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OL course will be also used to study their structure and At 1300 pani. to find enough ~ 3 107 L. high-z sources. we need to cover a surface of at [east 5 square degrees at a Se level of about 0.1 mJy (Lable 8)). | Of course will be also used to study their structure and At 1300 $\mu$ m, to find enough $\sim$ 3 $^{11}$ $_{\odot}$ high-z sources, we need to cover a surface of at least 5 square degrees at a $\sigma$ level of about 0.1 mJy (Table \ref{sfce_pred}) ). |
Such a survey will resolve about 50% of the CLB. | Such a survey will resolve about $\%$ of the CIB. |
To resolve 7804 of the CLB. one needs to reach a Se level of about 0.02 mv. (Lable 9)). | To resolve $\sim$ $\%$ of the CIB, one needs to reach a $\sigma$ level of about 0.02 mJy (Table \ref{resol_pred}) ). |
Therefore two kinds of surveys could be considered: a area ( 5 Sq. | Therefore two kinds of surveys could be considered: a large-area $\sim$ 5 Sq. |
deg.) | deg.) |
and an ultra deep (~ 10 arcmin?) survey. | and an ultra deep $\sim$ 10 $^2$ ) survey. |
For both surveys. the compact configuration has enough resolution not to be limited by the confusion. | For both surveys, the compact configuration has enough resolution not to be limited by the confusion. |
At 1300 jam. a 5o detection of 2.3 mv is reached in 1 see for a beam area of 0.16 aremin? proposal for phase 2 and Blain 2001). | At 1300 $\mu$ m, a $\sigma$ detection of 2.3 mJy is reached in 1 sec for a beam area of 0.16 $^2$ proposal for phase 2' and Blain 2001). |
This gives for the two types of surveys: In conclusion. if we want to achieve the two goals: (i) detect. enough. carly mergers mace of building blocks not allectec much. vet by star formation. ancl evolution and (it) probe most of the CLB source population at [arge wavelengths. we will have to do extragalactic surveys with using a substantial fraction of the time to find the sources. | This gives for the two types of surveys: In conclusion, if we want to achieve the two goals: (i) detect enough early mergers made of building blocks not affected much yet by star formation and evolution and (ii) probe most of the CIB source population at large wavelengths, we will have to do extragalactic surveys with using a substantial fraction of the time to find the sources. |
Such large surveys including the whole collaboration would. be much. more ellicient in terms of scientific progress than smaller area surveys conducted. by incliviclual smaller We have developped a phenomenological model. that constrains in a simple way the LR. luminosity function evolution with the redshift. ancl fits all the existing source counts ancl redshift clistribution. CLB intensity and for the inst time ς2 fluctuations observations from the απ-ας to he submuim range. | Such large surveys including the whole collaboration would be much more efficient in terms of scientific progress than smaller area surveys conducted by individual smaller We have developped a phenomenological model that constrains in a simple way the IR luminosity function evolution with the redshift, and fits all the existing source counts and redshift distribution, CIB intensity and for the first time CIB fluctuations observations from the mid-IR to the submm range. |
Phe model has been used to give some wedietions for futureLersehed deep survey observations and the albskyPlanck survey. | The model has been used to give some predictions for future deep survey observations and the all-sky survey. |
10 comes out that. the xdanned experiments(SIRESLerscehet Planch) will be mostly limited by the confusion. | It comes out that the planned experiments, ) will be mostly limited by the confusion. |
To find out a large number of objects that dominate the LE at high redshift. (272). uture experiments need both the angular resolution and sensitivity. | To find out a large number of objects that dominate the LF at high redshift $>$ 2), future experiments need both the angular resolution and sensitivity. |
Fhis can be achieved in the submnm only thanks o interferometres such asLLALA. | This can be achieved in the submm only thanks to interferometres such as. |
However. mapping large ractions of the skv with high signal-to-noise ratio will take | However, mapping large fractions of the sky with high signal-to-noise ratio will take |
As detailed in the literature (Ixistleretal.2008:Wistler2009:Qinetal. 2010)) we also consider an increase in GIUS rate as (1|z)"s where 0=0.2.0.5.0.8. | As detailed in the literature \cite{kistler08,kistler09,qin10}) ) we also consider an increase in GRB rate as $(1+z)^\delta$, where $\delta= 0.2,0.5,0.8$. |
We consider all of the SELL models with and without metallicity enhancements (no GSME evolution) ancl find that a few of these models are able to pass the £ and z constraints but fail to pass the BATSE and logN—P? constraints (Lable 4). | We consider all of the SFH models with and without metallicity enhancements (no GSMF evolution) and find that a few of these models are able to pass the $L$ and $z$ constraints but fail to pass the BATSE and $\log N-\log P$ constraints (Table 4). |
Lastly we consider the evolution of the luminosity function reals as detailed above. | Lastly we consider the evolution of the luminosity function break as detailed above. |
We see some consistency. with the CN model and evolution with 50.51.5. | We see some consistency with the CN model and evolution with $\gamma \sim 0.5-1.5$. |
The 3e regions or the €N models show areas of consistency. with a few showing 2e significance (i.c. 5=1.0. 1.3) (Figure 5)). | The $3\sigma$ regions for the CN models show areas of consistency, with a few showing $2\sigma$ significance (i.e. $\gamma=1.0, 1.3$ ) (Figure \ref{cosmo}) ). |
The general trend is again for shallow luminosity function slopes. he best models occurring in the area of (64.09.55.5)=(0.5.2.2.3107ergs +. 1.0). | The general trend is again for shallow luminosity function slopes, the best models occurring in the area of $(\alpha_1, \alpha_2, L_b,\gamma)=(0.5,2.2,3\times 10^{52} \rm ~erg~s^{-1}$ , 1.0). |
The HD models show some consistency to 36 in the same regions. although not as oacdlyv as the €N model (Lable 4 | The HB models show some consistency to $\sigma$ in the same regions, although not as broadly as the CN model (Table 4). |
Our work supports the idea that the CIUD rate is enhanced at higher redshift (Daigneetal.2006:Le2007:Guetta2009:Qinetal.2010:Wanderman&Piran 20100). | Our work supports the idea that the GRB rate is enhanced at higher redshift \cite{daigne06,le07, guetta07, li08, kistler08, kistler09, salvaterra09,salvaterra09b,salvaterra07, qin10, wanderman10}) ). |
Phe form of this increase. however. is still unclear. | The form of this increase, however, is still unclear. |
We have tested various SELL models anc enhancements to the GiB rate. reflecting possible clleets from changing cosmic metallicity and other evolutionary ellects. with a Monte. Carlo codo. | We have tested various SFH models and enhancements to the GRB rate, reflecting possible effects from changing cosmic metallicity and other evolutionary effects, with a Monte Carlo code. |
The resulting output was then tested for consistency. with a variety of available and. BATSE data. including the £. 2. and peak photon [lux distributions. | The resulting output was then tested for consistency with a variety of available and BATSE data, including the $L$, $z$, and peak photon flux distributions. |
Even when considering a numerical simulation model that takes into account a variety of realistic galactic evolution ellects. both with and without metal cuts. and a metallicity relation based on the GSME (Langer&Norman 2006)) our models do not show strong consistency. with the observed. sample. although we believe this is the right direction for this tvpe of study. | Even when considering a numerical simulation model that takes into account a variety of realistic galactic evolution effects, both with and without metal cuts, and a metallicity relation based on the GSMF \cite{langerandnorm06}) ) our models do not show strong consistency with the observed sample, although we believe this is the right direction for this type of study. |
This may indicate that metallicity is not solely responsible for the increased. rate and that. perhaps some other tvpe of enhancement is needed. | This may indicate that metallicity is not solely responsible for the increased rate and that perhaps some other type of enhancement is needed. |
To this end. we test both GRB rate evolution ane luminosity function. (break luminosity) evolution with redshift. finding that the latter is allowed within the constraints of the BATSE and cata with moderate (xLy,(1|wold?a ο] evolution. | To this end, we test both GRB rate evolution and luminosity function (break luminosity) evolution with redshift, finding that the latter is allowed within the constraints of the BATSE and data with moderate $\propto L_b\times(1+z)^{\sim 0.8-1.2}$ ) evolution. |
H This statement has. of course. a few caveats. | This statement has, of course, a few caveats. |
Embedded in the metallicity relation are a variety of assumptions about the GSME and the observed mass-metallicity relation. | Embedded in the metallicity relation are a variety of assumptions about the GSMF and the observed mass-metallicity relation. |
Laskar ct al. ( | Laskar et al. ( |
2011) show. using HIST observations of GRB host galaxies. that the metallicity relationship likely evolves between redshifts of 3.5. which would further allect the results. | 2011) show, using HST observations of GRB host galaxies, that the metallicity relationship likely evolves between redshifts of $-$ 5, which would further affect the results. |
Lt is possible that other combinations of parameters or assumptions müght vield a more realistic relation. ancl we sugeest further work on how the GSME and. stellar IME work in tandem to alleet the problem at hand. | It is possible that other combinations of parameters or assumptions might yield a more realistic relation, and we suggest further work on how the GSMF and stellar IMF work in tandem to affect the problem at hand. |
In addition. recent works have studied the AJZ relation of Type HE GIDs and found that the hosts lie below the SDSS AZrelation (Ixocevski&West2011:Mannuccietal.Campisietal.2011)). | In addition, recent works have studied the $M-Z$ relation of Type II GRBs and found that the hosts lie below the SDSS $M-Z$ relation \cite{kocevski11,mannucci11,campisi11}) ). |
This adds further evidence to the fact ju the assumption of this relation for these types of bursts is likely not. valid. ancl perhaps a consequence of the active μαar-formation environment instead of a strict metallicity cut. | This adds further evidence to the fact that the assumption of this relation for these types of bursts is likely not valid, and perhaps a consequence of the active star-formation environment instead of a strict metallicity cut. |
We have explored some basic changes. such as the evolution of the GSME faint-end slope. but a comprehensive study of this relation or a realistic alternative. are needed. | We have explored some basic changes, such as the evolution of the GSMF faint-end slope, but a comprehensive study of this relation or a realistic alternative, are needed. |
We have detailed a numerical and statistical approach aimed at understanding the properties of the CRB rate in the context of the cosmic star-formation history. inclucling 1 constraints [rom newly discovered high-z bursts ancl the possibly cllects of metallicity and various types of evolution. | We have detailed a numerical and statistical approach aimed at understanding the properties of the GRB rate in the context of the cosmic star-formation history, including the constraints from newly discovered $z$ bursts and the possibly effects of metallicity and various types of evolution. |
Recent works have addressed this problem in similar (Qinοἱal.2010)) and fully analytical (Manderman&Piran2010)) wavs. and share some common points. although both call on fairly strong evolution of the GRB rate (1|sy?~0.62) which we do not find. | Recent works have addressed this problem in similar \cite{qin10}) ) and fully analytical \cite{wanderman10}) ) ways, and share some common points, although both call on fairly strong evolution of the GRB rate $(1+z)^\delta \sim 0.6-2)$ which we do not find. |
Our work also benefits from. the inclusion of a fully numerical star-formation history model (Choi&aeamine 2010)) as well as a probing of the metallicity relation and cosmological considerations that may allect the GRB rate which are not included. in contemporary works on the subject. | Our work also benefits from the inclusion of a fully numerical star-formation history model \cite{choiandnag10}) ) as well as a probing of the metallicity relation and cosmological considerations that may affect the GRB rate which are not included in contemporary works on the subject. |
Butler et al. ( | Butler et al. ( |
2010) do not find evidence of strong luminosity function or CRB rate evolution and find that a smoothecl metallicity cut o£ ZiZ.=0.205.following the metallicity considerations of LN. can account for the observations of the current sample. although they acknowledge that there are large errors bars. | 2010) do not find evidence of strong luminosity function or GRB rate evolution and find that a smoothed metallicity cut of $Z/Z_\odot = 0.2-0.5$, following the metallicity considerations of LN, can account for the observations of the current sample, although they acknowledge that there are large errors bars. |
They also do not include evolution of the GSAIF. which may account for the dillerences with this work. | They also do not include evolution of the GSMF, which may account for the differences with this work. |
In addition. we analyse most components separately. and it is possible that the observed. distribution is a superposition of a variety. of effects. | In addition, we analyse most components separately, and it is possible that the observed distribution is a superposition of a variety of effects. |
With enough. computational time the various combinations of ellects can and should be tested. | With enough computational time the various combinations of effects can and should be tested. |
ὃν fitting the recshilt distribution and logNoP? distribution. of D.VESIS ancl bursts. Canipisi et. al. ( | By fitting the redshift distribution and $\log N-\log P$ distribution of BATSE and bursts, Campisi et al. ( |
2010) have reached the similar conclusion that Tvpe LL GRBs are unbiased. tracers of the star-lormation history. | 2010) have reached the similar conclusion that Type II GRBs are unbiased tracers of the star-formation history. |
Their analysis supports two possible scenarios: (i) a model with no metal cuts and a stronely evolving luminosity function or (ii) a non-evolving luminosity function. with a metal cut of Z/Z.«0.8. | Their analysis supports two possible scenarios: (i) a model with no metal cuts and a strongly evolving luminosity function or (ii) a non-evolving luminosity function with a metal cut of $Z/Z_\odot < 0.3$. |
Both scenarios assume and Lit a Schechter luminosity function. | Both scenarios assume and fit a Schechter luminosity function. |
This results are similar to the results presented here. although the luminosity evolution is stronger for the non-metal cut case and the authors claim such large changes in GRB properties with redshift as unrealistic. favoring a model with a metal cut and. no luminosity function evolution. | This results are similar to the results presented here, although the luminosity evolution is stronger for the non-metal cut case and the authors claim such large changes in GRB properties with redshift as unrealistic, favoring a model with a metal cut and no luminosity function evolution. |
This work is supported by NSE through grant ASTT-0008362. and by NASA through. grants NNXIOAXDASC: and NNNIOAPS3G. KN is supported. in part. bv. the NSE grant. AST-0807491. NASA erant. LIST-AR-12143-01- National Aeronautics and Space Administration under Crant/C'ooperative Agreement No. | This work is supported by NSF through grant AST-0908362, and by NASA through grants NNX10AD48G and NNX10AP53G. KN is supported in part by the NSF grant AST-0807491, NASA grant HST-AR-12143-01-A, National Aeronautics and Space Administration under Grant/Cooperative Agreement No. |
NNNOSALSTA issued by the Nevada NASA EPSCoR. program. and the President's Infrastructure Aware from UNLV. | NNX08AE57A issued by the Nevada NASA EPSCoR program, and the President's Infrastructure Award from UNLV. |
This. research. is. also | This research is also |
2000 periastron passages and have been described in detail in Johnston et al. ( | 2000 periastron passages and have been described in detail in Johnston et al. ( |
1996) and Johnston et al. ( | 1996) and Johnston et al. ( |
2001). | 2001). |
Generally. observations were made at 5 cillerent frequencies. 1.2. 1.4. 1.5. 4.8 and SA Cllz. | Generally, observations were made at 5 different frequencies, 1.2, 1.4, 1.5, 4.8 and 8.4 GHz. |
The three frequencies near 1.4 Cllz were obtained simultaneously aid. the observation would wen be followed. by one at either 4.8 or δ Cllz before eain observing at the lower frequencies. | The three frequencies near 1.4 GHz were obtained simultaneously and the observation would then be followed by one at either 4.8 or 8.4 GHz before again observing at the lower frequencies. |
1n order to measure an accurate DAL we assume that the DAL contribution from the 1.5 kpe along the line of sight through the interstellar medium is 146.8 pe em. (Connorsetal.2 2002). | In order to measure an accurate DM we assume that the DM contribution from the 1.5 kpc along the line of sight through the interstellar medium is 146.8 pc $^{-3}$ \cite{cjmm02}. |
. Then the extra DAL contribution from the clisk of the Be star could be computed from the relative difference in the timing residuals between the dilferent. frequencies. | Then the extra DM contribution from the disk of the Be star could be computed from the relative difference in the timing residuals between the different frequencies. |
Typical errors using this method are ~0.1 peem | Typical errors using this method are $\sim0.1$ $\pc$. |
The timing properties were analysed using the pulsar timing program“TEAIPO*.. which provides. least-squares fitting to the pulsar rotation ancl orbital parameters and the dispersion measure. | The timing properties were analysed using the pulsar timing program, which provides least-squares fitting to the pulsar rotation and orbital parameters and the dispersion measure. |
The MSS binary model of Wex (1998) was used for all fits except those in Section ??.. | The MSS binary model of Wex (1998) was used for all fits except those in Section \ref{sec:btj}. |
Po fit for steps in the orbital parameters. a new binary model (DT) based on the DT model (Blandford&Teukolskv1976). was implemented. | To fit for steps in the orbital parameters, a new binary model (BTJ) based on the BT model \cite{bt76} was implemented. |
This allows cumulative steps in longitude of periastron (uw). projected semi-major axis (à). eccentricity (c) and binary period. (15) to be inserted at specified: times ancl also allowed setting or solving for jumps in pulsar phase at the specified times. | This allows cumulative steps in longitude of periastron $\omega$ ), projected semi-major axis $x$ ), eccentricity $e$ ) and binary period $P_b$ ) to be inserted at specified times and also allowed setting or solving for jumps in pulsar phase at the specified times. |
We note that our results dilfer slightly from those of Wex et al. ( | We note that our results differ slightly from those of Wex et al. ( |
1998) when analvsing the same data set. | 1998) when analysing the same data set. |
ης is because new standard. profile templates were used and the dispersion analysis re-done for this work. | This is because new standard profile templates were used and the dispersion analysis re-done for this work. |
Fig. | Fig. |
1 shows that the pulse profile evolves substantially over the observed. range from 660 to 13600 MIIz. | \ref{fg:std} shows that the pulse profile evolves substantially over the observed range from 660 to 13600 MHz. |
At first elance it appears that. the profile gets narrower. (when plotted as in Fig. 1)) | At first glance it appears that the profile gets narrower (when plotted as in Fig. \ref{fg:std}) ) |
at lower frequencies. the opposite trend to that observed: in most. conal profiles. | at lower frequencies, the opposite trend to that observed in most conal profiles. |
However. a closer examination suggests that this apparent width: variation is largely due to dillering spectral index of components making up each peak of the profile. | However, a closer examination suggests that this apparent width variation is largely due to differing spectral index of components making up each peak of the profile. |
The outer components are stronger at higher frequencies whereas components on the inner edge of cach peak become strong at lower frequencies and dominate the profile at. [frequencies below 2. 11. | The outer components are stronger at higher frequencies whereas components on the inner edge of each peak become strong at lower frequencies and dominate the profile at frequencies below 2 GHz. |
The latter. spectral index for outer conal components. is commonly observed. in other conal pulsars (Rankin1983:Lyne&Manchester1988). and reinforces the interpretation that the steep edges define the boundary of a wide emission cone from a single polar region (cf. | The flatter spectral index for outer conal components is commonly observed in other conal pulsars \cite{ran83,lm88} and reinforces the interpretation that the steep edges define the boundary of a wide emission cone from a single polar region (cf. |
Manchester 1996)). | Manchester \nocite{man96}) ). |
As PSR. 1259.63 approaches periastron and. moves into the Be-star disk. electron density variations along the ine of sight result in DM changes as a function of orbital hase. | As PSR $-$ 63 approaches periastron and moves into the Be-star disk, electron density variations along the line of sight result in DM changes as a function of orbital phase. |
Fig. | Fig. |
2aa shows the observed. variations from. 1992. covering the 1994. 1997 and 9000 periastrons. and Fie. | \ref{fg:dm}a a shows the observed variations from 1992, covering the 1994, 1997 and 2000 periastrons, and Fig. |
2bb shows expanded. plots for +75 davs about the periastrons. | \ref{fg:dm}b b shows expanded plots for $\pm75$ days about the periastrons. |
DAL variations for the 1994 periastron were reported by Johnston et al. (1996). | DM variations for the 1994 periastron were reported by Johnston et al. \nocite{jml+96}, |
. ancl for the LOOT periastron these were monitored extensively and reported by Johnston ct al. ( | and for the 1997 periastron these were monitored extensively and reported by Johnston et al. ( |
2001). while the 2000 variations have not been previously reported. | 2001), while the 2000 variations have not been previously reported. |
At the 1994 periastron. a peak of 10.7E0.2 peem7 for ADM was observed at 730: after this it decreased with some Iluctuations until the pulsar was eclipsed. | At the 1994 periastron, a peak of $10.7\pm0.2$ $\pc$ for $\Delta$ DM was observed at $\tau-30$; after this it decreased with some fluctuations until the pulsar was eclipsed. |
In 1997. the peak was 7.7 poemat 728. slightly later than in 1994. and again the DM crops with some Iluctuations until the pulsar is eclipsed at ~718. | In 1997, the peak was $7.7$ $\pc$ at $\sim\tau-28$, slightly later than in 1994, and again the $\Delta$ DM drops with some fluctuations until the pulsar is eclipsed at $\sim\tau-18$. |
There are no multi-[requency data prior to the 2000 periastron. | There are no multi-frequency data prior to the 2000 periastron. |
The points of DM in Fie? are derived by extrapolating a timing solution based on observations in the 60 days before these pre-periastron points shown in Fig. | The points of $\Delta$ DM in \ref{fg:dm} are derived by extrapolating a timing solution based on observations in the 60 days before these pre-periastron points shown in Fig. |
2bb. The latest pre-periastron point is 52 davs before the periastron. | \ref{fg:dm}b b. The latest pre-periastron point is 52 days before the periastron. |
Based on the 1994 and 1997 behaviour. the peak of DÀ was not observed. | Based on the 1994 and 1997 behaviour, the peak of $\Delta$ DM was not observed. |
The DÁM fluctuations are well observed after the 1997 ancl 2000 periastrons. | The $\Delta$ DM fluctuations are well observed after the 1997 and 2000 periastrons. |
The ADAL values are much smaller than before periastron ancl approach zero about 30. days after periastron. | The $\Delta$ DM values are much smaller than before periastron and approach zero about 30 days after periastron. |
In fact. Fig. | In fact, Fig. |
2. (a) shows that the ADM values are slightly negative in müd-orbit. suggesting the reference. DAL value. of 146.8. peem. is too large. | \ref{fg:dm}~ (a) shows that the $\Delta$ DM values are slightly negative in mid-orbit, suggesting the reference DM value of 146.8 $\pc$ is too large. |
The filterbank observations from 7|43 after the 2000 poriastron to the end of the cata set gives a mean ollset of 0.9+ peem from the reference value: corresponding to an interstellar dispersion of 146.640.) pecm | The filterbank observations from $\tau+43$ after the 2000 periastron to the end of the data set gives a mean offset of $-0.2\pm0.1$ $\pc$ from the reference value; corresponding to an interstellar dispersion of $\pm$ 0.1 $\pc$. |
We first. used. simulations to determine the shape and amplitude of timing residuals corresponding {ο various xwanmeters used in the fitting procedure. | We first used simulations to determine the shape and amplitude of timing residuals corresponding to various parameters used in the fitting procedure. |
Fig. | Fig. |
3. shows us he timing signatures of the various parameters and the size of the resultant phase perturbations. | \ref{fg:many} shows us the timing signatures of the various parameters and the size of the resultant phase perturbations. |
In the figure we have attempted to ensure that the resultant resicluals were of he order of tens of milliseconds. similar to the residuals seen in the real data. | In the figure we have attempted to ensure that the resultant residuals were of the order of tens of milliseconds, similar to the residuals seen in the real data. |
Apart from the w and ο simulations. we introduced. steps at the second. and. third. periastron massages. with the change at the third periastron being twice he magnitude ancl of opposite sign to the change at the | Apart from the $\dot{\omega}$ and $\dot{x}$ simulations, we introduced steps at the second and third periastron passages, with the change at the third periastron being twice the magnitude and of opposite sign to the change at the |
“broad-band filter” (BBF). | “broad-band filter” (BBF). |
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