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In other words. the mereing of smaller chumps within the short timescale causes quick erowth of the nass of star clusters. | In other words, the merging of smaller clumps within the short timescale causes quick growth of the mass of star clusters. |
As a result. lwpermassive star clusters fori. | As a result, hypermassive star clusters form. |
Iu Fig. 5.. | In Fig. \ref{mass_evolution}, |
we show the evolution of the masses of these wperlassive star clusters (upper panels). aud their distances from the center of the galaxy. (lower pancls). | we show the evolution of the masses of these hypermassive star clusters (upper panels), and their distances from the center of the galaxy (lower panels). |
Tere. the definition of the distance is the smaller of the wo distances frou the two galactic centers. | Here, the definition of the distance is the smaller of the two distances from the two galactic centers. |
The clusters first erow quickly. reaching about a half of the final mass in around 20 Myr. After that. the erowth of the mass slows down. aud the clusters fall to the ceuter of the ealaxv due to the dvuamical friction, | The clusters first grow quickly, reaching about a half of the final mass in around $20~\rm{Myr}$ After that, the growth of the mass slows down, and the clusters fall to the center of the galaxy due to the dynamical friction. |
As a result. the merger remnant has a very compact and luminous core simular to those observed iu a large number of elliptical ealaxies (c.2..Ποσοetal. 2009).. | As a result, the merger remnant has a very compact and luminous core similar to those observed in a large number of elliptical galaxies \citep[e.g.,][]{kor09}. . |
Figue 6— shows the evolution of the SER. aud | Figure \ref{sfr} shows the evolution of the SFR and |
evolution of the diffusive infrared-to-ultraviolet background radiation field. | evolution of the diffusive infrared-to-ultraviolet background radiation field. |
Below. in section 2. based on optical aud X-ray observations we reconstruct and comment 'ivpical electron energy. distribution of FR I. jets. | Below, in section 2, based on optical and X-ray observations we reconstruct and comment `typical' electron energy distribution of FR I jets. |
We estimate different’ photon fields contribuüng to the inverse-Compton scattering of these electrons. and next evaluate and cdiseuss (he resulting fIuxes aid photon break energies in the Thomson regime. | We estimate different photon fields contributing to the inverse-Compton scattering of these electrons, and next evaluate and discuss the resulting fluxes and photon break energies in the Thomson regime. |
Section 3 contains detailed analvsis of the 5-rav. emission from the nearby sources Centaurus A and AL ST. | Section 3 contains detailed analysis of the $\gamma$ -ray emission from the nearby sources Centaurus A and M 87. |
Final conclusions are presented in section 4. | Final conclusions are presented in section 4. |
X-ray. observatory detected: about. LO jets clisplaving the FR I lage scale morphology. | X-ray observatory detected about 10 jets displaying the FR I large scale morphology. |
Thev include the ones in radio galaxies M87 (Marshallοἱal.2002:WilsonandYang 2002).. Centaurus A νταetal.2000.2002).. ο 129 (Llarrisetal.2002a).. 3C erastleetal. 2002).. PINS 0521-365 (Birkinshawοἱal.2002).. 3C* 270mM " Q3).. aM 84 metal.2002b).. 3C 66D (ILucdcastleetal.2001).. D2 n02064-35 and p2 n37 (Worrallοἱal.2001). | They include the ones in radio galaxies M87 \citep{mar02,wil02}, , Centaurus A \citep{kra00,kra02}, 3C 129 \citep{har02a}, 3C 31 \citep{har02}, , PKS 0521-365 \citep{bir02}, 3C 270 \citep{chi03}, M 84 \citep{har02b}, 3C 66B \citep{har01}, B2 0206+35 and B2 0755+37 \citep{wor01}. |
. All of these objects are located relatively with a ranging from 3.4 Mpe (Centaurus A) to more than 300 Mpc (PINS 0521-315 al redshift 2= 0.055). | All of these objects are located relatively nearby, with a distance ranging from $3.4$ Mpc (Centaurus A) to more than $300$ Mpc (PKS 0521-315 at redshift $z = 0.055$ ). |
All of them share many spectral and morphological similarities briefly sumnmarised below. | All of them share many spectral and morphological similarities briefly summarised below. |
X-ray jets in FR I radio sources are usually quite short (projected BN ~]—4kpc). | X-ray jets in FR I radio sources are usually quite short (projected length $\sim 1 - 4$ kpc). |
Typically. (μεν are composed of diffusive knots will a spatial scale MN2~0.1 kpe. although a strong inter-knol emission is also sometimes present. | Typically, they are composed of diffusive knots with a spatial scale $R \sim 0.1$ kpc, although a strong inter-knot emission is also sometimes present. |
X-raw jel corresponds to the radio morphology. anc 1 cases of M. 87. PINS0521-315. ος 66D. ος οἱ and D2 07554-37 observed byLST to the optical one. | X-ray jet morphology corresponds to the radio morphology, and – in cases of M 87, PKS 0521-315, 3C 66B, 3C 31 and B2 0755+37 observed by – to the optical one. |
A noted difference between these picturesare spatial offsets of some knot maxima as nieasured at N-ravs and al racdio/optical frequencies (up to ~0.008 kpe in M 87. ~0.08 kpe in Centaurus A and ~0.2 kpe in 3C 66D). | A noted difference between these picturesare spatial offsets of some knot maxima as measured at X-rays and at radio/optical frequencies (up to $\sim 0.008$ kpc in M 87, $\sim 0.08$ kpc in Centaurus A and $\sim 0.2$ kpc in 3C 66B). |
Also. X-ray jets (knots) seem to be narrower than their racdio/optieal counterparts (M 87. Centamrus A). | Also, X-ray jets (knots) seem to be narrower than their radio/optical counterparts (M 87, Centaurus A). |
Except of the weakest and the smallest objects (3C 270 and M 84). the observed X-ray Iuminosities of the discussed jets are Lyc10—107 erg/s. In most cases. multiwavelength observations of knot regions in FR 1Chandra jets allow one to construct broad band radio-to-X-ray. spectral energy. distribution. | Except of the weakest and the smallest objects (3C 270 and M 84), the observed X-ray luminosities of the discussed jets are $L_X \sim 10^{39} - 10^{42}$ erg/s. In most cases, multiwavelength observations of knot regions in FR I jets allow one to construct broad band radio-to-X-ray spectral energy distribution. |
All of the inferred knot spectra seem {ο be similar. and in addition diffieult to be explained in a framework ol standard jet models. | All of the inferred knot spectra seem to be similar, and in addition difficult to be explained in a framework of standard jet models. |
A spectral index of the radio emission (delined in a way 5,xva where 5, is (he energy flux spectraldensity) is always close to a, 0.0. | A spectral index of the radio emission (defined in a way $S_{\nu} \propto
\nu^{- \alpha}$ , where $S_{\nu}$ is the energy flux spectraldensity) is always close to $\alpha_R \sim 0.6$ . |
The opticalemission | The opticalemission |
provide a weighted mean column density at the position of the GRB of 5.8x1030cm~? (Kalberla et al. | provide a weighted mean column density at the position of the GRB of $5.8\times 10^{20}\cmdue$ (Kalberla et al. |
2005) and 6.5x10??cm? (Dickey Lockman 1990). | 2005) and $6.5\times 10^{20}\cmdue$ (Dickey Lockman 1990). |
To fit the absorption component we adopted the Wilms et al. ( | To fit the absorption component we adopted the Wilms et al. ( |
2000) model within the X-ray fitting program XSPEC) and abundance pattern. | 2000) model within the X–ray fitting program XSPEC) and abundance pattern. |
To check the Galactic absorption we fit the data of a closeby (8' from the GRB) AGN (R.A.: 19h35m57.5 Dec.: +78d21m27.0) with a power law model. | To check the Galactic absorption we fit the data of a closeby $8'$ from the GRB) AGN (R.A.: 19h35m57.5 Dec.: +78d21m27.0) with a power law model. |
The resulting column density is 9.5*?5x107?cm-? (90% confidence level). | The resulting column density is $9.5^{+4.8}_{-3.8}\times 10^{20}\cmdue$ $90\%$ confidence level). |
This is in line with the predictions by HI maps even if somewhat on the high side. | This is in line with the predictions by HI maps even if somewhat on the high side. |
Despite its small influence, we fix the Galactic contribution to 6.5107°cm7?. | Despite its small influence, we fix the Galactic contribution to $6.5\times 10^{20}\cmdue$. |
We perform the first fit grouping the RGS data to 20 counts per spectral bin and the pn data to 50 counts per bin in order to have a good tracing of the continuum at high energies. | We perform the first fit grouping the RGS data to 20 counts per spectral bin and the pn data to 50 counts per bin in order to have a good tracing of the continuum at high energies. |
A fit with the metallicity fixed to the solar values gives an intrinsic column density at z—0.54 of Ng(z)=2.3404x10?!cm? with a x?=1288.6 with 1262 degrees of freedom (dof, i.e. a reduced x2,,= 1.02) and 29% null hypothesis probability (nhp). | A fit with the metallicity fixed to the solar values gives an intrinsic column density at $z=0.54$ of $N_H(z)=2.3^{+0.1}_{-0.1}\times 10^{21}\cmdue$ with a $\chi^2=1288.6$ with 1262 degrees of freedom (dof, i.e. a reduced $\chi_{\rm red}^2=1.02$ ) and $29\%$ null hypothesis probability (nhp). |
The power law photon index is T=1.91*002 (see Table 2). | The power law photon index is $\Gamma=1.91^{+0.02}_{-0.02}$ (see Table 2). |
We also checked that the column density did not decrease by splitting into two parts the entire observation with almost the same number of counts. | We also checked that the column density did not decrease by splitting into two parts the entire observation with almost the same number of counts. |
Given the relatively large number of counts collected we leave free to vary the absorbing medium metallicity (keeping the solar abundance pattern). | Given the relatively large number of counts collected we leave free to vary the absorbing medium metallicity (keeping the solar abundance pattern). |
The fit do not improve significantly with x?=1286.8 with 1261 dof (x2,= 1.02) and 30% nhp. | The fit do not improve significantly with $\chi^2=1286.8$ with 1261 dof $\chi_{\rm red}^2=1.02$ ) and $30\%$ nhp. |
An error search shows that the metallicity of the absorbing medium is 0.279 (see Fig. | An error search shows that the metallicity of the absorbing medium is $\gsim 0.2\,Z_\odot$ (see Fig. |
2). | 2). |
The intrinsic column density is Λη2%(2)=5.0733x10?!cm7~?. | The intrinsic column density is $N_H(z)=5.0^{+2.3}_{-3.0}\times 10^{21}\cmdue$. |
For the lower limiting metallicity the column density is <7.7x10?!cm?. | For the lower limiting metallicity the column density is $<7.7\times 10^{21}\cmdue$. |
The number of counts is not high enough to afford an element by element analysis. | The number of counts is not high enough to afford an element by element analysis. |
To circumvent this we leave free the abundance of the single elements, one at a time, keeping the others all tied together. | To circumvent this we leave free the abundance of the single elements, one at a time, keeping the others all tied together. |
For many of them we just find upper limits (see Table 3). | For many of them we just find upper limits (see Table 3). |
In particular, a few elements show an abundance not consistent with zero, namely Ne, S and Fe. | In particular, a few elements show an abundance not consistent with zero, namely Ne, S and Fe. |
Motivated by this analysis we fit the same spectra leaving free the abundances of Ne, S and Fe and linking together all the other elements (mainly driven by Oxygen). | Motivated by this analysis we fit the same spectra leaving free the abundances of Ne, S and Fe and linking together all the other elements (mainly driven by Oxygen). |
The fit improves with a x?=1263.1 with 1258 dof (x2,= 1.00) and 4596 nhp. | The fit improves with a $\chi^2=1263.1$ with 1258 dof $\chi_{\rm red}^2=1.00$ ) and $45\%$ nhp. |
Formally, an F-test gives a random probability for the improvement of 3.3x10”?, equivalent to a significance of 4.16. | Formally, an F-test gives a random probability for the improvement of $3.3\times 10^{-5}$, equivalent to a significance of $4.1\,\sigma$. |
The new fit provides a column density of Ny=1.2768x1033cm"? and just an upper limit on the metallicity of the remaining elements of <0.6Zo. | The new fit provides a column density of $N_H=1.2^{+0.6}_{-0.8}\times 10^{22}\cmdue$ and just an upper limit on the metallicity of the remaining elements of $<0.6\,Z_\odot$. |
The abundances of the three elements left free to vary together lie within Ne=0.4—1.8Neo, 8=1.5—19S6, and Fe«0.3Feo, respectively. | The abundances of the three elements left free to vary together lie within ${\rm Ne}=0.4-1.8\,{\rm Ne}_\odot$, ${\rm S}=1.5-19\,{\rm S}_\odot$, and ${\rm Fe}<0.3\,{\rm Fe}_\odot$, respectively. |
A contour plot of S and Ne abundances is shown in Fig. | A contour plot of S and Ne abundances is shown in Fig. |
3. | 3. |
We also searched for structures due to absorption of different ionization stages and/or to dust around the edges of the above elements at the known GRB redshift of z—0.54, finding none. | We also searched for structures due to absorption of different ionization stages and/or to dust around the edges of the above elements at the known GRB redshift of $z=0.54$, finding none. |
Given the presence of possible uncertainties in the Galactic column density determination, in the adopted solar abundance pattern we left free to vary the Galactic column density in a +50% range around the selected value. | Given the presence of possible uncertainties in the Galactic column density determination, in the adopted solar abundance pattern we left free to vary the Galactic column density in a $\pm 50\%$ range around the selected value. |
We then repeat the fit with a constant metallicity and the fit with Ne, S and Fe free to vary. | We then repeat the fit with a constant metallicity and the fit with Ne, S and Fe free to vary. |
The fit with constant metallicity is characterized by a large Galactic column density (at the edge of the allowed distribution) and with a x2,q=1.02 (1260 dof). | The fit with constant metallicity is characterized by a large Galactic column density (at the edge of the allowed distribution) and with a $\chi_{\rm red}^2=1.02$ (1260 dof). |
The fit with variable Ne, S and Fe abundances gives a x2,=1.00 (1257 dof). | The fit with variable Ne, S and Fe abundances gives a $\chi_{\rm red}^2=1.00$ (1257 dof). |
In this case the F-test gives a a random probability for the improvement of 6.0x10~*, equivalent to a significance of 3.40. | In this case the F-test gives a a random probability for the improvement of $6.0\times 10^{-4}$, equivalent to a significance of $3.4\,\sigma$. |
We also search for possible emission lines. | We also search for possible emission lines. |
We first search for an iron line looking for an unresolved emission line in the 6.4-6.9 keV (rest frame) energy interval. | We first search for an iron line looking for an unresolved emission line in the 6.4–6.9 keV (rest frame) energy interval. |
This range is covered by pn data only. | This range is covered by pn data only. |
We do not find any significant line in this range and we are able to put an upper limit (90% confidence) of 37 eV to any line. | We do not find any significant line in this range and we are able to put an upper limit $90\%$ confidence) of 37 eV to any line. |
We also inspect the RGS data. | We also inspect the RGS data. |
One of the strongest (narrow) emission lines is OVIII at 654 eV (see, e.g., Bertoneet al. | One of the strongest (narrow) emission lines is OVIII at 654 eV (see, e.g., Bertoneet al. |
2010 for the list of line searched). | 2010 for the list of line searched). |
We find a hint for the presence of this line at an energy of 432 eV (665 eV rest frame) even if its significance is between 90% and 99% (see Fig. | We find a hint for the presence of this line at an energy of 432 eV (665 eV rest frame) even if its significance is between $90\%$ and $99\%$ (see Fig. |
4). | 4). |
The line equivalent width is 4.8 eV. The velocity difference of the line energy with respect to the quoted redshift is on the blue side of the spectrum, possibly indicating an outflow, but amounts to 5,000+650 km s~!, which seems somewhat high. | The line equivalent width is 4.8 eV. The velocity difference of the line energy with respect to the quoted redshift is on the blue side of the spectrum, possibly indicating an outflow, but amounts to $5,000\pm650$ km $^{-1}$ , which seems somewhat high. |
significantly lower FoAL compared to the optimal survey (151. compared to 328 for 10.000 fibres). and so the optimum. would be a value around. 3000-4000 fibres. | significantly lower FoM compared to the optimal survey (151, compared to 328 for 000 fibres), and so the optimum would be a value around 3000-4000 fibres. |
Note that if we were to include the elfect of reconstructing the power spectra on small scales (Eisenstein et al. | Note that if we were to include the effect of reconstructing the power spectra on small scales (Eisenstein et al. |
2006b). this may increase he required. number of fibres. | 2006b), this may increase the required number of fibres. |
Reconstruction increases the »erformance in measuring the BAQOs for any number density. out works especially well at high number densities. and may herefore increase the optimal number density. | Reconstruction increases the performance in measuring the BAOs for any number density, but works especially well at high number densities, and may therefore increase the optimal number density. |
We find that including measurements made by other xwvon oscillation and supernovae surveys in the total error analysis does not change the optimal survey. indicating that he analysis we do now should remain valid into the medium "uture. | We find that including measurements made by other baryon oscillation and supernovae surveys in the total error analysis does not change the optimal survey, indicating that the analysis we do now should remain valid into the medium future. |
Finally. we have demonstrated how the IPSO technique can be applied to “realistic” simulations of redshift surveys. including instrumental limitations such as fibre. number. repositioning overheads and galaxy number counts models. | Finally, we have demonstrated how the IPSO technique can be applied to “realistic” simulations of redshift surveys, including instrumental limitations such as fibre number, repositioning overheads and galaxy number counts models. |
Since the measurement of barvon acoustic oscillations is one method to measure the dark energv. IPSO could. also be applied to the survey instrumental design that use other methods such as weak lensing. integrated Sachs-Wolfe or cluster number counts. or even those with other science goals. | Since the measurement of baryon acoustic oscillations is one method to measure the dark energy, IPSO could also be applied to the survey instrumental design that use other methods such as weak lensing, integrated Sachs-Wolfe or cluster number counts, or even those with other science goals. |
We thank Sam Barden. Arjun Dex. Daniel Eisenstein. Andrew AleCrath. ane the rest. of WEALOS team A or helpful advice. ancl comments. | We thank Sam Barden, Arjun Dey, Daniel Eisenstein, Andrew McGrath, and the rest of WFMOS team A for helpful advice and comments. |
DP acknowledges wlplul comments from Pia Mukherjee. | DP acknowledges helpful comments from Pia Mukherjee. |
CB acknowledges unding from the Izaak Walton Willam Memorial. Fund or Advanced Studies anc the Canadian: Institute. for "Theoretical. Astrophysics. | CB acknowledges funding from the Izaak Walton Killam Memorial Fund for Advanced Studies and the Canadian Institute for Theoretical Astrophysics. |
δν is supported by. the Swiss E. | MK is supported by the Swiss NSF. |
DP is supported. by PPARC. | DP is supported by PPARC. |
RCN thanks the EW or support via a Marie Curic Chair. | RCN thanks the EU for support via a Marie Curie Chair. |
We thank Gomini for 'unding part of this work via the WEMOS feasibility stuc. | We thank Gemini for funding part of this work via the WFMOS feasibility study. |
We acknowledge the use of multiprocessor machines at the LCG. University of Portsmouth. | We acknowledge the use of multiprocessor machines at the ICG, University of Portsmouth. |
FE 26 keV. and LO d 1| keV. forthe ο, 190, and ?"Ne ines. respectively. | $\pm$ 26 keV, and 40 $\pm$ 10 keV forthe $^{12}$ C, $^{16}$ O, and $^{20}$ Ne lines, respectively. |
These are ~3% of the line energies. | These are $\sim$ of the line energies. |
Earlier caleulatious (ο,ο, ADurphny(1985))) sugeested vidt sof ~ aud for the carbou liue produce by direct excitation and spallation. respectively, roni proton interactions. | Earlier calculations (e.g. \citet{murphy85}) ) suggested widths of $\sim$ and for the carbon line produced by direct excitation and spallation, respectively, from proton interactions. |
The widths are expectedà to be even broader if accelerated a particles lav a significaut role iu line production (Ramaty&Craunecll1976:Lang&Werntz 1991). | The widths are expected to be even broader if accelerated $\alpha$ particles play a significant role in line production \citep{ram76,lang91}. |
. We rave caleulated thie shape of the 1?C de-excitation line for diffevent interacting particle distributions based on the recent work of Iicner.deSéréville.&Tatiscleff (2001).. | We have calculated the shape of the $^{12}$ C de-excitation line for different interacting particle distributions based on the recent work of \citet{kiener01}. . |
Line shapes were calculated for the following auguar distiibutious: dowuwzud beam (57 exponeutial falloff): fau bea following πιο where 0 is the angle between the particle velocity aud sun center (Murphy.Ίνο-zlovskv.&Ramaty 1988): downward isotropic: and upward isotropic. | Line shapes were calculated for the following angular distributions: downward beam $^{\circ}$ exponential fall-off); fan beam following $^{6} \theta$ where $\theta$ is the angle between the particle velocity and sun center \citep{murphy88}; downward isotropic; and upward isotropic. |
These caleulatious were performed for flares at the heliocentric angles of the five erouped$A/M observations. | These calculations were performed for flares at the heliocentric angles of the five grouped observations. |
From a conrparison of the fluxes in the 1.63 MeV. (22Ne} iid 6.13 MeV (290) lines for the flares grouped by heliocentric angle. we obtained power-law spectral indices from 7.9 to 1.5 for accelerated. particles having inipulsive-fizre abTuidanices aud au ambient Ne/O ratio of 0.25 (Ramaty.Maudzhavidze. | From a comparison of the fluxes in the 1.63 MeV $^{20}$ Ne) and 6.13 MeV $^{16}$ O) lines for the flares grouped by heliocentric angle, we obtained power-law spectral indices from $\sim$ 3.9 to 4.5 for accelerated particles having impulsive-flare abundances and an ambient Ne/O ratio of 0.25 \citep{ram96}. |
&Ilozlovskv— 1996).. In Fieue 17 we plot the calculated line xofiles fior a dowmward isotropic aneular distribution of accelerated particles following a power-law woH.h spectral index [.0 at the five heliocentric augles. | In Figure \ref{fig9} we plot the calculated line profiles for a downward isotropic angular distribution of accelerated particles following a power-law with spectral index 4.0 at the five heliocentric angles. |
The solid. ancl dotted: curves show the shapes for an assume accelerated. à /p ratio of 0.5 and 0.1. respectivelv. | The solid and dotted curves show the shapes for an assumed accelerated $\alpha$ /p ratio of 0.5 and 0.1, respectively. |
It is clear that the a particles plav a doninaut role iu determining the line shapes when the spectra are this steep. | It is clear that the $\alpha$ particles play a dominant role in determining the line shapes when the spectra are this steep. |
We note that the line width produced by a particles is significantly broader than that produced by protolis. | We note that the line width produced by $\alpha$ particles is significantly broader than that produced by protons. |
We then folded the calculated line profiles for the particle distributions and a power-law contiuuuni hnrough the SALAL/CORS iunstiuueut response function and fit the couut data fro 3.8 to Ls MeV. The ος)oduess of fit. measurecl using. the 479 statistic. 1s the probabilitv that a fit to a random distribution of muubers about a mean would eive. a higher. value of--- V7 for the same munber of degrees of freedom. | We then folded the calculated line profiles for the particle distributions and a power-law continuum through the /GRS instrument response function and fit the count data from 3.8 to 4.8 MeV. The goodness of fit, measured using the $\chi ^2$ statistic, is the probability that a fit to a random distribution of numbers about a mean would give a higher value of $\chi ^2$ for the same number of degrees of freedom. |
We found that the derived probabilities were not veorv sensitive to either the iudex of the power law or the à /p ratio. | We found that the derived probabilities were not very sensitive to either the index of the power law or the $\alpha$ /p ratio. |
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