source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
. The availability of deep observations of distaut galaxies bv the aud receutly bySpitzer (Wernerctal.2001). showed that the correlation between the IR and radio enüssion is not limited to the local universe but also exends in galaxies at higher redshitts (Ctarrett2002:Appοἵοιetal.2001). | The availability of deep observations of distant galaxies by the and recently by \citep{Werner04} showed that the correlation between the IR and radio emission is not limited to the local universe but also extends in galaxies at higher redshifts \citep{Garrett02,Appleton04}. |
.. It was also revealed that iu additkn to the FIR. the iul-IR cinission at —21jmu correates with the 20cm radio continui. though with mere scatter (Elbazetal.2002: 20051. | It was also revealed that in addition to the FIR, the mid-IR emission at $\sim$ $\mu$ m correlates with the 20cm radio continuum, though with more scatter \citep{Elbaz02,Gruppioni03,Appleton04,Wu05}. |
.. A nmuuber of very deepSpitzer mid-IR and FIR survevs ean now probe a population of ealaxics with low infrared huninosities for which aucillary data. iucludiug deep radio Huaging. are becoming available (c.g.Jauuuzi&Dey1999:Saundersetal.2007:Rosenberg 2006). | A number of very deep mid-IR and FIR surveys can now probe a population of galaxies with low infrared luminosities for which ancillary data, including deep radio imaging, are becoming available \citep[e.g.][]{Jannuzi99,Sanders07,Rosenberg06}. |
. This is particularly interesting since in the near future 20jnu is the longest wavelength which will likely be xobed bv the James Web Space Telescope. | This is particularly interesting since in the near future $\sim20\mu$ m is the longest wavelength which will likely be probed by the James Web Space Telescope. |
It is thus instructive to examine the mid-IR to radio correlation im nore detail in these low luuinosity nearby svstenis. for which little is known to date. | It is thus instructive to examine the mid-IR to radio correlation in more detail in these low luminosity nearby systems, for which little is known to date. |
As ubiquitous as the FIR/radio correlation appears o be. there are a few significant deviations fro it. | As ubiquitous as the FIR/radio correlation appears to be, there are a few significant deviations from it. |
Radio excess exists im some instances. such as galaxies josting an active ealactic nucleus (ACN). | Radio excess exists in some instances, such as galaxies hosting an active galactic nucleus (AGN). |
External nagnetic field. compression due to the interaction with rearby galaxies could also produce extra cluission iu he radio contiuuuau (Miller&Owen2001). | External magnetic field compression due to the interaction with nearby galaxies could also produce extra emission in the radio continuum \citep{Miller01}. |
.. Conversely. svuchrotrou deficieucvhas been fouud di some nasceut starburst ealaxies studied by Rousseletal.(2003.2006) which was attributed to the lack of time for massive | Conversely, synchrotron deficiencyhas been found in some nascent starburst galaxies studied by \citet{Roussel03,Roussel06} which was attributed to the lack of time for massive |
tesults are shown for both binary components A and D. out. since only the binary mass ratio changes between the wo. the results are very. similar. | Results are shown for both binary components A and B, but since only the binary mass ratio changes between the two, the results are very similar. |
Again. we see growth up o 10 km in the inner region of the disc. | Again, we see growth up to $70$ km in the inner region of the disc. |
Due to the fact hat the svstem is more strongly perturbed than 5.Cephei. he aceretion-[riendly region has shifted. inward. compared o Fig. 2.. | Due to the fact that the system is more strongly perturbed than $\gamceph$, the accretion-friendly region has shifted inward compared to Fig. \ref{figae}. |
The rough periodicity seen at early times in the »»ttom. panel of Fig. | The rough periodicity seen at early times in the bottom panel of Fig. |
5.r is due to different generations of Xdanetesimals. | \ref{figsizemassalphacen} is due to different generations of planetesimals. |
The Iargest objects are formed inside 1.43 AU with significant growth inside habitable zone <1 AU. | The largest objects are formed inside $1.4$ AU with significant growth inside habitable zone $< 1$ AU. |
We present the first study of planet. formation in. binaries aking into account the physical effects of collisions. | We present the first study of planet formation in binaries taking into account the physical effects of collisions. |
We show that these effects. tend. to. [favour growth owards large planetesimals in the perturbecl system. | We show that these effects tend to favour growth towards large planetesimals in the perturbed system. |
Two main mechanisms have been identified. | Two main mechanisms have been identified. |
First. frequent collisions tend. to. prevent planctesimeals [rom reaching heir equilibrium orbits. | First, frequent collisions tend to prevent planetesimals from reaching their equilibrium orbits. |
Lf collision rates are high enough. his means that the aceretion-hostile environment due o dillerential orbital phasing is never reached. | If collision rates are high enough, this means that the accretion-hostile environment due to differential orbital phasing is never reached. |
Second. ragments produced. by high-velocity collisions make up a arge reservoir of material that is very easily reacereted by he remaining planetesimals as they sweep through the disc on highly eccentric orbits. | Second, fragments produced by high-velocity collisions make up a large reservoir of material that is very easily reaccreted by the remaining planetesimals as they sweep through the disc on highly eccentric orbits. |
We have chosen parameters that in some respects should. be unfavourable for planetesimal growth. | We have chosen parameters that in some respects should be unfavourable for planetesimal growth. |
Gas drag is strong throughout the. disc in the models presented here. | Gas drag is strong throughout the disc in the models presented here. |
However. since gas drag plavs only a minor role. changing the gas density to a more realistic power law in radius does not change the results. | However, since gas drag plays only a minor role, changing the gas density to a more realistic power law in radius does not change the results. |
New planetesimals are formed. on circular orbits and are. therefore. immediately capable of destroving larger bodies that are on eccentric orbits. | New planetesimals are formed on circular orbits and are, therefore, immediately capable of destroying larger bodies that are on eccentric orbits. |
Forming new planetesimals on eccentric equilibrium orbits would favour more accreting collisions. | Forming new planetesimals on eccentric equilibrium orbits would favour more accreting collisions. |
On the other hand. we have assumed dust accretion to be efficient. | On the other hand, we have assumed dust accretion to be efficient. |
This ellicieney. depends on the size distribution that is produced in collisions and the vertical extent over which the fragments are distributed. | This efficiency depends on the size distribution that is produced in collisions and the vertical extent over which the fragments are distributed. |
Some of the fragments may. be lost due to racial drift. which is ignored in the current model. but it maw also help bringing more mass into the accretion friendly inner reglons. | Some of the fragments may be lost due to radial drift, which is ignored in the current model, but it may also help bringing more mass into the accretion friendly inner regions. |
We have worked in à 2D geometry for computational reasons. | We have worked in a 2D geometry for computational reasons. |
Although we have tried. to scale. the surface density in such a way to eet realistic. collision time scales. it is important to realise that του evolves in a different way with particle size than 7. | Although we have tried to scale the surface density in such a way to get realistic collision time scales, it is important to realise that $\tauctd$ evolves in a different way with particle size than $\tauc$. |
Lo is. therefore. not possible to have realistic collision time scales at all ines. | It is, therefore, not possible to have realistic collision time scales at all times. |
Furthermore. inclinations of the planetesimals nay x excited by collisions. resulting in 72. which decreases vw collision time scale. | Furthermore, inclinations of the planetesimals may be excited by collisions, resulting in $i=e/2$, which decreases the collision time scale. |
However. since most. collisions are estructive. it is likely that the remaining planetesimals will ssentially orbit in the plane of the disc. | However, since most collisions are destructive, it is likely that the remaining planetesimals will essentially orbit in the plane of the disc. |
Other important iee-dimensional effects were pointed out by 2? and 2.. who garowed that if the disc is inclined with respect to the binary plane. this may favour planetesimal accretion. | Other important three-dimensional effects were pointed out by \citet{xie09} and \citet{xie10}, who showed that if the disc is inclined with respect to the binary plane, this may favour planetesimal accretion. |
Clearly. three-imensional simulations are the way forward. and the results in this letter should be interpreted as a proof of principle. wt planetesimal accretion might be possible in tight binary systems under the right. circumstances. | Clearly, three-dimensional simulations are the way forward, and the results in this letter should be interpreted as a proof of principle, that planetesimal accretion might be possible in tight binary systems under the right circumstances. |
We thank Philippe Phébhault for useful comments that greatly improved the manuscript. Debra Fischer for getting us interested (again)) and the Isaac Newton Institute for their hospitality. | We thank Philippe Thébbault for useful comments that greatly improved the manuscript, Debra Fischer for getting us interested (again), and the Isaac Newton Institute for their hospitality. |
SJP and ZAIL are supported by STEC Postdoctoral Fellowships. | SJP and ZML are supported by STFC Postdoctoral Fellowships. |
excluding the 4-8 keV spectral range. | excluding the 4-8 keV spectral range. |
The line is broad and strongly redshifted as expected for an origin in the inner parts of an accretion disk. | The line is broad and strongly redshifted as expected for an origin in the inner parts of an accretion disk. |
We fitted each of the flux selected spectra with a model consisting of the best fit low state spectrum as discussed in section 3.. a power law. and the model. | We fitted each of the flux selected spectra with a model consisting of the best fit low state spectrum as discussed in section \ref{lowstate}, a power law, and the model. |
The model calculates the profile of a line originating from a Keplerian disk surrounding a Schwarzschild black hole. | The model calculates the profile of a line originating from a Keplerian disk surrounding a Schwarzschild black hole. |
The rest energy of the line was fixed at 6.4 keV. The geometric origin of the fluorescence line is detined by three parameters: the inner radius Ih. the outer radius /7,, and the power law index q which describes he radial variation of the emissivity. | The rest energy of the line was fixed at 6.4 keV. The geometric origin of the fluorescence line is defined by three parameters: the inner radius $R_{\rm i}$, the outer radius $R_{\rm o}$ and the power law index $q$ which describes the radial variation of the emissivity. |
We set /?/; to 6 των the radius of the innermost stable orbit for a Schwarzschild black hole. | We set $R_{\rm i}$ to 6 $r_{\rm g}$, the radius of the innermost stable orbit for a Schwarzschild black hole. |
The outer radius 75, was set to 400 αν | The outer radius $R_{\rm o}$ was set to 400 $r_{\rm g}$. |
The emissivity index q was left Tee tfo vary. | The emissivity index $q$ was left free to vary. |
Table 5 summarizes the results of the disk line modelling. | Table \ref{diskline} summarizes the results of the disk line modelling. |
For each spectrum we give both the equivalent widths of the narrow fluorescence lines from the fixed narrow line component and the best fit values for the broad disk line. | For each spectrum we give both the equivalent widths of the narrow fluorescence lines from the fixed narrow line component and the best fit values for the broad disk line. |
The values for the narrow ines are derived from the total fluxes of the lines at 6.4 keV. 7.06 keV. and 7.477 keV (see section 3). | The values for the narrow lines are derived from the total fluxes of the lines at 6.4 keV, 7.06 keV, and 7.477 keV (see section \ref{lowstate}) ). |
Note that the narrow line fluorescence from the torus can be a major contribution to the total ine flux. in particular during lower flux states. | Note that the narrow line fluorescence from the torus can be a major contribution to the total line flux, in particular during lower flux states. |
Hence the protile of the disk line can only be determined if the contribution from he torus ean be reliably measured. | Hence the profile of the disk line can only be determined if the contribution from the torus can be reliably measured. |
The best fit values of the disk inclination confirm the values of /~30° resulting from observations (Nandraeti.19973. | The best fit values of the disk inclination confirm the values of $i\sim 30^{\circ}$ resulting from observations \cite{Nandra97}. |
. The values of q in most cases are yoorly constrained. but generally the emissivity drops steeper than with q3. | The values of $q$ in most cases are poorly constrained, but generally the emissivity drops steeper than with $q=-3$. |
This indicates a concentration of the line emission in he innermost regions ofthe disk. as suggested by the broad profile ofthe iron line. | This indicates a concentration of the line emission in the innermost regions of the disk, as suggested by the broad profile of the iron line. |
Since the model tits described in section 4+ required no disk reflection (see Figs 34-4) we do not include disk reflection in the disk line model tits. | Since the model fits described in section \ref{specvar} required no disk reflection (see Figs \ref{mon_cont}+ \ref{dec_cont}) ) we do not include disk reflection in the disk line model fits. |
However. we repeated the model fits using a model with reflected fraction /?=1 instead of the power law model. | However, we repeated the model fits using a model with reflected fraction $R=1$ instead of the power law model. |
Including reflection to the fits did not change the line fluxes or disk line parameters significantuc | Including reflection to the fits did not change the line fluxes or disk line parameters significantly. |
The best fit photon indices. line fluxes. and line equivalent widths are plotted in Fig. 7.. | The best fit photon indices, line fluxes, and line equivalent widths are plotted in Fig. \ref{lineflux}. |
The best fit values of the fluorescence line flux from the long term monitoring observations show a clear correlation with the 2-10 keV flux. | The best fit values of the fluorescence line flux from the long term monitoring observations show a clear correlation with the 2-10 keV flux. |
However. the correlation. is not linear. in the continuum flux range range (2.0..5.5)10tereem.7s+ the line flux is nearly constant. | However, the correlation is not linear, in the continuum flux range range $(2.0..5.5) 10 ^{-11} {\rm
erg\; cm^{-2} s^{-1}}$ the line flux is nearly constant. |
The variation of the equivalent width is not very pronounced. the lowest and the three highest flux bins are consistent with AV=200eV. | The variation of the equivalent width is not very pronounced, the lowest and the three highest flux bins are consistent with $EW=200{\rm eV}$. |
On the shorter time scales of the December 1996 observation the line equivalent width decreases with continuum flux. although the line flux and continuum flux are still correlated. | On the shorter time scales of the December 1996 observation the line equivalent width decreases with continuum flux, although the line flux and continuum flux are still correlated. |
We show that the variable X-ray spectrum of NGC 4051 can be well described by a model that includes two principal components: | We show that the variable X-ray spectrum of NGC 4051 can be well described by a model that includes two principal components: |
computer power. | computer power. |
Note Chat we include it in the calculation of the radiation spectrum. as described in section ??.. | Note that we include it in the calculation of the radiation spectrum, as described in section \ref{spectrum}. |
The svnchrotron cooling5$ 4..,,(4./)init, is 5given bvMENO (e.g.Ivbicki&Lightman51979) where o is the Thomson cross section and Up(/) is the magnetic field energy. density calculated from equation (8)). | The synchrotron cooling $\dot{\gamma}_{\mathrm{syn}}(\gamma,t)$ is given by \citep[e.g.][]{rl79}
where $\sigma_{\mathrm{T}}$ is the Thomson cross section and $U_{\mathrm{B}}(t)$ is the magnetic field energy density calculated from equation \ref{eq8}) ). |
The inverse Compton cooling 3160(5) is given bv (e.g.Blumenthal&Gould1970) where h is the Planck's constant. ijj aud £g, are the frequency of the CAIB photons and that of scattered. photons. respectively. rear(ing) Is the distribution of the CMD described in equation (19)) below. D,=4s/vgi/Gnoc). g¢=hvg/(DP(5moc—hvg)). ο=241nq-(1+290—q)+050ΠΠΓΕ an | The inverse Compton cooling $\dot{\gamma}_{\mathrm{IC}}(\gamma)$ is given by \citep[e.g.][]{bg70} where $h$ is the Planck's constant, $\nu_{\rm{ini}}$ and $\nu_{\rm{fin}}$ are the frequency of the CMB photons and that of scattered photons, respectively, $n_{\rm{CMB}}(\nu_{\rm{ini}})$ is the distribution of the CMB described in equation \ref{eq19}) ) below. $\Gamma_{\rm{\epsilon}} = 4 \gamma h \nu_{\rm{ini}} / (m_{\rm{e}} c^2)$, $q = h \nu_{\rm{fin}} /(\Gamma_{\rm{\epsilon}} (\gamma m_{\rm{e}} c^2 - h \nu_{\rm{fin}}))$, $f(q, \Gamma_{\rm{\epsilon}}) = 2q \ln q + (1+2q)(1-q) + 0.5 (1-q) (\Gamma_{\epsilon}q)^2 / (1+\Gamma_{\epsilon}q)$, |
d 8 is the step function. | and $\theta$ is the step function. |
Finally. the adiabatic coolingS 544(5./)dad is 5given bv where we use equation (1)). | Finally, the adiabatic cooling $\dot{\gamma}_{\mathrm{ad}}(\gamma, t)$ is given by where we use equation \ref{eq1}) ). |
Note that the adiabatie cooling is independent of the expansion velocity Ppwx. | Note that the adiabatic cooling is independent of the expansion velocity $v_{\rm PWN}$. |
As was slated in section ??.. we consider that it is more reasonable to treat the acdiabatic loss rather than the escape loss. | As was stated in section \ref{intro}, we consider that it is more reasonable to treat the adiabatic loss rather than the escape loss. |
Zhangοἱal.(2008) considered an escape of the particles instead of an acdiabatie loss. | \citet{zet08} considered an escape of the particles instead of an adiabatic loss. |
They treated the escape of the particles based on the | They treated the escape of the particles based on the |
Thermal conduction mav play an important role iu the evolution of the intrachister luccliu (e.c. ?)). | Thermal conduction may play an important role in the evolution of the intracluster medium (e.g., \citet{voigt04}) ). |
Sufficieutlv stroug conductiou Cali offset radiative cooling losses iu massive clusters and reduce the cucrey requirenmens on active ealactie nuclei (AGN) eedback that is required to prevent overcooling iu less Massive clusters aud eroups (see ? and 7. for reviews). | Sufficiently strong conduction can offset radiative cooling losses in massive clusters and reduce the energy requirements on active galactic nuclei (AGN) feedback that is required to prevent overcooling in less massive clusters and groups (see \citet{mcnamara07} and \citet{norman} for reviews). |
It has also been sugeested that not only can thermal coxductiou serve as a mechanism for cool core heating but that it is vorv nuportaut for the stability of these svsteimis (2.. 7.. Ruszkowski Oh. in prep.). | It has also been suggested that not only can thermal conduction serve as a mechanism for cool core heating but that it is very important for the stability of these systems \citet{ruszkowski02}, \citet{guo08b}, Ruszkowski Oh, in prep.). |
It may also be responsible for setting a critical central eutropv threshold below which star formation is possible iu cluster cool cores (?).. | It may also be responsible for setting a critical central entropy threshold below which star formation is possible in cluster cool cores \citep{voit08}. |
The receutlv discovered bimodality iu the distribution of the cluster central cutropy (77) may be due to the combination of ACN feedback from the brightest cluster galaxies that stabilizes low cutropy clusters aud a colbination of mergers and thermal couduction that stabilize higher central cutropy clusters (ολε, | The recently discovered bimodality in the distribution of the cluster central entropy \citep{cavagnolo09, sanderson09}
may be due to the combination of AGN feedback from the brightest cluster galaxies that stabilizes low entropy clusters and a combination of mergers and thermal conduction that stabilize higher central entropy clusters \citep{guo08b, ruszkowski10, parrish10}. |
As such. hermal couduction may be important for uuderstanding of the feeding of the most massive black holes iu the Universe. | As such, thermal conduction may be important for understanding of the feeding of the most massive black holes in the Universe. |
In cluster outskirts thermal conduction may Hatten the temperature distributions (?).. which may lave consequences for the custer naass estimates. | In cluster outskirts thermal conduction may flatten the temperature distributions \citep{lemaster}, which may have consequences for the cluster mass estimates. |
This nav have possible impact Or precision cosmology as it relies on accurate lass neasurelents iu the most nassive Asa plasma transport process. thermal conduction is closely linked to gas viscosity. | This may have possible impact for precision cosmology as it relies on accurate mass measurements in the most massive As a plasma transport process, thermal conduction is closely linked to gas viscosity. |
Both types of transport processes nav expla various X-rav observations. | Both types of transport processes may explain various X-ray observations. |
For example. recent observatious of AIST (2). show that the temperature iu the shells centered ou the cluster center is remarkably isothermal. | For example, recent observations of M87 \citep{werner10} show that the temperature in the shells centered on the cluster center is remarkably isothermal. |
They suggest that such à ligh degree of isothermalty Is consistent with effective. heat conduction iu the tangential direction. | They suggest that such a high degree of isothermality is consistent with effective heat conduction in the tangential direction. |
velocity. | velocity. |
Iu other words. we shall have to make sure that theabsence of main sequence objects at a certain magnitude is real and not the product of poor statistics. | In other words, we shall have to make sure that the of main sequence objects at a certain magnitude is real and not the product of poor statistics. |
We shall cousider the CAL diagrams obtained for star siuuples with increasing values of |V| (Figs. 2.. | We shall consider the CM diagrams obtained for star samples with increasing values of $\vert V \vert$ (Figs. \ref{fig2}, |
3 aud 6)). exanuning them in relation one to the other. | \ref{fig5} and \ref{fig7}) ), examining them in relation one to the other. |
We start from the first velocity bin (0 < V| « 10 dan S1 ). | We start from the first velocity bin (0 $<$ $\vert V \vert$ $<$ 10 km $\rm
s^{-1}$ ). |
In this sample the main sequence is well populated from well below the turu-off of the oldest stars preseut (about uaenitude 1) to a huninosity corresponding to a iuninuun age of few 10' vr. | In this sample the main sequence is well populated from well below the turn-off of the oldest stars present (about magnitude 4) to a luminosity corresponding to a minimum age of few $10^7$ yr. |
Such age iuav be overestimated. since our sanuple cau iiss some of the voungest stars. | Such age may be overestimated, since our sample can miss some of the youngest stars. |
We proceed examining the CM diagrams for decreasing values of the V. component. | We proceed examining the CM diagrams for decreasing values of the $V$ component. |
The Επο region for the oldest stars aud the lower main sequence remain larecly simular down to W about -60 -s0. and the higher main sequence retains well populated wp to about few 105 yr as long as V. 2 -30 ian s | The turn-off region for the oldest stars and the lower main sequence remain largely similar down to $V$ about -60 $-$ -80, and the higher main sequence remains well populated up to about few $10^7$ yr as long as $V$ $>$ -30 km $\rm s^{-1}$. |
We interprete the CAL diagrams in the velocity bius V > -30 lan st as describing a population with the same UI (c101 vr) and maxima age (= 1010 vr). | We interprete the CM diagrams in the velocity bins $V$ $>$ -30 km $\rm
s^{-1}$ as describing a population with the same minimum $\leq 10^7$ yr) and maximum age $\geq$ $10^{10}$ yr). |
Most of the objecs are found a the right of the Hsochrouc. sugecstingOO the absence of a substantial nunber of inetal poor stars. | Most of the objects are found at the right of the isochrone, suggesting the absence of a substantial number of metal poor stars. |
Exi of the total sample with a,/7 < 1 docs not change this interpretation. | Exam of the total sample with $\sigma_\pi$ $\pi$ $<$ 1 does not change this interpretation. |
Iu fact we repeat the CAL diagram relative o 30 «V< -20 kins considering objects with o,/7 «1 o show that it actually carries tle sale mformation fouud in the 04/3 < 0.15 case. both for he high and the low regions of the main sequence. | In fact we repeat the CM diagram relative to -30 $< V <$ -20 km $\rm s^{-1}$ considering objects with $\sigma_\pi$ $\pi$ $<$ 1 to show that it actually carries the same information found in the $\sigma_\pi$ $\pi$ $\leq$ 0.15 case, both for the high and the low regions of the main sequence. |
The CAL diagram relative to the sample -l0<V.< -30 xls is the most ambiguous. | The CM diagram relative to the sample -40 $< V <$ -30 km $\rm s^{-1}$ is the most ambiguous. |
When the stars with σι π < 0.15 only are considered. the CM diagram (uot show) exhibits a stroue increase iu the age for a large part of tlie »opulation. up to about 1 2109. | When the stars with $\sigma_ \pi$ $\pi$ $\leq$ 0.15 only are considered, the CM diagram (not shown) exhibits a strong increase in the age for a large part of the population, up to about 1 $-$ 2 $10^9$ yr. |
Ou the coutrary. there ire apparently many vouug stars ↴⋅↴⋅in the sample↴ with- ez/z TN« lL (Fig. 2)). | On the contrary, there are apparently many young stars in the sample with $\sigma_ \pi$ $\pi$ $<$ 1 (Fig. \ref{fig2}) ), |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.