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T1ο clean componcuts of this nuage were hen subtracted in the wv-plaue by using the AIPS task UVSUD. | The clean components of this image were then subtracted in the uv-plane by using the AIPS task UVSUB. |
The image with the discrete sources subtracted confirms the xesence of a low-surface briehtuess radio alo at the chster center connected to a brighter pach of xxiplieral enission to the south-east. | The image with the discrete sources subtracted confirms the presence of a low-surface brightness radio halo at the cluster center connected to a brighter patch of peripheral emission to the south-east. |
Iu the botOl xuel of Fig. | In the bottom panel of Fig. |
3 we show a zoo «f the "Mad custer iu wluch the total inteusitv radio couQUIS. after subraction of discrete sources, are overlaid on the red tage of he Skxu Digital Sky Survey (left) aud ou the NAIA N-rav image (vielit). | \ref{radio_sub} we show a zoom of the “Main” cluster in which the total intensity radio contours, after subtraction of discrete sources, are overlaid on the red image of the Sloan Digital Sky Survey (left) and on the XMM X-ray image (right). |
In this figure it is ftither clear that the cloneation of the radio cussion versus west Is located iu οςjucideuce with the ~Subchister”. | In this figure it is further clear that the elongation of the radio emission versus west is located in coincidence with the “Subcluster”. |
We note that this feature is visible at 1100 MIIZ in the D configuration data-set only. whie iu the higher resolution nuages presented hero. and in f16 FIRST image (Becker. White IIelfaud. 1f395). any discrete source seenis to be present. | We note that this feature is visible at 1400 MHz in the D configuration data-set only, while in the higher resolution images presented here, and in the FIRST image (Becker, White Helfand 1995), any discrete source seems to be present. |
Ou the otrer haud. di this location a source. classified as ao discrete soleo. Is detected iun CAIRT nuages at 610 MITz (Venturi ct al. | On the other hand, in this location a source, classified as a discrete source, is detected in GMRT images at 610 MHz (Venturi et al. |
2008) and at 327 \ITIz (Cüaciutucci et al. | 2008) and at 327 MHz (Giacintucci et al. |
2nO). | 2010). |
We uote that. hints of possible diffuse emission (indicated by arrow" in the top panel of Fig. 3)). | We note that, hints of possible diffuse emission (indicated by arrows in the top panel of Fig. \ref{radio_sub}) ), |
left after the subtractio1 process. are preseut on east of the πλain” cluster. | left after the subtraction process, are present on east of the “Main” cluster. |
The: diffuse radio enuüssionus nüeht trace the process of a large-scale structure formation. where cosmic shocks origirated by complex merger eveuts are able to amplity mag1Cic fields aud accelerate svuchrotron clectrous along a €Tister filament. although a deeper observation is required to coufirm the. presence. of these faiut eiisslous. | These diffuse radio emissions might trace the process of a large-scale structure formation, where cosmic shocks originated by complex merger events are able to amplify magnetic fields and accelerate synchrotron electrons along a cluster filament, although a deeper observation is required to confirm the presence of these faint emissions. |
Iu Fig. | In Fig. |
Lo we show the radio iso-coutours of ATSL at 325 Mz. | \ref{radio2} we show the radio iso-contours of A781 at 325 MHz. |
This nuage has been obtained by combining the VLA data in D. C. and D configuration. | This image has been obtained by combining the VLA data in B, C, and D configuration. |
The resulting iuage hasa EWIIA beau of «995.00. | The resulting image hasa FWHM beam of $\times$. |
The VLA nuage has an aneuleuw resolution similar to the tapered GAIRT image at 327 AMIIz presented by Cüaciutucci et al. ( | The VLA image has an angular resolution similar to the tapered GMRT image at 327 MHz presented by Giacintucci et al. ( |
2010). and displays he same structures, | 2010), and displays the same structures. |
Most of the features Visible at 1100 MIIz are also preseut at 325 ALII. | Most of the features visible at 1400 MHz are also present at 325 MHz. |
In particular. the soutli-east peripheral patch aie all the discrete sources. with the ouly exception of the source IT. are clearly detected. | In particular, the south-east peripheral patch and all the discrete sources, with the only exception of the source H, are clearly detected. |
There is a lint of «iffuse cuuission on the right of source D. although most of tlhe* radio halo eiissiou visible at 1100 AIIIz is nüssius at 325 MIIEz. likely )ecause of the lower sensitivity of he low-frequency image. | There is a hint of diffuse emission on the right of source D, although most of the radio halo emission visible at 1400 MHz is missing at 325 MHz, likely because of the lower sensitivity of the low-frequency image. |
The only feature which appears to be comparatively brighter at 325 MIIz is the emission iu coincidence with the "Subceluste. | The only feature which appears to be comparatively brighter at 325 MHz is the emission in coincidence with the “Subcluster”. |
As we will «ee iu the next Section. this feature is characterized NX oa very steep radio spectra. | As we will see in the next Section, this feature is characterized by a very steep radio spectrum. |
By excluding the peripheral patch. the radio halo has a flux deusitv at LlO0 MITzZ of Syjoo,20.5. iid. | By excluding the peripheral patch, the radio halo has a flux density at 1400 MHz of $_{\rm 1400 MHz}\simeq$ 20.5 mJy. |
In lie sune area of about 1Mjpe?. the upper limit to the flux cleusity at 325 AIIIZ is Ssosagyeπλ iuJw. | In the same area of about $^2$, the upper limit to the flux density at 325 MHz is $_{325 MHz}<$ 137 mJy. |
This nut has been calcuated by considering that the surface xiehtuess of the halo. in the primary beam corrected nuage. is evervwwelherc| Jower than the 36 noise level (i.c. 6.6 wJy +). | This limit has been calculated by considering that the surface brightness of the halo, in the primary beam corrected image, is everywhere lower than the $\sigma$ noise level (i.e. 6.6 mJy $^{-1}$ ). |
Thus. we derive an upper Πιτ to the elobal halo spectral index of oj; «1.3. | Thus, we derive an upper limit to the global halo spectral index of $\alpha_{tot}<$ 1.3. |
Iu this section we prescut he spectral iudex nuage of ATSL between 325 IIIz iux 1100 MIIz. | In this section we present the spectral index image of A781 between 325 MHz and 1400 MHz. |
Iu order to be xoperlv compared. the two nuages have been corrected or the primary beam atteation of the VLA anuteuuas. reeridded to the sanue οσοuetry. and convolved to a common resohtion of «5: | In order to be properly compared, the two images have been corrected for the primary beam attenuation of the VLA antennas, regridded to the same geometry, and convolved to a common resolution of $\times$. |
We stress. however. that he original resoluticmus of the two images were already very close not just due to he tapering of the uv-data mut because o| the very siluilar intrinsic coverage of he relevant spatial freueucies. | We stress, however, that the original resolutions of the two images were already very close not just due to the tapering of the uv-data but because of the very similar intrinsic coverage of the relevant spatial frequencies. |
We do not preseit thesλοςral nmCX nuage ater subtractio1 of discrete sources because of he low seuxitiviY of the 325 MITz nage. | We do not present the spectral index image after subtraction of discrete sources because of the low sensitivity of the 325 MHz image. |
Iu the top panel of Fie. 5.. | In the top panel of Fig. \ref{spix}, |
WO presen the s)ectral lucex (left) and t1ο spectral index 1necrtainty (right) nuages between 32M5 zud 1100 ΑΠ with the 110) MIIZ radio lso-contours (primary bean correced) overlaid. | we present the spectral index (left) and the spectral index uncertainty (right) images between 325 and 1400 MHz with the 1400 MHz radio iso-contours (primary beam corrected) overlaid. |
They are calcilated. «ulv on those pixels whose brightucss is above the 36 level at both frequencies. | They are calculated only on those pixels whose brightness is above the $3\sigma$ level at both frequencies. |
The s)ectral index values ranges |)etwveeen a 20.5 aud à 2. while the COYTCSDOlLliug οτον are in the range 20.02 0.25. | The spectral index values ranges between $\alpha \simeq$ 0.5 and $\alpha \simeq$ 2, while the corresponding errors are in the range $\simeq$ $-$ 0.25. |
The discrete sources have a typical spectral index value of ac06 0.7. oulv the source F has a steep spectrin with a 1.00.2. | The discrete sources have a typical spectral index value of $\alpha\simeq 0.6-0.7$ , only the source F has a steep spectrum with $\alpha$ $\pm$ 0.2. |
Iu the bottom panel of Fig. | In the bottom panel of Fig. |
5 we show | \ref{spix} we show |
~5.QUM. | $\sim 5 \times 10^8
\msun$. |
On the other hand a groupof ~LO"AL. can only be the result of a pair of subhalos of 5«.10"AL. in our simulation. | On the other hand a groupof $\sim 10^7 \msun$ can only be the result of a pair of subhalos of $5 \times 10^6 \msun$ in our simulation. |
Our limited resolution also prevents us from quantifying the mass function inside the groups. | Our limited resolution also prevents us from quantifying the mass function inside the groups. |
Nevertheless. and for our largest groups we find that these are dominated by a few massive subhalos and many small ones. | Nevertheless, and for our largest groups we find that these are dominated by a few massive subhalos and many small ones. |
The group infall that we have been detecting in our simulation may well be related to the ghostly streams reported by Lynden-Bell&Lynden-Bell (1995). | The group infall that we have been detecting in our simulation may well be related to the ghostly streams reported by \citet{lyndenbell95}. |
The presence of satellites (dwarf galaxies and globular clusters) sharing a common orbital plane seems rather plausible in the context discussed here. | The presence of satellites (dwarf galaxies and globular clusters) sharing a common orbital plane seems rather plausible in the context discussed here. |
Instead of the disruption of a large progenitor (Lynden-Bell&Lynden-Bell1995). or the tidal formation of satellites within gas-rich major mergers (Kroupa 1997).. we would be witnessing the disruption by the tidal tield of the Milky Way of a "sub-group-size object composed by dwarf galaxies. | Instead of the disruption of a large progenitor \citep{lyndenbell95} or the tidal formation of satellites within gas-rich major mergers \citep{kroupa97}, we would be witnessing the disruption by the tidal field of the Milky Way of a “sub-group”-size object composed by dwarf galaxies. |
The possible implications of this finding are discussed in the Conclusions. | The possible implications of this finding are discussed in the Conclusions. |
The present-time distribution of angular momentum orientations of subhalos reflects both the anisotropy of the accretion pattern and the dynamical processes that affect subhalos while orbiting he MW-like halo. | The present-time distribution of angular momentum orientations of subhalos reflects both the anisotropy of the accretion pattern and the dynamical processes that affect subhalos while orbiting the MW-like halo. |
Fig. | Fig. |
6 shows the orientation of the angular momentum of subhalos accreted in the last 13 snapshots. from he present-day (top left) τος~1.08 (bottom left). | \ref{dist_of_angularmoment_orientation_accminus1} shows the orientation of the angular momentum of subhalos accreted in the last 13 snapshots, from the present-day (top left) to $z \sim 1.08$ (bottom left). |
Here the angular momentum is calculated using the position and velocity of a subhalo in the simulation box frame right before it was accreted. | Here the angular momentum is calculated using the position and velocity of a subhalo in the simulation box frame right before it was accreted. |
ote that only a fraction of these subhalos will have survived until he present-time. | Note that only a fraction of these subhalos will have survived until the present-time. |
The small scale clustering visible in this figure once again highlights the group infall. | The small scale clustering visible in this figure once again highlights the group infall. |
Note the presence of larger-scale patterns lasting over several snapshots (in particular in the op row. which corresponds to the last 2.4 Gyr). | Note the presence of larger-scale patterns lasting over several snapshots (in particular in the top row, which corresponds to the last 2.4 Gyr). |
This presumably implies that the infall patterns are related with persistent larger scale structures (filaments) in the tidal field. | This presumably implies that the infall patterns are related with persistent larger scale structures (filaments) in the tidal field. |
To understand this in more detail. we proceed to trace the evolution of the tidal field around the main halo in our simulation. | To understand this in more detail, we proceed to trace the evolution of the tidal field around the main halo in our simulation. |
To this end we select "field" particles. i.e. those that do not belong to the FOF group of the Milky Way-like halo. | To this end we select “field” particles, i.e. those that do not belong to the FOF group of the Milky Way-like halo. |
The projected spatial distribution of these particles within a 25. Mpe on a side box is shown in grey in Fig. 7.. | The projected spatial distribution of these particles within a $2h^{-1}$ Mpc on a side box is shown in grey in Fig. \ref{tidal_fields}. . |
Like in Fig. 6.. | Like in Fig. \ref{dist_of_angularmoment_orientation_accminus1}, |
each panel corresponds to a different redshift. starting from ο~1.08 in the bottom left panel © the present day in the top left. | each panel corresponds to a different redshift, starting from $z \sim 1.08$ in the bottom left panel to the present day in the top left. |
The distributions of surviving subhalos accreted at the corresponding epoch are overplotted in black. | The distributions of surviving subhalos accreted at the corresponding epoch are overplotted in black. |
Fig. | Fig. |
7 shows that the Milky Way like halo is embedded in a arger-scale tilamentary pattern. | \ref{tidal_fields} shows that the Milky Way like halo is embedded in a larger-scale filamentary pattern. |
These filaments are comparable in extent to the halo itself (as e.g. traced by the accreted subhalos). | These filaments are comparable in extent to the halo itself (as e.g. traced by the accreted subhalos). |
The lumpy nature of the filaments is also clearly visible. showing hat the infall is not a continuous flow. but is in groups as discussed above. | The lumpy nature of the filaments is also clearly visible, showing that the infall is not a continuous flow, but is in groups as discussed above. |
Note that the global orientation of the tidal fields near the main halo has not changed much over the last four snapshots. in agreement with what is observed in the top row of Fig. 6.. | Note that the global orientation of the tidal fields near the main halo has not changed much over the last four snapshots, in agreement with what is observed in the top row of Fig. \ref{dist_of_angularmoment_orientation_accminus1}. |
Furthermore. this large scale pattern is more or less aligned with the major axis of the main halo. shown by the dashed line in each panel (asinBailin&Steinmetz2005). | Furthermore, this large scale pattern is more or less aligned with the major axis of the main halo, shown by the dashed line in each panel \citep[as
in][]{bs05}. |
. Kroupa.Theis&Boily(2005) and have recently argued that the highly anisotropic distribution of TW satellites could not have been drawn from a nearly spherically distributed subhalo population. | \citet{kroupa05} and have recently argued that the highly anisotropic distribution of MW satellites could not have been drawn from a nearly spherically distributed subhalo population. |
Motivated by their claim and the results presented above. we wish to test here under what conditions such a configuration is likely in à ACDM simulation like ours. | Motivated by their claim and the results presented above, we wish to test here under what conditions such a configuration is likely in a $\Lambda$ CDM simulation like ours. |
There are many possible ways to define the degree of fattening of a distribution. | There are many possible ways to define the degree of flattening of a distribution. |
We shall here concentrate on the ollowing two measures: In what follows the positions are defined with respect to the centroid of the satellites tor subhalos). rather than with respect to the centre of the MW(-like) halo. | We shall here concentrate on the following two measures: In what follows the positions are defined with respect to the centroid of the satellites (or subhalos), rather than with respect to the centre of the MW(-like) halo. |
For the first measure (1). we use the inertia tensor defined as Note this inertia tensor differs from that previously used in Eq. €1)) | For the first measure (i), we use the inertia tensor defined as Note this inertia tensor differs from that previously used in Eq. \ref{inertia_tensor_ell_eq}) ) |
in which the positions were normalized by their ellipsoidal distance. | in which the positions were normalized by their ellipsoidal distance. |
Our preference for this new definition is based on the fact that the determination of the ellipsoidal distance is simultaneous to the determination of the eigenvalues of the inertia tensor Z;;. | Our preference for this new definition is based on the fact that the determination of the ellipsoidal distance is simultaneous to the determination of the eigenvalues of the inertia tensor $I_{ij}$ . |
This means that aniterative algorithm is used. in which outliers are successively discarded. until the desired level of convergence is | This means that aniterative algorithm is used, in which outliers are successively discarded, until the desired level of convergence is |
(e.g..Hamannetal.2004).. (e.g..Reynolds1997;Giustinietal.2010). | \citep[e.g.,][]{ham04}. \citep[e.g.,][]{rey97,giustini10}. |
. 5000kms! (Gangulyetal.2001) (Mathur.Elvis&Wilkes1995:Brandt. | $5000\mbox{\ km\ s}^{-1}$ \citep{gan01} \citep{mat95,bra00}. |
LaorWills2000).. X500 (seeHamann&Sabra2004). | $\lesssim 500$ \citep[see][]{hamsab04}. \\citep{tru06}. |
. studies of BALs. since they probe different regions in the quasar surroundings. | studies of BALs, since they probe different regions in the quasar surroundings. |
Although the intrinsic NALs are more difficult to identify. several practical considerations make the derivation of their physical conditions more straightforward than for BALs. (Hamann&Ferland1999). | Although the intrinsic NALs are more difficult to identify, several practical considerations make the derivation of their physical conditions more straightforward than for BALs. \citep{ham99}. |
. Because NALs are often unsaturated and resolved. we can measure directly the NAL coverage fractions and column densities of various ions. | Because NALs are often unsaturated and resolved, we can measure directly the NAL coverage fractions and column densities of various ions. |
In BAL systems. doublet transitions are often blended. makingsuch measurement more difficult. | In BAL systems, doublet transitions are often self--blended, makingsuch measurement more difficult. |
Approximately of all quasars show evidence of outflows (Misawaetal.2007;Ganguly&Brotherton2008). | Approximately of all quasars show evidence of outflows \citep{mis07,gan08}. |
. Although statistical methods can be used to determine this. partial coverage of doublets and multiplets and time variability analysis are the two most commonly used ways to determine decisively if a particular NAL system ts intrinsic. | Although statistical methods can be used to determine this, partial coverage of doublets and multiplets and time variability analysis are the two most commonly used ways to determine decisively if a particular NAL system is intrinsic. |
These signatures are only seen in low ionization transitions of intervening absorbers. and only rarely in the case of small molecular clouds that could be smaller than the projected size of the background quasar broad emission line region. ( | These signatures are only seen in low ionization transitions of intervening absorbers, and only rarely in the case of small molecular clouds that could be smaller than the projected size of the background quasar broad emission line region. ( |
e.g.. Jones et al. 2010.. | e.g., Jones et al. \nocite{jon10}, |
Ivanchik et al. 2010)). | Ivanchik et al. \nocite{iva10}) ). |
Partial coverage is described by the coverage fraction Cy. which is the fraction of photons from the background source that pass through the absorber (Barlow.Hamann.&Sargent 1997). | Partial coverage is described by the coverage fraction $C_{\rm f}$ , which is the fraction of photons from the background source that pass through the absorber \citep{bhs97}. |
. It can be estimated using the residual flux ratio. of resonance doublets (e.g..Barlow&Sargent1997;Gangulyetal. 1999). | It can be estimated using the residual flux ratio of resonance doublets \citep[e.g.,][]{barsar97,gan99}. |
. Variability of the absorption lines could be caused by transverse motion of the absorbing material or by changes in its ionization state (Hamann1997;Misawaetal. 2005). | Variability of the absorption lines could be caused by transverse motion of the absorbing material or by changes in its ionization state \citep{ham97,mis05}. |
. Using UV spectra of z€1.5 quasars observed at two different epochs separated by 4-10 years with theTelescope. Wiseetal.(2004) concluded that a minimum of 21% of the AALs are variable. | Using UV spectra of $z\leqslant1.5$ quasars observed at two different epochs separated by 4–10 years with the, \citet{wis04} concluded that a minimum of $21\%$ of the AALs are variable. |
A’ similar conclusion was reached for z—2 quasars by Narayananet (2004). | A similar conclusion was reached for $z\sim 2$ quasars by \citet{nar04}. . |
Early work investigating metal abundances using narrow associatedabsorption lines indicated supersolar metallicities. Le.. Z=Z. | Early work investigating metal abundances using narrow associatedabsorption lines indicated supersolar metallicities, i.e., $Z\geqslant
Z_\odot$ |
he combined elfect ofCoulomb and hadronic losses as well as the higher compressibility of composite CR. plus thermal eas leads to a less ellicient CR. bubble feedback. as will e discussed. in more detail in Section. ??.. | the combined effect of Coulomb and hadronic losses as well as the higher compressibility of composite CR plus thermal gas leads to a less efficient CR bubble feedback, as will be discussed in more detail in Section \ref{Importance}. |
However. even hough this allows the gas to cool somewhat more elficiently owards the centre the amount of stars formed is roughly. he same in the thermal and relativistic cases. | However, even though this allows the gas to cool somewhat more efficiently towards the centre the amount of stars formed is roughly the same in the thermal and relativistic cases. |
For ¢>Olfie. the star formation rate begins to »v more suppressed in the run with CR bubbles. because at this stage Lop starts to be comparable to the local Dy. | For $t > 0.1\,t_{\rm Hubble}$, the star formation rate begins to be more suppressed in the run with CR bubbles, because at this stage $P_{\rm
CR}$ starts to be comparable to the local $P_{\rm th}$. |
In fact. for the run with a=2.1. this transition occurs somewhat earlier. than in the simulation with a a. given steeperthat for a=2.1 the cosmic ray pressure reaches values significantsooner. | In fact, for the run with $\alpha = 2.1$, this transition occurs somewhat earlier than in the simulation with a steeper $\alpha$, given that for $\alpha = 2.1$ the cosmic ray pressure reaches significant values sooner. |
Lt can be seen that [or [omOdmaaae the BILAL starts to oscillate dramatically as à consequence of the presence of a dominant relativistic particle component. | It can be seen that for $t > 0.1t_{\rm Hubble}$ the BHAR starts to oscillate dramatically as a consequence of the presence of a dominant relativistic particle component. |
A similar behaviour has been observed in the case of star formation in small dwarl galaxies in the work by 7). | A similar behaviour has been observed in the case of star formation in small dwarf galaxies in the work by \cite{Jubelgas2007}. |
. When a bubbleis injected into the ICM and its Poy is comparable or higher than the local iy it | When a bubble is injected into the ICM and its $P_{\rm CR}$ is comparable or higher than the local $P_{\rm th}$ it |
is consistent with an origin of the broad components well above the stellar surface. | is consistent with an origin of the broad components well above the stellar surface. |
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