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In [act the anisotropic stress perturbation is crucial to the understanding of evolution of inhomogeneities in the early. radiation dominated universe (27).. | In fact the anisotropic stress perturbation is crucial to the understanding of evolution of inhomogeneities in the early, radiation dominated universe \citep{Hu:1998kj,koivisto:2005}. |
Pherefore an obviously interesting question is whether present observational data could allow for an anisotropic stress perturbation in the late universe which is dominated by the mysterious dark energy [uic (27) | Therefore an obviously interesting question is whether present observational data could allow for an anisotropic stress perturbation in the late universe which is dominated by the mysterious dark energy fluid \citep{ichiki:2007,barrow}. |
The cosmological effects of e due to possible viscosity of clark energy are however quite neglected in the literature. | The cosmological effects of $\sigma$ due to possible viscosity of dark energy are however quite neglected in the literature. |
The main reason for disregarding the anisotropic stress in the dark energy uid might be that conventional dark energy candidates. such as the cosmological constant or canonical scalar fields. are perfect [lüids with @=0. | The main reason for disregarding the anisotropic stress in the dark energy fluid might be that conventional dark energy candidates, such as the cosmological constant or canonical scalar fields, are perfect fluids with $\sigma=0$. |
Llowever. since there is no fundamental theoretical model to describe dark energy. there are no strong reasons tostick to such assumption. | However, since there is no fundamental theoretical model to describe dark energy, there are no strong reasons tostick to such assumption. |
Moreover. coupled. scalar fields have indeed a non-negligible anisotropic stress(2)..and dark energy vector field. candidates (which have been proposed in (??? ))) also have 0zx 0. | Moreover, coupled scalar fields have indeed a non-negligible anisotropic \citep{Schimd:2006pa}, ,and dark energy vector field candidates (which have been proposed in \citep{Armendariz-Picon:2004pm,Kiselev:2004py,Zimdahl:2000zm}) ) also have $\sigma\neq 0$ . |
The observational circumstances are given in Table 1. | The observational circumstances are given in Table 1. |
For the observations with the NOT. the telescope was set at (he comet's rate motion. while sidereal tracking was used [or both WHT and GTC. | For the observations with the NOT, the telescope was set at the comet's rate motion, while sidereal tracking was used for both WHT and GTC. |
In all cases (he images were obtained in service mode bv telescope stall and were reduced by subtracting the overscan level ancl flat fielded using standard procedures, | In all cases the images were obtained in service mode by telescope staff and were reduced by subtracting the overscan level and flat fielded using standard procedures. |
The observations with OSIRIS at the GTC were made with the Sloan ea. and i! filters. | The observations with OSIRIS at the GTC were made with the Sloan $^\prime$ $^\prime$ , and $^\prime$ filters. |
OSIRIS provides a field of view 7.8x 7.8 with a gap of 9.2" in the middle ancl a pixel scale of 0.125"pixel !. | OSIRIS provides a field of view $\arcmin\times$ $\arcmin$ with a gap of $\arcsec$ in the middle and a pixel scale of $\arcsec$ $^{-1}$. |
To increase the signal-to-noise the data were binned in 2x2 pixels. | To increase the signal-to-noise the data were binned in $\times$ 2 pixels. |
The images were calibrated using photometric zeropoints determined [rom standard star observations. | The images were calibrated using photometric zeropoints determined from standard star observations. |
The observations with AC'AM at the WILT were made with the Sloan ο” and x’ filters. | The observations with ACAM at the WHT were made with the Sloan $^\prime$ and $^\prime$ filters. |
ACAM is mounted permanently at a lolcded-Casseerain focus of the telescope and has a circular field of view of 8 diameter with a pixel scale of 0.25 “pixel +. | ACAM is mounted permanently at a folded-Cassegrain focus of the telescope and has a circular field of view of $\arcmin$ diameter with a pixel scale of 0.25 $\arcsec$ $^{-1}$. |
The resulting calibrated WIIT images show brightness levels in agreement with those obtained at GTC within 104. | The resulting calibrated WHT images show brightness levels in agreement with those obtained at GTC within . |
.. The observations with ALFOSC at the NOT were made with standard Johnson-Cousins V and A fillers. | The observations with ALFOSC at the NOT were made with standard Johnson-Cousins $V$ and $R$ filters. |
ALFOSC: provides a field of view of 6.5x6.5 with a pixel scale of 0.19 "pixel!. | ALFOSC provides a field of view of $\arcmin\times$ $\arcmin$ with a pixel scale of 0.19 $\arcsec$ $^{-1}$. |
These images were calibrated with stars in the field of view using magnitudes from the USNO-DI.O catalog (Monetetal.2003).. which provides a photometric accuracy of 20.3 J)mag. | These images were calibrated with stars in the field of view using magnitudes from the USNO-B1.0 catalog \citep{Monet03}, which provides a photometric accuracy of $\sim$ 0.3mag. |
Usine the (transformation equations of Fugukitaetal.(1996) and the magnitude of the Sun in the standard Johnson-Cousins V filler (V.—26.75.Cox2000).. we derive r= 26.96. | Using the transformation equations of \cite{Fugukita96} and the magnitude of the Sun in the standard Johnson-Cousins $V$ filter \citep[$V_\sun$=--26.75,][]{Cox00}, we derive $r^\prime_\sun$ =–26.96. |
Since A. -—27.29. the 7 images of the object obtained al GTC and WIIT were (Gransformed to HR standard Johnson-Cousins magnitudes by subtracting 0.33 mag. where we assume lor the object Che same spectral dependence as lor (he Sun within the bandpasses of these two red filters. | Since $R_\sun$ =–27.29, the $r^\prime$ images of the object obtained at GTC and WHT were transformed to $R$ standard Johnson-Cousins magnitudes by subtracting 0.33 mag, where we assume for the object the same spectral dependence as for the Sun within the bandpasses of these two red filters. |
The resulting GTC. WITT. and NOT 2 magnitudes are consistent with (heir errors and we conclude that the object did not change brightness significantly during our observations. | The resulting GTC, WHT, and NOT $R$ magnitudes are consistent with their errors and we conclude that the object did not change brightness significantly during our observations. |
The available images in each night were combined in order to improve the ratio ancl converted (o solar disk intensity. units appropriate for the analvsis in ternis of dust tail models. | The available images in each night were combined in order to improve the signal-to-noise ratio and converted to solar disk intensity units appropriate for the analysis in terms of dust tail models. |
Figure 1 depicts some representative images obtained on (he different dates with the instruments mentioned above. | Figure 1 depicts some representative images obtained on the different dates with the instruments mentioned above. |
In this figure one can clearly see the inactive nucleus (marked with an arrow) Iving outside the dust cloud that looks like a cometary tail. | In this figure one can clearly see the inactive nucleus (marked with an arrow) lying outside the dust cloud that looks like a cometary tail. |
The absence of a dust cloud surrounding the nucleus indicates Chat the cust. emission has stopped before the observations. | The absence of a dust cloud surrounding the nucleus indicates that the dust emission has stopped before the observations. |
Using aperture photometry. we determined (he nucleusH magnitudes for (he GTC. WIT. and NOT images whenever possible (see Table1). | Using aperture photometry, we determined the nucleus$R$ magnitudes for the GTC, WHT, and NOT images whenever possible (see Table1). |
The aperture size was of 1.5". and the nucleus profiles were stellar-like. | The aperture size was of $\arcsec$ , and the nucleus profiles were stellar-like. |
For the GTC images. | For the GTC images, |
observed. LAEs increases with decreasing redshift. whereas the escape fraction of Lvo emission becomes larger at higher redshift (Nilssonetal.2009;Llavesct2010). | observed LAEs increases with decreasing redshift, whereas the escape fraction of $\alpha$ emission becomes larger at higher redshift \citep{Nilsson2009, Hayes10}. |
.. Lt is certainly interesting ancl important to test whether or not our model can reproduce not only the physical properties of observed. LAIs but also their evolution. | It is certainly interesting and important to test whether or not our model can reproduce not only the physical properties of observed LAEs but also their evolution. |
We will present the properties of simulated: LES atNE redshifts in a forthcoming paper (Shimizu ct al. | We will present the properties of simulated LAEs at various redshifts in a forthcoming paper (Shimizu et al. |
2011 in prep). | 2011 in prep). |
We will also study the epoch of hydrogen reionization using our theoretical model. | We will also study the epoch of hydrogen reionization using our theoretical model. |
We are. grateful. to AL Ouchi. ἂν Matsuda ane T. Uavashino for providing their observational datas. | We are grateful to M. Ouchi, Y. Matsuda and T. Hayashino for providing their observational datas. |
sumerical simulations have been performed with the EUP and PRIMO cluster svstem installed at Institute for he Physics and. Mathematies of the Universe. University of Tokyo. | Numerical simulations have been performed with the EUP and PRIMO cluster system installed at Institute for the Physics and Mathematics of the Universe, University of Tokyo. |
Ες work was partially supported by Cirant-in-Vid or Young Scientists (8) (20674003). | This work was partially supported by Grant-in-Aid for Young Scientists (S) (20674003). |
TO acknowledges the inancial support of Cirant-in-Aid. for Scientific Research (5) (20224002) and of Crant-in-Vid. for Young Scientists (21840015) by JSIPS. | TO acknowledges the financial support of Grant-in-Aid for Scientific Research (S) (20224002) and of Grant-in-Aid for Young Scientists (21840015) by JSPS. |
the data gets better than the case of p=2. | the data gets better than the case of $p=2$. |
Besides. the source density distribution in the nearby universe could be nouunifori. | Besides, the source density distribution in the nearby universe could be nonuniform. |
Since our Galaxy locates inside the Local eroup and Local supercluster. we examine the effect of local overdeusitv of sources on the propagated spectra. | Since our Galaxy locates inside the Local group and Local supercluster, we examine the effect of local overdensity of sources on the propagated spectrum. |
Panel e shows the προς for a local overdensity iu the source umber by a factor of 2 relative to the average within the size of 30 Mpc (Blantonetal.2001). | Panel $c$ shows the spectrum for a local overdensity in the source number by a factor of 2 relative to the average within the size of 30 Mpc \citep{Blanton01}. |
. Due to the attenuation by background photons. higher energv cosmic ravs are mainly contributed bv closer sources. | Due to the attenuation by background photons, higher energy cosmic rays are mainly contributed by closer sources. |
Thus a local excess of the source umber density increases the fux of higher energy cosmic ravs and thus hardens the propagated spectiun at the hiehest energy end. | Thus a local excess of the source number density increases the flux of higher energy cosmic rays and thus hardens the propagated spectrum at the highest energy end. |
We should note that. although the deviation between the propagated spectrmm and observational data ds relatively huge in the wind composition scenario. the situation can also be improved if some of the above effects are taken iuto account. | We should note that, although the deviation between the propagated spectrum and observational data is relatively large in the wind composition scenario, the situation can also be improved if some of the above effects are taken into account. |
Panel d of Fig. | Panel d of Fig. |
À6 prescuts the result after considering a harder injection cosmic-ray spectrin of p=1.6 as well as a local overdensity iu the source nuniber by a factor of 2. | \ref{fig3}
presents the result after considering a harder injection cosmic-ray spectrum of $p=-1.6$ as well as a local overdensity in the source number by a factor of 2. |
The theoretical energy «ροζσπα agrees reasonably well with the observational data. | The theoretical energy spectrum agrees reasonably well with the observational data. |
Iu Fig. Απ. | In Fig. \ref{fig4}, |
we preseut how the average mass umber of cosmic ravs evolves with cucrey for differcut scenarios discussed. above. | we present how the average mass number of cosmic rays evolves with energy for different scenarios discussed above. |
uecded. Oue cau see that. for the lyvpernova ejecta composition scenarios. the average mass number increases eracduallv with euergv. which is consistent with the fiudiug bv PAO that the composition of CITECRs becomes increasingly heavy with cucrey. | One can see that, for the hypernova ejecta composition scenarios, the average mass number increases gradually with energy, which is consistent with the finding by PAO that the composition of UHECRs becomes increasingly heavy with energy. |
In the wind composition scenarios. the average imass nunber increases more slowly with energv. | In the wind composition scenarios, the average mass number increases more slowly with energy. |
However. since the hadronic interaction models at such high euergies are not well understood. such a composition is still cousisteut with the measuremeuts within the uncertainties of theoretical expectations. | However, since the hadronic interaction models at such high energies are not well understood, such a composition is still consistent with the measurements within the uncertainties of theoretical expectations. |
We lave shown that. with some fraction of heavy uuclei such as iron in the source composition. the seni-relativistic hwperuova inodel can explain the energv spectrum and composition of ultra-high euergv cosnic rays as ieasured by PAO. | We have shown that, with some fraction of heavy nuclei such as iron in the source composition, the semi-relativistic hypernova model can explain the energy spectrum and composition of ultra-high energy cosmic rays as measured by PAO. |
The heavy elements may originate from the lapernova ejecta aud are brought ite the forward shock region via the Ravleigh-Tavlor (R-T) instabilitv. | The heavy elements may originate from the hypernova ejecta and are brought into the forward shock region via the Rayleigh-Taylor (R-T) instability. |
R-T fingers are seen to be able to reach orward shock frout or overtake the forward shocks (c.c.Junetal.1996:Blondin&Ellison2001:Bloudinetal.2001:Rakowskict 2011).. so they should be able to bring heavy elements iuto the ambient moediuna. | R-T fingers are seen to be able to reach the forward shock front or overtake the forward shocks \citep[e.g.][]{Jun96, BE01, Blondin01,Rakowski11}, so they should be able to bring heavy elements into the ambient medium. |
Observations of some Calactie supernova remnants such as Cas A and SN 1000,Cloud. have provided direct and indirect evidences for protrusion of heavy elements iu the material bevoud the blast wave2000). | Observations of some Galactic supernova remnants such as Cas A and SN 1006, have provided direct and indirect evidences for protrusion of heavy elements in the material beyond the blast wave. |
. When the ejecta is being decelerated by the swept-aip aüubient wind material reverse shocks are also foriuing. | When the ejecta is being decelerated by the swept-up ambient wind material, reverse shocks are also forming. |
The reverse shock may accelerate particles iu the shocked ejecta, which provides an alternative site where CUE [umeavv uuclei are accelerated. | The reverse shock may accelerate particles in the shocked ejecta, which provides an alternative site where UHE heavy nuclei are accelerated. |
The reverser shock in the ivperuova case ds typically iildlv relativistic with a speed about A44220.5. as shown in the appendix. | The reverser shock in the hypernova case is typically mildly relativistic with a speed about $\beta_{\rm rs}\simeq 0.5$, as shown in the appendix. |
Since ie duternal energv density iu the shocked ejecta aud shocked wind are equal. the magnetic field in the reverse VAwek is Be=(egyερ)7B. whore ep, is the magnetic cheld cnerey equipartition factor iureverse shock aud ep is 1e se factor for the forward shock. | Since the internal energy density in the shocked ejecta and shocked wind are equal, the magnetic field in the reverse shock is $B_{\rm rs}=(\epsilon_{B\rm
r}/\epsilon_B)^{1/2}B$, where $\epsilon_{B\rm r}$ is the magnetic field energy equipartition factor in reverse shock and $\epsilon_B$ is the same factor for the forward shock. |
Then we obtain the Παπ eunerev of particles accelerated by the reverse VArock which can also reach ultra-high energies for heavy nuclei. | Then we obtain the maximum energy of particles accelerated by the reverse shock which can also reach ultra-high energies for heavy nuclei. |
Tu sunuusrv. hwpernovae reninauts are shown to be potential sources of UITECTs. | In summary, hypernovae remnants are shown to be potential sources of UHECRs. |
They have been discovered to be present within the GZI& horizon. 100 Alpe. | They have been discovered to be present within the GZK horizon, $\sim 100$ Mpc. |
Due to the semi-relativistic ejecta produced. by the explosion. they can accelerate particles to ultra-high enereies. | Due to the semi-relativistic ejecta produced by the explosion, they can accelerate particles to ultra-high energies. |
Ibvperuovae naturally provide intermediate lass or heavy elements either in the form of stellar wind or the hwperuova ejecta itself. | Hypernovae naturally provide intermediate mass or heavy elements either in the form of stellar wind or the hypernova ejecta itself. |
Tn this paper. we lave calculated the propagated spectrin aud composition of UIIECTs arriving at the Earth elven au initial cosmic rav spectrum and composition at the sources produced bv liverpnovac. | In this paper, we have calculated the propagated spectrum and composition of UHECRs arriving at the Earth given an initial cosmic ray spectrum and composition at the sources produced by hyerpnovae. |
We fud that with the source cosnic-rav colposition equal to that of the Lywpernova ejecta. the propagated energy spectrum aud composition are colmpatible with the measurements by PAO. | We find that with the source cosmic-ray composition equal to that of the hypernova ejecta, the propagated energy spectrum and composition are compatible with the measurements by PAO. |
We are grateful to the auouvinous referee for the valuable report. aud to A. M. Taylor. P. Mésszios. S. Razzaque aud Z. C. Dai for useful discussions. | We are grateful to the anonymous referee for the valuable report, and to A. M. Taylor, P. Mésszárros, S. Razzaque and Z. G. Dai for useful discussions. |
This work is supported bv the NSFC under grants 10972008. and 11033002. the 973 program uuder eraut 2009CBa?1500. the program ofNCET. aud the Fok Ying Tung Education Foundation lom When the ejecta is being decelerated by the swept-up ambient wind material. reverse shocks are also forming. | This work is supported by the NSFC under grants 10973008 and 11033002, the 973 program under grant 2009CB824800, the program of NCET, and the Fok Ying Tung Education Foundation 1cm When the ejecta is being decelerated by the swept-up ambient wind material, reverse shocks are also forming. |
feature appears to be sharper (e.g. see Fig. | feature appears to be sharper (e.g. see Fig. |
4 and Fig. | 4 and Fig. |
AX very interesting feature can be seen when the [LIare is receding (ὁ2 90°). | A very interesting feature can be seen when the flare is receding $\phi\approx 90^{o}$ ). |
The main redshifted. Εις maximuntr broadens (see Fie. | The main redshifted flux maximum broadens (see Fig. |
4 and Fig. | 4 and Fig. |
6) or even splits into two maxima (see Fig. | 6) or even splits into two maxima (see Fig. |
3. Fig. | 3, Fig. |
5 or Fig. | 5 or Fig. |
7). | 7). |
In order to illustrate this elfect better. appropriate cross-sections through the reverberation maps shown in Fig. | In order to illustrate this effect better, appropriate cross-sections through the reverberation maps shown in Fig. |
3 and Fig. | 3 and Fig. |
4 are given (see Fig. | 4 are given (see Fig. |
9 and Fig. | 9 and Fig. |
10 respectively). | 10 respectively). |
Note that in the case of Fig. | Note that in the case of Fig. |
9 three maxima are seen. | 9 three maxima are seen. |
This thickening or splitting is à purely. eeneral relativistic ellect. | This thickening or splitting is a purely general relativistic effect. |
Phe middle maximum is à consequence of strong bending of light which is focused on the opposite side of the black hole relative to the actual position of the flare ancl [eads to the enhancement of the illumination of the disc in this region. | The middle maximum is a consequence of strong bending of light which is focused on the opposite side of the black hole relative to the actual position of the flare and leads to the enhancement of the illumination of the disc in this region. |
The inlluence of this ellect is even more important. for the flare. located: closer to the black hole (Grp. 6.5m) as shown in Fig. | The influence of this effect is even more important for the flare located closer to the black hole $r_{flare}=6.5m$ ) as shown in Fig. |
7 and Fie. | 7 and Fig. |
11. | 11. |
Fig. | Fig. |
1l shows three distinct. maxima with the middle maximum duc to light focusing being dominant. | 11 shows three distinct maxima with the middle maximum due to light focusing being dominant. |
As in the previous cases. such an οσο could. be observable by future high throughput spectrometers. | As in the previous cases, such an effect could be observable by future high throughput spectrometers. |
The magnitude of the energy. shift. of the main maximum is related. to the inclination of the accretion cise (e.g. compare Fig. | The magnitude of the energy shift of the main maximum is related to the inclination of the accretion disc (e.g. compare Fig. |
3 and Fie. | 3 and Fig. |
4) and also to the distance of the flare [rom the centre. | 4) and also to the distance of the flare from the centre. |
“Phe latter elect is demonstrated. for example in Fie.8. Fig. | The latter effect is demonstrated for example in Fig.3, Fig. |
7 and Fig. | 7 and Fig. |
8 where the bright A-shapecl feature extends gradually to lower energies as the distance μι of the [lare from the centre decreases. | 8 where the bright $\Lambda$ -shaped feature extends gradually to lower energies as the distance $r_{flare}$ of the flare from the centre decreases. |
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