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Differential gains are available in MeqTrees: pointing selfcal is implemented in an experimental version of CASA (Bhatnagar priv. | Differential gains are available in MeqTrees; pointing selfcal is implemented in an experimental version of CASA (Bhatnagar priv. |
comm.). | comm.), |
but is not publicly available at time of writing. | but is not publicly available at time of writing. |
This makes a quantitative comparison impossible. but the algorithms may be compared in principle. | This makes a quantitative comparison impossible, but the algorithms may be compared in principle. |
The peeling approach and differential gains are very similar in that they attempt to solve for the same effect: a direction-dependent complex gain term. | The peeling approach and differential gains are very similar in that they attempt to solve for the same effect: a direction-dependent complex gain term. |
In essence. peeling approximates a full-sky RIME as: where X,,, is the model coherency PMof 5source(GT. s (typically à phase-shifted delta function. for a point source model. but Gaussian sources are also possible in e.g. NEWSTAR). | In essence, peeling approximates a full-sky RIME as: where $\coh{X}{spq}$ is the model coherency of source $s$ (typically a phase-shifted delta function, for a point source model, but Gaussian sources are also possible in e.g. NEWSTAR). |
Peeling consists of a least-squares solution for for one set of gams at a time (as in regular selfcal). followed by "temporary" subtraction of sources for which a solution has been obtained. | Peeling consists of a least-squares solution for for one set of gains at a time (as in regular selfcal), followed by “temporary” subtraction of sources for which a solution has been obtained. |
Differential gains uses an equation like (16)). | Differential gains uses an equation like \ref{eq:de}) ). |
First. a regular selfcal step is done to obtain G,, solutions on short time/frequency scales. | First, a regular selfcal step is done to obtain $\jones{G}{p}$ solutions on short time/frequency scales. |
This is followed by à simultaneous least-squares solution for all the AE,, terms. on longer time/frequency scales. | This is followed by a simultaneous least-squares solution for all the $\jones{\Delta E}{sp}$ terms, on longer time/frequency scales. |
Peeling 1s subject to selfeal contamination at each stage of the process. due to the as-yet-unsolved-for contributions of fainter sources. | Peeling is subject to selfcal contamination at each stage of the process, due to the as-yet-unsolved-for contributions of fainter sources. |
This is especially severe when sources have comparable flux. | This is especially severe when sources have comparable flux. |
Differential gains overcomes this by solving for all sources simultaneously. | Differential gains overcomes this by solving for all sources simultaneously. |
In principle. it should | In principle, it should |
with the expression where SeoAvu is the limit on the velocity integratec ine [lux iu Jy km lop is the observing frequency iu GHz. and D, is the luminosity clistauce iu xc. | with the expression where $S_{CO} \Delta v$ is the limit on the velocity integrated line flux in Jy km $^{-1}$, $\nu_{obs}$ is the observing frequency in GHz, and $D_L$ is the luminosity distance in Mpc. |
The choice of cosmological parameters euters in Dy. aud we adopt. Ay=τὸ καν ο οτα i ) lor consistency. with most work in this field. ( | The choice of cosmological parameters enters in $D_L$ , and we adopt $H_0 = 75$ km $^{-1}$, $\Omega = 1$ and $\Omega_{\Lambda} = 0$ for consistency with most work in this field. ( |
An alternative cosmology with Hy=f)T Nlb 5 ] and O4—Q.T results in D, largere by a factor of 1.51 for this redshil.) | An alternative cosmology with $H_0 = 75$ km $^{-1}$, $\Omega = 1$ and $\Omega_{\Lambda} = 0.7$ results in $D_L$ larger by a factor of 1.54 for this redshift.) |
The effective lirewidth is 1ot known. but it likely falls iu the range 150 t0 550 kins + found for a large sample of tItralumiuous galaxies in the local universe (Solomon et al. | The effective linewidth is not known, but it likely falls in the range 150 to 550 km $^{-1}$ found for a large sample of ultraluminous galaxies in the local universe (Solomon et al. |
1997). | 1997). |
For the 3o flux limit obtained in he tore sesitive part of the CO J=2-1 spectrum. assuming a linewidth of 200 kim Lo Leo—]1)«5.1x10! N kin . Ρο | For the $3\sigma$ flux limit obtained in the more sensitive part of the CO J=2–1 spectrum, assuming a linewidth of 200 km $^{-1}$, $L^{'}_{CO}(2-1) < 5.1 \times 10^{10}$ K km $^{-1}$ $^2$. |
For the 3e flux. limit for CO J—5-I. agalu assuimniug a linewidth of 200 kns Ly ho.<eod)x10 Ix kin ! pc?. If obta | For the $3\sigma$ flux limit obtained for CO J=5–4, again assuming a linewidth of 200 km $^{-1}$, $L^{'}_{CO}(5-4) < 3.0 \times 10^{10}$ K km $^{-1}$ $^2$. |
inedthe assumed linewidth were two times larger. then these luminosity liinitis be V2 times higher. | If the assumed linewidth were two times larger, then these luminosity limits would be $\sqrt{2}$ times higher. |
Conversion of these CO luminosity limits to molecular gas mass limits is faught witl uncertaluties. | Conversion of these CO luminosity limits to molecular gas mass limits is fraught with uncertainties. |
But a simple conversiou factor from CO luminosity to Πο mass is commony taken to be L5AM. (Ix km 1 pc)> Lale value deteriniued for Milky Way molecular clouds (Sauders. Scoville Soifer 1991). | But a simple conversion factor from CO luminosity to $_2$ mass is commonly taken to be $4.5~M_{\odot}$ (K km $^{-1}$ $^2$ $^{-1}$, the value determined for Milky Way molecular clouds (Sanders, Scoville Soifer 1991). |
There is evidence [rom comparisous of luminosity based mass estinales with dynamical nass estimates tiat the couversiou factor may be perhaps five times lower in ultraluimious objects {Downes Soloi101 1998). | There is evidence from comparisons of luminosity based mass estimates with dynamical mass estimates that the conversion factor may be perhaps five times lower in ultraluminous objects (Downes Solomon 1998). |
Additioual corrections of order unity are also ueede| to accotut properly Or excitatiou from t1e elevated«cosmic background radiation at high recshilt. | Additional corrections of order unity are also needed to account properly for excitation from the elevated cosmic background radiation at high redshift. |
Adopting the Galactic conversion factor for CO J=2-1 liue luminosity gives a limit ou the molecular gas nass of~2.310H. iu the SDSS 1011-0122 system. | Adopting the Galactic conversion factor for CO J=2–1 line luminosity gives a limit on the molecular gas mass of $\sim2.3\times10^{11}~M_{\odot}$ in the SDSS 1044-0125 system. |
Using the same couversion factor for the CO J=5-1 inelmiosity gives a liiil ou the molecular gas mass of of ~1.3:10H.AL. in the SDSS 1011-0120 system. | Using the same conversion factor for the CO J=5–4 line luminosity gives a limit on the molecular gas mass of of $\sim1.3\times10^{11}~M_{\odot}$ in the SDSS 1044-0125 system. |
These mass limits are comparable to the mass indicated. [roin the detection of CO etuission [rom some z>| quasars. iucludiug at least two thought uot be amplified by eravitatioual leusine. | These mass limits are comparable to the mass indicated from the detection of CO J=5--4 emission from some $z>4$ quasars, including at least two thought not to be amplified by gravitational lensing. |
In particular. observatious of CO J=5-1 emission [rom BRI202-0722 at =L7 (Omont al. | In particular, observations of CO J=5–4 emission from BR1202-0725 at $z=4.7$ (Omont al. |
1996. Ohta et al. | 1996, Ohta et al. |
1996) and. BRIL335-O117 at 1.1 (Guilloteau et al. | 1996) and BRI1335-0417 at $z=4.4$ (Guilloteau et al. |
1997) indicate molecular gas masses in excess of LOM ML, (adjusted for the cosmology aud CO to H» conversion factor adopted here). | 1997) indicate molecular gas masses in excess of $10^{11}$ $_{\odot}$ (adjusted for the cosmology and CO to $_2$ conversion factor adopted here). |
There is uo clear physical argument to explain why some quasar environments show CO emission at this sensitivity level while others do not (Cuilloteatu et al. | There is no clear physical argument to explain why some quasar environments show CO emission at this sensitivity level while others do not (Guilloteau et al. |
1900). | 1999). |
In any case. the CO J=2-1 and J—5-1 luminosity limits suggest that the e1virouruent ol SDSS 1011-0125 does not possess an euo‘MOUS mass reservoir of either low excitation or high excitation molecular gas. | In any case, the CO J=2–1 and J=5–4 luminosity limits suggest that the environment of SDSS 1044-0125 does not possess an enormous mass reservoir of either low excitation or high excitation molecular gas. |
The CO 22-1 limit is comparable to the amount of molecular gas detected toware the leused quasar APM 08279-5255. wherePapacopoulos et al. ( | The CO J=2–1 limit is comparable to the amount of molecular gas detected toward the lensed quasar APM 08279+5255, wherePapadopoulos et al. ( |
2001) [ouud several CO —2-] emission features with total luminosity 6.6+3.1x10! lx kins | pe? attributed to (unleused) molecular gas rich companion galaxies to the αιasar host. | 2001) found several CO J=2–1 emission features with total luminosity $6.6\pm3.1 \times 10^{11}$ K km $^{-1}$ $^2$ attributed to (unlensed) molecular gas rich companion galaxies to the quasar host. |
For the SDSS 1011-0125 observations. such [features | For the SDSS 1044-0125 observations, such features |
(Estimate of A3). Sincee the integral. appearing. in. Ay is.done over Q?>\(Qt; we obtain⋅ Usiug this inequality together with the fact that we Cal estimate the term Ay as follows: where the above series converges for17>n4-2. (Estimate of A). Using Lemma 2.10... and the fact that ||.αςπα& on Qe\Qt the term Ay can be estimated as follows: where the aboveseries also converges for 10>n+2. From (2.26)). (2.29)) aud (2.30)). inequality (2.23)) directly follows with a coustaut C>0 indepeudent of j. | (Estimate of $A_3$ Since the integral appearing in $A_{3}$ isdone over $Q^{2-k}\setminus Q^{1-k}$, we obtain Using this inequality together with the fact that we can estimate the term $A_{3}$ as follows: where the above series converges for$\o{m}>n+2$ (Estimate of $A_4$ Using Lemma \ref{mean_est}, and the fact that $\|2^{(j-1)a}z\|^{\o{m}}\geq 2^{\o{m}(j-k)}$ on $Q^{2-k}\setminus
Q^{1-k}$, the term $A_{4}$ can be estimated as follows: where the aboveseries also converges for $\o{m}>n+2$ From \ref{key_eq-1}) ), \ref{key_eq8}) ) and \ref{key_eq9}) ), inequality \ref{key_eq1}) ) directly follows with a constant $C>0$ independent of $j$ . |
a The constauts that will appear may differ from line to line. but only depend on » aud 1. | $\hfill{\blacksquare}$ The constants that will appear may differ from line to line, but only depend on $n$ and $m$. |
The proof of this lemina combines somehow the proof of Lemimas 2.9 and 2.11.. | The proof of this lemma combines somehow the proof of Lemmas \ref{BSI_lelemme} and \ref{key_lemma}. . |
We write down ui as a finite stun of a telescopie sequencefor NV> 1: From Lemuna 2.10.. we deduce that: | We write down $u_{Q^{1}}$ as a finite sum of a telescopic sequencefor $N\geq 1$ : From Lemma \ref{mean_est}, , we deduce that: |
the mass ratio of 0.915 and emphasized the need for a new photometric analysis of the system to attain the absolute physical parameters. | the mass ratio of 0.915 and emphasized the need for a new photometric analysis of the system to attain the absolute physical parameters. |
Zascheetal.(2009) updated the light elements after having analyzed all photometric and astrometric data available for the system. | \citet{zasche09} updated the light elements after having analyzed all photometric and astrometric data available for the system. |
According to the observational indicators. MR Del has properties similar to stars of BY Dra type or of RS CVn stars. | According to the observational indicators, MR Del has properties similar to stars of BY Dra type or of short-period RS CVn stars. |
The results of our photometric analysis. based on updated spectroscopic elements of Pribullaetal.(2009b).. are given in Table 7.. | The results of our photometric analysis, based on updated spectroscopic elements of \citet{pribb09}, are given in Table \ref{TabMRDel}. |
Figure 3. shows the observed (LCO) and the synthetic (LCC) light curves 1n the B. V. and R filters (upper left). the B—V and V-B color indices (lower left). the O—C residuals (upper right) and the geometrical model of the system in representative phases 0.3 and 0.7 (lower right). | Figure \ref{fMRDel} shows the observed (LCO) and the synthetic (LCC) light curves in the B, V, and R filters (upper left), the $B-V$ and $V-B$ color indices (lower left), the $O-C$ residuals (upper right) and the geometrical model of the system in representative phases 0.3 and 0.7 (lower right). |
Table 7 lists parameter uncertainties estimated by combining the formal nonlinear least-squared fitting errors with the errors arising from the uncertainty of the spectroscopic mass ratio (q=0.915+ 0.012). as described in Section 3.. | Table \ref{TabMRDel} lists parameter uncertainties estimated by combining the formal nonlinear least-squared fitting errors with the errors arising from the uncertainty of the spectroscopic mass ratio $q=0.915 \pm 0.012$ ), as described in Section \ref{analysis}. |
Our model includes two cool spots on the more-massive. hotter componet. | Our model includes two cool spots on the more-massive, hotter component. |
The spotted model ts supported by the X-ray observations. | The spotted model is supported by the X-ray observations. |
Another activity indicator is the flare event observed by Clausenetal.(2001) which was most pronounced in the u band. | Another activity indicator is the flare event observed by \citet{clau01} which was most pronounced in the u band. |
In addition. there are night-to-night differences in the light curves. increasing in strength from the y to the u band. so cool spots can be expected on one or both components: however. the uniqueness of the spot locations obtained in our solution is questionable to some degree. | In addition, there are night-to-night differences in the light curves, increasing in strength from the y to the u band, so cool spots can be expected on one or both components; however, the uniqueness of the spot locations obtained in our solution is questionable to some degree. |
A good fit could not be | A good fit could not be |
of BBLPO2. | of BBLP02. |
The models cau be sununuanzed as follows: the dominant dark matter component. which is unaffected bv the enerev injection. collapses and virializes to form. bound halos. | The models can be summarized as follows: the dominant dark matter component, which is unaffected by the energy injection, collapses and virializes to form bound halos. |
The distribution of the dark matter iu such halos is assunued to be the same as for the selfsimular clusters described above. | The distribution of the dark matter in such halos is assumed to be the same as for the self-similar clusters described above. |
While the dark component is unaffected by energy injection. the collapse of the barvouic conrponeut is hiudered by the pressure forces induced x eutropy injection. | While the dark component is unaffected by energy injection, the collapse of the baryonic component is hindered by the pressure forces induced by entropy injection. |
If the maxima iufall velocity due surely to eravity of the dark halo is subsonic. the flow will )o stronely affected by the pressure aud it will not uudereo accretion shocks. | If the maximum infall velocity due purely to gravity of the dark halo is subsonic, the flow will be strongly affected by the pressure and it will not undergo accretion shocks. |
[It is assmmed that the barvous will accumulate outo the halosZsentropicallg at the adiabatic Bouc accretion rate (as described iu Balogh et al. | It is assumed that the baryons will accumulate onto the halos at the adiabatic Bondi accretion rate (as described in Balogh et al. |
1999). | 1999). |
This treatineut. however. is only appropriate for low mass alos. | This treatment, however, is only appropriate for low mass halos. |
Ifthe eravity of the dark halos is strong enough (as it is expected to be in the hot clusters being considered vere) that the maxima iufall velocity is frausonulc or supersonic. the eas will experience an additional (generally dominant) cutropy increase due to accretion shocks. | If the gravity of the dark halos is strong enough (as it is expected to be in the hot clusters being considered here) that the maximum infall velocity is transonic or supersonic, the gas will experience an additional (generally dominant) entropy increase due to accretion shocks. |
Iu order to trace the shock history of the eas. a detailed knowledge of the merger history of the cluster/eroup is required but is not considered by BDBLDPU2. | In order to trace the shock history of the gas, a detailed knowledge of the merger history of the cluster/group is required but is not considered by BBLP02. |
ILustead. it js asstuned that at some earlier time the most massive cluster progenitor will have had a mass low euouch such that shocks were uceleible in its formation. simular to the low mass halos discussed above. | Instead, it is assumed that at some earlier time the most massive cluster progenitor will have had a mass low enough such that shocks were negligible in its formation, similar to the low mass halos discussed above. |
This progenitor forms au iscutropic gas core of radius rat the cluster ceuter. | This progenitor forms an isentropic gas core of radius$r_{c}$ at the cluster center. |
The cutropy of eas outside of the core. however. will be affected by shocks. | The entropy of gas outside of the core, however, will be affected by shocks. |
Receut high resolution numerical simulatious sugecst that the "entropy profile for gas outside this core can be adequately represented by a simple analytic expression given bv luA(r)=InAy|olu(re£r.) (bewis et al. | Recent high resolution numerical simulations suggest that the “entropy” profile for gas outside this core can be adequately represented by a simple analytic expression given by $\ln{K(r)} = \ln{K_0} + \alpha \ln{(r/r_c)}$ (Lewis et al. |
2000). where AN—AT,3* | 2000), where $K \equiv
kT_e n_e^{-2/3}$. |
For the massive. hot clusters (Ty23 keV) of interest here. a~1.1 (Tozzi Norman 2001: DBLDPO2). | For the massive, hot clusters $T_X \gtrsim 3$ keV) of interest here, $\alpha \sim 1.1$ (Tozzi Norman 2001; BBLP02). |
Following this prescription aud specitvine the paraluctors ον Pyas(e). aud a (as discussed in BBLPO2) colmpletely determines the models. | Following this prescription and specifying the parameters $r_c$, $\rho_{gas}(r_c)$, and $\alpha$ (as discussed in BBLP02) completely determines the models. |
Under all conditions. the gas is assuned to be iu hverostatic equilibria within the dark halo potential. | Under all conditions, the gas is assumed to be in hydrostatic equilibrium within the dark halo potential. |
The complicated: effects of radiative cooling are neglected by these models. | The complicated effects of radiative cooling are neglected by these models. |
The amplitude of the SZ effect is directly proportional to the “Compton parameter” (y) which is given by where 0 is the projected position from the cluster center. στ is the Thomson cross-section. and {σε2o(yal(GP) is the electron pressure of the ICAL at the 3-cimensioual position ©. | The amplitude of the SZ effect is directly proportional to the “Compton parameter” $y$ ) which is given by where $\theta$ is the projected position from the cluster center, $\sigma_T$ is the Thomson cross-section, and $P_e(\vec{r}) \equiv n_e(\vec{r}) kT_e(\vec{r})$ is the electron pressure of the ICM at the 3-dimensional position $\vec{r}$. |
The iutegral is performed over the line-ofsight (7) through the cluster. | The integral is performed over the line-of-sight $l$ ) through the cluster. |
All of the plivsics of the SZ effect is contained within the Compton parameter. | All of the physics of the SZ effect is contained within the Compton parameter. |
It is the SZ effect analog of the vay surface brightuess of a cluster and is a measure of the average fractional energy. gain of a photon due to iuverse-Compton scattering while passing through a cloud of gas (n this case. the ICM) with an electron pressure profile of PAF). | It is the SZ effect analog of the X-ray surface brightness of a cluster and is a measure of the average fractional energy gain of a photon due to inverse-Compton scattering while passing through a cloud of gas (in this case, the ICM) with an electron pressure profile of $P_e(\vec{r})$. |
As discussed by BBELPO2 aud MDD02. the presence of excess cutropy will modify both a clusters density aud temperature profiles. | As discussed by BBLP02 and MBB02, the presence of excess entropy will modify both a cluster's density and temperature profiles. |
In the case where it is preheating 0.1iu that eives rise to an eutropy core. as iu the present study. the temperature of the gas near the center of the cluster is increased and. therefore. so is the global emission-woeiehted temperature of the cluster (e.9.. Fie. | In the case where it is preheating 0.1in that gives rise to an entropy core, as in the present study, the temperature of the gas near the center of the cluster is increased and, therefore, so is the global emission-weighted temperature of the cluster (e.g., Fig. |
1 of NDBDB02). | 1 of MBB02). |
At the sale time. the density of the eas at the cluster center is dramatically reduced (e.g.. Fig. | At the same time, the density of the gas at the cluster center is dramatically reduced (e.g., Fig. |
2 of NBD02). | 2 of MBB02). |
It turns out that. relatively speaking. preheating las a stronger influence on the density than it docs ou the temperature. at least at the centers of massive clusters. | It turns out that, relatively speaking, preheating has a stronger influence on the density than it does on the temperature, at least at the centers of massive clusters. |
The result is that he eas pressure in central regious of a cluster is reduced by xelieatiug and. cousequeutly. so is the clusters Compton xuanueter. | The result is that the gas pressure in central regions of a cluster is reduced by preheating and, consequently, so is the cluster's Compton parameter. |
To demonstrate this. we plot cluster pressure xofiles (2= 0.2) for several values of the cutropy floor iu Figure 1 CR, is the radius of the cluster). | To demonstrate this, we plot cluster pressure profiles $z = 0.2$ ) for several values of the entropy floor in Figure 1 $R_{halo}$ is the radius of the cluster). |
The addition of an cutropy floor leads to a decrease in the eas pressure acar the claster core. | The addition of an entropy floor leads to a decrease in the gas pressure near the cluster core. |
The eas pressure in the outer regions of the clusters. however. remains relatively uuchauged as he eutropv increase due to eravitational shock heating dominates the nou-eravitational eutropy injection. | The gas pressure in the outer regions of the clusters, however, remains relatively unchanged as the entropy increase due to gravitational shock heating dominates the non-gravitational entropy injection. |
Also of rote is that the difference between the various nodels is greatest for the lower mass cluster. | Also of note is that the difference between the various models is greatest for the lower mass cluster. |
This is expected since the lower mass cluster has a shallower votential well aud. thus. is more stronely mfüuenced by he presence of an entropy floor. | This is expected since the lower mass cluster has a shallower potential well and, thus, is more strongly influenced by the presence of an entropy floor. |
With an cutropy floor significantly affecting the pressure of the ICAL near the center of a cluster. the Comptou waralucter will be most stronely modified if it is evaluated within the smallestpossible projected radius [ic the central Compton parameter. g(0=0) yy. | With an entropy floor significantly affecting the pressure of the ICM near the center of a cluster, the Compton parameter will be most strongly modified if it is evaluated within the smallestpossible projected radius [i.e., the Compton parameter, $y(\theta = 0) \equiv y_0$ ]. |
Tutegrating (or averaging) the Compton parameter within larger projected radi (for example. Rpg. the radius of the | Integrating (or averaging) the Compton parameter within larger projected radii (for example, $R_{halo}$ , the radius of the |
It has been known for niuiv vears that radio pulses frou he Crab pulsar are affected both by a variable delay due to changes im dispersion and by a variable pulse xoadeniues due to scattering along the line of sight (?:: ?)). | It has been known for many years that radio pulses from the Crab pulsar are affected both by a variable delay due to changes in dispersion and by a variable pulse broadening due to scattering along the line of sight \cite{rc73}; ; \cite{ir77}) ). |
Both phenomena vary ou a typical time scale of about 100 days. but iu previous observations their variations lave appeared to be impertectly correlated anc possibly even uneorrelated. | Both phenomena vary on a typical time scale of about 100 days, but in previous observations their variations have appeared to be imperfectly correlated and possibly even uncorrelated. |
At our two observatories. we have naintained for several vears two series of observations o separately monitor these two phenomena. and can iow report a discrete event that shows a remarkably good correlation between variations in scattering and in dispersion measure. | At our two observatories, we have maintained for several years two series of observations to separately monitor these two phenomena, and can now report a discrete event that shows a remarkably good correlation between variations in scattering and in dispersion measure. |
Observations of dispersion iueasure are made at least once a week at Jodrell Bank Observatory as part of the Crab pulsar timine ephemeris which has Όσοι xoduced. aud made generally available since 1982. | Observations of dispersion measure are made at least once a week at Jodrell Bank Observatory as part of the Crab pulsar timing ephemeris which has been produced and made generally available since 1982. |
The ephemeris is based on daily observations of time of arrival of pulses at 610 MIIz. while the dispersion delay is ueasured by comparison with similar observations at 1100 MIIz. | The ephemeris is based on daily observations of time of arrival of pulses at 610 MHz, while the dispersion delay is measured by comparison with similar observations at 1400 MHz. |
Observations at Pusheching Badio Astronomi Observatorv monitoring the pulse shape at 111 MIIz iive continued since 20014. | Observations at Pushchino Radio Astronomy Observatory monitoring the pulse shape at 111 MHz have continued since 2004. |
Both before and durius the event the pulse is broadened with a steep rise aud au approximately exponeutial decay with a time constant of several nuülliseconds: this characteristic decay time is monitored almost daily. | Both before and during the event the pulse is broadened with a steep rise and an approximately exponential decay with a time constant of several milliseconds; this characteristic decay time is monitored almost daily. |
A distinctive property of the Crab pulsar low frequency observations ds that the scatter broadening iav he comparable with or ereater than the pulsar period. | A distinctive property of the Crab pulsar low frequency observations is that the scatter broadening may be comparable with or greater than the pulsar period. |
To avoid the resulting confusion we use for observations the eiut pulses of this pulsar. which stand out of the regular pulses as rare. strong. well defined sinele pulses. | To avoid the resulting confusion we use for observations the giant pulses of this pulsar, which stand out of the regular pulses as rare, strong, well defined single pulses. |
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