamma$ -rays in the Galaxy.

In general. SSs are the detached interacting binaries in which the WDs accrete the matter of the RCs via stellar winds.

In general, SSs are the detached interacting binaries in which the WDs accrete the matter of the RGs via stellar winds.

ὃν a population synthesis method. Lüetal.(2006). czuried out a detailed investigation of SSs.

By a population synthesis method, \cite{Lu2006} carried out a detailed investigation of SSs.

They found. that. the occurrence rate of the novae in ος is greatly. alleeteck by common-envelope evolution and the stellar wind velocity Vi of theBC.

They found that the occurrence rate of the novae in SSs is greatly affected by common-envelope evolution and the stellar wind velocity $V_{\rm w}$ of theRG.

Following Lüetal.(2006) and Zhuetal.(2010).. [or common-envelope evolution in dilflerent. simulations we use an GQeeAce=0.5 in o-algorithm and 5=1.75 in a 5- algorithm. respectively: for the stellar wind. Vio=ieu where rece is the escape velocity and Vy is determined. by he relation between the mass-loss rates and the terminal wind velocities fitted by Wintersetal.(2003) às: In this work we consider three cases with different input xwameters: ((1) iin case 1. a4.=0.5 and V=NE (ii) lin case 2.5=1.75 and V—S6 ((iii) in case 3. eX.=0.5 and Vi taken as(5).

Following \cite{Lu2006} and \cite{Zhu2010}, , for common-envelope evolution in different simulations we use an $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ in $\alpha$ -algorithm and $\gamma=1.75$ in a $\gamma$ -algorithm, respectively; for the stellar wind, $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ where $v_{\rm esc}$ is the escape velocity and $V_{\rm w}$ is determined by the relation between the mass-loss rates and the terminal wind velocities fitted by \cite{Winters2003} as: In this work we consider three cases with different input parameters: (i) in case 1, $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ and $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ (ii) in case 2, $\gamma=1.75$ and $V_{\rm w}=\frac{1}{2}v_{\rm esc}$ (iii) in case 3, $\alpha_{\rm ce}\lambda_{\rm ce}=0.5$ and $V_{\rm w}$ taken as.

Using the model of SSs and erid for novae in Yaronetal. (2005).. we can estimate £7)luax in which 6=0" and η=0.01 or every nova.

Using the model of SSs and grid for novae in \cite{Yaron2005}, we can estimate $E^{\rm p}_{\rm max}$ in which $\theta=0^{\rm o}$ and $\eta=0.01$ for every nova.

According to Kamaeetal. (2006).. p.p interaction can occur when the energv. of proton is higher han 100 eV.κε which. results in. -rav emission.

According to \cite{Kamae2006}, , $p-p$ interaction can occur when the energy of proton is higher than $10^9$ eV, which results in $\gamma$ -ray emission.

.. sTherefore.. we assume that the novae in SSs are 5-rav sources if the £jULES.p in Eq. (1))

Therefore, we assume that the novae in SSs are $\gamma$ -ray sources if the $E^{\rm p}_{\rm max}$ in Eq. \ref{eq:epmax}) )

is higher than 10" eV. Using a population svnthesis method described in Luetal. 2008).. we model10" binary svstems which gives a statistical error for our Monte Carlo simulation lower than 5 percent for the svmbiotic novae.

is higher than $10^9$ eV. Using a population synthesis method described in \cite{Lu2006, Lu2008}, , we model$10^6$ binary systems which gives a statistical error for our Monte Carlo simulation lower than 5 percent for the symbiotic novae.

In order to estimate the occurrence rate of the x-ray sources like V407 Cvgni. we

In order to estimate the occurrence rate of the $\gamma$ -ray sources like V407 Cygni, we

iaenetosphere above the polar cap.

magnetosphere above the polar cap.

Low-level accretion. possibly of supernova fallback material. has been mvoked to ease the problem iu both cases.

Low-level accretion, possibly of supernova fallback material, has been invoked to ease the problem in both cases.

This would point to the existence of a debris disk surrounding the INS (lich could be detected in the optical-intrared range). a long sought astrophysical object. so far possibly observed ouly in the case of the Anomalous N-rav Pulsar (ANP) 0112|61 (?).

This would point to the existence of a debris disk surrounding the INS (which could be detected in the optical-infrared range), a long sought astrophysical object, so far possibly observed only in the case of the Anomalous X-ray Pulsar (AXP) 0142+61 .

. Very deep imaging of the field of hhave been performed both from the grouud with the ESO (VLT)) and with the (IEST)).

Very deep imaging of the field of have been performed both from the ground with the ESO ) and with the ).

Optical oobservatious (De Luca et al.

Optical observations (De Luca et al.

2001) did not reveal any potential counterpart down to R~27.1 aud V~27.3 while observations iu the optical with the and in the near infrares (NIR) with the showed the preseuce of a faint source (hereafter "source Z) close to the Nav position. with maguitudes Wess026.14 and Ks~20.7(77).

2004) did not reveal any potential counterpart down to $\sim$ 27.1 and $\sim$ 27.3 while observations in the optical with the and in the near infrared (NIR) with the showed the presence of a faint source (hereafter “source Z”) close to the X-ray position, with magnitudes $_{F555W}\sim26.4$ and $\sim$ 20.7.

Ilowever. the association with wwas soon after questioned by ou the basis of precise absolute astrometry of he Ππμασος, which showed a positional offset of source Z with respect to the coordinates.

However, the association with was soon after questioned by on the basis of precise absolute astrometry of the images, which showed a positional offset of source Z with respect to the coordinates.

The same source was observed in the NIR by?.. who reported very red colours. consistent with au AL dwarf. aud also questioned its possible association to bbecause of the inconsistency with the position.

The same source was observed in the NIR by, who reported very red colours, consistent with an M dwarf, and also questioned its possible association to because of the inconsistency with the position.

also observed the field withSpitzerat L5jan aud at δρα, but did not detect auv source at the target position.

also observed the field withat $\mu$ m and at $\mu$ m, but did not detect any source at the target position.

Tere we repor onu a different. iudependent test to assess the association of source Z to5209.. using multi-epoch data collected with the ((Sect. 2)).

Here we report on a different, independent test to assess the association of source Z to, using multi-epoch data collected with the (Sect. \ref{hst}) ).

The same ddataset. completed by erouncd based data collected with theVLT.. is also used to derive strinecut constraints ou the optical/imtrared cussion from ((Sect. 3)).

The same dataset, completed by ground based data collected with the, is also used to derive stringent constraints on the optical/infrared emission from (Sect. \ref{phot}) ).

Results ave discussed in Sect. L..

Results are discussed in Sect. \ref{disc}. .

1296.5|10.0 has a remarkable. well defined bilatera sviunetre(27).

G296.5+10.0 has a remarkable, well defined bilateral symmetry.

Very likely. the explosion site lies on the svinmnetry axis of the SNR. but the current position of 1s sigmificautly offset from the apparent ceuter of the host SNR.

Very likely, the explosion site lies on the symmetry axis of the SNR, but the current position of is significantly offset from the apparent center of the host SNR.

Iudeed. the geometrical ceuter position evaliatec by is ~δ to the south west of the N-rav source(2)..

Indeed, the geometrical center position evaluated by is $\sim8\arcmin$ to the south west of the X-ray source.

Assuming for the system au age of TOOOO years. such a displacement would tuaply a proper motion of ~TO amas ol. corresponding to a projected velocity of ~GIO kan 4. consistent with the observed. velocity distribution for radio pulsus(7).

Assuming for the system an age of 000 years, such a displacement would imply a proper motion of $\sim70$ mas $^{-1}$, corresponding to a projected velocity of $\sim640$ km $^{-1}$, consistent with the observed velocity distribution for radio pulsars.

. This offers a natura wav to test the putative identification: if iudeed associated with5209.. source Z should show a significant proper motion.

This offers a natural way to test the putative identification: if indeed associated with, source Z should show a significant proper motion.

Thus. we used iuulti-epochi oobservations to search for an angular displacement of source Z. We observed the field of with the on May sth 2007 (Proeranune 10791).

Thus, we used multi-epoch observations to search for an angular displacement of source Z. We observed the field of with the on May 8th 2007 (Programme 10791).

Our observations were originally scheduled for execution with the (WEC) of the (CACS))(22).

Our observations were originally scheduled for execution with the ) of the ).

. Unfortunately. theACS/WEC was put in idle state on January 2007 due to a failure of the on-board electronics.

Unfortunately, the was put in idle state on January 2007 due to a failure of the on-board electronics.

Our observations were re-scheduled aud executed with the C(WEPC'2)).

Our observations were re-scheduled and executed with the ).

A set of four 500s exposures were obtained duriug one spacecraft orbit. through the sliW filter (A=8OL2AÀ: AA=1539 Aj).

A set of four 500s exposures were obtained during one spacecraft orbit, through the 814W filter $\lambda= 8012$ ; $\Delta \lambda=1539$ ).

In order to exploit the maxinuun spatial resolution for the proper motion measurement. was placed at the ceutre of the (PC) Tip (00015/pixel).

In order to exploit the maximum spatial resolution for the proper motion measurement, was placed at the centre of the ) chip 045/pixel).

Our new data add up to observations collected with the oon July 28th and August 7th 2003 (Proeranuue 9872) and available in the public aarchives.

Our new data add up to observations collected with the on July 28th and August 7th 2003 (Programme 9872) and available in the public archives.

This first-epoch ddataset allowed to pick up Source Z as a possible couuterpart to5209.

This first-epoch dataset allowed to pick up Source Z as a possible counterpart to.

. TheWEC (070050/pixel) was used in both visits.

The 050/pixel) was used in both visits.

Two SOGQUOLCOR of L aud 5 exposures were obtained through the broad-banud filters 555WN (A=531|GÀ:: AA=1193 Aj) and SIIW (A= S333A: AA=2511 ÀJ) for a total iuteeration time of 12800 s and 10200 x. respectively.

Two sequences of 4 and 5 exposures were obtained through the broad-band filters 555W $\lambda= 5346$; $\Delta \lambda=1193$ ) and 814W $\lambda= 8333$ ; $\Delta \lambda=2511$ ), for a total integration time of 12800 s and 10200 s, respectively.

The complete dataset spans a time baselime of ~3.75 vears.

The complete dataset spans a time baseline of $\sim3.75$ years.

We downloaded the data from the Space Telescope European Coordinating Facility (ST-ECF) Scieuce DataArchivel.

We downloaded the data from the Space Telescope European Coordinating Facility (ST-ECF) Science Data.

Ou-the-fly data reduction (bias and flat- correction) and flux. calibration were applied usine theSoftware (STSDAS) through the ST-EC'F Data Archive pipeline.

On-the-fly data reduction (bias and flat-field correction) and flux calibration were applied using the ) through the ST-ECF Data Archive pipeline.

To filter cosimic rav hits. sinele eexposures were combined and averaged using the taskcombine. while single exposures were combined using which also produces a mosaic mage of the two cchips aud applies the correction for the geometric distortions of the camera.

To filter cosmic ray hits, single exposures were combined and averaged using the task, while single exposures were combined using which also produces a mosaic image of the two chips and applies the correction for the geometric distortions of the camera.

the surface density of galaxies at the inner edge of the extended radio relic.

Remarkably, the transverse extent of the wall is also comparable to the relic, ~2 Mpc.

Remarkably, the transverse extent of the wall is also comparable to the relic, $\sim$ 2 Mpc.

Fig.

7 shows a 1-D slice across the Coma halo and relic in radio (1.4 GHz) and brightnesses.

\ref{slice} shows a 1-D slice across the Coma halo and relic in radio (1.4 GHz) and brightnesses.

There is no significant X-ray emission beyond the relic.

We also plot the surface density of SDSS galaxies in three velocity bins along the same slice, with a width of 2°.

We also plot the surface density of SDSS galaxies in three velocity bins along the same slice, with a width of $^{\circ}$.

Moving away from the cluster, we find a sharp drop in the number of galaxies with 6600 < v < 8200 just as the relic radio emission increases.

Moving away from the cluster, we find a sharp drop in the number of galaxies with 6600 $<$ v $<$ 8200 just as the relic radio emission increases.

The other velocity bins, as expected, are dominated by the cluster itself and drop off gradually without any special behaviour at the relic position.

Several infalling groups of galaxies have been identified near this region (e.g., Adami et al.

2005), but this is the first identification of a very broad transverse feature in the infall pattern into Coma.

Rines et al.

2003 estimate a virial radius of 2.8 Mpc (using ~1.3rag9, Eke, Cole Frenk 1996, and converting to Hp=70), so the galaxy wall is already within the virialized region but can still be isolated in space and velocity.

2003 estimate a virial radius of 2.8 Mpc (using $\sim$ $_{200}$, Eke, Cole Frenk 1996, and converting to $_0$ =70), so the galaxy wall is already within the virialized region but can still be isolated in space and velocity.

The correspondence between the wall of galaxies and the inner edge of the radio relic suggests a causal link.

This poses a problem for the association between the radio relic and the NAT radio source NGC 4789, which is a possible source of seed relativistic electrons (EnBlin et al.

This poses a problem for the association between the radio relic and the NAT radio source NGC 4789, which is a possible source of seed relativistic electrons lin et al.

1998).

Given its radial velocity of 8365 km/s and the morphology of the bent jet, NGC 4789 is apparently on the back side of the cluster (moving away), while the wall of galaxies, if

finds a maximum of 770+110rad m-?.

finds a maximum of $770\pm110$rad $^{-2}$.

The maps are similar moving north across where the RM switches from positive to negative (0?15,0?4),values; Tsuboietal.(1986) measure RM—250rad m? while the present survey finds —220+130 rad m~?.

The maps are similar moving north across $(0\ddeg15,0\ddeg4)$, where the RM switches from positive to negative values; \citet{t86} measure $\approx-250$rad $^{-2}$ while the present survey finds $-220\pm130$ rad $^{-2}$.

There is some agreement at the northwestern edge of the polarized emission of the Radio Arc, shown in Figure 5,, where the RM switches back to positive values.

There is some agreement at the northwestern edge of the polarized emission of the Radio Arc, shown in Figure \ref{rmcomp}, where the RM switches back to positive values.

The exact location of this second RM sign change is slightly different and may reflect the different RM depths each survey is sensitive to.

Yusef-Zadehetal.(1997) presented a detailed study of the polarization properties of the nonthermal filament G359.544-0.18 (RF-C3) at 6 and 3.6 cm.

\citet{y97} presented a detailed study of the polarization properties of the nonthermal filament G359.54+0.18 (RF-C3) at 6 and 3.6 cm.

Figure 3 of that work has a similar 6 cm brightness and RM distribution as the present work, both presented here and in Lawetal.

Figure 3 of that work has a similar 6 cm brightness and RM distribution as the present work, both presented here and in \citet{gcl_vla}.

The RM map of the filament shows three distinct, (2008a)..bright clumps each having relatively uniform values.

The RM map of the filament shows three distinct, bright clumps each having relatively uniform values.

The morphology seen in the present survey is similar to that of Yusef-Zadehetal.(1997),, although it had roughly three times better resolution ccompared to iin the present work).

The morphology seen in the present survey is similar to that of \citet{y97}, although it had roughly three times better resolution compared to in the present work).

The first clump, at RA, Dec = -29:12:30) has RM&—2700 rad m-?, (B1950)compared to (17:40:41,—3960+1100 rad m~? in the present survey.

The first clump, at RA, Dec (B1950) = (17:40:41, –29:12:30) has $\approx-2700$ rad $^{-2}$ , compared to $-3960\pm1100$ rad $^{-2}$ in the present survey.

The second clump, at (17:40:43, —29:12:40), has RM—2000 rad m~?, compared to —2200+440 rad m? in the present survey.

The second clump, at (17:40:43, –29:12:40), has $\approx-2000$ rad $^{-2}$, compared to $-2200\pm440$ rad $^{-2}$ in the present survey.

The third clump, at (17:40:44,29:12:45), has RM—1500 rad m?, compared to —15404:660 rad m~? in the present survey.

The third clump, at (17:40:44,--29:12:45), has $\approx-1500$ rad $^{-2}$, compared to $-1540\pm660$ rad $^{-2}$ in the present survey.

We conclude that, in general, there is good agreement between the RM of the present survey and that of Yusef-Zadehetal. (1997).

We conclude that, in general, there is good agreement between the RM of the present survey and that of \citet{y97}.

. In summary, the polarized intensity and RM of the present 6 cm survey shows good agreement with those of other surveys.

In summary, the polarized intensity and RM of the present 6 cm survey shows good agreement with those of other surveys.

This is consistent with the fact the polarimetric leakgage is expected to have relatively little frequency structure for the VLA feed design (Cotton1994,1999);; any systematic errors in the polarization angle are subtracted when forming the iimage.

This is consistent with the fact the polarimetric leakgage is expected to have relatively little frequency structure for the VLA feed design \citep{c94,c99}; any systematic errors in the polarization angle are subtracted when forming the image.

It also shows that histogram fitting of the vvalues is a reasonable estimate of the RM and its uncertainty at 6 cm in this region.

It also shows that histogram fitting of the values is a reasonable estimate of the RM and its uncertainty at 6 cm in this region.

'The 6 cm polarized continuum intensity of the northern extension of the Radio Arc is several mJy beam~! and spans the entire eastern edge of the survey up to a latitude of b~0?8.

The 6 cm polarized continuum intensity of the northern extension of the Radio Arc is several mJy $^{-1}$ and spans the entire eastern edge of the survey up to a latitude of $b\sim0\ddeg8$.

To test for frequency structure in the polarized intensity, the polarized intensity maps in the two bands were differenced.

The lack of diffuse emission in the difference map shows that the two maps have similar diffuse emission within roughly 1 mJy.

The comparable 20 cm mosaic of polarized continuum shows no extended emission down to a level of about 0.1 mJy beam! detailinLawetal.2008a).

The comparable 20 cm mosaic of polarized continuum shows no extended emission down to a level of about 0.1 mJy $^{-1}$ \citep[more detail in][]{gcl_vla}.

. For latitudes(more up to b=0°3, the polarized continuum emission seen in the 6 cm interferometric maps (Fig. 1))

For latitudes up to $b=0\ddeg3$, the polarized continuum emission seen in the 6 cm interferometric maps (Fig. \ref{poln_polc}) )

has a total intensity counterpart in the same data.

However, north of b=0?3, the total intensity counterpart is too extended to be detected by the VLA 6 cm observations.

However, north of $b=0\ddeg3$, the total intensity counterpart is too extended to be detected by the VLA 6 cm observations.

Since the polarized emission is broken into small spatial scales (as shown in Figure 2)), it is detected throughout the region and the apparent polarization fraction often exceeds100%.

Since the polarized emission is broken into small spatial scales (as shown in Figure \ref{canal}) ), it is detected throughout the region and the apparent polarization fraction often exceeds.

. 'To estimate the polarization fraction without the effect of missing flux, we compare the VLA polarized-intensity maps to continuum maps from the Green Bank Telescope (Lawetal.2008b)..

To estimate the polarization fraction without the effect of missing flux, we compare the VLA polarized-intensity maps to continuum maps from the Green Bank Telescope \citep{gcsurvey_gbt}.

We convolve the VLA maps to the GBT resolution to estimate the polarization fraction; this will be a lower limit, since the VLA emission is laced with depolarized canals.