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We anodel the photon spectrum (photous 1 cu? keV i) usine the standard deconvolution and the Levenbere-Alarquardt uoulinear least-squares fitting algorithia that incorporates model variances. | We model the photon spectrum (photons $^{-1}$ $^{2}$ $^{-1}$ ) using the standard deconvolution and the Levenberg-Marquardt nonlinear least-squares fitting algorithm that incorporates model variances. |
The spectra were modeled using CONT channels 2-114. which covered the enerev rauge of ~ 30-1800 keV. Count spectra frou the two brightest detectors (i.c.. the two detectors with the simallest source angles to the LAD normal vector) were eenerallv used to make the fits. | The spectra were modeled using CONT channels 2-14, which covered the energy range of $\sim 30$ -1800 keV. Count spectra from the two brightest detectors (i.e., the two detectors with the smallest source angles to the LAD normal vector) were generally used to make the fits. |
Iu some cases. the source aneles differed substantially (2 20°) resulting in a normalization offset iu the count spectra between the two detectors. | In some cases, the source angles differed substantially $>20^{\circ}$ ) resulting in a normalization offset in the count spectra between the two detectors. |
The data from the two detectors were fit jointly with a inultiplicative effective area correction terni in the spectral model. | The data from the two detectors were fit jointly with a multiplicative effective area correction term in the spectral model. |
For bursts in which the effective area correction was small (E 554). we sunuaed the CONT count rates from the individual detectors to maximize the count statistics. | For bursts in which the effective area correction was small $\lax 5\%$ ), we summed the CONT count rates from the individual detectors to maximize the count statistics. |
These were cases in which the source angles of the two detectors differed by only a few degrees. | These were cases in which the source angles of the two detectors differed by only a few degrees. |
Detectors with angles to the source exceeding 60° or with strong signal from sources such as Vela N-1 or (νο XN-1 were excluded fro the fit. | Detectors with angles to the source exceeding $60^{\circ}$ or with strong signal from sources such as Vela X-1 or Cyg X-1 were excluded from the fit. |
The free paraicters of the power-law spectral model are the amplitude aud the power-law iudex o,. | The free parameters of the power-law spectral model are the amplitude and the power-law index $\alpha_{p}$. |
The free paralucters of the BPL and SBPL are the amplitude. low-energv Index oj. ligh-cnereyv index Ayien. and the break energv {νι | The free parameters of the BPL and SBPL are the amplitude, low-energy index $\alpha_{\rm low}$, high-energy index $\alpha_{\rm high}$, and the break energy $E_{b}$. |
The slopes of the spectral energy flux. £z, are readily obtained from the simple relations: a=ayy11 aud a’=Agen|1. | The slopes of the spectral energy flux, $F_{\nu}$ are readily obtained from the simple relations: $\alpha = \alpha_{\rm low} + 1$ and $\alpha^{\prime} = \alpha_{\rm high} + 1$. |
For the decay emission of each burst. we modeled the time-inteerated spectrum defined over a time interval that was either the same as or shorter in leneth than the time iuterval used in making the temporal fits. | For the decay emission of each burst, we modeled the time-integrated spectrum defined over a time interval that was either the same as or shorter in length than the time interval used in making the temporal fits. |
Iu all cases. the time interval was restricteddecag. | In all cases, the time interval was restricted. |
Shorter time intervals were used for events with weaker signal-to-noise. | Shorter time intervals were used for events with weaker signal-to-noise. |
For the PF class bursts. however. we selected the eutire burst enüssion (spectra of the burst intervals starting at the peak eave ucarly ideutical paramcter values). | For the PF class bursts, however, we selected the entire burst emission (spectra of the burst intervals starting at the peak gave nearly identical parameter values). |
For the imajoritv of events. the SBPL model was preferred over the single power-law model. | For the majority of events, the SBPL model was preferred over the single power-law model. |
Sunuuarzed in Table 2 are the spectral fit paramecters for eveuts where the SBPL was the better choice of model based on the 4? statistic. | Summarized in Table 2 are the spectral fit parameters for events where the SBPL was the better choice of model based on the $\chi^{2}$ statistic. |
Iu 6 bursts. however. the fitted value of the ligh-ΟΠΟΙΟΥ slope was unusually steep (SLO). | In 6 bursts, however, the fitted value of the high-energy slope was unusually steep $\lax -4.0$ ). |
Preece et ((1998) poiuted out that spectral models with curvature cal soinetines overestimate the steepness of the ligh-energv slope. depending on how well the data tolerate curvature (sce figure Lin Preeceetal.1998)). | Preece et (1998) pointed out that spectral models with curvature can sometimes overestimate the steepness of the high-energy slope, depending on how well the data tolerate curvature (see figure 1 in \cite{preece98}) ). |
The broken power-law model was therefore used for these 6 events. resulting in slightly better reduced. 47. values aud better constrained values of the fitted higli-energv slope. | The broken power-law model was therefore used for these 6 events, resulting in slightly better reduced $\chi^{2}$ values and better constrained values of the fitted high-energy slope. |
This resulted im spectral parameters for a total of 20 bursts. | This resulted in spectral parameters for a total of 20 bursts. |
The reduced κ values are reasonable. althoneh a few bursts for which joiut fts were made tended to eive sliebitlv. larger values (A2/d.o.f.22). | The reduced $\chi^{2}$ values are reasonable, although a few bursts for which joint fits were made tended to give slightly larger values $\chi^{2}/d.o.f. \gax 2$ ). |
Given in the table are the fitted values of the spectral indices. the break energy. aud their uncertainties from the covariance matrix. | Given in the table are the fitted values of the spectral indices, the break energy, and their uncertainties from the covariance matrix. |
Also given is he difference in spectral slope across the break energy. A—|a'a|. aud the value of p calculated from the iel-enerev spectral slope. | Also given is the difference in spectral slope across the break energy, $\Delta = \mid\alpha^{\prime} - \alpha\mid$, and the value of $p$ calculated from the high-energy spectral slope. |
Of the remaining bursts. 9 events were best represented x the sinele power-law function. | Of the remaining bursts, 9 events were best represented by the single power-law function. |
The best fit paramcters and the corresponding value of p are preseuted in Table 3. | The best fit parameters and the corresponding value of $p$ are presented in Table 3. |
The spectral fits for the remaining 12 eveuts resulted iu )»or 4? values and poorly coustrained parameter values. regardless of the choice of spectral model. | The spectral fits for the remaining 12 events resulted in poor $\chi^{2}$ values and poorly constrained parameter values, regardless of the choice of spectral model. |
These are clearly cases when the counting statistics are too poor to constrain he model parameters aud therefore were excluded. | These are clearly cases when the counting statistics are too poor to constrain the model parameters and therefore were excluded. |
The results in Table 3 should be interpreted with some degree of caution. | The results in Table 3 should be interpreted with some degree of caution. |
These eveuts may be cases in which the flux evel was to low. causing the break in the spectrum to be washed out in the counting noise. | These events may be cases in which the flux level was to low, causing the break in the spectrum to be washed out in the counting noise. |
Iu such cases. the sinele vower-law will often be adequate to model the spectrin. even though the true burst spectruu may coutain a break. | In such cases, the single power-law will often be adequate to model the spectrum, even though the true burst spectrum may contain a break. |
A careful inspection of Table 2 inunediately allows us to identify lieh-cnereyv afterelow candidates based ou three characteristic signatures of the svuchrotron spectra which we categorize as the following: (1) iu the fast-cooling mode. the spectral slope below the high-cncreyv break (144) is 1/2 for radiative or adiabatie evolution. as secu from equations 3 and L. (3) i the slow-cooling mode. the change in spectral slope across the high-cucrey break G4.) 1s 1/2. as secon from equation 5. aud (111) the electron energv index. p. calculated frou, the measured spectral slope should have a value in the range 2.0<p<2.5. the typical raneederived froin afterelows observed at N-rav. optical. aud radio wavelengths. | A careful inspection of Table 2 immediately allows us to identify high-energy afterglow candidates based on three characteristic signatures of the synchrotron spectrum which we categorize as the following: (i) in the fast-cooling mode, the spectral slope below the high-energy break $\nu_{\rm m}$ ) is $-1/2$ for radiative or adiabatic evolution, as seen from equations 3 and 4, (ii) in the slow-cooling mode, the change in spectral slope across the high-energy break $\nu_{\rm c}$ ) is $1/2$, as seen from equation 5, and (iii) the electron energy index, $p$, calculated from the measured spectral slope should have a value in the range $2.0 \le p \le 2.5$, the typical rangederived from afterglows observed at X-ray, optical, and radio wavelengths. |
Applying these criteria. we label events with these properties iu the last colunn of Table 2. | Applying these criteria, we label events with these properties in the last column of Table 2. |
We thus iunuediatelv ideutify several fast-cooliug candidates: CRDB920622. GCRDBO910119b. CRDB960530(2). GRD9850301. | We thus immediately identify several fast-cooling candidates: GRB920622, GRB940419b, GRB960530(2), GRB980301. |
Each of these events has a value of α within sje1ia of 0.5 aud a value of p simular to those fouud for afterelows. | Each of these events has a value of $\alpha$ within one-sigma of $-0.5$ and a value of $p$ similar to those found for afterglows. |
GRDB970111. GRDB971208.. GRDB990316. and CGRDB9905158 are only iareinally cousisteut with fast-cooling. lavine larecy p values aud reduced 47x. Noue ofthe events iu Table 2 are consistent (within one-sigina) with A= 0.5. suggesting no slow-cooliug | GRB970411, GRB971208, GRB990316, and GRB990518 are only marginally consistent with fast-cooling, having larger $p$ values and reduced $\chi^2$ 's. None ofthe events in Table 2 are consistent (within one-sigma) with $\Delta = 0.5$ , suggesting no slow-cooling |
To tighten constraints on the ages of the GC's. we show the positions of our composite metal-poor and. metal-rich GCs in Figure 6.. | To tighten constraints on the ages of the GCs, we show the positions of our composite metal-poor and metal-rich GCs in Figure \ref{fig:ssp2}. |
We find that from the and L9i indices. the metal rich composite GC appears somewhat vounger than the metal-poor GC. | We find that from the $\beta$ and $\delta_A$ indices, the metal rich composite GC appears somewhat younger than the metal-poor GC. |
The difference varies between the cilferent age diagnostics. but isin the range 25 Civr at ~ 29. and is model dependent. | The difference varies between the different age diagnostics, but is in the range 2–5 Gyr at $\sim$ $\sigma$, and is model dependent. |
Needless to say. this is consistent with the GCs being coeval (and old) in these current data. | Needless to say, this is consistent with the GCs being coeval (and old) in these current data. |
In summary. the individual S/N of the GC's are generally too low to place constraints upon individual GC ages. | In summary, the individual S/N of the GCs are generally too low to place constraints upon individual GC ages. |
“Pherefore we have co-added: the individual. spectra (separated by metallicity) and. find that. these composite NGC 524 GC's have line-streneths consistent with old stellar populations. | Therefore we have co-added the individual spectra (separated by metallicity) and find that these composite NGC 524 GCs have line-strengths consistent with old stellar populations. |
The metal-poor and metal-rich sub-populations ave old and coeval at the 20 level of confidence. | The metal-poor and metal-rich sub-populations are old and coeval at the $\sigma$ level of confidence. |
lt has been known for some time that luminous elliptical ealaxiecs exhibit non-solar abundance ratios from their integrated. light. O'Connell 1976: Peletier 1950: Worthev.Faber.&Gonzalez 1902). | It has been known for some time that luminous elliptical galaxies exhibit non-solar abundance ratios from their integrated light \citeANP{OConnell76} 1976; \citeANP{Peletier89} 1989; \citeANP{Worthey92} 1992). |
Integrated light studies also suggest that GCs associated. with both ellipticals (NGC 1899: Ixisslor-DPatigetal. 1998: Forbesetal. 2001) and spirals CM 381: deFreitasPacheco L997: the Sombrero: Larsenetal. 2002: the Alilky Way: Borgesetal. 1995 and see later this section) also exhibit similar abundance To assess the a /Ee] ratios in our data. following | Integrated light studies also suggest that GCs associated with both ellipticals (NGC 1399: \citeANP{KisslerPatig98} 1998; \citeANP{Forbes01} 2001) and spirals (M 31: \citeANP{deFreitasPacheco97} 1997; the Sombrero: \citeANP{Larsen02} 2002; the Milky Way: \citeANP{Borges95} 1995 and see later this section) also exhibit similar abundance To assess the $\alpha$ /Fe] ratios in our data, following |
on ΑΕΛ][Vs] plane. where LY,[ and. [Vi] are projected distance and velocity ofa GC. respectively. can provide some information on the orbital eccentricity (ον) of the previous host galaxy of LOGCs. | on $|X_{\rm p}|$ $|X_{\rm p}| \times |V_{\rm p}|$ plane, where $|X_{\rm p}|$ and $|V_{\rm p}|$ are projected distance and velocity of a GC, respectively, can provide some information on the orbital eccentricity $e_{\rm p}$ ) of the previous host galaxy of ICGCs. |
ο(ος can be distributed. in the upper right region with larger .X,|[Vi| and larger LV, in the low ey model. | ICGCs can be distributed in the upper right region with larger $|X_{\rm p}| \times |V_{\rm p}|$ and larger $|X_{\rm p}|$ in the low $e_{\rm p}$ model. |
Phe lower panel of Fig. | The lower panel of Fig. |
S can provide a clearer cillerence. and demonstrates that the high e, model has the narrower distribution in the histogram of δρ[Vi] and a larger number of GCs in the smaller. LY]η bins. | 8 can provide a clearer difference, and demonstrates that the high $e_{\rm p}$ model has the narrower distribution in the histogram of $|X_{\rm p}| \times |V_{\rm p}|$ and a larger number of GCs in the smaller $|X_{\rm p}| \times |V_{\rm p}|$ bins. |
From these results. we expect that if a cluster consists mostly of galaxies with higher orbital eccentricities. it shows narrower δρ][V] histogram of ICCGCs and their {ος are located: preferentially in the lower region of the LV,|- δρο) plane of LOGCs. | From these results, we expect that if a cluster consists mostly of galaxies with higher orbital eccentricities, it shows narrower $|X_{\rm p}| \times |V_{\rm p}|$ histogram of ICGCs and their ICGCs are located preferentially in the lower region of the $|X_{\rm p}|$ $|X_{\rm p}| \times |V_{\rm p}|$ plane of ICGCs. |
We can investigate the above just basecl on the projected distance ancl radial velocity of each LOGC in a cluster. | We can investigate the above just based on the projected distance and radial velocity of each ICGC in a cluster. |
Fig. | Fig. |
9 suggests that if we know the mass profile of a cluster and thus its rotation curve through X-ray hot gas observations anc kinematical properties of cluster member galaxies. we can provide even clearer evidence on orbital properties of cluster galaxies by combining the LCCC kinematics and the cluster mass profile. | 9 suggests that if we know the mass profile of a cluster and thus its rotation curve through X-ray hot gas observations and kinematical properties of cluster member galaxies, we can provide even clearer evidence on orbital properties of cluster galaxies by combining the ICGC kinematics and the cluster mass profile. |
This figure demonstrates that if the orbits of cluster galaxies are more eccentric as a whole. LOGCs (stripped from these) populate mostly in the upper left region on the Nu mns plane. where Vi is the circular velocity of a ICCGC at LV, and estimated. from the adopted mass profile of a cluster. | This figure demonstrates that if the orbits of cluster galaxies are more eccentric as a whole, ICGCs (stripped from these) populate mostly in the upper left region on the $|X_{\rm
p}| \times |V_{\rm p}|$ $1-|V_{\rm p}/V_{\rm c}|$ ) plane, where $V_{\rm c}$ is the circular velocity of a ICGC at $|X_{\rm p}|$ and estimated from the adopted mass profile of a cluster. |
‘This is essentially because these LOGCs also have highly eccentric orbits anc thus smaller. projected. velocity [Vj] for the V. at rely positions. | This is essentially because these ICGCs also have highly eccentric orbits and thus smaller projected velocity $|V_{\rm p}|$ for the $V_{\rm c}$ at their positions. |
Furthermore we can more clearly see the dillerence between the mocdels with cilferent orbital eccentricities in the histogram οἱ 1πιτοι | Furthermore we can more clearly see the difference between the models with different orbital eccentricities in the histogram of $1-|V_{\rm p}/V_{\rm c}|$. |
This cilference suggests that if the orbits of cluster galaxies are more eccentric as a whole. such a cluster may show the peak at high values of 1.[V5/M.]. | This difference suggests that if the orbits of cluster galaxies are more eccentric as a whole, such a cluster may show the peak at high values of $1-|V_{\rm p}/V_{\rm c}|$. |
H should be noted here that the observed. difference in the above properties. of ICCGCS may not be clearly distinguished as proposed because of the mixture of cillerent orbits of galaxies in a cluster. | It should be noted here that the observed difference in the above properties of ICGCs may not be clearly distinguished as proposed because of the mixture of different orbits of galaxies in a cluster. |
However future observations on LOGC's in different nearby clusters of galaxies can still provide valuable constraints on galaxy evolution in clusters. | However future observations on ICGCs in different nearby clusters of galaxies can still provide valuable constraints on galaxy evolution in clusters. |
We have numerically investigated: the roles of the cluster ical field in dynamical evolution of the GC system of NGC 1404. | We have numerically investigated the roles of the cluster tidal field in dynamical evolution of the GC system of NGC 1404. |
We summarise our principle results as follows. ( | We summarise our principle results as follows. ( |
1) Final Sx depends both on ey (orbital eccentricity of GC. 1404) and on es (the ratio of the scale length of the GC system to the elfective radius of stellar component of GC 1404) in such à way that Sp is smaller in the moclels with larger ey, and larger duc. | 1) Final $S_{\rm N}$ depends both on $e_{\rm p}$ (orbital eccentricity of NGC 1404) and on $a_{\rm gc}$ (the ratio of the scale length of the GC system to the effective radius of stellar component of NGC 1404) in such a way that $S_{\rm N}$ is smaller in the models with larger $e_{\rm p}$ and larger $a_{\rm gc}$. |
Sp can be significantly reduced ον tidal stripping of GCs to become as low as the observed value (~ 2). only if the orbit of NCC 1404 is highly eccentric (with orbital eccentricity of 2 0.5) and if the initial scale ength of the GCs distribution is about twice as large as the ellective radius of NGC 1404. ( | $S_{\rm N}$ can be significantly reduced by tidal stripping of GCs to become as low as the observed value $\sim$ 2), only if the orbit of NGC 1404 is highly eccentric (with orbital eccentricity of $>$ 0.5) and if the initial scale length of the GCs distribution is about twice as large as the effective radius of NGC 1404. ( |
2) The radial number density profile of the GC system comes. steeper after the stripping of GC's. because the outer GC's are more easily stripped during tidal interaction. | 2) The radial number density profile of the GC system becomes steeper after the stripping of GCs, because the outer GCs are more easily stripped during tidal interaction. |
This result. implies that since the observed slope of the power-law density. profile of the GC system in NGC 1401 is shallower (~ —1.3) compared. with the typical value or cluster ellipticals 1.9). the initial slope (before tidal interaction) was rather shallow (1o. larger than 1.9). if the observed Low Spx is due to tidal stripping of GC's. ( | This result implies that since the observed slope of the power-law density profile of the GC system in NGC 1404 is shallower $\sim$ $-1.3$ ) compared with the typical value for cluster ellipticals $-1.9$ ), the initial slope (before tidal interaction) was rather shallow (i.e. larger than $-1.3$ ), if the observed low $S_{\rm N}$ is due to tidal stripping of GCs. ( |
3) One of the observable characteristics of a cluster elliptical with a low Sp. (« 2) resulting from tidal stripping is the larger ratio of Sx within 12H. 10 that within 510 Do where A. is the ellective radius of a cluster elliptical galaxy. | 3) One of the observable characteristics of a cluster elliptical with a low $S_{\rm N}$ $<$ 2) resulting from tidal stripping is the larger ratio of $S_{\rm N}$ within $1-2R_{\rm e}$ to that within 5–10 $R_{\rm e}$, where $R_{\rm e}$ is the effective radius of a cluster elliptical galaxy. |
We also suggest that if tidal stripping of GC's via a cluster tidal field is a main cause for the evolution of Sx of cluster ellipticals. there should. be a positive correlation between the distance of a cluster elliptical from the centre of a cluster ancl y of the galaxy (as may be present in the Fornax cluster). ( | We also suggest that if tidal stripping of GCs via a cluster tidal field is a main cause for the evolution of $S_{\rm N}$ of cluster ellipticals, there should be a positive correlation between the distance of a cluster elliptical from the centre of a cluster and $S_{\rm N}$ of the galaxy (as may be present in the Fornax cluster). ( |
4) Stripped GCs are found to become intracluster GC's (CGCs) orbiting the centre of Fornax cluster (ic... NGC 1399) and their physical properties (e.g... number. racial distribution. and kinematics with respect to the cluster centre) depend. on the orbit of NGC 1404 and the initial distribution of the GCs in NCC 1404. | 4) Stripped GCs are found to become intracluster GCs (ICGCs) orbiting the centre of Fornax cluster (i.e., NGC 1399) and their physical properties (e.g., number, radial distribution, and kinematics with respect to the cluster centre) depend on the orbit of NGC 1404 and the initial distribution of the GCs in NGC 1404. |
For example. the LOGCs within the cluster core have a projected. number density profile with the power-law slope of ~ 0.9 and rather large velocity dispersion ( 340 kms 13 for the highly eccentric orbit model of NGC 1404. ( | For example, the ICGCs within the cluster core have a projected number density profile with the power-law slope of $\sim$ $-0.9$ and rather large velocity dispersion $\sim$ 340 km $^{-1}$ ) for the highly eccentric orbit model of NGC 1404. ( |
5) Our numerical results suggest wat not only structural properties but. the kinematical properties of LOGCs formecl from tidal stripping can depend: strongly on the orbits of their previous host. galaxies. | 5) Our numerical results suggest that not only structural properties but the kinematical properties of ICGCs formed from tidal stripping can depend strongly on the orbits of their previous host galaxies. |
This implies that the detailed investigation ofICCGC kinematies by multi-object spectrographs on S-l0nmreclass telescopes can shed new insight into galaxy dynamics in the cluster as a whole. | This implies that the detailed investigation of ICGC kinematics by multi-object spectrographs on 8-10m-class telescopes can shed new insight into galaxy dynamics in the cluster as a whole. |
We are grateful to the referee. Oleg Cnedin. for. valuable comments. which contribute to improve the present. paper. | We are grateful to the referee Oleg Gnedin for valuable comments, which contribute to improve the present paper. |
KB and. WJC acknowledge the financial support. of the Australian Research Council throughout the course of this work. | KB and WJC acknowledge the financial support of the Australian Research Council throughout the course of this work. |
ALB would like to thank the Swinburne Research and Development Grants Scheme. | MB would like to thank the Swinburne Research and Development Grants Scheme. |
S(v)« v), where the relativistic electrons have suffered large energy losses after their ejection from the parent galaxy. | $S(\nu) \propto \nu^{\alpha}$ ), where the relativistic electrons have suffered large energy losses after their ejection from the parent galaxy. |
At 25 cm the total extent of the Beaver is ~240" (360 kpc), and it is polarized at16.1%,,13.3%,, and at 18, 21, and 25 cm, respectively. | At 25 cm the total extent of the Beaver is $\sim240{\arcsec}$ (360 kpc), and it is polarized at, and at 18, 21, and 25 cm, respectively. |
In our high-frequency observations, the three filaments are clearly detected both in total intensity and polarization. | In our high-frequency observations, the three filaments are clearly detected both in total intensity and polarization. |
In the following, we refer to them using the convention adopted in Fig. | In the following, we refer to them using the convention adopted in Fig. |
1 in ?.. | 1 in \citet{gov}. |
They lie near the cluster center and are located at the edges of the halo, which therefore has an uncommon rectangular shape. | They lie near the cluster center and are located at the edges of the halo, which therefore has an uncommon rectangular shape. |
Between 18 cm and 25 cm, their morphology and size do not change. | Between 18 cm and 25 cm, their morphology and size do not change. |
The angular extent of F1, F2, and F3 is 323" (485 kpc), 330" (500 kpc), and 370" (550 kpc), respectively. | The angular extent of F1, F2, and F3 is ${\arcsec}$ (485 kpc), ${\arcsec}$ (500 kpc), and ${\arcsec}$ (550 kpc), respectively. |
Their fractional polarization ranges between (see Table 6)) with an uncertainty of ~2%. | Their fractional polarization ranges between (see Table \ref{polarisationpercentages}) ) with an uncertainty of $\sim 2$. |
. To obtain the rotation measure maps of the radio galaxies and the filaments, we produced masks of the sources from the total intensity image and applied them to the high-frequency RM cube (18 cm + 21 cm + 25 cm). | To obtain the rotation measure maps of the radio galaxies and the filaments, we produced masks of the sources from the total intensity image and applied them to the high-frequency RM cube (18 cm + 21 cm + 25 cm). |
To increase the signal-to-noise ratio for the weakest structures, for the analysis we decided to use the RM cube at half resolution. | To increase the signal-to-noise ratio for the weakest structures, for the analysis we decided to use the RM cube at half resolution. |
The RM value for each pixel within a source was obtained by fitting a Gaussian profile to the observed RM distribution. | The RM value for each pixel within a source was obtained by fitting a Gaussian profile to the observed RM distribution. |
For the brightest sources in the field (Double, Goldfish, and TRG) the instrumental polarization is higher than the thermal noise. | For the brightest sources in the field (Double, Goldfish, and TRG) the instrumental polarization is higher than the thermal noise. |
For the Double radio galaxy, for example, the instrumental polarization is ~80—90 ywJy. | For the Double radio galaxy, for example, the instrumental polarization is $\sim80-90~\mu$ Jy. |
Therefore, as detection limit we have chosen 100 yJy (1007) for these sources, and 50 uJy (5c) for the others. | Therefore, as detection limit we have chosen 100 $\mu$ Jy $\sigma$ ) for these sources, and 50 $\mu$ Jy $\sigma$ ) for the others. |
The Faraday spectra of the radio galaxies and the filaments have different levels of complexity. | The Faraday spectra of the radio galaxies and the filaments have different levels of complexity. |
We show this property in Fig. 11.. | We show this property in Fig. \ref{spectraradiogalaxies_I}. |
In each panel we present the total intensity image of one radio galaxy, including a few examples of Faraday spectra extracted at the specified positions (A and B). | In each panel we present the total intensity image of one radio galaxy, including a few examples of Faraday spectra extracted at the specified positions (A and B). |
The profiles in Faraday space along one direction within the source are also given. | The profiles in Faraday space along one direction within the source are also given. |
From this image it is evident that the sources that lie in projection near the cluster center (Double, TRG, and Goldfish) have Faraday spectra characterized by one main peak at a specific Faraday depth, plus significant secondary peaks (above 507) at different Faraday depths. | From this image it is evident that the sources that lie in projection near the cluster center (Double, TRG, and Goldfish) have Faraday spectra characterized by one main peak at a specific Faraday depth, plus significant secondary peaks (above $\sigma$ ) at different Faraday depths. |
This property reflects on the complexity of the RM distributions of these radio galaxies, that are characterized by a complex and non-Gaussian profile (see Fig. 14)). | This property reflects on the complexity of the RM distributions of these radio galaxies, that are characterized by a complex and non-Gaussian profile (see Fig. \ref{histogramscentralradiogalaxies}) ). |
On the other hand, the radio galaxies which lie at large projected distance from the cluster center (Bean, Embryo, and Beaver in Fig. 11)) | On the other hand, the radio galaxies which lie at large projected distance from the cluster center (Bean, Embryo, and Beaver in Fig. \ref{spectraradiogalaxies_I}) ) |
and the radio filaments (see Fig. 15)) | and the radio filaments (see Fig. \ref{spectrafilaments}) ) |
show Faraday spectra with only one significant peak. | show Faraday spectra with only one significant peak. |
For this reason the RM images could only be produced for the three external radio galaxies and for the 3 filaments. | For this reason the RM images could only be produced for the three external radio galaxies and for the 3 filaments. |
Figure 16 presents the results. | Figure \ref{rmimagescomposite} presents the results. |
The RM images have a resolution of 28"x 30". | The RM images have a resolution of $28{\arcsec} \times 30{\arcsec}$ . |
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