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Fifty X-ray sources were detected. above a Lux limit of ον 1(0.5-2 keV) using the same source detection algorithm. as Hasinger (1998a). | Fifty X-ray sources were detected above a flux limit of $^{-15}$ (0.5-2 keV) using the same source detection algorithm as Hasinger (1998a). |
Of the 42 sources identified. 3 were clusters or groups of galaxies with up to 3 galaxy recshilts per svstem. | Of the 42 sources identified, 3 were clusters or groups of galaxies with up to 3 galaxy redshifts per system. |
None ofthe unidentified sources are coincident with an excess of galaxies on lt band CCD images at Ret. making their identification with clusters or eroups at z« 0.7 very unlikely. | None of the unidentified sources are coincident with an excess of galaxies on R band CCD images at $\lesssim$ 24, making their identification with clusters or groups at $<$ 0.7 very unlikely. |
One of the 3 clusters has no spectroscopic redshift: we estimate a redshift of zzzl.1 based on the magnitude of R=22.1 of the brightest galaxy given bv Zamorani (1999). | One of the 3 clusters has no spectroscopic redshift; we estimate a redshift of $\approx$ 1.1 based on the magnitude of R=22.1 of the brightest galaxy given by Zamorani (1999). |
In the deep (79 ksec) NEP survey of Bower (1996) identifications were made for 18 out of 20 X-ray sources brighter than +! (0.5-2 keV). including one cluster. | In the deep (79 ksec) NEP survey of Bower (1996) identifications were made for 18 out of 20 X-ray sources brighter than $^{-14}$ (0.5-2 keV), including one cluster. |
Phe ciffuse extended source in this. field identified by Bure (1992) as a nearby cluster (NEP Xl) at z=0.088 has also been included. | The diffuse extended source in this field identified by Burg (1992) as a nearby cluster (NEP X1) at z=0.088 has also been included. |
We note that by minimizing the size of the PSPC PSE (by using the hard 0.5-2 keV band) and including the extra exposure of Bower (1996). NEP ΝΤ breaks up into several components. | We note that by minimizing the size of the PSPC PSF (by using the hard 0.5-2 keV band) and including the extra exposure of Bower (1996), NEP X1 breaks up into several components, |
Clearly the small scale spiky structure increases the locating ability of Ες substantially and allows astronomers to search of clectromagnetic counterparts of the GAY source. | Clearly the small scale spiky structure increases the locating ability of PTAs substantially and allows astronomers to search of electromagnetic counterparts of the GW source. |
Thanks to the pulsar timing parallax signal. we are able to include the pulsar term in our scheme to detect non-evolutionary CAV sources. | Thanks to the pulsar timing parallax signal, we are able to include the pulsar term in our scheme to detect non-evolutionary GW sources. |
Our method. first presented in. 2.. was complemented recently by the work of 2 who also include the pulsar term in their analysis pipeline. | Our method, first presented in \cite{KS10}, was complemented recently by the work of \cite{CC10} who also include the pulsar term in their analysis pipeline. |
Their approach to use the pulsar term however differs. as they study the chirped signal and ignore the timing parallax signal. while we use the pulsar term in combination with the timing parallax signal. | Their approach to use the pulsar term however differs, as they study the chirped signal and ignore the timing parallax signal, while we use the pulsar term in combination with the timing parallax signal. |
The pulsar term leads to many potential important applications. | The pulsar term leads to many potential important applications. |
For example. ?. recently. point out that the curvature of the wavefronts introduces a GW parallax elect. which can be used to measure the SMDBIID distance. | For example, \cite{DF10} recently point out that the curvature of the wavefronts introduces a GW parallax effect, which can be used to measure the SMBHB distance. |
Other possible applications are very interesting to study also. | Other possible applications are very interesting to study also. |
For instance. while the pulsar distance measurement helps the GW cletection. he CAV signal. on the other hand. also helps the pulsar distance measurement. | For instance, while the pulsar distance measurement helps the GW detection, the GW signal, on the other hand, also helps the pulsar distance measurement. |
Due to the interference between the pulsar erm and the Earth term. the GW signal is sensitive to the pulsar distance. | Due to the interference between the pulsar term and the Earth term, the GW signal is sensitive to the pulsar distance. |
In this wav. the detected. CAV signal can be used to increase the accuracy of our distance measurements. | In this way, the detected GW signal can be used to increase the accuracy of our distance measurements. |
As we expected. the overall measurement accuracy for pulsar distances increases with GW amplitude. as shown in Figure 5.. | As we expected, the overall measurement accuracy for pulsar distances increases with GW amplitude, as shown in Figure \ref{fig:dmea}. |
Phe improvement in accuracy increases faster with respect to he GAY amplitude for pulsars with smaller distances (2).=0.5 kkpc) than the ones with larger distances (D,=2kkpe). | The improvement in accuracy increases faster with respect to the GW amplitude for pulsars with smaller distances $D_{\rm psr} = 0.5$ kpc) than the ones with larger distances $D_{\rm psr} = 2$ kpc). |
‘This is because the distance measurement usingὃν the timinge parallax has highere precision for near-by pulsars than for lar-away oulsars. | This is because the distance measurement using the timing parallax has higher precision for near-by pulsars than for far-away pulsars. |
A smaller number of multiple distance solutions occurs lor the nearby pulsars. for which the accuracy of distance measurements increases faster when the GW amplitude becomes larger. | A smaller number of multiple distance solutions occurs for the nearby pulsars, for which the accuracy of distance measurements increases faster when the GW amplitude becomes larger. |
Beside pulsar timing parallax techniques. one can also determine pulsar distances from other methods. such as the Very Long Baseline Interferometry (VLBI). pulsar orbital parallax. and relative acceleration (2).. | Beside pulsar timing parallax techniques, one can also determine pulsar distances from other methods, such as the Very Long Baseline Interferometry (VLBI), pulsar orbital parallax, and relative acceleration \citep{STWKS10}. |
Since in this paper we mainly ocus on extracting GW information and pulsar distances from pulsar timing array data. investigation on incorporating these extra information into PTA GW data analysis will be presented in future studies. | Since in this paper we mainly focus on extracting GW information and pulsar distances from pulsar timing array data, investigation on incorporating these extra information into PTA GW data analysis will be presented in future studies. |
Finally. we have demonstrated how the errors in the parameter estimation can be determined reliably. | Finally, we have demonstrated how the errors in the parameter estimation can be determined reliably. |
We have calculated he parameter estimation error using both the CRB and the ZZ bound. | We have calculated the parameter estimation error using both the CRB and the ZZ bound. |
Both of the methods give nearly identical results in he high SNR. region. as we expected. | Both of the methods give nearly identical results in the high SNR region, as we expected. |
But the two methocs deviate from each other when the SNR is low. | But the two methods deviate from each other when the SNR is low. |
The CRB is known o predict. unreachable accuracy in the Iow-SNIU regime. while the ZZ bound gives more trustful results in that region. | The CRB is known to predict unreachable accuracy in the low-SNR regime, while the ZZ bound gives more trustful results in that region. |
For most applications involving CAV detection. the noise contributions dominate above the GW signals. so that the ZZ bound urns out to be superior in reliability compared to the CRB. | For most applications involving GW detection, the noise contributions dominate above the GW signals, so that the ZZ bound turns out to be superior in reliability compared to the CRB. |
We gratefully. acknowledge support from ERC Advanced. Grant “LEAP. Crant Agreement. Number 227947 (PL Michael Ixramoer). | We gratefully acknowledge support from ERC Advanced Grant “LEAP”, Grant Agreement Number 227947 (PI Michael Kramer). |
We thank Fredrick Jenet for his suggestions ancl help. | We thank Fredrick Jenet for his suggestions and help. |
We also thank Alberto Sesana for illuminating discussion and Joris. Verbiest for reading the manuscript ancl his detailed. suggestions. | We also thank Alberto Sesana for illuminating discussion and Joris Verbiest for reading the manuscript and his detailed suggestions. |
We would like to thank Helge Rottmann for the assistance of using the VLBI cluster to perform most of the numerical computation in this paper. | We would like to thank Helge Rottmann for the assistance of using the VLBI cluster to perform most of the numerical computation in this paper. |
We also thank the anonymous referee for helpful comments. | We also thank the anonymous referee for helpful comments. |
We have carefully cousidered the concept of upper μπιτς in the context of uudetected sources. aud have eveloped a rigorous formalisii to undoersaud aud express the concept. | We have carefully considered the concept of upper limits in the context of undetected sources, and have developed a rigorous formalism to understand and express the concept. |
Despite its sccming simplicity. upper nuits are not treated iun a uniforiu faslikon in astronomical literature. leaciug to considerable variatious iu uecaning and value. | Despite its seeming simplicity, upper limits are not treated in a uniform fashion in astronomical literature, leading to considerable variations in meaning and value. |
We formally define an upper luit to the sotree intensity as the miaxiununi mteusitv --- can have without exceeding a specifier detection threshold at a even probability. | We formally define an upper limit to the source intensity as the maximum intensity it can have without exceeding a specified detection threshold at a given probability. |
This is defined by the y.tatistical power of the detection aleoritlini. | This is defined by the statistical power of the detection algorithm. |
This is equivalent to «lefiniug it as the largest source inteusiv hat remains undetected at the specified probabilitv. aid is defined x the probability of Type ID eror. | This is equivalent to defining it as the largest source intensity that remains undetected at the specified probability, and is defined by the probability of Type II error. |
Tliρα oe| the detection probability is computed or a varietv of source intensities. the upper limit is then ideutifixd x determining the intercept of the required probabilitv with this curve. | Thus, if the detection probability is computed for a variety of source intensities, the upper limit is then identified by determining the intercept of the required probability with this curve. |
Thus. an upper limit is depeudoeut ouly on the detection criterion. which is eeuer:uly a fuiction oilv of the backeround. aud independent of the source counts. | Thus, an upper limit is dependent only on the detection criterion, which is generally a function only of the background, and independent of the source counts. |
This is different frou the upper: bound (i.c. the upper edge of a confidence interval). which is obtained when the probability distribution of the souree mtesity is computed eiven that some counts are observed in the putative source region. | This is different from the upper bound (i.e, the upper edge of a confidence interval), which is obtained when the probability distribution of the source intensity is computed given that some counts are observed in the putative source region. |
We distiuguis1 betwec3 the upper bouud of the confidence imterval and the wpper limit of source detectability. | We distinguish between the upper bound of the confidence interval and the upper limit of source detectability. |
Uunlike a coufideuce interval (or Bavesiui credible mterval). au upper uit is a function of the deection procedure alone aux does uot necessarily depend ou the observed source counts. | Unlike a confidence interval (or Bayesian credible interval), an upper limit is a function of the detection procedure alone and does not necessarily depend on the observed source counts. |
The primary goals of this paper are to clearly define au τιoper limit. to sharpen the distiuction between an upper limit aud an upper bound. aud to av out a detailed xocedure to compute the former for amv detection process. | The primary goals of this paper are to clearly define an upper limit, to sharpen the distinction between an upper limit and an upper bound, and to lay out a detailed procedure to compute the former for any detection process. |
Iu particular. we have slow1 low fo colmpute 4»per limits for the simple Poissou case. | In particular, we have shown how to compute upper limits for the simple Poisson case. |
We also provide a step-by-step procedure for deriv1e it when a simplifed significance-based detection method is eniploved. | We also provide a step-by-step procedure for deriving it when a simplified significance-based detection method is employed. |
To extract the most science frou caalogs. we areue for using a consistent. statistically reasonable recipe of an upper Πιτ beiug related to the satistical power of a test. | To extract the most science from catalogs, we argue for using a consistent, statistically reasonable recipe of an upper limit being related to the statistical power of a test. |
In addition. we illustrate the peril of using au upper bound iu place of an upper limit and of only reporting a frequentist confidence iuterval when a source is detected. | In addition, we illustrate the peril of using an upper bound in place of an upper limit and of only reporting a frequentist confidence interval when a source is detected. |
Conversely. includiug cowfidkence bouuds. even for nou-detectious. Is a wav to avoid the Eddington bias aud increase the scieutific useluess of large ccatalogs. | Conversely, including confidence bounds, even for non-detections, is a way to avoid the Eddington bias and increase the scientific usefulness of large catalogs. |
We also describe a general recipe for calctating an upper Iit for auv well-defined detection algoritlan. | We also describe a general recipe for calculating an upper limit for any well-defined detection algorithm. |
Bricfly. the detection threshold should be first defined σος ou an acceptable probability of a false detection (the a-level threshold). aud an intensity that esures that the source will be detected at a specitfed probability (the level detection probability) should be computed: this latter iuensitv is identified with the upper ματ, | Briefly, the detection threshold should be first defined based on an acceptable probability of a false detection (the $\alpha$ -level threshold), and an intensity that ensures that the source will be detected at a specifed probability (the $\beta$ -level detection probability) should be computed; this latter intensity is identified with the upper limit. |
We recommend that when upper linits are reported in the literature. both the corresponding o aud > values should also be reported. | We recommend that when upper limits are reported in the literature, both the corresponding $\alpha$ and $\beta$ values should also be reported. |
This work was supported by NASA-AISRP exaut. NNGOGGELI7C (AC). CXC NASA contract (VLK. AS). NSF erauts DAIS 0106085 and DAIS 09-07522 (DvD). | This work was supported by NASA-AISRP grant NNG06GF17G (AC), CXC NASA contract NAS8-39073 (VLK, AS), NSF grants DMS 04-06085 and DMS 09-07522 (DvD). |
We acknowledge useful discussions with Rick Harucden. Frau Primini. Jeff Scarele. Tom Loredo. Tom Alderoft. Paul Green. ασ Drake. aud waticipauts aud oreanizersOo of tho SAMSI/SaFeDoe ProgramC» on Astrostatistes. | We acknowledge useful discussions with Rick Harnden, Frank Primini, Jeff Scargle, Tom Loredo, Tom Aldcroft, Paul Green, Jeremy Drake, and participants and organizers of the SAMSI/SaFeDe Program on Astrostatistcs. |
The total numbers of -ray emitting radio quasars/FR IIs as functions of sensitivity are plotted in Fig. 2.. | The total numbers of $\gamma$ -ray emitting radio quasars/FR IIs as functions of sensitivity are plotted in Fig. \ref{fig2}. |
About 1200 4- radio quasars will be detected by GLAST based on the EC model, if its sensitivity is 30 times higher than that of EGRET at 100 MeV (Gehrels&Michelson1999). | About 1200 $\gamma$ -ray radio quasars will be detected by GLAST based on the EC model, if its sensitivity is 30 times higher than that of EGRET at 100 MeV \citep[][]{gm99}. |
. Our calculations show that no FR II radio galaxies will be detected by GLAST as y-ray emitters either for EC or SSC models. | Our calculations show that no FR II radio galaxies will be detected by GLAST as $\gamma$ -ray emitters either for EC or SSC models. |
We find that almost all y-ray quasars (~99%)) to be detected by GLAST will be FSRQs for the EC model, and the remainder (~1%)) will be SSRQs. | We find that almost all $\gamma$ -ray quasars $\sim$ ) to be detected by GLAST will be FSRQs for the EC model, and the remainder $\sim$ ) will be SSRQs. |
For the SSC model, ~1800 quasars will be detected by GLAST, of which ~80% will be FSRQs (see Fig. 2)). | For the SSC model, $\sim 1800$ quasars will be detected by GLAST, of which $\sim$ will be FSRQs (see Fig. \ref{fig2}) ). |
We use the derived y-ray LF 8)) | We use the derived $\gamma$ -ray LF (Eq. \ref{gammalf}) ) |
to calculate the contribution of all radio quasars/FR IIs (Eq.to the EGRB (listed in Table 1). | to calculate the contribution of all radio quasars/FR IIs to the EGRB (listed in Table 1). |
1.0cm We find that the redshifts of almost all EGRET BL Lac objects are <1, which implies that BL Lac objects may have different space density and evolutionary behaviors from quasars. | 1.0cm We find that the redshifts of almost all EGRET BL Lac objects are $\la 1$, which implies that BL Lac objects may have different space density and evolutionary behaviors from quasars. |
The LF of BL Lac objects was derived from the DXRBS by Padovanietal.(2007),, however, the results for BL Lac objects are more uncertain than those for FSRQs, because of the small number statistics and ~30% of them having no redshift. | The LF of BL Lac objects was derived from the DXRBS by \citet{pg07}, however, the results for BL Lac objects are more uncertain than those for FSRQs, because of the small number statistics and $\sim$ of them having no redshift. |
In this work, we use a parent radio LF of radio quasars/FR II galaxies derived from the FSRQ LF. | In this work, we use a parent radio LF of radio quasars/FR II galaxies derived from the FSRQ LF. |
The derived redshift distributions of y-ray quasars are similar for different models (EC or SSC), which are roughly consistent with that of the EGRET quasars. | The derived redshift distributions of $\gamma$ -ray quasars are similar for different models (EC or SSC), which are roughly consistent with that of the EGRET quasars. |
It was suggested that the y-ray radiative mechanisms are different for and BL Lac objects, i.e., the EC mechanism may be quasarsresponsible for quasars, while the SSC is for BL Lac objects (e.g.,Dondi&Ghisellini1995). | It was suggested that the $\gamma$ -ray radiative mechanisms are different for quasars and BL Lac objects, i.e., the EC mechanism may be responsible for quasars, while the SSC is for BL Lac objects \citep*[e.g.,][]{dg95}. |
. If the EC mechanism is indeed responsible for y-ray quasars, the predicted y-ray quasars to be detected by GLAST will be ~1200. | If the EC mechanism is indeed responsible for $\gamma$ -ray quasars, the predicted $\gamma$ -ray quasars to be detected by GLAST will be $\sim1200$. |
Our results are roughly consistent with the estimate given by Dermer(2007) based on a simplified blazar model. | Our results are roughly consistent with the estimate given by \citet{d07} based on a simplified blazar model. |
The SSC model predicts a simple relation between -ray luminosity and radio luminosity of the jets (see Eq. 3)). | The SSC model predicts a simple relation between $\gamma$ -ray luminosity and radio luminosity of the jets (see Eq. \ref{l_sscrad}) ). |
The EGRET flux limit vf™Gppr at 100 MeV can be converted to a radio flux density limit f;ra~10 Jy at 5GHz by using Eq. (3)). | The EGRET flux limit ${\nu}f_{\nu,\rm
EGRET}^{\rm min}$ at 100 MeV can be converted to a radio flux density limit $f_{\nu,\rm rad}\sim 10$ Jy at 5GHz by using Eq. \ref{l_sscrad}) ). |
This means that all EGRET quasars should have their radio flux densities higher than ~10 Jy, which is obviously inconsistent with most EGRET quasars having νιZ;1 Jy (e.g.,Steckeretal.1993;Zhou1997). | This means that all EGRET quasars should have their radio flux densities higher than $\sim 10$ Jy, which is obviously inconsistent with most EGRET quasars having $f_{\nu,\rm
rad}\ga 1$ Jy \citep*[e.g.,][]{s93,zhou97}. |
. Thus, the SSC model is unlikely to be responsible for EGRET quasars, unless the physical properties of the jets are significantly different for individual sources, i.e., the values of Cssc for most sources deviate significantly from a constant value. | Thus, the SSC model is unlikely to be responsible for EGRET quasars, unless the physical properties of the jets are significantly different for individual sources, i.e., the values of ${\cal C}_{\rm SSC}$ for most sources deviate significantly from a constant value. |
Our results show that no FR II galaxies will be detected by GLAST. | Our results show that no FR II galaxies will be detected by GLAST. |
Most GLAST quasars will be FSRQs (~99% for the EC model), which implies that FSRQs will still be good candidates for identifying the y-ray sources even for | Most GLAST quasars will be FSRQs $\sim$ for the EC model), which implies that FSRQs will still be good candidates for identifying the $\gamma$ -ray sources even for |
and in agreement with the fairly circularly-symmetric ray contours centered on the BCG (seeMann&Ebeling 2011). | and in agreement with the fairly circularly-symmetric X-ray contours centered on the BCG \citep[see][]{MannEbeling2011}. |
. The discovery of a high-redshift galaxy in the field of MACS0329 adds to several known high-reshift galaxies lensed by galaxy clusters (e.g.,Egamietal.2005;Bradleyetal.2008,2011;Zheng2009;Richard 2011). | The discovery of a high-redshift galaxy in the field of MACS0329 adds to several known high-reshift galaxies lensed by galaxy clusters \citep[e.g.,][]{Egami2005on2218highz,Bradley2008,Bradley2011,Zheng2009,Richard2011}. |
. Here we summarize the properties of this unique source. | Here we summarize the properties of this unique source. ( |
It is one of the highest-redshift multiply lensed objects(1) known to date, and lensed into four separate images. | 1) It is one of the highest-redshift multiply lensed objects known to date, and lensed into four separate images. |
The angular separation between arcs 1.2 and 1.4 is ~ 1’, considerably larger than previously reported cases (Egamietal.2005;Richard2011).. ( | The angular separation between arcs 1.2 and 1.4 is $\sim1\arcmin$ , considerably larger than previously reported cases \citep{Egami2005on2218highz, Richard2011}. ( |
2) The source is one of the brightest at z>6: its Ji25 magnitude is 24.0 AB, making it a viable candidate for follow-up spectroscopy. ( | 2) The source is one of the brightest at $z>6$: its $J_{125}$ magnitude is 24.0 AB, making it a viable candidate for follow-up spectroscopy. ( |
3) The galaxy is consistent with being a dwarf galaxy. | 3) The galaxy is consistent with being a dwarf galaxy. |
Its intrinsic (delensed) magnitude of Jj25=27.1 AB makes it a sub-L, galaxy at this redshift. | Its intrinsic (delensed) magnitude of $J_{125} = 27.1$ AB makes it a $L_{*}$ galaxy at this redshift. |
It occupies a source-plane area of ~2.2 kpc?, similar to previously deduced sizes of high-z lensed galaxies (e.g.Zitrinetal.2011b). | It occupies a source-plane area of $\sim$ 2.2 $^{2}$, similar to previously deduced sizes of $z$ lensed galaxies \citep[e.g.][]{Zitrin2011b}. |
. Due to the hierarchical growth of structure, galaxies are expected to be small at high redshifts, with dwarf galaxies constituting the building material of larger structures. | Due to the hierarchical growth of structure, galaxies are expected to be small at high redshifts, with dwarf galaxies constituting the building material of larger structures. |
Our source-plane reconstruction shows at least three (possibly knots, consistent with several other reports star-forming)of high-redshift galaxies with multiple components (Franxetal.1997;Bradleyetal. possibly as the result of merging. | Our source-plane reconstruction shows at least three (possibly star-forming) knots, consistent with several other reports of high-redshift galaxies with multiple components \citep{Franx1997on1358highz,Bradley2008,Bradley2011, Zheng2009,Oesch2010highz,Zitrin2011b}, possibly as the result of merging. |
In addition, we 2011b),measure an overall half-light radius of 0.12”, consistent with that found in Bouwens2006) and (Oeschetal.2010).. ( | In addition, we measure an overall half-light radius of $\sim0.12\arcsec$ , consistent with that found in \citet{Bouwens2004size,Bouwens2006size} and \citep{Oesch2010highz}. ( |
4) The SED fits to the multiband photometry of the source suggest a demagnified stellar mass of ~10?Mo, a SFR-weighted age of ~180 Myr, subsolar metallicity (Z/Ze~ low dust content (Ay~0.1 mag), and a demagnified 0.5),SFR of ~3.2Mo yr7!. | 4) The SED fits to the multiband photometry of the source suggest a demagnified stellar mass of $\sim10^{9}~\mathcal{M}_{\sun}$, a SFR-weighted age of $\sim180$ Myr, subsolar metallicity $Z/Z_{\sun}\sim0.5$ ), low dust content $A_{V}\sim0.1$ mag), and a demagnified SFR of $\sim3.2~\mathcal{M}_{\sun}$ $^{-1}$. |
The specific SFR of ~3.4 Gyr7!, which is slightly higher than that found by other recent studies (Gonzalezetal.2011;Stark2009;Labbéetal.2010;McLure 2011),, implies a mass-doubling time of just 600 Myr and therefore vigorous ongoing star formation considering its low mass. ( | The specific SFR of $\sim3.4$ $^{-1}$, which is slightly higher than that found by other recent studies \citep{Gonzalez2011highz,Stark2009highz,Labbe2010highz,McLure2011MNRAShighz}, , implies a mass-doubling time of just $600$ Myr and therefore vigorous ongoing star formation considering its low mass. ( |
5) The UV continuum is blue, with a UV-slope 8— —2.5+0.06, consistent with measurements of other faint z~6 galaxies and suggests that these sources are largely dust free (Bouwensetal.2009a,2011;Finkelstein2011;Vanzellaetal. 2011).. | 5) The UV continuum is blue, with a UV-slope $\beta = -2.5 \pm 0.06$ , consistent with measurements of other faint $z\sim6$ galaxies and suggests that these sources are largely dust free \citep{Bouwens2009uvslopes,Bouwens2011uvslope,Finkelstein2011highzUV,Vanzella2011uvslope}. |
'The discovery of the galaxy presented here shows once more the novelty and tremendous potential of galaxy clustersforobservationally accessing the faint early universe. | The discovery of the galaxy presented here shows once more the novelty and tremendous potential of galaxy clustersforobservationally accessing the faint early universe. |
'The authors thank Saurabh Jha for useful discussions, and the anonymous referee for valuable comments. | The authors thank Saurabh Jha for useful discussions, and the anonymous referee for valuable comments. |
and protons within the shocks. as well as place stringent constraints on fundamental neutrino properties. | and protons within the shocks, as well as place stringent constraints on fundamental neutrino properties. |
Neutrinos Irom inelastic nuclear collisions typically have energies of the order of LO GeV. Unless the phototube density in and other next egeneration neutrino telescopes is veher (han currently planned. (hese detectors will have a small effective detector area Lor such low energy neutrinos and will not be sensitive to them (Ixarle2002).. | Neutrinos from inelastic nuclear collisions typically have energies of the order of 10 GeV. Unless the phototube density in and other next generation neutrino telescopes is higher than currently planned, these detectors will have a small effective detector area for such low energy neutrinos and will not be sensitive to them \citep{bigice}. |
However. there is a growing indication that. like neutrinos from p—5 interactions. these neutrinos may be an important clue (o conditions in ancl around (he GRB central engine. | However, there is a growing indication that, like neutrinos from ${\rm p-\gamma}$ interactions, these neutrinos may be an important clue to conditions in and around the GRB central engine. |
Baheall& discussed (he neutrino signal from inelastic collisions occurring during the dvnanmic decoupling of neutrons in the acceleration stage of the evolution of a GRD fireball. | \cite{bah00} discussed the neutrino signal from inelastic collisions occurring during the dynamic decoupling of neutrons in the acceleration stage of the evolution of a GRB fireball. |
These jeutrinos offer an indication of the composition and Lorentz lactor of the GRB jet. | These neutrinos offer an indication of the composition and Lorentz factor of the GRB jet. |
Neutron diffusion in GRD environments was considered in a more general context by Mészáros (2000). | Neutron diffusion in GRB environments was considered in a more general context by \cite{mes00}. |
. Those authors showed that internal shocks at radii r~3+L0!*em and transverse diffusion of neutrons into the jet at much larger radii can lead (o an appreciable neutrino V.ional. | Those authors showed that internal shocks at radii $r\sim 3\cdot 10^{11}{\rm cm}$ and transverse diffusion of neutrons into the jet at much larger radii can lead to an appreciable neutrino signal. |
The possibility of using neutrinos to distinguish between collapsars aud supranovae Vietri&Stella1998.1999). was discussed by Guetta&Granot(2002). | The possibility of using neutrinos to distinguish between collapsars and supranovae \citep{vie98,vie99} was discussed by \cite{gue02}. |
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