source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
However. N-rays. could. still increase the ionisation fraction. and hence the associated ellective o. in the active laver. | However, X-rays could still increase the ionisation fraction, and hence the associated effective $\alpha$, in the active layer. |
Igea&Classgold(1999). showed that the X-ray ionisation rate near the surface is significantly larger than that of cosmic ravs. | \cite{igea99} showed that the X-ray ionisation rate near the surface is significantly larger than that of cosmic rays. |
The accretion rate through the laver may be higher than predieted in equation (29)) if this is taken into account and perhaps could help to explain the Tauri accretion rates for higher critical magnetic Ievnolds number. | The accretion rate through the layer may be higher than predicted in equation \ref{mdoteq}) ) if this is taken into account and perhaps could help to explain the T Tauri accretion rates for higher critical magnetic Reynolds number. |
Observed EU Orionis svstems may be one manifestation of the time-dependent aceretion phenomenon. | Observed FU Orionis systems may be one manifestation of the time-dependent accretion phenomenon. |
These luminous voung stellar objects are found in star [orming regions (Llartmann&Wenvon1996). | These luminous young stellar objects are found in star forming regions \citep{hartmann96}. |
.. Their optical brightness increases by live magnitudes or more on a timescale of vears and decays on a timescale of 50-100 vears (Llerbig1977). | Their optical brightness increases by five magnitudes or more on a timescale of years and decays on a timescale of 50-100 years \citep{herbig77}. |
. The timescale and. brightness. of accretion outbursts from the eravo-magneto instability are similar to the observed. EU. Orionis outbursts (Armitage.Livio&Pringle2001:Zhuetal. 2009b). | The timescale and brightness of accretion outbursts from the gravo-magneto instability are similar to the observed FU Orionis outbursts \citep{armitage01,zhu09b}. |
. Time-dependent numerical simulations are needed to investigate these effects more fully in a disc with a dead zone determined by the critical magnetic Ievnolds number. | Time-dependent numerical simulations are needed to investigate these effects more fully in a disc with a dead zone determined by the critical magnetic Reynolds number. |
By comparing the resulting outbursts with FU Orionis observations the value of the critical magnetic Revnolds number may be further constrained. | By comparing the resulting outbursts with FU Orionis observations the value of the critical magnetic Reynolds number may be further constrained. |
A typical assumption used in time-dependent simulations of accretion disces with dead zones is that the surface density in the MIU active laver above the dead zone is constant with radius. | A typical assumption used in time-dependent simulations of accretion discs with dead zones is that the surface density in the MRI active layer above the dead zone is constant with radius. |
However. when the dead zone is identified by the value of the eritical magnetic Revnolcls number including the elfects of recombination. the active laver surface density is eencrally found to increase with radius. | However, when the dead zone is identified by the value of the critical magnetic Reynolds number including the effects of recombination, the active layer surface density is generally found to increase with radius. |
The constant laver is best reproduced. with a low critical magnetic Itevnolds number. | The constant layer is best reproduced with a low critical magnetic Reynolds number. |
llowever. MIID. simulations suggest the critical magnetic ltevnolds number may be much higher. | However, MHD simulations suggest the critical magnetic Reynolds number may be much higher. |
For higher critical magnetic Revnolds numbers. estosX100 we have found an analytical fit to the active surface density. (equations 27 and 28)) that will be a useful approximation in future time-dependent. calculations. | For higher critical magnetic Reynolds numbers, $Re_{\rm M,crit}\gtrsim 100$ we have found an analytical fit to the active surface density (equations \ref{fit} and \ref{fit2}) ) that will be a useful approximation in future time-dependent calculations. |
The dead: zone structure is very sensitive to the value of the critical magnetic Itevnolds number. as scen in Fig. 4.. | The dead zone structure is very sensitive to the value of the critical magnetic Reynolds number, as seen in Fig. \ref{layer}, |
but the value is still uncertain and needs to be clarified in further work. | but the value is still uncertain and needs to be clarified in further work. |
However. this is complicated by the fact that the a parameter in the viscosity is also uncertain. | However, this is complicated by the fact that the $\alpha$ parameter in the viscosity is also uncertain. |
The metallicity variation between our galaxy. the LAIC and the SAIC is not significant enough to alfect the expected size of the dead. zone in a disc (ignoring possible dillerences in dust abundances). | The metallicity variation between our galaxy, the LMC and the SMC is not significant enough to affect the expected size of the dead zone in a disc (ignoring possible differences in dust abundances). |
When comparing accretion disces in our galaxy with those in the LMC and SMC for example. the high metallicity limit will be appropriate. | When comparing accretion discs in our galaxy with those in the LMC and SMC for example, the high metallicity limit will be appropriate. |
In order to explain the observed cilference in disc lifetimes with metallicity (e.g.DeMarchi.Panagia&Itomaniello2010.2011). some other elfect must be taken into account. | In order to explain the observed difference in disc lifetimes with metallicity \citep[e.g.][]{demarchi10,demarchi11} some other effect must be taken into account. |
We acknowledge useful. comments from the anonymous referee. | We acknowledge useful comments from the anonymous referee. |
RGM thanks the Space Telescope Science Institute [or a Giaceoni Fellowship. | RGM thanks the Space Telescope Science Institute for a Giacconi Fellowship. |
SUL acknowledges support from NASA erant NNXOYTALT2CG. JEP thanks the Collaborative Visitor Program at STScl for its support and hospitality. | SHL acknowledges support from NASA grant NNX07AI72G. JEP thanks the Collaborative Visitor Program at STScI for its support and hospitality. |
eas. sugeesting another form of heating is required to allow (he removal of eas from clwarls with larger perigalacticons (e.g.(heperigalacticonof250kpeLuxοἱal. 2010). | gas, suggesting another form of heating is required to allow the removal of gas from dwarfs with larger perigalacticons \citep[e.g. the perigalacticon of Sextans and Draco is expected to be $>50$~kpc][]{Lux2010}. |
. A dwarl consisting of cold. dense gas is resistant. to tidal and ram-pressure stripping with only a thin skin. ionized bv the Galactic of extragalactie UV fields being removed. | A dwarf consisting of cold, dense gas is resistant to tidal and ram-pressure stripping with only a thin skin, ionized by the Galactic of extragalactic UV fields being removed. |
Early star formation however. will heat this cold gas. raising it in the potential well aud making it more easily stripped. | Early star formation however, will heat this cold gas, raising it in the potential well and making it more easily stripped. |
This warm ionized gas. at the same pressure as cold gas will occupy ~60 Gines the volume. and being much less dense will become even easier {ο strip than its height in the potential well would indicate. | This warm ionized gas, at the same pressure as cold gas will occupy $\sim60$ times the volume, and being much less dense will become even easier to strip than its height in the potential well would indicate. |
The star formation in dwarfs is considered to consist of periods of low-level star formation. during which short bursts are induced which increase (he star formation by a factor of 3. consistent wilh clwarls observed by Leeetal.(2009). | The star formation in dwarfs is considered to consist of periods of low-level star formation, during which short bursts are induced which increase the star formation by a factor of $3$, consistent with dwarfs observed by \citet{Lee2009}. |
.. We assume these bursts are triggered bv perigalacticon passages. induced by shocks created through tidal interactions. elal. 2010). | We assume these bursts are triggered by perigalacticon passages, induced by shocks created through tidal interactions \citep{Pasetto2010}. |
. There is also evidence (hat bursts may. he triggered by the re-accretion of heated. and expanded gas (Valekeetal.2008).. however. much of (his gas will be stripped while in the warm phase and [ar from the centre preventing the re-accretion and subsequent starburst [rom occurring. | There is also evidence that bursts may be triggered by the re-accretion of heated and expanded gas \citep{Valcke2008}, however, much of this gas will be stripped while in the warm phase and far from the centre preventing the re-accretion and subsequent starburst from occurring. |
The base star formation rate is taken to be similar to the dwarfs that surround. M31. with a star formation rate (IxXaisin&Narachentsey2006) where / is a constant calculated by assuming that the gas will be completely depleted αἱ a (me /=5/1,~70 Gyr. | The base star formation rate is taken to be similar to the dwarfs that surround M31, with a star formation rate \citep{Kaisin2006}
where $k$ is a constant calculated by assuming that the gas will be completely depleted at a time $t=5/H_0\sim70~$ Gyr. |
For an initial gas mass of 5x101 ALL. /=—13.62. | For an initial gas mass of $5\times10^7$ $_\odot$, $k=-13.62$. |
The UV.X-ray spectrum produced by the star formation in the dwarls—which is responsible for the majority of heatingwas calculated with (Leithererοἱal.1999): when a burst occured the change in the spectrum was caleulated bx summing many short bursts together to produce an approximately continuous change in the spectrum. | The UV–X-ray spectrum produced by the star formation in the dwarfs—which is responsible for the majority of heating—was calculated with \citep{Leitherer1999}; when a burst occured the change in the spectrum was calculated by summing many short bursts together to produce an approximately continuous change in the spectrum. |
Above the Lyman limit.. (his. spectrum Gin. eres sH H > ©loge 1) and its. changes during. a burst. was well fitted by | Above the Lyman limit, this spectrum (in ergs $s^{-1}$ $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ ) and its changes during a burst was well fitted by |
which lines-o[-sight intersect a cluster with a velocity dispersion large enough to plausibly account for the weak lensing peak (Section 4.2)). | which lines-of-sight intersect a cluster with a velocity dispersion large enough to plausibly account for the weak lensing peak (Section \ref{clusters}) ). |
In this section we evaluate candidate svstems along the line-of-sight toward the 6 robust GTO2dee? weak lensing peaks of the revised Subaru weak lensing map (Section 3.1)). | In this section we evaluate candidate systems along the line-of-sight toward the 6 robust $^2$ weak lensing peaks of the revised Subaru weak lensing map (Section \ref{newmap}) ). |
Figures 1l and 12 show the redshift z distribution within 3’ of the central position of each of the candidate halos. | Figures \ref{fig:velhisto1.ps} and \ref{fig:velhisto2.ps} show the redshift $z$ distribution within $^\prime$ of the central position of each of the candidate halos. |
The dark histogram shows the redshilt distribution in bius of 0.002(14-: ). | The dark histogram shows the redshift distribution in bins of $z$ ). |
The thin histogram shows the redshift distribution in a concentric cone with a 6 radius in each of the candidate halo directions. | The thin histogram shows the redshift distribution in a concentric cone with a $6^\prime$ radius in each of the candidate halo directions. |
The probes toward peaks 0 and 1 each contain an impressive peak. | The probes toward peaks 0 and 1 each contain an impressive peak. |
These peaks correspond to obvious “fingers” in redshift space (Figure 11)). | These peaks correspond to obvious “fingers” in redshift space (Figure \ref{fig:velhisto1.ps}) ). |
The rest frame line-ol-sieht velocity dispersions of these svstems within a 6 probe are. respectively 576464 kins | and 598458 kn 1 in agreement with previous measures (Dressler οἱ al 1999 (peak 0): Hamana οἱ al. | The rest frame line-of-sight velocity dispersions of these systems within a $^\prime$ probe are, respectively $\pm$ 64 km $^{-1}$ and $\pm$ 58 km $^{-1}$ in agreement with previous measures (Dressler et al 1999 (peak 0); Hamana et al. |
2008 (peak 1): see Section 2.1)). | 2008 (peak 1); see Section \ref{history}) ). |
Table 3. provides the number of galaxies ancl (he rest frame line-of-sight velocity dispersion for both 3' and 6' probes. | Table \ref{tbl:VDisp} provides the number of galaxies and the rest frame line-of-sight velocity dispersion for both $^{\prime}$ and $^{\prime}$ probes. |
The rest. frame velocity dispersions lor these (wo clusters places both of (hem very near the v=3.7 detection threshold. | The rest frame line-of-sight velocity dispersions for these two clusters places both of them very near the $\nu = 3.7$ detection threshold. |
Counts of faint galaxy. within these probes show a >Poy excess as one might expect (see Table 3)). | Counts of faint galaxy within these probes show a $> 2\sigma_S$ excess as one might expect (see Table \ref{tbl:VDisp}) ). |
We note that the weak lensing map detects the clusters corresponding to peaks 0 and 1 at hieh significance ν» 6. | We note that the weak lensing map detects the clusters corresponding to peaks 0 and 1 at high significance $\nu \sim 6$ . |
IIowever lor ν~6. our predicted velocity clispersions for these peaks exceed the measured. values by only2 | However for $\nu \sim 6$, our predicted velocity dispersions for these peaks exceed the measured values by only. |
0-30%.. The lo errors in the velocity dispersions are =10%. | The $\sigma$ errors in the velocity dispersions are $\gtrsim 10$. |
. Thus we cannot draw any definitive conclusion for this apparent difference between the measured. velocity. dispersion from (he redshift survey and the weak lensing sensitivity analvsis. | Thus we cannot draw any definitive conclusion for this apparent difference between the measured velocity dispersion from the redshift survey and the weak lensing sensitivity analysis. |
The probe toward weak lensing peak 3 shows a more complex situation. | The probe toward weak lensing peak 3 shows a more complex situation. |
There are two significant peaks in the histogram. one al 2=0.297 and the other at 2=0.673. | There are two significant peaks in the histogram, one at $z = 0.297$ and the other at $z = 0.673$. |
The velocity dispersions are 567 km | and 558 km 1+. respectively within 3! probes: the rest. [rame line-of-sight velocity dispersions ave svstematically smaller in the 6 probes. but (he errors are large. | The line-of-sight velocity dispersions are 567 km $^{-1}$ and 558 km $^{-1}$, respectively within $^{\prime}$ probes; the rest frame line-of-sight velocity dispersions are systematically smaller in the $^{\prime}$ probes, but the errors are large. |
The faint galaxy. counts show an excess (Table 3)). probably resulting from the more distant svstem. | The faint galaxy counts show an excess (Table \ref{tbl:VDisp}) ), probably resulting from the more distant system. |
Both svstems max contribute to the lensing signal. | Both systems may contribute to the lensing signal. |
Probes toward weak lensine peaks 7 and 9 reveal no candidate svstems. | Probes toward weak lensing peaks 7 and 9 reveal no candidate systems. |
Figure shows (hat neither of these peaks overlaps awell-populated probe through the reclshilt survey. | Figure \ref{fig:sigmamap.scaled.ps} shows that neither of these peaks overlaps awell-populated probe through the redshift survey. |
Furthermore. (here is no excess in (he faint galaxy. counts. | Furthermore, there is no excess in the faint galaxy counts. |
One possibility is (hat the | One possibility is that the |
Now. in the case of a spatially fat FRW universe with no cosmological coustaut (only). the αναο equations (5)) allow us to replace 8zGp/23 with (6/a). | Now, in the case of a spatially flat FRW universe with no cosmological constant (only), the dynamic equations \ref{eq:reldyn}) ) allow us to replace $8 \pi G \rho/3 $ with $(\dot{a}/a)^2$. |
This gives us and we recover the formmla for the cosinological redshift. | This gives us and we recover the formula for the cosmological redshift. |
Iu this oue case the cosmological redshift can be broken down iuto a factor of velocity aloue aud one of static eravitation alone. | In this one case the cosmological redshift can be broken down into a factor of velocity alone and one of static gravitation alone. |
We may in this case attribute its cause to a simple combination of both static mass andmotion’. | We may in this case attribute its cause to a simple combination of both static mass and. |
.. There are several reasons to be uneasy over this derivation. | There are several reasons to be uneasy over this derivation. |
For one. a static pressurcless universe is iupossible: there iust either be pressure (enough to slow up in the cncrev-momentiu tensor. or possibly as a cosmological coustant) or motion. | For one, a static pressureless universe is impossible; there must either be pressure (enough to show up in the energy-momentum tensor, or possibly as a cosmological constant) or motion. |
More worrviughv. we have both assumed a spatially Hat universe (in substituting the dvnamics equatious in Equation (12)). and one that is curved (Gu using the equation for eravitational redshift): thus the r variables in cach part of the derivation are not necessarily tle same (though at large distances they certainly approach cach other). | More worryingly, we have both assumed a spatially flat universe (in substituting the dynamics equations in Equation \ref{eq:twoshift}) ), and one that is curved (in using the equation for gravitational redshift); thus the $r$ variables in each part of the derivation are not necessarily the same (though at large distances they certainly approach each other). |
The result should be taken as no more than indicative th:t both motion and eravity coutribute to the cosmological redshift. | The result should be taken as no more than indicative that both motion and gravity contribute to the cosmological redshift. |
Moreover. other situations are not even this simple: so we still seek a more general wav of thiukiug of the phenomenon. | Moreover, other situations are not even this simple; so we still seek a more general way of thinking of the phenomenon. |
For this. we consider a different derivation. | For this, we consider a different derivation. |
Note that a light rav follows a null geodesic. so that all along it LEvervwhlere aloug this rax. then. If we integrate this equation aloug the path of propagation from the cussion to the observation of the licht ray we ect which is just the comoving distance between the euütter aud the observer. | Note that a light ray follows a null geodesic, so that all along it Everywhere along this ray, then, If we integrate this equation along the path of propagation from the emission to the observation of the light ray we get which is just the comoving distance between the emitter and the observer. |
Now we integrate the expression again. this time starting at a time just one wavelength later (0f) aud eudiug at a time just one wavelcneth later (0£,). | Now we integrate the expression again, this time starting at a time just one wavelength later $\delta t_e$ ) and ending at a time just one wavelength later $\delta t_o$ ). |
The difference between this integral and the previous one will be the change in comoving distance. | The difference between this integral and the previous one will be the change in comoving distance. |
But if the emitter and observer are both stationary with respect to the comoving coordinates. there is no change m comoving distance. | But if the emitter and observer are both stationary with respect to the comoving coordinates, there is no change in comoving distance. |
So Aud if anv change in e is small diving the period of one light wave. | So And if any change in $a$ is small during the period of one light wave, |
Iu our model we know the time at which the QS collapses to a DIE (the time of the steep decay). | In our model we know the time at which the QS collapses to a BH (the time of the steep decay). |
The calculations above assumed that this occurred at =Ts | The calculations above assumed that this occurred at $t_{\rm collapse}=\tau$. |
However. it could also occur at fep<Τ. which teollapseimplies that the magnetic field is weaker thau found above. | However, it could also occur at $t_{\rm collapse}<\tau$, which implies that the magnetic field is weaker than found above. |
Hence. the magnetic field. found. above is the maximum possible magnetic field. and therefore the spin-down liinosity. accretion rate and prompt eamunua rav energy are also miaxiniuna. | Hence, the magnetic field found above is the maximum possible magnetic field, and therefore the spin-down luminosity, accretion rate and prompt gamma ray energy are also maximum. |
The OS magnetic field needed to explain the flattening observed in GRD 070110 is B—Gs«1015 GG (seo Table 23). | The QS magnetic field needed to explain the flattening observed in GRB 070110 is $B=6.8\times10^{14}$ G (see Table \ref{calculatedtable}) ). |
The correspoucding spiu-down bhmuuinositv is fouud to be 1.G«10/5 eye/s. We can compare this to the observed eugime Iuniuositv assuninug an opening auele of 10 deerees for this outflow. | The corresponding spin-down luminosity is found to be $1.6\times10^{48}$ erg/s. We can compare this to the observed engine luminosity assuming an opening angle of 10 degrees for this outflow. |
If ve assume an efficiency. of 1054 in converting kinetic enerev to photons we see that we lave an order of magnitude more eucrey than needed. | If we assume an efficiency of $10\%$ in converting kinetic energy to photons we see that we have an order of magnitude more energy than needed. |
Comparing the observed prompt eauuna rav enerev to what we find from the jet launched by the QS. we again find that the jet energy is higher (by a factor 1) than the observed gana ray energy. | Comparing the observed prompt gamma ray energy to what we find from the jet launched by the QS, we again find that the jet energy is higher (by a factor 4) than the observed gamma ray energy. |
The QS magnetic field needed to explain the flattening observed in CRB 0606074 ix B21«107 C (see Table 2) | The QS magnetic field needed to explain the flattening observed in GRB 060607A is $1\times10^{15}$ G (see Table \ref{calculatedtable}) ). |
The corresponding spin-down luminosity is found. to be 3.6sLOS ore/s, Assunüng LOY efficiency im producing N-rav photons. we find (as for GRD 070110) that the estimated Iuninositv is lieher than the observed. | The corresponding spin-down luminosity is found to be $3.6\times10^{48}$ erg/s. Assuming $10\%$ efficiency in producing X-ray photons, we find (as for GRB 070110) that the estimated luminosity is higher than the observed. |
The Sauna rav enerev released during the prompt phase is also higher than the observed. gamuna ray enerev. | The gamma ray energy released during the prompt phase is also higher than the observed gamma ray energy. |
The higher luminosities can be because the estimate for the magnetic field is too high. 1ieauing that 7 is lareer and that the QS collapsed to a BIT before £27. | The higher luminosities can be because the estimate for the magnetic field is too high, meaning that $\tau$ is larger and that the QS collapsed to a BH before $t=\tau$. |
A lower magnetic field iuplies that the accretion rate is lower. | A lower magnetic field implies that the accretion rate is lower. |
Alternatively. we have overestimated the efficiencies. or the opening auele of the outflow is larger. | Alternatively, we have overestimated the efficiencies, or the opening angle of the outflow is larger. |
Iu GRBOGOGTA there are several N-rav flares observed uutil about 300 seconds (about 75 seconds when corrected for redshift). | In GRB06067A there are several X-ray flares observed until about 300 seconds (about 75 seconds when corrected for redshift). |
/ If we explain these flares bx accretion outo the QS as well. that means that the accretion process lasts for about 75 secouds. | If we explain these flares by accretion onto the QS as well, that means that the accretion process lasts for about $75$ seconds. |
The derived accretion rates imply the necessity of a debris disk with a duas of the order of ~10ΤΑΗν, which is reasonable since the QN. goes off inside a collapsar. where such a laree fallback disk is iu principle allowed. | The derived accretion rates imply the necessity of a debris disk with a mass of the order of $\sim 10^{-1}M_\odot$, which is reasonable since the QN goes off inside a collapsar, where such a large fall-back disk is in principle allowed. |
We lave presented a model to explain the fattening and occasional sharp drop-off ποσα im X-ray afterglows of some CRBs. | We have presented a model to explain the flattening and occasional sharp drop-off seen in X-ray afterglows of some GRBs. |
Our model borrows the framework of the 3 stage model presented in SODOT which makes use of an intermediate QS stage between the NS and the DIL | Our model borrows the framework of the 3 stage model presented in SOB07 which makes use of an intermediate QS stage between the NS and the BH. |
By appealing to a secoudary outflow. from the QS spin-down due to maenetic braking. our model seenis to explain the GRD itself νο, prompt cussion). the observed flat segment G.c. plateau). aud the subsequent sharp or gradual decay following the plateau. | By appealing to a secondary outflow, from the QS spin-down due to magnetic braking, our model seems to explain the GRB itself (i.e. prompt emission), the observed flat segment (i.e. plateau), and the subsequent sharp or gradual decay following the plateau. |
The sharp or gradual decay. depends ou whether the OS collapses to a DII or not durius spin-down. | The sharp or gradual decay depends on whether the QS collapses to a BH or not during spin-down. |
During spiu-down. a break will be seen after a characteristic time 7 given by Eq. | During spin-down, a break will be seen after a characteristic time $\tau$ given by Eq. |
3) followed by a power law with power of 5/3 to 3 (Panaitescu2007). | \ref{eq:charac_time} followed by a power law with power of $-5/3$ to $-3$ \citep{panaitescu07}. |
.. A very sharp drop-off will be seen if the QS collapses to a BIT during spin-down. | A very sharp drop-off will be seen if the QS collapses to a BH during spin-down. |
We uote that. if there was a wav for launching ultrarelativistic jets from accretion onto NSs. then it would be tempting to not include the QS phase in our model aud appeal ouly to NS to DII transition. | We note that, if there was a way for launching ultrarelativistic jets from accretion onto NSs, then it would be tempting to not include the QS phase in our model and appeal only to NS to BH transition. |
Towever. we are not aware of any such mechanism for αςήπιο au ultrarelativistie jet from accretion outo a NS. aud from an chergctics perspective it secius unlikely. | However, we are not aware of any such mechanism for launching an ultrarelativistic jet from accretion onto a NS, and from an energetics perspective it seems unlikely. |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.