Datasets:
ReadingTimeMachine
/
rtm-sgt-ocr-v1

Dataset card Data Studio Files Community
2
source
stringlengths
1
2.05k
⌀
target
stringlengths
1
11.7k
The center has risen a maximum of about 179 pc at 8 Myrs.
The center has risen a maximum of about 179 pc at 8 Myrs.
Without a magnetic field to restrain it, Model A could have been expected to rise more than was observed.
Without a magnetic field to restrain it, Model A could have been expected to rise more than was observed.
Buoyancy,
Buoyancy,
This is a good assumption given that non-nesligible irradiation sets in with the ouset of RLOF aud we have seeu that this phase has lasted for z10*10% yr.
This is a good assumption given that non-negligible irradiation sets in with the onset of RLOF and we have seen that this phase has lasted for $\simless10^7-10^8$ yr.
Therefore any possible donor star has filled the Roche lobe following a standard evolutionary path that docs not involve radiation of the outer envelope.
Therefore any possible donor star has filled the Roche lobe following a standard evolutionary path that does not involve irradiation of the outer envelope.
The effect of radiation müsht instead be important curing the RLOF phase when the outermost lavers of the cuvelope expand after absorbing the energv deposited. by the N-rav radiation.
The effect of irradiation might instead be important during the RLOF phase when the outermost layers of the envelope expand after absorbing the energy deposited by the X-ray radiation.
The star beime coustraimed to fit the Roche lobe will increase its mass loss through the ner Lagrangian poit aud boost the mass transfer rate.
The star being constrained to fit the Roche lobe will increase its mass loss through the inner Lagrangian point and boost the mass transfer rate.
The average X-ray power absorbed bv the companion during zu outburst is with £L,~L«LPerest being the average X-rav huuinosity during an outburst.
The average X-ray power absorbed by the companion during an outburst is with $\langle\,L_{x}\rangle\sim4\times10^{37}\rm\,erg\,s^{-1}$ being the average X-ray luminosity during an outburst.
Given a duty cycle of Az0.005. the long term average power absorbed bv the donor will be z10°Loebores1.
Given a duty cycle of $\Delta\simless0.005$, the long term average power absorbed by the donor will be $\simless\,10^{33}-10^{34}\rm\,erg\,s^{-1}$.
This value corresponds to S10%=LOO% of the outgoing power produced by a low mass stars in the sub-giaut brauch aud cannot be considered negligible.
This value corresponds to $\simless10\%-100\%$ of the outgoing power produced by a low mass stars in the sub-giant branch and cannot be considered negligible.
The quantitative effect of irradiation can be calculated with stellar evolutionary codes aud is bevoud the scope of this work.
The quantitative effect of irradiation can be calculated with stellar evolutionary codes and is beyond the scope of this work.
We have investigated the evolutionary listory of discussing three distinct evolutionary epochs: a dipole cominated spin-down epoch. a wind epoch and the currently observed RLOF phase.
We have investigated the evolutionary history of discussing three distinct evolutionary epochs: a dipole dominated spin-down epoch, a wind epoch and the currently observed RLOF phase.
By takine iuto account the effect of timune noise im the N-rav pulsations we are able to confirm the loue-term average spiucup of 1.1:5101?IIzs.! observed during the 2010 outburst (Papittoetal.2011:Cavecclict 2011)).
By taking into account the effect of timing noise in the X-ray pulsations we are able to confirm the long-term average spin-up of $1.4\times10^{-12}\rm\,Hz\,s^{-1}$ observed during the 2010 outburst \citealt{pap11,cav11}) ).
If the N-vav flux represents most of the bolometric fiux. the short term spin frequency derivatives 7 do not scale with the N-rav flux in the expected wav.
If the X-ray flux represents most of the bolometric flux, the short term spin frequency derivatives $\dot{\nu}$ do not scale with the X-ray flux in the expected way.
The long-teiu D. however. leads to exclude a scenario in which has already reached spin equilibrimu.
The long-term $\dot{\nu}$, however, leads to exclude a scenario in which has already reached spin equilibrium.
A long-term spin equilibrium with magnetic dipole spin down in quiescence balancing the accretion torques iu outburst is also excluded based ou binary evolution considerations.
A long-term spin equilibrium with magnetic dipole spin down in quiescence balancing the accretion torques in outburst is also excluded based on binary evolution considerations.
The spiu-up timescale suggests that has been spuu-up iu the RLOF phase for a ΠΠfew τ1δ10* vr.
The spin-up timescale suggests that has been spun-up in the RLOF phase for a few $10^{7}$ yr.
Àn interpretation of the observed QPOs as the equilibrimm frequencies of indicates a siluilar timescale for the spiu-up epoch up to the present.
An interpretation of the observed QPOs as the equilibrium frequencies of indicates a similar timescale for the spin-up epoch up to the present.
The spin-down epochs suggest that the NS in had reached a spin frequency of the order of Uz at the onset of the RLOF phase. if we assume that the newborn NSs had typical initial values for Qy and 449.
The spin-down epochs suggest that the NS in had reached a spin frequency of the order of $\sim1$ Hz at the onset of the RLOF phase, if we assume that the newborn NSs had typical initial values for $\Omega_0$ and $\mu_0$.
This reinforces the idea that the time elapsed in RLOF has been very short. since the time required to spiu-up al Uz NS to 11 IIz with a spin frequency derivative of 10Πες aud a duty evele of 0.01 is a few 10 vr.
This reinforces the idea that the time elapsed in RLOF has been very short, since the time required to spin-up a 1 Hz NS to 11 Hz with a spin frequency derivative of $10^{-12}\rm\,Hz\,s^{-1}$ and a duty cycle of 0.01 is a few $10^7$ yr.
Based on these findings we conclude that is in an exceptionally carly RLOF phase.
Based on these findings we conclude that is in an exceptionally early RLOF phase.
The total spin-up timescale to transform from a slow pulsar ( 2) iuto a millisecond one (vc LOOTIz) ix 210%LO? ve,
The total spin-up timescale to transform from a slow pulsar $\sim\,1\rm\,Hz$ ) into a millisecond one $\nu>100\rm\,Hz$ ) is $\simgreat10^8-10^{9}$ yr.
Since the current RLOF phase will last for ~109yy. we would expect to observe today ~110 accreting millisecoud pulsars that have followed a simular evolutionary listory as.
Since the current RLOF phase will last for $\sim10^9$ yr, we would expect to observe today $\sim1-10$ accreting millisecond pulsars that have followed a similar evolutionary history as.
. The accreting millisecond pulsar SAN J1T15.9-2021 (vz1121Iz) in the elobular cluster NGC GLLO has a companion with mass and orbital period conrpatible with beime in a slightly evolved. post main sequence phase (Altaimiranoetal.2008).
The accreting millisecond pulsar SAX J1748.9-2021 $\nu\simeq442\rm\,Hz$ ) in the globular cluster NGC 6440 has a companion with mass and orbital period compatible with being in a slightly evolved post main sequence phase \citep{alt08}.
. It is likely hat the companion of SAN J17I8.9-2021 has followed an evolutionary history simular to with he difference that it spent a louger time in the RLOF yhase thus turning its NS iuto a imillisecond pulsar.
It is likely that the companion of SAX J1748.9-2021 has followed an evolutionary history similar to with the difference that it spent a longer time in the RLOF phase thus turning its NS into a millisecond pulsar.
The scenario outlined above is compatible with the observed »pulatiou of accreting pulsars. although it is dithicult ο assess the problem in a robust wav elven the small Munber statistics.
The scenario outlined above is compatible with the observed population of accreting pulsars, although it is difficult to assess the problem in a robust way given the small number statistics.
It is also possible that ανασα] interactions have plaved a role by decreasing the number of simular svstenis in elobulay clusters via ionization interactions.
It is also possible that dynamical interactions have played a role by decreasing the number of similar systems in globular clusters via ionization interactions.
This however requires further mvestieations along with detailed binary evolution calculations.
This however requires further investigations along with detailed binary evolution calculations.
The prior spin-down history gives results which are not compatible with a binary having an age of several 10? yr as expected if is prinordial or if it has not suffered exchange interactions iu the last few 10? vr.
The prior spin-down history gives results which are not compatible with a binary having an age of several $10^9$ yr as expected if is primordial or if it has not suffered exchange interactions in the last few $10^9$ yr.
Different formation scenarios are available to explain the apparcut discrepancy between the age of the cluster and the binary.
Different formation scenarios are available to explain the apparent discrepancy between the age of the cluster and the binary.
À first possibility is that a recent dyvnauical counter has plaved a role in forming the binary or in accelerating the ouset of the RLOF phase.
A first possibility is that a recent dynamical encounter has played a role in forming the binary or in accelerating the onset of the RLOF phase.
If an exchange interaction has taken place in the last few Loo105 ye then this would explain why the appareut age of is so short compared with the age of the luster.
If an exchange interaction has taken place in the last few $10^7-10^8$ yr then this would explain why the apparent age of is so short compared with the age of the cluster.
The location of in the high mass and hieh central density cluster Terzan 5 suggests that the interaction rate nüght be particularly high iu this enviromment.
The location of in the high mass and high central density cluster Terzan 5 suggests that the interaction rate might be particularly high in this environment.
Dynamical simulations to calculate the typical interaction rates in the core of Terzan 5 are required to assess this question.
Dynamical simulations to calculate the typical interaction rates in the core of Terzan 5 are required to assess this question.
A second possibility for the origin of is formation of the accreting pulsar via accretion induced collapse (AIC) of a massive white dwarf (ALivajietal. 1980... Nomoto 1987)).
A second possibility for the origin of is formation of the accreting pulsar via accretion induced collapse (AIC) of a massive white dwarf \citealt{miy80}, \citealt{nom87}) ).
Iu. this scenario a ONeMg white dwarf accretes matter until its mass exceeds the Chandrasekhar Παπά aud it collapses to forma an NS via electron captures on Mg aud Ne unclei.
In this scenario a ONeMg white dwarf accretes matter until its mass exceeds the Chandrasekhar limit and it collapses to form an NS via electron captures on Mg and Ne nuclei.
Iu this case the binary nist have gone through a preliminary coutact phase during which the white dwarf was accreting from a donor star in RLOF.
In this case the binary must have gone through a preliminary contact phase during which the white dwarf was accreting from a donor star in RLOF.
During the collapse approximately AZ... are ejected from the binary (see for example vandenHeuvel2011 for a discussion) which causes a sudden expansion of the orbit turning the svstem mto a detached binary.
During the collapse approximately $\msun$ are ejected from the binary (see for example \citealt{van11} for a discussion) which causes a sudden expansion of the orbit turning the system into a detached binary.
At this point the binary follows the three evolutionary epochs described in Section ?? with the NS age set by the ouset of AIC.
At this point the binary follows the three evolutionary epochs described in Section \ref{evolution} with the NS age set by the onset of AIC.
The formation of NSs from AIC has been discussed by Lyuectal.(1996). to explain the voung radio pulsa population iu globular clusters.
The formation of NSs from AIC has been discussed by \citet{lyn96} to explain the young radio pulsar population in globular clusters.
Lyne et al.
Lyne et al.
note that the formation rate of vouus pulsars iu elobular clusters is larecr than that of millisecond pulsus iu elobulay clusters.
note that the formation rate of young pulsars in globular clusters is larger than that of millisecond pulsars in globular clusters.
The voung pulsus are found in iietal rich globular clusters with biel mass and lieh central densitics. like Terzan 5 (Laneetal. 2011)).
The young pulsars are found in metal rich globular clusters with high mass and high central densities, like Terzan 5 \citealt{lyn96,boy11}) ).
Further evidence that elobular clusters nüeht contain vouug NSs was also preseuted by
Further evidence that globular clusters might contain young NSs was also presented by
↕↴∖↴↴∖↴∐∶↴∙⊾∐↑↕⋅↖↽↕∐∶↴⋁∐↸∖↥⋅↑∐⋜⋯↕⊔∪↴∖↴↑↸∖⋜∐⋅∐↸∖↥⋅↸∖↴∖↴↑↕⋯⋜↧↑↸∖↴∖↴↕≯↥⋅∪⋯↑∐↸∖ literature (although those are takeu at higher frequeucies. see Sect. 3.1).
is slightly higher than most earlier estimates from the literature (although those are taken at higher frequencies, see Sect. \ref{s9:lit_ps}) ).
llowever. note that the Auriga and Ilorologiun regions were selected for their conspicuous structure in P. so we expect these regious to show more structure ou large (degree) scales than the "average" ISM. aud thus exhibit a steeper spectimm.
However, note that the Auriga and Horologium regions were selected for their conspicuous structure in $P$, so we expect these regions to show more structure on large (degree) scales than the “average” ISM, and thus exhibit a steeper spectrum.
The power spectra of Q aud U in the Auriga region are somewhat stecper than iu Uorologimm. indicating that the Iorologiuu reeion probably contains more sinall-scale structure in the Faraday screen than the Auriga region.
The power spectra of $Q$ and $U$ in the Auriga region are somewhat steeper than in Horologium, indicating that the Horologium region probably contains more small-scale structure in the Faraday screen than the Auriga region.
The power spectra of RAL in Fig.
The power spectra of $RM$ in Fig.
5 are shallower thau the Q. © or P power spectra.
\ref{f9:ps_p_obs} are shallower than the $Q$, $U$ or $P$ power spectra.
In fact. we do not expect a direct correspoudence between the multipole spectral indices of RAL and P (or Q. U). as the former describes very directly the electron content and maguetic feld iu the ISAL (integrated over the line of sight). whereas in the latter case the polarized radiation is modulated by Faraday rotation and depolarization.
In fact, we do not expect a direct correspondence between the multipole spectral indices of $RM$ and $P$ (or $Q$, $U$ ), as the former describes very directly the electron content and magnetic field in the ISM (integrated over the line of sight), whereas in the latter case the polarized radiation is modulated by Faraday rotation and depolarization.
Iu the WENSS polarization region. power spectra were evaluated for subfields. to study possible dependences of the multipole spectral iudex on Galactic longitude and/or latitude.
In the WENSS polarization region, power spectra were evaluated for subfields, to study possible dependences of the multipole spectral index on Galactic longitude and/or latitude.
The 11 subfields are shown in Fie. 9..
The 11 subfields are shown in Fig. \ref{f9:wenssboxes}, ,
superimposed on ereyv scale maps of P.
superimposed on grey scale maps of $P$.
The power spectra of polarized intensity P are shown in Fig. 10..
The power spectra of polarized intensity $P$ are shown in Fig. \ref{f9:ps_wr_pi},
where the subfields are arranged as in Fig. 9..
where the subfields are arranged as in Fig. \ref{f9:wenssboxes}.
The power spectra iu subfields 9. 10 and 11. at high Galactic latitude b. have a lower amplitude than the power spectra at lower 5. which is consistent with the decreasing amount of P at higher 5 visible in Fie. 3..
The power spectra in subfields 9, 10 and 11, at high Galactic latitude $b$, have a lower amplitude than the power spectra at lower $b$, which is consistent with the decreasing amount of $P$ at higher $b$ visible in Fig. \ref{f9:wenss}.
The multipole spectral indices of the power spectra of P. (Q and C are given in Table 2.. and the depeudenuce of ap on Calactic longitude aud latitude is) shown in Fig. 11..
The multipole spectral indices of the power spectra of $P$, $Q$ and $U$ are given in Table \ref{t9:a_wenss}, and the dependence of $\alpha_P$ on Galactic longitude and latitude is shown in Fig. \ref{f9:ps_wr_lb}.
The observed decrease of spectral iudex with increasing latitude ppower spectra become flatter with increasing latitude) indicates a decreasein the amount of large-scale structure with increasing latitude.
The observed decrease of spectral index with increasing latitude power spectra become flatter with increasing latitude) indicates a decreasein the amount of large-scale structure with increasing latitude.
The dependence of spectral iudex on longitude
The dependence of spectral index on longitude
2 at 40 GIlz for the northern source is CO(2-1) line emission (Figs.
2 at 40 GHz for the northern source is CO(2-1) line emission (Figs.
la.b.c). and hence that the velocity profiles for the CO Lines [rom the northern and southern sources are clilferent.
1a,b,c), and hence that the velocity profiles for the CO lines from the northern and southern sources are different.
Figure 2 shows the results from the B array observations with a resolution of about 0.25".
Figure 2 shows the results from the B array observations with a resolution of about $0.25''$.
The southern source in IF1 (Fig 2b) appears as two unresolved. roughly equal components (Ilux densities ~0.41£0.12 mJ»).separated by 0.3".
The southern source in IF1 (Fig 2b) appears as two unresolved, roughly equal components (flux densities $\sim 0.41\pm0.12$ mJy),separated by $''$.
The implied lower limit to the (redshilt corrected) brightness temperature for these components is about 25 kx. The northern source in IFs 1 and 2 (Fig 2a) is marginally detected: with a peak surface brightness of 0.2530.08 mJv +.
The implied lower limit to the (redshift corrected) brightness temperature for these components is about 25 K. The northern source in IFs 1 and 2 (Fig 2a) is marginally detected with a peak surface brightness of $0.25 \pm 0.08$ mJy $^{-1}$.
This implies two. or more. Compact components with flux densities below our detection threshold. or diffuse emission on a scale >0.5".
This implies two, or more, compact components with flux densities below our detection threshold, or diffuse emission on a scale $\ge 0.5''$.
Figure 3 shows the radio continuum image of 12020725 at 1.4 GlIz. with a resolution ol about 2".
Figure 3 shows the radio continuum image of 1202–0725 at 1.4 GHz, with a resolution of about $2''$.
Both components are detected. with flux densities as listed in Table 2.
Both components are detected, with flux densities as listed in Table 2.
The total 1.4 Gllz flux density of 12020725 has been measured three times over the course of three vears with values of 240440 pJv (Yun et al.
The total 1.4 GHz flux density of 1202–0725 has been measured three times over the course of three years with values of $240\pm40$ $\mu$ Jy (Yun et al.
2000). 305+60. μ.]ν. and 390+40 piJv.
2000), $305\pm60$ $\mu$ Jy, and $390\pm40$ $\mu$ Jy.
In our analvsis we adopt the mean value of 315 μ.]ν.
In our analysis we adopt the mean value of 315 $\mu$ Jy.
Ht is possible that the source is variable. although making accurate (ix density measurements al this level is difficult due to confusion problems arising in wide-field imaging at 1.4 GHz.
It is possible that the source is variable, although making accurate flux density measurements at this level is difficult due to confusion problems arising in wide-field imaging at 1.4 GHz.
The low resolution image of the CO(2-1) emission from 13350417. as reproduced [rom Carilli et al. (
The low resolution image of the CO(2-1) emission from 1335–0417, as reproduced from Carilli et al. (
1999). is shown in Figure ta. along with the off-line channel in Figure th.
1999), is shown in Figure 4a, along with the off-line channel in Figure 4b.
For this image the UV-data were naturally weighted in order (o maximize sensitivity. at the expense of resolution.
For this image the UV-data were naturally weighted in order to maximize sensitivity, at the expense of resolution.
The rms noise on the image is 0.11 mJv 1 and the resolution is about 1.6".
The rms noise on the image is 0.11 mJy $^{-1}$ and the resolution is about $1.6''$.
The CO emission appears extended north-south in this image by about 1".
The CO emission appears extended north-south in this image by about $''$ .
In order to investigate this extension in more detail. we re-imaged the data using
In order to investigate this extension in more detail, we re-imaged the data using
uminosity bins. there are huge sample-to-sample variations with. or example. the ERS sample delivering very blue values of (7 in the Ariipc19.5 bin at 2—Y compared to either of he deeper HUDF and HUDFO9-? samples.
luminosity bins, there are huge sample-to-sample variations with, for example, the ERS sample delivering very blue values of $\langle \beta \rangle$ in the $M_{UV,AB} \simeq -19.5$ bin at $z \simeq 7$ compared to either of the deeper HUDF and HUDF09-2 samples.
This is basically the effect of the plume to low values of ή seen at the ERS flux limit as shown in Fig.
This is basically the effect of the plume to low values of $\beta$ seen at the ERS flux limit as shown in Fig.
4.
4.
To reconcile the results from the different samples in the luminosity bins in which they overlap. we found it necessary to Insist on a minimum signal:noise ratio requirement.
To reconcile the results from the different samples in the luminosity bins in which they overlap, we found it necessary to insist on a minimum signal:noise ratio requirement.
To avoid introducing any further colour bias we simply chose to insist that every object retained in the final. retined sample was detected in at least one WFC3/IR near-infrared passband at a minimum level of σ.
To avoid introducing any further colour bias we simply chose to insist that every object retained in the final, refined sample was detected in at least one WFC3/IR near-infrared passband at a minimum level of $\sigma$.
The impact of this further level of quality control is then shown in the bottom row of plots in Fig.
The impact of this further level of quality control is then shown in the bottom row of plots in Fig.
5.
5.
No longer do the different samples deliver substantially different average values of (22. and it can be seen that the very low values of 1:2? were indeed largely resulting from the lowest signal:noise ratio sources.
No longer do the different samples deliver substantially different average values of $\langle \beta \rangle$, and it can be seen that the very low values of $\langle \beta \rangle$ were indeed largely resulting from the lowest signal:noise ratio sources.
Importantly. with this level of further quality control. we are left with only one HUDF object in the Miiig218.5 bin at >&7. and hence cannot plot a meaningful average value of (6:25.
Importantly, with this level of further quality control, we are left with only one HUDF object in the $M_{UV,AB} \simeq -18.5$ bin at $z \simeq 7$, and hence cannot plot a meaningful average value of $\langle \beta \rangle$.
At >=6. where we can still probe this luminosity bin. the evidence for (.7) being significantly bluer than 1/2?=—2 has disappeared.
At $z = 6$, where we can still probe this luminosity bin, the evidence for $\langle \beta \rangle$ being significantly bluer than $\langle \beta \rangle = -2$ has disappeared.
This final result is summarized in Fig.
This final result is summarized in Fig.
6. where we overplot he dependence of 6:7) on Mirag as a function of redshift.
6, where we overplot the dependence of $\langle \beta \rangle$ on $M_{UV,AB}$ as a function of redshift.
The derived datapoints shown in Fig.
The derived datapoints shown in Fig.
6 are tabulated in Table 1.
6 are tabulated in Table 1.
Clearly. hese results are consistent with (3)=—2 over the full redshift and luminosity range which can be probed with these data.
Clearly, these results are consistent with $\langle \beta \rangle = -2$ over the full redshift and luminosity range which can be probed with these data.
At the bright end they are also in good agreement with the results derived by both Bouwens et al. (
At the bright end they are also in good agreement with the results derived by both Bouwens et al. (
2010b) and Finkelstein et al. (
2010b) and Finkelstein et al. (
2010). so any disagreement is really contined to My:720.
2010), so any disagreement is really confined to $M_{UV} > -20$.
Of course. it might be argued that by insisting on rejecting he lowest signal:noise ratio sources. we have effectively "thrown away" the "evidence" for how .7 behaves at the faintest luminosities
Of course, it might be argued that by insisting on rejecting the lowest signal:noise ratio sources, we have effectively “thrown away” the “evidence” for how $\beta$ behaves at the faintest luminosities
in (he timing residuals of BIS28—11 that is accompanied by. correlated pulse shape changes is explained by free-precession of the neutron star (Stairsetal.2000).
in the timing residuals of $-$ 11 that is accompanied by correlated pulse shape changes is explained by free-precession of the neutron star \citep{sta00}.
. In the case of PSR 1642—023. multiple low frequency components in the power spectrum of timing residuals can be explained well as a result of continuous generation of slow glitches. (he amplitude of which is correlated with the time interval following the gliteh (Shabanova2009h)..
In the case of PSR $-$ 03, multiple low frequency components in the power spectrum of timing residuals can be explained well as a result of continuous generation of slow glitches, the amplitude of which is correlated with the time interval following the glitch \citep{sha09b}. .
As reported by Yuanetal.(2010).. the slow glitehes have been identified vet [or two pulsars J0G31--1036 and BLOOT+10.
As reported by \citet{yua10}, the slow glitches have been identified yet for two pulsars J0631+1036 and B1907+10.
The first pulsar is a voung pulsar with 7~4.4x10! years. and the second one is substantially older with 7~L.7x109 vears.
The first pulsar is a young pulsar with $\tau \sim 4.4\times 10^{4}$ years, and the second one is substantially older with $\tau \sim 1.7\times 10^{6}$ years.