diff --git "a/batch_s000017.csv" "b/batch_s000017.csv"
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+source,target
+ Iu addition. we make svuthesized beam sizes the same as possible at both wavelengths. using weighting aud tapering schemes. in order to münunuize the beam size effect on the flux comparison.," In addition, we make synthesized beam sizes the same as possible at both wavelengths, using weighting and tapering schemes, in order to minimize the beam size effect on the flux comparison."
+ After proper weighting and taperie schemes. we could match the beam sizes to within," After proper weighting and tapering schemes, we could match the beam sizes to within."
+ The details of applied weighting aud tapering schemes are listed in Table 2. with final svuthesized bemus., The details of applied weighting and tapering schemes are listed in Table \ref{tab_beam} with final synthesized beams.
+ Driees robust paraiueter is used (?7).. which is a knob to provide intermediate weighting between natural aud uniform weighting.," Briggs' robust parameter is used \citep{briggs1995}, which is a knob to provide intermediate weighting between natural and uniform weighting."
+ The parameter of 2 gives a weieliting close to natural weieliting aud 2 close to nuiform weighting., The parameter of $2$ gives a weighting close to natural weighting and $-2$ close to uniform weighting.
+" Total iux CF.) of the thermal dust coutiuumau emission represents the total mass (Mr) of the source. if the source is optically thin at the observational frequeucies. where wy. DQ(DI4) Mp. aud D are opacity (nass absorption cocficient) of the dust serais. blackbody radiation intensity of a dust temperature Dy. total mass. and distance to the source. respectively,"," Total flux $F_\nu$ ) of the thermal dust continuum emission represents the total mass $M_T$ ) of the source, if the source is optically thin at the observational frequencies, where $\kappa_\nu$, $B_\nu(T_{d})$, $M_T$ , and $D$ are opacity (mass absorption coefficient) of the dust grains, blackbody radiation intensity of a dust temperature $T_{d}$, total mass, and distance to the source, respectively."
+" The opacity of dust erains (&,) depeuds on dust properties such as sizes. colmpoucuts. aud shapes."," The opacity of dust grains $\kappa_{\nu}$ ) depends on dust properties such as sizes, components, and shapes."
+" Hf the dependence is simple. for example a power law (E,cr7). the dust eran propertics can be studied by observations at two frequencies."," If the dependence is simple, for example a power law $\kappa_\nu \varpropto \nu^\beta$ ), the dust grain properties can be studied by observations at two frequencies."
+ In addition. in the case that the Ravleigh-Jeans approxination of blackbody radiation is applicable (ivmKT). the relationship between spectral indexes of the observed flux deusities (0) aud spectral indexes of the dust exin opacity (2) is simply. Note that this relation is valid only iu the optically thin asstuuption aud the Ravleigh-Jeaus approximation.," In addition, in the case that the Rayleigh-Jeans approximation of blackbody radiation is applicable $h\nu \ll kT$ ), the relationship between spectral indexes of the observed flux densities $\alpha$ ) and spectral indexes of the dust grain opacity $\beta$ ) is simply, Note that this relation is valid only in the optically thin assumption and the Rayleigh-Jeans approximation."
+ ? showed that ./ mainly depends on the size distribution of dust erains rather than their coupoucuts and shapes: sunall (j| (~ 1) is likely indicating dust erain size distribution up to 3A., \citet{draine2006} showed that $\beta$ mainly depends on the size distribution of dust grains rather than their components and shapes; small $\beta$ $\sim 1$ ) is likely indicating dust grain size distribution up to $3 \lambda$.
+ Since our observations are up to 3 numi 9~l would sugeest a grain size distribution up to about 1 cm.," Since our observations are up to 3 mm, $\beta \sim 1$ would suggest a grain size distribution up to about 1 cm."
+ Figure d preseuts maps of Lilis IRS 2. Litls IRS 3. and L1157.," Figure \ref{fig_betamap} presents maps of L1448 IRS 2, L1448 IRS 3, and L1157."
+" Dust coutimaiun maps at and hhave been separately constructed using different weiehtiues and taperimes as described in refsoc,bsand Table2inordertohaveassimilarsuyithesitedbeamsaspos v"," Dust continuum maps at and have been separately constructed using different weightings and taperings as described in \\ref{sec_obs} and Table \ref{tab_beam}
+ in order to have as similar synthesized beams as possible at the two wavelengths."
+alues of cach source have been calculated using the two contimmun images., Afterwards $\beta$ values of each source have been calculated using the two continuum images.
+ Only regions above three signal-to-noise ratio (SNR) levels on the both maps have Όσοι used to derive ) assunmniue where g4 and my are frequencies correspouding to and ddata. as listed in Table 2..," Only regions above three signal-to-noise ratio (SNR) levels on the both maps have been used to derive $\beta$ assuming where $\nu_1$ and $\nu_0$ are frequencies corresponding to and data, as listed in Table \ref{tab_beam}."
+ Note that the Ravleigh-Jeans approximation aud the optically thin assuuption are used., Note that the Rayleigh-Jeans approximation and the optically thin assumption are used.
+ In the case of an average dust temperature of about 30 I&. the upper lait of frequencies to which the Ravleigh-Jeans approximation cau be applied is about 625 GIIz.," In the case of an average dust temperature of about 30 K, the upper limit of frequencies to which the Rayleigh-Jeans approximation can be applied is about 625 GHz."
+ Since the higher frequencyof our data is about 230 GIIz. the assumption is valid for this study.," Since the higher frequencyof our data is about 230 GHz, the assumption is valid for this study."
+ However. caution should be taken in comparison at sibmullimeter wavelengths for cold objects such as the Class 0 YSO envelopes.," However, caution should be taken in $\beta$ comparison at submillimeter wavelengths for cold objects such as the Class 0 YSO envelopes."
+The rate of supernova explosions is astrophysically important because it determines the rate at which heavy elements are dispersed into the interstellar medium. thereby constraming galactic chemical evolution.,"The rate of supernova explosions is astrophysically important because it determines the rate at which heavy elements are dispersed into the interstellar medium, thereby constraining galactic chemical evolution."
+ Since the progenitors of core-collapse supernovae (SNec) are believed to be short-lived massive stars. the SNee rate is expected to reflect the star-formation rate. increasing with redshift like ~(1+2°° for zo«0.5 (Hopkins&Beacom2006).," Since the progenitors of core-collapse supernovae (SNcc) are believed to be short-lived massive stars, the SNcc rate is expected to reflect the star-formation rate, increasing with redshift like $\sim(1+z)^{3.6}$ for $z<0.5$ \citep{hopkins}."
+. Thermonuclear type Ia supernovae (SNIa) have both long- and short-lived progenitors so the SNIa rate has a delayed component making the SNla rate rise more slowly with redshift. ~(1+2) (Pritchetetal.2008)..," Thermonuclear type Ia supernovae (SNIa) have both long- and short-lived progenitors so the SNIa rate has a delayed component making the SNIa rate rise more slowly with redshift, $\sim(1+z)^{2}$ \citep{pritchet}."
+ The SNla rate is now known to a precision of about 20%., The SNIa rate is now known to a precision of about $20\%$.
+ Measurements have profited from the high luminosity of SNIa which make them relatively easy to detect and identify., Measurements have profited from the high luminosity of SNIa which make them relatively easy to detect and identify.
+ Furthermore. their utility as cosmological distance indicators has motivated intense searches.," Furthermore, their utility as cosmological distance indicators has motivated intense searches."
+ An example is the Supernova Legacy Survey (SNLS) at the Canada-France- Telescope (CFHT) performed between 2003 and 2008., An example is the Supernova Legacy Survey (SNLS) at the Canada-France-Hawaii Telescope (CFHT) performed between 2003 and 2008.
+" Using early SNLS data. Neilletal.(2006) derived a SNIa rate at a redshift z~0.5 of where /is=Ho/70kmsec""!Mpc'."," Using early SNLS data, \citet{neill2006}
+ derived a SNIa rate at a redshift $z\sim 0.5$ of where $h_{70}=H_0/70\km\,\second^{-1}\Mpc^{-1}$."
+ The rate for SNee is more difficult to measure because observed SNec have a magnitude distribution that peaks roughly 1.5mag fainter than SNIa and covers a range of more than Smag (Richardsonetal.," The rate for SNcc is more difficult to measure because observed SNcc have a magnitude distribution that peaks roughly $1.5\,{\rm mag}$ fainter than SNIa and covers a range of more than $5\,{\rm mag}$ \citep{richardson}."
+2002)... The local rate was measured by Cappellaroetal.(1999) using. 137 supernovae discovered by eye and photographically.," The local rate was measured by \citet{Capp1999}
+ using 137 supernovae discovered by eye and photographically."
+ Most had spectroscopic identification. about half being SNIa and half SNee (SNIb/e and SNID.," Most had spectroscopic identification, about half being SNIa and half SNcc (SNIb/c and SNII)."
+ After etficiency corrections. the SNec rate was found to be a factor 2.441.3 greater than the SNIa rate.," After efficiency corrections, the SNcc rate was found to be a factor $2.4\pm1.3$ greater than the SNIa rate."
+ The SNee rate at z~0.3 was measured by Cappellaroal.(2005) and more recently by Botticellaetal.(2005)., The SNcc rate at $z\sim0.3$ was measured by \citet{Capp2005} and more recently by \citet{Bott2008}.
+ The latter used images taken over a six year period with typically four months between images., The latter used images taken over a six year period with typically four months between images.
+ They found 18 SNee candidates and 13 SNla candidates (of which a total of 25 are spectroscopically confirmed) to find a SNee rate at z~0.26 a factor 4+2 greater than the SNla rate., They found 18 SNcc candidates and 13 SNIa candidates (of which a total of 25 are spectroscopically confirmed) to find a SNcc rate at $z\sim0.26$ a factor $4\pm2$ greater than the SNIa rate.
+ Finally. Dahleretal.(2004) used the Advanced Camera for Surveys οἱ the Hubble Space Telescope to obtain images for five epochs separated by ~45days.," Finally, \citet{dahlen} used the Advanced Camera for Surveys on the Hubble Space Telescope to obtain images for five epochs separated by $\sim45\,{\rm days}$."
+" For redshifts <1. they found 17 SNla candidates (with some spectroscopic identification) anc 16 SNee candidates (no spectroscopic identification) which allowed them to derive R,./R;,=3.6x2.0 at z-0. and RR,22.5E1.0 at z~0.8."," For redshifts $<1$, they found 17 SNIa candidates (with some spectroscopic identification) and 16 SNcc candidates (no spectroscopic identification) which allowed them to derive $R_{cc}/R_{Ia}=3.6\pm2.0$ at $z\sim0.4$ and $R_{cc}/R_{Ia}=2.5\pm1.0$ at $z\sim0.8$."
+ All existing measurements of the SNee rate suffer from the fact that the discovery procedure involved the comparison of images separated in time by intervals comparable to or greater than the characteristic ~|month time scales of supernovae.," All existing measurements of the SNcc rate suffer from the fact that the discovery procedure involved the comparison of images separated in time by intervals comparable to or greater than the characteristic $\sim 1\,{\rm month}$ time scales of supernovae."
+ Consequently. well-sampled light curves for most candidates are not available. complicatir£& the type identification and efficiency caleulations.," Consequently, well-sampled light curves for most candidates are not available, complicating the type identification and efficiency calculations."
+" The SLS ""rolling search"" avoids this problem because of its high cadence monitoring of four |deg- fields in the e. 7. /' and z' bands over a total of5 years."," The SNLS “rolling search” avoids this problem because of its high cadence monitoring of four $1\,deg^2$ fields in the $g^\prime$, $r^\prime$, $i^\prime$ and $z^\prime$ bands over a total of 5 years."
+ During each 6 month observing seaso1 for each field. typically four observations perlunation were obtained in the 7’ and /' bands. three in the z' band and two inthe ¢’ band.," During each 6 month observing season for each field, typically four observations perlunation were obtained in the $r^\prime$ and $i^\prime$ bands, three in the $z^\prime$ band and two in the $g^\prime$ band."
+ This strategy yields well-sampled light curves (e.g. Figs. 1.. 5," This strategy yields well-sampled light curves (e.g. Figs. \ref{lcIa}, ,"
+ and 3)) with high efficiency for all events occurring during the observing season, \ref{lcplateau} and \ref{lcCC}) ) with high efficiency for all events occurring during the observing season
+The source count models can be accessed through the world wicle web at Alip:££www.irdsas.ac.jpicppícounts CPP is supported. by a Japan Society [or the Promotion of Science (JSPS) fellowship.,The source count models can be accessed through the world wide web at $http://www.ir.isas.ac.jp/\sim cpp/counts/$ CPP is supported by a Japan Society for the Promotion of Science (JSPS) fellowship.
+ CPP thanks the referee Steve Eales and Lideo Alatsuhara lor significant. comments and suggestions., CPP thanks the referee Steve Eales and Hideo Matsuhara for significant comments and suggestions.
+ CPP would like to thank Professor Haruyuki Okuda for providing him with the chance to spend a very fruitful ancl eventful time at the Institute of Space and Astronautical Science. Japan.," CPP would like to thank Professor Haruyuki Okuda for providing him with the chance to spend a very fruitful and eventful time at the Institute of Space and Astronautical Science, Japan."
+metallicity to 47 Tuc with [M/H]=—0.6.,metallicity to 47 Tuc with $\mh=-0.6$.
+" Given the errors, the isochrone fit for the Aquarius stream is consistent with the Sgr dwarf."," Given the errors, the isochrone fit for the Aquarius stream is consistent with the Sgr dwarf."
+ We thus investigated a possible link between the Aquarius stream and the Sagittarius dwarf debris., We thus investigated a possible link between the Aquarius stream and the Sagittarius dwarf debris.
+" The details of this investigation are given in Appendix A. The overall result is that the Aquarius stream's kinematics match those of the Sagittarius dwarf debris, calculated using a variety of potential models (oblate, spheroid, prolate, triaxial) from Law (2005, 2009)."," The details of this investigation are given in Appendix A. The overall result is that the Aquarius stream's kinematics match those of the Sagittarius dwarf debris, calculated using a variety of potential models (oblate, spheroid, prolate, triaxial) from Law (2005, 2009)."
+" The oblate model shows a potential match for a small section of nearby debris when considering the line-of-sight velocity in the Galactic rest-frame, V, alone."," The oblate model shows a potential match for a small section of nearby debris when considering the line-of-sight velocity in the Galactic rest-frame, $\vgal$, alone."
+" However, the full kinematics of Vy,Vr,Vz displays that the kinematics of the Aquarius stream and this nearby section are actually quitedifferent?°."," However, the full kinematics of $\VPHI,\ \VR,\ \VZ$ displays that the kinematics of the Aquarius stream and this nearby section are actually quite."
+. The possible connection is further ruled out by the fact that the oblate halo potential model does not compare well with other observational data for Sgr dwarf debris., The possible connection is further ruled out by the fact that the oblate halo potential model does not compare well with other observational data for Sgr dwarf debris.
+" Since the Aquarius stream lies in the southern part of the RAVE data, it could not be discovered in the main, northern SDSS survey."," Since the Aquarius stream lies in the southern part of the RAVE data, it could not be discovered in the main, northern SDSS survey."
+" Thus the stream is far removed from the SDSS-discovered substructures, including the Canis Major overdensity at (1,b)=(240°,—8?) (Martinezetal.2005) and the Virgo overdensity at (1,b)=(300°,+60°) (Juriéetal.2008)."," Thus the stream is far removed from the SDSS-discovered substructures, including the Canis Major overdensity at $(l,\ b)=(240^\circ,\ -8^\circ)$ \citep{Martinez2005} and the Virgo overdensity at $(l,\ b)=(300^\circ,\ +60^\circ)$ \citep{Juric2008}."
+". Further, the stream is located between the southern SEGUE SDSS stripes so it unsurprising that this has not been detected in this survey."," Further, the stream is located between the southern SEGUE SDSS stripes so it unsurprising that this has not been detected in this survey."
+ The stream’s Galactic latitude of b=—60° rules out a relation to the more planar Monoceros stream (b< 40°) (Penarrubiaetal.2005)., The stream's Galactic latitude of $b=-60^\circ$ rules out a relation to the more planar Monoceros stream $b<40^\circ$ ) \citep{Penarrubia2005}.
+". Its velocities and latitude are also inconsistent with the thick disk asymmetries detected by Parkeretal.(2003, 2004).."," Its velocities and latitude are also inconsistent with the thick disk asymmetries detected by \citet{Parker2003, Parker2004}. ."
+" The Hercules-Aquila cloud, again detected using SDSS photometry, is located at |=40° and extends above and below the plane by 50° Belokurovetal. (2007).."," The Hercules-Aquila cloud, again detected using SDSS photometry, is located at $l=40^\circ$ and extends above and below the plane by $50^\circ$ \citet{Belokurov2007}. ."
+ The velocities of the b>0° segment are =+180kmst and the structure ranges over heliocentricVga) distances of d=10—20kpc.," The velocities of the $b>0^\circ$ segment are $\vgal = +180\,\kms$ and the structure ranges over heliocentric distances of $d=10 - 20\,\kpc$."
+ The Hercules-Aquila cloud is near the Aquarius stream on the sky., The Hercules-Aquila cloud is near the Aquarius stream on the sky.
+" However, despite the lack of velocity data below the plane, it can be clearly seen that the two entities are separate: the centering in (1, b) for the two are shifted from each other and their distance ranges are clearly incompatible."," However, despite the lack of velocity data below the plane, it can be clearly seen that the two entities are separate: the centering in (l, b) for the two are shifted from each other and their distance ranges are clearly incompatible."
+" Additionally, in Section 6.1 we trace the orbit of a simple model for the Aquarius stream and the resulting region of phase-space that the debris inhabits does not overlap with the Hercules-Aquila cloud in (1,b, και)."," Additionally, in Section \ref{subsec:sat} we trace the orbit of a simple model for the Aquarius stream and the resulting region of phase-space that the debris inhabits does not overlap with the Hercules-Aquila cloud in $l,\ b,\ \vgal$ )."
+" We have calculated orbits for candidate stars in the potential of Helmi et al (2006), which has contributions from a disk, bulge, and dark halo."," We have calculated orbits for candidate stars in the potential of Helmi et al (2006), which has contributions from a disk, bulge, and dark halo."
+" Table 4 gives averages for various quantities derived from these orbits as well as the median quantities for the overall kinematics, using both sets of distances."," Table \ref{tab4} gives averages for various quantities derived from these orbits as well as the median quantities for the overall kinematics, using both sets of distances."
+" Note that we chose the median as it gives more consistent results, and for this reason we also excluded the two most distant stars with d>9 kpc as their kinematics differed greatly from the others."," Note that we chose the median as it gives more consistent results, and for this reason we also excluded the two most distant stars with $d> 9$ kpc as their kinematics differed greatly from the others."
+" Also, the values for the pericentre and apocentre only include non-radial orbits."," Also, the values for the pericentre and apocentre only include non-radial orbits."
+" Figure 6 shows the L,-Lyerp and L,-Energy (Lindblad) planes for orbits based on both distance estimates, where to aid comparison to other studies we use here energies as calculated in Dinescuetal. (1999)."," Figure \ref{f6} shows the $L_z$ $L_\mathrm{perp}$ and $L_z$ -Energy (Lindblad) planes for orbits based on both distance estimates, where to aid comparison to other studies we use here energies as calculated in \citet{Dinescu1999}."
+". Note that the scatter of the isochrone distance results is large so the majority of these points lie off the plot, as do some of the RPM distance results."," Note that the scatter of the isochrone distance results is large so the majority of these points lie off the plot, as do some of the RPM distance results."
+" We plot for reference stars in the Geneva Copenhagen survey (Nordstrometal.2004), which is comprised mainly of thin and some thick disk stars."," We plot for reference stars in the Geneva Copenhagen survey \citep{Nordstrom2004}, which is comprised mainly of thin and some thick disk stars."
+ The circular orbit loci for this potential are also shown in the Lindblad diagram., The circular orbit loci for this potential are also shown in the Lindblad diagram.
+ To show the typical error covariance we also ran a Monte Carlo (MC) simulation for each star (with the RPM , To show the typical error covariance we also ran a Monte Carlo (MC) simulation for each star (with the RPM distances).
+"Fromthe errors in distances,proper motion and radial distances).velocity we generated a sample of 1000 points representative of each distribution, which"," Fromthe errors in distances,proper motion and radial velocity we generated a sample of 1000 points representative of each distribution, which"
+So far we have ignored the effects of extinction iu our conrparisou of observations and models,So far we have ignored the effects of extinction in our comparison of observations and models.
+ Frou detailed studies of obscured Calactic giant regions there is evidence that the extinction to the stars is highly variable with values between Ay=0 and Ay=L5inasg (Bhuin. Danuneli. Conti 1999: Bhun. Couti. Diuauineli 2000) for the very voung regions.," From detailed studies of obscured Galactic giant regions there is evidence that the extinction to the stars is highly variable with values between $A_K=0$ and $A_K=4.5\,$ mag (Blum, Damineli, Conti 1999; Blum, Conti, Damineli 2000) for the very young regions."
+ If high extinction is prescut in the regious of LIRGs. then the predicted fractious of regions. coiucidences aud older star clusters will be affected.," If high extinction is present in the regions of LIRGs, then the predicted fractions of regions, coincidences and older star clusters will be affected."
+ Althouehl we have no way derive the extinctions from the current Pan observations. we cau obtain some estimates from optical spectroscopy.," Although we have no way derive the extinctions from the current $\alpha$ observations, we can obtain some estimates from optical spectroscopy."
+ AATIOO and Lippari et al. (, AAH00 and pari et al. (
+2000) have measured extinctions of up to Ay& L2amag using the Bahuer decrement for a few regions in Arp 299 and NCC 3256 respectively.,"2000) have measured extinctions of up to $A_V \simeq 4.2\,$ mag using the Balmer decrement for a few regions in Arp 299 and NGC 3256 respectively."
+" This is equivalent to extinctions at the observed wavelength of Pad of up to Apa,zm OSimae."," This is equivalent to extinctions at the observed wavelength of $\alpha$ of up to $A_{{\rm Pa}\alpha} \simeq 0.5\,$ mag."
+ The ages of these regions from the observed equivalent widtls of Πα are of between 3 and GNINDIVY using Leitherer et al. (, The ages of these regions from the observed equivalent widths of $\alpha$ are of between 3 and Myr using Leitherer et al. (
+1999) models.,1999) models.
+ Although the measured extinctions are not high enough to compromise the detection of regions at nearinfrared wavelengths. we cannot rule out the possibility that verv voung regious («23 Myr) suffer from elevated extinctions. even at nem-iufrared waveleneths.," Although the measured extinctions are not high enough to compromise the detection of regions at near-infrared wavelengths, we cannot rule out the possibility that very young regions $<3\,$ Myr) suffer from elevated extinctions, even at near-infrared wavelengths."
+ 1f this were the cases we would be missing the vounecst regions. and hence the observed fraction of regions and coimcidences will be lower lits.," If this were the case, we would be missing the youngest regions, and hence the observed fraction of regions and coincidences will be lower limits."
+ This would iu turi translate in even vounger age distribution of the detected population of star clusters., This would in turn translate in even younger age distribution of the detected population of star clusters.
+ The observed luminosity and mass functious of old elobular clusters around galaxies and voung star clusters observed in interacting galaxies appear to be distinctivelv different. with the former having a log-normal shape. aud the latter a power law form with no evidence for a turnover (sce for instance Ehueercen Efremov 1997: Whitinore et al.," The observed luminosity and mass functions of old globular clusters around galaxies and young star clusters observed in interacting galaxies appear to be distinctively different, with the former having a log-normal shape, and the latter a power law form with no evidence for a turnover (see for instance Elmegreen Efremov 1997; Whitmore et al."
+ 1999: Zepf et al., 1999; Zepf et al.
+ 1999)., 1999).
+ This poses an interesting problem if the clusters observed iu salaxies are vounger versious of todav's globular clusters., This poses an interesting problem if the clusters observed in galaxies are younger versions of today's globular clusters.
+ Oue of the proposed solutions is that the mass distribution of old elobular clusters was initially a power law that was later modified by selective destruction of low mass star clusters to become a log-norimat like wass fiction (see Fall Zhang 2001 for a detailed discussion. and also Whitmore 2000).," One of the proposed solutions is that the mass distribution of old globular clusters was initially a power law that was later modified by selective destruction of low mass star clusters to become a log-normal like mass function (see Fall Zhang 2001 for a detailed discussion, and also Whitmore 2000)."
+ Although including a detailed treatineut of the problem of cluster destruction is bevoud the scope of this paper. we can attempt to see its effects on our calculations.," Although including a detailed treatment of the problem of cluster destruction is beyond the scope of this paper, we can attempt to see its effects on our calculations."
+ If clusters in Arp 299 and NGC 3256 have been forming at a coustaut rate for approximately LOO Nr. then ~50% of all clusters with masses above 5«104M. should be destroyed during that period to account for the observed fraction of star clusters.," If clusters in Arp 299 and NGC 3256 have been forming at a constant rate for approximately $100\,$ Myr, then $\simeq 50\%$ of all clusters with masses above $5\times 10^4\,{\rm M}_\odot$ should be destroyed during that period to account for the observed fraction of star clusters."
+ Note that Zepf et al. (, Note that Zepf et al. (
+1999) proposed selective destruction of low mass clusters as oue possibility to account for the observed optical colors and Iuinosities of clusters in NGC 3256 (the other possibility was a very voung aee for the clusters).,1999) proposed selective destruction of low mass clusters as one possibility to account for the observed optical colors and luminosities of clusters in NGC 3256 (the other possibility was a very young age for the clusters).
+ Finally we cousider the effects of the formation of Pan shells ou the measured uuuber of coincidences., Finally we consider the effects of the formation of $\alpha$ shells on the measured number of coincidences.
+ Iu the Auteunae galaxy there is evidence that many of the slightly older. nore massive clusters (1.0. 5-LOADNIwr) have blown large Te shells around themselves. hence there is no lounger a good correspondence between the distribution of Ila and the f-hand ceuter of the cluster (Whitinore ct al.," In the Antennae galaxy there is evidence that many of the slightly older, more massive clusters (i.e., Myr) have blown large $\alpha$ shells around themselves, hence there is no longer a good correspondence between the distribution of $\alpha$ and the $I$ -band center of the cluster (Whitmore et al."
+ 1999)., 1999).
+ Such Pana shells ave diffuse and will not be identifies as reeious. causing the προς of coincidences to be underestimated iu the 57 Myv age range.," Such $\alpha$ shells are diffuse and will not be identified as regions, causing the number of coincidences to be underestimated in the $5-7\,$ Myr age range."
+ For ages 2TMyr the region endussion will not be detecte as its ΕΕ will be below our detection threshok (Section 5.2).," For ages $>7\,$ Myr the region emission will not be detected as its luminosity will be below our detection threshold (Section 5.2)."
+ The result of missing some coiucideuces because of the presence of Pan shells would be an age distribution of the clusters shelthy vouuger than iuferre in Section 5.3., The result of missing some coincidences because of the presence of $\alpha$ shells would be an age distribution of the clusters slightly younger than inferred in Section 5.3.
+ Note that this effect is ouly relevant to the most imassive clusters in Arp 299 iud NGC 3256 where we expect to see coincidences with the preseut detection threshold., Note that this effect is only relevant to the most massive clusters in Arp 299 and NGC 3256 where we expect to see coincidences with the present detection threshold.
+ Iu this paper we have presented Z$T/NICMOS broad-baud and uarrow-baud Pao tuagine of a sample of 5 LIRGs., In this paper we have presented /NICMOS broad-band and narrow-band $\alpha$ imaging of a sample of 8 LIRGs.
+" The sample galaxies exhibit a range of infrared huninosities (logLy,=10.91.I1L82L.). as well as a varicty of dynamical stages: isolated galaxies interacting ealaxies aud mergers."," The sample galaxies exhibit a range of infrared luminosities $\log L_{\rm IR} = 10.94-11.82\,{\rm L}_\odot$ ), as well as a variety of dynamical stages: isolated galaxies interacting galaxies and mergers."
+ The Pan images have allowed us to identify the location of regions. whereas the Z7-baud contimmiun nuages have revealed the star clusters.," The $\alpha$ images have allowed us to identify the location of regions, whereas the $H$ -band continuum images have revealed the star clusters."
+ In all ealaxies iu our sample except NGC 6210 aud Zw 019.057 we have detected a large number of star clusters and regions., In all galaxies in our sample except NGC 6240 and Zw 049.057 we have detected a large number of star clusters and regions.
+ The absolute Z£-baud magnitudes of the identified SSCs range up to MyL7.2inag.," The absolute $H$ -band magnitudes of the identified SSCs range up to $M_H=-17.2\,$ mag."
+ A huge fraction of the region population (exchiding the nuclear emission) shows huuimosities iu excess of that of 30 Doradus. the prototypical giant region.," A large fraction of the region population (excluding the nuclear emission) shows luminosities in excess of that of 30 Doradus, the prototypical giant region."
+ The main characteristic of the Pao emissiou. iu the isolated LIRGs iu our sample is the lack of stroug uuclear cluission., The main characteristic of the $\alpha$ emission in the isolated LIRGs in our sample is the lack of strong nuclear emission.
+ Most of the regions are distributed along the spiral rius of the galaxies., Most of the regions are distributed along the spiral arms of the galaxies.
+ The interacting/imereine LIRGs ou the other hand. show bright unclear Pan enission together with widely spread star formation along the the spiral iris and at the interface of interacting ealaxies.," The interacting/merging LIRGs on the other hand, show bright nuclear $\alpha$ emission together with widely spread star formation along the the spiral arms and at the interface of interacting galaxies."
+ The fraction of nuclear Pan eiissiou to the total cnussion varies from system to system. as it depends ou factors sch as the age of the interaction process. the initial eas content of the ealaxics and the relative masses of the ealaxies.," The fraction of nuclear $\alpha$ emission to the total emission varies from system to system, as it depends on factors such as the age of the interaction process, the initial gas content of the galaxies and the relative masses of the galaxies."
+ We have analyzed the properties — bIuninosities. sizes aud huuimositv functions — of regions in LIRGs at au uuprecedeuted spatial resolution (between 15 and 78 pe).," We have analyzed the properties – luminosities, sizes and luminosity functions – of regions in LIRGs at an unprecedented spatial resolution (between 15 and 78 pc)."
+ Cdant reeious are ubiquitous in LIRGs aud are located not only near the nuclei of interacting galaxies. but also at the interface of iuteracting ealaxies and along the spiral arms.," Giant regions are ubiquitous in LIRGs and are located not only near the nuclei of interacting galaxies, but also at the interface of interacting galaxies and along the spiral arms."
+ This population of lughiy luminous regions is uot observed iu uormal galaxies., This population of highly luminous regions is not observed in normal galaxies.
+ We lave fitted power laws to the region LFs of Arp 299 and NGC 3256 and found that the iudices are within the values previously measured iu the disks of normal galaxies., We have fitted power laws to the region LFs of Arp 299 and NGC 3256 and found that the indices are within the values previously measured in the disks of normal galaxies.
+ We lave compared the properties of the regions in LIRCGs with a small sample of normal galaxies observed with the same spatial resolution and found that eiaut veeious are more common in LIRGs than in normal ealaxies., We have compared the properties of the regions in LIRGs with a small sample of normal galaxies observed with the same spatial resolution and found that giant regions are more common in LIRGs than in normal galaxies.
+ The iueasured sizes of eiut regious in, The measured sizes of giant regions in
+assumptions. ie. the self-similarity of the solutions).,"assumptions, i.e. the self-similarity of the solutions)."
+ PelleGer&ον(1992) further developed the general theory. of non-sell-similar solutions of hyvdromagnetie disk winds., \citet{Pelletier92} further developed the general theory of non-self-similar solutions of hydromagnetic disk winds.
+ In these models the transler of excess energy (o (he escaping particles is made at the expense ol the rotational energv of the matter in the disk and it is mediated by the magnetic field., In these models the transfer of excess energy to the escaping particles is made at the expense of the rotational energy of the matter in the disk and it is mediated by the magnetic field.
+ While this (vpe of model gained popularity. the issue of jet/outlfow formation took a different turn. with the introduction of acdvection-dominated accretion flow (ADAF) (ο...Naravan&Yi1994:Mlanmotoetal. 1997).," While this type of model gained popularity, the issue of jet/outlfow formation took a different turn with the introduction of advection-dominated accretion flow (ADAF) \citep[e.g.,][]{Narayan94,Manmoto97}."
+. These radiativelv inefficient. accretion [lows (RIAF) were found to have positive Bernoulli integral of the flow and could therefore fulfill the condition necessary for the launching of jet/wind outflows: as such they present potentially interesting sites for the oriein of such oulllows., These radiatively inefficient accretion flows (RIAF) were found to have positive Bernoulli integral of the flow and could therefore fulfill the condition necessary for the launching of jet/wind outflows; as such they present potentially interesting sites for the origin of such outflows.
+ The positivity of the Bernoulli integral has been discussed and analysed by Blanclord&Begelman(1999) who pointed out Chat 1 is due to the combination of energy. transfer by (he viscous torques [rom the inner to the outer sections of the flow (the gas of the flow becomes bounded al its inner edge) and the local dissipation of the flows azimuthal kinetic energy which is not radiated away but stored in Che fluid to increase its internal energv., The positivity of the Bernoulli integral has been discussed and analysed by \citet{BB99} who pointed out that it is due to the combination of energy transfer by the viscous torques from the inner to the outer sections of the flow (the gas of the flow becomes bounded at its inner edge) and the local dissipation of the flow's azimuthal kinetic energy which is not radiated away but stored in the fluid to increase its internal energy.
+ The latter authors then argued Chat ἰ excess energv can be carried away (o infinity (along with some fraction of the accreti mass and angular moment) to produce continuous outflows from all radii to infinitv. whi leaving (he remaining flow with negative Bernoulli constant to naturally acerete onto the compact object: advection-dominated. inflow-outllow solution (ADIOS)., The latter authors then argued that the excess energy can be carried away to infinity (along with some fraction of the accreting mass and angular momentum) to produce continuous outflows from all radii to infinity while leaving the remaining flow with negative Bernoulli constant to naturally accrete onto the compact object: advection-dominated inflow-outflow solution (ADIOS).
+ In (his case. while (he necessary excess energy is (transferred by (the viscous torques from the flows more highly bound inner section. (he necessary separation of mass to components with positive and negalive total energv is still left unspecified.," In this case, while the necessary excess energy is transferred by the viscous torques from the flow's more highly bound inner section, the necessary separation of mass to components with positive and negative total energy is still left unspecified."
+ An altogether different model (hat offers a simplified picture of such a separation was presented by Subramanianetal.(1999) who proposed that in the tenuous. collisionless plasma of an ADAF particles (protons) could be accelerated. via a second-orcler Fermi acceleration bv (he shear motions of the underlying quasi-Ixeplerian azimuthial flow.," An altogether different model that offers a simplified picture of such a separation was presented by \citet{Subramanian99} who proposed that in the tenuous, collisionless plasma of an ADAF particles (protons) could be accelerated via a second-order Fermi acceleration by the shear motions of the underlying quasi-Keplerian azimuthal flow."
+ Thev ren argued that if sulliciently laree pressure is built in (he accelerated. particle proton population (the electrons generally lose energy on time scales short compared to their transit me through the svstem and cannot build an energy densitv that could be dynamically nuportant) aud for favorable geometries of the disk magnetic field (large scale poloidal loops wal open up above the disk) (he relativistic particle population could naturally (through je action of the gravitational field) segregate itself [rom the non-relativistic one. carrving ME to infinity only the accelerated (Ezi imc) portion of the disk plasma.," They then argued that if sufficiently large pressure is built in the accelerated particle proton population (the electrons generally lose energy on time scales short compared to their transit time through the system and cannot build an energy density that could be dynamically important) and for favorable geometries of the disk magnetic field (large scale poloidal loops that open up above the disk) the relativistic particle population could naturally (through the action of the gravitational field) segregate itself from the non-relativistic one, carrying off to infinity only the accelerated $E \gsim m_pc^2$ ) portion of the disk plasma."
+ In this case the eneine is a combination of the particle acceleration and the action of the gravitational fiel., In this case the engine is a combination of the particle acceleration and the action of the gravitational field.
+ Finally. a model along the same lines was proposed by Contopoulos&Ixazanas(1995) who suggested that even in the case of a completev turbulent magnetic field. a separation of (he relativistic and non-relativistic particle populaAions is possible through the production οἱ relativistic neutrons in the collisions of the relativistic protons with the ambient plasma ancl the ensuing production of relativistic neutrons.," Finally, a model along the same lines was proposed by \citet{Contopoulos95} who suggested that even in the case of a completely turbulent magnetic field, a separation of the relativistic and non-relativistic particle populations is possible through the production of relativistic neutrons in the collisions of the relativistic protons with the ambient plasma and the ensuing production of relativistic neutrons."
+" T1ο subsequent decay of neutrons back into protons produces then a proton fluid in regions o[ space devoid of inertia whose ratio (and hence its asvanptotie Lorentz Fact) depends only on the ratio ήστι. where Ris the size of the svstem and 7, the neutron lie time and can lead to highly relativistic flows for black hole masses Af210A.."," The subsequent decay of neutrons back into protons produces then a proton fluid in regions of space devoid of inertia whose energy-to-mass ratio (and hence its asymptotic Lorentz factor) depends only on the ratio $R/c \tau_{\rm
+n}$, where $R$ is the size of the system and $\tau_{\rm n}$ the neutron life time and can lead to highly relativistic flows for black hole masses $M \gsim 10^8 \Msun$."
+ In the present note we lollow a similar simplifiel view to study outflows in objects powered bv accretion: We consider the presence of (2-dimensional) shocks as a means of dissipation of (he accretion kinetic energy in a [ashion similar to that considered by, In the present note we follow a similar simplified view to study outflows in objects powered by accretion: We consider the presence of (2-dimensional) shocks as a means of dissipation of the accretion kinetic energy in a fashion similar to that considered by
+the group regime. where fossils fall on the Lj;—c relation of non-fossil galaxy groups.,"the group regime, where fossils fall on the $L_R-\sigma$ relation of non-fossil galaxy groups."
+ We will discuss these features in Section 6 in the light of other scaling relations., We will discuss these features in Section 6 in the light of other scaling relations.
+ Group velocity dispersions are based on our spectroscopic observations of the sample in Table | and NED for the rest of the systems., Group velocity dispersions are based on our spectroscopic observations of the sample in Table 1 and NED for the rest of the systems.
+ They are calculated using the following relation. also used in the comparison sample of Osmond&Ponman(20043:3/2). +.," They are calculated using the following relation, also used in the comparison sample of \citet{osmond04}:, ."
+ This estimator corrects for a statistical bias. which results if one uses the normal unbiased estimator for o7 and then takes the square root to obtain e (which is then not unbiased).," This estimator corrects for a statistical bias, which results if one uses the normal unbiased estimator for $\sigma^2$ and then takes the square root to obtain $\sigma$ (which is then not unbiased)."
+ This correction is the origin of the term 3/2 (rather than |) in the denominator of the above equation., This correction is the origin of the term 3/2 (rather than 1) in the denominator of the above equation.
+ If excess X-ray luminosity in fossils were the only difference between the X-ray properties of fossils and non-fossil groups. then hey wouldbe expected to deviate from the Ly7 relation known or non-fossil groups and clusters.," If excess X-ray luminosity in fossils were the only difference between the X-ray properties of fossils and non-fossil groups, then they wouldbe expected to deviate from the $L_X-T$ relation known for non-fossil groups and clusters."
+ This appeared to be the case rom an earlier ROSAT study (Jonesetal.2003). based on very imited statistics., This appeared to be the case from an earlier ROSAT study \citep{jones03} based on very limited statistics.
+ Two fossil groups. for which Jonesetal.(2003) could measure the temperature. were found to have high. X-ray uminosity for their gas temperature.," Two fossil groups, for which \citet{jones03} could measure the temperature, were found to have high X-ray luminosity for their gas temperature."
+ Based on this finding. it was argued that fossils are low-entropy systems due to their higher gas density. in comparison to non-fossils.," Based on this finding, it was argued that fossils are low-entropy systems due to their higher gas density, in comparison to non-fossils."
+ However we have shown in Khosroshahietal.(2006). that the X-ray temperature of the RX 11416.442315 was underestimated in the analysis., However we have shown in \citet{kmpj06} that the X-ray temperature of the RX J1416.4+2315 was underestimated in the analysis.
+ The present wider study shows that the above system is not an exception and. as it is seen in Fig 4.. fossils fall on the conventional byT. relation of non-fossil groups and clusters.," The present wider study shows that the above system is not an exception and, as it is seen in Fig \ref{LT}, fossils fall on the conventional $L_X-T$ relation of non-fossil groups and clusters."
+ Hence. if from our earlier arguments we assert that is enhanced in fossils. then it follows that they must also have £Lelevatedx. mean temperature values. such that they remain on the standard group Ly7 relation.," Hence, if from our earlier arguments we assert that $L_X$ is enhanced in fossils, then it follows that they must also have elevated mean temperature values, such that they remain on the standard group $L_X-T$ relation."
+ Mahdavi&Geller(2001):XueWu(9000). have presented Lx70 relations for clusters and groups.," \citet{mahdavi01,xue00} have presented $L_X-\sigma$ relations for clusters and groups."
+ Osmond(2004) found a slope of 2.31 + 0.61 in Lxσ for heir sample of galaxy groups with intergalactic X-ray emission. flatter than the value of 4.5+1.1 found by Helsdon&Ponman (2000)..," \citet{osmond04} found a slope of 2.31 $\pm$ 0.61 in $L_X-\sigma$ for their sample of galaxy groups with intergalactic X-ray emission, flatter than the value of $4.5 \pm 1.1$ found by \citet{helsdon00}. ."
+ There is a good deal of scatter in the relation. which may in »art account for the disagreement between various studies.," There is a good deal of scatter in the relation, which may in part account for the disagreement between various studies."
+ While Ponmanetal. (1996). Mulchaey&Zabludoff 1908): Helsdon&Pon-man(2000). and Mahdavi&Geller(2001). find that groups are consistent with the cluster-relation slope of z+ . Mahdavi (1997.2000) and Xue&Wu(2000). find significantly flatter relations in groups with a slope similar to the finding of Osmond (2004). ," While \citet{ponman96}, \citet{mz98}; \citet{helsdon00} and \citet{mahdavi01} find that groups are consistent with the cluster-relation slope of $\approx 4$ , Mahdavi (1997,2000) and \citet{xue00} find significantly flatter relations in groups with a slope similar to the finding of \citet{osmond04}. ."
+Figure 5. shows the distribution of fossil groups inthe plane of Lx.—0 along with the non-fossil groups and clusters., Figure \ref{lxsigma} shows the distribution of fossil groups inthe plane of $L_X-\sigma$ along with the non-fossil groups and clusters.
+ Fossils appear more X-ray luminous thannon-fossil groups for a given, Fossils appear more X-ray luminous thannon-fossil groups for a given
+reliability.,reliability.
+" In refobserved,, the direct observations are compared with the images and in refresults,, values measured on the sharpened images are reported."," In \\ref{observed}, the direct observations are compared with the de-convolved images and in \\ref{results}, values measured on the sharpened images are reported."
+" In refring,, one-dimensional cuts along the major axis of the observed and de-convolved iimage are displayed."," In \\ref{ring}, one-dimensional cuts along the major axis of the observed and de-convolved image are displayed."
+" It would have been more natural to select the observation atum,, having the highest spatial resolution, but the sscan map data are of considerably higher S/N, outweighing the apparent resolution advantage of the shorter wavelength data."," It would have been more natural to select the observation at, having the highest spatial resolution, but the scan map data are of considerably higher S/N, outweighing the apparent resolution advantage of the shorter wavelength data."
+" In the analysed sub-frame, the flux was conserved within by the MEM routine."," In the analysed sub-frame, the flux was conserved within by the MEM routine."
+" Prior to the de-convolution, a stellar point source with photospheric flux of mmJy at refresults)) was subtracted from the PACS image."," Prior to the de-convolution, a stellar point source with photospheric flux of mJy at \\ref{results}) ) was subtracted from the PACS image."
+" The resulting sharpened image reveals a central and a central depression with a depth of about2%,, which is consistent with the debris residing in a ring or belt around the star."," The resulting sharpened image reveals a central and a central depression with a depth of about, which is consistent with the debris residing in a ring or belt around the star."
+" With standard assumptionsregarding the emitting grains (astronomical silicates, a blow-out size limit ag,=0.6 ffor the VV star and a —3.5 power law index for the size distribution) we find that the ssurface brightness profiles along the major and minor axes are well reproduced assuming a disk inclination of about aand a two-parameter model for the surface density, X(r)."," With standard assumptionsregarding the emitting grains (astronomical silicates, a blow-out size limit $a_{\rm min}= 0.6$ for the V star and a $-3.5$ power law index for the size distribution) we find that the surface brightness profiles along the major and minor axes are well reproduced assuming a disk inclination of about and a two-parameter model for the surface density, $\Sigma(r)$."
+" These parameters are the peak density position, rmax, and the power law index of the surface density profile for r>rmax."," These parameters are the peak density position, $r_{\rm max}$, and the power law index of the surface density profile for $r > r_{\rm max}$."
+" The best fit to the surface brightness profiles is consistent with a ring-like disc, having values of rmax~85 AAU and X(r>rg)οςr3, respectively."," The best fit to the surface brightness profiles is consistent with a ring-like disc, having values of $r_{\rm max} \sim 85$ AU and $\Sigma(r > r_{\rm max}) \propto r^{-3}$, respectively."
+" A more elaborate model, with the size distribution computed self-consistently and taking into account the profiles also at other wavelengths, will be presented by Augereau et al. ("," A more elaborate model, with the size distribution computed self-consistently and taking into account the profiles also at other wavelengths, will be presented by Augereau et al. ("
+in prep.).,in prep.).
+ This roughly AAU wide ring or belt at about AAU from the star appears similar to the EKB of the Solar System., This roughly AU wide ring or belt at about AU from the star appears similar to the EKB of the Solar System.
+" Based on an analogy with the debris disc around Fomalhaut and on theoretical expectations, it is quite possible that another gas giant planet, ,cc, could be orbiting the star inside the inner belt edge."," Based on an analogy with the debris disc around Fomalhaut and on theoretical expectations, it is quite possible that another gas giant planet, c, could be orbiting the star inside the inner belt edge."
+" Given the age of the system, GGyr, the direct detection of ,cc, for instance by means of coronography, can be expected to be hard (see,e.g.,Beichmanetal. 2006).."," Given the age of the system, Gyr, the direct detection of c, for instance by means of coronography, can be expected to be hard \citep[see, e.g.,][]{beichman2006}. ."
+" Based on imaging observations with PACS in the three photometric bands at um,, aand wwe find that"," Based on imaging observations with PACS in the three photometric bands at , and we find that"
+"more possibly ambipolar diffusion. cau become dominant in the high density cold gas. the turbulent diffusion in the carly stages of accretion is able to form a light aud large rotationally supported disk very quickly. iun only a few LO! yy,","more possibly ambipolar diffusion, can become dominant in the high density cold gas, the turbulent diffusion in the early stages of accretion is able to form a light and large rotationally supported disk very quickly, in only a few $10^4$ yr."
+ Finally. we should remark that other mechanisms to remove or reduce the effects of the magnetic braking in the inner regions of protostellar cores have been also investigated in the literature receutly.," Finally, we should remark that other mechanisms to remove or reduce the effects of the magnetic braking in the inner regions of protostellar cores have been also investigated in the literature recently."
+ Heunebelle&Cia-rdi(2009). verified that the maeuetic brakiug efficiency may decrease significantly when the rotation axis of the core is unisaligued with the direction of the regular magnetic field., \citet{hennebelle_ciardi_2009} verified that the magnetic braking efficiency may decrease significantly when the rotation axis of the core is misaligned with the direction of the regular magnetic field.
+" They claim that even for small angeles of the order of 10.20"" there are significant differences with respect to the aligned case;", They claim that even for small angles of the order of $10-20^o$ there are significant differences with respect to the aligned case.
+ Also. in a concoiitaut work to the present one. Wrasnopolskyetal.(2011) have exanuned the Wall effect on disk formation.," Also, in a concomitant work to the present one, \citet{krasnopolsky_etal_2011} have examined the Hall effect on disk formation."
+ They found that a ITall-àuduced magnetic torque can diffuse magnetic flux outward aud generate a rotationally supported disk in the collapsing flow. even when the core is initially however the spun-up material remains too sub-Keplerian (Lictal.2011).," They found that a Hall-induced magnetic torque can diffuse magnetic flux outward and generate a rotationally supported disk in the collapsing flow, even when the core is initially non-rotating, however the spun-up material remains too sub-Keplerian \citep{li_etal_2011}."
+. Of course. in the uear future. these imechamisus must be tested along with the just proposed turbulent magnetic reconnection aud even with ambipolar diffusou. iu order to assess the relative müportance of cach effect ou disk formation aud evolution.," Of course, in the near future, these mechanisms must be tested along with the just proposed turbulent magnetic reconnection and even with ambipolar diffusion, in order to assess the relative importance of each effect on disk formation and evolution."
+ Nonetheless. «λος ATID turbuleuce is expected to be present iu these magnetic cores (e... Dallesteros-Paredes&:MacLimactal 2010.. ind references therein). turbulent reconnection arises as a natural wuechauisia for removing magnetic flux excess and allowing the formation of these disks.," Nonetheless, since MHD turbulence is expected to be present in these magnetic cores (e.g., \citealt{ballesteros-paredes_maclow_2002, melioli_etal_2006, leao_etal_2009, santos-lima_etal_2010}, and references therein), turbulent reconnection arises as a natural mechanism for removing magnetic flux excess and allowing the formation of these disks."
+ lu recent nunierical study in SNLO. we showed that magnetic reconnection iu a turbulent cloud cau efficicutly transport magnetic flux from the iuner denser regions to the periphery of the cloud thus enabling the cloud. to collapse to form a star.," In recent numerical study in SX10, we showed that magnetic reconnection in a turbulent cloud can efficiently transport magnetic flux from the inner denser regions to the periphery of the cloud thus enabling the cloud to collapse to form a star."
+ Tere. also by iuneaus of fully 3D ATID sinmlatious. we have investigated the same mechanisni acting in a votating collapsing cloud core and shown that the magnetic flux excess of the inner regions of the system can be effectively removed allowing the formation of a rotationally sustained protostellar disk.," Here, also by means of fully 3D MHD simulations, we have investigated the same mechanism acting in a rotating collapsing cloud core and shown that the magnetic flux excess of the inner regions of the system can be effectively removed allowing the formation of a rotationally sustained protostellar disk."
+ Another cupirical finding in SNIO is that the eficiency ofthe magnetic field expulsion via recounectiou diffisivity increases with the source gravitational field., Another empirical finding in SX10 is that the efficiency of the magnetic field expulsion via reconnection diffusivity increases with the source gravitational field.
+ This is a natural cousequeuce ofdiffusion iu the presence of the gravitational field which pulls oue component (gas and does not act on the other weightless courponeut anaenetie feld)., This is a natural consequence of diffusion in the presence of the gravitational field which pulls one component (gas) and does not act on the other weightless component (magnetic field).
+ Di terms of the problem im haud. this implies that more massive protostars can induce magnetic field segregation even for weaker level of turbulence.," In terms of the problem in hand, this implies that more massive protostars can induce magnetic field segregation even for weaker level of turbulence."
+ We plan to explore nuucerically this issue iu a forthcoming work., We plan to explore numerically this issue in a forthcoming work.
+ Tn this paper we showed that the concept of reconnection diffusion (L05) successfully works iu the formation of protostellar disks., In this paper we showed that the concept of reconnection diffusion (L05) successfully works in the formation of protostellar disks.
+ Together with our earlier testing of magnetic field removal through recounection diffusion from collapsing clouds this paper supports a considerable change of the paradieui of star formation., Together with our earlier testing of magnetic field removal through reconnection diffusion from collapsing clouds this paper supports a considerable change of the paradigm of star formation.
+ Tudeed. in the presence of reconnection diffusion. there is no necessity to appeal to ambipolar diffusion.," Indeed, in the presence of reconnection diffusion, there is no necessity to appeal to ambipolar diffusion."
+ The latter wav still be nuportaut iu low ionization. low turbulence cuvirouments. but. iu anv case. the domain of its applicability is seriously challeuged.," The latter may still be important in low ionization, low turbulence environments, but, in any case, the domain of its applicability is seriously challenged."
+ The application of the reconnection diffusion concept to protostellar disk formation aud. i a more eencral framework. to accretion disks in seueral is natural as the disks are expected to be turbulent. chabling our appeal to LV99 model of fast recounection.," The application of the reconnection diffusion concept to protostellar disk formation and, in a more general framework, to accretion disks in general, is natural as the disks are expected to be turbulent, enabling our appeal to LV99 model of fast reconnection."
+" An Huportant accepted source of turbulence iu accretion disks is the well known magueto-rotational instability (MRI) (Chandrasekhar 1960. Balbus ITawleyv 1991)""s but at carlicy stages turbulence can be induced by the ντοςπασα motions associated with the disk. formation."," An important accepted source of turbulence in accretion disks is the well known magneto-rotational instability (MRI) (Chandrasekhar 1960, Balbus Hawley 1991), but at earlier stages turbulence can be induced by the hydrodynamical motions associated with the disk formation."
+ Turbulence is ubiquitous in astrophysical environments as it follows from theoretical considerations based on the high Revuolds uuubers of astrophysical flows and is stronely supported bv studies of spectra of the interstellar electron ceusitv fluctuations (sec Arvinstrongetal.1995:Chepurnov&Lazarian 2010)) as well as of WT (Lazarian2009. for a review aud references therein: Chepurnovctal. 2010)) and CO lines (see Padoanetal. 200901).," Turbulence is ubiquitous in astrophysical environments as it follows from theoretical considerations based on the high Reynolds numbers of astrophysical flows and is strongly supported by studies of spectra of the interstellar electron density fluctuations (see \citealt{armstrong_etal_1995, chepurnov_lazarian_2010}) ) as well as of HI \citealt{lazarian_2009} for a review and references therein; \citealt{chepurnov_etal_2010}) ) and CO lines (see \citealt{padoan_etal_2009}) )."
+ The application of the recomnection diffusion mechanisia to already formed accretion disks will be investigated im detail elsewhere., The application of the reconnection diffusion mechanism to already formed accretion disks will be investigated in detail elsewhere.
+ It should be noted however that former studies of the injection of turbulence in accretion disks have shown that at this stage turbulence may be ineffective to maenetic fiux diffusion outward (Rothstein&Lovelace2008)., It should be noted however that former studies of the injection of turbulence in accretion disks have shown that at this stage turbulence may be ineffective to magnetic flux diffusion outward \citep{rothstein_lovelace_2008}.
+. Appealing to the LV99 inodel of fast maguetic reconnection and inspired by the successful demonstration of removal of magnetic feld through reconnection diffusion from nunerical models of nolecular clouds in SX10 we have performed uuucerical simulations and demonstrated that: l., Appealing to the LV99 model of fast magnetic reconnection and inspired by the successful demonstration of removal of magnetic field through reconnection diffusion from numerical models of molecular clounds in SX10 we have performed numerical simulations and demonstrated that: 1.
+ The concept of reconnection diffusion is applicable o the formation of protostellar disks with radius ~100 AU., The concept of reconnection diffusion is applicable to the formation of protostellar disks with radius $\sim 100$ AU.
+ The extension of this concept to accretion disks is OYeseen., The extension of this concept to accretion disks is foreseen.
+ 2., 2.
+ In the eravitational field. reconnection diffusion uitigates magnetic breaking allowing the formation of xotostellar disks.," In the gravitational field, reconnection diffusion mitigates magnetic breaking allowing the formation of protostellar disks."
+ 3., 3.
+ The removal of magnetic field through recounectiou diffusion is fast cuough to explain observatious without he necessity of appealing to cuhanced fluid resistivity., The removal of magnetic field through reconnection diffusion is fast enough to explain observations without the necessity of appealing to enhanced fluid resistivity.
+stellar masses inferred from the CBO7 library are on average 0.12 dex lower than the BCO3-based ones. with a large (0.17 dex) scatter.,"stellar masses inferred from the CB07 library are on average 0.12 dex lower than the BC03-based ones, with a large (0.17 dex) scatter."
+ The lack of a clear trend with stellar mass or redshift of the ratio of the two estimates translates into a lack of systematic difference between the best-fit Schechter parameters in the two cases., The lack of a clear trend with stellar mass or redshift of the ratio of the two estimates translates into a lack of systematic difference between the best-fit Schechter parameters in the two cases.
+ The largest disagreement was found at 2— 3. where the effect of the TP-AGB phase is expected to be the most important.," The largest disagreement was found at $z \sim 2 - 3$ , where the effect of the TP-AGB phase is expected to be the most important."
+ The main result of this study is the steepening of the faint- slope: the value of« increases from —1.4420.03 at z~0.8 to —1.56£0.16 at z~3. and then flattens up to z~4.," The main result of this study is the steepening of the faint-end slope: the value of $\alpha$ increases from $-1.44 \pm 0.03$ at $z\sim 0.8$ to $-1.86 \pm 0.16$ at $z\sim 3$, and then flattens up to $z\sim 4$."
+ We have confirmed the steepening of the low-mass end. which had been pointed out by previous authors. with deeper and higher quality photometry.," We have confirmed the steepening of the low-mass end, which had been pointed out by previous authors, with deeper and higher quality photometry."
+ Our results are unaffected by degeneracies in the M parameter. and they are insensitive to the choice of either the stellar templates or the functional shape fitted to the GSMF. as well as to the limitations of the small area covered by ERS observations.," Our results are unaffected by degeneracies in the $M^*$ parameter, and they are insensitive to the choice of either the stellar templates or the functional shape fitted to the GSMF, as well as to the limitations of the small area covered by ERS observations."
+ We computed the SMD as a function of redshift and compared it with the integrated star formation histories derived by Hopkins&Beacom(2006) and Reddy&Steidel(2009)., We computed the SMD as a function of redshift and compared it with the integrated star formation histories derived by \cite{hopkins06} and \cite{reddy09}.
+ The finer sampling of the GSMF at low masses and the steep inferred faint-end slopes determine the higher SMD estimates at >2 than most previous works. solving the disagreement observed by previous authors between the SMD and the integrated SFRD at these redshifts.," The finer sampling of the GSMF at low masses and the steep inferred faint-end slopes determine the higher SMD estimates at $z>2$ than most previous works, solving the disagreement observed by previous authors between the SMD and the integrated SFRD at these redshifts."
+ However. despite the steep GSMF that we find. the integrated star formation history still exceeds the direct measure of the SMD at z~2 by a factor of ~2—3. even when our data are analysed together with the results of previous large surveys to ensure a good sampling of also the bright-end tail of the GSMF.," However, despite the steep GSMF that we find, the integrated star formation history still exceeds the direct measure of the SMD at $z\sim 2$ by a factor of $\sim 2-3$, even when our data are analysed together with the results of previous large surveys to ensure a good sampling of also the bright-end tail of the GSMF."
+ Finally. we compared our GSMF and SMD estimates with the predictions of four models of galaxy formation and evolution.," Finally, we compared our GSMF and SMD estimates with the predictions of four models of galaxy formation and evolution."
+ All models predict a larger abundance of low mass galaxies than observations. at least up to z~3.," All models predict a larger abundance of low mass galaxies than observations, at least up to $z\sim3$."
+ They also underestimate the stellar mass of high mass galaxies in. the highest redshift bin. although cosmic variance effects prevent us from drawing firm conclusions. at these redshifts.," They also underestimate the stellar mass of high mass galaxies in the highest redshift bin, although cosmic variance effects prevent us from drawing firm conclusions at these redshifts."
+ The overabundance of low mass galaxies translates into a general overestimation of the total SMD with respect to the data up to zo2. while this density is underestimated at z>3 owing to the dearth of massive galaxies at these redshifts.," The overabundance of low mass galaxies translates into a general overestimation of the total SMD with respect to the data up to $z\sim 2$, while this density is underestimated at $z\gtrsim 3$ owing to the dearth of massive galaxies at these redshifts."
+ The exact degree of disagreement depends on the particular model., The exact degree of disagreement depends on the particular model.
+ Future CANDELS data will cover a larger sky area and allow a finer sampling of both the bright-end of the GSMF and its normalization. and at the same time they will be deep enough to accurately probe the GSMF faint-end.," Future CANDELS data will cover a larger sky area and allow a finer sampling of both the bright-end of the GSMF and its normalization, and at the same time they will be deep enough to accurately probe the GSMF faint-end."
+ These. together with spectroscopic follow-up campaigns. will reduce the uncertainties in the stellar masses. and they will significantly improve our results and our understanding of the stellar mass assembly process.," These, together with spectroscopic follow-up campaigns, will reduce the uncertainties in the stellar masses, and they will significantly improve our results and our understanding of the stellar mass assembly process."
+in the immediate vicinity of the black hole.,in the immediate vicinity of the black hole.
+ Furthermore. at millimeter wavelengths the blurring due to interstellar electron scattering ts subdominant.," Furthermore, at millimeter wavelengths the blurring due to interstellar electron scattering is subdominant."
+ Thus. at wavelengths of 1.3mm and below it is possible to image the emitting region surrounding Ser A*.," Thus, at wavelengths of $1.3\,\mm$ and below it is possible to image the emitting region surrounding Sgr A*."
+ Even with the strong gravitational lensing in the vicinity of the horizon. imaging the immediate vicinity of the black hole requires extraordinary resolutions.," Even with the strong gravitational lensing in the vicinity of the horizon, imaging the immediate vicinity of the black hole requires extraordinary resolutions."
+ The silhouette cas= by the horizon on the surrounding emission is roughly 53+2 ," The silhouette cast by the horizon on the surrounding emission is roughly $53\pm2\,\muas$ ."
+At the present time. this resolution is accessible only via asl...millimeter-wavelength very-long baseline interferometry (mm-VLBI).," At the present time, this resolution is accessible only via millimeter-wavelength very-long baseline interferometry (mm-VLBI)."
+ VLBI observations of Ser A* at 1.4mm using the Institut de Radioastronomie Milliméttrique (IRAM) 30m telescope at Pico Veleta and one of the 15m dishes at Plateau de Bure. produced the size estimate of 110—60j/as. with the large uncertainties due to limited calibration accuracy (?)..," VLBI observations of Sgr A* at $1.4\,\mm$ using the Institut de Radioastronomie Milliméttrique (IRAM) $30\,\m$ telescope at Pico Veleta and one of the $15\,\m$ dishes at Plateau de Bure, produced the size estimate of $110\pm60\,\muas$, with the large uncertainties due to limited calibration accuracy \citep{Kric_etal:98}."
+ The first successful mm-VLBI observation of Ser A* with Earth-scale baselines was performed in April. 2007. during which visibilities were measured on the 4.6«10°km baseline between Mauna Kea. Hawaii to Mount Graham. Arizona (?)..," The first successful mm-VLBI observation of Sgr A* with Earth-scale baselines was performed in April, 2007, during which visibilities were measured on the $4.6\times10^3\,\km$ baseline between Mauna Kea, Hawaii to Mount Graham, Arizona \citep{Doel_etal:08}."
+ By fitting these with a gaussian model. ?— found a typical intrinsic source size of 3773jas (after correcting for the sub-dominant broadening due to interstellar electron scattering). smaller than the black hole silhouette.," By fitting these with a gaussian model, \citet{Doel_etal:08} found a typical intrinsic source size of $37^{+5}_{-3}\,\muas$ (after correcting for the sub-dominant broadening due to interstellar electron scattering), smaller than the black hole silhouette."
+ Since that time a number of groups have analyzed the 2007 mm-VLBI data using various physically motivated accretion models for the emission region (????).. inferring from these efforts the black hole spin vector.," Since that time a number of groups have analyzed the 2007 mm-VLBI data using various physically motivated accretion models for the emission region \citep{Brod_etal:09,Huan-Taka-Shen:09,Mosc_etal:09,Dext-Agol-Frag-McKi:10}, inferring from these efforts the black hole spin vector."
+ Despite finding generally similar results. these have been limited by the lack of multiple long baseline observations and the limited north-south coverage obtained.," Despite finding generally similar results, these have been limited by the lack of multiple long baseline observations and the limited north-south coverage obtained."
+ Recently. a second. and considerably set of mm-VLBI observations have been (?).. largerproviding the opportunity to revisit. and substantially reportedimprove. constraints the black hole spin and accretion physics.," Recently, a second, and considerably larger set of mm-VLBI observations have been reported \citep{Fish_etal:10}, providing the opportunity to revisit, and substantially improve, constraints upon the black hole spin and accretion physics."
+ Here we uponreport upon the first effort to do this using a physically motivated accretion model. similar to that described in ?.. that fits the known spectral and polarization properties of Ser A*.," Here we report upon the first effort to do this using a physically motivated accretion model, similar to that described in \citet{Brod_etal:09}, that fits the known spectral and polarization properties of Sgr A*."
+ In addition to improving the resulting parameter estimation. it is now possible to identify statistical signatures of both the asymmetry of the image and the importance of the underlying physics that governs the image morphology.," In addition to improving the resulting parameter estimation, it is now possible to identify statistical signatures of both the asymmetry of the image and the importance of the underlying physics that governs the image morphology."
+ Section 2. summarizes the full set of mm-VLBI observations we consider.," Section \ref{sec:SoO}
+ summarizes the full set of mm-VLBI observations we consider."
+ Section 3. describes the models we consider and how the resulting visibility data is produced., Section \ref{sec:VM} describes the models we consider and how the resulting visibility data is produced.
+ How models are compared and the parameter estimates are produced ts discussed in Section 4.., How models are compared and the parameter estimates are produced is discussed in Section \ref{sec:BDA}.
+ The fitting process and results are presented in Section 5.. and our best estimates for the black hole spin vector can be found in Section 6..," The fitting process and results are presented in Section \ref{sec:MF}, and our best estimates for the black hole spin vector can be found in Section \ref{sec:EBHS}."
+ Section 7. describes the implications for different potential future observations.," Section \ref{sec:OFO}
+ describes the implications for different potential future observations."
+ Finally concluding remarks are collected in Section 8.., Finally concluding remarks are collected in Section \ref{sec:C}.
+ In the analysis presented. here we make full use of the recent observations described in?) and ?.., In the analysis presented here we make full use of the recent observations described in \citet{Fish_etal:10} and \citet{Doel_etal:08}.
+ In both cases. observations targeting Ser A¥ were made at 1.3mm using the Submillimeter Telescope (SMT)) on Mt. Graham in Arizona. 10m dishes in the Combined Array for Research in Millimeter-wave Astronomy (CARMA)) at Cedar Flat. California. and the James Clerk Maxwell Telescope (JCMT)) located on Mauna Kea. Hawaii.," In both cases, observations targeting Sgr A* were made at $1.3\,\mm$ using the Submillimeter Telescope ) on Mt. Graham in Arizona, $10\,\m$ dishes in the Combined Array for Research in Millimeter-wave Astronomy ) at Cedar Flat, California, and the James Clerk Maxwell Telescope ) located on Mauna Kea, Hawaii."
+ 9? report upon measurements obtained on the nights of the April. 11 12. 2007. using theJCMT.. and a single dish.," \citet{Doel_etal:08} report upon measurements obtained on the nights of the April, 11 12, 2007, using the, and a single dish."
+ 19 visibility amplitudes were obtained on the and baselines. with an upper limit on April 11th. 2007 along the baseline.," 19 visibility amplitudes were obtained on the and baselines, with an upper limit on April 11th, 2007 along the baseline."
+ The locations of these observations on the u—v plane are indicated in the lower-left panel of Figure 1.. labeled 2007.," The locations of these observations on the $u$ $v$ plane are indicated in the lower-left panel of Figure \ref{fig:Vobs}, labeled 2007."
+ Signal-to-noise ratios typical of the short and long baselines are 8 and 4. respectively.," Signal-to-noise ratios typical of the short and long baselines are 8 and 4, respectively."
+ During this time. observations the single-dish flux was estimated via the full array. operating as a stand-alone instrument. to be 2.4—+0.25Jy.," During this time, observations the single-dish flux was estimated via the full array, operating as a stand-alone instrument, to be $2.4\pm0.25\,\Jy$."
+ This is similar to the visibility obtained on the baselines andamplitudes consistent with a single. compact gaussian component (?)..," This is similar to the visibility amplitudes obtained on the baselines and consistent with a single, compact gaussian component \citep{Doel_etal:08}."
+ This flux is anomalously low in comparison to the typical I.3mm flux of ~3Jy. and was taken as evidence for Ser A* appearing in a quiescent state.," This flux is anomalously low in comparison to the typical $1.3\,\mm$ flux of $\sim3\,\Jy$, and was taken as evidence for Sgr A* appearing in a quiescent state."
+ This interpretation is supported by the lack of a significant difference between analyses of each day separately (?).., This interpretation is supported by the lack of a significant difference between analyses of each day separately \citep{Brod_etal:09}.
+ Full details of the observations. calibration and. data processing can be found in ?..," Full details of the observations, calibration and data processing can be found in \citet{Doel_etal:08}."
+ 9? report upon more recent observations performed on the nights of April. 5-7. 2009. to the 95. 96. and 97 days of 2009.," \citet{Fish_etal:10} report upon more recent observations performed on the nights of April, 5–7, 2009, corresponding to the 95, 96, and 97 days of 2009."
+ These made correspondinguse of theJCMT..SMT.. and two dishes. operated as independent VLBI stations.," These made use of the, and two dishes, operated as independent VLBI stations."
+ 54 visibility amplitudes were obtained on and baselines on all days. and to both of the baselines on days 96 and 97.," 54 visibility amplitudes were obtained on and baselines on all days, and to both of the baselines on days 96 and 97."
+ Positions of the observations on each day are indicated in the upper panels of Figure 1.. labeled 2009.95. 2009.96. and 2009.97.," Positions of the observations on each day are indicated in the upper panels of Figure \ref{fig:Vobs}, labeled 2009.95, 2009.96, and 2009.97."
+ Signal-to-noise ratios typical of the short and long baselines are 17 and 5. respectively.," Signal-to-noise ratios typical of the short and long baselines are 17 and 5, respectively."
+ Thus. this second data set represents a significant improvement in both the number and precision of the data obtained.," Thus, this second data set represents a significant improvement in both the number and precision of the data obtained."
+" In addition to the VLBI baselines. the presence of two independent dishes in the array allowed the measurement of very-short baseline visibilities. probing angular scales ~10""."," In addition to the VLBI baselines, the presence of two independent dishes in the array allowed the measurement of very-short baseline visibilities, probing angular scales $\sim10''$."
+ These found substantially more correlated flux density than the baselines did. inconsistent with a single compact gaussian component.," These found substantially more correlated flux density than the baselines did, inconsistent with a single compact gaussian component."
+— The interpretation of the difference in correlated flux density between the baselines and the baselines is presently —unclear. and it may be possible for geometric models annular rings. extended double multiplesource) to fit the data.," The interpretation of the difference in correlated flux density between the baselines and the baselines is presently unclear, and it may be possible for multiple geometric models (e.g., annular rings, extended double source) to fit the data."
+ (e.g..Within the context of our analysis. we will assume that this difference is due to a separate large-scale component not present during the 2007 observations.," Within the context of our analysis, we will assume that this difference is due to a separate large-scale component not present during the 2007 observations."
+ This is indirectly supported by the fact that the source sizes inferred from the mid and long baseline data are unchanged despite the variations in the visibility magnitudes, This is indirectly supported by the fact that the source sizes inferred from the mid and long baseline data are unchanged despite the variations in the visibility magnitudes
+for Si).,for Si).
+ As noted. those temperatures agree well with the (ranges of) temperatures ΟΕΕ Ες of these IT- and We-like ious (this is also true for O. see below).," As noted, those temperatures agree well with the (ranges of) temperatures of maximum emissivities of these H- and He-like ions (this is also true for O, see below)."
+" The G value that we derive for indicates simular log(77) for the different elements: 6.95+0.25, 6,82-EO.2. and G.87250.5 for S. Si ancl Me. respectively."," The $\cal G$ value that we derive for indicates similar $\log(T)$ for the different elements: $\pm$ 0.25, $\pm$ 0.2 and $\pm$ 0.5 for S, Si and Mg, respectively."
+ While the G ratios for Me and Si are sinular to those of OD stars (@vhatever heir nature} reported bv2008).. the G ratio derived. for S can only ο Compared to those of the magnetic objects 7 SSco aud 01 CC. as well as to those of giant and main-sequence stars (only supereiants having ueher ratios).," While the $\cal G$ ratios for Mg and Si are similar to those of OB stars (whatever their nature) reported by, the $\cal G$ ratio derived for S can only be compared to those of the magnetic objects $\tau$ Sco and $\theta^1$ C, as well as to those of giant and main-sequence stars (only supergiants having higher ratios)."
+" If may be noted that the above eniperatfures corresponds to ~0.7keV. which agrees well with the ""cool temperature of the ""hot"" elobal model (see next section)."," It may be noted that the above temperatures corresponds to $\sim$ 0.7keV, which agrees well with the “cool” temperature of the “hot” global model (see next section)."
+ At temperatures los(D) of 6.8.7.0. Ry (calculated using ATOMDD as above) is about 2 for S. 2.65 or Si aud 2.12 for Meg.," At temperatures $\log(T)$ of 6.8–7.0, ${\cal R}_0$ (calculated using ATOMDB as above) is about 2 for S, 2.65 for Si, and 2.42 for Mg."
+" It is well known that R-i-mR,ΕΣ”:ο where we can neglect he density depeudence for massive stars."," It is well known that ${\cal R} = \frac{f}{i}= {\cal R}_0\, \frac{1}{1 + \phi / \phi_c + n_e / n_c}$, where we can neglect the density dependence for massive stars."
+ The UV flux o depends on the stellar output but also he dilution factor., The UV flux $\phi$ depends on the stellar output but also the dilution factor.
+" We derive the UW flix for uusing the model (C£;, 2IOkIS aud log(g)2L0) iu he grid of O-star atmosphere models calculated with CMEGEN which is closest to the parameters derived from a dedicated atinosphere fitting of 6||uazüs.", We derive the UV flux for using the model $T_{eff}$ =40kK and $\log(g)$ =4.0) in the grid of O-star atmosphere models calculated with CMFGEN which is closest to the parameters derived from a dedicated atmosphere fitting of .
+ After averaging the flix iu⋅ the velocity interval ποσαthe rest waveleneths of the > transitions⋅⋅ (following. a simular⋅⋅ path as (2006).. see also Tables)). we thenà derivedWyre a adoati radiusm1: ©tq LOO.02 LR.2 frouy. the R ratios of Si. and formation«3.2 R« for Me (the ratio itself corresponding to R.):," After averaging the flux in the velocity interval nearthe rest wavelengths of the $\rightarrow$ transitions (following a similar path as , see also \ref{wav}) ), we then derived a formation radius of $\pm$ $R_*$ from the $\cal R$ ratios of Si, and $<$ $R_*$ for Mg (the ratio itself corresponding to $R_*$."
+ The lines are too noisy to provide a meamnefil sulphurcoustraiut on the formation radius., The sulphur lines are too noisy to provide a meaningful constraint on the formation radius.
+ Such rather close radii are similar to those ecuerally found for O-stars2009). inchiding 05 CC2005): therefore. it cannot be used to discriminate between various N-ray enudssion models.," Such rather close radii are similar to those generally found for O-stars, including $\theta^1$ C: therefore, it cannot be used to discriminate between various X-ray emission models."
+" It may be worth ποιο, though. that our values fit well the temperature vs formation radius relation found by2008)."," It may be worth noting, though, that our values fit well the temperature vs formation radius relation found by."
+. Following(2009).. we derive an abundance ratio Mle/Si. relative to the solar abundance of(1989).. of IZEO.11 using the U-like aud IHe-like resonance lines.," Following, we derive an abundance ratio Mg/Si, relative to the solar abundance of, of $\pm$ 0.11 using the H-like and He-like resonance lines."
+ Χοποια abundances in CNO elements ave uot cutirely surprising since ddisplavs both a nitrogen overabundance aud an enriched surrouudiug uebulathereim)., Non-solar abundances in CNO elements are not entirely surprising since displays both a nitrogen overabundance and an enriched surrounding nebula.
+. However. changes iu he Mg/Si ratio are not expected (see also nex section).," However, changes in the Mg/Si ratio are not expected (see also next section)."
+ Finally. we may try to use the ratio of IT-to-IIe ines for oxvecu to check whether our Lypotlesis of a cooler plasiia dis seusible.," Finally, we may try to use the ratio of H-to-He lines for oxygen to check whether our hypothesis of a cooler plasma is sensible."
+ The triplet is very roisv in the RCS data. but the flux in the r liue cau 6 coustrained to 2.L40.6\10 *ppleem 7 (observed) or 69⋅+E19«105 Ds (corrected. for absorption).," The triplet is very noisy in the RGS data, but the flux in the r line can be constrained to $\pm$ $\times10^{-5}$ $^{-2}$ $^{-1}$ (observed) or $69\pm19\times10^{-5}$ $^{-2}$ $^{-1}$ (corrected for absorption)."
+ This corresponds to a IL-to-ITe ratio of 0.672:0.20. viclding a temperature of Ιου).~6.35.," This corresponds to a H-to-He ratio of $\pm$ 0.20, yielding a temperature of $\log(T)\sim6.35$."
+ This is much lower than the eniperatures fouul above: it would be difficult /CN kFGepxoktugeo this measurement iu the contextof an isothermal asma. thereby confinnünug our wpothesis of a multicteniperature plasma.," This is much lower than the temperatures found above: it would be difficult to reproduce this measurement in the contextof an isothermal plasma, thereby confirming our hypothesis of a multi-temperature plasma."
+ This result is certainly far from perfect because of its high uncertainty. but it constitutes a lint owards the non-uuiforuni teniperature of the X- cluitting regions.," This result is certainly far from perfect because of its high uncertainty, but it constitutes a hint towards the non-uniform temperature of the X-ray emitting regions."
+ Clobal fits were simultaneously made onu the IIEC MEG and 0th order spectra. using. a binning. ensuring at least 20 counts per bin (see Sect.," Global fits were simultaneously made on the HEG, MEG and 0th order spectra, using a binning ensuring at least 20 counts per bin (see Sect."
+ 2)., 2).
+ Results of the fits are preseuted iu. Table 9.., Results of the fits are presented in Table \ref{tabglobal}.
+ Note again that the quoted leo errors are sometimes axvinetricali the value shown here always is the larecst value., Note again that the quoted $\sigma$ errors are sometimes asymetrical: the value shown here always is the largest value.
+ We. used two sets of modelsthat both assume thermal plasimaiu collisional ionization equilibria (CTE)., We used two sets of modelsthat both assume thermal plasmain collisional ionization equlibrium (CIE).
+ First. we used models with discrete temperature cononents or ap," First, we used models with discrete temperature components or in"
+ First. we used models with discrete temperature cononents or ape," First, we used models with discrete temperature components or in"
+ First. we used models with discrete temperature cononents or apec," First, we used models with discrete temperature components or in"
+ First. we used models with discrete temperature cononents or apec ," First, we used models with discrete temperature components or in"
+ First. we used models with discrete temperature cononents or apec M," First, we used models with discrete temperature components or in"
+"We have investigated the variability. of Sevtert 1 and 2 galaxies on short aud long timescales and ‘ound iudicatious for variability in three Sevtert 2.0 ealaxies on short timescales (NCC 1068. IRAS 0117-0710. NGC ὠδδ),","We have investigated the variability of Seyfert 1 and 2 galaxies on short and long timescales and found indications for variability in three Seyfert 2.0 galaxies on short timescales (NGC 1068, IRAS 0147-0740, NGC 4388)."
+ À possible explanation for this variability uieht be the presence of boreholes iu the absorbing nolecular torus around the central black hole region., A possible explanation for this variability might be the presence of boreholes in the absorbing molecular torus around the central black hole region.
+ Significant N-rav variability during the pointed and survey observations were detected for 58 percent of he Sevfert 1 ealaxics., Significant X-ray variability during the pointed and survey observations were detected for 58 percent of the Seyfert 1 galaxies.
+lower than the true ones.,lower than the true ones.
+ This accounts for the than-expected normalisation of the WL-derived mass-concentration relation evident in Fig. 2.., This accounts for the lower-than-expected normalisation of the WL-derived mass-concentration relation evident in Fig. \ref{mcplot}.
+ We aim to provide an explanation for these biases in Section 5 below., We aim to provide an explanation for these biases in Section \ref{sec:errorsources} below.
+" Note that we have chosen to use medians, rather than means to quantify bias."," Note that we have chosen to use medians, rather than means to quantify bias."
+" The rationale behind this is that, in a non-Gaussian distribution as is the case here (log-normal), the mean, unlike the median, depends on the scatter as overpredictions can be arbitrarily high, whereas values can clearly not be underpredicted by more than100%."," The rationale behind this is that, in a non-Gaussian distribution as is the case here (log-normal), the mean, unlike the median, depends on the scatter as overpredictions can be arbitrarily high, whereas values can clearly not be underpredicted by more than."
+. From the amount of scatter evident in our results (see Figs., From the amount of scatter evident in our results (see Figs.
+" 3 and 5)) one should not be surprised to find that the mean mass and concentration show, in general, a positive bias with respect to their median counterparts."," \ref{histogram} and \ref{fig:scatter}) ) one should not be surprised to find that the mean mass and concentration show, in general, a positive bias with respect to their median counterparts."
+" For completeness, we show the difference between mean and median bias in Fig."," For completeness, we show the difference between mean and median bias in Fig."
+ C1 in appendix C.., \ref{fig:mean} in appendix \ref{sec:mean}.
+ The bias and scatter in cluster masses derived from weak, The bias and scatter in cluster masses derived from weak
+(0.2-0.5 )) in the ONC and evolve these stars forward using theoretical stellar evolution models auc different angular mormentuui loss rates.,(0.2–0.5 ) in the ONC and evolve these stars forward using theoretical stellar evolution models and different angular momentum loss rates.
+ In. 822 we present the theoretical framework of these uodels aud how they were applied to to the observatioual data., In 2 we present the theoretical framework of these models and how they were applied to to the observational data.
+ Iu 833 we will apply these moclels o data from both the Pleiades and the ONC., In 3 we will apply these models to data from both the Pleiades and the ONC.
+ These calculations will coustrain the two parameters of the angular momentum loss mechanisms. the saturation threshold and the disk-lockiug lifetime. eadiug to a single preferred. model.," These calculations will constrain the two parameters of the angular momentum loss mechanisms, the saturation threshold and the disk-locking lifetime, leading to a single preferred model."
+ Iufrared excess data for stars in the ONC is also used as au additional initial condition to test this model., Infrared excess data for stars in the ONC is also used as an additional initial condition to test this model.
+ In 811 we discuss the implications of these results. inclucling the uniqueness of our model aud how it cau be applied to further observational data iu he pre-MS.," In 4 we discuss the implications of these results, including the uniqueness of our model and how it can be applied to further observational data in the pre-MS."
+ To construct inodels of low-mass stars. we used the Yale Rotating Evolution Code (YREC. Guenther et al.," To construct models of low-mass stars, we used the Yale Rotating Evolution Code (YREC, Guenther et al."
+ 1992)., 1992).
+ YREC is a Heuyey code which solves tle equations of stellar structure in one dimeusioun., YREC is a Henyey code which solves the equations of stellar structure in one dimension.
+ YREC uses the nuclear reaction rates of Gruzinov Balicall (1998) aud the equation ol state from Saumon. Chabrier van Horn (1995).," YREC uses the nuclear reaction rates of Gruzinov Bahcall (1998) and the equation of state from Saumon, Chabrier van Horn (1995)."
+ Our models have a metallicity of Z = 0.0176 aud a mixing length of a = L.815. calibrated such that a 1.0 inodel will reproduce the solar radius aud luminosity at the solar age.," Our models have a metallicity of $Z$ = 0.0176 and a mixing length of $\alpha$ = 1.845, calibrated such that a 1.0 model will reproduce the solar radius and luminosity at the solar age."
+ The input physics for these models is discussed iu Sills. Pinsonueault. 'Terudrup (2000).," The input physics for these models is discussed in Sills, Pinsonneault, Terndrup (2000)."
+ Although the theoreticalmodels for magnetic star-disk interaction discussed in 31 are sophisticated aud complex in their formulation. the treatment of disk-lockiug for our angular momentum moclels is quite simple. since the disk aud star corotate at a fixed angular velocity.," Although the theoretical models for magnetic star-disk interaction discussed in 1 are sophisticated and complex in their formulation, the treatment of disk-locking for our angular momentum models is quite simple, since the disk and star corotate at a fixed angular velocity."
+ When the star is its period is held coustant over the lifetime of the clisk. 77:4. aud the angular momentuui Change is then a function of stars moment of inertia.," When the star is disk-locked, its period is held constant over the lifetime of the disk, $\tau_{disk}$, and the angular momentum change is then a function of star's moment of inertia."
+ When the age of the stellar model reaches τικ. (he star ds released. [rou disk-Iockiug aud evolves uuder equatious (1) aud (2) below.," When the age of the stellar model reaches $\tau_{disk}$, the star is released from disk-locking and evolves under equations (1) and (2) below."
+ The disk lifetime. 75:44. is relative to the birthliue (0 Myr) aid not to the observed age of the star. τε.," The disk lifetime, $\tau_{disk}$, is relative to the birthline (0 Myr) and not to the observed age of the star, $\tau_\star$."
+ We use the birthline of Palla Staller (1991). which is the cdeuterium-burniug main sequence and corresponds to the upper envelope of T Tauri stars in the H-R diagram.," We use the birthline of Palla Staller (1991), which is the deuterium-burning main sequence and corresponds to the upper envelope of T Tauri stars in the H-R diagram."
+ The time evolution of the stellar mome of inertia was takeu from the YREC stellar inodels. in which solic-bocdy rotation was enforced.," The time evolution of the stellar moment of inertia was taken from the YREC stellar models, in which solid-body rotation was enforced."
+ Stellar mocels of this mass rauge that included iuterual angular momentuu trausport were nearly kleutical to solid body models (Sills. Pinsouneault. Terndrup 2000).," Stellar models of this mass range that included internal angular momentum transport were nearly identical to solid body models (Sills, Pinsonneault, Terndrup 2000)."
+ Rotating stars that lose mass through maguetized stellar winds will also lose angular mouentum., Rotating stars that lose mass through magnetized stellar winds will also lose angular momentum.
+ To quantify this loss rate. we used a prescription for angular momentui loss adopted from Ixawaler (1988) and MacGregor Brennan (1991). aud described in Wrishuamurthi et al. (," To quantify this loss rate, we used a prescription for angular momentum loss adopted from Kawaler (1988) and MacGregor Brennan (1991), and described in Krishnamurthi et al. ("
+1997).,1997).
+ We write, We write
+"from equation (A4)) and taking only terms linear in à, and (2.",from equation \ref{eqn:quad}) ) and taking only terms linear in $\delta z_c$ and $\hat Q$.
+ Here. we note that OQ~©(d2) so that the second term in the bracket is lower order than (27.," Here, we note that $\hat Q\sim\mathcal{O}(d_c^2)$ so that the second term in the bracket is lower order than $\hat Q^2$."
+" In order to find the magnification for this image. we differentiate equation (A4)). and find. substituting : =:y|ὅτε, for which 8: is given by equation (A5)). into equation (A6)). we obtain. alt Note that the total magnification for the given source position is usually dominated by one or two images found close to the critical curve."," In order to find the magnification for this image, we differentiate equation \ref{eqn:quad}) ), and find, Then substituting $z=z_0+\delta z_c$ , for which $\delta z_c$ is given by equation \ref{eqn:delc}) ), into equation \ref{eqn:part}) ), we obtain, Note that the total magnification for the given source position is usually dominated by one or two images found close to the critical curve."
+ Thus. we consider only the case for which the non-perturbed PSPL images lie close to the unit circle. so we have leu]=1|A and A<1.," Thus, we consider only the case for which the non-perturbed PSPL images lie close to the unit circle, so we have $|z_0|=1+\Delta$ and $\Delta\ll 1$."
+" Then. we find the expression for the inverse magnification for the QL approximation (up to the order of d? ). extreme wide-binary case. one can rewrite the lens equation (À 1)) as1 ""m position and the mass of the first lens component are the origin and the unit mass so that £4 =O. +2= dy.ey= 1. €)=qu. and dq=(L|qu):/74,.."," Then, we find the expression for the inverse magnification for the QL approximation (up to the order of $d_c^3$ ), For the extreme wide-binary case, one can rewrite the lens equation \ref{eqn:lenseq}) ) as Here the position and the mass of the first lens component are the origin and the unit mass so that $z_1=0$, $z_2=-d_1$, $\epsilon_1=1$, $\epsilon_2=q_w$, and $d_1=(1+q_w)^{1/2}d_w$."
+ We note that. apart from the constant translation. the first non-PSPL term here is essentially the shear. >=qud472quedA⋅↽|du)! for the CRL4 approximation.," We note that, apart from the constant translation, the first non-PSPL term here is essentially the shear, $\gamma=q_w d_1^{-2}=q_w d_w^{-2}(1+q_w)^{-1}$ for the CRL approximation."
+". Analogous to the extreme close binary.. ifp ~<J1. the image position for the CRL approximation can be found by the perturbative approach. but here the corresponding PSPL source position would be&|qud,=τυly! "," Analogous to the extreme close binary, if $\gamma\ll 1$, the image position for the CRL approximation can be found by the perturbative approach, but here the corresponding PSPL source position would be $\zeta+q_w d_w^{-1}=z_0-\bar z_0^{-1}$."
+"Then. the image deviation ὅτι—2.:9 of CRL from PSPL is Using this result and the derivative of equation (A 10)). we find From the same argument used for the QL approximation of the extreme close binary. wecan set |:y|=1|A. and then the inverse magnification for the CRL approximation (up to the order of d.) is By comparing equations (A8)) and (A14)). we therefore establish the magnification correspondence (up to the order of d dz) between the close binary with Q=d24,(1|q.)? and the wide binarywith >=πμle)| when Q~-& 1."," Then, the image deviation $\delta z_w=z-z_0$ of CRL from PSPL is Using this result and the derivative of equation \ref{eqn:crl}) ), we find From the same argument used for the QL approximation of the extreme close binary, wecan set $|z_0|=1+\Delta$, and then the inverse magnification for the CRL approximation (up to the order of $d_w^{-3}$ ) is By comparing equations \ref{eqn:magc}) ) and \ref{eqn:magw}) ), we therefore establish the magnification correspondence (up to the order of $d_c^2\sim d_w^{-2}$ ) between the close binary with $\hat Q=d_c^2 q_c (1+q_c)^{-2}$ and the wide binarywith $\gamma=d_w^{-2} q_w (1+q_w)^{-1}$ when $\hat Q\simeq\gamma\ll 1$ ."
+and absence of € IIT absorption therefore sugeecsts a ~BO51I classification with au uncertainty of roughly half a spectral subtype aud lack of strong constraints ou the luminosity class.,and absence of C III absorption therefore suggests a $\sim$ B0.5–1I classification with an uncertainty of roughly half a spectral subtype and lack of strong constraints on the luminosity class.
+ However. we caution that if C TIT is weak due to abundance anomalies or uear-critical rotation then this limit may not apply. and the weakness of the Ue I lines permit a classification of O9.5DOI that is broadly consistent with the Pascheu-series line streustlis.," However, we caution that if C III is weak due to abundance anomalies or near-critical rotation then this limit may not apply, and the weakness of the He I lines permit a classification of O9.5–B0I that is broadly consistent with the Paschen-series line strengths."
+ Figue 3 shows RV curves for the two conmpoucuts of the system., Figure \ref{fig:rv} shows RV curves for the two components of the system.
+ Taking the 9.20-dav period reported by Donanos(2007)— as a starting point. an crror-weighted. 472 fitgp to the radia. velocitics of the absorptiou-lne componcut vieldecd bes-fit values for the orbital period of days. cousistent wih an iude)ondenut deteriuuation using a Lomb-Scarele periodogran (Press&Reybicki1989).. a xvstende velocity of and πο]auplitude |.," Taking the 9.20-day period reported by \cite{bonanos} as a starting point, an error-weighted $\chi^2$ fit to the radial velocities of the absorption-line component yielded best-fit values for the orbital period of days, consistent with an independent determination using a Lomb-Scargle periodogram \citep{press}, a systemic velocity of and semi-amplitude ."
+ The corresponding fit to the cussion line RV has a systemic velocity ando semuieuuplitude, The corresponding fit to the emission line RV has a systemic velocity of and semi-amplitude.
+ Exrors are derived οι he fitting residuals using the bootstrap method (Efron&Tibshirawi1991)., Errors are derived from the fitting residuals using the bootstrap method \citep{efron}.
+. We note tlat systemic velocity ¢erved from the emission line fit js somewhat ower than that derived from the companion. and is in closer aerecment with the mean radial velocities of other W«] supereiautsOo (see Paper D.," We note that systemic velocity derived from the emission line fit is somewhat lower than that derived from the companion, and is in closer agreement with the mean radial velocities of other Wd1 supergiants (see Paper I)."
+ Discrepaucies iu lis parameter are commonly observed in early-typespectroscopic binaries (e.g. the 98-day ΟΠ(Ώ|Os.51 iuarv1191014: Ramctal. 2001)) although the effect is simall iun comparison with some οher evolved Svstenas (cec.228766: Massey&€'onti1977:Dauwetal. 20023) in which wind contaminaion stronely affects derived systemic velocities.," Discrepancies in this parameter are commonly observed in early-typespectroscopic binaries (e.g. the 9.8-day O7III(f)+O8.5I binary; \citealt{rauw01}) ) although the effect is small in comparison with some other evolved systems (e.g.; \citealt{massey,rauw02}) ) in which wind contamination strongly affects derived systemic velocities."
+ Taking these values vields a nass ratio aud masses or the two colmponents of: and Finally. micasuremeuts of blended lvdrogen lires with Gaussian fits tend to vield svstematically lower values of A4 and A» than methods such as spectral «liseutauelimg (Sinon&Sturm1991) that are less affected by lending (Andersen1975:ποιον&Clausen2007).," Taking these values yields a mass ratio and masses for the two components of: and Finally, measurements of blended hydrogen lines with Gaussian fits tend to yield systematically lower values of $K_1$ and $K_2$ than methods such as spectral disentangling \citep{simon} that are less affected by blending \citep{andersen,southworth}."
+. The paucity of stroug lines free from siguificaut interstellar. tellure and wind contamination iu fjo ᾖ- alc I-baud spectra of ~BO supergiauts makes the extent oftjs effect on our determination of A4 aud Wap. hard to quautity. and we therefore note that our method nav uncderestinate the masses of the two components ofW13.," The paucity of strong lines free from significant interstellar, telluric and wind contamination in the $R$ - and $I$ -band spectra of $\sim$ B0 supergiants makes the extent of this effect on our determination of $K_\text{em}$ and $K_\text{abs}$ hard to quantify, and we therefore note that our method may underestimate the masses of the two components of."
+". To coustrai Εκ, we folded the R-baud photonetric data reported by Boanos(2007) on to the 9.27l dav period determined frou the RV data."," To constrain $\text{sin}^3i$, we folded the $R$ -band photometric data reported by \cite{bonanos} on to the 9.271 day period determined from the RV data."
+ The data were binned ο reduce he considerable scatter prescut in the light curve. which is most probably a consequence of iutriusic aperiodic variaΗΠΑ iu one or both compoιομι low-evel photometric aud spectroscopic variabilit vis a feature of all tranusitiojii supereiauts in Well. witi the blue ivpergiants disλανι rapid phnotoimetrie variability. at he ~0.1 maguitude level and the earhk-B supereiauts also variable (Bonanos2007:Clarkctal. 2010a).," The data were binned to reduce the considerable scatter present in the light curve, which is most probably a consequence of intrinsic aperiodic variability in one or both components: low-level photometric and spectroscopic variability is a feature of all transitional supergiants in Wd1, with the blue hypergiants displaying rapid photometric variability at the $\sim$ 0.1 magnitude level and the early-B supergiants also variable \citep{bonanos, clark10}. ."
+. Therefore. eivoen the linited dataset and shallow ~0.15r machiΡοude," Therefore, given the limited dataset and shallow $\sim$ 0.15 magnitude"
+For marginalisation we instead use a method that is based on subtracting the magnitude of an arbitrarily chosen low-redshift (so its magnitude depends only on the IIubble constant) andanchor SN from the magnitudes of SNe in a data set then mareinalisine over magnitude of anchor SN.,For marginalisation we instead use a method that is based on subtracting the magnitude of an arbitrarily chosen low-redshift (so its magnitude depends only on the Hubble constant) anchor SN from the magnitudes of SNe in a data set and then marginalising over magnitude of anchor SN.
+ The resulting likelihood. function is derived in the Appendix., The resulting likelihood function is derived in the Appendix.
+ Εις is equivalent to marginalising over the nuisance parameter {fy with a Gaussian prior centered around its value derived from the anchor SN. alongwith the corresponding standard deviation.," This is equivalent to marginalising over the nuisance parameter $H_0$ with a Gaussian prior centered around its value derived from the anchor SN, alongwith the corresponding standard deviation."
+ The method can be easily eeneralized. to à case where the priors on Lfy are specified separately as described in the Appenclix., The method can be easily generalized to a case where the priors on $H_0$ are specified separately as described in the Appendix.
+ We now compare the analytically marginalised. likelihood function in Eq X5 to the results of Bavesian marginalisation over fy., We now compare the analytically marginalised likelihood function in Eq \ref{eq:marglik} to the results of Bayesian marginalisation over $H_0$.
+ For this purpose we consider parameter estimation for GDOA., For this purpose we consider parameter estimation for GD04.
+ We assume a Hat ACDAL universe for this exercise., We assume a flat $\Lambda$ CDM universe for this exercise.
+ The dimensionless Hubble parameter is given by and [latness implies O4=1Qi., The dimensionless Hubble parameter is given by and flatness implies $\Ol = 1 - \Om$.
+ Thus. the only free parameters are Z/ü and Oxy.," Thus, the only free parameters are $H_0$ and $\Om$."
+ The normalized likelihood function is given by llere à;=quofhbeeav(oziHo.O31). where the subscript Pompe.V and “ij=jc? is the covariance matrix.," The normalized likelihood function is given by Here $x_i = \mu_i - \mu_{\rm theory} (z_i; H_0,\Om)$, where the subscript $i=1,\ldots,N$ and $\Sigma_{ij} = \delta_{ij}\,\sigma_i^2$ is the covariance matrix."
+ The superscript T denotes the matrix operation of taking the transpose of a matrix., The superscript 'T' denotes the matrix operation of taking the transpose of a matrix.
+" The posterior probability for parameters Ox; ancl Lfy is given by We choose a uniform prior for Oy; in the range 0x Llane for ff,=1005kms!Mpe in the range 04x:1."," The posterior probability for parameters $\Om$ and $H_0$ is given by We choose a uniform prior for $\Om$ in the range $0 \le \Om \le 1$ and for $H_0 = 100h\,\rm km \,s^{-1}\,Mpc^{-1}$ in the range $0.4 \le h \le 1$."
+ The mareinalised probability clistribution for the matter density is given by The probability density is normalized after carrving out the integration., The marginalised probability distribution for the matter density is given by The probability density is normalized after carrying out the integration.
+ A similar probability density function for the matter density can be obtained by the Bayesian inversion of the marginalised likelihood. function. given in I., A similar probability density function for the matter density can be obtained by the Bayesian inversion of the marginalised likelihood function given in Eq.
+ AD., A5.
+ ligure 1. plots a comparison between the two probability densities ancl shows that the two clistributions are nearly identical., Figure \ref{fig:comparison} plots a comparison between the two probability densities and shows that the two distributions are nearly identical.
+ For completeness some details are repeated. here. from GSLOS., For completeness some details are repeated here from GSL08.
+ For our analysis we have considered a [at .ACDM universe. which can be easily generalized to à more general model of dark energy.," For our analysis we have considered a flat $\Lambda$ CDM universe, which can be easily generalized to a more general model of dark energy."
+redwards and to higher luminosities away [rom the zero-age main sequence (ZAAIS). due to the increasing helium content of their cores.,"redwards and to higher luminosities away from the zero-age main sequence (ZAMS), due to the increasing helium content of their cores."
+ This movement continues until the point of core hydrogen exhaustion. when the star has reached the reached the terminal age MS CEXMS or MIS turn-oll).," This movement continues until the point of core hydrogen exhaustion, when the star has reached the reached the terminal age MS (TAMS or MS turn-off)."
+ Finally. after the turn-olf. the post-main-sequence evolution is driven by the burning of heavier elements which leads to much more rapid movement in the CMD.," Finally, after the turn-off, the post-main-sequence evolution is driven by the burning of heavier elements which leads to much more rapid movement in the CMD."
+ Vhis relatively high velocity in the CMD means that post-main-sequence evolution has the potential to give. precise ages., This relatively high velocity in the CMD means that post-main-sequence evolution has the potential to give precise ages.
+ Llowever. for voung galactic clusters the paucity of stars in this region of the CAID means such an age can depend on just one star. and such ages are rightly treated with some scepticism.," However, for young galactic clusters the paucity of stars in this region of the CMD means such an age can depend on just one star, and such ages are rightly treated with some scepticism."
+ Conversely. the main-sequence evolution (from the ZAMS to the turn-oll) has a larger number of stars. but the movement is often subtle. ancl using the normal technique of simv plotting isochrones over the data leads to large uncertainties in age. and to questions over objectivity.," Conversely, the main-sequence evolution (from the ZAMS to the turn-off) has a larger number of stars, but the movement is often subtle, and using the normal technique of simply plotting isochrones over the data leads to large uncertainties in age, and to questions over objectivity."
+ However. we have been developing a method of making objective fits to colour-magnituce data. which should allow us to unlock the information in this stage ofa star's evolution.," However, we have been developing a method of making objective fits to colour-magnitude data, which should allow us to unlock the information in this stage of a star's evolution."
+ The technique. called 77 fitting. can be viewed as an extension of X7 to data points with uncertainties in two or more observables. and to models which are distributions (not just lines) in the data space.," The technique, called $\tau^2$ fitting, can be viewed as an extension of $\chi^2$ to data points with uncertainties in two or more observables, and to models which are distributions (not just lines) in the data space."
+ The ain of this paper is to apply the 7? fitting technique o the main-sequence evolution of voung stars. and use he resulting ages to create à revised age scale for. PAIS stars.," The aim of this paper is to apply the $\tau^2$ fitting technique to the main-sequence evolution of young stars, and use the resulting ages to create a revised age scale for PMS stars."
+ Surprisingly. this leads to a significantly older ages han the commonly. used. contraction ages. a result. which we will discuss in Section 11..," Surprisingly, this leads to a significantly older ages than the commonly used contraction ages, a result which we will discuss in Section \ref{discuss}."
+ To derive this result. we irst have to update our statistical techniques. originally. described. in ?.. since. as we discuss in Section 4. the echnique will not. work for the isochrones we wish to it.," To derive this result we first have to update our statistical techniques originally described in \cite{2006MNRAS.373.1251N}, , since, as we discuss in Section \ref{stats}, the technique will not work for the isochrones we wish to fit."
+ We therefore lav out the changes which need to be made by following an example through fitting (Section 5)). esting the goodness of fit (Section 6)) and determining the uncertainties in the derived. parameters (Section 7)).," We therefore lay out the changes which need to be made by following an example through fitting (Section \ref{fit}) ), testing the goodness of fit (Section \ref{goodness}) ) and determining the uncertainties in the derived parameters (Section \ref{uncer}) )."
+ Before doing so. however. we discuss the data ane models we use (Sections 2. ancl Sections 3)).," Before doing so, however, we discuss the data and models we use (Sections \ref{data} and Sections \ref{models}) )."
+ We deal with the ellects of interstellar extinction in Section S.. and the details of each cluster in Section 9..," We deal with the effects of interstellar extinction in Section \ref{extin}, and the details of each cluster in Section \ref{individual}."
+ We draw all the results together in our discussion in Section 11.., We draw all the results together in our discussion in Section \ref{discuss}.
+ To compare a set of ages derived. from MS. evolution with contraction ages we need a sample of clusters aud associations which have contraction ages. ancl for cach of which data are available for AIS fitting.," To compare a set of ages derived from MS evolution with contraction ages we need a sample of clusters and associations which have contraction ages, and for each of which data are available for MS fitting."
+ Our sample. is. therefore. based on the groups we placed. in age order using the PATS in ?..," Our sample is, therefore, based on the groups we placed in age order using the PMS in \cite{2008MNRAS.386..261M}."
+ Clearly. for cach of these groups we require stars in the appropriate mass range to show significant ALS evolution. but we also require. extinetions and reliable distance measurements.," Clearly, for each of these groups we require stars in the appropriate mass range to show significant MS evolution, but we also require extinctions and reliable distance measurements."
+ CY. photometry can provide all three of these., $UBV$ photometry can provide all three of these.
+ First the €D/ D.V. diagram provides extinctions., First the $U-B$ $B-V$ diagram provides extinctions.
+ Second. the upper part of VsV. diagram is age sensitive. tracing the evolution of stars from the ZAMS to the turn-oll.," Second, the upper part of $V$ $B-V$ diagram is age sensitive, tracing the evolution of stars from the ZAMS to the turn-off."
+. Finally. in the age range of interest the lower mass stars are still close to the ZAAIS. and the sequence turns redwards. making it ideal as a distance measure.," Finally, in the age range of interest the lower mass stars are still close to the ZAMS, and the sequence turns redwards, making it ideal as a distance measure."
+ Furthermore. the CBV. photo-electric system is very consistent and. well characterised.," Furthermore, the $UBV$ photo-electric system is very consistent and well characterised."
+ However. to ensure we maintain the highest level of consistency we have restricted. ourselves as far as possible to the data of Johnson and collaborators. primarily taken in the 1950s and 1960s.," However, to ensure we maintain the highest level of consistency we have restricted ourselves as far as possible to the data of Johnson and collaborators, primarily taken in the 1950s and 1960s."
+ As we shall show later. the quality of these data when combined with the transformations of 7. is impressive. eiving 77 values which mean the model is a good fit to the data.," As we shall show later, the quality of these data when combined with the transformations of \cite{1998A&A...333..231B} is impressive, giving $\tau^2$ values which mean the model is a good fit to the data."
+ Clearly we wish to avoid PAIS stars contaminating our sample at faint maenituces and red colours. and so for most objects we apply a cut in observed 2BV which roughly corresponds to (23τους0.0.," Clearly we wish to avoid PMS stars contaminating our sample at faint magnitudes and red colours, and so for most objects we apply a cut in observed $B-V$ which roughly corresponds to $(B-V)_0 < 0.0$."
+ Alost of the datasets we use have robust. uncertainties derived from. comparisons of many measurements of stars., Most of the datasets we use have robust uncertainties derived from comparisons of many measurements of stars.
+ ‘This presents us with a problem. as the quoted uncertainties in colour are always smaller than those in maenitucle.," This presents us with a problem, as the quoted uncertainties in colour are always smaller than those in magnitude."
+ Conventional error analysis vieles a correlation between. sav. V and £2V. and in previous work we have always xen careful to include that correlation when modeling the uncertainties.," Conventional error analysis yields a correlation between, say, $V$ and $B-V$, and in previous work we have always been careful to include that correlation when modeling the uncertainties."
+ The starting point for such an analvsis is hat V and. D are measured. independently. ancl so. the uncertainties in | and D.V are 81 and dl?|8/7 respectively.," The starting point for such an analysis is that $V$ and $B$ are measured independently, and so the uncertainties in $V$ and $B-V$ are $\delta V$ and $\sqrt{\delta V^2+\delta B^2}$ respectively."
+ Such an analysis also leads to the conclusion hat the uncertainty in D.V must be larger than that in V. in direct. contradiction to the quoted uncertainties for most of the data presented here.," Such an analysis also leads to the conclusion that the uncertainty in $B-V$ must be larger than that in $V$, in direct contradiction to the quoted uncertainties for most of the data presented here."
+ “Phis is because it is not shoton statistics which are the driver of the uncertainties. rut changes in the transpareney.," This is because it is not photon statistics which are the driver of the uncertainties, but changes in the transparency."
+ In this work. we therefore model the uncertainties as uncorrelatect.," In this work, we therefore model the uncertainties as uncorrelated."
+ Although we will tery other models later. we begin by using “CGeneva-Bessell” isochrones.," Although we will try other models later, we begin by using “Geneva-Bessell” isochrones."
+" For the stellar interior we follow the suggestion of 2.. and use the ""basic mocel set” (Le. set πο} of the Geneva. isochrones (?).."," For the stellar interior we follow the suggestion of \cite{2001A&A...366..538L}, and use the “basic model set” (i.e. set “c”) of the Geneva isochrones \citep{1992A&AS...96..269S}."
+ Temporal interpolation is a much more significant issue for post-MS isochrones than the PALS isochrones we have fitted in the past. as there are sharp cliscontinuities in the rate of change of magnitude and colour with time. as exemplified by the MS turn-olf.," Temporal interpolation is a much more significant issue for post-MS isochrones than the PMS isochrones we have fitted in the past, as there are sharp discontinuities in the rate of change of magnitude and colour with time, as exemplified by the MS turn-off."
+. We therefore use the code provided on the website to interpolate the isochrones to the appropriate age., We therefore use the code provided on the website to interpolate the isochrones to the appropriate age.
+ We then convert from Iuminosity and cllective temperature to colours and magnitudes using the tables of ?.. assuming the colours of Vega are zero (though V= 0.03).," We then convert from luminosity and effective temperature to colours and magnitudes using the tables of \cite{1998A&A...333..231B}, assuming the colours of Vega are zero (though $V=0.03$ )."
+ We also use Bessell et al's colour dependent extinction vectors., We also use Bessell et al's colour dependent extinction vectors.
+ For some of the most luminous stars the gravities are rather low. and fall just outside the range of gravities given by 7..," For some of the most luminous stars the gravities are rather low, and fall just outside the range of gravities given by \cite{1998A&A...333..231B}."
+ In these cases we extrapolate the models by. simply setting the colour to that for the lowest available gravity., In these cases we extrapolate the models by simply setting the colour to that for the lowest available gravity.
+ In these cases a linear extrapolation would be cilferent by less than 0.001 mags. implving that the overall error due to the extrapolation is much smaller than the uncertainties in colour.," In these cases a linear extrapolation would be different by less than 0.001 mags, implying that the overall error due to the extrapolation is much smaller than the uncertainties in colour."
+ For reasons explained in Section 9.2 we used the Tveho-2 photometry for σ Ori., For reasons explained in Section \ref{sori} we used the Tycho-2 photometry for $\sigma$ Ori.
+ In this case we have used the conversion givenin 7. to convert the Geneva-Dessell isochrones into the Tvcho svstem.," In this case we have used the conversion givenin \cite{2000PASP..112..961B}
+ to convert the Geneva-Bessell isochrones into the Tycho system."
+ (2.statethattheTvcho- We used the reddening vector derived in ?..," \citep[][state that the Tycho-1 and Tycho-2 systems should be identical.]
+{2000A&A...357..367H}
+ We used the reddening vector derived in \cite{2008MNRAS.386..261M}. ."
+2,.
+006).. has a relatively low abundance in (Glasseold1996).. and is presumably below the current detection limit even if it might be present in6.," $^+$ has a relatively low abundance in \citep{glassgold96}, and is presumably below the current detection limit even if it might be present in."
+. We observed abundant carbon chains and radicals in6.. including CO. SiCs. CN. HON. CS. Coll. Can. Cyl. HICSN. and CII4CN. all of which are linear.," We observed abundant carbon chains and radicals in, including CO, $_2$, CN, HCN, CS, $_2$ H, $_3$ N, $_4$ H, $_3$ N, and $_3$ CN, all of which are linear."
+ This characteristic feature is similar to those of ancl (seeCernicharoetal.2000). although (hese lines are much fainter in6.," This characteristic feature is similar to those of and \citep[see][]{cernicharo00}
+ although these lines are much fainter in."
+. The most intriguing characteristic of is (he strong CN emission., The most intriguing characteristic of is the strong CN emission.
+ The integrated intensity ratio of the CN (2.1) eroup aud the CO (21) transition is 4.6. a factor of 2.2 larger than the value in 2008)..," The integrated intensity ratio of the CN (2–1) group and the $^{13}$ CO (2–1) transition is 4.6, a factor of 2.2 larger than the value in \citep{he08}. ."
+ CN is mainly formed (τος the photocissociation of IICN. According to Heetal.(2008).. the IIPCN 2)/PCO ὢ1) integrated intensity ratio in is 4.5. a [actor of 3.2 larger than that in6.," CN is mainly formed through the photodissociation of HCN, According to \citet{he08}, the $^{13}$ CN $^{13}$ CO (2–1) integrated intensity ratio in is 4.5, a factor of 3.2 larger than that in."
+.. Therefore. our observations provide strong evidence (hat reaction (3)) dominates the chemistry of CN and ICN in AGB stars and the photodissociation is more efficient in the more evolved C-rich envelopeG.," Therefore, our observations provide strong evidence that reaction \ref{hcn}) ) dominates the chemistry of CN and HCN in AGB stars and the photodissociation is more efficient in the more evolved C-rich envelope."
+. The above discussion also suggests that about 30% CN formed from IICN has been destroved., The above discussion also suggests that about $\%$ CN formed from HCN has been destroyed.
+ On the other hand. CN can be reprocessed into ΕΠ through the reaction We do find that the ILC4N line intensities relative to the @CO (21) transition in ave a factor of ~3 larger than those inIRC+10216.. indicating efficient formation of 1IC4N in6.," On the other hand, CN can be reprocessed into $_3$ N through the reaction We do find that the $_3$ N line intensities relative to the $^{13}$ CO (2–1) transition in are a factor of $\sim3$ larger than those in, indicating efficient formation of $_3$ N in."
+. We did not find evidence for the enhancement of (he C4N radical. suggesting that photoclissociation of 11CN into C4N is insignificant in this object.," We did not find evidence for the enhancement of the $_3$ N radical, suggesting that photodissociation of $_3$ N into $_3$ N is insignificant in this object."
+ shows strong Cll emission., shows strong $_2$ H emission.
+ The CSI radical is dominantly produced through the photodissociation reaction Our observations show that the Coll line intensities relative to the CO (1)transition in, The $_2$ H radical is dominantly produced through the photodissociation reaction Our observations show that the $_2$ H line intensities relative to the $^{13}$ CO (2–1)transition in
+that in the B band light curve the Dux varies from under 9 mJy to over 24 mJy. so the majority of the tux measured must be nuclear (assuming that the star light is constant).,"that in the B band light curve the flux varies from under 9 mJy to over 24 mJy, so the majority of the flux measured must be nuclear (assuming that the star light is constant)."
+ In the V band the maximum to minimum variation is slightlv smaller but still more than a factor of 2., In the V band the maximum to minimum variation is slightly smaller but still more than a factor of 2.
+ This dillerence in variability amplitude can be caused by stronger star light contamination in the V band but it is also possible that the AGN V emission is intrinsically less variable., This difference in variability amplitude can be caused by stronger star light contamination in the V band but it is also possible that the AGN V emission is intrinsically less variable.
+ The power density spectrum (PDS) can be used. to quantify the variability amplitude as a function of the time-scale of the variations. or correspondingly. of their Fourier requency.," The power density spectrum (PDS) can be used to quantify the variability amplitude as a function of the time-scale of the variations, or correspondingly, of their Fourier frequency."
+ The PDS is constructed. through the modulus squared. of the discrete. Fourier. Transform. (DET). (Pressetal. 1992)., The PDS is constructed through the modulus squared of the discrete Fourier Transform (DFT) \citep{Press}.
+. For the normalisation used in our calculations. he integral of the PDS over frequency equals the normalised variance of the light curve.," For the normalisation used in our calculations, the integral of the PDS over frequency equals the normalised variance of the light curve."
+ For most AGN X-ray light curves. the PDS has a power aw shape of slope 1 bending to a steeper slope at. high requencies (c.g. Summons— et iin prep... Mellardyctal.2004. 2005)).," For most AGN X-ray light curves, the PDS has a power law shape of slope $\sim -1$ bending to a steeper slope at high frequencies (e.g. Summons et in prep., \citealt{McHardy4051,McHardyMCG}) ),"
+ which is similar to the PDS found in stellar mass black hole binaries in the soft sta| (seo[ον2007.ora review)., which is similar to the PDS found in stellar mass black hole binaries in the soft state \citep[see][for a review]{uttleyreview}.
+ Lt is customary to multiply the variability »ower by frequency when plotting the PDS. to highlight deviations of the power law slope from 1 as. in this case. he low frequency. part of the PDS appears approximately lat and the breaks are more noticeable.," It is customary to multiply the variability power by frequency when plotting the PDS, to highlight deviations of the power law slope from –1 as, in this case, the low frequency part of the PDS appears approximately flat and the breaks are more noticeable."
+ We use this παπατα of presentation in Fig. 4.., We use this standard of presentation in Fig. \ref{pds}.
+" ln Summonsetal.(2007) we show the X-ray PDS of aancl explore the significance of an apparent cuasi-periodic oscillation (QPO) at a frequency of ~510""Lz.", In \citet{summons} we show the X-ray PDS of and explore the significance of an apparent quasi-periodic oscillation (QPO) at a frequency of $\sim5 \times 10 ^{-6}$ Hz.
+ The new. intensively sampled light curve obtained for this object. shown in Fig. 3..," The new, intensively sampled light curve obtained for this object, shown in Fig. \ref{intensive},"
+ with a sampling rate of three times claily and a length. of four months. covers the frequency. range 1034. llz.," with a sampling rate of three times daily and a length of four months, covers the frequency range $10^{-7}-3.4\times 10^{-5}$ Hz."
+ This range covers the time-scales corresponding to the peak frequency. of the possible OPO very well and allows us to test its significance conclusively., This range covers the time-scales corresponding to the peak frequency of the possible QPO very well and allows us to test its significance conclusively.
+ We caleulated the PDS using the long term Hlight curves and short term Iligght. curves discussed in Summonsetal.(2007) and added the new intensive cata., We calculated the PDS using the long term light curves and short term light curves discussed in \citet{summons} and added the new intensive data.
+ Phe resulting PDS is shown in solid lines in bie. 4..," The resulting PDS is shown in solid lines in Fig. \ref{pds},"
+ where the segments correspond to the cillerent X-ray light curves used., where the segments correspond to the different X-ray light curves used.
+ We fitted a bending power law model defined as to the PDS using the Monte. Carlo fitting technique of Uttleyetal.(2002)., We fitted a bending power law model defined as to the PDS using the Monte Carlo fitting technique of \citet{psresp}.
+". Phe low-frequency slope op. the high frequency slope ag. the bend. frequeney fi, and the normalisation ;À were allowed to vary."," The low-frequency slope $ \alpha_L$, the high frequency slope $\alpha_H$, the bend frequency $f_b$ and the normalisation $A$ were allowed to vary."
+" The fitting parameters are ap=OS. ag=2.2. fy10"" IIz. consistent with the values obtained by Summonsetal.(2007). for the same bending power law model."," The best-fitting parameters are $\alpha_L=0.8$, $\alpha_H=2.2$, $f_b=5.8\times 10^{-6}$ Hz, consistent with the values obtained by \citet{summons} for the same bending power law model."
+ The corresponding model is shown by the dashed line in Fig. 4.., The corresponding model is shown by the dashed line in Fig. \ref{pds}.
+ The simple bending power law provided an excellent.fit to the new data acceptance probability). making the possible QPO feature unnecessary.," The simple bending power law provided an excellentfit to the new data acceptance probability), making the possible QPO feature unnecessary."
+ We also computed the PDS of the B band data. shown by the dotted line in Fig. 4..," We also computed the PDS of the B band data, shown by the dotted line in Fig. \ref{pds}."
+" Phe long term light curve was used to constrain the PDS at frequencies 35107.2 Uz and the intensive light curve. covered the range 2.]0'7.10"" Lz.", The long term light curve was used to constrain the PDS at frequencies $3\times10^{-8}-2\times 10^{-7}$ Hz and the intensive light curve covered the range $2\times10^{-7}-7\times 10^{-6}$ Hz.
+ We fixed az= 0.8. tthe best- value found for the X-ray PDS. for direct comparison with those data.," We fixed $\alpha_L=0.8$ , the best-fitting value found for the X-ray PDS, for direct comparison with those data."
+ Fie., Fig.
+ 5. shows the (solid lines) anc, \ref{contours} shows the (solid lines) and
+ 5. shows the (solid lines) ancl, \ref{contours} shows the (solid lines) and
+The history of star formation and chemical enrichment of galaxies is encoded in the ages and chemical compositions of their stellar populations.,The history of star formation and chemical enrichment of galaxies is encoded in the ages and chemical compositions of their stellar populations.
+ In particular. powerful insights on the processes leading to the assembly of the Galactic halo are gained by studies of the chemical abundances of their constituent. populations of field stars and elobular clusters (GC's).," In particular, powerful insights on the processes leading to the assembly of the Galactic halo are gained by studies of the chemical abundances of their constituent populations of field stars and globular clusters (GCs)."
+ It is only natural to extend such studies to the nearest giant spiral galaxy. M 31.," It is only natural to extend such studies to the nearest giant spiral galaxy, M 31."
+even cooler blackbodies (see DeLucaetal.2001).,even cooler blackbodies (see \citealt{del04}) ).
+ Thus. there is no single category of rotatiou-powered pulsar iuto which ffits neatly.," Thus, there is no single category of rotation-powered pulsar into which fits neatly."
+ The cussion of many intermecdiate-aged pulsars is dominated bv high-cucerey 5-ravs. probably even those not vet detected because of the lanited sensitivity of EGRET.," The emission of many intermediate-aged pulsars is dominated by high-energy $\gamma$ -rays, probably even those not yet detected because of the limited sensitivity of EGRET."
+ The spin parameters of aare not inconsistent with those of known y-ray pulsars., The spin parameters of are not inconsistent with those of known $\gamma$ -ray pulsars.
+" However. its location is confused with the EGRET source 3EC J1856|0111 that is about 1° fromτὸ, and is coincident with the supernova remnant Wl."," However, its location is confused with the EGRET source 3EG J1856+0114 that is about $1^{\circ}$ from, and is coincident with the supernova remnant W44."
+ This EGRET source is hard but variable (Nolanetal.2003)., This EGRET source is hard but variable \citep{nol03}.
+ We can assume that the fux of ο] J185610111. zm3.6.1019 Cres οσα? lo Gs a conservative upper lit on the eamuna-rav flux of PSR J1852|0010.," We can assume that the flux of 3EG J1856+0114, $\approx 3.6\times 10^{-10}$ ergs $^{-2}$ $^{-1}$, is a conservative upper limit on the gamma-ray flux of PSR J1852+0040."
+ Since the upper limit on the spin-down flux {πι of PSR J1852|OOI0 is ες1019 eres cin? sb. it could be an as-vot undetected 5-ray pulsar with au cficieucy of a few percent. typical of vouug or middle-aged pulsars.," Since the upper limit on the spin-down flux $\dot E/4\pi d^2$ of PSR J1852+0040 is $4\times 10^{-10}$ ergs $^{-2}$ $^{-1}$, it could be an as-yet undetected $\gamma$ -ray pulsar with an efficiency of a few percent, typical of young or middle-aged pulsars."
+ Even though the X-ray huninosity of lis consistent with miuinual NS cooling curves for au age of 10? tyr (Pageetal.2001). its blackbody temperature inplies an cmitting area that is just z0.5% of the NS surface.," Even though the X-ray luminosity of is consistent with minimal NS cooling curves for an age of $10^{3-4}$ yr \citep{pag04}, its blackbody temperature implies an emitting area that is just $\approx
+0.5\%$ of the NS surface."
+ This is cousisteut with the hiehlv modulated pulse profile comiug from a πα] rotating hot spot whose measured temperature falls well above any reasonable NS cooling curve., This is consistent with the highly modulated pulse profile coming from a small rotating hot spot whose measured temperature falls well above any reasonable NS cooling curve.
+ The most likelv region for localized heating of the NS surface is at the magnetic poles., The most likely region for localized heating of the NS surface is at the magnetic poles.
+" The canonical area for tle polar cap is A,=2287RO/Pe=~1410)? αμ”."," The canonical area for the polar cap is $A_{pc} =
+{2\pi^2 R^3 / {P c}} \approx 1 \times 10^{10}$ $^2$."
+ This is only 1054 of the area implied bv the fit to 1ο X-ray προςτι using the blackbody model., This is only $10\%$ of the area implied by the fit to the X-ray spectrum using the blackbody model.
+ In the outer-eap model for στα pulsars (Wangetal.19098).. 1ο N-rav luminosity of the hot polar cap is Dnuited n ⋝⋅↖↽↑∐↸∖≼∶≺≻↕≼⊔⋅↸∖↕↸⊳∐≓⋅↧∏∐⋜⋯↙↸⊳∏∐⋅↸∖∐↑∢∪↖↖⇁∪↕≯⊀∖⊽⋃≈⊇∖ ↕∣⋮⋝⇉⋖∫≽∣⋅∐∣⋅↱≻↴∖↴⋝⇉∐⋚↕∩⊔≼∶⋟↴∖↴↓≼∐↰⋯↴∖↴↕↑↕∐∶↴⋁⋜⋯⋜↧↖⇁↸∖↥⋅⋜↧∶↴∙⊾↸∖ ↸∖∐∖↥⋅∶↴∙⊾⋅↖↽↻↸∖↥⋅↻⋜∐⋅↑↕↸⊳↕↸∖∪↕⋟⊏⊺≈↓∙∶≩↸∖↥⋅∶↴∙∷∖↴∙↽∕∏∐∖⋯⋜⋯↕∐∐∐⊔ ∐∐∐∪↴∖↴↕⋅↖↽↕↴∖↴⊓↴⋝∪↻≈∙↗⊔∑⊢∖⊽⋃↓∖ ↽∩⋮⋝⇉↸∖↥⋅∶↴∙⋱∖↴↴∖↴↓∙∐↸∖↥⋅," In the outer-gap model for $\gamma$ -ray pulsars \citep{wan98}, the X-ray luminosity of the hot polar cap is limited by the Goldreich-Julian $e^{\pm}$ current flow of $\dot N_0 \approx 2
+\times 10^{32}\,(P/0.105\,{\rm s})^{-2}\,(B/10^{12}\,{\rm
+G})$ $^{-1}$ depositing an average energy per particle of$E_f
+\approx 4.3$ ergs."
+↸∖ ∐∖↕≯↥⋅⋜↧↸⊳↑↕∪∐⋅↗↳⊽∪↕≯↑∐↸∖↸⊳↿∐⋅↥⋅↸∖∐↑↥⋅↸∖⋜↧↸⊳∐↕∐∶↴∙⊾∐↸∖↴∖↴↿∐⋅↕⋟⋜↧↸⊳↸∖↕↴∖∷∖↴↸∖↑↑∪ ∣⊇∙↑∐∖⋯⋜⋯↕⋯⋯⊔↻∪↴∖↴↴∖↴∏," The maximum luminosity is $L({\rm bol}) \approx f
+E_f \dot N_0 < 4 \times 10^{32}$ ergs $^{-1}$."
+⋝↕↸∖↸∖↴∖↴↑↕↕⊔⋜↧↑↸∖≼↕≯∪↥⋅⋜↧↷↴≓↥⋅⋜↧⋅↖⇁↻∏↕↴∖↴⋜∐⋅ ∐∖⋜∐⋅↕↴∖↴≼∐∖⋜↧↑∐∐∐↸∖∙∖↖⊽∐∐↸∖↕≯⋜↧∐↕∐∶↴⋁↴∖↴∐∪↥⋅↑∪↕⋟↑∐↸∖≺⋔↴∖↴↸∖↥⋅↖↽↸∖≼⇂ ⊸∖⊽≓↥⋅⋜↧⋅↖↽↕⋯⊔↕∐∪↴∖↴↕↑⋅↖↽∪↕⋟↕⋟≋↕⊰⋅∐≺∖∖⋅↱↗⊇⊔∣∩∐∣," Here the fraction $f$ of the current reaching the surface is set to 1/2, the maximum possible estimated for a $\gamma$ -ray pulsar near its death line."
+ ⋝↴⋝∙↖↽⋜⋃∪↥⋅≺∐∖↥⋅⊣≻↕∟⋯⋜↧∶↴∙⊾∐↕↑⋯∐∖∙↑↕∐↴∖↴↻↥⋅↸∖≼↕↸⊳↑↕∪∐⋜∏∏≻∐↸∖↴∖↴∪∐↕⋅↖⇁ ∪⋜↧⊔⋜⋯↕⋯⋜↧∐⋅↖⇁↸∖↕−⊔↸⊳↕↸∖∐↑↷↴≓↥⋅⋜↧⋅↖↽↻∏↕↴∖↴⋜∐⋅∙⊏↕↑↕∐∖↥⋅∏∐∖↷≓ ↥⋅⋜↧⋅↖⇁↸∖↨∟↴∎↸⊳↕↸∖∐↸⊳⋅↖⇁∪↥⋅↑∐↸∖⊈⋚↕⋝∐↸∖↕≼⊔," While falling short of the observed X-ray luminosity of by an order-of-magnitude, this prediction applies only to a maximally efficient $\gamma$ -ray pulsar."
+↴∖↴↕∐↘↽↸∖↕∙↖⇁↑∪↴⋝↸∖↕∪↖↖⇁↸∖↥⋅∙↴∖↴∪↕∏↴∖↴⋯↸∖↸⊳∐⋜⋯↕↴∖↴⋯↕↴∖↴∐⋜∐⋅≼↧↻↥⋅↸∖↴∖∷∖↴↸∖≺↧↑∪⋜⋯⊳∪∏∐↑↕≯∪↥⋅↑∐↸∖⊸∖↕ ↥⋅⋜↕⋅↖↽↕∏↕∐∐∪↴∖↴↕↑⋅↖↽∪↕≯↕⋟≋↕⊰⋅∐≺∖∖⋅↱⊐⊇⊔∣∩∐∣⋈↕⋟∪↕⋜∐⋅⊣⊳⋜⋯∐↸∖⋜↧↑↕∐∶↴," Either the $\gamma$ -ray efficiency or the $B_{\rm p}$ field is likely to be lower, so this mechanism is hard pressed to account for the X-ray luminosity of."
+⋁ ⊔∪≼∐∖↕↴∖↴∪↕≯∐⋜∐⋅≺∐∐∶↴∙⊾∙∖↽⋀∖↕∏↴∖↴∐⋯∪↖↽⊔∩∩↕∙⊇∩∩⊇⋝↻↥⋅↸∖≼∐↸⊳↑↸∖↖↽↸∖∐ ↸∖↴∖↴↴∖↴⊸∖⊽≓↥⋅⋜↧⋅↖⇁↕∏∐∐∐∪↴∖↴↕↑⋅↖⇁↑∐⋜," Polar-cap heating models of \citet{har01,har02} predict even less X-ray luminosity than \citet{wan98}."
+⋯↖↖⊽⋜⋯∶↴∙⊾↸∖↑⋜↧↕∙∐∩∩≺∖∖⋝∙∙⊺⋜∐↘↽↸∖∐⋜↧↑ ace value. all such 1nodols fall short of predictiug the apparent area. feniperature. aud ποπ]τν of the N-arax enission from," Taken at face value, all such models fall short of predicting the apparent area, temperature, and luminosity of the X-ray emission from."
+" While the temperature and huuinositv of aaro ereater than those of middle-aged pulsars. its Ποπεν is less than those of ANPs. which have £L,—1079 eyes «T and thermal spectral components of kTupm0.1 keV (Mereghetti2002)."," While the temperature and luminosity of are greater than those of middle-aged pulsars, its luminosity is less than those of AXPs, which have $L_x \sim 10^{34-35.5}$ ergs $^{-1}$ and thermal spectral components of $kT_{\rm BB} \simgt 0.4$ keV \citep{mer02}."
+. The spectrum of lis sugecstive of a maguctar of low X-ray luuinosity. perhaps like the quiescent state of the transient. ANP NTE Jlslo197 (Ilalperu&Cotthelf2005).," The spectrum of is suggestive of a magnetar of low X-ray luminosity, perhaps like the quiescent state of the transient AXP XTE J1810–197 \citep{hal05}."
+". According to the magnetar theory. the X-ray ciission ultimately derives from the decay of an enormous magnetic feld (Bo>Lis10175 G: Dunean&""Thompson1996))."," According to the magnetar theory, the X-ray emission ultimately derives from the decay of an enormous magnetic field $B \simgt 4.4 \times 10^{13}$ G; \citealt{dun96}) )."
+ Although ccould be an “anomalous.” fast ANP. the implied magnetic field strength is insufficient to power the observed XN-rav Iuuinositv over the lifetime of the pulsar. estimated as Lytangc8s10H eres. since the available lmaguctic euerev is onlv ~BPR?/6=3«107(B/105)? eres.," Although could be an “anomalous,” fast AXP, the implied magnetic field strength is insufficient to power the observed X-ray luminosity over the lifetime of the pulsar, estimated as $L_x \tau_{\rm SNR} \sim
+8 \times 10^{44}$ ergs, since the available magnetic energy is only $\approx B^2 R^3/6 = 3 \times 10^{43} (B/10^{13})^2$ ergs."
+ More detailed predictions invoking the magnetar theory (e.g.. currents on twisted magnetic field lines external to the star: Thompson.Lvutikov.&νακατά20023) are μιαν] insufficient to sustain the observed N-rav enassion.," More detailed predictions invoking the magnetar theory (e.g., currents on twisted magnetic field lines external to the star; \citealt*{tho02}) ) are similarly insufficient to sustain the observed X-ray emission."
+" Although binary NS N-vay transicuts in quiescence often have huuinositics similar to that ofJ1852]0010.. their spectra. as sununnuuidzed. eg. bv ""Toiisickal.(2001). are characterized as softer blackbodies covering the full NS surface. rather thau a sanall hot spot."," Although binary NS X-ray transients in quiescence often have luminosities similar to that of, their spectra, as summarized, e.g., by \citet{tom04}, are characterized as softer blackbodies covering the full NS surface, rather than a small hot spot."
+ Even if the hotter enuüssiou frou: is hypothesized to come from residual accroetion. the current observations distavor a binary scenario based ou its steady. long-teri flix. lack of orbital Doppler delay. and absence of characteristic red noise in its timing spectrum.," Even if the hotter emission from is hypothesized to come from residual accretion, the current observations disfavor a binary scenario based on its steady long-term flux, lack of orbital Doppler delay, and absence of characteristic red noise in its timing spectrum."
+" The unclassified star «1"" from the ppositiou. while unlikely to be a binary companion of J1852]0010.. prevents us frou deriving a constrainius upper limit on optical emission from either the pulsar itself or a fall-back accretiou disk."," The unclassified star $<1^{\prime\prime}$ from the position, while unlikely to be a binary companion of , prevents us from deriving a constraining upper limit on optical emission from either the pulsar itself or a fall-back accretion disk."
+" Even if ppossesses a fossil accretion disk. it may be unable to accrete because the maeuetospheric radius is 5,—23«1oply(QUIM.CLB77 ca which is 3«10? cm for an assinued inaguetic moment y=B,,S/2zm100 C au? aud an observed L=GADQB/R3.7<010? eres s+."," Even if possesses a fossil accretion disk, it may be unable to accrete because the magnetospheric radius is $r_{\rm m} = 3 \times
+10^8\,\mu_{30}^{4/7}\,(M/M_{\odot})^{1/7}\,
+L_{37}^{-2/7}\,R_6^{-2/7}$ cm, which is $3 \times 10^9$ cm for an assumed magnetic moment $\mu = B_{\rm p}\,R^3/2 \approx 10^{30}$ G $^3$ and an observed $L = GM\dot m/R = 3.7 \times 10^{33}$ ergs $^{-1}$."
+" Therefore. the iaguetie dipole pressure ejects any potential accreting matter well outside the leht evlinder radius. rj,=eP/2x5«105 em."," Therefore, the magnetic dipole pressure ejects any potential accreting matter well outside the light cylinder radius, $r_{\ell} = cP/2\pi = 5
+\times 10^8$ cm."
+" Only in the case of By as small as τον105 G could bbea ""slow rotator.” with Pz since the equilibriun (or απ). period for disk P4.accretion is Pog=3.6pho)(AL/AL.)PTBoPΠρο κ, While such a value of By is common among low-niass X-ray binaries. it would be surprising for such a voung NS."," Only in the case of $B_{\rm p}$ as small as $7
+\times 10^8$ G could be a “slow rotator,” with $P \approx
+P_{\rm eq}$, since the equilibrium (or minimum) period for disk accretion is $P_{\rm eq} = 3.6\,\mu_{30}^{6/7}\,(M/M_{\odot})^{-2/7}\,
+L_{37}^{-3/7}\,R_6^{-3/7}$ s. While such a value of $B_{\rm p}$ is common among low-mass X-ray binaries, it would be surprising for such a young NS."
+ For an intermediate value of the maguetic feld streneth. ccould be iu the propeller regime. P—Lg. ta which matter is flune out from the maenuetospheric radius :t a rate ip. which causes it to spin down at a rate P (ogMenonetal.1999:Zavli- 2001). where Zz107 ο em? is the NS momeut of inertia.," For an intermediate value of the magnetic field strength, could be in the propeller regime, $P < P_{\rm eq}$, in which matter is flung out from the magnetospheric radius at a rate $\dot m$, which causes it to spin down at a rate $\dot P \approx 2\,\dot m\,r_{\rm
+m}^2\,I^{-1}\,P\,(1-P/P_{\rm eq})$ \citep[e.g.,][]{men99,zav04}, , where $I
+\approx 10^{45}$ g $^2$ is the NS moment of inertia."
+" The observed upper IitP«7«10 ll sol sets an upper mit of ii<3.7«101(4,/105011)2 gs lin the propeller scenario."," The observed upper limit$\dot P < 7 \times 10^{-14}$ s $^{-1}$ sets an upper limit of $\dot m < 3.7 \times 10^{16}\,(r_{\rm m}/10^8\,{\rm
+cm})^{-2}$ g $^{-1}$ in the propeller scenario."
+ But in that case. it is not clear how the hiehlv pulsed. thermal ταν enission is produced.," But in that case, it is not clear how the highly pulsed, thermal X-ray emission is produced."
+ Even if the bulk of the fall-back material is ejected. as," Even if the bulk of the fall-back material is ejected, as"
+Cataclysmic variable stars (CVs) are short-perioc binary systems which typically consist of a cool main-sequence star transferring mass via a gas stream ancl accretion disc to a white dwarf primary.,Cataclysmic variable stars (CVs) are short-period binary systems which typically consist of a cool main-sequence star transferring mass via a gas stream and accretion disc to a white dwarf primary.
+" Phe impact of the stream. with the accretion disce forms a so-called ""bright spot’. which in systems that are significantly. inclined to our line of sight can cause a rise in the observed Hux as this region rotates into view. resulting in an ""orbital hump' in the lightcurve."," The impact of the stream with the accretion disc forms a so-called `bright spot', which in systems that are significantly inclined to our line of sight can cause a rise in the observed flux as this region rotates into view, resulting in an `orbital hump' in the lightcurve."
+ In high-inclination svstems eclipses of the white ewarf. bright spot ancl dise by the red. cwarl secondary can also occur.," In high-inclination systems eclipses of the white dwarf, bright spot and disc by the red dwarf secondary can also occur."
+ Analysis of these eclipses can vield determinations of system parameters such as the mass ratioq. the orbital inclination and the radius of the accretion disc Ly citealtwoodsOa)).," Analysis of these eclipses can yield determinations of system parameters such as the mass ratio, the orbital inclination and the radius of the accretion disc $R_{d}$ \\citealt{wood89a}) )."
+ Eclipsing systems are therefore. valuable sources of data on CVs., Eclipsing systems are therefore valuable sources of data on CVs.
+ Dwarf novae are a sub-tvpe of CVs which show intermittent luminosity increases of 25 magnitudes. known as outbursts.," Dwarf novae are a sub-type of CVs which show intermittent luminosity increases of 2–5 magnitudes, known as outbursts."
+ X. Further. sub-tvpe. of dwarf. novae. are the SU. UAla stars. which exhibit) superoutbursts. at regular intervals. during which the luminosity increases bv 0.7 magnitudes over the normal outburst maximum.," A further sub-type of dwarf novae are the SU UMa stars, which exhibit superoutbursts at regular intervals, during which the luminosity increases by $\sim0.7$ magnitudes over the normal outburst maximum."
+ These superoutbursts are characterised by the presence of superhumps increases in brightness that usually recur at a slightly longer period than the orbital evcle., These superoutbursts are characterised by the presence of superhumps – increases in brightness that usually recur at a slightly longer period than the orbital cycle.
+ Phere is found to be a relationship between this superhump period excess ¢ and the mass ratio C2).., There is found to be a relationship between this superhump period excess $\epsilon$ and the mass ratio \citep{patterson98}.
+ Determinations of the mass ratios of SU UAla stars are therefore useful to calibrate this relation. which can then be used to determine the mass ratios ofother SU UMa stars.," Determinations of the mass ratios of SU UMa stars are therefore useful to calibrate this relation, which can then be used to determine the mass ratios of other SU UMa stars."
+ OU Vir is a faint (V —Ls: 7)) eclipsing CV with a period. of 1.75 hr which has been seen in outburst and probably superoutburst (2).. marking it as à SU UMa dwarf nova.," OU Vir is a faint (V $\sim18$; \citealt{mason02}) ) eclipsing CV with a period of 1.75 hr which has been seen in outburst and probably superoutburst \citep{vanmunster00}, , marking it as a SU UMa dwarf nova."
+ ?. presented. time-resolved. multi-colour. photometry ancl spectroscopy of OU Vir. concluding that the eclipse is of the bright spot and disc. but not the white dwarf.," \citet{mason02} presented time-resolved, multi-colour photometry and spectroscopy of OU Vir, concluding that the eclipse is of the bright spot and disc, but not the white dwarf."
+ In this paper we present lighteurves of OU Vir. obtained with ULTILACAM. an ultra-fast. triple-beam CCD camera: for more details see ?:: Dhillon et al..," In this paper we present lightcurves of OU Vir, obtained with ULTRACAM, an ultra-fast, triple-beam CCD camera; for more details see \citet{dhillon01b}; Dhillon et al.,"
+ in preparation., in preparation.
+ OU. Vir was observed on the nights of 16 May. 2002 and 19. 20. 22 and 25 May 2003 usingULTILACAM on the," OU Vir was observed on the nights of 16 May 2002 and 19, 20, 22 and 25 May 2003 usingULTRACAM on the"
+but to reconstruct the visibilities on each bolometer separately and combine the visibilities afterwards.,but to reconstruct the visibilities on each bolometer separately and combine the visibilities afterwards.
+ Such a strategy would increase the length of the phase-shifting sequences. but in a reasonable (and tractable) way thanks to the intrinsic shortness of our proposed phase-shifting scheme.," Such a strategy would increase the length of the phase-shifting sequences, but in a reasonable (and tractable) way thanks to the intrinsic shortness of our proposed phase-shifting scheme."
+~1210Hs 4.,$\sim 1 \times 10^{-14}$ $^{-1}$.
+ Phe model results appear to be consistent with the observational data., The model results appear to be consistent with the observational data.
+ We consider then in particular for the cxtragalactic comparison. galaxies where AGN - and/or starburst - activities may have enhanced the cosmic ray ionization rates as well as. in some cases. created ULIICGs.," We consider then in particular for the extragalactic comparison, galaxies where AGN - and/or starburst - activities may have enhanced the cosmic ray ionization rates as well as, in some cases, created ULIRGs."
+ In ? we selected: well-known galaxies. such as Arp 220 or M S2 as examples of sources with active nuclei. and where therefore the cosmic ray ionization rate may be enbanced.," In \citet{Baye09a} we selected well-known galaxies, such as Arp 220 or M 82 as examples of sources with active nuclei, and where therefore the cosmic ray ionization rate may be enhanced."
+ Arp 220 is the prototypical ultraluminous galaxy while M82 is the prototypical starburst., Arp 220 is the prototypical ultraluminous galaxy while M82 is the prototypical starburst.
+ Recently (2). ο has been discovered in these two regions and its fractional abundance has been estimated to be ~ 10 7., Recently \citep{VanderTak08} $_{3}$ $^+$ has been discovered in these two regions and its fractional abundance has been estimated to be $\sim$ $\times$ $^{-9}$.
+" ""These authors fou that observations of M. 82 are matched by a hieh-¢ PDR. Le. an evolved starburst while N-ray models are best a reproducing the observations of Arp 220."," These authors found that observations of M 82 are matched by a $\zeta$ PDR, i.e., an evolved starburst while X-ray models are best at reproducing the observations of Arp 220."
+ In fact. our models indicate that one can obtain high abundance of this ion à low (i.e in à PDR) as well as high (ie dense star forming eas) extinction as long as the ¢ is ~ 13 l and the metallicity is solar.," In fact, our models indicate that one can obtain high abundance of this ion at low (i.e in a PDR) as well as high (i.e dense star forming gas) extinction as long as the $\zeta$ is $\sim$ $^{-13}$ $^{-1}$ and the metallicity is solar."
+ At high extinction. which could represen the nuclear part of the galaxy. the abundance is higher. providing possibly a better match for the observations.," At high extinction, which could represent the nuclear part of the galaxy, the abundance is higher, providing possibly a better match for the observations."
+ We note here that we are not attempting to model Arp 220 that. since as 2? pointed out. it has a quite unusual geometry.," We note here that we are not attempting to model Arp 220 that, since as \citet{VanderTak08}
+ pointed out, it has a quite unusual geometry."
+ Another interesting object which has been recently studied in molecular emission is Mrk 231. a ULIBCG.," Another interesting object which has been recently studied in molecular emission is Mrk 231, a ULIRG."
+ A hieh resolution SPIRE FILS spectrum reveals the presence of ions such as and H120 (?)..," A high resolution SPIRE FTS spectrum reveals the presence of ions such as $^+$, $^+$ and $_2$ $^+$ \citep{VanderWerf10}."
+ While abuncances are not derived. we can use our Table 2. to determine what tvpe of model is able to produce high fractional abundances 10.| 27) of these three ions.," While abundances are not derived, we can use our Table \ref{tab:3} to determine what type of model is able to produce high fractional abundances $\ge$ $^{-10}$ ) of these three ions."
+. We ⇁⋅find that the only. regime. is. an environment with low metallicity (0.1. solar). very hieh cosmic. rav ionisation.- rates (2 1017? 1 *) and low visual extinction.," We find that the only regime is an environment with low metallicity (0.1 solar), very high cosmic ray ionisation rates $\ge$ $^{-16}$ $^{-1}$ ) and low visual extinction."
+ ? explained the high abundances of these ions bv involving XDIt-chemistry., \citet{VanderWerf10} explained the high abundances of these ions by involving XDR-chemistry.
+ We find that. in agreement with 7.. determining the origin of molecular emission from λος such as Mrk 231 is not trivial when both sources of energy (CLR ane N-ravs) are present.," We find that, in agreement with \citet{Papa10a}, determining the origin of molecular emission from ULIRGs such as Mrk 231 is not trivial when both sources of energy (CR and X-rays) are present."
+ Finally. our results are necessarily indicative rather than specific in that we do not estimate molecular line intensities.," Finally, our results are necessarily indicative rather than specific in that we do not estimate molecular line intensities."
+ Our study was motivated. by the recent. investigation of filaments around the central galaxies of clusters of galaxies bv ? and by the considerations of the ellects of high cosmic rav ionization rates in ULIRGs by ? as well as recent Lerschel results., Our study was motivated by the recent investigation of filaments around the central galaxies of clusters of galaxies by \citet{Baye10a} and by the considerations of the effects of high cosmic ray ionization rates in ULIRGs by \citet{Papa10a} as well as recent Herschel results.
+ We find that several species. many detected in extragalactic environments. are in fact tracers of very high ionization fractions.," We find that several species, many detected in extragalactic environments, are in fact tracers of very high ionization fractions."
+ The general conclusions that we can draw from this study are as follows:, The general conclusions that we can draw from this study are as follows:
+come back to considering how these characteristics should be interpreted below.,come back to considering how these characteristics should be interpreted below.
+" In our observations, B19444-17 nulls about 2/3 of the time, somewhat higher than the value given by DCHR, but closer to null percentage reported by Rankin, (1986)."," In our observations, B1944+17 nulls about 2/3 of the time, somewhat higher than the value given by DCHR, but closer to null percentage reported by Rankin, (1986)."
+" The majority of these null pulses can readily be distinguished from the bursts; however, there is a small portion of weak pulses that are difficult to identify as either nulls or pulses."," The majority of these null pulses can readily be distinguished from the bursts; however, there is a small portion of weak pulses that are difficult to identify as either nulls or pulses."
+" Interestingly, the distinction between nulls and pulses is easier to define at L band, as can be seen in the respective null histograms of Fig. 2.."," Interestingly, the distinction between nulls and pulses is easier to define at L band, as can be seen in the respective null histograms of Fig. \ref{nullhistograms}."
+" Given that the nulls and pulses cannot be fully distinguished, we can choose an intensity threshold that will be conservative and reliable either in selecting pulses or nulls, but not both."," Given that the nulls and pulses cannot be fully distinguished, we can choose an intensity threshold that will be conservative and reliable either in selecting pulses or nulls, but not both."
+" In Fig. 2,,"," In Fig. \ref{nullhistograms},"
+" we have taken the latter option—that is, using low thresholds that will tend to slightly underestimate the null population."," we have taken the latter option—that is, using low thresholds that will tend to slightly underestimate the null population."
+" Then, using this conservative discriminator of nulls, we have computed the burst- and null-length histrograms in Figure 3.."," Then, using this conservative discriminator of nulls, we have computed the burst- and null-length histrograms in Figure \ref{burstnullfreq}."
+" These show that 1-pulse bursts and nulls have the highest frequency, but we see that very long bursts and nulls also occur."," These show that 1-pulse bursts and nulls have the highest frequency, but we see that very long bursts and nulls also occur."
+" In the 7000- 327-MHz observation, for instance, a small number of bursts of 40-50 P, and two of 80-90 P, were encountered alongside the more frequent long nulls ranging up to 300 Pi."," In the 7000-pulse 327-MHz observation, for instance, a small number of bursts of 40-50 $P_1$ and two of 80-90 $P_1$ were encountered alongside the more frequent long nulls ranging up to 300 $P_1$."
+ Even qualitatively we immediately see that the nulls in B1944+17 are distributed within the PS in a very non-random manner., Even qualitatively we immediately see that the nulls in B1944+17 are distributed within the PS in a very non-random manner.
+" Recent investigations into pulsar nulling have raised two important new questions about their distributions: a) whether they are randomly distributed(e.g.,, Redman Rankin 2009; Rankin Wright 2007); and b) whether they are periodic (HR07/09)."," Recent investigations into pulsar nulling have raised two important new questions about their distributions: a) whether they are randomly distributed, Redman Rankin 2009; Rankin Wright 2007); and b) whether they are periodic (HR07/09)."
+" With such a large null fraction, one would expect to see few long sequences in any given observation."," With such a large null fraction, one would expect to see few long sequences in any given observation."
+ The tendency of B1944+17’s bursts and nulls to clump into sequences of roughly 20-100 pulses immediately indicates a non-random distribution., The tendency of B1944+17's bursts and nulls to clump into sequences of roughly 20-100 pulses immediately indicates a non-random distribution.
+ Application of the, Application of the
+"The average amplitude of quasar variability has been seen to depend on several factors: time lag between measurements, Iuninositv of the quasar (2.. 2.. 2.. 2 εν 7)). wavoleugth of observation (?2.. 7.. 24). aud black hole mass (7... 23).","The average amplitude of quasar variability has been seen to depend on several factors: time lag between measurements, luminosity of the quasar \citealt{vandenberk04}, , \citealt{devries05}, , \citealt{wilhite08}, , \citealt{bauer09a}, \citealt{macleod10}, , \citealt{meusinger11}) ), wavelength of observation \citealt{vandenberk04}, \citealt{devries05}, \citealt{meusinger11}) ), and black hole mass \citealt{wilhite08}, \citealt{bauer09a}) )."
+ The dependence of variability amplitude on redshift is less obvious: ? measured a slight increase in the variability with redshift. while ?— measured a slight decrease.," The dependence of variability amplitude on redshift is less obvious; \cite{vandenberk04} measured a slight increase in the variability with redshift, while \cite{devries05} measured a slight decrease."
+ More receuthl ? and? have measured no significant dependeuce of variability amplitude ou redshift.," More recently, \cite{macleod10} and \cite{meusinger11} have measured no significant dependence of variability amplitude on redshift."
+ Iu practice. the vauiabilitv-Iuninositv trend measured is of the form: A linear relation has indeed been observed (?.. ?)). although there are conflicting results for the value of he power-law slope o. perhaps due to selection effects.," In practice, the variability-luminosity trend measured is of the form: A linear relation has indeed been observed \citealt{vandenberk04}, \citealt{bauer09a}) ), although there are conflicting results for the value of the power-law slope $\alpha$, perhaps due to selection effects."
+ If faint quasars are included in the analysis for which one cannot observe the full exteut of the variability. the ucasured slope will become artificially shallow: this effect nost clearly manifests itself as a flattening in the relation at the lowest observable Iuminosities. as is illustrated iu feure 5b of ?..," If faint quasars are included in the analysis for which one cannot observe the full extent of the variability, the measured slope will become artificially shallow; this effect most clearly manifests itself as a flattening in the relation at the lowest observable luminosities, as is illustrated in figure 5 of \cite{bauer09a}."
+ The low-huninosity luit of the binning scheme in this work is chosen to exclude this regime frou he data set., The low-luminosity limit of the binning scheme in this work is chosen to exclude this regime from the data set.
+ The value of the coustaut C depeuds ou he details of the normalization of the data. as described yolow.," The value of the constant $C$ depends on the details of the normalization of the data, as described below."
+ When studyiug how the variability of a large quasar sample depends ou one of the quasars’ properties. we uust treat the parameters as iudependcutly as possible.," When studying how the variability of a large quasar sample depends on one of the quasars' properties, we must treat the parameters as independently as possible."
+ To this end. we use a mnethod introduced bv 7 and adopted in ὧν," To this end, we use a method introduced by \cite{vandenberk04} and adopted in \cite{bauer09a}."
+ Four basic quantities are kuown or all of the quasars iu our sample: tine lag between neasnreimients 7. Iniinositv L. estimated black hole mass M. aud redshift :.," Four basic quantities are known for all of the quasars in our sample: time lag between measurements $\tau$, luminosity $L$, estimated black hole mass $M$, and redshift $z$."
+ There are kuown correlations between all of these parameters. due to plivsical relationships or artificial effects such as detection biases in fux-lTiuüted survers.," There are known correlations between all of these parameters, due to physical relationships or artificial effects such as detection biases in flux-limited surveys."
+ To avoid these complications aud study only the dependence of variability on hiinosity. we would like to identify a set of quasars with identical properties except for their luminosity. aud then examine how the variability differs between them.," To avoid these complications and study only the dependence of variability on luminosity, we would like to identify a set of quasars with identical properties except for their luminosity, and then examine how the variability differs between them."
+ To approximate this procedure. we have split cach parameters range uto bius: 8 bins in r. G bius in AM. 6 bins in z. and 1 bins iu £.," To approximate this procedure, we have split each parameter's range into bins: 8 bins in $\tau$, 6 bins in $M$, 6 bins in $z$, and 4 bins in $L$."
+ The biu limits are given in table 1: quasars with properties outside the eiven ranges are not used in the analysis., The bin limits are given in table \ref{bin_limits}; quasars with properties outside the given ranges are not used in the analysis.
+ To iieasure quasar variability we use a quantity simular to that of the structure function., To measure quasar variability we use a quantity similar to that of the structure function.
+ We define: where Am is the magnitude difference between two incdependecut observations of au object. aud c6 ids the error on those imeasurements.," We define: where $\Delta m$ is the magnitude difference between two independent observations of an object, and $\sigma$ is the error on those measurements."
+ This is similar to the structure function as used in 7? ancl ?: however. here instead of beiug an ensemble measurement. oue V ids measured for cach pair of magnitude measurements of a quasar.," This is similar to the structure function as used in \cite{vandenberk04} and \cite{bauer09a}; however, here instead of being an ensemble measurement, one $V$ is measured for each pair of magnitude measurements of a quasar."
+ Four measurements of a quasar will vield six Arm ineasurcmecnts. and therefore six different V nieasureients for the single quasar.," Four measurements of a quasar will yield six $\Delta m$ measurements, and therefore six different $V$ measurements for the single quasar."
+ As V is imaginary when Ar is less than the measurement error σ. we ouly use data which show significant (7 lo) variability.," As $V$ is imaginary when $\Delta m$ is less than the measurement error $\sigma$, we only use data which show significant $>1 \sigma$ ) variability."
+ This cut on the data is described further im section 5.1.., This cut on the data is described further in section \ref{datacuts_section}.
+ For each multidimensional biu. a mean variability anplitude V is determined bv taking the mean of all V. values measured for that biu.," For each multi-dimensional bin, a mean variability amplitude $\overline{V}$ is determined by taking the mean of all $V$ values measured for that bin."
+ Then. holding coustaut the iudices for time lag. mass. redshift. aud waveleneth. one can compare the V. values across the £L bins.," Then, holding constant the indices for time lag, mass, redshift, and wavelength, one can compare the $\overline{V}$ values across the 4 $L$ bins."
+ This procedure vields 8&ος6=288 possible Lpoint plots of mean variability amplitude Wo versus hDmunuinositv. or 1152 possible V. values.," This procedure yields $8 \times 6 \times 6 = 288$ possible 4-point plots of mean variability amplitude $\overline{V}$ versus luminosity, or 1152 possible $\overline{V}$ values."
+ Most of the imulti-diieusional bins are not well populated by the quasar sample (for example. ligh-redshift low-luninosity bins).," Most of the multi-dimensional bins are not well populated by the quasar sample (for example, high-redshift low-luminosity bins)."
+ Iu fact. we obtain 103 bins with at least 50 measurement pairs. which is the mini. we require in order to aclequatcly determine V.," In fact, we obtain 403 bins with at least 50 measurement pairs, which is the minimum we require in order to adequately determine $\overline{V}$."
+ To examine the overall behavior of V with respect to L one can normalize the [-poiut measured trends together and average the resulting uoriualized data in each £ bin to fud a simple. mecamineful result of how the variability scales with the quasar Iuninosity.," To examine the overall behavior of $\overline{V}$ with respect to $L$ one can normalize the 4-point measured trends together and average the resulting normalized data in each $L$ bin to find a simple, meaningful result of how the variability scales with the quasar luminosity."
+ The normalization consists of an additive constant in log(V}. ic. each L-poiut measured trend has its own constant C as defined in equation 2..," The normalization consists of an additive constant in $\overline{V}$ ), i.e. each 4-point measured trend has its own constant $C$ as defined in equation \ref{v_vs_l_equation}."
+ The 7. M. aud :onudtidunensonal bin that has the best statistics is chosen to be the standard. aud the log(V) versus log(L) treuds from all other 7. M. aud : bins are normalized to that standard using one constant offset per r. M. 2 combination.," The $\tau$, $M$, and $z$ multi-dimensional bin that has the best statistics is chosen to be the standard, and the $\overline{V}$ ) versus $L$ ) trends from all other $\tau$, $M$, and $z$ bins are normalized to that standard using one constant offset per $\tau$, $M$, $z$ combination."
+ The constant is determined by ininimizine the chi square difference. between the V values frou the two datasets in the same £ bin. for the L bius where there exist data from both sets;," The constant is determined by minimizing the chi square difference between the $\overline{V}$ values from the two datasets in the same $L$ bin, for the $L$ bins where there exist data from both sets."
+ For a visual representation of the normalization procedure. see fieure [d in ?..," For a visual representation of the normalization procedure, see figure 4 in \cite{bauer09a}."
+" After averaging the normalized data. we are left with one Vom, versus £ troud with arbitrary y axis normalization but meauinetul slope."," After averaging the normalized data, we are left with one $V_{\mathrm{norm}}$ versus $L$ trend with arbitrary $y$ axis normalization but meaningful slope."
+ This nonualization technique has been shown to give results for variability amplitude versus timelag 7 that are consistent with indepeudent measurements (see table liu ?.., This normalization technique has been shown to give results for variability amplitude versus timelag $\tau$ that are consistent with independent measurements (see table 4 in \cite{bauer09a}.
+ In this wav. we study how the variability scales with bhuunmositv. comparing only objects with simular values of the other parameters.," In this way, we study how the variability scales with luminosity, comparing only objects with similar values of the other parameters."
+" After the normalization. deviations from the mean Vu,£ relation will not be caused by known. but lensine-indepencdent. correlations such as that between variability aud time lag."," After the normalization, deviations from the mean $V_{\mathrm{norm}}-L$ relation will not be caused by known, but lensing-independent, correlations such as that between variability and time lag."
+ Using the normalization constants calculated in this wav. cach measured variability amplitude V. is normalized according to its r. AZ. and : bin: the resulting Vac then canbe usedto estimatethequasars lensing magnification.as detailed below.," Using the normalization constants calculated in this way, each measured variability amplitude $V$ is normalized according to its $\tau$ , $M$ , and $z$ bin; the resulting $V_{\mathrm{norm}}$ then canbe usedto estimatethequasar's lensing magnification,as detailed below."
+ We note that. for data taken in a single pass-band. the," We note that, for data taken in a single pass-band, the"
+calculated small-scale density spectra in turbulent interstellar plasmas.,calculated small-scale density spectra in turbulent interstellar plasmas.
+ They also considered compressible turbulence in plasmas with ;2 less than unity that mateh those found in the magnetic tubes of T Tauri stars., They also considered compressible turbulence in plasmas with $\beta$ less than unity that match those found in the magnetic tubes of T Tauri stars.
+ Thev argued that the dvnamies of the eascade is roughly independent of 3., They argued that the dynamics of the cascade is roughly independent of $\beta$.
+ In the above mentioned papers. il was assumed (hat ATID wave-packets propagate al the Alfvénn speed. in a direction either parallel or antiparallel with respect to the local mean magnetic field and that the nonlinear interactions are restricted to collisions between oppositelv. directed wave-packets.," In the above mentioned papers, it was assumed that MHD wave-packets propagate at the Alfvénn speed, in a direction either parallel or antiparallel with respect to the local mean magnetic field and that the nonlinear interactions are restricted to collisions between oppositely directed wave-packets."
+ Maron&Goldreich(2001) made simulations of ihe interaction between oppositely directed. Alfvénn waves., \citet{maron..01} made simulations of the interaction between oppositely directed Alfvénn waves.
+" In our previous paper (Paper D)). instead. of a wave spectrum. we assumed a mean frequency, 2=Fue; while in this paper. we (reat different. [requency values."," In our previous paper \citeauthor{vasc..00}) ), instead of a wave spectrum, we assumed a mean frequency, $\varpi = F w_i$, while in this paper, we treat different frequency values."
+ The wave frequency obtained. from laboratory experiments (e.g..Burkeetal.1998)... ον0.101. could be taken as the upper limit for the wave spectrum.," The wave frequency obtained from laboratory experiments \citep[e.g.,][]{moralez..98}, $\Omega_{sup} \sim 0.1 \Omega_i$ , could be taken as the upper limit for the wave spectrum."
+ On the other hand. Scheurwater showed that Allvénn wave [requencies can be greater than 107j.," On the other hand, \citet{scheur..88} showed that Alfvénn wave frequencies can be greater than $10^{-5} \Omega_i$."
+ We. therefore. study the frequency interval 10.20;<Q<0.10;," We, therefore, study the frequency interval $10^{-5} \Omega_i \leq \Omega
+\leq 0.1 \Omega_i$."
+ Assuming this frequency spectrum. we caleulate the damping lengths and corresponding f.—(002)2/e4 parameter for all the damping mechanisms treated in section 3..," Assuming this frequency spectrum, we calculate the damping lengths and corresponding $f = \langle \delta v^2 \rangle^{1/2}
+/ v_A$ parameter for all the damping mechanisms treated in section \ref{dampmech}."
+ Thus. hereafter. f is not a [ree parameter anvmore.," Thus, hereafter, $f$ is not a free parameter anymore."
+ We assume that the energy released by the damping of the Alfvénn waves in a given volume is radiated away., We assume that the energy released by the damping of the Alfvénn waves in a given volume is radiated away.
+ Thus. we have πι = T1032) where I is given by Figure 2.. 7 is the optical clepth. σ is the Stefan-Doltzmann constant. and 2d. (he area.," Thus, we have dV = T^4 A, where $T$ is given by Figure \ref{fig2}, , $\tau$ is the optical depth, $\sigma$ is the Stefan-Boltzmann constant, and $A$, the area."
+ For a given 7 and temperature profile. we know the amount of energy that is radiated away.," For a given $\tau$ and temperature profile, we know the amount of energy that is radiated away."
+ This must be equal to the energy dissipated by the Alfvénn waves. which is determined by / and the damping mechanism usec.," This must be equal to the energy dissipated by the Alfvénn waves, which is determined by $f$ and the damping mechanism used."
+ Figure 7 shows the results of our caleulations. which were made for three different values of the optical depth 7 and for 10?<F0.1.," Figure \ref{fig7} shows the results of our calculations, which were made for three different values of the optical depth $\tau$ and for $10^{-5}
+\leq F \leq 0.1$."
+ Since the star is nol strongly obscured bv the tube. it is Likely that the mean optical depth is less than 1.," Since the star is not strongly obscured by the tube, it is likely that the mean optical depth is less than 1."
+ Taking into account the values of the adiabatic heating (Martin1996) and calculating the associated optical depth. we obtain 7~107.," Taking into account the values of the adiabatic heating \citep*{martin} and calculating the associated optical depth, we obtain $\tau \sim 10^{-3}$."
+ However. lor the sakeof completeness. we made the calculations [or 7T=[0? and 7= Las well.," However, for the sakeof completeness, we made the calculations for $\tau = 10^{-2}$ and $\tau = 1$ as well."
+ Figure 7 shows the damping lengths divided by the length of the tube. caleulated according to equation (32)). for the nonlinear and turbulent. damping mechanisms mentioned above.," Figure \ref{fig7} shows the damping lengths divided by the length of the tube, calculated according to equation \ref{consist}) ), for the nonlinear and turbulent damping mechanisms mentioned above."
+ The left column of Figure 7. shows the nonlinear dampinglengths., The left column of Figure \ref{fig7} shows the nonlinear dampinglengths.
+ We note that even [or 7=10? with the lowest frequeney (F=10 7). the damping length is less than the," We note that even for $\tau = 10^{-3}$ with the lowest frequency $(F = 10^{-5})$ , the damping length is less than the"
+We have calculated the linear torque exerted by a planet on a circular orbit on a disc containing à toroidal magnetic field.,We have calculated the linear torque exerted by a planet on a circular orbit on a disc containing a toroidal magnetic field.
+ In contrast to the nonmagnetic case. there is no singularity at the corotation radius. where the frequency of the perturbation matches the orbital frequency.," In contrast to the nonmagnetic case, there is no singularity at the corotation radius, where the frequency of the perturbation matches the orbital frequency."
+ However. there are two new singularities on both sides of corotation. where the frequency of the perturbation in a frame rotating with the uid matches that of à slow MIID wave propagating along the field line.," However, there are two new singularities on both sides of corotation, where the frequency of the perturbation in a frame rotating with the fluid matches that of a slow MHD wave propagating along the field line."
+ These socalledresonances are Closer to the planet than the Lindblad resonances., These so–called are closer to the planet than the Lindblad resonances.
+ In addition. on each side of the planet. there are two turning points located bevond the magnetic resonance.," In addition, on each side of the planet, there are two turning points located beyond the magnetic resonance."
+ One of them coinncides with the Lindblad resonance., One of them ncides with the Lindblad resonance.
+ For values of m larger than some critical value moa. there is a third turning point between the magnetic resonance and corotation. on each side of the planet.," For values of $m$ larger than some critical value $m_{\rm
+crit}$, there is a third turning point between the magnetic resonance and corotation, on each side of the planet."
+ Like in the nonmagnetic case. waves propagate outside the outermost turning points.," Like in the nonmagnetic case, waves propagate outside the outermost turning points."
+ But here they also propagate inside the intermediate turning points when mbmg. or between the intermediate and innermost turning points when mmo (see figure 1)).," But here they also propagate inside the intermediate turning points when $m<m_{\rm crit}$, or between the intermediate and innermost turning points when $m>m_{\rm crit}$ (see figure \ref{fig1}) )."
+ The singular modes excited at the magnetic resonances can therefore propagate., The singular modes excited at the magnetic resonances can therefore propagate.
+ There is a significant torque exerted on the region of the disc inside the outermost turning points (which comnnceide with the Lindblad resonances for the values of m of interest)., There is a significant torque exerted on the region of the disc inside the outermost turning points (which ncide with the Lindblad resonances for the values of $m$ of interest).
+ Like the Lindblad torque. this torque is negative inside the planet's orbit and. positive outside its orbit.," Like the Lindblad torque, this torque is negative inside the planet's orbit and positive outside its orbit."
+ Phe whole region around corotation. not just à narrow zone around the magnetic resonances. contribute to this torque.," The whole region around corotation, not just a narrow zone around the magnetic resonances, contribute to this torque."
+ In other words. the magnetic resonances contribute to a global. not a point-like. torque.," In other words, the magnetic resonances contribute to a global, not a point-like, torque."
+ Since these resonances are closer to the planet than the Lindblad resonances. they couple more strongly to the tidal potential.," Since these resonances are closer to the planet than the Lindblad resonances, they couple more strongly to the tidal potential."
+ Pherefore. the torque exerted around the magnetic resonances dominate over the Lindblad torque if the magnetic field is large enough.," Therefore, the torque exerted around the magnetic resonances dominate over the Lindblad torque if the magnetic field is large enough."
+ Lf in addition 3—c/e3 increases fast enough with radius. the outer magnetic resonance becomes less important (it disappears altogether when there is no magnetic field outside the planets orbit) and the total torque is then negative. dominated by the inner magnetic resonance.," If in addition $\beta \equiv c^2/v_A^2$ increases fast enough with radius, the outer magnetic resonance becomes less important (it disappears altogether when there is no magnetic field outside the planet's orbit) and the total torque is then negative, dominated by the inner magnetic resonance."
+ This corresponds to à positive torque on the planet. which leads to outward migration.," This corresponds to a positive torque on the planet, which leads to outward migration."
+ The amount by which 3 has to increase outward for the total torque exerted on the cise to be negative depends mainly on the magnitude of 3., The amount by which $\beta$ has to increase outward for the total torque exerted on the disc to be negative depends mainly on the magnitude of $\beta$.
+ We have found that for 51 at corotation. the cumulative torque (obtained by summing up the contributions from all the values of m) exerted on the dise is negative when 2 increases at least as fast as r7.," We have found that for $\beta =1$ at corotation, the cumulative torque (obtained by summing up the contributions from all the values of $m$ ) exerted on the disc is negative when $\beta$ increases at least as fast as $r^2$."
+ Wo x£c. the cumulative. torque becomes negative. for⋅ values of⋅ ή) between 107⊐ and 10. whereas it⋠⋠ is negative. for⋅ ην)=107Dp if 47Xa.1 ," If $\beta
+\propto r^3$, the cumulative torque becomes negative for values of $\beta(r_p)$ between $10^2$ and 10, whereas it is negative for $\beta(r_p)=10^2$ if $\beta \propto r^4$."
+The migration timescales that correspond to the torques calculated above are rather short., The migration timescales that correspond to the torques calculated above are rather short.
+" The orbital decay timescale of a planet of mass AM, at radius ry, is 7=M,FO,EAL). where 1) is the cumulative torque exerted by the planet on the disc."," The orbital decay timescale of a planet of mass $M_p$ at radius $r_p$ is $\tau = M_p r_p^2 \Omega_p / |T|$, where $T$ is the cumulative torque exerted by the planet on the disc."
+" Remembering that Pox(AL,/AL Νο get: In a standard dise model. X~100. 10 & em2 at Lau (see. for instance. Papaloizou Terquem. 1999)."," Remembering that $T \propto (M_p/M_\ast)^2$, we get: In a standard disc model, $\Sigma \sim
+100$ $10^3$ g $^{-2}$ at 1 au (see, for instance, Papaloizou Terquem 1999)."
+" Pherefore. 7~10 10"" vr for a one earth mass planet at 1 au in à nonmagnetic disc. as 7~10? in that case (sce fig. 113)."," Therefore, $\tau \sim 10^5$ $10^6$ yr for a one earth mass planet at 1 au in a nonmagnetic disc, as $\tilde{T} \sim 10^3$ in that case (see fig. \ref{fig6}) )."
+ This is in agreement with Ward (1986. 1997).," This is in agreement with Ward (1986, 1997)."
+ In a magnetic disc. |Z] may become larger. leading to an even shorter migration timescale (note that given the limited accuracy of the numerical scheme used. as indicated in section 7.2.2.. only orders of magnitude for the migration timescale are obtained here).," In a magnetic disc, $|\tilde{T}|$ may become larger, leading to an even shorter migration timescale (note that given the limited accuracy of the numerical scheme used, as indicated in section \ref{sec:caseBne0}, only orders of magnitude for the migration timescale are obtained here)."
+ However. it: is important to keep in mind that these timescales arefecal.," However, it is important to keep in mind that these timescales are."
+ Once the planet. migrates outward out of the region where 2 increases with radius. it may enter a region where 3 behaves dillerently and then resume inward migration for instance.," Once the planet migrates outward out of the region where $\beta$ increases with radius, it may enter a region where $\beta$ behaves differently and then resume inward migration for instance."
+ Such a situation would be expected to occur in a turbulent magnetized cise in which the large scale field structure changes sullicicntly slowly., Such a situation would be expected to occur in a turbulent magnetized disc in which the large scale field structure changes sufficiently slowly.
+ Unless the magnetic field is above equipartition. it is unstable by the magnetorotational instability (Balbus Haley 1991. 1998 and references therein). the saturated nonlinear outcome of which is MIID turbulence.," Unless the magnetic field is above equipartition, it is unstable by the magnetorotational instability (Balbus Hawley 1991, 1998 and references therein), the saturated nonlinear outcome of which is MHD turbulence."
+ In a turbulent magnetized cisc. the main component of the field is toroidal. as shown by numerical simulations (Llawley. Ganmmic Balbus 1995: Brandenburg et al.," In a turbulent magnetized disc, the main component of the field is toroidal, as shown by numerical simulations (Hawley, Gammie Balbus 1995; Brandenburg et al."
+ 1995)., 1995).
+ Global simulations have also shown that the turbulence saturates at a level corresponding to 3LOO if there is no mean flux or to lower values of 3 if there is a mean flux (Llawley 2001: Steinacker Papaloizou 2002)., Global simulations have also shown that the turbulence saturates at a level corresponding to $\beta \sim 100$ if there is no mean flux or to lower values of $\beta$ if there is a mean flux (Hawley 2001; Steinacker Papaloizou 2002).
+ Also. substantial spatial inhomogeneities are created by racial variations of the Maxwell stress (Lawlev 2001: Steinacker Papaloizou 2002). so that the field may display significant eradients.," Also, substantial spatial inhomogeneities are created by radial variations of the Maxwell stress (Hawley 2001; Steinacker Papaloizou 2002), so that the field may display significant gradients."
+ On the basis of the work presented here. we are Led to speculate that in such a disc the planet would sulfer alternately inward and outward migration. or even no migration at all.," On the basis of the work presented here, we are led to speculate that in such a disc the planet would suffer alternately inward and outward migration, or even no migration at all."
+ lt would then oscillates back and forth in some region of the disc. or undergo some kind of dilfusive migration. either outward or inward depending on the field gradients encountered.," It would then oscillates back and forth in some region of the disc, or undergo some kind of diffusive migration, either outward or inward depending on the field gradients encountered."
+ Note however that in a turbulent cise the magnetic field may. vary locally on timescales shorter than the timescales needed to establish the tvpe of tidal response assumed in this paper., Note however that in a turbulent disc the magnetic field may vary locally on timescales shorter than the timescales needed to establish the type of tidal response assumed in this paper.
+ lt has been pointed out that protoplanctary dises may be ionized enough for the magnetic field to couple to the matter only in their innermost and outermost parts (Cammie 1996: Eromang. TFerquem Balbus 2002).," It has been pointed out that protoplanetary discs may be ionized enough for the magnetic field to couple to the matter only in their innermost and outermost parts (Gammie 1996; Fromang, Terquem Balbus 2002)."
+ A planet forming at around one astronomical unit. where the ionization fraction is very low. would then migrate inward on the timescale calculated by Ware (10560. L997) until it reaches smaller radii where the field is coupled to the matter.," A planet forming at around one astronomical unit, where the ionization fraction is very low, would then migrate inward on the timescale calculated by Ward (1986, 1997) until it reaches smaller radii where the field is coupled to the matter."
+ At this point. there would be a magnetic field inside the planet's orbit but not outside its orbit.," At this point, there would be a magnetic field inside the planet's orbit but not outside its orbit."
+ Ehe torque on the planet would then reverse. and. outward migration would occur.," The torque on the planet would then reverse, and outward migration would occur."
+ However. as soon as the planet would. reenter the nonmagnetic region. inward. migration would resume.," However, as soon as the planet would re–enter the nonmagnetic region, inward migration would resume."
+ Hence. the planet would stall at the border between the magnetic and nonmagnetic regions.," Hence, the planet would stall at the border between the magnetic and nonmagnetic regions."
+ lnwarcl migration in a magnetized disc may then either be very significantly slowed down. occur only on limited. scales. or not occur at all.," Inward migration in a magnetized disc may then either be very significantly slowed down, occur only on limited scales, or not occur at all."
+ Phe planet would then be able to grow to become a terrestrial planet or the core of a giant planet., The planet would then be able to grow to become a terrestrial planet or the core of a giant planet.
+ When the planet becomes massive enough (about 10 earth masses). the interaction with the dise becomes nonlinear. with a eap being opened. up around the planets orbit.," When the planet becomes massive enough (about 10 earth masses), the interaction with the disc becomes nonlinear, with a gap being opened up around the planet's orbit."
+" Because a rather strong torque is exerted in the vicinity of the planet in the presence of a magnetic field. a gap should open up more easily and be ""cleaner in a turbulent magnetic disc than in a"," Because a rather strong torque is exerted in the vicinity of the planet in the presence of a magnetic field, a gap should open up more easily and be 'cleaner' in a turbulent magnetic disc than in a"
+instability is shown to exist in the linear regime.,instability is shown to exist in the linear regime.
+ To test the magnetic field amplification model against observations one needs to determine the saturated magnetic field accurately., To test the magnetic field amplification model against observations one needs to determine the saturated magnetic field accurately.
+ The saturation of the instability cannot be determined in linear theory as the reaction of the instability on the Ch momentun distribution is not automatically included in the linear calculation., The saturation of the instability cannot be determined in linear theory as the reaction of the instability on the CR momemtum distribution is not automatically included in the linear calculation.
+ Although there are numerical simulations of the instability that extend to the nonlinear regime. different: saturation levels have been predicted: (Niemiec. 2008).," Although there are numerical simulations of the instability that extend to the nonlinear regime, different saturation levels have been predicted \citep{netal08,rs08}."
+. In this paper. the instability is discussed in the quasilinear formalism in which the reaction of the instability on the CR. distribution can be included. self-consistentlv.," In this paper, the instability is discussed in the quasilinear formalism in which the reaction of the instability on the CR distribution can be included self-consistently."
+ So. in this formalism one can estimate the saturation analvticallv. with both nonresonant and resonant clüffusion orocesses considered.," So, in this formalism one can estimate the saturation analytically, with both nonresonant and resonant diffusion processes considered."
+ The. treatment of the nonresonant diffusion presented. here is similar to that used. for the irehose instability. (Davidson1972)., The treatment of the nonresonant diffusion presented here is similar to that used for the firehose instability \citep{d72}.
+. We emphasize the major difference between the Cl streaming instability and he firehose instability., We emphasize the major difference between the CR streaming instability and the firehose instability.
+ Phe former is caused by streaming motion and the latter is due to a pressure anisotropy with excess of parallel pressure over. the perpendicular oessure (with respect to the mean magnetic field)., The former is caused by streaming motion and the latter is due to a pressure anisotropy with excess of parallel pressure over the perpendicular pressure (with respect to the mean magnetic field).
+ To some extent. the streaming instability also resembles the Weibel instabilitya nonresonant. purely growing mocdoe driven by anisotropy in the particle distribution (Weibel1959).," To some extent, the streaming instability also resembles the Weibel instability—a nonresonant, purely growing mode driven by anisotropy in the particle distribution \citep{w59}."
+. In Sec 2 the kinetic theory of CR streaming instabilities is discussed with emphasis on the nonresonant instability., In Sec 2 the kinetic theory of CR streaming instabilities is discussed with emphasis on the nonresonant instability.
+ Quasilinear cilfusion driven by the nonresonant instability is discussed. in both the short anc lone wavelength approximations in the nonresonant regime in Sec 3 and in the resonant regime in Sec 4., Quasilinear diffusion driven by the nonresonant instability is discussed in both the short and long wavelength approximations in the nonresonant regime in Sec 3 and in the resonant regime in Sec 4.
+ Application to SN shocks is discussed in Sec 5., Application to SN shocks is discussed in Sec 5.
+ We outline the kinetic theory of CR streaming instabilities including both the usual resonant instability ancl the nonresonant instability and focus particularly on the latter., We outline the kinetic theory of CR streaming instabilities including both the usual resonant instability and the nonresonant instability and focus particularly on the latter.
+ Our treatment builds on other recent. discussions of linear kinetic theory of the CR-induced nonresonant instability (Reville.Wirk&Dully2006:AmatoBlasi2008).," Our treatment builds on other recent discussions of linear kinetic theory of the CR-induced nonresonant instability \citep{retal06,ab08}."
+.. For convenience we assume a single species of CRs with charge q and mass m., For convenience we assume a single species of CRs with charge $q$ and mass $m$ .
+ To model CR streaming at velocity cj; along the mean magnetic field we consider a class of streamingdistributions in momentum space eiven by where and vy are the nondimensional momenta (normalized wy)by me) parallel andperpendicular to the mean magnetic field. respectively. σ is an integer =Lo meu: is the CR number density. e is the ClUs velocity written as a function⋅. of ej and ny. w=(32|2AL/2uc. and For the standard. DSA one has p=2 (Bell1978:Drury1983) and when the nonlinear ellect on DSA is included. p deviates from this canonical value (Kichler1984).," To model CR streaming at velocity $v_{CR}$ along the mean magnetic field we consider a class of streamingdistributions in momentum space given by where $u_\parallel$ and $u_\perp$ are the nondimensional momenta (normalized by $mc$ ) parallel andperpendicular to the mean magnetic field, respectively, $\sigma$ is an integer $\geq1$, $n_{CR}$ is the CR number density, $v$ is the CR's velocity written as a function of $u_\parallel$ and $u_\perp$, $u=(u^2_\perp+u^2_\parallel)^{1/2}$, and For the standard DSA one has $p=2$ \citep{b78,d83} and when the nonlinear effect on DSA is included, $p$ deviates from this canonical value \citep{e84}."
+.. Ehe distribution with σ=1 corresponds to that used in (1986)., The distribution with $\sigma=1$ corresponds to that used in \citet{m86}.
+. The distribution (1)) implies where cosaoyfi , The distribution \ref{eq:fcr}) ) implies where $\cos\alpha\equiv u_\parallel/u$.
+In the second expression in (3)). one chooses (8.0) in place of (yori) as independent variables.," In the second expression in \ref{eq:ncr}) ), one chooses $(u,\alpha)$ in place of $(u_\parallel,u_\perp)$ as independent variables."
+ Ho can be verified. that averaging the parallel velocity ey; over the distribution (1)) gives the streaming velocity. neg. which is independent. of the choice of the xvwameter o.," It can be verified that averaging the parallel velocity $v_\parallel$ over the distribution \ref{eq:fcr}) ) gives the streaming velocity $v_{CR}$, which is independent of the choice of the parameter $\sigma$."
+" The Cl current is then given by dey,=quent.", The CR current is then given by $J_{CR}=qn_{CR}v_{_{\rm CR}}$ .
+ Phe presence of streaming. CRs allects the rvackeround plasma in two wavs. due to their charge density and their. current. density. respectively.," The presence of streaming CRs affects the background plasma in two ways, due to their charge density and their current density, respectively."
+ The background Xdasma must have a charge density and a current. density hat are equal and opposite to those of the ος., The background plasma must have a charge density and a current density that are equal and opposite to those of the CRs.
+ This requires that the electrons and ions (assumed to be protons) rave dilferent charge densities. n.zny. and that they move relative to cach other with streaming velocities οxzey). which are assumed to be along the guiding magnetic field.," This requires that the electrons and ions (assumed to be protons) have different charge densities, $n_e\ne n_p$, and that they move relative to each other with streaming velocities $v_e\ne v_p$, which are assumed to be along the guiding magnetic field."
+ The neutralization conditions require (Achterberg1983) These properties of the background. plasma drive the nonresonant instability attributed to the Clts., The neutralization conditions require \citep{a83} These properties of the background plasma drive the nonresonant instability attributed to the CRs.
+ A formal procedure to derive the dispersion relation involves separating the plasma response tensor into that for a background. plasma. denoted by νεο plus that for the CR component. denoted by νε.," A formal procedure to derive the dispersion relation involves separating the plasma response tensor into that for a background plasma, denoted by $K_{ij}$, plus that for the CR component, denoted by $\Delta K_{ij}$."
+ The background plasma can be regarded as a cold. magnetizecl plasma. while the CR component is described by the distribution. (1)).," The background plasma can be regarded as a cold, magnetized plasma, while the CR component is described by the distribution \ref{eq:fcr}) )."
+ Assume that the gvrofrequeney ofCRs is 2=[q|D/m. where B is the mean background magnetic field.," Assume that the gyrofrequency of CRs is $\Omega=|q|B/m$, where $B$ is the mean background magnetic field."
+ A useful approximation is Ακ1. where Ay is the perpendicular wave number. OQ=Ofs and ~=(1|u7)-7 is the Lorentz factor. of CRs.," A useful approximation is $k_\perp v_\perp/\tilde{\Omega}\ll1$, where $k_\perp$ is the perpendicular wave number, $\tilde{\Omega}=\Omega/\gamma$ and $\gamma=(1+u^2)^{1/2}$ is the Lorentz factor of CRs."
+ The approximation implies thatin the response tensor. only the first evroharmonies terms are important.," The approximation implies thatin the response tensor, only the first gyroharmonics terms are important."
+" For the background. plasma. one assumes (4A«C and the low- approximations.c« Όλα)JApey| and o«x ο, where ©; is the evrolrequeney of ions in the background plasma."," For the background plasma, one assumes $v^2_A\ll c^2$ and the low-frequency approximations,$\omega\ll\Omega_i$ , $\omega\ll |k_\parallel v_\parallel|$ and $\omega\ll
+\tilde{\Omega}$ , where $\Omega_i$ is the gyrofrequency of ions in the background plasma."
+ The background: magnetic field. is assumed. to be along the 3-axis., The background magnetic field is assumed to be along the 3-axis.
+ Since AsusxOFfw2 can be set to ox. only the 2;2 components of the response tensor are relevant.," Since $K_{33}\propto \Omega^2_i/\omega^2$ can be set to $\infty$ , only the $2\times2$ components of the response tensor are relevant."
+ For the background plasma. these components can be written as," For the background plasma, these components can be written as"
+"Ἐνοςvif? spectrum derived for the range around GGHz: The spectral index d£,/dv between GGHz and the K band amounts to <0.39.",$F_\nu \propto \nu^{1/3}$ spectrum derived for the range around GHz: The spectral index ${\rm d}F_\nu / {\rm d}\nu$ between GHz and the K band amounts to $\le 0.39$.
+" We note that this value is very similar to that of other galactic nuclei, like Sgr A* (BDM96), 881 (Reuter and Lesch 1996); 1104 (Jauch and Duschl, in prep.),"," We note that this value is very similar to that of other galactic nuclei, like Sgr A* (BDM96), 81 (Reuter and Lesch 1996); 104 (Jauch and Duschl, in prep.),"
+ where a=1/3., where $\alpha = 1/3$.
+" If NGC 1068 has the same spectral shape as these other galactic nuclei, then a fraction of F7?""** could indeed be contributed from the nucleus of NGC 1068."," If NGC 1068 has the same spectral shape as these other galactic nuclei, then a fraction of $F_{\rm K}^{\rm 30\,mas}$ could indeed be contributed from the nucleus of NGC 1068."
+" However, one has to admit that very little is known about the true nuclear spectrum of 11068 in the intermediate frequency range."," However, one has to admit that very little is known about the true nuclear spectrum of 1068 in the intermediate frequency range."
+" To persue our speculation, we assume — as a working hypothesis — that also between GGHz and the IR range, the spectrum goes like v!/?."," To persue our speculation, we assume – as a working hypothesis – that also between GHz and the IR range, the spectrum goes like $\nu^{1/3}$."
+ We then follow BDM96 and interpret this as optically thin synchrotron radiation of quasi-monoenergetic electrons., We then follow BDM96 and interpret this as optically thin synchrotron radiation of quasi-monoenergetic electrons.
+" The mean electron energy then is fairly well constrained since the maximum of ΓΡ has to be at frequencies above the K band, but not much higher as otherwise the total nuclear flux from the center of 11068 would be too large."," The mean electron energy then is fairly well constrained since the maximum of $F_\nu^{\rm nuc}$ has to be at frequencies above the K band, but not much higher as otherwise the total nuclear flux from the center of 1068 would be too large."
+ The situation is less clear with the SSA frequency., The situation is less clear with the SSA frequency.
+ We cannot rule out that SSA in fact occurs at frequencies even smaller than GGHz., We cannot rule out that SSA in fact occurs at frequencies even smaller than GHz.
+" As a consequence of this, the source radius discussed below is only a lower limit."," As a consequence of this, the source radius discussed below is only a lower limit."
+ For details we refer the reader toBDM96!.., For details we refer the reader to.
+" If we assume that the maximum of P7""* is indeed achieved around um, and that SSA of the source becomes important for frequencies below GGHz, we find as emitting region a homogeneous sphere of radius R~ 2105?cm (~0.7mpc~130AU0.01 mas) with a magnetic field B~ 11G (assumed to be the same everywhere in this region)."," If we assume that the maximum of $F_\nu^{\rm nuc}$ is indeed achieved around $\mu$ m, and that SSA of the source becomes important for frequencies below GHz, we find as emitting region a homogeneous sphere of radius $R \sim 
+2\,10^{15}\,$ cm $\sim 0.7\,{\rm 
+mpc}\sim 130\,{\rm AU\sim 0.01\,mas}$ ) with a magnetic field $B \sim 11\,$ G (assumed to be the same everywhere in this region)."
+" The relativistic electrons have a number density ne1.110? cm-?, a mean energy E~2.7 GeV and a width of the energy distribution AE/E~1."," The relativistic electrons have a number density $n_{\rm e} \sim 1.1\,10^3\,{\rm cm}^{-3}$ , a mean energy $E \sim 2.7\,$ GeV and a width of the energy distribution $\Delta E / 
+E \sim 1$."
+ In reffigspect wwe show a comparison of the observed fluxes of 11068 core and our model spectrum using the above parameters., In \\ref{figspect} we show a comparison of the observed fluxes of 1068 core and our model spectrum using the above parameters.
+" If our speculation applies, it turns out that the main difference between 11068 and other galactic centers analysed on the basis of the same interpretation (Sgr A*: BDM96; 881: Reuter and Lesch 1996; 1104: Jauch and Duschl, in prep.)"," If our speculation applies, it turns out that the main difference between 1068 and other galactic centers analysed on the basis of the same interpretation (Sgr A*: BDM96; 81: Reuter and Lesch 1996; 104: Jauch and Duschl, in prep.)"
+ are the source radius and - especially - the energy of the relativistic electrons., are the source radius and - especially - the energy of the relativistic electrons.
+" The above size of mmas means that our resolved mmas object is not the synchrotron source itself but rather a larger object, most likely the nuclear torus and/or a circumnuclear scattering halo."," The above size of mas means that our resolved mas object is not the synchrotron source itself but rather a larger object, most likely the nuclear torus and/or a circumnuclear scattering halo."
+ We have resolved a compact source with a diameter of dGaussv30 mas in the core of 11068.," We have resolved a compact source with a diameter of $d_{\rm Gauss}\sim 
+30\,$ mas in the core of 1068."
+ This object is most likely a nuclear torus and/or a circumnuclear scattering halo., This object is most likely a nuclear torus and/or a circumnuclear scattering halo.
+ Part of the radiation from the central mmas may be light from the nucleus scattered in the halo., Part of the radiation from the central mas may be light from the nucleus scattered in the halo.
+" Under this assumption, we were able to determine physical parameters of the nucleus and compare them to other galactic centers, active and non-active ones."," Under this assumption, we were able to determine physical parameters of the nucleus and compare them to other galactic centers, active and non-active ones."
+ One then is tempted to speculate that the higher efficiency of the acceleration mechanism for the electrons may be the true difference between an active galactic center and a normalone., One then is tempted to speculate that the higher efficiency of the acceleration mechanism for the electrons may be the true difference between an active galactic center and a normalone.
+emission was found in the radio survey of Filipovic et al. (,emission was found in the radio survey of Filipović et al. (
+1998) and they are candidates for AGNs.,1998) and they are candidates for AGNs.
+ The source No., The source No.
+ 103 (RX J0054.9-7226) is identified with the XTE J0055-724 = 1SAX J0054.9-7226 source (Marshall et al., 103 (RX J0054.9-7226) is identified with the XTE J0055-724 = 1SAX J0054.9-7226 source (Marshall et al.
+" 1998, Israel 1998)."," 1998, Israel 1998)."
+ The detection of pulsations with a period of 59 s withBeppoSAX confirm the X-ray binary nature of this source., The detection of pulsations with a period of 59 s with confirm the X-ray binary nature of this source.
+ Stevens et al. (, Stevens et al. (
+"1998) identified early type emission-line stars through colour indices and Ha emission for the sources with catalogue indices 3, 69, 103, and 158.","1998) identified early type emission-line stars through colour indices and $\alpha$ emission for the sources with catalogue indices 3, 69, 103, and 158."
+ The source No., The source No.
+ 153 (RX J0100.7-7211) was considered to be consistent with a background AGN shining through the SMC bulge (Paper I)., 153 (RX J0100.7-7211) was considered to be consistent with a background AGN shining through the SMC bulge (Paper I).
+ Sources No., Sources No.
+ 157 and 160 were found to coincide with detector struts and were rejected accordingly (Paper I)., 157 and 160 were found to coincide with detector struts and were rejected accordingly (Paper I).
+ Weak hard X-ray binaries are an interesting class of objects as they have been predicted to exist and their number is expected to be large especially in galaxies of low metallicity like the SMC., Weak hard X-ray binaries are an interesting class of objects as they have been predicted to exist and their number is expected to be large especially in galaxies of low metallicity like the SMC.
+ In previous work (Bruhweiler et al., In previous work (Bruhweiler et al.
+ 1987; Wang Wu 1992) candidates for such sources have been found and either classified as low luminosity Be systems or as background objects., 1987; Wang Wu 1992) candidates for such sources have been found and either classified as low luminosity Be systems or as background objects.
+" Here, we are searching for candidates of this class by applying the same selection criteria as for the strong (or higher luminosity) hard X-ray binaries as outlined in Paper I:0.5,, and extent likelihood<50."," Here, we are searching for candidates of this class by applying the same selection criteria as for the strong (or higher luminosity) hard X-ray binaries as outlined in Paper I:, and extent likelihood."
+". The only difference is to select objects with count rates ofs-!,, i.e. with luminosities below assuming a standard spectral model for the source flux (Paper I)."," The only difference is to select objects with count rates of, i.e. with luminosities below assuming a standard spectral model for the source flux (Paper I)."
+ There are 60 such objects and we tentatively classify these sources as class=Bw., There are 60 such objects and we tentatively classify these sources as class=Bw.
+ This class is a substantial fraction (25%)) of the total catalogue entries and turns out to be the class with most members., This class is a substantial fraction ) of the total catalogue entries and turns out to be the class with most members.
+ We find 15 of these objects which coincide with ddetections., We find 15 of these objects which coincide with detections.
+ This may reflect that we are considerably, This may reflect that we are considerably
+for LIIS 147 and LIIS 542 are taken from BLR. and for LIIS 4033 from Dahnοἱal.(2004).,"for LHS 147 and LHS 542 are taken from BLR, and for LHS 4033 from \citet{dahn04}."
+. Out of the 32 objects listed in Table 1. 22 have J///¥ measurements. 5 only have J and IH. while 2 have no infrared data.," Out of the 32 objects listed in Table 1, 22 have $JHK$ measurements, 8 only have $J$ and $H$, while 2 have no infrared data."
+ Also reported in Table 1 ave the infrared. photometric uncertainties (in parentheses) aud (he number of independent observations., Also reported in Table 1 are the infrared photometric uncertainties (in parentheses) and the number of independent observations.
+ The model atmospheres used in (his analysis are described al length in Bergeronetal.BRLandBL) with the collision-induced opacities Hom molecular hydrogen updated [rom the work of Jorgensenetal.(2000) and Borvsowetal.(2001)., The model atmospheres used in this analysis are described at length in \citet[][see also BRL and BLR]{bsw95} with the collision-induced opacities from molecular hydrogen updated from the work of \citet{jorgensen} and \citet{borysow01}.
+. These models are in local thermodynamic equilibrium. they allow energv transport by convection. aud they can be caleulated with arbitrary mixed hydrogen ancl helium compositions.," These models are in local thermodynamic equilibrium, they allow energy transport by convection, and they can be calculated with arbitrary mixed hydrogen and helium compositions."
+ Synthetic are obtained using the procedure outlined in Bergeronοἱal.(1995b) but with the new Vega [Iuxes taken from Bohlin&Gilliland(2004) and the Vega magnitudes from Table Al of Besselletal.(1998)., Synthetic are obtained using the procedure outlined in \citet{bwb95} but with the new Vega fluxes taken from \citet{bohlin04} and the Vega magnitudes from Table A1 of \citet{bessell98}.
+". Similarly. in order to compare the photometric observations with the model atmosphere predictions. we convert (see also BRL) the optical and infrared magnitudes mz into observed [fluxes averaged over (he transmission function S4,CLÀX)HEN usinge the followinge equation where is the averaged observed fIux received at Earth."," Similarly, in order to compare the photometric observations with the model atmosphere predictions, we convert (see also BRL) the optical and infrared magnitudes $m$ into observed fluxes averaged over the transmission function $S_m(\lambda)$ using the following equation where is the averaged observed flux received at Earth."
+" The transmission functions $,(À) are taken from Bessell(1990). flor the DVRJ fillers on the Johnson-Cousins photometric system. ancl from Bessell&Brett(1988) [or the 7L fillers on the Johnson-Glass svstem."," The transmission functions $S_m(\lambda)$ are taken from \citet{bessell90} for the $BVRI$ filters on the Johnson-Cousins photometric system, and from \citet{bessell88} for the $JHK$ filters on the Johnson-Glass system."
+" The constants €,, loreach passband using the new fluxes and zero points [or Vega are ej=—20.4761. cy=—21.00708. eg,=—21.6300. e;=—22.3480. e,=--ο411. ej;=—24.8281. and Cxdy=—25.9877."," The constants $c_m$ foreach passband using the new fluxes and zero points for Vega are $c_B=-20.4761$, $c_V=-21.0798$, $c_R=-21.6300$, $c_I=-22.3480$, $c_J=-23.7417$, $c_H=-24.8387$, and $c_K=-25.9877$."
+ These constants differ slightlvo [rom those used by BRL and BLR. whichwere based on older Vega fluxes.," These constants differ slightly from those used by BRL and BLR, whichwere based on older Vega fluxes."
+" Note also that with this new calibration. the +0.05 mag correction determined empirically ancl applied by BRL to the J. 1, and A constants is nol required here (see 5.2.1 of BRL)."," Note also that with this new calibration, the +0.05 mag correction determined empirically and applied by BRL to the $J$ , $H$ , and $K$ constants is not required here (see 5.2.1 of BRL)."
+"same model outlined in section 4.1 applied to Ser A. it would predict an outer radius of the outflow of about may22.510? Schwarzschild radii (~ 0.0257 at the distance of the galactic center) and an inner accretion rate ALAa)Alas""4.10"" AL.f yr. also consistent with the gas density in the inner accretion flow implied by linear polarization measurements (Aitkenetal.2000:Boweretal.2003: 2005).","same model outlined in section \ref{sec:model} applied to Sgr $^{*}$, it would predict an outer radius of the outflow of about $r_{\rm out}\approx 2.5 \times
+10^{3}$ Schwarzschild radii $\sim 0.025$ ” at the distance of the galactic center) and an inner accretion rate $\dot M_{\rm a}(r_{\rm
+in})=\dot M_{\rm out} \xi^{p} \simeq 4 \times 10^{-9}$ $M_{\odot}/$ yr, also consistent with the gas density in the inner accretion flow implied by linear polarization measurements \cite{aitken:00,bower:03,bower:05}."
+". Altogether. this would indicate that the SMBH at the galactic center is the source of about ο1077 ergs ! of mechanical power (similar to what estimated within jet model fits to Ser A* SED. see e.g. Falcke Biermann 1999, Falcke Markoff 2000). or about 1.5 supernovae every 10 years."," Altogether, this would indicate that the SMBH at the galactic center is the source of about $5\times
+10^{38}$ ergs $^{-1}$ of mechanical power (similar to what estimated within jet model fits to Sgr $^{*}$ SED, see e.g. Falcke Biermann 1999, Falcke Markoff 2000), or about 1.5 supernovae every $^5$ years."
+ Such a mechanical power input into the galactic center could play a significant role in the production of the TeV ~-rays recently observed by the HESS (High Energy Stereoscopic System) collaboration (Aharonianetal.2004: 2004)... as well as in the heating of the hot (>S keV) diffused plasma detected by (Munoetal.2004).," Such a mechanical power input into the galactic center could play a significant role in the production of the TeV $\gamma$ -rays recently observed by the HESS (High Energy Stereoscopic System) collaboration \cite{aharonian:04,atoyan:04}, as well as in the heating of the hot $> 8$ keV) diffused plasma detected by \cite{muno:04}."
+. In section 3.. we have concentrated our attention on the assessment of the effects of relativistic beaming on the measured slope of the Lyin-Lp correlation.," In section \ref{sec:rad}, we have concentrated our attention on the assessment of the effects of relativistic beaming on the measured slope of the $L_{\rm kin}$ $L_{\rm R}$ correlation."
+ However. one should expect a second source of scatter in any correlation between kinetic power and nuclear luminosity tin any waveband).," However, one should expect a second source of scatter in any correlation between kinetic power and nuclear luminosity (in any waveband)."
+ The large scale power is an average over the typical age of the cavities and bubbles observed in the atmosphere of the galaxies in our sample., The large scale power is an average over the typical age of the cavities and bubbles observed in the X-ray atmosphere of the galaxies in our sample.
+ Such an age may be assumed to be of the order of either the buoyant rise time. or the sound crossing time. or the refill time of the radio lobes (Birzan 2004).," Such an age may be assumed to be of the order of either the buoyant rise time, or the sound crossing time, or the refill time of the radio lobes \cite{birzan:04}."
+. These estimates typically differ by about a factor of 2. and. for the sample in question. lie in the range {μον10τοῦ yeurs.," These estimates typically differ by about a factor of 2, and, for the sample in question, lie in the range $t_{\rm
+age} \sim 10^7 - 10^8$ years."
+ On the other hand. the core power is variable on time scale much shorter than that.," On the other hand, the core power is variable on time scale much shorter than that."
+ The bias introduced by this fact can be quantitied as follows., The bias introduced by this fact can be quantified as follows.
+ First of all. we notice that AGN X-ray lighteurves. where most of the accretion poweremerges. show a characteristic rms-flux relation which implies they have a formally non-linear. exponential form (Uttley.McHardy&Vaughan2005).. and the luminosity follows a log-normal distribution (Nipoti&Binney2005).," First of all, we notice that AGN X-ray lightcurves, where most of the accretion power emerges, show a characteristic rms-flux relation which implies they have a formally non-linear, exponential form \cite{uttley:05}, and the luminosity follows a log-normal distribution \cite{nipoti:05}."
+.. Let us. for the sake of simplicity. define as L(/) the luminosity of the AGN. be it bolometric. Kinetic or radio. under the implicit assumption that very close to the central engine jet and accretion are strongly coupled. so that all of them follow a log- distribution.," Let us, for the sake of simplicity, define as $L(t)$ the luminosity of the AGN, be it bolometric, kinetic or radio, under the implicit assumption that very close to the central engine jet and accretion are strongly coupled, so that all of them follow a log-normal distribution."
+ Thus. logL is normally distributed with mean jii and variance στ.," Thus, $\log L$ is normally distributed with mean $\mu_l$ and variance $\sigma_l^2$."
+" The measured kinetic power Z1, discussed in this work being a long-term time average. it should be determined by the mean of the log-normally distributed £L. ie. (£)=ppexptr|07/2)."," The measured kinetic power $L_{\rm kin}$ discussed in this work being a long-term time average, it should be determined by the mean of the log-normally distributed $L$, i.e. $\langle L \rangle=\mu=\exp{(\mu_l^2+\sigma_l^2/2)}$ ."
+ Because the log-normal distribution is positively skewed. the mode nz of the distribution of £L. i.e. the most likely value of a measurement of it. Lou. 18 than the mean: Lasom=expí(qu07).," Because the log-normal distribution is positively skewed, the mode $m$ of the distribution of $L$, i.e. the most likely value of a measurement of it, $L_{\rm obs}$, is than the mean: $L_{\rm
+obs}=m=\exp{(\mu_l-\sigma_l^2)}$."
+ Therefore the most likely value of the ratio of the mean to the observed luminosity is given by (Nipoti Binney 2005: Uttley et al., Therefore the most likely value of the ratio of the mean to the observed luminosity is given by (Nipoti Binney 2005; Uttley et al.
+ 2005): where we have introduced the rms (fractional) variability of the observed lightcurve σue=oxpayl.an easily measurable quantity.," 2005): where we have introduced the rms (fractional) variability of the observed lightcurve $\sigma_{\rm rms}^2=\exp{\sigma_l^2}-1$, an easily measurable quantity."
+ The higher the rms variability of a lighteurve. the higher is the probability that an instantaneous measurement of the luminosity yields a value smaller than (£5 (Nipoti& 2005). and also the higher the most likely ratio between the two values.," The higher the rms variability of a lightcurve, the higher is the probability that an instantaneous measurement of the luminosity yields a value smaller than $\langle L \rangle$ \cite{nipoti:05}, and also the higher the most likely ratio between the two values."
+ Before discussing what are the appropriate values of σημ.2 to be used in Eq. (14)).," Before discussing what are the appropriate values of $\sigma_{\rm
+rms}^2$ to be used in Eq. \ref{eq:rms}) ),"
+ we also note that the observed slope of the £L5- Lai; correlation may deviate from unity for a log-normal distribution. thus skewing any observed correlation between mean and instantaneous power. like those we have discussed so far.," we also note that the observed slope of the $\langle L \rangle$ $L_{\rm obs}$ correlation may deviate from unity for a log-normal distribution, thus skewing any observed correlation between mean and instantaneous power, like those we have discussed so far."
+ It is easy to show that from Eq. (1+)), It is easy to show that from Eq. \ref{eq:rms}) )
+ we obtain: Log-normal AGN variability may thus skew the observed relation between jet average kinetic power and instantaneous core luminosity much in the same way relativistic beaming does., we obtain: Log-normal AGN variability may thus skew the observed relation between jet average kinetic power and instantaneous core luminosity much in the same way relativistic beaming does.
+ Inspection of Eqs. (14)), Inspection of Eqs. \ref{eq:rms}) )
+" and (15)) suggests that the measured slopes of any correlation of the (average) kinetic vs. core power will be substantially affected by variability if. and only if. both στmas, and logevus/0logLa are at least of the order of unity."," and \ref{eq:rms-slope}) ) suggests that the measured slopes of any correlation of the (average) kinetic vs, core power will be substantially affected by variability if, and only if, both $\sigma_{\rm rms}^2$ and $\partial \log \sigma_{\rm rms}^2 /\partial
+\log L_{\rm obs}$ are at least of the order of unity."
+ Unfortunately. very little is known observationally about the variability amplitude of AGN on very long timescales. especially or AGN of low luminosity as those considered here.," Unfortunately, very little is known observationally about the variability amplitude of AGN on very long timescales, especially for AGN of low luminosity as those considered here."
+ Brighter AGN (Seyferts}). on shorter timescales (1-10 years). have indeed rms variability amplitudes that rise steeply with decreasing luminosity. rom about a5.c107 for LxeLol up to ayy.3«LO1 or Lsez1077 (Nandraetal.1997:Markowitz&Edelson2001).," Brighter AGN (Seyferts), on shorter timescales (1-10 years), have indeed rms variability amplitudes that rise steeply with decreasing luminosity, from about $\sigma_{\rm rms}^2 \approx 10^{-2}$ for $L_{\rm X} \approx
+10^{44}$ up to $\sigma_{\rm rms}^2 \approx 10^{-1}$ for $L_{\rm X} \approx
+10^{42}$ \cite{nandra:97,markowitz:01}."
+ However. no evidence is vet found of such a trend continuing down o lower luminosities.," However, no evidence is yet found of such a trend continuing down to lower luminosities."
+" On the contrary. suggestions have been made hat 22, may flattens out at lower Lx (Papadakis2004:Paolillo at values of a few times + for the typical X-ray uminosity of the objects in our sample."," On the contrary, suggestions have been made that $\sigma_{\rm rms}^2$ may flattens out at lower $L_{\rm X}$ \cite{papadakis:04,paolillo:04} at values of a few times $^{-1}$ for the typical X-ray luminosity of the objects in our sample."
+ This would imply that the observed correlations between large scale and core powers are not skewed by variability bias., This would imply that the observed correlations between large scale and core powers are not skewed by variability bias.
+ However. we should also consider the possibility that. if," However, we should also consider the possibility that, if"
+causality. however. and the link can be indirect.,"causality, however, and the link can be indirect."
+ For instance. 7 recently pointed out the strong correlation between excess racidus ancl the incident stellar Hux. supporting the scenario of ?..," For instance, \citet{fre07} recently pointed out the strong correlation between excess radius and the incident stellar flux, supporting the scenario of \citet{gui06}."
+ Incident (ux. like tidal effects also scales with eAt.," Incident flux, like tidal effects also scales with $a/R$."
+ Fig., Fig.
+ 2 shows that. in the present sample. the correlation of excess radius is stronger with aff than with the strength of tidal effects.," \ref{tides} shows that, in the present sample, the correlation of excess radius is stronger with $a/R$ than with the strength of tidal effects."
+ The present sample of transiting planet is not sullicient. to distinguish. between the two types of explanations without further modelling. but it does suggest wt tidal effects are not the main factor in determining the size of close-in gas giant planets and that the intensity of vw incident [ux is a better candidate.," The present sample of transiting planet is not sufficient to distinguish between the two types of explanations without further modelling, but it does suggest that tidal effects are not the main factor in determining the size of close-in gas giant planets and that the intensity of the incident flux is a better candidate."
+ Ifthe tentative indications provided by the present sample of transiting planets are confirmed. the following overall picture - close-in planets end up the formation phase on orbits with a wide distribution of - tidal interaction with the star circularise the planetary orbit. beginning around {οσαfA)=1.5 for 2-Aly planets. then circularizes it.," If the tentative indications provided by the present sample of transiting planets are confirmed, the following overall picture - close-in planets end up the formation phase on orbits with a wide distribution of - tidal interaction with the star circularise the planetary orbit, beginning around $log(a/R)=1.5$ for $_{\rm J}$ planets, then circularizes it."
+ Early signs of circularisation are visible alreacly at larger - for planets lighter than Jupiter. the circularisation is not accompanied. hy significant orbital decay. ancl effect. on the - for planets with mass comparable to Jupiter or slightly heavier. circularisation is accompanied by significant orbital decay. leading to a mass-periocd relation. ancl possible destruction for Al~1.2A. - still heavier planet produce a marked. stellar excess rotation. and corresponding orbital decay. during. the process of tidal evolution and cireularisation.," Early signs of circularisation are visible already at larger - for planets lighter than Jupiter, the circularisation is not accompanied by significant orbital decay and effect on the - for planets with mass comparable to Jupiter or slightly heavier, circularisation is accompanied by significant orbital decay, leading to a mass-period relation, and possible destruction for $M\sim 1-2 M_{\rm J}$ - still heavier planet produce a marked stellar excess rotation, and corresponding orbital decay, during the process of tidal evolution and circularisation."
+ LW their initial position exceeds a critical value. they may reach a spin-orbit svnchronous state with the star. - às à result of this evolution. the orbital characteristics of very close-in planets. (Ποι planets) are. profounclly modified by tidal interactions. including the orbital distance ancl period.," If their initial position exceeds a critical value, they may reach a spin-orbit synchronous state with the star, - as a result of this evolution, the orbital characteristics of very close-in planets (“hot” planets) are profoundly modified by tidal interactions, including the orbital distance and period."
+" Planets lighter than 2 AZ, ancl closer than a few ανν are ""sheperded along a mass-period trend. controlled. by tidal angular momentum exchange ancl possibly tidal inflation of the planet."," Planets lighter than 2 $M_J$ and closer than a few days are ""sheperded"" along a mass-period trend controlled by tidal angular momentum exchange and possibly tidal inflation of the planet."
+ The period-mass relation of close-in planets cannot be used at face value to constrain formation and migration If this interpretation of the data ds. correct. we can roughly divide the mass-period. plane for. exoplanet according to their sensitivity to tidal evolution.," The period-mass relation of close-in planets cannot be used at face value to constrain formation and migration If this interpretation of the data is correct, we can roughly divide the mass-period plane for exoplanet according to their sensitivity to tidal evolution."
+ Figure 6 sums up the observable elfects of tidal evolution as inferred here from transiting planets., Figure \ref{fig5} sums up the observable effects of tidal evolution as inferred here from transiting planets.
+ Phe possible transition zone where orbital circularisation ancl orbital decay occur is indicated. as well as the Roche limit (with. reasonable assumptions on the mass-radius relation of planets).," The possible transition zone where orbital circularisation and orbital decay occur is indicated, as well as the Roche limit (with reasonable assumptions on the mass-radius relation of planets)."
+ Planets at the weak end. of the scale. of tidal effects can conserve eccentric and. misaligned orbits., Planets at the weak end of the scale of tidal effects can conserve eccentric and misaligned orbits.
+ Jupiter-mass planets placed near the 3-davs limit will substantially alfect the rotation of their parent star. moving closer in the process by orbital decay. down to 1-2 αν periods.," Jupiter-mass planets placed near the 3-days limit will substantially affect the rotation of their parent star, moving closer in the process by orbital decay, down to 1-2 day periods."
+ Planets near the 2-3 Jupiter mass range will undergo strong orbital decay. [acing potential destruction if they start on a close orbit.," Planets near the 2-3 Jupiter mass range will undergo strong orbital decay, facing potential destruction if they start on a close orbit."
+ Large planets or brown cdwarfs in the 5-15.r Jupiter-mass range will spin up their host star substantially if they orbi close enough. all the way to synchronous rotation.," Large planets or brown dwarfs in the 5-15 Jupiter-mass range will spin up their host star substantially if they orbit close enough, all the way to synchronous rotation."
+ Lighter planets will not allect their star detectably., Lighter planets will not affect their star detectably.
+ Lf they are close enough. they will spiral inwards to the Roche limit in a shor time.," If they are close enough, they will spiral inwards to the Roche limit in a short time."
+ As more objects. Dill. the diagram. this eloba interpretation can be tested and refined.," As more objects fill the diagram, this global interpretation can be tested and refined."
+ In Figure 6.. the two known transiting stars with the lowest masses are also plotted. OGLE-PR-122 (7) anc OGLE-TPR-123 (2)...," In Figure \ref{fig5}, the two known transiting stars with the lowest masses are also plotted, OGLE-TR-122 \citep{pon05} and OGLE-TR-123 \citep{pon06}."
+ Phese objects show that the connection with stellar binaries is coherent. with the first object having an orbit still eccentric but clear evidence of excess rotation of the primary. and the second a circular and svachronous orbit.," These objects show that the connection with stellar binaries is coherent, with the first object having an orbit still eccentric but clear evidence of excess rotation of the primary, and the second a circular and synchronous orbit."
+ Figure 6. shows. for HD 189733. 7 Boo and OCLE-TR-123. the initial orbital distance assuming that all the angular momentum of the star is put back in the orbit.," Figure \ref{fig5} shows, for HD 189733, $\tau$ Boo and OGLE-TR-123, the initial orbital distance assuming that all the angular momentum of the star is put back in the orbit."
+ The initial positions will be placed near the limit of tidal spin-up. which provides accitional support to this interpretation.," The initial positions will be placed near the limit of tidal spin-up, which provides additional support to this interpretation."
+ It also suggests one of the reasons for the inaceuraey. of tidal timescales calculated from present. orbital parameters: the initial orbit could have been very dillerent., It also suggests one of the reasons for the inaccuracy of tidal timescales calculated from present orbital parameters: the initial orbit could have been very different.
+ For OGLE-PR- for instance. the initial position is above the frame of the plot.," For OGLE-TR-122 for instance, the initial position is above the frame of the plot."
+(in particular massive halos) competes with the increase in the comoving size of the error boxes.,(in particular massive halos) competes with the increase in the comoving size of the error boxes.
+ Our simple calculations show that in the scenario of CC10.. resolved PLA sources with AL=10AZ. and zx0.5 are likely to have at worst dozens of interlopers in the error box.," Our simple calculations show that in the scenario of \citetalias{CorCor10}, resolved PTA sources with $M\ga 10^{9}\Msol$ and $z\ltsim 0.5$ are likely to have at worst dozens of interlopers in the error box."
+ With this low number. one could conceivably perform follow-up observations of cach individual candidate.," With this low number, one could conceivably perform follow-up observations of each individual candidate."
+ 1f. on the other hand. luminosity distances to the source cannot be determined. this number increases to 107. suggesting that it will become extremely cillicult to clectromagnetically identify the source in the absence of an obvious. tell-tale EAL signature.," If, on the other hand, luminosity distances to the source cannot be determined, this number increases to $\sim 10^{3}$, suggesting that it will become extremely difficult to electromagnetically identify the source in the absence of an obvious, tell-tale EM signature."
+ In practice. the number of interloping galaxies may be somewhat larger than the value computed by equations 11- 16.," In practice, the number of interloping galaxies may be somewhat larger than the value computed by equations \ref{eq:NhSV}$ $-$ \ref{eq:NACC}."
+ hehalomeassofanggivencandidatehostsystemmwillnoll lnowpbu)gd dod edtlicadilhy appeb CMauinau Cheat) relation will lower the minimum halo mass threshold. for candidacy.," The halo mass of any given candidate host system will not be known a priori, and the intrinsic scatter in the $M_{\rm SMBH}-M_{\rm halo}$ $M_{\rm SMBH}-\sigma_{\rm
+ host}$ ) relation will lower the minimum halo mass threshold for candidacy."
+ On the other hand. the simple calculations presented. here. co not. consider detailed: demographic properties of resolved. PLA) sources and. plausible hosts. such as the presence of a nuclear stellar core (Makino1997:Hawvindranath.Ho&Filippenko2002:Milosavljeviéetal.2002:Volonteri.Macau&Llaardt2003) or galaxy morphology.," On the other hand, the simple calculations presented here do not consider detailed demographic properties of resolved PTA sources and plausible hosts, such as the presence of a nuclear stellar core \citep[]{Makino97, Ravind+02, Milos+02, VMH03} or galaxy morphology."
+ Including such factors in the analysis will narrow the field of candidate hosts, Including such factors in the analysis will narrow the field of candidate hosts.
+" As we argue in 3.3. candidate ACN counterparts may be further vetted by examining their UV. and. X-ray emission for features indicative of a central SAIBLI binary (see also Sesanaetal.2011. for an in-depth: discussion of possible high-energy signatures for. pre-decoupling Le, low>by PPA sources)."," As we argue in \ref{subsec:emission}, candidate AGN counterparts may be further vetted by examining their UV and X-ray emission for features indicative of a central SMBH binary (see also \citealt{Sesana+11} for an in-depth discussion of possible high-energy signatures for pre-decoupling — i.e., $t_{\rm GW}>t_{\nu}$ — PTA sources)."
+ Lo addition. PLA sources are sullicienthy nearby that it should be possible to observe an interloping AGN together with its host. galaxy.," In addition, PTA sources are sufficiently nearby that it should be possible to observe an interloping AGN together with its host galaxy."
+ Lt should therefore be possible to combine the ACN emission. the ealaxy luminosity and the inferred. SMDII mass to. check candidate counterparts.," It should therefore be possible to combine the AGN emission, the galaxy luminosity and the inferred SMBH mass to cross-check candidate counterparts."
+ Alotivated by the results of the previous section that the number of plausible host. galaxies in the PTA error box may be tractable for follow-up LEAT searches. we next mocel he EAL emission. properties of SAIBLL binaries detectable ον P'PAXs.," Motivated by the results of the previous section that the number of plausible host galaxies in the PTA error box may be tractable for follow-up EM searches, we next model the EM emission properties of SMBH binaries detectable by PTAs."
+ We focus our attention on SMDLIILD binaries that are undergoing luminous accretion. as these are the most omising Class of objects for LEAL identification.," We focus our attention on SMBH binaries that are undergoing luminous accretion, as these are the most promising class of objects for EM identification."
+" Normalizing the binary mass AZ and rest-frame period P to the typical orders of magnitude expected. of resolved ""PA sources. AL=10""AZ.Mo and P=loveF4. we write he semi-major axis for the source binary as Binaries detectable by P""PAs have long overcome the so- “final parsec problem."," Normalizing the binary mass $M$ and rest-frame period $P$ to the typical orders of magnitude expected of resolved PTA sources, $M=10^{9}\Msol M_{9}$ and $P=1\yr~P_{1}$, we write the semi-major axis for the source binary as Binaries detectable by PTAs have long overcome the so-called “final parsec” problem."
+ The rest-frame time to merger for a binary with mass AZ and semi-major axis e. driven by CN emission alone. is (Peters.1964).," The rest-frame time to merger for a binary with mass $M$ and semi-major axis $a$, driven by GW emission alone, is \citep{Peters64}."
+".. Beeause tvpical resolved: sources have AMD=0g773. we normalize the svnimetric mass ralio g—(MaS/MQ|AdsALJ to the value ae—WAL/M,=0.25)0.16."," Because typical resolved sources have $M/\Mch=\eta^{-3/5}\sim 3$, we normalize the symmetric mass ratio $\eta\equiv (M_{2}/M_{1})/[1+M_{2}/M_{1}]^{2}$ to the value $\eta_{1:4}\equiv\eta(M_{2}/M_{1}=0.25)=0.16$ ."
+" Note that our ad hoc translation between AL and LM is not very sensitive to the value of q: the ratio AZ/,ME varies by less than a [actor of two in the Aste ", Note that our ad hoc translation between $M$ and $\Mch$ is not very sensitive to the value of $q$; the ratio $M/\Mch$ varies by less than a factor of two in the range $0.1\le M_{2}/M_{1} \le 1$.
+shad dunn equation. LS corresponds to binaries in circular orbits. with eccentric orbits merging Faster.," The upper bound in equation \ref{eq:tmerge} corresponds to binaries in circular orbits, with eccentric orbits merging faster."
+ Recent work has shown that binarics nw have eccentricities as high as 0.6 at decoupling (Iltoedigetal.2011: see also Armitage&Natarajan2005:Cuadraetal.2009).," Recent work has shown that binaries may have eccentricities as high as $\sim 0.6$ at decoupling \citealt{Roedig+11}; see also \citealt{AN05, Cuadra+09}) )."
+ Thus. typical P'EAX-resolved: sources will coalesce on scales of 10° vears.," Thus, typical PTA-resolved sources will coalesce on scales of $\sim 10^{3}$ years."
+ However. exceptionally compact sources will coalesce on scales of several vears: for example. a binary with P?=0.1vr approximately the lowest binary period that is expected to be observable with IAAS will merge in Fuere~4vr.," However, exceptionally compact sources will coalesce on scales of several years; for example, a binary with $P=0.1\yr$ — approximately the lowest binary period that is expected to be observable with PTAs — will merge in $t_{\rm merge}\sim 4\yr$."
+ The tidal torques of the compact SML binary provide a particularly promising mechanism for producing a tell-tale observable feature., The tidal torques of the compact SMBH binary provide a particularly promising mechanism for producing a tell-tale observable feature.
+ Theoretical caleulations (Goldreichjevié2008:6πανταetal.2009:Chang2010). robustly predict that in. geometrically thin circumbinary accretion discs. binary torques can open an annular. low-clensity gap around the orbit of the secondary.," Theoretical calculations \citep{GT80, Artymo+91, ArtLub94, AN02,
+ Bate+03, Hayasaki+07,MM08, Cuadra+09, Chang+10} robustly predict that in geometrically thin circumbinary accretion discs, binary torques can open an annular, low-density gap around the orbit of the secondary."
+ Ehe σας inside the gap accretes onto the individual SALBIIs while the gas outside is pushed. outward by the tidal torques., The gas inside the gap accretes onto the individual SMBHs while the gas outside is pushed outward by the tidal torques.
+ “Phe binary's tidal torques transfer orbital angular momentum into the outer disc. causing the binarv orbit to shrink &radually while maintaining a roughly axisvnunetrics cireumbinary. gap.," The binary's tidal torques transfer orbital angular momentum into the outer disc, causing the binary's orbit to shrink gradually while maintaining a roughly axisymmetric circumbinary gap."
+ The gap opens near the resonance radius /?237/702.08. CXrtvimowiczetal.1991)., The gap opens near the resonance radius $R\approx 3^{2/3}a \approx 2.08 a$ \citep{Artymo+91}.
+. Numerical. simulations of thin circumbinary disces (see refs., Numerical simulations of thin circumbinary discs (see refs.
+ in above paragraph) produce gaps with an azimuthally averaged radius of 1.5.3 times the binary semimajor axis., in above paragraph) produce gaps with an azimuthally averaged radius of $1.5-3$ times the binary semimajor axis.
+ Phe exact size ancl shape of the gap is not easily characterized: the geometry depends on the binary masses and orbital cecentricity. as well as the ellicieney of angular momentum transport within the disc.," The exact size and shape of the gap is not easily characterized; the geometry depends on the binary masses and orbital eccentricity, as well as the efficiency of angular momentum transport within the disc."
+ Following Milosavljevió&Phinney(2005)... we parametrize the size of the gap as Ry=2Àe. where A~| is a dimensionless parameter.," Following \cite{MP05}, we parametrize the size of the gap as $R_{\lambda}\equiv 2\lambda a$, where $\lambda\sim 1$ is a dimensionless parameter."
+" We are interested in cireumbinary discs that are truncated inside Ay~20034,2Hp""GMAO (equation 17))."," We are interested in circumbinary discs that are truncated inside $R_{\lambda}\sim 200 M_{9}^{-2/3}P_{1}^{2/3}
+GM/c^{2}$ (equation\ref{eq:aPTA}) )."
+ Below. we model surface density profiles and thermal emission spectra of such disces. as well as the thermal emission due to leakage of gas into the cavity and onto incliviclual SAIBLIs.," Below, we model surface density profiles and thermal emission spectra of such discs, as well as the thermal emission due to leakage of gas into the cavity and onto individual SMBHs."
+ Adopting a geometrically thin. thermal grav-body.— disc," Adopting a geometrically thin, thermal gray-body disc"
+10 previous two decades a wealth of observational elfort has oen invested in studying the redshift réseime 0«21. --nvestigating whether the dominant. factors are linked. to 1¢ Alpe-seale cluster environments (eg.,"the previous two decades a wealth of observational effort has been invested in studying the redshift réggime $0<z<1$, investigating whether the dominant factors are linked to the Mpc-scale cluster environments (eg."
+ Prestage Peacock 1988: Lll Lilly 1991: Ellingson. Yee του 1991: Wold 6 al.," Prestage Peacock 1988; Hill Lilly 1991; Ellingson, Yee Green 1991; Wold et al."
+ 2000: Best 2000) or the kpe-scale properties of the iost galaxies (e.g. Smith Ileckman 1989: Best. Longair Rotttgering 1998: MeLure Dunlop 2000: Dunlop ct al.," 2000; Best 2000) or the kpc-scale properties of the host galaxies (e.g. Smith Heckman 1989; Best, Longair Rötttgering 1998; McLure Dunlop 2000; Dunlop et al."
+ 003: Zirm. Dickinson Dev 2003)," 2003; Zirm, Dickinson Dey 2003)."
+ One of the incontrovertible observational facts which las emerged from these studies. at least at 2<1L. is that 1¢ host galaxies of all powerful racio-Ioud XN are massive L L' ellipticals (eg.," One of the incontrovertible observational facts which has emerged from these studies, at least at $z<1$, is that the host galaxies of all powerful radio-loud AGN are massive $L>L^{\star}$ ellipticals (eg."
+ Taylor ct al., Taylor et al.
+ 1996: MeLure ct al., 1996; McLure et al.
+ 1999: Dunlop et al., 1999; Dunlop et al.
+ 2003)., 2003).
+" The apparent uniformity of racio-oud AGN host. galaxies has taken on added: importance over the last few vears. following the discovery in nearby (distance z; 150 Alpe) inactive galaxies that a reasonably accurate estimate (AAS,20.3 dex) of the central black-vole mass can be obtained via its correlation with the mass ofthe host spheroidal component (Ixormendsy Lichstone 1995: Magorrian et al."," The apparent uniformity of radio-loud AGN host galaxies has taken on added importance over the last few years, following the discovery in nearby (distance $\ltsim$ 150 Mpc) inactive galaxies that a reasonably accurate estimate $\Delta M_{bh}\simeq 0.3$ dex) of the central black-hole mass can be obtained via its correlation with the mass of the host spheroidal component (Kormendy Richstone 1995; Magorrian et al."
+ 1998: Gebhardt et al., 1998; Gebhardt et al.
+ 2000: Ferrarese Merritt 2000: MeLure Dunlop 2002: Marconi Llunt 2003: Tremaine et al., 2000; Ferrarese Merritt 2000; McLure Dunlop 2002; Marconi Hunt 2003; Tremaine et al.
+ 2002)., 2002).
+" Moreover. recent. progress has also indicated: that a similarly accurate black-hole mass estimate (NAM,c0.4 dex) can be obtained for broad-line AGN using emission-lino. widths to derive the virial mass estimate (eg."," Moreover, recent progress has also indicated that a similarly accurate black-hole mass estimate $\Delta M_{bh}\simeq 0.4$ dex) can be obtained for broad-line AGN using emission-line widths to derive the virial mass estimate (eg."
+ Ixaspi et al., Kaspi et al.
+ 2000: AleLure Dunlop 2002: AleLure Jarvis 2002: Vestergaard 2002)., 2000; McLure Dunlop 2002; McLure Jarvis 2002; Vestergaard 2002).
+ Consequently. à arge body of work has appeared in the recent literature investigating the possible link between racio luminosity and Xack-hole mass in raclio-loucl AGN (eg.," Consequently, a large body of work has appeared in the recent literature investigating the possible link between radio luminosity and black-hole mass in radio-loud AGN (eg."
+ Laor 2000: Lacy οἱ al., Laor 2000; Lacy et al.
+ 2001: AleLure Dunlop 2001a: MeLure Dunlop 2002: Dettoni et al., 2001; McLure Dunlop 2001a; McLure Dunlop 2002; Bettoni et al.
+ 2003: Dunlop ct al., 2003; Dunlop et al.
+ 2003)., 2003).
+ Unfortunately. observational studies have traclitionally oen subject to a degeneracy between radio. luminosity and redshift produced as a by-produet of Hux-limited racio samples.," Unfortunately, observational studies have traditionally been subject to a degeneracy between radio luminosity and redshift produced as a by-product of flux-limited radio samples."
+ In order to study the properties of radio-Ioud GN separated by a large dynamic range in radio luminosity. it has previously been necessary to select. samples. consisting of objects covering a wide range of redshifts.," In order to study the properties of radio-loud AGN separated by a large dynamic range in radio luminosity, it has previously been necessary to select samples consisting of objects covering a wide range of redshifts."
+ This has led to dillicullics in interpreting the data due to the complication of potentially significant evolutionary cllects., This has led to difficulties in interpreting the data due to the complication of potentially significant evolutionary effects.
+ However. by selecting our sample of objects from four complete. Iow-frequeney selected: radio samples with successively fainter Uux-clensity limits. it has been possible to construct a sample of radio galaxies which spans three decades in radio Luminosity at a virtually constant cosmic epoch (0.4>< 0.6).," However, by selecting our sample of objects from four complete, low-frequency selected radio samples with successively fainter flux-density limits, it has been possible to construct a sample of radio galaxies which spans three decades in radio luminosity at a virtually constant cosmic epoch $0.4<z<0.6$ )."
+ In this respect the motivation for this study is identical to that pursued by Hill Lilly (1991). who investigated the cluster. environments of a sample of 45 racio galaxies spanning a similar range in racio bIuminosities ab 200.5.," In this respect the motivation for this study is identical to that pursued by Hill Lilly (1991), who investigated the cluster environments of a sample of 45 radio galaxies spanning a similar range in radio luminosities at $z\simeq 0.5$."
+ ὃν successfully disentanegling the ellects of luminosity and redshift [or the first time. Lill Lilly (1991) were able to demonstrate that z20.5 marks an apparent epoch-dependent. change in cluster environment. with radio galaxies tending to inhabit significantly richer cluster environments than their low-redshift. counterparts.," By successfully disentangling the effects of luminosity and redshift for the first time, Hill Lilly (1991) were able to demonstrate that $z\simeq 0.5$ marks an apparent epoch-dependent change in cluster environment, with radio galaxies tending to inhabit significantly richer cluster environments than their low-redshift counterparts."
+ Furthermore. the wide range of radio luminosities in the Lill Lilly sample enabled. them to determine that extended radio Luminosity did not appear to be a strong function of cluster richness.," Furthermore, the wide range of radio luminosities in the Hill Lilly sample enabled them to determine that extended radio luminosity did not appear to be a strong function of cluster richness."
+ Developments over. the intervening decade have provided our new studs with two crucial advantages over the previous work of Hill Lilly (1991)., Developments over the intervening decade have provided our new study with two crucial advantages over the previous work of Hill Lilly (1991).
+ Firstly. the availability of the TORS (Willott et al.," Firstly, the availability of the 7CRS (Willott et al."
+ 2003: Lacy. ct al., 2003; Lacy et al.
+ 1999). and 'TexOx-1000 (Lil Rawlings 2003) racio samples has now allowed. us to compile a sample of 41 radio galaxies at 2oc0.5. which is complete over the full iee-decadoes range in radio power.," 1999) and TexOx-1000 (Hill Rawlings 2003) radio samples has now allowed us to compile a sample of 41 radio galaxies at $z\simeq 0.5$, which is complete over the full three-decades range in radio power."
+ Secondly. with the high-resolution imaging provided by the Hubble Space Telescope (HIST). and the improvement in follow-up. observations available with ground-based telescopes. we are now in a position o obtain a wide range of high quality data-sets for our new sample which were not obtainable only ten vears ago.," Secondly, with the high-resolution imaging provided by the Hubble Space Telescope (HST), and the improvement in follow-up observations available with ground-based telescopes, we are now in a position to obtain a wide range of high quality data-sets for our new sample which were not obtainable only ten years ago."
+ Consequently. our new sample of z20.5 radio galaxies is he basis for an ambitious project. which will for the first ime svstematically investigate the connection between the radio luminosity of radio galaxies. the properties of their rost galaxies. their location on the fundamental plane. the richness and mass of their cluster environments. and the masses of their central black holes.," Consequently, our new sample of $z\simeq 0.5$ radio galaxies is the basis for an ambitious project, which will for the first time systematically investigate the connection between the radio luminosity of radio galaxies, the properties of their host galaxies, their location on the fundamental plane, the richness and mass of their cluster environments and the masses of their central black holes."
+ The results presented. in this paper represent the first μαage of this wide-ranging project. ancl concentrate solely on the analysis of our deep £ band LIS ΛΕΡΟΣ imaging ata.," The results presented in this paper represent the first stage of this wide-ranging project, and concentrate solely on the analysis of our deep $I-$ band HST WFPC2 imaging data."
+ The results of our extensive follow-up observations of 1e sample are deferred to a series of future papers., The results of our extensive follow-up observations of the sample are deferred to a series of future papers.
+ The structure of the paper is as follows. in Section 2 10 details of the radio-galaxy sample are described. while in Section 3 the LIST observations and data reduction are iscussed.," The structure of the paper is as follows, in Section 2 the details of the radio-galaxy sample are described, while in Section 3 the HST observations and data reduction are discussed."
+ In Section 4 the modelling of the LIST imaging is escribed. with the main results presented and compared to literature results on low-recshilt radio galaxies in Section 5.," In Section 4 the modelling of the HST imaging is described, with the main results presented and compared to literature results on low-redshift radio galaxies in Section 5."
+ In Section 6 the host-ealaxy properties of our οςΙΙ sub-sample are compared to literature. results to explore. the evidence for dynamical evolution of the most powerful radio ealaxies within the redshift range 0.0<z«OLS., In Section 6 the host-galaxy properties of our 3CRR sub-sample are compared to literature results to explore the evidence for dynamical evolution of the most powerful radio galaxies within the redshift range $0.0<z<0.8$.
+ In Section 7 the properties of the radio-galaxy hosts are compared to those of local brightest. cluster galaxies., In Section 7 the properties of the radio-galaxy hosts are compared to those of local brightest cluster galaxies.
+ In. Section S the masses of the radio galaxies central black holes. are estimated. and in Section 9 the relationship between black-hole mass and radio luminosity is investigated.," In Section 8 the masses of the radio galaxies' central black holes are estimated, and in Section 9 the relationship between black-hole mass and radio luminosity is investigated."
+ The main conclusions are summarized in Section LO., The main conclusions are summarized in Section 10.
+" Unless otherwise stated. throughout this paper the following cosmology is assumed: //,=TO tMpe 1.0,203, 04=07."," Unless otherwise stated, throughout this paper the following cosmology is assumed: $H_{0}=70$ $^{-1}$ $^{-1}$, $\Omega_{m}=0.3$, $\Omega_{\Lambda}=0.7$."
+ The full 4l-object z=0.5 radio-galaxy sample (hereafter the ZP5 sample) consists of all the narrow-line radio galaxies in the redshift interval 0.4<z0.6 from four. completo. low-[requeney selected radio survevs: θα (Laing. Riley Longair 1983). GCE (Eales et al.," The full 41-object $z\simeq 0.5$ radio-galaxy sample (hereafter the ZP5 sample) consists of all the narrow-line radio galaxies in the redshift interval $0.4<z<0.6$ from four, complete, low-frequency selected radio surveys; 3CRR (Laing, Riley Longair 1983), 6CE (Eales et al."
+ 1997: Rawlings. Eales," 1997; Rawlings, Eales"
+reached only within the innermost disk. well inside 1 AU.,"reached only within the innermost disk, well inside 1 AU."
+ Crystalline and amorphous silicate erains inmost protoplanetary disks may Caius have experienced the same large-scale transport phases as the refractory particles [ος in Wild 2., Crystalline and amorphous silicate grains in most protoplanetary disks may thus have experienced the same large-scale transport phases as the refractory particles found in Wild 2.
+ While accretion disk models driven by a generic turbulent viscosity have long been invoked as a means to explain large-scale transport (Gail 2001. 2002. 2004: Tschirnuter Gail 2007; Ciesla 2007. 2003. 2009. 2010a.b: Birnstiel et al.," While accretion disk models driven by a generic turbulent viscosity have long been invoked as a means to explain large-scale transport (Gail 2001, 2002, 2004; Tscharnuter Gail 2007; Ciesla 2007, 2008, 2009, 2010a,b; Birnstiel et al."
+ 2009; IIughes Armitage 2010: ]leinzeller οἱ al., 2009; Hughes Armitage 2010; Heinzeller et al.
+ 2011: Jacquet et al., 2011; Jacquet et al.
+ 2011). the detailed physics behind a-viscosity remains unclear. especially considering that the magneto-rotational instability (MBRI) often assumed to be the source of the a-viseosity is unable to drive disk evolution in the magnetically dead midplane regions (e.g.. Matsumura Pudritz 2006) of most interest for planetary formation.," 2011), the detailed physics behind $\alpha$ -viscosity remains unclear, especially considering that the magneto-rotational instability (MRI) often assumed to be the source of the $\alpha$ -viscosity is unable to drive disk evolution in the magnetically dead midplane regions (e.g., Matsumura Pudritz 2006) of most interest for planetary formation."
+ Objections have also been raised to the assumption that angular momentum transport in a MBI-driven disk can be described by the standard model for a-viscosity (Pessah. Chan. Psalüs 2008).," Objections have also been raised to the assumption that angular momentum transport in a MRI-driven disk can be described by the standard model for $\alpha$ -viscosity (Pessah, Chan, Psaltis 2008)."
+ In contrast. a MGQU disk presents a sell-consistent mechanism for studyiug mixing and (rausport in protoplanetary disks. with no ου parameters bevond the initial choice of a disk massive enough. and cold enough. to be MGU.," In contrast, a MGU disk presents a self-consistent mechanism for studying mixing and transport in protoplanetary disks, with no free parameters beyond the initial choice of a disk massive enough, and cold enough, to be MGU."
+ We present here a new set of three-dimensional MGU «disk models similar to those studied previously (e.g.. Boss 2004a. 2006. 2007. 2008). bul with several variations intended to test the robustness of the conclusions about mixing5 and (transport of the previous models.," We present here a new set of three-dimensional MGU disk models similar to those studied previously (e.g., Boss 2004a, 2006, 2007, 2008), but with several variations intended to test the robustness of the conclusions about mixing and transport of the previous models."
+ Given that a MGU disk appears to be a likely requirement for the formation of eas eiant planets. by either core accretion (e.g.. Inaba οἱ al.," Given that a MGU disk appears to be a likely requirement for the formation of gas giant planets, by either core accretion (e.g., Inaba et al."
+ 2003: Chambers 2008) or by clisk instability (e.g.. Boss 2010). these MGU models are intended to learn what cosmochenmical consequences might also derive [rom such phases of rapid disk evolution.," 2003; Chambers 2008) or by disk instability (e.g., Boss 2010), these MGU models are intended to learn what cosmochemical consequences might also derive from such phases of rapid disk evolution."
+ The disk evolution calculations were perlormed with a munerical code (hat uses finite differences to solve the three-dimensional equations of hydrodsnamies. radiative (ransler. ancl the Poisson equation for (he gravitational potential.," The disk evolution calculations were performed with a numerical code that uses finite differences to solve the three-dimensional equations of hydrodynamics, radiative transfer, and the Poisson equation for the gravitational potential."
+ The code is (he same as (hat used in the previous studies of mixing and transport in disks (Boss 2004a. 2006. 2007. 2003).," The code is the same as that used in the previous studies of mixing and transport in disks (Boss 2004a, 2006, 2007, 2008)."
+ The code has been shown to be second-order-accurate in both space and time through convergence testing (Boss Alvhill 1992)., The code has been shown to be second-order-accurate in both space and time through convergence testing (Boss Myhill 1992).
+ The equations are solved on a spherical coordinate grid., The equations are solved on a spherical coordinate grid.
+" The number of grid points in each spatial direction is: NV,= 51. .Nj=23 in $/2>00. and N,,=256."," The number of grid points in each spatial direction is: $N_r = 51$ , $N_\theta = 23$ in $\pi/2 \ge \theta \ge 0$, and $N_\phi = 256$."
+ This relatively low degree of numerical spatial resolution (compared to hieh resolution disk instabilitv models. e.g.. Boss 2010) was chosen in order to evolve the disks as lar Forward in (nme as possible in several vears of computing on a dedicated workstation.," This relatively low degree of numerical spatial resolution (compared to high resolution disk instability models, e.g., Boss 2010) was chosen in order to evolve the disks as far forward in time as possible in several years of computing on a dedicated workstation."
+ The radial grid is uniformly spaced between 1l and 10 AU. with boundary conditions at," The radial grid is uniformly spaced between 1 and 10 AU, with boundary conditions at"
+"of m, and therefore variations with larger amplitudes than 5530 (see Fig. 3)).","of m, and therefore variations with larger amplitudes than 530 (see Fig. \ref{fig:lc}) )."
+" We performed Chi-Square-tests to characterize the significance of the variability, following e.g., ?.."," We performed Chi-Square-tests to characterize the significance of the variability, following e.g., \citet{2003A&A...401..161K}."
+" For the reduced X2 we obtain values of 32.2 at 22GGHz, 7.5 at 43GGHz, and 0.9 at 86GGHz."," For the reduced $\chi_\nu^2$ we obtain values of 32.2 at GHz, 7.5 at GHz, and 0.9 at GHz."
+" The corresponding probabilities for the source not being variable are far less than Although being formally insignificant, the variations at GGHz appear to correlate with the variations seen at the two lower frequencies."," The corresponding probabilities for the source not being variable are far less than Although being formally insignificant, the variations at GHz appear to correlate with the variations seen at the two lower frequencies."
+" To describe the strength of the variability, the modulation index m and the variability amplitude Y (defined as 3 x nm?—m, where mo is the modulation index of the calibrator NRAO5530) are summarized in Table 4,, where Y corresponds to a 3 c variability amplitude, from which systematic variations mg, which are still seen in the calibrator, are subtracted (?).."," To describe the strength of the variability, the modulation index m and the variability amplitude Y (defined as 3 $\times$ $\sqrt{m^2-m_{0}^2}$, where $_{0}$ is the modulation index of the calibrator 530) are summarized in Table \ref{tab:flux}, , where Y corresponds to a 3 $\sigma$ variability amplitude, from which systematic variations $m_0$, which are still seen in the calibrator, are subtracted \citep{1987AJ.....94.1493H}."
+ For AA* those from systematic bias corrected Y-amplitudes range between The observed day-to-day variations compare well with similar variations seen by other authors at other times., For A* those from systematic bias corrected Y-amplitudes range between The observed day-to-day variations compare well with similar variations seen by other authors at other times.
+" Using the VLA, ? found an increase in the flux density at a level of Since for AA* each VLBI track lasted about 6-7 hours, we are Sgrnot able to detect flux density variations on shorter timescales than this."," Using the VLA, \citet{2006ApJ...650..189Y} found an increase in the flux density at a level of Since for A* each VLBI track lasted about 6-7 hours, we are not able to detect flux density variations on shorter timescales than this."
+" A splitting of the VLBI coverage at shorter intervals, e.g. in two or three coverages of equal duration, does not allow measuring the total source flux with sufficient accuracy (main limitation: uv-coverage and lack of secondary calibrator scans) and therefore prevents the significant detection of variability on timescales shorter than 6hhrs."," A splitting of the VLBI coverage at shorter intervals, e.g. in two or three coverages of equal duration, does not allow measuring the total source flux with sufficient accuracy (main limitation: uv-coverage and lack of secondary calibrator scans) and therefore prevents the significant detection of variability on timescales shorter than hrs."
+" A nonstationary source, which would vary with a large amplitude during the time of the VLBI experiment, however, would cause significant image degradation, leading to a reduced dynamical range in the CLEAN maps and the appearance of side lobes."," A nonstationary source, which would vary with a large amplitude during the time of the VLBI experiment, however, would cause significant image degradation, leading to a reduced dynamical range in the CLEAN maps and the appearance of side lobes."
+" Since this is not observed, we can exclude variations that are much larger than our typical amplitude calibration errors of 10—20 From the measured total flux densities, we calculated a 3-frequency spectral index between 22 and 86GGHz (defined as S,« vy”)."," Since this is not observed, we can exclude variations that are much larger than our typical amplitude calibration errors of $10-20$ From the measured total flux densities, we calculated a 3-frequency spectral index between 22 and GHz (defined as $_{\nu} \propto \nu^{\alpha}$ )."
+" We obtained an inverted spectrum, with spectral indices ranging between (0.44 + 0.04) and (0.64 + 0.05)."," We obtained an inverted spectrum, with spectral indices ranging between (0.44 $\pm$ 0.04) and (0.64 $\pm$ 0.05)."
+" For AA*, a frequency break in the spectrum was suggested between ~20—100 GGHz (???).. "," For A*, a frequency break in the spectrum was suggested between $\sim 20 - 100$ GHz \citep{1998ApJ...499..731F, 2003ApJ...586L..29Z,
+2005ApJ...634L..49A}."
+"Below this break frequency, the spectral slope is much shallower (lower) than at higher frequencies, where the so-called sub-mm excess causes an increase in the inverted spectral index (??).. "," Below this break frequency, the spectral slope is much shallower (lower) than at higher frequencies, where the so-called sub-mm excess causes an increase in the inverted spectral index \citep{1997ApJ...490L..77S, 2006JPhCS..54..328K}."
+Our observing frequencies are just in the transition region between cm- and sub-mm range., Our observing frequencies are just in the transition region between cm- and sub-mm range.
+" Therefore the measured spectral indices are slightly higher than previously reported spectra, resulting from VLBI at cm-wavelengths."," Therefore the measured spectral indices are slightly higher than previously reported spectra, resulting from VLBI at cm-wavelengths."
+ ? made simultaneous multi-wavelength observations of AA* in 2003., \citet{2005ApJ...634L..49A} made simultaneous multi-wavelength observations of A* in 2003.
+ They describe the spectrum from short centimeter (3.6ccm) to millimeter mmm) wavelengths by a power law of the form S , They describe the spectrum from short centimeter cm) to millimeter mm) wavelengths by a power law of the form S $\propto \nu^{0.43}$.
+ος99. ? measured a spectral index of 0.52 between mmm and 2mmm wavelength., \citet{1998ApJ...499..731F} measured a spectral index of 0.52 between mm and mm wavelength.
+" The observed spectral indices reported here are fullyconsistent with these previous studies, and confirm the onset of a sub-mm excess over the more shallow power-law shape seen at cm-wavelength."," The observed spectral indices reported here are fullyconsistent with these previous studies, and confirm the onset of a sub-mm excess over the more shallow power-law shape seen at cm-wavelength."
+its maxinmn briehtness.,its maximum brightness.
+ The average post-imipact maximum-light spectral distribution is shifted to match the average flix [rom 625 nm to 3875 nm at other times., The average post-impact maximum-light spectral distribution is shifted to match the average flux from 625 nm to 875 nm at other times.
+ The deviations from this fits al times prior to the impact. and in the days after the impact when the inner coma of comet Tempel 1. had essentially returned to its normal condition. show that the Εκ at short wavelenths was lower al those times.," The deviations from this fits at times prior to the impact, and in the days after the impact when the inner coma of comet Tempel 1 had essentially returned to its normal condition, show that the flux at short wavelenths was lower at those times."
+ The 375 mm data point is about 0.2 mag lower (han during the peak post-impact brightness., The 375 nm data point is about 0.2 mag lower than during the peak post-impact brightness.
+ This is particularly well seen in the 7/5.∕ 7/7. and 7/8 data points in Fig.," This is particularly well seen in the 7/5, 7/7, and 7/8 data points in Fig."
+ Th., 7b.
+ For Che quiescent state of the comet. the long ∖∖⊽≀↧↴∖↽≼↲↥≼↲∐≸≟⊔↥↕↽≻∪↕∐↥⋝∖⊽∐≼↲⊳∖⇁∡∖⇁⊳∖⊽∩↲↕∐≀↧↴∐≺∢≀↧↴∐⋡∖↽≀↧↴∣↽≻∪∖⇁≼↲⊔∐↲↓∎↓⊔≼↲≼⊔↕∐↕↽≻≀↧↴≺∢↥≼↲∙↿≼↲≺∢↥≀↧↴⊳∖⇁↕↽≻≼↲≺∢∏⋅∏∐↓⋅⊔∐↲≼↲∐≱≼↲≺∢↥↕⊳∖⇁ small (220.05 mag).," For the quiescent state of the comet, the long wavelength points lie systematically above the fitted impact ejecta spectrum, the effect is small $\approx$ 0.05 mag)."
+ Over the full range [rom 375 nm to 925 nm. we observe that the impact eeneraled ejecta were 220.25 mag bluer (han the quiescent comet coma.," Over the full range from 375 nm to 925 nm, we observe that the impact generated ejecta were $\approx$ 0.25 mag bluer than the quiescent comet coma."
+ This corresponds to awd per 100 nm change in the slope of the spectrum averaged over (he wavelength range from 375 nm to 925 nm., This corresponds to a $\approx$ per 100 nm change in the slope of the spectrum averaged over the wavelength range from 375 nm to 925 nm.
+ This change in color can be explained by an unspecifie combination of two effects., This change in color can be explained by an unspecific combination of two effects.
+ First. the particle size distribution of the material ejected after (he impact may contain a lareer fraction of very small particles. much smaller (han (he wavelengths of visible light that leads to a bluer color of the Raleigh scattered light.," First, the particle size distribution of the material ejected after the impact may contain a larger fraction of very small particles, much smaller than the wavelengths of visible light that leads to a bluer color of the Raleigh scattered light."
+ Second. it may point to large quantities of pure water ice crvstals that are known to have blue optical rellection spectra (Liev Clark 1985).," Second, it may point to large quantities of pure water ice crystals that are known to have blue optical reflection spectra (Lucey Clark 1985)."
+ Infrared observations of the post-impact material have also found indications that the size distribution of impact-ejected material contained more small particles than were released bv (he comet outside of the impact event (IIarker Woodward and Wooden 2005) and (Sugita et al., Infrared observations of the post-impact material have also found indications that the size distribution of impact-ejected material contained more small particles than were released by the comet outside of the impact event (Harker Woodward and Wooden 2005) and (Sugita et al.
+ 2005)., 2005).
+ In combination. the changes in (he comet's color in the hours after impact indicate that the material ejected by (the impact contains smaller particles aud more ice. and is therefore probably more pristine than the material released from the surface of the comet under normal conditions.," In combination, the changes in the comet's color in the hours after impact indicate that the material ejected by the impact contains smaller particles and more ice, and is therefore probably more pristine than the material released from the surface of the comet under normal conditions."
+potential in a rotating reference [rame centered. on the cluster centre-of-mass. with the .r-axis directed. away from. the galactie centre and the y-axis in the direction of motion (see Ciersz Llegeic 1997: Vesperini Legeie 1997).,potential in a rotating reference frame centered on the cluster centre-of-mass with the $x$ -axis directed away from the galactic centre and the $y$ -axis in the direction of motion (see Giersz Heggie 1997; Vesperini Heggie 1997).
+ A tidal radius (also called the Jacobi radius: Gieles Baumearelt 2008). can then be defined corresponding to the saddle point on the aeaxis of the ellective cluster potential.," A tidal radius (also called the Jacobi radius: Gieles Baumgardt 2008), can then be defined corresponding to the saddle point on the $x$ -axis of the effective cluster potential."
+ Within an /-bodvy simulation the user has the freedom to set a length-scale Ro. with one possible choice being that the outermost stars of the initial density. profile are scaled to sit at à. in which case the cluster is said to be Itoche-Iobe filling (e.g. Fanikawa Fukushige 2005).," Within an $N$ -body simulation the user has the freedom to set a length-scale $R_{\rm sc}$ with one possible choice being that the outermost stars of the initial density profile are scaled to sit at $r_{\rm t}$, in which case the cluster is said to be Roche-lobe filling (e.g. Tanikawa Fukushige 2005)."
+ Phe Plummoer profile formally. extends to infinite radius so in practice a cut-olf at à radius of is applied to avoid rare cases of large distance.," The Plummer profile formally extends to infinite radius so in practice a cut-off at a radius of $\sim 10 \, r_{\rm h}$ is applied to avoid rare cases of large distance."
+ For our Plunimer models described below this leads to Fia2Srn. where rugas ds the position of the outermost star. and. we define κε as the tidal-radius filling factor.," For our Plummer models described below this leads to $r_{\rm max} \simeq 8 \, r_{\rm h}$, where $r_{\rm max}$ is the position of the outermost star, and we define $r_{\rm max} / r_{\rm t}$ as the tidal-radius filling factor."
+ The related ratio rj/ri is an important quantity to describe the structure ol star cluster models., The related ratio $r_{\rm h} / r_{\rm t}$ is an important quantity to describe the structure of star cluster models.
+" In this work we model parent galaxies with two distinct masses: Ma=9010""M. to represent a dwarl galaxy such as NOGCG6822 and Ak=9.103134. to model a more substantial galaxy such as M31."," In this work we model parent galaxies with two distinct masses: $M_{\rm g} = 9 \times 10^{9} \, M_\odot$ to represent a dwarf galaxy such as $\,6822$ and $M_{\rm g} = 9 \times 10^{10} \, M_\odot$ to model a more substantial galaxy such as M31."
+" As discussed. above all clusters are set to orbit at 76,= LOkpe in these model galaxies."," As discussed above all clusters are set to orbit at $R_{\rm gc} = 10\,$ kpc in these model galaxies."
+ Six distinct. models are. performed., Six distinct models are performed.
+ These are listed in ‘Table 1.., These are listed in Table \ref{t:table1}.
+. Models NI. N2 and N3 are all evolved. in. the 6822-like tidal field.," Models N1, N2 and N3 are all evolved in the $\,6822$ -like tidal field."
+ They are identical in all respects except for using ιο=7. 14 and 21 so that they have initial ticlal-raclii filling factors of 0.33. 0.66 ancl 1.00. respectively (the corresponding rm/r; ratios are given in Table 1)).," They are identical in all respects except for using $R_{\rm sc} = 7$, 14 and 21 so that they have initial tidal-radii filling factors of 0.33, 0.66 and 1.00, respectively (the corresponding $r_{\rm h} / r_{\rm t}$ ratios are given in Table \ref{t:table1}) )."
+ \loclel ΝΟΕ is the same as N2 except that it starts with 95000 single stars and 5000 binaries rather than 100000. single. stars.," Model N2b is the same as N2 except that it starts with $95\,000$ single stars and $5\,000$ binaries rather than $100\,000$ single stars."
+ We then have Model MI whieh is evolved in the stronger AIS1-like tical field for comparison., We then have Model M1 which is evolved in the stronger M31-like tidal field for comparison.
+ All of these models start with a Plummer density. profile whereas Model Al2 starts with a Wine profile and is also evolved in the M31-like tidal field., All of these models start with a Plummer density profile whereas Model M2 starts with a King profile and is also evolved in the M31-like tidal field.
+ Both MI and M2 start with tidal-racii filling factors of 1.00 (rs matches the tidal radius)., Both M1 and M2 start with tidal-radii filling factors of 1.00 $r_{\rm max}$ matches the tidal radius).
+ The initial mass of each model is AM258000A. which gives an initial ri of 129 pe for models NI. N2. N2b and N3 compared to 60 pc for models ALL ancl M2.," The initial mass of each model is $M \simeq 58\,000 \, M_\odot$ which gives an initial $r_{\rm t}$ of $129\,$ pc for models N1, N2, N2b and N3 compared to $60\,$ pc for models M1 and M2."
+ Each model is evolved. to an age of 20 vr or until of the stars remain. whichever occurs first 1e latter only happens for M2 at an age of 17.4Gvr.," Each model is evolved to an age of $20\,$ Gyr or until of the stars remain, whichever occurs first – the latter only happens for M2 at an age of $17.4\,$ Gyr."
+ The simulations are performed using Tesla SLOTO GPUs at Swinburne University., The simulations are performed using Tesla S1070 GPUs at Swinburne University.
+ We will also lean on the results of some previous iN- simulations when evaluating our results., We will also lean on the results of some previous $N$ -body simulations when evaluating our results.
+ These include models. IX100-00a. and Ix100-00b. of LIurley. (2007) which both featurecl 100000 single stars evolved within a standard Galactic tidal field (Ciersz Llegeic 1997): an orbital speed of 220kms tat Reo=S 5kpe (with corresponding AL.=9.LOM ALY.," These include models K100-00a and K100-00b of Hurley (2007) which both featured $100\,000$ single stars evolved within a standard Galactic tidal field (Giersz Heggie 1997): an orbital speed of $220 \, {\rm km} \, {\rm s}^{-1}$ at $R_{\rm gc} = 8.5\,$ kpc (with corresponding $M_{\rm g} = 9 \times 10^{10} \, M_\odot$ )."
+ Phe two models were setup in the same wav and the difference of note was the formation of a long-lived binary composed. of two stellar-mass black holes (12115) in Ix100-00b which altered the central structure of the cluster compared to W100-00a., The two models were setup in the same way and the difference of note was the formation of a long-lived binary composed of two stellar-mass black holes (BHs) in K100-00b which altered the central structure of the cluster compared to K100-00a.
+ Along the same lines we will use the results of Alackey ct al. (, Along the same lines we will use the results of Mackey et al. (
+2008) who looked at the clfect ofa population of DII-BII binaries on cluster evolution. as well as models including intermeciate-mass black-holes (e.g. Cull et al.,"2008) who looked at the effect of a population of BH-BH binaries on cluster evolution, as well as models including intermediate-mass black-holes (e.g. Gill et al."
+ 2008)., 2008).
+ Also mentioned will be models of 30000 single stars from Liurley et al. (," Also mentioned will be models of $30\,000$ single stars from Hurley et al. ("
+2004). mainly [or illustrative purposes.,"2004), mainly for illustrative purposes."
+ lt is important to emphasize that unless. otherwise μα»ecified the radii quoted will be based on three-dimoensional ata., It is important to emphasize that unless otherwise specified the radii quoted will be based on three-dimensional data.
+ The core radius. r«. comes from a cdensitv-weighted aleulation (C'asertano μι 1985) that is. traditionally used in. N-body models. and. is not. comparable to. the uantitv derived. by observational methods.," The core radius, $r_{\rm c}$, comes from a density-weighted calculation (Casertano Hut 1985) that is traditionally used in $N$ -body models and is not comparable to the quantity derived by observational methods."
+ This has been iscussecl in the past (e.g. Wilkinson et al., This has been discussed in the past (e.g. Wilkinson et al.
+ 2003: Llurley 2007)., 2003; Hurley 2007).
+ For our purposes this is fine as we use re as an indicator of the cluster cdvnamical state. in. particular to etermine if core-collapse has been reached. rather than to compare to observed results for actual clusters.," For our purposes this is fine as we use $r_{\rm c}$ as an indicator of the cluster dynamical state, in particular to determine if core-collapse has been reached, rather than to compare to observed results for actual clusters."
+ We also use ry as the three-cimensional half-miass racius., We also use $r_{\rm h}$ as the three-dimensional half-mass radius.
+ Here we are mainly interested in the relative values between mocels., Here we are mainly interested in the relative values between models.
+ Llowever. we will also provide the hall-leht radius. yo. calculated. from a two-dimensional projection of the data. to give a reference point for comparing the size of the mocel clusters to real clusters.," However, we will also provide the half-light radius, $r_{\rm h,l}$, calculated from a two-dimensional projection of the data, to give a reference point for comparing the size of the model clusters to real clusters."
+ The cilferent initial filling factors of models NI. N2 and N3 leac to cülferent. initial half-mass radii (see Table 1)) and allow us to investigate the elfect this has on the long-term evolution of star clusters residing within a weak tidal field.," The different initial filling factors of models N1, N2 and N3 lead to different initial half-mass radii (see Table \ref{t:table1}) ) and allow us to investigate the effect this has on the long-term evolution of star clusters residing within a weak tidal field."
+ lt is well established. that the evolution of a star cluster is intricately Linked to the two-body relaxation timescale which is typically characterized by the value at the racius.," It is well established that the evolution of a star cluster is intricately linked to the two-body relaxation timescale which is typically characterized by the value at the half-mass radius,"
+Blandford 2010; Wang et al.,Blandford 2010; Wang et al.
+ 2010; Shcherbakov Huang 2011)., 2010; Shcherbakov Huang 2011).
+" Budden eqution, i.e., the second-order equation to the complex function actually corresponds to the same approach (Budden 1972; Zheleznyakov 1977)."," Budden eqution, i.e., the second-order equation to the complex function actually corresponds to the same approach (Budden 1972; Zheleznyakov 1977)."
+" On the other hand, both the standard and the Zheleznyakov-Budden approaches are not quite for quantitative estimates of the polarization of the escaping emission in general case."," On the other hand, both the standard and the Zheleznyakov-Budden approaches are not quite for quantitative estimates of the polarization of the escaping emission in general case."
+" But since we are going to describe the propagation of originally fully polarized waves, not theensemble of waves, we actually need only two equations for observable parameters, i.e., the position angle and the Stokes parameter V."," But since we are going to describe the propagation of originally fully polarized waves, not the of waves, we actually need only two equations for observable parameters, i.e., the position angle and the Stokes parameter $V$."
+" There exists a different approach that allows us immediately write down the equations for these observable quantities, namely, the Stokes parameter V, defining the circular polarization and the position angle p.a., characterizing the orientation of polarization ellipse (Kravtsov Orlov 1990)."," There exists a different approach that allows us immediately write down the equations for these observable quantities, namely, the Stokes parameter $V$, defining the circular polarization and the position angle $p.a.$, characterizing the orientation of polarization ellipse (Kravtsov Orlov 1990)."
+" This approach is valid in the quasi-isotropic case, i.e., in the case when the dielectric tensor can be presented as where the anisotropic part χι is small as compared to isotropic one."," This approach is valid in the quasi-isotropic case, i.e., in the case when the dielectric tensor can be presented as where the anisotropic part $\chi_{ij}$ is small as compared to isotropic one."
+" In this case we have two small parameters — general WKB parameter 1/kL and As a result, the solution can be found by expansion over this two small parameters."," In this case we have two small parameters — general WKB parameter $1/kL$ and As a result, the solution can be found by expansion over this two small parameters."
+" As one can check, these conditions are just realized in the pulsar magnetosphere (Andrianov Beskin 2010)."," As one can check, these conditions are just realized in the pulsar magnetosphere (Andrianov Beskin 2010)."
+" Indeed, in the region r~res;10”R the value of v=wp/w* is much smaller than unity."," Indeed, in the region $r \sim r_{\rm esc} \sim 10^{3}R$ the value of $v = \omega_{\rm p}^2/\omega^2$ is much smaller than unity."
+" Accordingly, the deviation of the refractive indices from unity, |n1,2—1|ev, is also very small here, so we can neglect the wave refraction in the polarization formation region."," Accordingly, the deviation of the refractive indices from unity, $|n_{1,2}-1| \sim v$, is also very small here, so we can neglect the wave refraction in the polarization formation region."
+" 'The Kravtsov-Orlov equation is the equation for the complex angle O=0)+i0», where O is a position angle and O5 determines the circular polarization by the relation Here I is the intensity of the wave."," The Kravtsov-Orlov equation is the equation for the complex angle $\Theta = \Theta_1 + i \Theta_2$, where $\Theta_1$ is a position angle and $\Theta_2$ determines the circular polarization by the relation Here $I$ is the intensity of the wave."
+ The components of the dielectric tensor xi; are to be written in a frame of unitary vectors a and b in the picture plane where a is determined by the projection of the vector Ve., The components of the dielectric tensor $\chi_{ij}$ are to be written in a frame of unitary vectors ${\bf a}$ and ${\bf b}$ in the picture plane where ${\bf a}$ is determined by the projection of the vector $\nabla \varepsilon$.
+" Finally, is the ray torsion (see Kravtsov Orlov 1990 for more detail)"," Finally, is the ray torsion (see Kravtsov Orlov 1990 for more detail)."
+ It can be easily understood that the rotation of position angle described by the ray torsion is fictious and describes only the rotation of coordinate system., It can be easily understood that the rotation of position angle described by the ray torsion is fictious and describes only the rotation of coordinate system.
+" As a result, we can write down Here is a coordinate along the ray propagation, and the angle Bg(1) defines the orientation of the external magnetic field in the picture plane 3)."," As a result, we can write down Here $l$ is a coordinate along the ray propagation, and the angle $\beta_{B}(l)$ defines the orientation of the external magnetic field in the picture plane )."
+" Further, where the signs correspond to the regions before/after the cyclotron resonance and Finally, εν} are the components of plasma dielectric tensor in the frame where the z-axis directs along the wave propagation and the external magnetic field lies in the xz- plane (see Appendix C)."," Further, where the signs correspond to the regions before/after the cyclotron resonance and Finally, $\varepsilon_{i'j'}$ are the components of plasma dielectric tensor in the frame where the $z$ -axis directs along the wave propagation and the external magnetic field lies in the $xz$ -plane (see Appendix C)."
+ We would like to note that in these equations the circular polarization is defined as it is common in radio astronomy (positive V corresponds to LHC polarization)., We would like to note that in these equations the circular polarization is defined as it is common in radio astronomy (positive $V$ corresponds to LHC polarization).
+ Nonrelativistic version of the above equations is given in Czyz et al. (, Nonrelativistic version of the above equations is given in Czyz et al. (
+2007).,2007).
+"2008). As one can see on Fig. 10,,"," As one can see on Fig. \ref{angles},"
+ in the geometrical optics region Eqns. (70))-(71)), in the geometrical optics region Eqns. \ref{t1}) \ref{t2}) )
+ describe of the angle Θι near the value 0;=fg-Γδ., describe of the angle $\Theta_{1}$ near the value $\Theta_{1} = \beta_{B} + \delta$.
+" As the ray moves into the region of rarefied plasma, the length of the spatialoscillations L~c/(wAn) increases and in the region becomes larger than the characteristic length r."," As the ray moves into the region of rarefied plasma, the length of the spatial $L \sim c/(\omega \Delta \, n)$ increases and in the region becomes larger than the characteristic length $r$."
+" As a result, the angles Θι and O» become constant for r>>resc."," As a result, the angles $\Theta_{1}$ and $\Theta_{2}$ become constant for $r \gg r_{\rm esc}$."
+ They are the values that characterize the outgoing radiation., They are the values that characterize the outgoing radiation.
+" Thus, the basic equations (70))-(71)) generalize ones obtained by Andrianov Beskin (2010) for zero drift velocity U=0 when Re[e,,]=0 and, hence, 6= 0."," Thus, the basic equations \ref{t1}) \ref{t2}) ) generalize ones obtained by Andrianov Beskin (2010) for zero drift velocity ${\bf U}=0$ when ${\rm Re}\,[\varepsilon_{x'y'}] = 0$ and, hence, $\delta = 0$ ."
+" In particular, they now include into consideration the aberration effect considered by Blaskiewicz et al. ("," In particular, they now include into consideration the aberration effect considered by Blaskiewicz et al. ("
+1991).,1991).
+" This effect was also considered by Petrova Lyubarskii (2000), but for the infinite magnetic field only."," This effect was also considered by Petrova Lyubarskii (2000), but for the infinite magnetic field only."
+ It is important that in Eqns. (70))-(71)), It is important that in Eqns. \ref{t1}) \ref{t2}) )
+ the angle ©; is measured relative to the laboratory frame because these equations contain the difference between ©; and fp only., the angle $\Theta_1$ is measured relative to the laboratory frame because these equations contain the difference between $\Theta_1$ and $\beta_{B}$ only.
+ Equations above have the following important property., Equations above have the following important property.
+" For homogeneous media (fg= const, ei;= const) the parameters of polarization ellipse O1 and Θο remain constant if the following conditions are valid: Here (see the definition of e»;; in Appendix D)"," For homogeneous media $\beta_{B} = $ const, $\varepsilon_{ij} =$ const) the parameters of polarization ellipse $\Theta_1$ and $\Theta_2$ remain constant if the following conditions are valid: Here (see the definition of $\epsilon_{i'j'}$ in Appendix D)"
+Our * value of 2.1+0.2 is consistent with that of the CMF in the OMC-1 region (Ikeda&Kitamura2009)..,Our $\gamma$ value of $\pm$ 0.2 is consistent with that of the CMF in the OMC-1 region \citep{ike09b}.
+ This agreement confirms that our observations could resolve star-forming cores even at the large distance of ~ 1 kpc., This agreement confirms that our observations could resolve star-forming cores even at the large distance of $\sim $ 1 kpc.
+ A poor spatial resolution is one of the major causes of the underestimation of y., A poor spatial resolution is one of the major causes of the underestimation of $\gamma$.
+" To examine the dependence of y on the spatial resolution, we created the smoothed OMC-1 data cubes by changing the effective resolutions, as described in 84.2,, and derived the y values for the smoothed cubes, as shown in Figure 16.."," To examine the dependence of $\gamma$ on the spatial resolution, we created the smoothed OMC-1 data cubes by changing the effective resolutions, as described in \ref{coreProperties}, and derived the $\gamma$ values for the smoothed cubes, as shown in Figure \ref{gamma-res}. ."
+" It is clearly shown that our resolution of 22"", corresponding to 0.097 pc for $140, can correctly estimate the y value within the uncertainties."," It is clearly shown that our resolution of $''$, corresponding to 0.097 pc for S140, can correctly estimate the $\gamma $ value within the uncertainties."
+" Actually, our ¥ value is consistent with that in the study by Tachiharaetal.(2002),, having a spatial resolution enough to resolve 0.1 pc-scale cores."," Actually, our $\gamma$ value is consistent with that in the study by \citet{tac02}, having a spatial resolution enough to resolve 0.1 pc-scale cores."
+" In addition, Figure 16 predicts that γ is considerably underestimated for the case that becomes larger than 0.1 pc, which is the minimum radius of the cores in the OMC-1 region at the highest resolution of 0.061 pc."," In addition, Figure \ref{gamma-res} predicts that $\gamma$ is considerably underestimated for the case that becomes larger than 0.1 pc, which is the minimum radius of the cores in the OMC-1 region at the highest resolution of 0.061 pc."
+" Therefore, the small 7 value of ~ 1.7 in RCW 106 (Wongetal.2008) is likely due to the coarse spatial resolution of 0.78 pc at the distance of 3.6 kpc (Lockman1979)."," Therefore, the small $\gamma$ value of $\sim$ 1.7 in RCW 106 \citep{won08}
+ is likely due to the coarse spatial resolution of 0.78 pc at the distance of 3.6 kpc \citep{loc79}."
+". This study concludes that the power-law shape with y > 2 in CMF holds even in tenuous structures with the densities of 1054 of the S140 region, in addition to our recent work by Ikeda&Kitamura(2009) in OMC-1."," This study concludes that the power-law shape with $\gamma$ $>$ 2 in CMF holds even in tenuous structures with the densities of $^{3\mbox{\scriptsize --}4}$ of the S140 region, in addition to our recent work by \citet{ike09b} in OMC-1."
+" Furthermore, the y values of the CMFs in $140 and OMC-1 are quite consistent with that of the Galactic field-averaged IMF of 2.340.7 (Kroupa2001).."," Furthermore, the $\gamma$ values of the CMFs in S140 and OMC-1 are quite consistent with that of the Galactic field-averaged IMF of $\pm$ 0.7 \citep{kro01a}."
+ These observational facts lead us to the hypothesis that the power-law nature in the IMF originates in molecular cloud structures with densities of less thancm?., These observational facts lead us to the hypothesis that the power-law nature in the IMF originates in molecular cloud structures with densities of less than.
+" Our conclusion of the resemblance between the CMF and the IMF is consistent with the recent theoretical works showing that such a resemblance should be understood as a statistical relation, rather than aone-to-one correspondence between a core and a star to be formed within it."," Our conclusion of the resemblance between the CMF and the IMF is consistent with the recent theoretical works showing that such a resemblance should be understood as a statistical relation, rather than aone-to-one correspondence between a core and a star to be formed within it."
+ Smithetal.(2009) examined the formation and evolution of, \citet{smi09} examined the formation and evolution of
+é] frequencies.,all frequencies.
+ The latter point is mostly based ou nousimitancous observations by studving the energy clistribution at different frequencies (e. Lundgreu ct i, The latter point is mostly based on non-simultaneous observations by studying the energy distribution at different frequencies (e.g. Lundgren et al.
+ 1995). but simultaneous dual-frequeucy observations have been mace for the Crab pulsar (see Salluen et al.," 1995), but simultaneous dual-frequency observations have been made for the Crab pulsar (see Sallmen et al."
+ 1999 aud references therein) where oulv of all eiat pulses are secu both at GOO and 1100 MITz., 1999 and references therein) where only of all giant pulses are seen both at 600 and 1400 MHz.
+ The results of these studies for the spectral index of giant pulses are somewhat iucouclusive., The results of these studies for the spectral index of giant pulses are somewhat inconclusive.
+ While carler observations conclude that. on average. he spectra index of giant pulses is flatter than he average main pulse spectral index. this fiudiug is not supported by the most recent study where the average spectral index of the giant pulscs is comparable to that of the average main pulse (see discussion by Sallmen et al.," While earlier observations conclude that, on average, the spectral index of giant pulses is flatter than the average main pulse spectral index, this finding is not supported by the most recent study where the average spectral index of the giant pulses is comparable to that of the average main pulse (see discussion by Sallmen et al."
+ 1999)., 1999).
+ Tlowever. the Crab pulsar has a very steep radio spectruni. so little may be learut from the relative streneth of its normal aud eiut pulses at either extreme eud of the radio spectrin.," However, the Crab pulsar has a very steep radio spectrum, so little may be learnt from the relative strength of its normal and giant pulses at either extreme end of the radio spectrum."
+ However. we can compare other aspects of the strong pulses frou PSR B1133|16 pulses to the known propertics of giant pulses or those of so-called “elaut nicropulses.," However, we can compare other aspects of the strong pulses from PSR B1133+16 pulses to the known properties of giant pulses or those of so-called “giant micropulses”."
+ The latter have been discovered for the Vela pulsar as very strong narrow pulses of very stall width occurring at fixed. narrow phases (Jolustou et al.," The latter have been discovered for the Vela pulsar as very strong narrow pulses of very small width occurring at fixed, narrow phases (Johnston et al."
+ 2001). also showing a power-law cherey distribution.," 2001), also showing a power-law energy distribution."
+ While a power-law visible τι the cuuimlative xobabilitv function of elaut pulse flux deusities clearly separates them from normal pulses. this appears not to o the case for PSR B1133116 in Fie. 16..," While a power-law visible in the cumulative probability function of giant pulse flux densities clearly separates them from normal pulses, this appears not to be the case for PSR B1133+16 in Fig. \ref{fluxdist1133}."
+" It is possible hat a amber of LO ""eiant pulses is too small to produce a recognisable feature. but a rate of one strong pulse in every 100 (at [850 MITZ) is a ιο. ligher rate than or PSRs 21 aud D19372] and is comparable to hat of the Crab pulsar."," It is possible that a number of 40 “giant” pulses is too small to produce a recognisable feature, but a rate of one strong pulse in every 100 (at 4850 MHz) is a much higher rate than for PSRs $-$ 24 and B1937+21 and is comparable to that of the Crab pulsar."
+ The noticeable chauge in slope of the hieh cucrey cud of PSR D1133|16's cumulative xobabilitv functions. when come from low to high radio requencies. nay be au indication of an cierging powcr-aw compoucut.," The noticeable change in slope of the high energy end of PSR B1133+16's cumulative probability functions, when going from low to high radio frequencies, may be an indication of an emerging power-law component."
+ It appears that the stroug pulses in PSR D11533|16 xeferablv occur at the pulse phase of the leading component., It appears that the strong pulses in PSR B1133+16 preferably occur at the pulse phase of the leading component.
+ This is verified by computing the average xofile from the 10 strong pulses at 1850 ΛΠΣ., This is verified by computing the average profile from the 40 strong pulses at 4850 MHz.
+ This profile is shown as the inset to Fig., This profile is shown as the inset to Fig.
+ 17. where we compare 1 to the average pulse profile., \ref{giants} where we compare it to the average pulse profile.
+ It becomes obvious that the strong oilses appear to be narrower aud indeed appear mostly at he trailing edee of the leading component. being shelthy offset from its centre.," It becomes obvious that the strong pulses appear to be narrower and indeed appear mostly at the trailing edge of the leading component, being slightly offset from its centre."
+ This implies that anv correlation iu he flux densities of the two normal profile compoucuts, This implies that any correlation in the flux densities of the two normal profile components
+Equation 28).,Equation 28).
+" As noted emlier. since this scenario is nearly face-on. fopr Is a factor ~3.1 higher than the detection fraction =Locos20;520, for isotropic eniission."," As noted earlier, since this scenario is nearly face-on, $f_{\rm opt}$ is a factor $\sim 3.4$ higher than the detection fraction $\approx
+1-\cos 2\theta_j\approx 2\bar{\theta_j}^{2}$ for isotropic emission."
+ Equation 3 shows that ifthe average opening anele is ~0.12. which is the value inferred for 00512214 0j(Dirvowsetal.2006:Soderberg2006).. as well as the typical opening angle required to recoucile the observed SCRB rate with the best-bet NS-NS ierecr rate refsecGRB)}. then up to fopr~OL of CW: eveuts will be accompanied by potentially detectable optical afterelows.," Equation \ref{eq:fopt} shows that if the average opening angle is $\bar{\theta}_j\simeq 0.12$, which is the value inferred for 051221A \citep{Burrows+06,Soderberg+06}, as well as the typical opening angle required to reconcile the observed SGRB rate with the best-bet NS-NS merger rate \\ref{sec:GRB}) ), then up to $f_{\rm opt}
+\sim 0.1$ of GW events will be accompanied by potentially detectable optical afterglows."
+" This result is consistent with the rate of a few afterglows per vear mferred by Cowardetal.(2011) for their assmued total ALIGO/Virgo merecr rate of ~135 vrἩ,", This result is consistent with the rate of a few afterglows per year inferred by \citet{Coward+11} for their assumed total ALIGO/Virgo merger rate of $\sim 135$ $^{-1}$.
+ On the other laud. if O; is much larger. 2Ub fee. as found for 0050721 by Grupeet 2006)). then fopr is of order uuity. but the overall CW event rate may be lower than the best-bet ALIGO/Vireo rate.," On the other hand, if $\bar{\theta_j}$ is much larger, $\gtrsim 0.4$ (e.g., as found for 050724 by \citealt{Grupe+06}) ), then $f_{\rm opt}$ is of order unity, but the overall GW event rate may be lower than the best-bet ALIGO/Virgo rate."
+" Bevond cousiderations of depth and cadeuce. a ""uique optical identification of CW events also requires discrimination between off-axis afterelows aud potential contanunauts."," Beyond considerations of depth and cadence, a unique optical identification of GW events also requires discrimination between off-axis afterglows and potential contaminants."
+ We discuss this issue in refsec:compare.., We discuss this issue in \\ref{sec:compare}.
+ NS-NS/NS-BIT imiersers inav also be accompaied by non-thermal radio afterglow cuussion. which cau originate either from the ultravelativistic jet (as in the case of the optical afterglow). or from more spherical. sub-relativistic ejecta (Nakar&Piran901111 hereafter NPIL).," NS-NS/NS-BH mergers may also be accompanied by non-thermal radio afterglow emission, which can originate either from the ultra-relativistic jet (as in the case of the optical afterglow), or from more spherical, sub-relativistic ejecta \citealt{Nakar&Piran11}; hereafter NP11)."
+ The latter includes matter ejected dynamically during the merger process (“tidal tails). or in outflows from the accretion disk (see Figure 1)).," The latter includes matter ejected dynamically during the merger process (“tidal tails”), or in outflows from the accretion disk (see Figure \ref{fig:cartoon}) )."
+ Adopting standard models for svuchrotron emission frou. a relativistic shock. NP11 estimate that the peals racio brightuess for these cases is: where Jog= 04/026. 14 is the observing frequency in Guz: aud dj=200d599 Alpe is the luminosity cistance. again normalized to the ALIGO/Virgo rauge for NS-NS morecrs.," Adopting standard models for synchrotron emission from a relativistic shock, NP11 estimate that the peak radio brightness for these cases is: where $\beta_{0.2}=v_{\rm ej}/0.2c$ , $\nu_1$ is the observing frequency in GHz; and $d_{L} = 200 d_{L,200}$ Mpc is the luminosity distance, again normalized to the ALIGO/Virgo range for NS-NS mergers."
+ Equation 5 also assumes characteristic values of p=2.5 for the electron distribution power law index. aud €.=€p0.1 for the fractions of enerev density impartec to relativistic clectrous aud maenetic fields. respectively.," Equation \ref{eqn:Fp} also assumes characteristic values of $p=2.5$ for the electron distribution power law index, and $\epsilon_{e}=
+\epsilon_{B}=0.1$ for the fractions of energy density imparted to relativistic electrons and magnetic fields, respectively."
+ The radio emission peaks at the deceleration time: The peak brightness depends scusitively on both the properties of the ejecta CE and 0) and on the cireumburst density., The radio emission peaks at the deceleration time: The peak brightness depends sensitively on both the properties of the ejecta $E$ and $\beta$ ) and on the circumburst density.
+ As we diseuss du detail below. the realistic detection threshold for à couvineing detection with the EVLA (even with ~30 hr per epoch) is about 0.5 iis.," As we discuss in detail below, the realistic detection threshold for a convincing detection with the EVLA (even with $\sim 30$ hr per epoch) is about 0.5 mJy."
+ This requirement therefore defines a figure of merit for a radio detection of: With the exception of the velocity paramcter. this figure of merit is identical to the case of off-axis optical afterglows in teris of the dependence ou Ej aud à.," This requirement therefore defines a figure of merit for a radio detection of: With the exception of the velocity parameter, this figure of merit is identical to the case of off-axis optical afterglows in terms of the dependence on $E_j$ and $n$."
+" For quasi-spherical ejecta. a characteristic mass of Ma~LO? AL. in tidal tails or disk winds has an EzmoMaenj2~10""10. ere for the expected range of velocities?)~ο0."," For quasi-spherical ejecta, a characteristic mass of $M_{\rm ej}\sim
+10^{-2}$ $_\odot$ in tidal tails or disk winds has an $E\approx M_{\rm ej}v_{\rm ej}^{2}/2\sim 10^{50}-10^{51}$ erg for the expected range of $\beta\sim
+0.1-0.3$."
+" This results in at most FOAM,G2Ola,ÜN requiring ay21 ? for a detection."," This results in at most $FOM_{\rm rad}\approx 0.4\,
+n_0^{7/8}$, requiring $n_0\gtrsim 1$ $^{-3}$ for a detection."
+ For more typical densities of =0.1 cem5m associated. with SCRBs (Bergeretal.2005:Soderbergetal. 2006).. the radio enuüssion frou quasi-spherical ejecta will be essentially undetectable uuless the energy seale is uch laveer than ~1073 ere (Figure 6)).," For more typical densities of $\lesssim 0.1$ $^{-3}$ associated with SGRBs \citep{Berger+05,Soderberg+06}, the radio emission from quasi-spherical ejecta will be essentially undetectable unless the energy scale is much larger than $\sim 10^{51}$ erg (Figure \ref{fig:fom}) )."
+ Equations 5- 7. can also be applied to the case of off-axis afterglow cuuission using c1d (Nakar&Pivan2011) along with values for the jet energv aud cireuniburst density inferred from the optical afterglow data (POALpronZ0.1: Equation 1)).," Equations \ref{eqn:Fp}- \ref{eqn:fomr} can also be applied to the case of off-axis afterglow emission using $\beta\approx 1$ \citep{Nakar&Piran11} along with values for the jet energy and circumburst density inferred from the optical afterglow data $FOM_{\rm opt,on}\lesssim 0.1$; Equation \ref{eqn:fomo}) )."
+ In Figure 6 we plot the region of Lo phase-space that is accessible to radio detections (FOAM2 0.2)., In Figure \ref{fig:fom} we plot the region of $E-n$ phase-space that is accessible to radio detections $FOM_{\rm rad}\gtrsim 0.2$ ).
+ As cau be seen from the Figure. none of the existingad SCRB optical afterglow intersect this region. indicating that radio detections of off-axis afterelows are likely to be rare despite the overall isotropy of the signal.," As can be seen from the Figure, none of the existing SGRB optical afterglow intersect this region, indicating that radio detections of off-axis afterglows are likely to be rare despite the overall isotropy of the signal."
+ We now address in detail the estimated mininuuu radio brightuess necessary for a successful detection., We now address in detail the estimated minimum radio brightness necessary for a successful detection.
+ Although faint radio enussion is in principle detectable with a deep integration. a significant challenge is the sinall feld of view of sensitive instruments such as the EVLA (z0.1 deg? at 1 GIIz) requiring ~100200 poiutings to cover a typical CAV error region of teus of square deerees.," Although faint radio emission is in principle detectable with a deep integration, a significant challenge is the small field of view of sensitive instruments such as the EVLA $\approx 0.4$ $^2$ at 1 GHz), requiring $\sim 100-200$ pointings to cover a typical GW error region of tens of square degrees."
+ Tarecting individual galaxies within the error region does not decrease the number of required pointings since there are 100 ealaxies with LoOL within a typical error region (to 200 Mpc).," Targeting individual galaxies within the error region does not decrease the number of required pointings since there are $\sim
+400$ galaxies with $L\gtrsim 0.1\,L^*$ within a typical error region (to 200 Mpc)."
+ Even with oulv 10 miu per pointing. 30 hrper epoch will be required to cover the full error already a substantial allocation of EVLA time.," Even with only 10 min per pointing, $\sim 30$ hrper epoch will be required to cover the full error , already a substantial allocation of EVLA time."
+ Multiple, Multiple
+Unusual spectral. polarization. and Uuctuation properties of the precursors make them puzzling.,"Unusual spectral, polarization, and fluctuation properties of the precursors make them puzzling."
+— To the best of our. knowledge. no attempts have previously. been aimed at explaining the physies of this phenomenon.," To the best of our knowledge, no attempts have previously been aimed at explaining the physics of this phenomenon."
+ Recently Dyvks.Zhang&Cil(2005) have suggested. a geometrical model for the profile of PSI. 1822-09., Recently \citet*{d05} have suggested a geometrical model for the profile of PSR B1822-09.
+ These authors assume that the main pulse and the precursor originate. independently at dillerent locations in. the magnetosphere ancl the precursor. emission. intermittently reverses its direction to form the interpulse., These authors assume that the main pulse and the precursor originate independently at different locations in the magnetosphere and the precursor emission intermittently reverses its direction to form the interpulse.
+ In that model. the mechanism of reversal of the emission direction remains obscure. but for any conceivable switching mechanism it is principally clillieult to explain its dependence on the main pulse intensity.," In that model, the mechanism of reversal of the emission direction remains obscure, but for any conceivable switching mechanism it is principally difficult to explain its dependence on the main pulse intensity."
+ In the present paper. we for the first time propose a physical mechanism of the precursor. formation.," In the present paper, we for the first time propose a physical mechanism of the precursor formation."
+ La our model. the precursor arises as a result of induced scattering of the main pulse emission. into the background by the particles of the ultrarelativistic highhy maenetizecl plasma ofa pulsar.," In our model, the precursor arises as a result of induced scattering of the main pulse emission into the background by the particles of the ultrarelativistic highly magnetized plasma of a pulsar."
+ Our mechanism naturally explains the observed polarization. spectral. and [uctuation properties of the precursor emission as well as suggests its connection to the main pulse.," Our mechanism naturally explains the observed polarization, spectral, and fluctuation properties of the precursor emission as well as suggests its connection to the main pulse."
+ Pulsar radio emission is generated deep in the magnetosphere inside of the open field line tube., Pulsar radio emission is generated deep in the magnetosphere inside of the open field line tube.
+ Lence. it originates ancl propagates in the Dow of the ultrarelativistic clectron-positron plasma. which streams alone the open magnetic lines.," Hence, it originates and propagates in the flow of the ultrarelativistic electron-positron plasma, which streams along the open magnetic lines."
+" As the brightness temperatures of the radio emission are extremely high. Tg~107%10°"" IK. the waves may be subject to elflicient induced scattering oll the plasma particles."," As the brightness temperatures of the radio emission are extremely high, $T_B\sim 10^{25}-10^{30}$ K, the waves may be subject to efficient induced scattering off the plasma particles."
+" According to the radio emission theories. based on the plasma instabilities. the frequeney of the generated waves is close to the local Lorentz-shifted: proper plasma frequency. Ww wp where wy=vn,c?£m. n, is the plasma number density.. ο and n are the electron. charge and mass. and 5 is the plasma Lorentz-Factor. (but.seeMelrose&Geclalin1999.forthecriticismofthis point).."," According to the radio emission theories based on the plasma instabilities, the frequency of the generated waves is close to the local Lorentz-shifted proper plasma frequency, $\omega\sim\omega_p\sqrt{\gamma}$ , where $\omega_p\equiv\sqrt{4\pi
+n_ee^2/m}$, $n_e$ is the plasma number density, $e$ and $m$ are the electron charge and mass, and $\gamma$ is the plasma Lorentz-factor \citep[but see][for the criticism of this
+point]{gm99}."
+ In the vicinity of the emission. region. where the above condition is still valid. induced scattering olf the plasma particles is a collective process.," In the vicinity of the emission region, where the above condition is still valid, induced scattering off the plasma particles is a collective process."
+ The transverse waves are involved. in the induced. three-wave interactions (Luo&Melrose 2006)., The transverse waves are involved in the induced three-wave interactions \citep{lm06}.
+. X particular case of inducecl Ramanscattering in application to the pulsar magnetosphere has been considered by Gangadhara&Ixrishan(1993) and Lyutikoy(1998)., A particular case of induced Ramanscattering in application to the pulsar magnetosphere has been considered by \citet{gk93} and \citet{l98}.
+. As the plasma number density decreases with distance from the neutron star along with the magnetic field strength. προBoxer well above the emission. region w and the collective cllects become negligible.," As the plasma number density decreases with distance from the neutron star along with the magnetic field strength, $n_e\propto
+B\propto r^{-3}$, well above the emission region $\omega\gg\omega_p\sqrt{\gamma}$ and the collective effects become negligible."
+ The cuexternal magnetic field. significantly allects the scattering process on condition that the radio wave frequency in the particle rest. frame is much. less than the electron gvrolrequeney. a!xwe;=eB/me.," The external magnetic field significantly affects the scattering process on condition that the radio wave frequency in the particle rest frame is much less than the electron gyrofrequency, $\omega^\prime\ll\omega_G\equiv eB/mc$."
+ This condition is valid up to the radius of evelotron resonance. which lies in the outer magnetosphere.," This condition is valid up to the radius of cyclotron resonance, which lies in the outer magnetosphere."
+ In the present paper. we examine the induced. scattering. which takes place well above the emission region and well below the evelotron resonance radius.," In the present paper, we examine the induced scattering, which takes place well above the emission region and well below the cyclotron resonance radius."
+ Then the magnetized induced Compton scattering is à single-particle process ancl the incident waves are approximately transverse. clectromagnetic waves polarized either in the plane of the wavevector. and. the ambient magnetic field or perpendicularly to this plane., Then the magnetized induced Compton scattering is a single-particle process and the incident waves are approximately transverse electromagnetic waves polarized either in the plane of the wavevector and the ambient magnetic field or perpendicularly to this plane.
+ In application to. pulsar magnetosphere. induced scattering in a superstrong magnetic field has first. been considered by Blandford&Sceharlemann(1976). anel found to be ellicient.," In application to pulsar magnetosphere, induced scattering in a superstrong magnetic field has first been considered by \citet{bs76} and found to be efficient."
+ Later on the process has been suggested. to explain a number of phenomena in pulsar radio emission (Lyubarskit&Petrova1996:2004a.b).," Later on the process has been suggested to explain a number of phenomena in pulsar radio emission \citep{lp96,p04a,p04b}."
+. In. the present paper. we consider the pulsar radio beam scattering into the background. and particularly concentrate on the erowth of the scattered. component. which is identified with the precursor component of the pulse profile.," In the present paper, we consider the pulsar radio beam scattering into the background and particularly concentrate on the growth of the scattered component, which is identified with the precursor component of the pulse profile."
+ In the preceding literature on the induced scattering in a superstrong magnetic field. the kinetic equation for photon occupation numbers is derived from. an analysis of the scattering by a single relativistic electron and does. not include the particle cüstribution function explicitly. so that it directly. corresponds to the cold plasma case.," In the preceding literature on the induced scattering in a superstrong magnetic field, the kinetic equation for photon occupation numbers is derived from an analysis of the scattering by a single relativistic electron and does not include the particle distribution function explicitly, so that it directly corresponds to the cold plasma case."
+ In the present oer. we generalize the kinetic equation for the more realistic case of a hot plasma.," In the present paper, we generalize the kinetic equation for the more realistic case of a hot plasma."
+ The corresponding formalism js also been outlined in Blandford&Scharlemann(1976).. out. we shall go through the derivation once more in order ο Correct a slight error in that. paper.," The corresponding formalism has also been outlined in \citet{bs76}, but we shall go through the derivation once more in order to correct a slight error in that paper."
+ In the approximation of an infinitely strong magnetic ield. the scattering cross-section in the electron frametakes he form where ris the classical cleetron radius. 6 and 6j are the propagation angles of the incident ancl scattered photons. respectively. OY is the elementary. solid. angle [or the scattered. photons. and the primes denote. the quantities of the electron rest frame.," In the approximation of an infinitely strong magnetic field, the scattering cross-section in the electron frametakes the form where $r_e$ is the classical electron radius, $\theta^\prime$ and $\theta_1^\prime$ are the propagation angles of the incident and scattered photons, respectively, $\rm d\Omega_1^\prime$ is the elementary solid angle for the scattered photons, and the primes denote the quantities of the electron rest frame."
+ The cross-section (1) corresponds to the scattering between the photons with the ordinary polarization. i.c with the electric vectors lving in he plane of the ambient magnetic field.," The cross-section (1) corresponds to the scattering between the photons with the ordinary polarization, i.e with the electric vectors lying in the plane of the ambient magnetic field."
+ The scattering involving the extraordinary polarization states is negligible. since any perturbecl motion of a particle (in the field. of he incident wave) perpendicular to the ambient. magnetic 101 is suppressed.," The scattering involving the extraordinary polarization states is negligible, since any perturbed motion of a particle (in the field of the incident wave) perpendicular to the ambient magnetic field is suppressed."
+ Let us consider the scattering in. the aboratory [rame between the two photon states. & and ky. involving the electrons with the momenta p and p|dp along he magnetic field.," Let us consider the scattering in the laboratory frame between the two photon states, $\bmath{k}$ and $\bmath
+{k_1}$, involving the electrons with the momenta $p$ and $p+\delta
+p$ along the magnetic field."
+ In the scattering act. the momentun xwallel to the magnetic field is conserved: The probability of generating spontaneously. scattered photons per electron per unit time is given by whereafk) is the photon occupation. number. Jcos8. 3 ds the particle velocity in units of e. and the argument of the delta-function signifies the equality of the," In the scattering act, the momentum parallel to the magnetic field is conserved: The probability of generating spontaneously scattered photons per electron per unit time is given by where$n(\bmath k)$ is the photon occupation number, $\eta\equiv
+1-\beta\cos\theta$ , $\beta$ is the particle velocity in units of $c$ , and the argument of the delta-function signifies the equality of the"
+power law is ruled out from the fits. unless the gas-to-dust ratio is very low.,"power law is ruled out from the fits, unless the gas-to-dust ratio is very low."
+ The broken power law high energy slope was fixed at à spectral index of 3=1.0 as found in its to the X-ray spectrum alone., The broken power law high energy slope was fixed at a spectral index of $\beta=1.0$ as found in fits to the X-ray spectrum alone.
+ The lower energy. slope was fixed at;=0.5 as expected. for a cooling break in he standard fireball model (Sari.Piran.&Naravan1998)., The lower energy slope was fixed at $\beta=0.5$ as expected for a cooling break in the standard fireball model \citep{Sari}.
+. Below z&LS the fit becomes unacceptably poor (P<2% Yom a X7 test)., Below $z\approx1.8$ the fit becomes unacceptably poor $P<2$ from a $\chi^2$ test).
+ The region shown is curtailed at the upper end at 2=4 based on the limits from the optical spectrum. discussed below., The region shown is curtailed at the upper end at $z=4$ based on the limits from the optical spectrum discussed below.
+ The vertical extent of the region represents he error range on chy., The vertical extent of the region represents the error range on $A_V$.
+ We note that the hatched area is essentially consistent with the requirement of a reddened afterglow (above the solid curve)., We note that the hatched area is essentially consistent with the requirement of a reddened afterglow (above the solid curve).
+ Phe part of the hatched region which does not coincide with the grey shaded. band corresponds to best fitting models where the curvature of the N-ray spectrum is partly due to a cooling break. hence reducing the required absorption.," The part of the hatched region which does not coincide with the grey shaded band corresponds to best fitting models where the curvature of the X-ray spectrum is partly due to a cooling break, hence reducing the required absorption."
+ In fact. the hardness ratio for the orbit 2 spectrum is consistent with that for orbit 3 and also for the combined spectrum: beyond. orbit. 3. providing no evidence for a moving spectral break. but the statistics are too poor to make a firm statement.," In fact, the hardness ratio for the orbit 2 spectrum is consistent with that for orbit 3 and also for the combined spectrum beyond orbit 3, providing no evidence for a moving spectral break, but the statistics are too poor to make a firm statement."
+ We could. of course. have allowed more freedom in the moclel fitting. for example by not fixing the gas-to-dust ratio.," We could, of course, have allowed more freedom in the model fitting, for example by not fixing the gas-to-dust ratio."
+ However. allowing this to be a free parameter finds models with lower ccolumn at a given extinction. whereas. for long CRB afterelows. the columns determined fron N-ravs are more often in excess of those determined. from the reddening (e.g.Starlingetal.," However, allowing this to be a free parameter finds models with lower column at a given extinction, whereas, for long GRB afterglows, the columns determined from X-rays are more often in excess of those determined from the reddening \citep[e.g.][]{Starling}."
+2007)... A Larger than expected. cooling break (ASον 0.5) could. in principle reduce the required extinction. but then we would not explain the reddening of the afterglow.," A larger than expected cooling break $\Delta\beta>0.5$ ) could in principle reduce the required extinction, but then we would not explain the reddening of the afterglow."
+ In any case. even with greater [freedom dt is hard to find any reasonable solution at ο«1.5.," In any case, even with greater freedom it is hard to find any reasonable solution at $z<1.5$."
+ We conclude that the burst likely occurred at 2L8. which. combined with the constraint +<2.8. requires 2x:Ae <5.," We conclude that the burst likely occurred at $z\gsim1.8$, which, combined with the constraint $z\lsim2.8$, requires $2\lsim A_V\lsim5$ ."
+ This extinction would be high by typical CRB standards. but quite moderate compared to sight lines close the plane of a disk galaxy or through large clust-enshrouclec star-forming regions.," This extinction would be high by typical GRB standards, but quite moderate compared to sight lines close the plane of a disk galaxy or through large dust-enshrouded star-forming regions."
+ We note that. whilst not common. optically dark bursts in blue galaxies at moderate redshifts. have been identifie before. for example. GRB 970828 (Djorgovskictal.2001).. CRB 000210. (CGorosabeletal.2003).. GRB 051022. (Ro and GRB 070306 (Claunsenetal.2005).," We note that, whilst not common, optically dark bursts in blue galaxies at moderate redshifts, have been identified before, for example, GRB 970828 \citep{Djorgovski01}, GRB 000210 \citep{Gorosabel03}, GRB 051022 \citep{Rol07} and GRB 070306 \citep{Jaunsen08}."
+. Nevertheless. it remains odd that the location of the burs in this case places it within 0.2 aresec. corresponding to at most a kiloparsec or two (at. any. plausible redshif of the optically brightest. part of the host.," Nevertheless, it remains odd that the location of the burst in this case places it within $0.2$ arcsec, corresponding to at most a kiloparsec or two (at any plausible redshift), of the optically brightest part of the host."
+ This strong spatial coincidence between burst position and the optically brightest regions of their hosts is seen in many other low- long GRBs (Fruchteretal. 200G6).. and may be a consequence of the very short. life-times of massive-star GRB progenitors. which do not move far [from the star-forming region of their birth (Larssonetal.," This strong spatial coincidence between burst position and the optically brightest regions of their hosts is seen in many other low-reddening long GRBs \citep{Fruchter06}, , and may be a consequence of the very short life-times of massive-star GRB progenitors, which do not move far from the star-forming region of their birth \citep{Larsson}."
+2007).. In the case of GRB 060923. unless there is some separation along the line-of-sight. we require the dust-attenuated GIU to be close to a relatively unreddened region that dominates the optical light of the galaxy.," In the case of GRB 060923A, unless there is some separation along the line-of-sight, we require the dust-attenuated GRB to be close to a relatively unreddened region that dominates the optical light of the galaxy."
+ A plausible &eometry. is one in which high optical-cdepth molecular clouds provide patchy obscuration of a Iarge star-forming complex2004).. and the sight-line to the GiB happens to intersect one of these.," A plausible geometry is one in which high optical-depth molecular clouds provide patchy obscuration of a large star-forming complex, and the sight-line to the GRB happens to intersect one of these."
+ There is great interest in GRBs with very. red. optical-nlhR colours since they could be at. very high. redshift., There is great interest in GRBs with very red optical-nIR colours since they could be at very high redshift.
+". GI 060923,X is arguably the most extreme example found. to-date. being detected in the Ix-band. but with only. deep limits in carly observations in all bluer filters."," GRB 060923A is arguably the most extreme example found to-date, being detected in the K-band, but with only deep limits in early observations in all bluer filters."
+ LE purely due to à Lya break in the near-LR. this would indicate a redshift’ beyond. 2~Ll.," If purely due to a $\alpha$ break in the near-IR, this would indicate a redshift beyond $z\sim11$."
+ However. our later-time optical observations revealed a faint galaxy. presumably. the host. at the same position. which must be at a more moderate redshift. probably 2<2.8 but certainly not above zzd.," However, our later-time optical observations revealed a faint galaxy, presumably the host, at the same position, which must be at a more moderate redshift, probably $z\lsim2.8$ but certainly not above $z\approx4$."
+ The morphology of the host: a bright knot coincident with the GRB itself and extended low surface brightness features may indicate a merger/interaction has produced this burst of star formation.," The morphology of the host: a bright knot coincident with the GRB itself, and extended low surface brightness features may indicate a merger/interaction has produced this burst of star formation."
+ A combined. analysis of the X-ray and optical/nllk data suggest that the burst is also likely to have 2LS in order to reconcile the absorption required in both bands., A combined analysis of the X-ray and optical/nIR data suggest that the burst is also likely to have $z\gsim1.8$ in order to reconcile the absorption required in both bands.
+ There is. of course. a slim chance that the GRB is coincident with an unrelated foreground. galaxy.," There is, of course, a slim chance that the GRB is coincident with an unrelated foreground galaxy."
+ Following the analysis of Piroetal.(2002) we calculate a probability that the afterglow would be found. coincident with an unrelated ealaxy as bright as 2=25.6 to be about a small but non-negligible figure., Following the analysis of \citet{Piro} we calculate a probability that the afterglow would be found coincident with an unrelated galaxy as bright as $R=25.6$ to be about – a small but non-negligible figure.
+ However given the close alignment of the afterglow with the brightest knot of the galaxy. and that the colour and magnitude of the putative host are otherwise very. plausible for a typical long-duration GRB. the conservative explanation remains that this is a case of dust rather than distance.," However given the close alignment of the afterglow with the brightest knot of the galaxy, and that the colour and magnitude of the putative host are otherwise very plausible for a typical long-duration GRB, the conservative explanation remains that this is a case of dust rather than distance."
+ (πο 060923 nonetheless is likely to be representative of some proportion of the dark GRB population. which optical surveys are biased against finding., GRB 060923A nonetheless is likely to be representative of some proportion of the dark GRB population which optical surveys are biased against finding.
+ It is therefore also a good example of the kind of interlopers which we must be able to reject in order to identifv very high redshift GRBs., It is therefore also a good example of the kind of interlopers which we must be able to reject in order to identify very high redshift GRBs.
+ This studs emphasises that carly. deep. photometry in a range of optical ancl nl filters is essential to reliably identify candidates. ancl followup spectroscopy is highly desirable where possible.," This study emphasises that early, deep photometry in a range of optical and nIR filters is essential to reliably identify candidates, and followup spectroscopy is highly desirable where possible."
+ We thank the Ulx Science ancl Technology Facilities Council or financial support. in particular NICE lor a Senior tescarch Fellowship and AJL for a Postdoctoral Fellowship.," We thank the UK Science and Technology Facilities Council for financial support, in particular NRT for a Senior Research Fellowship and AJL for a Postdoctoral Fellowship."
+ The research activities of JG are supported by the Spanish Ministry. of Science. through the progranunes ESP2005-τς00-09 anc AYAPO04-01515., The research activities of JG are supported by the Spanish Ministry of Science through the programmes ESP2005-07714-C03-03 and AYA2004-01515.
+ DAL acknowledges he Instrument Centre lor Danish Astrophysics., DM acknowledges the Instrument Centre for Danish Astrophysics.
+ 9) acknowledges support by ai Marie. Curie. Latra-European Fellowship within the 6th European Community Framework 'rosram. under contract number MISLHE-CE-2006-042001. and a Grant of Excellence from the Icelancdie Research Eund.," PJ acknowledges support by a Marie Curie Intra-European Fellowship within the 6th European Community Framework Program under contract number MEIF-CT-2006-042001, and a Grant of Excellence from the Icelandic Research Fund."
+ AO. acknowledges support from the Norwegian Research Council. grant. nr.," AOJ acknowledges support from the Norwegian Research Council, grant nr."
+ 166072.We also eratefully acknowledge the work of the wider tteam that makes this research possible., 166072.We also gratefully acknowledge the work of the wider team that makes this research possible.
+Let us imagine that the infall of dwarf galaxies and gas was really the dominating process for the buildiug-up of the ¢D halo aud the GCS.,Let us imagine that the infall of dwarf galaxies and gas was really the dominating process for the building-up of the cD halo and the GCS.
+ Tow many dyvarfs aud their transformed. gas would then have contributed to the ¢D halo elt aud how many GCs nueght belong to the ¢D halo?, How many dwarfs and their transformed gas would then have contributed to the cD halo light and how many GCs might belong to the cD halo?
+ NGC 1399 possesses about 5800 elobular clusters (sec Sect., NGC 1399 possesses about 5800 globular clusters (see Sect.
+ 3.2)., 3.2).
+ About 1300 of them would belong to the bulee. Mya21.5 nag (see Sect.," About 1300 of them would belong to the bulge, $M_{\rm V,gal} = -21.5$ mag (see Sect."
+ 3.3.1). if one asses an initial specific yequency of Sy=3.2. which is the mean value for the other ellipticals in the Fornax cluster. except NGC L10L aud NGC 1280.," 3.3.1), if one assumes an initial specific frequency of $S_N = 3.2$, which is the mean value for the other ellipticals in the Fornax cluster, except NGC 1404 and NGC 1380."
+ That means that 1500. OCs would belong o the ¢D halo aud its specific frequency would be about Sy=10+41., That means that 4500 GCs would belong to the cD halo and its specific frequency would be about $S_N = 10\pm1$.
+ Note that half of the total GCS (=2900 GCs) are assigned to the metal-poor peak around |Fo/HI| = —— 1.3 dex. aud therefore a least 1600 inetaliich GCs ([Fo/TI] = 0.6 dex) have to be explained by the infall scenario. if oue assunes that all 1300 remaining bulee GCs belong to the ποσαΊο sub-," Note that half of the total GCS $= 2900$ GCs) are assigned to the metal-poor peak around [Fe/H] $\simeq$ $-$ 1.3 dex, and therefore at least 1600 metal-rich GCs ([Fe/H] $\simeq$ $-$ 0.6 dex) have to be explained by the infall scenario, if one assumes that all 1300 remaining bulge GCs belong to the metal-rich sub-population."
+ Tow can dwarf galaxies account for such a high Sy:D, How can dwarf galaxies account for such a high $S_N$?
+ As presented in Sect., As presented in Sect.
+ 5. there are iaiuly three scenarios possible.," 5, there are mainly three scenarios possible."
+ Firstly. accreted gas-poor dwarfs possessed high GC frequencies themselves.," Firstly, accreted gas-poor dwarfs possessed high GC frequencies themselves."
+ In this case. the average Sy of all accreted dwarts aud GCs can have values between baud 22 depending on the iuitial couditious (see Table 5).," In this case, the average $S_N$ of all accreted dwarfs and GCs can have values between 4 and 22 depending on the initial conditions (see Table 5)."
+ Secondly. the infalhug ooeas of previously eas-rich cawarfs was effectively couverted iuto elobular clusters.," Secondly, the infalling gas of previously gas-rich dwarfs was effectively converted into globular clusters."
+ Regarding the starburst as au isolated eutitv its resulting systems of stars aud clusters eau have ον values between 10 and 90 (sec Table 6)., Regarding the starburst as an isolated entity its resulting systems of stars and clusters can have $S_N$ values between 40 and 90 (see Table 6).
+ Finally. the stripping of GCs from dwarf galaxies was more effective than the stripping of their field population.," Finally, the stripping of GCs from dwarf galaxies was more effective than the stripping of their field population."
+ That this is in xuucipal possible is indicated by the fact that the Sy value of the outer parts of galaxies that are primarily affected by stripping can be in the order of 30 (see Sect., That this is in principal possible is indicated by the fact that the $S_N$ value of the outer parts of galaxies that are primarily affected by stripping can be in the order of 30 (see Sect.
+ 5.1. case 1b).," 5.1, case 1b)."
+ Among these 3 possibilities the stripping of GCs from nAwart salaxies most probably plavs a minor role., Among these 3 possibilities the stripping of GCs from dwarf galaxies most probably plays a minor role.
+" Even : all 50 αιασ within the core radius of the istribution are renimnants. whose outer GC have bee1 stripped off. we calculate that maxiuallv some huucdred (ος have been captured bv this process. assuniünes an initia Sy=5.4, Sy=30 for the stripped stars and GCs, and a final Sy=3.0 for the remmaut (similar te the values for NGC. 1172. MeLaughlliu et al. 1991))."," Even if all 50 dE/dS0s within the core radius of the galaxy distribution are remnants, whose outer GCs have been stripped off, we calculate that maximally some hundred GCs have been captured by this process, assuming an initial $S_N = 5.5$, $S_N = 30$ for the stripped stars and GCs, and a final $S_N = 3.0$ for the remnant (similar to the values for NGC 4472, McLaughlin et al. \cite{mcla94b}) )."
+ Iu t1e following. we consider the case that at πιο 500 CC's have beeu What is the correct mixture of the two other processes that fulfill the following (1) The cD halo has been formed only by accreted aud rewly formed matter. aud its total luuinosity is Myp=21.65 ias (2) The speci&c frequency of the accreted and newly ornmed GCs with respect to the halo luminosity is 10. (=1500 CC's) (3) The [500 CC's in the ¢eD halo cousist of 2500 metal- (blue) aud 2000 metalrich (red) GCs (this implies hat the GCS of the bulge has 100 uetal-poor aud 900 netalaxich GCs} (1) GCS captured by accretion ai stripping of dwiurfs can only be Iu Table 7 we preseut the possible mixtures of the 3 Xrocesses. starting with cases for which CC accretion is dominant aud eudiug with cases in which most COCs have con formed from iufalliug gas.," In the following, we consider the case that at most 500 GCs have been What is the correct mixture of the two other processes that fulfill the following (1) The cD halo has been formed only by accreted and newly formed matter, and its total luminosity is $M_{\rm V,cD} = -21.65$ mag (2) The specific frequency of the accreted and newly formed GCs with respect to the halo luminosity is $S_N = 10$ , $= 4500$ GCs) (3) The 4500 GCs in the cD halo consist of 2500 metal-poor (blue) and 2000 metal-rich (red) GCs (this implies that the GCS of the bulge has 400 metal-poor and 900 metal-rich GCs) (4) GCs captured by accretion and stripping of dwarfs can only be In Table \ref{tmixgc} we present the possible mixtures of the 3 processes, starting with cases for which GC accretion is dominant and ending with cases in which most GCs have been formed from infalling gas."
+ Iu the first five cases we assumed that all metal-poor CC's were captured or stripped., In the first five cases we assumed that all metal-poor GCs were captured or stripped.
+ Asstuming a high Sy value or the accretion process (ον=9. cases 1 aud 2). the cluster formation efficiency (CFE) does not need to be as Heh as estimated for merecr aud starburst situations (see Table 6)).," Assuming a high $S_N$ value for the accretion process $S_N = 9$, cases 1 and 2), the cluster formation efficiency (CFE) does not need to be as high as estimated for merger and starburst situations (see Table \ref{tsimsf}) )."
+ ILoxcever. as discussed. in Sect.," However, as discussed in Sect."
+ 6. a high Sv requires a high accretion rate of cawarf galaxies. a steep initial slope of the faint eud of the ealaxy LE. and also very faint dwarf galaxies should lave possessed at least oue GC.," 6, a high $S_N$ requires a high accretion rate of dwarf galaxies, a steep initial slope of the faint end of the galaxy LF, and also very faint dwarf galaxies should have possessed at least one GC."
+ The faintest dwarf galaxies with a GCS observed so far are the Local Croup dSplis Foruax aud Sagittarius (My~12.5 mae)., The faintest dwarf galaxies with a GCS observed so far are the Local Group dSphs Fornax and Sagittarius $M_V \simeq -12.5$ mag).
+ The other way arouud. if starbursts frou stripped eas can produce a high Sy value (10<NÜ. cases 3-5) and have formed 2000 metal-rich GCs. the Sy for the accreted metal-poor GCs is in the order 5-6.," The other way around, if starbursts from stripped gas can produce a high $S_N$ value $40 < S_N < 80$, cases 3-5) and have formed 2000 metal-rich GCs, the $S_N$ for the accreted metal-poor GCs is in the order 5-6."
+ Such values can easily be achieved by the accretion scenario prescuted. in Sect., Such values can easily be achieved by the accretion scenario presented in Sect.
+ 6 under various reasonable initial couditious (see Table 5)., 6 under various reasonable initial conditions (see Table 5).
+ Tu the cases 6 to 8 we assumed that the majority of the GCs hac their origin from iufalliug gas with a low value of the estimated starburst CPEs (20<Sy I0)., In the cases 6 to 8 we assumed that the majority of the GCs had their origin from infalling gas with a low value of the estimated starburst CFEs $20 < S_N < 40$ ).
+ The specific requeney for the remaining 1500 accreted GCs then can be very low (8«Sy 5). very faint chwarts do not need to possess GCs. and the umubers of dissolved dwarts can be in the order of 250-500.," The specific frequency for the remaining 1500 accreted GCs then can be very low $3 < S_N < 5$ ), very faint dwarfs do not need to possess GCs, and the numbers of dissolved dwarfs can be in the order of 250-500."
+ ILowever. oue then has to assune that CCs formed from metalrich as well as metal-poor gas and that most of the original dwart ealaxies have been very gas-rich.," However, one then has to assume that GCs formed from metal-rich as well as metal-poor gas and that most of the original dwarf galaxies have been very gas-rich."
+ We have sununarized tle properties of the different coniponeuts of the ceutral ealaxy NGC 1399 (CCS. cD halo. bulge) and the dwarf ealaxy population iu the ceuter of the cluster.," We have summarized the properties of the different components of the central galaxy NGC 1399 (GCS, cD halo, bulge) and the dwarf galaxy population in the center of the cluster."
+ We have analysed under which circunstances the GCS and cD halo can be explained by the iufall and accretion of gas-poor as well as eas-rich dwarf galaxies., We have analysed under which circumstances the GCS and cD halo can be explained by the infall and accretion of gas-poor as well as gas-rich dwarf galaxies.
+and it does not feature the deep electron fraction. trough that develops behind the shock due to the copious electron captures.,and it does not feature the deep electron fraction trough that develops behind the shock due to the copious electron captures.
+ Furthermore. the neutrino parametrisation scheme does not feature any neutrino heating.," Furthermore, the neutrino parametrisation scheme does not feature any neutrino heating."
+ Hence. it is important to distinguish the results from the collapse phase and bounce. where we believe that the neutrino physics parametrisation is à viable and reasonably accurate approach. from the dynamies of the postbounce phase. where the neutrino physics parametrisation is not sufficiently accurate.," Hence, it is important to distinguish the results from the collapse phase and bounce, where we believe that the neutrino physics parametrisation is a viable and reasonably accurate approach, from the dynamics of the postbounce phase, where the neutrino physics parametrisation is not sufficiently accurate."
+ The data from the postbounce evolution in our models has to be understood as an idealised exploration of potential GW features after the bounce signal., The data from the postbounce evolution in our models has to be understood as an idealised exploration of potential GW features after the bounce signal.
+ This is currently the state-of-the-art in three-dimensional simulations and handled the same way in the models of Ottetal.(2007)., This is currently the state-of-the-art in three-dimensional simulations and handled the same way in the models of \citet{OttPhysRev07}.
+ In order to improve the models also in the postbounce phase. we have developed the isotropic diffusion. source approximation (IDSA) for the neutrino transport (Liebendórferetal.2007) for future models.," In order to improve the models also in the postbounce phase, we have developed the isotropic diffusion source approximation (IDSA) for the neutrino transport \citep{Liebendoerfer07}
+ for future models."
+ The scheme has been implemented and tested 1n spherical symmetry. but is not vet fully functional in three dimensions.," The scheme has been implemented and tested in spherical symmetry, but is not yet fully functional in three dimensions."
+ From basic considerations and comparison to the results in axisymmetry with accurate neutrino transport (Mülleretal.2004) we expect that the additional neutrino emission will lead to a more compact neutron star that is acereting matter from an envelope with stronger fluid asymmetries., From basic considerations and comparison to the results in axisymmetry with accurate neutrino transport \citep{M2004} we expect that the additional neutrino emission will lead to a more compact neutron star that is accreting matter from an envelope with stronger fluid asymmetries.
+ Compared to our current simulations. one could expect a stronger GW signal in the late postbounce phase with a shorter periodicity in the signal for comparable initial rotation rates of the progenitor.," Compared to our current simulations, one could expect a stronger GW signal in the late postbounce phase with a shorter periodicity in the signal for comparable initial rotation rates of the progenitor."
+ However. this will be explored in more detail with the next generation of postbounce supernova models.," However, this will be explored in more detail with the next generation of postbounce supernova models."
+ In this paper we presented the GW analysis of two 3D MHD core collapse supernova simulations. which differ only in the amount of initial rotation.," In this paper we presented the GW analysis of two 3D MHD core collapse supernova simulations, which differ only in the amount of initial rotation."
+ Model sI5h started with an angular velocity of =0.3 rad/s. model sl5g with rad/s. We incorporated progenitor stars from stellar evolution calculations. spherically symmetric GR. effects. the Lattimer-Swesty EoS. magnetic fields and a neutrino parametrisation scheme that is accurate until the first few ms after bounce.," Model s15h started with an angular velocity of =0.3 rad/s, model s15g with rad/s. We incorporated progenitor stars from stellar evolution calculations, spherically symmetric GR effects, the Lattimer-Swesty EoS, magnetic fields and a neutrino parametrisation scheme that is accurate until the first few ms after bounce."
+ The particular choice of the initial angular momentum allowed clear distinctions between gravitational wave features that stem from rotation or nonaxisymmetric motions., The particular choice of the initial angular momentum allowed clear distinctions between gravitational wave features that stem from rotation or nonaxisymmetric motions.
+ The amplitudes for the three directions and polarisations.Αν. are not very sensitive to rotation and show a similar size of several centimetres.," The amplitudes for the three directions and polarisations,, are not very sensitive to rotation and show a similar size of several centimetres."
+ Apparently. they initially couple only weakly to rotationally induced large scale asymmetries in the mass-energy distribution.," Apparently, they initially couple only weakly to rotationally induced large scale asymmetries in the mass-energy distribution."
+ The only GW amplitude that is strongly correlated to axisymmetric rotation at the time of core-bounce turns out to be in the π/2.. Ξ 0--direction.," The only GW amplitude that is strongly correlated to axisymmetric rotation at the time of core-bounce turns out to be in the /2, =0 -direction."
+ In model s15g. it exceeds all other amplitudes and in particular the corresponding one from model sI5h by more than one order of magnitude.," In model s15g, it exceeds all other amplitudes and in particular the corresponding one from model s15h by more than one order of magnitude."
+ It shows a clear type I characteristics and implies that a rotational core collapse stays axisymmetric in the early postbounce phase. as lately discussec in Ottetal.(2007)..," It shows a clear type I characteristics and implies that a rotational core collapse stays axisymmetric in the early postbounce phase, as lately discussed in \citet{2007CQGra..24..139O}."
+ At the bounce stage of the simulatio sl5g. the dominant band of emission is peaking around 893 Hz (s15g). which is in good agreement with the recent findings from Mülleretal.(2004).. but roughly 150 Hz higher tha in (Dimmelmeieretal.2007;Ott2007)..," At the bounce stage of the simulation s15g, the dominant band of emission is peaking around 893 Hz (s15g), which is in good agreement with the recent findings from \citet{M2004}, , but roughly 150 Hz higher than in \citep{Dimmelmeier07PhysRev,2007CQGra..24..139O}."
+ Within the investigated time window of about 70 ms duration. the channel in model 15g accounts during the bounce and ring-dow phases for 990% of the total energy release in GW emission.," Within the investigated time window of about 70 ms duration, the channel in model 15g accounts during the bounce and ring-down phases for 90 of the total energy release in GW emission."
+ The models of the later postbounce phase are yet atfectec by significant technical and physical uncertainties., The models of the later postbounce phase are yet affected by significant technical and physical uncertainties.
+ They shoulc be understood as a qualitative outlook to future models that will have to include better neutrino transport and a higher resolution of the protoneutron star., They should be understood as a qualitative outlook to future models that will have to include better neutrino transport and a higher resolution of the protoneutron star.
+ In the postbounce phase. nonaxisymmetric dynamics starts to play an important role.," In the postbounce phase, nonaxisymmetric dynamics starts to play an important role."
+ The wave trains are revived either through convective or through nonaxisymmetric instabilities in the protoneutron star., The wave trains are revived either through convective or through nonaxisymmetric instabilities in the protoneutron star.
+ The GW amplitudes from both models show a sustained signal of approximately constant size. but of different physical origin.," The GW amplitudes from both models show a sustained signal of approximately constant size, but of different physical origin."
+ In model s15g. the occurrence of a so-called low 7/|W| instability of dominant #=2 character around 20 ms post-bounce leads to the prolonged narrow band GW emission at a frequency of 905 Hz. which ts in good agreement with Ottetal.(2007) in all points. except their dominant mode ism =|.," In model s15g, the occurrence of a so-called low T/|W| instability of dominant m=2 character around 20 ms post-bounce leads to the prolonged narrow band GW emission at a frequency of 905 Hz, which is in good agreement with \citet{2007CQGra..24..139O} in all points, except their dominant mode is m=1."
+ On the other hand. the onset of nonaxisymmetric dynamics in model sI5h is triggered by convective motions due to a negative entropy gradient.," On the other hand, the onset of nonaxisymmetric dynamics in model s15h is triggered by convective motions due to a negative entropy gradient."
+ In our current models. the characteristic frequency of emission linked to convective features takes place below 250 Hz.," In our current models, the characteristic frequency of emission linked to convective features takes place below 250 Hz."
+ This number may require adaption after the inclusion of more accurate neutrino transport., This number may require adaption after the inclusion of more accurate neutrino transport.
+ Further. no sign of a decay. be it due to a dynamical instability or convection. is present in the long term evolution of the wave trains. as 1t was already observed in (Mülleretal.2004:Ott2007)..," Further, no sign of a decay, be it due to a dynamical instability or convection, is present in the long term evolution of the wave trains, as it was already observed in \citep{M2004,2007CQGra..24..139O}."
+ Finally we conclude that gravitational waves from the discussed dynamical features. namely the rotational core bounce and its subsequent ringdown and low 77|W] instability related to a core collapse supernova like model s15g could possibly be detected by a current LIGO instrument within a distance of 10 kpe.," Finally we conclude that gravitational waves from the discussed dynamical features, namely the rotational core bounce and its subsequent ringdown and low T/|W| instability related to a core collapse supernova like model s15g could possibly be detected by a current LIGO instrument within a distance of 10 kpc."
+ For model sI5h it seems more likely that LIGO would only catch signals from later stages of the supernova evolution. meaning frequency distributions caused by convective motion.," For model s15h it seems more likely that LIGO would only catch signals from later stages of the supernova evolution, meaning frequency distributions caused by convective motion."
+ The future improvement of detector sensitivities in combination with several adjustable frequency responses of Advanced LIGO will give the opportunity to tune the instruments to a particular event/source and therefore allow to look much deeper into space., The future improvement of detector sensitivities in combination with several adjustable frequency responses of Advanced LIGO will give the opportunity to tune the instruments to a particular event/source and therefore allow to look much deeper into space.
+ This will enhance the chance of observing core collapse supernovae similarto our models up to distances of | Mpe., This will enhance the chance of observing core collapse supernovae similarto our models up to distances of 1 Mpc.
+"wwas observed in 2008 March and April. when CARMA was in its D configuration giving a ivpical svnthesized beam of 3.6x2.1"".","was observed in 2008 March and April, when CARMA was in its D configuration giving a typical synthesized beam of $\times$."
+. Theu-i coverage of the observations gives projected baselines of 3.79-95.03 kA (4.8-121.1 m)., The coverage of the observations gives projected baselines of 3.79-95.03 $\lambda$ (4.8-121.1 m).
+ All three correlator bands were set wp with 62 MIIz banclwidths. giving a resolution of ~1.2 km/s. Two of the bands were configured to overlap al the edges to cover theJ = 0-I*Lill line in the lower sideband andJ = 0-1 °Lill in the upper sideband.," All three correlator bands were set up with 62 MHz bandwidths, giving a resolution of $\sim$ 1.2 km/s. Two of the bands were configured to overlap at the edges to cover the = $^7$ LiH line in the lower sideband and = 0-1 $^6$ LiH in the upper sideband."
+ The remaining band was configured to observe theJ — 3-4 CO line in (he lower sideband. this choice was driven bv (he Irequency. configuration.," The remaining band was configured to observe the = 3-4 $^{13}$ CO line in the lower sideband, this choice was driven by the frequency configuration."
+ Neptune ancl AIWC349 were used as fIux density calibrators. 3C273 and 3454.3 were used to correct the passbands of each window. and 1733-130 and 1911-201 were used to calibrate the antenna based gains.," Neptune and MWC349 were used as flux density calibrators, 3C273 and 3C454.3 were used to correct the passbands of each window, and 1733-130 and 1911-201 were used to calibrate the antenna based gains."
+ The source was then sel[-calibrated to refine the solution., The source was then self-calibrated to refine the solution.
+ We obtained only 4.5 hours of good data due to poor weather., We obtained only 4.5 hours of good data due to poor weather.
+" wwas observed in 2008 August and September. when CARALA was also in its D configuration. eiving a typical synthesized beam of L.7x1.5""."," was observed in 2008 August and September, when CARMA was also in its D configuration, giving a typical synthesized beam of $\times$."
+. The u—e coverage of the observations gave projected baselines of 6.7-129.2 kA (1.1-147.1 m)., The $u-v$ coverage of the observations gave projected baselines of 6.7-129.2 $\lambda$ (7.7-147.1 m).
+ Two of the correlator bands were set up in 62 MlIz mode. giving ~1.1 km/s resolution. and the third was in 31 MlIz mode. eiving ~0.55 km/s resolution.," Two of the correlator bands were set up in 62 MHz mode, giving $\sim$ 1.1 km/s resolution, and the third was in 31 MHz mode, giving $\sim$ 0.55 km/s resolution."
+ Band 1 in the lower sideband was centered on theJ = 3-4 CO line., Band 1 in the lower sideband was centered on the = 3-4 $^{13}$ CO line.
+ Bands 2 and 3 in the lower sideband were set to observe the = 0-1 “Lill transition., Bands 2 and 3 in the lower sideband were set to observe the = 0-1 $^7$ LiH transition.
+" Dand Lin the upper sideband was set to observe theJ = 0-1 ""LilI transition.", Band 1 in the upper sideband was set to observe the = 0-1 $^6$ LiH transition.
+ The upper sideband ol the other two bands were for continuum., The upper sideband of the other two bands were for continuum.
+ Uranus was used as the flux clensitw calibrator. 3Cm8RA. 3454.3. and the internal noise source were used (o correct the passbands of each window.," Uranus was used as the flux density calibrator, 3C84, 3C454.3, and the internal noise source were used to correct the passbands of each window."
+ The data were sell-calibrated for (he antenna based gains., The data were self-calibrated for the antenna based gains.
+ We only obtained about seven hours οἱ good data on Chis source due to poor weather., We only obtained about seven hours of good data on this source due to poor weather.
+ Figure 1. shows our continuum map ofD0213—357., Figure \ref{fig:bo-contin} shows our continuum map of.
+. The two point sources (A and D) and the “hot spot are labeled aud the svnthesized beam is in the lower left corner., The two point sources (A and B) and the ”hot spot” are labeled and the synthesized beam is in the lower left corner.
+ The contour levels are zx126. t240. £360. ... στ 4.5 mJv/beam aud the peak fIux is 585 mJv/beam.," The contour levels are $\pm$ $\sigma$, $\pm$ $\sigma$, $\pm$ $\sigma$, ..., $\sigma$ = 4.5 mJy/beam and the peak flux is 885 mJy/beam."
+ Note that we do not detect the “hot spot (even at our 2.0 cutoff)., Note that we do not detect the ”hot spot” (even at our 3 $\sigma$ cutoff).
+ This peak fIux value is within (he uncertainties of the expected value at our rest Irequeney [rom Figure 2 of (Combes&Wi, This peak flux value is within the uncertainties of the expected value at our rest frequency from Figure 2 of \citep{combes97}.
+klind1997).. Figure 2. shows our 3o detection of * Lill. The noise level of the spectrum is denoted by the error bar and the dashed line denotes the expected rest velocity ol the line., Figure \ref{fig:bo-lih} shows our $\sigma$ detection of $^7$ LiH. The noise level of the spectrum is denoted by the error bar and the dashed line denotes the expected rest velocity of the line.
+ This line was detected toward the À component. where all other line detection have been mace (o date and is at the expected rest. velocity.," This line was detected toward the A component, where all other line detection have been made to date and is at the expected rest velocity."
+ For absorption features the, For absorption features the
+on the order hundreds of eV and L>1076 erg s! appear not to have been observed before.,on the order hundreds of eV and $L > 10^{36}$ erg $^{-1}$ appear not to have been observed before.
+ By now. observations 1n several nearby galaxies. including M31 eet al.," By now, observations in several nearby galaxies, including M31 et al."
+ 2003a) and M104 eet al., 2003a) and M104 et al.
+ 2003b) verify the result., 2003b) verify the result.
+ Finally. investigations have also yielded interesting information about extragalactic SSSs. providing some first steps toward answering the 7 questions posed in SStefano Kong (20034).," Finally, investigations have also yielded interesting information about extragalactic SSSs, providing some first steps toward answering the $7$ questions posed in Stefano Kong (2003a)."
+ We summarize this progress below. (, We summarize this progress below. (
+1)sources? It is too soon to claim an answer to this question.,1) It is too soon to claim an answer to this question.
+ We can say. however. the SSSs and QSSs seem to exist in virtually every galaxy. (," We can say, however, the SSSs and QSSs seem to exist in virtually every galaxy. ("
+Ongoing analyses of additional galaxies Jend further support to this conclusion.),Ongoing analyses of additional galaxies lend further support to this conclusion.)
+ We have so far carried out only a crude population analysis of the 4 galaxies studied here., We have so far carried out only a crude population analysis of the $4$ galaxies studied here.
+" In each case though. the data is consistent with a population of at least ~500 VSSs with L>10°"" ere s! active at the time of the observation."," In each case though, the data is consistent with a population of at least $\sim 500$ VSSs with $L > 10^{37}$ erg $^{-1}$ active at the time of the observation."
+ The true underlying population of such luminous sources could be several times larger., The true underlying population of such luminous sources could be several times larger.
+ Observations so far provide evidence for galactic populations of low-luminosity VSSs (L«1075 erg s! ). but do not constrain its size.," Observations so far provide evidence for galactic populations of low-luminosity VSSs $L < 10^{37}$ erg $^{-1}$ ), but do not constrain its size."
+ With regard to the potential of VSSs to ionize galactic ISMs. we note that even the most conservative estimate of VSS population for these galaxies yields more than 107 erg s~! of highly ionizing radiation released by SSSs in typical galaxies.," With regard to the potential of VSSs to ionize galactic ISMs, we note that even the most conservative estimate of VSS population for these galaxies yields more than $10^{40}$ erg $^{-1}$ of highly ionizing radiation released by SSSs in typical galaxies."
+ Although the luminosity is smaller than that associated with young stars. photons from VSSs may dominate at high energies. (," Although the luminosity is smaller than that associated with young stars, photons from VSSs may dominate at high energies. ("
+2)populations? Among the spiral galaxies in this study. there is not enough variation among the values of Mp to allow a search for correlations between My and the true size of the SSS/QSS population.,"2) Among the spiral galaxies in this study, there is not enough variation among the values of $M_B$ to allow a search for correlations between $M_B$ and the true size of the SSS/QSS population."
+ The ratio of the number of detected SSSs and QSSs to the total number of detected X-ray sources does vary systematically with inclination (39%. 37%. and 30% for ΜΟΙ. M83. and M51. respectively).," The ratio of the number of detected SSSs and QSSs to the total number of detected X-ray sources does vary systematically with inclination $39\%$, $37\%$, and $30\%$ for M101, M83, and M51, respectively)."
+ If we limit consideration to sources selected by the HR conditions. which seem to be the softest sources selected. the trend appears even stronger (17%. 11%. and 7% for MIOI. M83. and M51. respectively).," If we limit consideration to sources selected by the HR conditions, which seem to be the softest sources selected, the trend appears even stronger $17\%$, $11\%$, and $7\%$ for M101, M83, and M51, respectively)."
+ This might suggest similar sized populations in each of these galaxies. (, This might suggest similar sized populations in each of these galaxies. (
+3)populations? The existence of SSSs and QSSs in NGC 4697 demonstrates that there are SSSs in ellipticals. (,3) The existence of SSSs and QSSs in NGC 4697 demonstrates that there are SSSs in ellipticals. (
+See also Sarazin. Irwin. Bregman 2001.),"See also Sarazin, Irwin, Bregman 2001.)"
+ It is difficult to estimate the size of the underlying SSS/QSS populations in galaxies like NGC 4697. largely because we do not yet have good models for the characteristics of the SSS/QSS population in these galaxies.," It is difficult to estimate the size of the underlying SSS/QSS populations in galaxies like NGC 4697, largely because we do not yet have good models for the characteristics of the SSS/QSS population in these galaxies."
+ We have demonstrated. however. that. if the sources we detect in NGC 4697 are similar to the SSSs in the Milky Way and Magellanic Clouds that were used to define the class. its population of SSSs is large-about 1250 “classical” SSSs per detected example of a “classical” SSS in NGC 4697.," We have demonstrated, however, that, if the sources we detect in NGC 4697 are similar to the SSSs in the Milky Way and Magellanic Clouds that were used to define the class, its population of SSSs is large–about 1250 “classical” SSSs per detected example of a “classical” SSS in NGC 4697."
+ It is possible. however. that many of the 14 detected VSSs are examples of new types of sources.," It is possible, however, that many of the $14$ detected VSSs are examples of new types of sources."
+ Finally. we note that the data so far are consistent with the hypothesis that a significant portion of the diffuse soft emission found in elliptical galaxies may be due to unresolved SSSs (Fabbiano. Kim. Trinchieri 1994). (," Finally, we note that the data so far are consistent with the hypothesis that a significant portion of the diffuse soft emission found in elliptical galaxies may be due to unresolved SSSs (Fabbiano, Kim, Trinchieri 1994). ("
+4)disks? In the 3 spirals we have studied. fsss4oss. the fraction of SSSs and QSSs located within | kpe of the galaxy center is the same as fi. the fraction of all X-ray sources located within | kpe of the galaxy center.,"4) In the $3$ spirals we have studied, $f_{SSS+QSS}$, the fraction of SSSs and QSSs located within $1$ kpc of the galaxy center is the same as $f_{all}$, the fraction of all X-ray sources located within $1$ kpc of the galaxy center."
+ The values of fsss;oss Vary significantly among the spiral galaxies. however. (," The values of $f_{SSS+QSS}$ vary significantly among the spiral galaxies, however. ("
+5)holes? To answer this question. we have to rely on the dynamical measurements of the central supermassive BH.,"5) To answer this question, we have to rely on the dynamical measurements of the central supermassive BH."
+ M51 is a Seyfert 2 galaxy with a Jow-luminosity AGN (see. e.g.. Ho. Filippenko. Sargent 1997) and it is very likely that there is a supermassive. ~10M. (Hagiwara et al.," M51 is a Seyfert 2 galaxy with a low-luminosity AGN (see, e.g., Ho, Filippenko, Sargent 1997) and it is very likely that there is a supermassive, $\sim 10^7 M_{\odot}$ (Hagiwara et al."
+ 2001) BH in the center., 2001) BH in the center.
+ M83 might also harbor a ~107M.. BH (Thatte et al., M83 might also harbor a $\sim 10^7 M_{\odot}$ BH (Thatte et al.
+ 2000)., 2000).
+ The concentration factors for SSSs/QSSs in both M83 and MSI are high. approximately 50 for each galaxy.," The concentration factors for SSSs/QSSs in both M83 and M51 are high, approximately $50$ for each galaxy."
+ The concentration factors for all sources in each galaxy is comparable. but somewhat smaller.," The concentration factors for all sources in each galaxy is comparable, but somewhat smaller."
+ MIOI might have a <106M.. supermassive BH (Moody et al., M101 might have a $< 10^6 M_{\odot}$ supermassive BH (Moody et al.
+ 1995)., 1995).
+ The relative values of the concentration. factors among the spiral galaxies suggests that the concentration. of SSSs/QSSs. and in fact the concentration of all X-ray sources. scales with the mass of the central supermassive BH.," The relative values of the concentration factors among the spiral galaxies suggests that the concentration of SSSs/QSSs, and in fact the concentration of all X-ray sources, scales with the mass of the central supermassive BH."
+ It is interesting. but not at this point statistically significant. that the concentration factors are somewhat higher for SSSs/QSSs.," It is interesting, but not at this point statistically significant, that the concentration factors are somewhat higher for SSSs/QSSs."
+ NGC 4697 appears to have the highest-mass central BH of any of the galaxies in our sample (1.2«105M ..)., NGC 4697 appears to have the highest-mass central BH of any of the galaxies in our sample $1.2 \times 10^8 M_\odot$ ).
+ It is also the galaxy that has the highest concentration factor for all X-ray sources (Jj=140). the highest concentration factor for SSSs (CFessross= 360). and the largest difference between the central concentration of SSSs and other X-ray sources.," It is also the galaxy that has the highest concentration factor for all X-ray sources ${\cal F}_{all} = 140),$ the highest concentration factor for SSSs ${\cal F}_{SSS+QSS} = 360$ ), and the largest difference between the central concentration of SSSs and other X-ray sources."
+ From this it seems that interactions near the galaxy center may be playing an important role in creating X-ray sources. and that this is especially true for SSSs/QSSs. (," From this it seems that interactions near the galaxy center may be playing an important role in creating X-ray sources, and that this is especially true for SSSs/QSSs. ("
+An alternative hypothesis is that the characteristics of the stellar population near the center differ from those elsewhere in the galaxy.),An alternative hypothesis is that the characteristics of the stellar population near the center differ from those elsewhere in the galaxy.)
+ A variety of interactions could enhance the formation of X-ray binaries near the center of the galaxy: it may be that the additional enhancement for SSSs is due to the presence of the cores of stars that have been tidally disrupted by the central BH. (, A variety of interactions could enhance the formation of X-ray binaries near the center of the galaxy; it may be that the additional enhancement for SSSs is due to the presence of the cores of stars that have been tidally disrupted by the central BH. (
+6)clusters? In the 3 spiral galaxies we have studied. the SSSs seem to be concentrated in the spiral arms.,"6) In the $3$ spiral galaxies we have studied, the SSSs seem to be concentrated in the spiral arms."
+ Ongoing work on SSSs and QSSs in MIOI indicates that they may be preferentially located near HII regions and SNRs. suggesting that some SSSs/QSSs in MIOI may be related to young populations. (," Ongoing work on SSSs and QSSs in M101 indicates that they may be preferentially located near HII regions and SNRs, suggesting that some SSSs/QSSs in M101 may be related to young populations. ("
+7) Are SSSs significant contributors to the rates of Type la SNe?,7) Are SSSs significant contributors to the rates of Type Ia SNe?
+ It is too soon to say., It is too soon to say.
+ To use the data to answer this question. we need to estimate the fraction of the SSSs in each galaxy that are accreting WDs.," To use the data to answer this question, we need to estimate the fraction of the SSSs in each galaxy that are accreting WDs."
+ This ts difficult to do. even in the Magellanic Clouds.," This is difficult to do, even in the Magellanic Clouds."
+ There are. however. some signatures which can help.," There are, however, some signatures which can help."
+ First. the WD models developed to date invoke donor stars that too old to be concentrated in the spiral arms: the fraction of sources located away from the spiral arms and from regions with recent star formation can provide clues about the size of the NBWD population.," First, the WD models developed to date invoke donor stars that too old to be concentrated in the spiral arms; the fraction of sources located away from the spiral arms and from regions with recent star formation can provide clues about the size of the NBWD population."
+ Second. some of the accreting WDs that may be Type Ia progenitors should be surrounded by," Second, some of the accreting WDs that may be Type Ia progenitors should be surrounded by"
+To search for variations in the strength of the absorption features with orbital phase. we compare the depth of the strongest. features in the velocity. corrected: spectrum. of RAW Fri with the values observed in the the template star spectrum.,"To search for variations in the strength of the absorption features with orbital phase, we compare the depth of the strongest features in the velocity corrected spectrum of RW Tri with the values observed in the the template star spectrum."
+ We considered the combined regions around. the Na 1 absorption feature - 22140A)) and the Ca E absorption feature -- 22700A)). and mask out the rest of the spectrum.," We considered the combined regions around the Na I absorption feature - ) and the Ca I absorption feature - ), and mask out the rest of the spectrum."
+" Figure shows a plot of the ratio of RAV ""Iri versus template star absorption feature deficit through the orbital evele.", Figure \ref{f8} shows a plot of the ratio of RW Tri versus template star absorption feature deficit through the orbital cycle.
+ Phere is some evidence that the secondary features are strongest near phase zero and weaker near phase 0.5., There is some evidence that the secondary features are strongest near phase zero and weaker near phase 0.5.
+ Taken at face value. this suggests that the centroic of the secondary: features is shifted to the hemisphere of the secondary. that [aces away from the disc.," Taken at face value, this suggests that the centroid of the secondary features is shifted to the hemisphere of the secondary that faces away from the disc."
+ Although the effect is marginal. it is in accordance with what we expect due to heating ellects.," Although the effect is marginal, it is in accordance with what we expect due to heating effects."
+ The best sinusoidal fit to the data has an aniplitucde of 27£0.18 (solid line in Figure 7))., The best sinusoidal fit to the data has an amplitude of $0.27\pm0.18$ (solid line in Figure \ref{f8}) ).
+ Vhis sinusoidal fit to Ye data implies that the secondary star contributes ~39% of the Ix-band. πας at phase 0.0. while at. phase 0.5 this »ercentage is reduced to 15%.," This sinusoidal fit to the data implies that the secondary star contributes $\sim39\%$ of the K-band flux at phase 0.0, while at phase 0.5 this percentage is reduced to $\sim15\%$."
+ Hence. the absorption in 1e hemisphere of the secondary star nearest to the primary star is O44 times the strength. of the absorption in the jemisphere facing away from the primary.," Hence, the absorption in the hemisphere of the secondary star nearest to the primary star is $\sim0.4$ times the strength of the absorption in the hemisphere facing away from the primary."
+ Averaged over orbital phase these results suggest that 10 secondary star contributes 20+1354 of the Ix-band Dux., Averaged over orbital phase these results suggest that the secondary star contributes $29\pm13\%$ of the K-band flux.
+ ‘This is considerably dillerent from the 6545% estimated by Dhillon (2000)., This is considerably different from the $65\pm5\%$ estimated by Dhillon (2000).
+" The ΑΝοσο quick look light curves of RW Tri at the time of our ΕΠΗ observations show that RAW ‘Tri was at a magnitude of V~13. but during the observations of Dhillon (2000) RA ""Eri appeared to be in a low state with a magnitude of ~13.5."," The AAVSO quick look light curves of RW Tri at the time of our UKIRT observations show that RW Tri was at a magnitude of $\sim13$, but during the observations of Dhillon (2000) RW Tri appeared to be in a low state with a magnitude of $\sim13.8$."
+ Thus. RW Tri was approximately a [actor of 2 brighter during our observations than when observed by Dhillon (2000). accounting for the cilferent estimates of the secondary star contribution.," Thus, RW Tri was approximately a factor of 2 brighter during our observations than when observed by Dhillon (2000), accounting for the different estimates of the secondary star contribution."
+ Assuming that the [lux of the secondary star is constant. the acerction disc and stream in RAW Tri increased. by a factor of ~+ in brightness. between the observations of Dhillon (2000) ancl ours.," Assuming that the flux of the secondary star is constant, the accretion disc and stream in RW Tri increased by a factor of $\sim 4$ in brightness, between the observations of Dhillon (2000) and ours."
+ We noted. previously that the secondary. star leatures in RA ‘Tri are significantly broader than in the template spectra., We noted previously that the secondary star features in RW Tri are significantly broader than in the template spectra.
+ This is likely to be due to broadening caused by the rotation of the (phase-Iocked) secondary star as it orbits the white cart., This is likely to be due to broadening caused by the rotation of the (phase-locked) secondary star as it orbits the white dwarf.
+" The rotational velocity of the secondary star can be estimated by artificially broadening the template star spectrum which is assumed to have low Vi,sin7. and fitting it to the RAV Tri spectra."," The rotational velocity of the secondary star can be estimated by artificially broadening the template star spectrum which is assumed to have low $V_{rot}\sin i$, and fitting it to the RW Tri spectra."
+ The continuum was removed from the RAW Tri and template star spectra using a low order polynomial. after masking out strong absorption features.," The continuum was removed from the RW Tri and template star spectra using a low order polynomial, after masking out strong absorption features."
+" Each template was then artificially broadened in the velocity range Vi,sin;=I0. 200km/s in steps of LOkm/s. assuming partial limb darkening (linear limb darkening cocllicicnt = 0.5. North 2000)."," Each template was then artificially broadened in the velocity range $V_{rot}\sin i
+=10-200$ km/s in steps of 10km/s, assuming partial limb darkening (linear limb darkening coefficient = 0.5, North 2000)."
+ These broadened template spectra were then compared with the orbital velocity corrected. RAV ‘Tri spectra on each night., These broadened template spectra were then compared with the orbital velocity corrected RW Tri spectra on each night.
+ Resiclual spectra were produced by subtracting a constant times the shifted broadened template from cach RW Eri spectrum: the constant was adjusted to minimise the scatter on cach residual spectrum., Residual spectra were produced by subtracting a constant times the shifted broadened template from each RW Tri spectrum; the constant was adjusted to minimise the scatter on each residual spectrum.
+ A boxcar average smoothing was applied to the residual spectrum to eliminate any large scale structure., A boxcar average smoothing was applied to the residual spectrum to eliminate any large scale structure.
+ Phe reduced chi-squared (AZ) was calculated between cach residual. and smootheel spectrum in the wavelength. regions containing the Na I absorption feature tto 22108A3) and the Ca Ll absorption feature tto 22680A))., The reduced chi-squared $\chi_{\nu}^2$ ) was calculated between each residual and smoothed spectrum in the wavelength regions containing the Na I absorption feature to ) and the Ca I absorption feature to ).
+ Phe average results can be seen in Figure where we plot night 1 and night 2 separately. and also combined.," The average results can be seen in Figure \ref{f7a} where we plot night 1 and night 2 separately, and also combined."
+" The best [fit (minimum V2) for night 1 and 2 is obtained with τοιsin; 90kmj/s and V,sin 140km/'s respectively (Figure. 8))."," The best fit (minimum $\chi_{\nu}^{2}$ ) for night 1 and 2 is obtained with $V_{rot}\sin i
+\sim 90$ km/s and $V_{rot}\sin i \sim 140$ km/s respectively (Figure \ref{f7a}) )."
+" Based on the X7 distribution in ligure δις we estimate a mean 1,0sin?=120+ 20km/s. All emplate stars gave the same order of Chi-squared. values. confirming again that the data are not template sensitive in ais wavelength band."," Based on the $\chi_{\nu}^2$ distribution in Figure \ref{f7a}, we estimate a mean $V_{rot}\sin i=120\pm20$ km/s. All template stars gave the same order of Chi-squared values, confirming again that the data are not template sensitive in this wavelength band."
+ The minimum Z2 values had a range ver all the templates of Ἐνsin7 from 58) l00km/s for night 1. and 130 150km/s for night 2.," The minimum $\chi_{\nu}^2$ values had a range over all the templates of $V_{rot}\sin i$ from $80-100$ km/s for night 1, and $130-150$ km/s for night 2."
+" We confirmed tha 1e intrinsic 15,sin; of each template was Consistent with Kms. by cross correlating the templates against each other."," We confirmed that the intrinsic $V_{rot}\sin i$ of each template was consistent with 0km/s, by cross correlating the templates against each other."
+ The analysis of the orbital racial velocity was no ‘hanged significantly when we broadened the template lines w 120kmj/s. This is because rotational broadening allects 1e profile of the absorption lines. but not the position of," The analysis of the orbital radial velocity was not changed significantly when we broadened the template lines by 120km/s. This is because rotational broadening affects the profile of the absorption lines, but not the position of"
+] have utilized existing data [rom deep[Hubble andSpilzer surveys to provide a measure of the coanoving luminositv density at uliraviolet. and. optical wavelengths at z~6.,I have utilized existing data from deep and surveys to provide a measure of the co-moving luminosity density at ultraviolet and optical wavelengths at $z\sim6$.
+ In parüceular. I provide a measure of the contribution to the optical Iuminositw. density from [aint galaxies which are below (heSpZzer detection limit.," In particular, I provide a measure of the contribution to the optical luminosity density from faint galaxies which are below the detection limit."
+ Even alter accounting [or [unl ealaxies. (he optical Iuminositw. density is a [factor of 2—3 below the ultraviolet. Iuminosity densitv.," Even after accounting for faint galaxies, the optical luminosity density is a factor of $2-3$ below the ultraviolet luminosity density."
+" I fit the resultant. luminosity density estimates wilh a stellar population svnthesis nodel to determine the maximal age and stellar mass density al z~ο,", I fit the resultant luminosity density estimates with a stellar population synthesis model to determine the maximal age and stellar mass density at $z\sim6$.
+ Assuming a Salpeter IMF. HE find that the stellar mass density at z~6 is 1.6x 107 and the maximal age of a single stellar population is S100 Myr.," Assuming a Salpeter IMF, I find that the stellar mass density at $z\sim6$ is $\times$ $^{7}$ and the maximal age of a single stellar population is $\lesssim$ 100 Myr."
+ By comparing the iunmnber of ionizing photons per barvon produced over the age of the starburst with the nunber of ionizing photons per barvon required to keep the IGM ionized. I find that reionization must have been a brief inhomogeneous process lasting C100 Myr. and. must rave been completed as late as z«7 if the stellar [AIF had a Salpeter slope.," By comparing the number of ionizing photons per baryon produced over the age of the starburst with the number of ionizing photons per baryon required to keep the IGM ionized, I find that reionization must have been a brief inhomogeneous process lasting $\lesssim$ 100 Myr and must have been completed as late as $z<7$ if the stellar IMF had a Salpeter slope."
+ Motivated by WMAP results which suggest an early epoch of reionization. I investigate the form of the stellar IME if reionizalion was a single continuous process al redshilts higher than 6.," Motivated by WMAP results which suggest an early epoch of reionization, I investigate the form of the stellar IMF if reionization was a single continuous process at redshifts higher than 6."
+" If the past history of star-lormation in 2~6 galaxies was responsible for reionization. I find that the the slope of the stellar IAIF has to be non-Salpeter with a slope of a=—1.65 i 24,4:4,,""99 and a=—1.5 il z,5;;,,—111 for an IMF extending up to 200 M.."," If the past history of star-formation in $z\sim6$ galaxies was responsible for reionization, I find that the the slope of the stellar IMF has to be non-Salpeter with a slope of $\alpha=-1.65$ if 9 and $\alpha=-1.5$ if 11 for an IMF extending up to 200 $_{\sun}$."
+ However. the exact slope is sensitive to the ratio of the clumping factor to escape fraction. the metallicity ol the stars and the value of the visible huminositv clensitv at z~6.," However, the exact slope is sensitive to the ratio of the clumping factor to escape fraction, the metallicity of the stars and the value of the visible luminosity density at $z\sim6$."
+" On the other hand. the IGM could have been ionized between 6<z«15 bv a population of massive stars between 10—100 MÀN, which evolve into stellar remnants like black holes by zo6."," On the other hand, the IGM could have been ionized between $6<z<15$ by a population of massive stars between $10-100$ $_{\sun}$ which evolve into stellar remnants like black holes by $z\sim6$."
+ Such a population would not be constrained by (the rest-lrame visible and ultraviolet luminosity densitv αἱ 2~6., Such a population would not be constrained by the rest-frame visible and ultraviolet luminosity density at $z\sim6$.
+ The resultant. comoving mass densitv of remnants ad z~6 would be 2x10* acid would therefore be comparable to the stellar mass density (hat we see in z~6 galaxies., The resultant comoving mass density of remnants at $z\sim6$ would be $2\times10^{7}$ and would therefore be comparable to the stellar mass density that we see in $z\sim6$ galaxies.
+ Finally. I consider the evolving stellar initial mass function Chat has been suggested bv Dave(2008) and vanDokkum(2008) to account for the discrepancy in specilic formation rate estimates and colors of early-(ype galaxies at lower redshifts respectively.," Finally, I consider the evolving stellar initial mass function that has been suggested by \citet{Dave}
+ and \citet{PvD} to account for the discrepancy in specific star-formation rate estimates and colors of early-type galaxies at lower redshifts respectively."
+ This IMF if evolved out to z4. would result in a bottom-light IME with a non-Salpeter slope al Af< 12MM...," This IMF if evolved out to $z\sim4$, would result in a bottom-light IMF with a non-Salpeter slope at $M<12$ $_{\sun}$."
+ Such an IME could produce sufficient ionizing photons to account lor late reionization of the IGM at zS7 only if the visible light luminosity density was at the high end of our assumed range., Such an IMF could produce sufficient ionizing photons to account for late reionization of the IGM at $z\lesssim7$ only if the visible light luminosity density was at the high end of our assumed range.
+0.1iurespectively. Redshifts for the clusters will be important to learn about ICAL preheating. but it is interesting that a iiminimn core radius is miposed bv preheating which is rot very sensitive to redshift.,"0.1in Redshifts for the clusters will be important to learn about ICM preheating, but it is interesting that a minimum core radius is imposed by preheating which is not very sensitive to redshift."
+" The simple existeuce of au approxinate nininmuni core radius would be a ""1]okiug eun for preheatius. aud the value of this core size is an indicator of the approximate value of the cutropy floor."," The simple existence of an approximate minimum core radius would be a “smoking gun” for preheating, and the value of this core size is an indicator of the approximate value of the entropy floor."
+" For example. using spiro,=100 keV cu? will result iu a iuininnuu core radius of roughly 307,"," For example, using $s_{floor}=100$ keV $^{2}$ will result in a minimum core radius of roughly $''$."
+ Redshift evolution of the amount of preheating should be observable in the SZ properties of the cluster catalog., Redshift evolution of the amount of preheating should be observable in the SZ properties of the cluster catalog.
+ Of particular importance is that preheating should not sjeuificautlv affect the wmmber of clusters discovered in au SZ survey. ouly the appearance of the discovered clusters.," Of particular importance is that preheating should not significantly affect the number of clusters discovered in an SZ survey, only the appearance of the discovered clusters."
+" Towever. as iudicated by the dotted line in Figure 2.. the change in expected vields is comparable to the change in vields expected from a shift iu ©,, bv1054..."," However, as indicated by the dotted line in Figure \ref{fig:mlims}, the change in expected yields is comparable to the change in yields expected from a shift in $\Omega_m$ by."
+ Clearly. uon-eravitational heatius will be an muportaut consideration when trving to make precise estimates of cosmological paralucters frou SZ surveys (c.e.. Taian. Moly. aud Holder 2001).," Clearly, non-gravitational heating will be an important consideration when trying to make precise estimates of cosmological parameters from SZ surveys (e.g., Haiman, Mohr, and Holder 2001)."
+ The structure of ligh-+vedshift clusters can have iuportant implications for studies of CAIB anisotropies., The structure of high-redshift clusters can have important implications for studies of CMB anisotropies.
+ At small anguhuw scales. the angulu power spectrum can become dominated by the thermal SZ effect.," At small angular scales, the angular power spectrum can become dominated by the thermal SZ effect."
+ The magnitude of the thermal SZ power spectrum. ds most affected by preheating at exactly the angular scales where it becomes dominant. as shown in Figure 3.. as also fouud in other studies (olderaudCarlstrom1999:IKoimnatsu 2000).," The magnitude of the thermal SZ power spectrum is most affected by preheating at exactly the angular scales where it becomes dominant, as shown in Figure \ref{fig:sz_cl}, as also found in other studies \citep{holder99a,komatsu99a,springel00}."
+ 0.1iu Subtracting out the bright clusters can siguificautlv reduce the power at low inultipoles but does uot sienificautly affect the angular power spectrum at high f. where the thermal SZ effect is expected to exceed. the primary audsotropies.," 0.1in Subtracting out the bright clusters can significantly reduce the power at low multipoles, but does not significantly affect the angular power spectrum at high $\ell$, where the thermal SZ effect is expected to exceed the primary anisotropies."
+ The power at high 6 is mainly duc to low-mass clusters at higher :. where nou-gravitational heating can be very significant.," The power at high $\ell$ is mainly due to low-mass clusters at higher $z$, where non-gravitational heating can be very significant."
+ The lieher deusities at lugher redshift lead to a lower sclf-similar prediction for the ceutral eutropy. making clusters at ligher redshift especially seusitive. to excess entropy.," The higher densities at higher redshift lead to a lower self-similar prediction for the central entropy, making clusters at higher redshift especially sensitive to excess entropy."
+ Deep SZ surveys will allow studies of the clusters which are contributing the majority of the signal at small angular scales., Deep SZ surveys will allow studies of the clusters which are contributing the majority of the signal at small angular scales.
+ Shallow surveys will detect the clusters which provide the bulk of the power at larger angular scales. where the primary anisotropies are iore than an order of maguitude larecr than the thermal SZ effect. but will have very little information on the siiall angular scales at (z2000.," Shallow surveys will detect the clusters which provide the bulk of the power at larger angular scales, where the primary anisotropies are more than an order of magnitude larger than the thermal SZ effect, but will have very little information on the small angular scales at $\ell \ga 2000$."
+ We have shown that the SZ effect is well-suited for studies of non-eravitational heating., We have shown that the SZ effect is well-suited for studies of non-gravitational heating.
+ A well-planned survey should not have a selection fiction which strouglv depends on the amount of non-eravitational heating. while SZ images of the resulting catalog will be very scusitive to this enerev injection.," A well-planned survey should not have a selection function which strongly depends on the amount of non-gravitational heating, while SZ images of the resulting catalog will be very sensitive to this energy injection."
+ As a redshift-indepeudent probe. the SZ is particularly well-suited for high-redshift work. where nuon-eravitational heating can be very importaut.," As a redshift-independent probe, the SZ is particularly well-suited for high-redshift work, where non-gravitational heating can be very important."
+ We have shown that the expected vield and redshift distribution of clusters for an SZ survey is only weakly sensitive to preleating. but that this seusifivitv is sienificant if one wishes to do “precision cosmologv with SZ survers.," We have shown that the expected yield and redshift distribution of clusters for an SZ survey is only weakly sensitive to preheating, but that this sensitivity is significant if one wishes to do “precision cosmology” with SZ surveys."
+ This is true for either deep or shallow survevs., This is true for either deep or shallow surveys.
+ While high-ass clusters are less scusitive to preheating. he survey vield is more sensitive to the mass hit at Neher masses.," While high-mass clusters are less sensitive to preheating, the survey yield is more sensitive to the mass limit at higher masses."
+ Thus we have a situation where it will o difüeult to use high-uass clusters to learn about xeheatime. but we must understand it to do cosimolosy.," Thus we have a situation where it will be difficult to use high-mass clusters to learn about preheating, but we must understand it to do cosmology."
+ For low-mass clusters. the mass lait is still importaut or the survey vield. but direct maging of the resulting catalog of clusters should vield a wealth of information on xelheatiug.," For low-mass clusters, the mass limit is still important for the survey yield, but direct imaging of the resulting catalog of clusters should yield a wealth of information on preheating."
+ The peak of the SZ angular power spectrum is very sensitive to the details of non-eravitational heating aud direct jmaees of high redshift. low mass clusters may be the best wav of muderstauding this signal.," The peak of the SZ angular power spectrum is very sensitive to the details of non-gravitational heating and direct images of high redshift, low mass clusters may be the best way of understanding this signal."
+ Approximately half of the power at (22000. where the thermal SZ signal is stronecr than the primary anisotropies. should be coming from clusters that could be detected in deep SZ surveys.," Approximately half of the power at $\ell \ga 2000$, where the thermal SZ signal is stronger than the primary anisotropies, should be coming from clusters that could be detected in deep SZ surveys."
+ Upcoming CAIB experiments such as MAP and PLANCK. with their laree beam sizes. are not well-suited for finding lieh redshift and low mass clusters.," Upcoming CMB experiments such as MAP and PLANCK, with their large beam sizes, are not well-suited for finding high redshift and low mass clusters."
+ Ground-based bolometer arrays and interferometers with sub-arcuunute resolution will be well equipped for finding and studving these clusters., Ground-based bolometer arrays and interferometers with sub-arcminute resolution will be well equipped for finding and studying these clusters.
+ Deep SZ surveys will happen within the next few vears and the imtormation that they vield will be verv valuable for coustraiming non-eravitational heatius through ealaxy formation or relonization., Deep SZ surveys will happen within the next few years and the information that they yield will be very valuable for constraining non-gravitational heating through galaxy formation or reionization.
+ This intormation should be remarkably mubiased. auc Huprovenients in SZ imaeine over the uext several wears should coutinue to improve our understanding of these processes.," This information should be remarkably unbiased, and improvements in SZ imaging over the next several years should continue to improve our understanding of these processes."
+ ," \nocite {seljak96}
+ "
+include ~ 60.000 redshifts.,"include $\sim$ 60,000 redshifts."
+ To date the survev includes ~42.000 recshilts.," To date the survey includes $\sim$ 42,000 redshifts."
+ We began observations in the HectoMAP strip in 2009 and they will continue at least through the spring of 2012., We began observations in the HectoMAP strip in 2009 and they will continue at least through the spring of 2012.
+ We acquired spectra Lor (he objects with the HHectospec (Fabricant el al., We acquired spectra for the objects with the Hectospec (Fabricant et al.
+ 1993. 2005). a 300-fiber robotic instrument mounted on the MMT.," 1998, 2005), a 300-fiber robotic instrument mounted on the MMT."
+ The IHectospec observation planning software (Roll et al., The Hectospec observation planning software (Roll et al.
+ 1993) enables efficient acquisition of a pre-selected sample of galaxies., 1998) enables efficient acquisition of a pre-selected sample of galaxies.
+ The software enables assignment of priorities as a functüon of galaxy properties., The software enables assignment of priorities as a function of galaxy properties.
+ The spectra cover the wavelength range 3500 10.000 wwith a resolution of ~6À..," The spectra cover the wavelength range 3500 — 10,000 with a resolution of $\sim$ 6."
+ Exposure times ranged from 0.75 1.5 hours., Exposure times ranged from 0.75 — 1.5 hours.
+" For galaxies with SDSS r4,«20.5 and rjj,«21.5. we can observe in gray (and even some bright) lime because the galaxies have relatively high surface brightness. the IHectospec fibers are small. and the strip we chose is nearly always far from the moon."," For galaxies with SDSS $_{petro} < 20.5$ and $_{fiber} < 21.5$, we can observe in gray (and even some bright) time because the galaxies have relatively high surface brightness, the Hectospec fibers are small, and the strip we chose is nearly always far from the moon."
+ We reduced the data with the standard Hectospec pipeline (Mink οἱ al., We reduced the data with the standard Hectospec pipeline (Mink et al.
+ 2007) and derived redshifis with RVSAO (Ixurtz Mink 1998) with templates constructed for this purpose (Fabricant et al., 2007) and derived redshifts with RVSAO (Kurtz Mink 1998) with templates constructed for this purpose (Fabricant et al.
+ 2005)., 2005).
+ Repeat observations vield robust estimates of (he median error in ος where z is the redshift and © is the speed of light., Repeat observations yield robust estimates of the median error in $cz$ where $z$ is the redshift and $c$ is the speed of light.
+ For emission line objects. the median error (normalized by (1+ 2)) is 27 km |: the median [or absorption line objects (again normalized by (1+ 2)) is 27 km | (Geller et al.," For emission line objects, the median error (normalized by $(1 + z)$ ) is 27 km $^{-1}$ ; the median for absorption line objects (again normalized by $(1 + z$ )) is 37 km $^{-1}$ (Geller et al."
+ 2011: Fabricant et al., 2011; Fabricant et al.
+ 2005)., 2005).
+ The median redshift of HlectoMAP. is z=0.34. ~3.8 times the depth of the Sloan redshifl survey.," The median redshift of HectoMAP is $ z = 0.34$, $\sim 3.8$ times the depth of the Sloan redshift survey."
+ If the HectoMAP region cuts through great walls al 0.2<z0.6. they could easily exceed the extent of the Sloan Great wall and still be contained within the survey boundaries.," If the HectoMAP region cuts through great walls at $0.2 < z < 0.6$, they could easily exceed the extent of the Sloan Great wall and still be contained within the survey boundaries."
+ Figure 4 is a cone diagram showing the curent status of HectoMAD., Figure \ref{fig:hectomap} is a cone diagram showing the current status of HectoMAP.
+ It is alwavs fascinating (ο see the wav each individual Hectospec field contributes to the definition of large-scale structure in the survey region., It is always fascinating to see the way each individual Hectospec field contributes to the definition of large-scale structure in the survey region.
+ Michael Kitz and Scott IXenvon made the movie that shows the data in Figure 4 in the order of acquisition., Michael Kurtz and Scott Kenyon made the movie that shows the data in Figure \ref{fig:hectomap} in the order of acquisition.
+ In the movie. each vellow dot represents a galaxy.," In the movie, each yellow dot represents a galaxy."
+ The gray arcs are spaced by 0.1 in redshift beginning al 2—0.1., The gray arcs are spaced by 0.1 in redshift beginning at $z = 0.1$.
+ Figure 4 shows the final frame of the movie along with an inset of the first slice of the CLA survey [rom Figure | on the same scale., Figure \ref{fig:hectomap} shows the final frame of the movie along with an inset of the first slice of the CfA survey from Figure \ref{fig:firstslice} on the same scale.
+ There are 42.147 redshifts in the current HectoMAP. sample (7343 are from the SDSS: the rest are new Ilectospec data).," There are 42,147 redshifts in the current HectoMAP sample (7343 are from the SDSS; the rest are new Hectospec data)."
+ The movie gives an idea of the excitement we feel as we watch the survey develop., The movie gives an idea of the excitement we feel as we watch the survey develop.
+ In contrast with the first slice of the CLA sturvey where the striking pattern was a surprise. we now expect to see well-delineated. large voids surrounded by the ealaxies that deline the cosmic web.," In contrast with the first slice of the CfA survey where the striking pattern was a surprise, we now expect to see well-delineated, large voids surrounded by the galaxies that define the cosmic web."
+ Like other surveys to its depth. IectoNLAP. will provide constraints on (he characteristicsof the süiructure as a function of cosmological epoch.," Like other surveys to its depth, HectoMAP will provide constraints on the characteristicsof the structure as a function of cosmological epoch."
+simple energy argument similar to 2? to estimate how the subcluster masses and half-mass radii evolve after the expulsion of the gas.,simple energy argument similar to \citet{hills80} to estimate how the subcluster masses and half-mass radii evolve after the expulsion of the gas.
+ It is insightful to consider the system at two key moments., It is insightful to consider the system at two key moments.
+ The evolution of the subclusters between these two moments can be quantified by evaluating the conservation of energy., The evolution of the subclusters between these two moments can be quantified by evaluating the conservation of energy.
+" For any subcluster. we can express the kinetic energy as where n, is the virial radius. and (v7)? denotes the mean square velocity. which as a result can be written as The total energy at the moment of instantaneous gas removal thus becomes where the relevant quantities have been marked with subscrip ""E to indicate the moment of gas expulsion."," For any subcluster, we can express the kinetic energy as where $r_{\rm v}$ is the virial radius, and $\langle v^2\rangle$ denotes the mean square velocity, which as a result can be written as The total energy at the moment of instantaneous gas removal thus becomes where the relevant quantities have been marked with subscript `1' to indicate the moment of gas expulsion."
+ Given a deviation from virial equilibrium. a subcluster wil respond by changing its radius and/or mass.," Given a deviation from virial equilibrium, a subcluster will respond by changing its radius and/or mass."
+ As can be verified from Fig. 5..," As can be verified from Fig. \ref{fig:mbound},"
+ most subclusters contain a certain number of unbound sink particles. which were either previously retained by the gas potential. or are randomly passing the subcluster close enough to be included by the cluster identitication algorithm.," most subclusters contain a certain number of unbound sink particles, which were either previously retained by the gas potential, or are randomly passing the subcluster close enough to be included by the cluster identification algorithm."
+ These unbounc sink particles will escape the subcluster upon gas expulsion and take away some of the kinetic energy., These unbound sink particles will escape the subcluster upon gas expulsion and take away some of the kinetic energy.
+ We now consider a second moment in time. at which the gas-rid subcluster has reached viria equilibrium (Qi»= 0.5) and the unbound sink particles have successfully escaped.," We now consider a second moment in time, at which the gas-rid subcluster has reached virial equilibrium $Q_{\rm vir,2}=0.5$ ) and the unbound sink particles have successfully escaped."
+" At this time. energy conservation dictates where E» is the total energy of the (virialised) subcluster. Ex. is the total energy of the escaped stars. and the relevant quantities have been marked with subscript 2"" to indicate the moment of virialisation."," At this time, energy conservation dictates where $E_2$ is the total energy of the (virialised) subcluster, $E_{\rm esc}$ is the total energy of the escaped stars, and the relevant quantities have been marked with subscript `2' to indicate the moment of virialisation."
+ The parameter «7 denotes the surplus energy per unit mass of the escaped stars after they clear the potential of the subcluster. in units of its mean square velocity.," The parameter $\beta$ denotes the surplus energy per unit mass of the escaped stars after they clear the potential of the subcluster, in units of its mean square velocity."
+ The values of «9 can be estimated from the simulation by computing CT;+Vif(GMymi/2n4) for each of the unbound sinkcles?.," The values of $\beta$ can be estimated from the simulation by computing $(T_i+V_i)/(GM_{\rm cl,2}m_i/2r_{\rm v,1})$ for each of the unbound sink."
+. For this. we use the relation between the virial and half- radius corresponding to a? potential. which is given by n=làn (e.g.2).," For this, we use the relation between the virial and half-mass radius corresponding to a \citet{plummer11} potential, which is given by $r_{\rm v}=1.3r_{\rm h}$ \citep[e.g.][]{heggiehut}."
+ The escaping sink particles are the tail of an approximatelyMaxwellian velocity distribution of the sink particles in the subcluster. and consequently the distribution of .7 declines exponentially as f67)xexpt2/09). with τν around unity.," The escaping sink particles are the tail of an approximatelyMaxwellian velocity distribution of the sink particles in the subcluster, and consequently the distribution of $\beta$ declines exponentially as $f(\beta)\propto {\rm exp}(-\beta/\beta_0)$, with $\beta_0$ around unity."
+ The mean of such a distribution equals o per definition. which illustrates that unbound stars typically retain velocities similar to the mean square velocity in the subcluster after they escape.," The mean of such a distribution equals $\beta_0$ per definition, which illustrates that unbound stars typically retain velocities similar to the mean square velocity in the subcluster after they escape."
+ Combining Eqs., Combining Eqs.
+ 8. and 10. one obtains an expression for the evolution of the gas-rid subcluster as it approaches virial equilibrium. which relates the half-mass radii. masses. initial virial ratio and «7.," \ref{eq:v2} and \ref{eq:energy2}, one obtains an expression for the evolution of the gas-rid subcluster as it approaches virial equilibrium, which relates the half-mass radii, masses, initial virial ratio and $\beta$."
+ It is given by where n» is the only unknown and all other variables given by the simulation.," It is given by where $r_{\rm h,2}$ is the only unknown and all other variables given by the simulation."
+" For 7=0. i.e. all escaping stars are only marginally unbound, the expression returns the basic result that when unbound stars escape from a virialised system (Οι=0.5). it contracts to reattain virialequilibriuny’."," For $\beta=0$ , i.e. all escaping stars are only marginally unbound, the expression returns the basic result that when unbound stars escape from a virialised system $Q_{\rm vir,1}=0.5$ ), it contracts to reattain virial."
+. Inserting typical values of :j=I (see above). Qi=0.59 and Mas/Maj0.95 (see Sect. 3.29) ," Inserting typical values of $\beta=1$ (see above), $Q_{\rm vir,1}=0.59$ and $M_{\rm cl,2}/M_{\rm cl,1}=0.95$ (see Sect. \ref{sec:virial}) )"
+in Eq., in Eq.
+ LI yields nyoft=1-04. which justifies the earlier assumption that the radius does not change much between instantaneous σας removal and virialisation (see footnote 8)).," \ref{eq:rh} yields $r_{\rm h,2}/r_{\rm h,1}=1.04$, which justifies the earlier assumption that the radius does not change much between instantaneous gas removal and virialisation (see footnote \ref{ft:assum}) )."
+ This minor expansion is driven by the slightly supervirial state of the subelusters. and inhibited by the escape of unbound stars. which have velocities larger than the escape velocity.," This minor expansion is driven by the slightly supervirial state of the subclusters, and inhibited by the escape of unbound stars, which have velocities larger than the escape velocity."
+ As discussed in Sect. 3.2..," As discussed in Sect. \ref{sec:virial},"
+ the virial ratios of the subclusters give an indication of their survival fraction after gas removal., the virial ratios of the subclusters give an indication of their survival fraction after gas removal.
+ Out of all 140 subclusters identified in the simulation. only 10 have virial ratios Qi;2| and are therefore unbound.," Out of all 140 subclusters identified in the simulation, only 10 have virial ratios $Q_{\rm vir}>1$ and are therefore unbound."
+" In the last snapshot of the simulation. only one of the 21 subclusters will disperse after the removal of thegas""."," In the last snapshot of the simulation, only one of the 21 subclusters will disperse after the removal of the."
+.. As a result. typically of all the identified subclusters survive gas expulsion.," As a result, typically of all the identified subclusters survive gas expulsion."
+ The fate of these survivors depends on whether they expand. and how their environment affects them.," The fate of these survivors depends on whether they expand, and how their environment affects them."
+ Expanded subclusters with lower densities are more susceptible to disruption by tidal shocks., Expanded subclusters with lower densities are more susceptible to disruption by tidal shocks.
+ The evolution of the half-mass radii after gas expulsion can be considered in more detail by evaluating Eq., The evolution of the half-mass radii after gas expulsion can be considered in more detail by evaluating Eq.
+ 11. for each of the subclusters in the simulation., \ref{eq:rh} for each of the subclusters in the simulation.
+ This enables a comparison of the distribution of half-mass radii of the current subelusters (at the moment of instantaneous gas removal) with the distribution of their half-mass radii when they have reached virial equilibrium. which is shown in Fig. 9..," This enables a comparison of the distribution of half-mass radii of the current subclusters (at the moment of instantaneous gas removal) with the distribution of their half-mass radii when they have reached virial equilibrium, which is shown in Fig. \ref{fig:rh}."
+ The distribution of half-mass radii changes remarkably little after gas removal. as the means of the lognormal functions that are fitted to both distributions differ by 0.035 dex.," The distribution of half-mass radii changes remarkably little after gas removal, as the means of the lognormal functions that are fitted to both distributions differ by 0.035 dex."
+ This implies mofma=1.08. very similar to the earlier. simple estimate of rn= 1.04.," This implies $r_{\rm h,2}/r_{\rm h,1}=1.08$, very similar to the earlier, simple estimate of $r_{\rm h,2}/r_{\rm h,1}=1.04$ ."
+ The subclusters in the last snapshot experienceroughly 1.5 times this expansion after gas removal. which is of the same order of magnitude as the expansion of the," The subclusters in the last snapshot experienceroughly 1.5 times this expansion after gas removal, which is of the same order of magnitude as the expansion of the"
+resolution. simply because radciating regions occupy a significant fraction of the volume.,"resolution, simply because radiating regions occupy a significant fraction of the volume."
+ For example. in our 128* simulations we get satisfactory results with ~4x10° ravs sampling ihe diffuse radiation.," For example, in our $128^3$ simulations we get satisfactory results with $\sim 4\times 10^6$ rays sampling the diffuse radiation."
+ Emissivity from radiative recombinations and bremsstralilung are computed using atomic models for hydrogenic and IHe-like ions., Emissivity from radiative recombinations and bremsstrahlung are computed using atomic models for hydrogenic and He-like ions.
+" For hvdrogenic ious the [requency-dependent recombination and [Iree-Iree emissivity is (IDunmer 1994) where llere a is the fine-structure constant. vy is the hydrogen Lyman limit. Z is the atomic number. Aj and Ap are the Boltzmann and Planck constants. and o,(Z.€) is the eross-section of photoionization lor level »."," For hydrogenic ions the frequency-dependent recombination and free-free emissivity is (Hummer 1994) where Here $\alpha$ is the fine-structure constant, $\nu_1$ is the hydrogen Lyman limit, $Z$ is the atomic number, $\kb$ and $\hp$ are the Boltzmann and Planck constants, and $\sigma_n(Z,{\cal E})$ is the cross-section of photoionization for level $n$."
+ Similarly. for the Bree-Iree coellicient we have introduced the Iree-Iree Gaunt [actor gg(u. A). the Compton wavelength Αι. and u=fv/hpT. (Hummer 1994).," Similarly, for the free-free coefficient we have introduced the free-free Gaunt factor $g_{\rm ff}(u,\lambda)$ , the Compton wavelength $\lambda_{\rm c}$ , and $u\equiv h\nu/\kb\te$ (Hummer 1994)."
+" The quantity 7j, is either my for hydrogen or 7j. for coubly ionizecl helium. while i. is the fraction of free electrons relative to hydrogen."," The quantity $n_{\rm ion}$ is either $n_{{\rm H}^+}$ for hydrogen or $n_{{\rm He}^{++}}$ for doubly ionized helium, while $x_{\rm e}$ is the fraction of free electrons relative to hydrogen."
+" We take the numerical values for the hycogenic photoionization cross-sections o,(Z.€) from Storev Thunimer (1991) ancl the free-lree Gaunt [actor οί.A) from Ibununer (1988). and integrate eq. (4))"," We take the numerical values for the hydrogenic photoionization cross-sections $\sigma_n(Z,{\cal E})$ from Storey Hummer (1991) and the free-free Gaunt factor $g_{\rm ff}(u,\lambda)$ from Hummer (1988), and integrate eq. \ref{eq:emissivity}) )"
+ numerically (o get for our single Lrequency group., numerically to get for our single frequency group.
+ The expression for enissivities due to recombinations to the n/S state of neutral helium is similar and ean be found in IIuminer Storey (1993)., The expression for emissivities due to recombinations to the $nlS$ state of neutral helium is similar and can be found in Hummer Storey (1998).
+ In our caleulations we neglectthe emissivity[rom di-electronie recombinations of helium., In our calculations we neglectthe emissivityfrom di-electronic recombinations of helium.
+Globular clusters (GCs) are very old stellar objects with typical masses between 107 Mc; and 10° Me teorresponding roughly to total lummosities between My2—5 to My= -10). having in general compact sizes with half-light radit of a few pe.,"Globular clusters (GCs) are very old stellar objects with typical masses between $10^4$ $_{\sun}$ and $10^6$ $_{\sun}$ (corresponding roughly to total luminosities between $M_V = -5$ to $M_V = -10$ ), having in general compact sizes with half-light radii of a few pc."
+ This morphology makes them easily observable also in external galaxies with modern telescopes (seeBrodie&Strader2006.andreferences therein).., This morphology makes them easily observable also in external galaxies with modern telescopes \citep[see][and references therein]{brodie06}.
+ The has a rich GC system containing 150 GCs (Harris1996)., The has a rich GC system containing 150 GCs \citep{harris}.
+. Most of them are compact with sizes of a few pe., Most of them are compact with sizes of a few pc.
+ Only 13 GCs (or 9%)) have an effective radius larger than 10 pe., Only 13 GCs (or ) have an effective radius larger than 10 pc.
+ Most of these extended clusters (ECs) are fainter than about My=-7. only2419.. having a half-light radius of about 20 pe. has a high luminosity of about My=-9.4 mag.," Most of these extended clusters (ECs) are fainter than about $M_{\rm V} = -7$, only, having a half-light radius of about 20 pc, has a high luminosity of about $M_{\rm V} = -9.4$ mag."
+ Further ECs in the vicinity of the Milky Way have been found in the and the (Mackey&GilmoreMarel 2005).," Further ECs in the vicinity of the Milky Way have been found in the and the \citep{mackey04,vandenbergh04,mvdm}."
+ Comparable objects have also been detected around other galaxies., Comparable objects have also been detected around other galaxies.
+ Huxoretal.(2005) found three ECs aroundM31.. which have very large radit above 30 pc.," \citet{huxor04} found three ECs around, which have very large radii above 30 pc."
+ These clusters were detected by chance as the automatic detection algorithms of the MegaCam Survey discarded such extended objects as likely background contaminations., These clusters were detected by chance as the automatic detection algorithms of the MegaCam Survey discarded such extended objects as likely background contaminations.
+ Follow-up observations by Mackeyetal.(2006).. using the ACS camera of the Hubble Space Telescope (HST). resolved the ECs into stars proving their nature as M31 clusters.," Follow-up observations by \citet{mackey06}, using the ACS camera of the Hubble Space Telescope (HST), resolved the ECs into stars proving their nature as M31 clusters."
+ They also detected a fourth EC around M31., They also detected a fourth EC around M31.
+ The M31 ECs have masses of the order of 10? Mo., The M31 ECs have masses of the order of $10^5$ $_{\sun}$.
+ Further observations increased the number of ECs in M31 to 13 (Huxoretal.2008)., Further observations increased the number of ECs in M31 to 13 \citep{huxor08}.
+. However. Huxoretal.(2011) showed that the previous estimates of the effective radit were considerably too large.," However, \cite{huxor11} showed that the previous estimates of the effective radii were considerably too large."
+ The new size estimate are well below 30 pe., The new size estimate are well below 30 pc.
+ Chandaretal.(2004) observed a part of the disks of the nearby galaxiesM81..M83..NGC6946.. MIOI.. and using HST and found ECs with effective radit larger than 10 pe in four of them.," \citet{chandar04} observed a part of the disks of the nearby galaxies, and using HST and found ECs with effective radii larger than 10 pc in four of them."
+ M51 showed a very high fraction of ECs in the observed area: 8 of 34 GCs (24%))., M51 showed a very high fraction of ECs in the observed area: 8 of 34 GCs ).
+ ECs are now detected in all types of galaxies from dwarfs to ellipticals (e.g.Larsenetal.2009;DaCosta 2009).," ECs are now detected in all types of galaxies from dwarfs to ellipticals \citep[e.g.][]{larbro00,harris02,lee05,peng,chisa,stonkute,georgiev,dacosta}."
+. Hilkeretal.(1999) and Drinkwateretal.(2000) discovered in the compact objects with luminosities above the brightest known GCs and which were not resolved by ground-based observations., \citet{hilker99} and \citet{drinkwater00} discovered in the compact objects with luminosities above the brightest known GCs and which were not resolved by ground-based observations.
+ These objects have masses between a few 10° Me and 10° Me and effective radii between rar=10 and 100 pe., These objects have masses between a few $10^{6}$ $_{\sun}$ and $10^{8}$ $_{\sun}$ and effective radii between $r_{\rm eff} = 10$ and 100 pc.
+ Drinkwateretal. interpreted these objects as a new type of galaxy and reflected this interpretation in the name “ultra-compact dwarf galaxy” (UCD)., \citet{drinkwater00} interpreted these objects as a new type of galaxy and reflected this interpretation in the name “ultra-compact dwarf galaxy” (UCD).
+ Bekkietal.(2001) suggested that UCDs are the remnants of dwarf galaxies which lost their dark matter halo and all stars except their nucleus., \citet{bekki01} suggested that UCDs are the remnants of dwarf galaxies which lost their dark matter halo and all stars except their nucleus.
+ Next to the interpretation as a galaxy. UCDs were also considered as high-mass versions of normal GCs (Mieskeetal.2002).. or as merged massive complexes of star clusters (Kroupa1998:Fellhauer& 2002a).," Next to the interpretation as a galaxy, UCDs were also considered as high-mass versions of normal GCs \citep{mieske02}, or as merged massive complexes of star clusters \citep{krou98,fellhauer02a}."
+. Aany UCDs have been found now., Many UCDs have been found now.
+ Next to the Fornax Cluster. they have been observed in the galaxy cluster (Mieskeetal. 2004).. around in the (Haseganetal.2005:Evstigneeva2007). the (Mieskeetal.2007).. the (Madridetal.2010).. and (Blakeslee&Barber 2008).," Next to the Fornax Cluster, they have been observed in the galaxy cluster \citep{mieske04}, , around in the \citep{hasegan,evstigneeva07}, the \citep{mieske07}, the \citep{madrid}, and \citep{blakeslee}."
+. While most known UCDs belong to galaxy clusters. they have also been observed in rather isolated environments. e.g. in the Sombrero galaxy 104)) by Hauetal.," While most known UCDs belong to galaxy clusters, they have also been observed in rather isolated environments, e.g. in the Sombrero galaxy ) by \citet{hau}."
+(2009).. Forbesetal.(2008) and Mieskeetal.(2008) analyzed larger samples of UCDs., \citet{forbes08} and \citet{mieske08} analyzed larger samples of UCDs.
+ They find that normal and extended star clusters and UCDs form à coherent data set where size and mass-to-light ratio increase continuously with their total mass and concluded that UCDs are more likely bright extended clusters than naked cores of stripped dwarf galaxies., They find that normal and extended star clusters and UCDs form a coherent data set where size and mass-to-light ratio increase continuously with their total mass and concluded that UCDs are more likely bright extended clusters than naked cores of stripped dwarf galaxies.
+ The marginally enhanced mass-to-lightratios of UCDs can be explained by slightly modified initial stellar mass functions (Mieske&Kroupa2008:Dabringhausenetal. 2009).," The marginally enhanced mass-to-lightratios of UCDs can be explained by slightly modified initial stellar mass functions \citep{miekrou08,dabringhausen09}. ."
+The importance of dust extinction in the Calaxyv has been recognized since carly in this century when star-counting surveys revealed absorption of optical light by dark clouds’ (Barnard 1919).,The importance of dust extinction in the Galaxy has been recognized since early in this century when star-counting surveys revealed absorption of optical light by `dark clouds' (Barnard 1919).
+ It is fortunate that extinction correlates relatively well with reddening in the Galaxy. because it is difficult chough to accurately measure either the distance to a typical astronomical source or its iutiusic lundnosityv let alone both.," It is fortunate that extinction correlates relatively well with reddening in the Galaxy, because it is difficult enough to accurately measure either the distance to a typical astronomical source or its intrinsic luminosity – let alone both."
+ But knowing the intrinsic color (2V) (using an uuobseured dine of sight) along with the observed DBV color and reddcning-extinction law ly=Ry{eB—V)(B.V4REQB| V). oue can correctly determine the distauce of an object from its distance modulis. or vice versa.," But knowing the intrinsic color $(B-V)_i$ (using an unobscured line of sight) along with the observed $B-V$ color and a reddening-extinction law $A_V =
+R_V[(B-V)-(B-V)_i] \equiv R_VE(B-V)$ , one can correctly determine the distance of an object from its distance modulus, or vice versa."
+" value of Ry varies markedly within the Milkv. Way ""izdeu 1990) aud amoug differeut galaxies aKeRhy<for : ruisFaleo et al."," The value of $R_V$ varies markedly within the Milky Way $ 3
+\la R_V \la 6$; Mathis 1990) and among different galaxies $1.5 \la
+R_V \la 7.2$ ; Falco et al."
+ 1999). but a value of Ry2 is useful y ια estimates of Galactic extinction.," 1999), but a value of $R_V \simeq 3.2$ is useful for many estimates of Galactic extinction."
+ Radiation from extragalactic objects is subject to extinction by dust both imside anc outside of the Galaxy., Radiation from extragalactic objects is subject to extinction by dust both inside and outside of the Galaxy.
+ However. while extragalactic dust has received attention. our knowledge of its amount aud propertics is rather limited. because methods useful for estimating dust density in galaxies have proved less effective when extended to intergalactic space.," However, while extragalactic dust has received attention, our knowledge of its amount and properties is rather limited, because methods useful for estimating dust density in galaxies have proved less effective when extended to intergalactic space."
+ For example. a umber of eroups have attempted to measure extinction bv dust in clusters using backeround-object (Bosscounting. aud several claius of intracluster dust Wagouer 1973: Bovle. Foug Shauks 1988: Romani Maoz 1992) and extragalactic dust cloud detections (Wsozlek et al.," For example, a number of groups have attempted to measure extinction by dust in clusters using background-object counting, and several claims of intracluster dust (Bogart Wagoner 1973; Boyle, Fong Shanks 1988; Romani Maoz 1992) and extragalactic dust cloud detections (Wsozlek et al."
+ 1988 aud references therein) have been made. but even now these claus remain controversial (see Maoz 1995).," 1988 and references therein) have been made, but even now these claims remain controversial (see Maoz 1995)."
+ The strongest proposed Huüts on a diffuse distribution of iutergalactie dust with a Galactic reddening law have come from studies of the redshift evolutiou of the mean quasar spectral index (c.g. Wright 1981: Wright \alsau Loss: Cheng. Gaskell Ioratkar 1991).," The strongest proposed limits on a diffuse distribution of intergalactic dust with a Galactic reddening law have come from studies of the redshift evolution of the mean quasar spectral index (e.g. Wright 1981; Wright Malkan 1988; Cheng, Gaskell Koratkar 1991)."
+ These studies lit uniforii dust of nn colmoving density to have Ay(.=1)<0.05 nag (from Wright Malkan 1988). aud are niost sensitive to dust at 2>>1. While our knowledge of it is poor. iutergalactie dust could have great cosmological importance. as it could affect results concerning the cosmic microwave (CMD) aud cosnüc mfared (CIB) backgrounds. galaxy aud quasar uubers at high 2. galaxy evolution. large-scale structure. ete.," These studies limit uniform dust of constant comoving density to have $A_V(z=1) \la 0.05$ mag (from Wright Malkan 1988), and are most sensitive to dust at $z > 1.$ While our knowledge of it is poor, intergalactic dust could have great cosmological importance, as it could affect results concerning the cosmic microwave (CMB) and cosmic infrared (CIB) backgrounds, galaxy and quasar numbers at high $z$, galaxy evolution, large-scale structure, etc."
+ This paper discusses intergalactic dust chiefly iu the context of its nmuportance iu imeasurenmients of the cosmolocical deceleration paranueter — a subject discussed nunerous times. first by BEieeusou (1919) aud most recently in Aguirre 1999 (A990).," This paper discusses intergalactic dust chiefly in the context of its importance in measurements of the cosmological deceleration parameter – a subject discussed numerous times, first by Eigenson (1949) and most recently in Aguirre 1999 (A99)."
+ Conditions iu the diffuse intergalactic medinm (ICAL) strouely distavor dust formation. so whatever interealactic dust exists is probably either the remmant of an carly Population III epoch. or is formed in galaxies aud removed by some wechanisin. (," Conditions in the diffuse intergalactic medium (IGM) strongly disfavor dust formation, so whatever intergalactic dust exists is probably either the remnant of an early Population III epoch, or is formed in galaxies and removed by some mechanism. ("
+The remaining possibility. that a substantial dust-forming population of extragalactic stars exists. is not considered here.),"The remaining possibility, that a substantial dust-forming population of extragalactic stars exists, is not considered here.)"
+ Previous investigations of intergalactic dust have almost iuvariablv assuiued that it has properties simular to that of Galactic dust: but this assunrptiou is not well font.justified., Previous investigations of intergalactic dust have almost invariably assumed that it has properties similar to that of Galactic dust; but this assumption is not well justified.
+ Even amoung galaxies. Ry varies by a factor of and interealactic dust may have creation. destruction. aud selection 1iechanisuisquite," Even among galaxies, $R_V$ varies by a factor of four, and intergalactic dust may have creation, destruction, and selection mechanismsquite"
+reasonable estimates in the presence of systematic effects that render the other approaches useless.,reasonable estimates in the presence of systematic effects that render the other approaches useless.
+ Since the approaches respoud cifferentv to systematic and random errors. a sensible strategy is to estimate rp with all of aud look for cousisteucy amoug tlie restIts.," Since the approaches respond differently to systematic and random errors, a sensible strategy is to estimate $r_0$ with all of and look for consistency among the results."
+ The sample analysis of the Lymau-break survey helps illustrate the papers main poius., The sample analysis of the Lyman-break survey helps illustrate the paper's main points.
+ An initial estimate of ry1141 lpcfrom ectation 1 disagreed badly with the estimate i)dM! Ape from the robust equation &.. suggestiug hat the initial analysis must have had large sysematic errors.," An initial estimate of $r_0\sim 11h^{-1}$ Mpcfrom equation \ref{eq:blain} disagreed badly with the estimate $r_0\sim 4h^{-1}$ Mpc from the robust equation \ref{eq:expk}, suggesting that the initial analysis must have had large systematic errors."
+ The largest systematic error caue roni inaccuracies in the assumed selection fnction., The largest systematic error came from inaccuracies in the assumed selection function.
+ heacing Ho with a beter model reduced tje estimated: values of ry to 7.2. 6.1. 5.7. and NUES Alpe from equations 1.. LL. 2.. aud 8.. respectvelv.," Replacing it with a better model reduced the estimated values of $r_0$ to $7.2$ , $6.4$ $5.7$ and $4.0h^{-1}$ Mpc from equations \ref{eq:blain}, \ref{eq:nexpgiventheta}, \ref{eq:nexpgiventhetaz}, and \ref{eq:expk}, respectively."
+ The differences were still not negligible coupa'ed to the raudoin tncertaiities 6))., The differences were still not negligible compared to the random uncertainties \ref{sec:uncertainties}) ).
+ The hieh value from equation 1. was ¢ue to artificial augular clustering of galaxies inuposed by the surveys spectroscopic selection criteria., The high value from equation \ref{eq:blain} was due to artificial angular clustering of galaxies imposed by the survey's spectroscopic selection criteria.
+ It alone amoue he esinnators cloes ιοί correct for this., It alone among the estimators does not correct for this.
+ The remaining systematic problems are not easy to trace., The remaining systematic problems are not easy to trace.
+ They could result from residual errors in the selection ftuction or from changes in the correlation fiunction sloye al large sedaratious., They could result from residual errors in the selection function or from changes in the correlation function slope at large separations.
+ In any case. since the effect of systematic errors is minimized when tlev are sinall compared o the signal. E imaximized 11e sledal by liniitiug the analysis to angular. pairs with smaller separations.," In any case, since the effect of systematic errors is minimized when they are small compared to the signal, I maximized the signal by limiting the analysis to angular pairs with smaller separations."
+ As equation 3. shows. tlje LULLaber of pairs with large angular sep:watious Is Ore sensiive to low level systematics tlali to he elistering strength &.," As equation \ref{eq:nexpxibar} shows, the number of pairs with large angular separations is more sensitive to low level systematics than to the clustering strength $\bar\xi$."
+" Restricting the analysis pairs with angular separation 0;;<00 300"". Lobalnec the estimates 7j225.14.|. Lon4. NUM j»€ from equations L. 5.. and δ.."," Restricting the analysis to pairs with angular separation $\theta_{ij}<300''$ , I obtained the estimates $r_0=5.1h^{-1}$, $4.9h^{-1}$, $4.4h^{-1}$ Mpc from equations \ref{eq:nexpgiventheta}, \ref{eq:nexpgiventhetaz}, and \ref{eq:expk}."
+ Since he raπο uucertainty is ~lht Mpe 6) ). these ünates agree well with each other aud with the value ryLOx0.6h.1 Mpc favored by the eular-clustering analysis of Adelberger e al. (," Since the random uncertainty is $\sim 1h^{-1}$ Mpc \ref{sec:uncertainties}) ), these estimates agree well with each other and with the value $r_0=4.0\pm 0.6h^{-1}$ Mpc favored by the angular-clustering analysis of Adelberger et al. ("
+2001).,2004).
+ This paper provides some support [or the common prejudice agalust estimates of ry derived rou sualb galaxy samples., This paper provides some support for the common prejudice against estimates of $r_0$ derived from small galaxy samples.
+ The iuille eft panel of Figure Loslows how large the raucom uncertainties are for a simulated sample of NV=20 galaxies with true «'orrelation leugth ry=3.5h1 Ipe in a 10x pencil-beam survey.," The middle left panel of Figure \ref{fig:gif_summary}
+ shows how large the random uncertainties are for a simulated sample of $N=20$ galaxies with true correlation length $r_0=3.5h^{-1}$ Mpc in a $10'\times 10'$ pencil-beam survey."
+ Figure 6 may 1lake the poitt more forcefully., Figure \ref{fig:westconfint} may make the point more forcefully.
+" I extracted MUNELOUS raliοι subsamples of 10 galaxies [rom the 170-object Ly""unau-break galaxy caalog iu he Westphal field (Steicel et al.", I extracted numerous random subsamples of 10 galaxies from the 170-object Lyman-break galaxy catalog in the Westphal field (Steidel et al.
+ 2003). caleulated ry for each subsanje with equation L using the rue LBG selection Luneion. and tabulaed the results.," 2003), calculated $r_0$ for each subsample with equation \ref{eq:blain} using the true LBG selection function, and tabulated the results."
+ The spreac ---1 estimated iy is enorLLOUS., The spread in estimated $r_0$ is enormous.
+ Iu realistic situations. uncertaiuty in the assumed selection fuuction is likely to be tle worst source of systematic error.," In realistic situations, uncertainty in the assumed selection function is likely to be the worst source of systematic error."
+ A skeptic mieht point out that this uncertainty will probably ly be large in the stall sample limit. where none of the approaches work well. aud that my stested alternaives are not much «X an improvenment when the uucertaintv in the selection furOl ls stnall (see. e.g..IC tle upper left pauel of Fieure [)).," A skeptic might point out that this uncertainty will probably only be large in the small sample limit, where none of the approaches work well, and that my suggested alternatives are not much of an improvement when the uncertainty in the selection function is small (see, e.g., the upper left panel of Figure \ref{fig:gif_summary}) )."
+ TUs Is true toa »oiut. but it would be foolish to reject je ~BOM ‘ecluction iuraudoim uucertaiuty that equation 5 provides relative to equ:dion L..," This is true to a point, but it would be foolish to reject the $\sim 30$ reduction inrandom uncertainty that equation \ref{eq:nexpgiventhetaz} provides relative to equation \ref{eq:blain}. ."
+ According toFigure 5.. a3QU decrease in the uicertainty in rg for the LBC sample requires a," According toFigure \ref{fig:varn_vs_n},,a $\sim 30$ decrease in the uncertainty in $r_0$ for the LBG sample requires a"
+where P;=P(A?) and w; is the weight of a data point.,where $P_i = P(\lambda_i ^2)$ and $w_i$ is the weight of a data point.
+" The RMSF is given by Here, Ao i-iis arbitrary and we chose the wavelength corresponding to the weighted average 4?."," The RMSF is given by Here, $\lambda_0$ is arbitrary and we chose the wavelength corresponding to the weighted average $\lambda^2$."
+" When planning a rotation-measure experiment, three main parameters are involved, namely the channel width 042, the width of the A? distribution AA’, and the shortest wavelength squared a "," When planning a rotation-measure experiment, three main parameters are involved, namely the channel width $\delta \lambda^2$ , the width of the $\lambda^2$ distribution $\Delta \lambda^2$, and the shortest wavelength squared $\lambda_{min} ^2$."
+"They respectively determine the maximum observable Faraday depth, the resolution in Faraday space, and the largest scale in Faraday space to which the observation is sensitive."," They respectively determine the maximum observable Faraday depth, the resolution in Faraday space, and the largest scale in Faraday space to which the observation is sensitive."
+" If we assume a top hat weight function that is 1 between 2 and 42,,, and zero elsewhere, the estimates of the FWHM of the main peak of the RMSF, the scale in Faraday space to which the sensitivity has dropped to50%,, and the maximum Faraday depth to which one has more than sensitivity are approximately In Table 4,, we list the parameters for the three ranges of wavelengths used for RM-synthesis."," If we assume a top hat weight function that is 1 between $\lambda_{min} ^2$ and $\lambda_{max} ^2$ and zero elsewhere, the estimates of the FWHM of the main peak of the RMSF, the scale in Faraday space to which the sensitivity has dropped to, and the maximum Faraday depth to which one has more than sensitivity are approximately In Table \ref{FWHM}, we list the parameters for the three ranges of wavelengths used for RM-synthesis."
+ In Fig., In Fig.
+" 3 we show the RMSF of the 18cm+21cm+25cm, 85 cm, and 2 m datasets."," \ref{RMSF} we show the RMSF of the 18cm+21cm+25cm, 85 cm, and 2 m datasets."
+" Due to the sparse sampling in 22, the RM cubes have sidelobes of —4096, ~15%—20%,, and ~60% at 2 m, 85 cm, and 18cm+21cm+25cm, respectively."," Due to the sparse sampling in $\lambda^2$, the RM cubes have sidelobes of $\sim40$, $\sim15$, and $\sim60$ at 2 m, 85 cm, and 18cm+21cm+25cm, respectively."
+" These sidelobes can be cleaned through a 1-dimensional deconvolution, which is an extension of ? CLEAN to the complex domain."," These sidelobes can be cleaned through a 1-dimensional deconvolution, which is an extension of \citet{1974A&AS...15..417H} CLEAN to the complex domain."
+ We used public available made available by George Heald to perform the deconvolution., We used public available made available by George Heald to perform the deconvolution.
+ The procedure follows., The procedure follows.
+ Table 5 lists the results of the analysis of the 30 background sources observed aroundA2255., Table \ref{30sourcestable} lists the results of the analysis of the 30 background sources observed around.
+. Seven out of the 30 sources have a doublestructure?., Seven out of the 30 sources have a double.
+". In this case, we computed the RM of each component (labeled A and B in Table 5))."," In this case, we computed the RM of each component (labeled A and B in Table \ref{30sourcestable}) )."
+ The histogram of the RM values for the 30 sources is presentedin Fig., The histogram of the RM values for the 30 sources is presentedin Fig.
+ ?? and Fig., \ref{rmhist} and Fig.
+ ?? shows the results., \ref{30sourcesfig} shows the results.
+" At the location of each source,"," At the location of each source,"
+Typically in the spectra of bright QSOs. several lines may be identified as beeing due to absorption of intervening material situated along the line of sight.,"Typically in the spectra of bright QSOs, several lines may be identified as beeing due to absorption of intervening material situated along the line of sight."
+ These lines originate when the continuum radiation meets a common atom or ion in ils eround state that partially absorbs it. leaving its imprint in the emiting QSO's spectrum.," These lines originate when the continuum radiation meets a common atom or ion in its ground state that partially absorbs it, leaving its imprint in the emiting QSO's spectrum."
+ However. if some excitation mechanism is present in the absorbing region. then a small fraction of atoms or ions will also be found. populated. in their Iowest-Iving excited levels.," However, if some excitation mechanism is present in the absorbing region, then a small fraction of atoms or ions will also be found populated in their lowest-lying excited levels."
+ Therefore. in. addition to absorption lines arising from the atom/ion's ground state. one mav also expect to detect weaker lines arising. [rom excited levels.," Therefore, in addition to absorption lines arising from the atom/ion's ground state, one may also expect to detect weaker lines arising from excited levels."
+ lt has been long pointed out that fine-structure absorption lines arising [rom the ground and. lowest-bing excited enerey levels of common atoms/ions may be used as an indicator of the physical conditions in the eas (Baheall&WolfLOGS:SmedingPot," It has been long pointed out that fine-structure absorption lines arising from the ground and lowest-lying excited energy levels of common atoms/ions may be used as an indicator of the physical conditions in the gas \cite{BW,SP}."
+"tasch 1979).. Η we model the absorbing region as a single. homogeneous cloud. then the ratio of the volumetric densities of atoms/ions populated in excited. states n to atoms/ions in the ground state n will match the corresponding column density ratios: For example. the column densities of tons populated inH theirH ground 2“PY,"" and first. excitedH 2“PY.D levels may be inferred from the equivalent widths of the corresponding ⊳↘ −−⊐⊐↴ ⇉⊳∖−⇉↓≻−↓↴⋡⊥⊐⋅−≽⊳∖−≻↓≻∐⊽↾⊁⊐⋜⋯∠∟≻⊳∖−≽↓≻↓↴⇡⋝⊐−⊐−⊐∣κ 7DEos "," If we model the absorbing region as a single, homogeneous cloud, then the ratio of the volumetric densities of atoms/ions populated in excited states $n^*$ to atoms/ions in the ground state $n$ will match the corresponding column density ratios: For example, the column densities of $^+$ ions populated in their ground $^2\mathrm{P}^o_{1/2}$ and first excited $^2\mathrm{P}^o_{3/2}$ levels may be inferred from the equivalent widths of the corresponding $^2$ 2p $^2\mathrm{P}^o_{1/2}\rightarrow$ $^2$ ${^2\mathrm{D}^e_{3/2}}$ and $^2$ 2p $^2\mathrm{P}^o_{3/2}\rightarrow$ $^2$ ${^2\mathrm{D}^e_{5/2}}$ "
+"following discussion, were strictly estimated using returned values for peak intensities while line widths were simply not considered.","following discussion, were strictly estimated using returned values for peak intensities while line widths were simply not considered."
+" To assure the selection of the most reliable emission-line profiles in our sample of over 0000 spectra, a series of conditions are proposed."," To assure the selection of the most reliable emission-line profiles in our sample of over 000 spectra, a series of conditions are proposed."
+ of them need to be fulfilled otherwise the given spectrum is rejected from the discussion to follow., of them need to be fulfilled otherwise the given spectrum is rejected from the discussion to follow.
+ The conditions are summarized as follow: Two conditions are proposed in order to confirm that, The conditions are summarized as follow: Two conditions are proposed in order to confirm that
+The Laser Latererometer Cravitational Wave Observatory (LIGO) achieves its design sensitivity cltwing its fifth science run S5: (Abbotectal. 2009c)]].,The Laser Interferometer Gravitational Wave Observatory (LIGO) achieved its design sensitivity during its fifth science run [S5; \citep{lscinstrument09}] ].
+. Anavsls of 85 data is progressing well. with new upper limits being placed on the strength. of various classes of burst sources (Abbottetal.2000d:Axwlie2010a:Abvolt 2010a).. stochastic backgrounds. (Campanis2WS:Abbott 2009a).. compa« ‘binary sources (Abbottetal.20096:Abaclicetal.200ο.) and continuous-wave sources (Abbottetal.2009£b.2Ob:Abaclic201 )»)..," Analysis of S5 data is progressing well, with new upper limits being placed on the strength of various classes of burst sources \citep{lscburst09, lscburst10, lscburst10b}, , stochastic backgrounds \citep{giampanis08, lscstochastic09}, compact binary sources \citep{lsccbc09, lscinspiral10, lsccbc10} and continuous-wave sources \citep{allsky09, lsceah09, knownpul09, casa10}."
+ In some cases. the LIGO limits on astrophysical parameters beat those inferre from electromagnetic astronomy. e.g. the maximum ellipticity and internal magnetic field sreneth of the Crab pulsar (Abbottetal.2008.2010b).," In some cases, the LIGO limits on astrophysical parameters beat those inferred from electromagnetic astronomy, e.g. the maximum ellipticity and internal magnetic field strength of the Crab pulsar \citep{crab08, knownpul09}."
+. Recenth. an 85 search was completed which placed upper limis on the amplitucle of r-mode oscillations of the neutron star in the supernova remnant C'assiopela A (Abaclicetal.2010b)..Aspherical. isolated neutron stars. constitute one promising class of continuous-wave source. candidates (Ostriker&Gunn 1969)..," Recently, an S5 search was completed which placed upper limits on the amplitude of $r$ -mode oscillations of the neutron star in the supernova remnant Cassiopeia A \citep{casa10}.Aspherical, isolated neutron stars constitute one promising class of continuous-wave source candidates \citep{ostriker69}. ."
+ The origin of the semi-permanent cquacdrupole in these objects can be thermoelastic (MelatosHaskell2008) or hyedromagnetic (Bonazzola&CGourgoul-Akeün&Wasserman2008:Mastrano 2010).," The origin of the semi-permanent quadrupole in these objects can be thermoelastic \citep{melatos00, ushomirsky00, nayyar06, haskell08} or hydromagnetic \citep{bonazzola96, cutler02, haskellb08, haskell08, akgun08, mastrano10}."
+.. Phermoclastic deformations arise due to uneven electron capture rates in the neutron star crust., Thermoelastic deformations arise due to uneven electron capture rates in the neutron star crust.
+ A persistent temperature gracdien at the base of the crust produces a mass quadrupole momen of  LOee emer (c~10':2000)., A persistent temperature gradient at the base of the crust produces a mass quadrupole moment of $\sim 10^{38}$ g $^{2}$ \citep[$\epsilon \sim 10^{-7}$.
+ Lvclromagnetic deformations. on the other hand. are produced by large internal magnetic fields. and. misalignec magnetic and spin axes.," Hydromagnetic deformations, on the other hand, are produced by large internal magnetic fields, and misaligned magnetic and spin axes."
+" For cxample. a neutron star with spin frequency 300111. anc internal toroidal field D, GG has an ellipticity «€~10""(BA10) (Cutler 2002)."," For example, a neutron star with spin frequency Hz and internal toroidal field $B_t \gtrsim 3.4 \times 10^{12}$ G has an ellipticity $\epsilon \sim 10^{-6} \left(\langle B_t \rangle / 10^{15}\,\text{G}\right)$ \citep{cutler02}."
+. Phe deformation of an ideal [uid starwith an arbitrary magnetic field. distributionancl a barotropic, The deformation of an ideal fluid starwith an arbitrary magnetic field distributionand a barotropic
+ccan serve as a diagnostic tool to distinguish AGN from starburst activity.,can serve as a diagnostic tool to distinguish AGN from starburst activity.
+ In sections 2... 3. and 4. the sample. observations and their results are presented.," In sections \ref{s:sample}, , \ref{s:obs} and \ref{s:res} the sample, observations and their results are presented."
+ lline parameters are presented in section4.1 and in section 4.2 ccolumn densities and fractional abundances are calculated., line parameters are presented in section\ref{s:line} and in section \ref{s:column} column densities and fractional abundances are calculated.
+ In sections 5.1 and 5.2 aabundances in the context of X-ray and UV irradiated models are discussed and in 5.3. the potential importance of grain chemistry., In sections \ref{s:xdr} and \ref{s:hcop} abundances in the context of X-ray and UV irradiated models are discussed and in \ref{s:grain} the potential importance of grain chemistry.
+ In the last section. 5.7.. we present a brief future outlook.," In the last section, \ref{s:future}, we present a brief future outlook."
+ We have selected a sample consisting of seven nearby luminous starburst and AGNs - and one distant ULIRG., We have selected a sample consisting of seven nearby luminous starburst and AGNs - and one distant ULIRG.
+ The galaxies are selected from their bright HCN line emission., The galaxies are selected from their bright HCN line emission.
+ From our previous experience with extragalactic wwe knew that the line is weaker than the standard high density gas tracers such as HCN and sso we restricted ourselves to a relatively small sample of seven objects (coordinates. FIR luminosities and distances are presented in Table 1)): is anearby barred Sed galaxy of moderate lumimosity (central 400 pe has {μμ of 6xLOS La) with a central starburst.," From our previous experience with extragalactic we knew that the line is weaker than the standard high density gas tracers such as HCN and so we restricted ourselves to a relatively small sample of seven objects (coordinates, FIR luminosities and distances are presented in Table \ref{t:sample}) is anearby barred Scd galaxy of moderate luminosity (central 400 pc has $L_{\rm FIR}$ of $6 \times 10^8$ $_{\odot}$ ) with a central starburst."
+ Within its central 300. pe. (30°)) two molecular arms enc in a clumpy central ring of dense gas (e.g.?) which surrounds a young star cluster., Within its central 300 pc ) two molecular arms end in a clumpy central ring of dense gas \citep[e.g.][]{downes92} which surrounds a young star cluster.
+ The ring is suggested to outline the X2 orbits in a larger-scale bar., The ring is suggested to outline the $X2$ orbits in a larger-scale bar.
+ The chemistry of IC 342 has been investigated in detail in a high-resolution study by ?.., The chemistry of IC 342 has been investigated in detail in a high-resolution study by \citet{meier05}.
+ They find that the chemistry in the ring is à mixture of PDR-dominated regions and regions of younger star-forming clouds., They find that the chemistry in the ring is a mixture of PDR-dominated regions and regions of younger star-forming clouds.
+ The chemistry in the bar/arms ts dominated by shocks as shown by CH30OH (?) and SiO (?).., The chemistry in the bar/arms is dominated by shocks as shown by $_3$ OH \citep{meier05} and SiO \citep{usero06}.
+ Five (A-E) giant molecular clouds (GMCs) are found in the ring and arms., Five (A-E) giant molecular clouds (GMCs) are found in the ring and arms.
+ The JJCMT beam of our oobservations is pointed towards the region of GMC B- but also includes GMCs A and E. The dust temperature of IC 342 is estimated to 42 K (e.g.?).., The JCMT beam of our observations is pointed towards the region of GMC B- but also includes GMCs A and E. The dust temperature of IC 342 is estimated to 42 K \citep[e.g.][]{becklin80}.
+ is also a nearby barred galaxy located in the Sculptor group with a compact nuclear starburst and a IR luminosity that appears to originate in regions of intense massive star formation within its central few hundred parsees (?).., is also a nearby barred galaxy located in the Sculptor group with a compact nuclear starburst and a IR luminosity that appears to originate in regions of intense massive star formation within its central few hundred parsecs \citep{strickland04}.
+ From their 2 mm spectral scan ? suggest that the chemistry of NGC 253 shows strong similarities to that of the Galactic Center molecular region. which is thought to be dominated by low-velocity shocks.," From their 2 mm spectral scan \citet{martin06} suggest that the chemistry of NGC 253 shows strong similarities to that of the Galactic Center molecular region, which is thought to be dominated by low-velocity shocks."
+ High resolution SiO observations show bright emission resulting from large scale shocks as well as gas entrained in a nuclear outflow (?).., High resolution SiO observations show bright emission resulting from large scale shocks as well as gas entrained in a nuclear outflow \citep{garcia-burillo00}.
+ Note also that it is suggested that NGC 253 is a galaxy in which a strong starburst and a weak AGN coexist (e.g.22)..," Note also that it is suggested that NGC 253 is a galaxy in which a strong starburst and a weak AGN coexist \citep[e.g.][]{weaver02,muller-sanchez10}."
+ High resolution observations of HCN and 11-0 (?). show strongly centrally concentrated emission., High resolution observations of HCN and 1–0 \citep{knudsen07} show strongly centrally concentrated emission.
+ The JJCMT beam covers the bulk of the nuclear emission from HCN and., The JCMT beam covers the bulk of the nuclear emission from HCN and.
+. The central dust temperature of NGC 253 ts estimated to 50 K (?).., The central dust temperature of NGC 253 is estimated to 50 K \citep{melo02}.
+ is the nearest example of a type 2 Seyfert galaxy luminous in the infrared., is the nearest example of a type 2 Seyfert galaxy luminous in the infrared.
+ Surrounding the AGN there ts a ccircumnuclear molecular ring or -disk (CND) and on a larger scale there is a NIR stellar bar 2.3 kpe long., Surrounding the AGN there is a circumnuclear molecular ring or -disk (CND) and on a larger scale there is a NIR stellar bar 2.3 kpc long.
+ This bar is connected to a large-scale. molecular starburst ring that contributes about half the bolometric luminosity of the galaxy (e.g.22222).," This bar is connected to a large-scale, molecular starburst ring that contributes about half the bolometric luminosity of the galaxy \citep[e.g.][]{telesco84, scoville88, tacconi94, helfer95, tacconi97}."
+ Bright HCN 1-0 line emission is observed towards the CND while the 11-0 emission is relatively fainter by a factor of 1.5 (e.g.?).., Bright HCN 1–0 line emission is observed towards the CND while the 1–0 emission is relatively fainter by a factor of 1.5 \citep[e.g.][]{kohno01}.
+ Only the CND and a fraction of the bar ts covered by our JCMT beam and we adopt a source size of ffor the CND., Only the CND and a fraction of the bar is covered by our JCMT beam and we adopt a source size of for the CND.
+" The dust of the inner aappears to show a strong temperature gradient - from about 800 K in the very inner region to 150 - 275 K at a distance of 0.""8 or greater (2)..", The dust of the inner appears to show a strong temperature gradient - from about 800 K in the very inner region to 150 - 275 K at a distance of 8 or greater \citep{tomono06}.
+ 2 find temperatures of about 150 K 200 pc from the nucleus., \citet{alloin00} find temperatures of about 150 K 200 pc from the nucleus.
+ There is a also radio jet from the nucleus which falls into our JCMT beam (e.g.?).., There is a also radio jet from the nucleus which falls into our JCMT beam \citep[e.g.][]{wilson83}.
+ is a luminous. edge-on. Sa-type galaxy with a deeply dust-enshrouded nucleus (?)..," is a luminous, edge-on, Sa-type galaxy with a deeply dust-enshrouded nucleus \citep{spoon01}. ."
+ NGC 4418 is a FIR- galaxywith a logarithmic IR-to-radio continuum ratio (6) of 3 (2). , NGC 4418 is a FIR-excess galaxywith a logarithmic IR-to-radio continuum ratio $q$ ) of 3 \citep{roussel03}. .
+This excess may be caused by either a young pre-supernova starburst or a buried AGN (???)..," This excess may be caused by either a young pre-supernova starburst or a buried AGN \citep{aalto07b, roussel03, imanishi04}. ."
+ Unusually, Unusually
+Tjo lnass of a star is the simele inost iuportaut paralucter in deteriuuing how it will interact with its environnent. how lone it will live. aud the nature of its ¢cath.,"The mass of a star is the single most important parameter in determining how it will interact with its environment, how long it will live, and the nature of its death."
+ Therefore the distribution of masses of newly formred starsthe initial mass function (IME)has far reacing duplications for the evolution of the cosmos., Therefore the distribution of masses of newly formed stars—the initial mass function (IMF)—has far reaching implications for the evolution of the cosmos.
+ The TALF is usually asstuned to be universal. with a shape deseribed by a power-lav above ~1krod2.. and a loe-normal below2003).," The IMF is usually assumed to be universal, with a shape described by a power-law above $\sim 1$, and a log-normal below."
+ Stars form from molecular clouds. aid therefore hei nature and characteristics define the initial condi10115 for star formation.," Stars form from molecular clouds, and therefore their nature and characteristics define the initial conditions for star formation."
+ The hierarchical desity and velcity structure of molecular clouds is iudicaIve of supoersonic turbuleuce2000)., The hierarchical density and velocity structure of molecular clouds is indicative of supersonic turbulence.
+.. However. cold. dense regiois of quiescent eas. or “cores.” found withiji Galactic clouds are believed to he the sites of future low-nuass star formation1991).," However, cold, dense regions of quiescent gas, or “cores,” found within Galactic clouds are believed to be the sites of future low-mass star formation."
+. A lareeOo hun continu survey of deuse cores in the p OOphiuchus star-forming region by revealed a core mass distribution with a declining powcr-law slope siuilar to the IAIF., A large mm continuum survey of dense cores in the $\rho$ Ophiuchus star-forming region by revealed a core mass distribution with a declining power-law slope similar to the IMF.
+ This led the authors to conclude. that the cores observed iu thermal dust enission are the direct progenitors of individual stars or stellar svstenis., This led the authors to conclude that the cores observed in thermal dust emission are the direct progenitors of individual stars or stellar systems.
+ Tie deuse core mass fiction (DCAIF) in other regions iis also exhibited similarities to the IME2O07)., The dense core mass function (DCMF) in other regions has also exhibited similarities to the IMF.
+. Cores in the Pipe Nebula ideutified throπο dust extinction show a turnover at masses about a factor of 3 higier than the IAIF of the Trapezimu cluscr leading to the interpretation that cores evolve with a nearly cousaut star formation efficienev of ~30%2007).," Cores in the Pipe Nebula identified through dust extinction show a turnover at masses about a factor of 3 higher than the IMF of the Trapezium cluster leading to the interpretation that cores evolve with a nearly constant star formation efficiency of $\sim
+30$."
+. As first suggested by(L991). he deusitv proability. distribution function (PDF) iu isothermal turbulent flows 1s expected to be lognormal providing heoretical meai* to prodce the low-mass ος of the DCME.," As first suggested by, the density probability distribution function (PDF) in isothermal turbulent flows is expected to be log-normal providing theoretical means to produce the low-mass end of the DCMF."
+ The power-law tai of the DCAIF can also be explained in terus of post-shock eas within a turbulent ποσα2002).. or bv deviations rou Isothermality1998).," The power-law tail of the DCMF can also be explained in terms of post-shock gas within a turbulent medium, or by deviations from isothermality."
+. In addition. the star formation cficiency of cores is theoretically expected o be nearly consaut and lie between if: outflows from xotostars are the primary imeciating Actor20001.," In addition, the star formation efficiency of cores is theoretically expected to be nearly constant and lie between if outflows from protostars are the primary mediating factor."
+ These results zipport a one-to-one or nearly ouce-to-o relatiouship betwi(en dense cores and the future stars to ori from them., These results support a one-to-one or nearly one-to-one relationship between dense cores and the future stars to form from them.
+ However. the similarity of the DCAIF to he IME remains the only piece ofobservational evidence or this kiud of reiiouship.," However, the similarity of the DCMF to the IMF remains the only piece of evidence for this kind of relationship."
+ \leamwhile. there are several reasons to thins that this one-to-one relationship may vot hold.," Meanwhile, there are several reasons to think that this one-to-one relationship may not hold."
+ It is clear ha sole deuse cores must fragment to xoduce the large fraction of observed inultiple stellar systems2007).. aud it is possible that several fragieuts per core nay be necessary to explain close binary svstenis2005).," It is clear that some dense cores must fragment to produce the large fraction of observed multiple stellar systems, and it is possible that several fragments per core may be necessary to explain close binary systems."
+. Unfortunately. uost observaions of pre-stellar cores are dHunited to spatial resolutkAs orders of magnitude ereater than characteristic binary separations.," Unfortunately, most observations of pre-stellar cores are limited to spatial resolutions orders of magnitude greater than characteristic binary separations."
+" It is also cff""ult to determine what fraction of cores ideutified i sumον data will evolve iuto stars.", It is also difficult to determine what fraction of cores identified in survey data will evolve into stars.
+ Iu a recent y.udyv of cores iu the Pipe Nebula. found that he majority of cores are eravitationally uubound. with only the highest mass cores appearing destined to forma stars.," In a recent study of cores in the Pipe Nebula, found that the majority of cores are gravitationally unbound, with only the highest mass cores appearing destined to form stars."
+ The most massive core in the Pipe Nebula. Barnard 59. is the sole active core in the nebula aud hartOTS ALL aSSOCIatioN of ~20 τοις starsWOT).," The most massive core in the Pipe Nebula, Barnard 59, is the sole active core in the nebula and harbors an association of $\sim 20$ young stars."
+ Despite these open questions. the iutrigung similarity," Despite these open questions, the intriguing similarity"
+the cluster) prompted us to use it to measure the distauce to the Pleiades.,the cluster) prompted us to use it to measure the distance to the Pleiades.
+ In thisLetter we present accurate new B aud V ophotoclectric photometry aud radial velocities of TD 23612. and use them to derive an orbital solution ancl iufer the distance to the star. aud therefore to the cluster.," In this we present accurate new $B$ and $V$ photoelectric photometry and radial velocities of HD 23642, and use them to derive an orbital solution and infer the distance to the star, and therefore to the cluster."
+ We observed IID 23612 in B and V. (standard Jolnsou filters) from a private obscrvatory near Cembra (Trento). Italy.," We observed HD 23642 in $B$ and $V$ (standard Johnson filters) from a private observatory near Cembra (Trento), Italy."
+ The instrament was a 28 ocn Schinidt-Casscerain telescope equipped with an Optec SSP5 photometer., The instrument was a 28 cm Schmidt-Cassegrain telescope equipped with an Optec SSP5 photometer.
+ It proved already to be a very accurate and reliable iustruineut (cf, It proved already to be a very accurate and reliable instrument (cf.
+ Siviero et al., Siviero et al.
+ 2001) perfectly suited to deal with the low amplitude of WD 23612 eclipses (0.1 mae)., 2004) perfectly suited to deal with the low amplitude of HD 23642 eclipses $\sim$ 0.1 mag).
+ 223568 117661. Ve=6.82 140.011. Br=6.8[2 0.015. spectrum D9.5Vj was chosen as comparison star aud 223763 117791. Ve =6.96340.011. Br=7.109+ 0.016. spectrum AIV) as a check star.," 23568 17664, $V_T$ $\pm$ 0.011, $B_T= 6.842\pm0.015$ , spectrum B9.5V) was chosen as comparison star and 23763 17791, $V_T$ $\pm$ 0.011, $B_T=7.109\pm 0.016$ , spectrum A1V) as a check star."
+ Following the Bessell (2000) rausformatious jetween Tyeho aud Johnson plotometric svstenis. we aopted. V —6.830 and B-6.851 for the comparison star.," Following the Bessell (2000) transformations between Tycho and Johnson photometric systems, we adopted $V$ =6.830 and $B$ =6.851 for the comparison star."
+" The comparison sTab Was LHileasurec against the check star oa least OLnco OCVCOYV observing run. and fouud constaut with standard deviatious of 0.005 mag in and 0.006. mag in D. coufirmine the Ilipparcos photometric results,"," The comparison star was measured against the check star at least once every observing run, and found constant with standard deviations of 0.005 mag in $V$ and 0.006 mag in $B$, confirming the Hipparcos photometric results."
+ Tn all 192 measurements iu V. aud 132 in D were collected of ΠΟ 23612 between Aue. 19. 2003 aud Feb. 16. 2001.," In all, 492 measurements in $V$, and 432 in $B$ were collected of HD 23642 between Aug. 19, 2003 and Feb. 16, 2004."
+ All the observations were corrected for émospleric extinction and imstrmucutal color equations (via calibration ou Laudolt's equatorial fields). and the iustruineutal differential macuitucles were transformed iuto the standard Johuson UDV system.," All the observations were corrected for atmospheric extinction and instrumental color equations (via calibration on Landolt's equatorial fields), and the instrumental differential magnitudes were transformed into the standard Johnson UBV system."
+ The variable. comparison and check stars are similar iu spectral ype. hie verv close ou the sky (15 arcnün) aud were always observed at zenith distances <607. which argues for a lieh cousisteney of our photometry.," The variable, comparison and check stars are similar in spectral type, lie very close on the sky (15 arcmin) and were always observed at zenith distances $<60^\circ$, which argues for a high consistency of our photometry."
+ Ilieh resolution spectra of TD 23612 were obtained at five distinct epochs with the ELODIE echelle spectrograph at the 1.9312 telescope of the Maute-Provence Observatory (Baranne ct al., High resolution spectra of HD 23642 were obtained at five distinct epochs with the ELODIE echelle spectrograph at the 1.93m telescope of the Haute-Provence Observatory (Baranne et al.
+ 1996)., 1996).
+ ELODIE covers the spectral ranec 3900-6800. lu a single exposure as 67 orders at a lean resolving power of 12000., ELODIE covers the spectral range 3900-6800 in a single exposure as 67 orders at a mean resolving power of 42000.
+ Optimal extraction and waveleugth calibration were performed with the on-line automatic reduction pipeline., Optimal extraction and wavelength calibration were performed with the on-line automatic reduction pipeline.
+ Racial velocities for the Wo conmmponueuts are reported in Table 1., Radial velocities for the two components are reported in Table 1.
+ They were measured by both classical eross-correlation against suitable A-type templates aud by the Least-Square Decouvolution method of Donati et al (1997)., They were measured by both classical cross-correlation against suitable A-type templates and by the Least-Square Deconvolution method of Donati et al (1997).
+ The errors given in Table 1 reflect the scatter of the measurements obtained with the Wo methods run with different parameters., The errors given in Table 1 reflect the scatter of the measurements obtained with the two methods run with different parameters.
+ One parameter. the temperature of the primary (Ti). cannot be directly modeled in SB2 EB analysis aud uiust be determined independently.," One parameter, the temperature of the primary $T_1$ ), cannot be directly modeled in SB2 EB analysis and must be determined independently."
+ Tu the case of ΠΟ 23612. the spectral classification reported in literature is too sparse. ranging from B9 to Al.," In the case of HD 23642, the spectral classification reported in literature is too sparse, ranging from B9 to A1."
+ To determine 7i we hen turned to photometry aud searched the literature aud the General Catalogue of Photometric Data (OCPD. AMeruiulliod et al.," To determine $T_1$ we then turned to photometry and searched the literature and the General Catalogue of Photometric Data (GCPD, Mermilliod et al."
+ 1997b) for observations of ΠΟ 23612., 1997b) for observations of HD 23642.
+" Data were found in the WalravenVBELUW. Strouunercavbey WDUVR. Geneva, DDO. Vilnius. Johusou.CDVRI. Tycho. 2\LASS and ASN 7Ll photometric svsteiis (Crawford Perry 1976. Ikoruilov et al."," Data were found in the Walraven, Strömmgren, Geneva, DDO, Vilnius, Johnson, Tycho, 2MASS and ASN 74 photometric systems (Crawford Perry 1976, Kornilov et al."
+ 1991. Morel Magnuenat 1978. Persinecr Castelaz 1990. van οσοι ct al.," 1991, Morel Magnenat 1978, Persinger Castelaz 1990, van Leeuwen et al."
+ 1986. Wesselius et al.," 1986, Wesselius et al."
+ 1982. GCPD).," 1982, GCPD)."
+ Using the relevant parameters for these, Using the relevant parameters for these
+population). throughout the Local Croup.,"population), throughout the Local Group."
+ The second is near-infrarecl single-star photometry using 2MASS data., The second is near-infrared single-star photometry using 2MASS data.
+" We selected all stars within a radius of 1"" of the CIIVC »ositionis for analysis.", We selected all stars within a radius of $^o$ of the CHVC positions for analysis.
+ The data are presented and cliscussed in section 3., The data are presented and discussed in section 3.
+ We show that the 2NLXSS data are sensitive o intermediate and. old. stellar populations in the dwarf Spheroidal companions of the Alilky Way., We show that the 2MASS data are sensitive to intermediate and old stellar populations in the dwarf Spheroidal companions of the Milky Way.
+ Since our observations do not reveal a stellar content in five LEVC's. we conclude in section 4 that the five objects studied: by us did not experience any star formation over cosmic time.," Since our observations do not reveal a stellar content in five HVCs, we conclude in section 4 that the five objects studied by us did not experience any star formation over cosmic time."
+ Even with such a small sample. we can alreacy place some interesting constraints on the putative stellar content of LIVC's.," Even with such a small sample, we can already place some interesting constraints on the putative stellar content of HVCs."
+ We submied a target list of southern. ος lor VLT service observing to the good. «0766. secing queue.," We submitted a target list of southern CHVCs for VLT service observing to the good, $<$ 6, seeing queue."
+ Lere. we report on the results of our first semester of observations. with a fall right ascension coverage.," Here, we report on the results of our first semester of observations, with a fall right ascension coverage."
+ The results of VET service mode observations to be conducted in the spring of 2003 will be the subject of our paper LL., The results of VLT service mode observations to be conducted in the spring of 2003 will be the subject of our paper II.
+ In order to minimise contamination of the images hy foreground stars in the Galaxy. we chose targets at Galactic latitucles above 30°].," In order to minimise contamination of the images by foreground stars in the Galaxy, we chose targets at Galactic latitudes above $|$ $^o$$|$."
+ Phe target list was drawn [rom Braun Burton (1999) and emphasises CIIVCs observed. with radio interferometers from Braun Burton (2000). Drünns et al. (," The target list was drawn from Braun Burton (1999) and emphasises CHVCs observed with radio interferometers from Braun Burton (2000), Brünns et al. ("
+2000). and from unpublished LIE observations by one ol us (Ixerp).,"2000), and from unpublished HI observations by one of us (Kerp)."
+ We added to our target list. LILPASS J1712-64.," We added to our target list, HIPASS J1712-64."
+ Ixilborn et al. (, Kilborn et al. (
+2000) crew attention to this Isolated. star-free cloud of neutral hydrogen. which they report could be at à distance of 3.2 Alpe (based on its HE velocity and Llubble’s law).,"2000) drew attention to this isolated, star-free cloud of neutral hydrogen, which they report could be at a distance of 3.2 Mpc (based on its HI velocity and Hubble's law)."
+ We obtained observations of five LVC's (including LUPASS J1712-64)., We obtained observations of five HVCs (including HIPASS J1712-64).
+ Details of the targets are presented. in ‘Table 1., Details of the targets are presented in Table 1.
+ The coordinates reflect. the J2000 centres of our VLT poitings., The coordinates reflect the J2000 centres of our VLT poitings.
+ According to the naming convention of Braun Burton (1999). the name consists of a three-cigit Galactic ongitude. followed. by a two-digit Galactic latitude. and a three-digit’ local-stancard-of-rest. velocity.," According to the naming convention of Braun Burton (1999), the name consists of a three-digit Galactic longitude, followed by a two-digit Galactic latitude, and a three-digit local-standard-of-rest velocity."
+ The Galactic coordinates lor LUPASS J1712-64 are 122326755. bzz- 14:90.," The Galactic coordinates for HIPASS J1712-64 are $\approx$ 5, $\approx$ 5."
+ eotice that this cloud is at a lower Galactic latitude than he remainder of the sample., Notice that this cloud is at a lower Galactic latitude than the remainder of the sample.
+" Phe Galactic absorptions were obtained usingNED"".", The Galactic absorptions were obtained using.
+.. NED provides Calactic extinctions rom Schlegel. Finkenbeiner Davis (1998). and assumes a Galactic extinction law with Ry-=3.1 (6ανάσα. Clavton Mathis 1989).," NED provides Galactic extinctions from Schlegel, Finkenbeiner Davis (1998), and assumes a Galactic extinction law with $_V$ =3.1 (Cardelli, Clayton Mathis 1989)."
+ The observations were carried out at the VET. UTI CANTU with the FORSI instrument.," The observations were carried out at the VLT UT1 “ANTU"" with the FORS1 instrument."
+ They. consisted. of dithered. exposures in Land. V. for each of the CIIVCS.," They consisted of dithered exposures in I and V, for each of the CHVCs."
+ The intregration times measured 5 300s in Land ο 300s in V. except for LEVC'039-33-260. for which they were δι 200s in ] and 3 200s in V. Each observation was accompanied bv a quality. file.," The intregration times measured 5 $\times$ 300s in I and 3 $\times$ 300s in V, except for HVC039-33-260, for which they were 5 $\times$ 200s in I and 3 $\times$ 200s in V. Each observation was accompanied by a quality file."
+ We checked. that the conditions during the observations were photometric. and that the. secing ∖∖⊽⋜↧⊳∖⊔∪∖∖⊽∪↓⋅⊳∖∢⋅↿," We checked that the conditions during the observations were photometric, and that the seeing was no worse than 6."
+↓⋯⊔∪⋅↻↻⊳∖∖⋖⊾⊔⊳∖⋯⇂⇂↓∐⋅≻↓↓≻∢⊾↓↓⊔⋖⊾−↓⋅⋖⊾∠⇂⋯⇍⋯ ∕∕⋅⋅⇁ ⋠⋠ἱ data in our analysis., We used the pipeline-reduced data in our analysis.
+ The individual exposures were shifted ancl added., The individual exposures were shifted and added.
+ The. hand cata needed: additional [lat-field correction bevond. that already applied in. the pipeline., The I-band data needed additional flat-field correction beyond that already applied in the pipeline.
+ Single-star photometry was carried out with DAODPLIIOT IL (Stetson 1992)., Single-star photometry was carried out with DAOPHOT II (Stetson 1992).
+ We used. isolated stars in. each field. in order to determine the point-spread function. (PSE)., We used isolated stars in each field in order to determine the point-spread function (PSF).
+ We find that the EWLIAL was between 07552 and 058. with a mean of ⋅∕∕⋅⋅07665.," We find that the FWHM was between 52 and 8, with a mean of 65."
+ Photometric. zero points. for⋅ cach night. and in each filter were obtained from the European Southern Observatory’s (ESO) homepage., Photometric zero points for each night and in each filter were obtained from the European Southern Observatory's (ESO) homepage.
+ Extinction. terms in. V and L and V-LI colour terms needed. to transform to the Johnson-Cousins photometric svstem. were gleaned from he ESO homepage as well.," Extinction terms in V and I, and V-I colour terms needed to transform to the Johnson-Cousins photometric system, were gleaned from the ESO homepage as well."
+ We used the PSE stars to find he zero points. determine aperture corrections. and. define he transformation of the data into the Johnson-C'ousins Vo] system.," We used the PSF stars to find the zero points, determine aperture corrections, and define the transformation of the data into the Johnson-Cousins V, I system."
+ Phe DAOPIIOT. parameters sharpness and roundness were used to clip extended: sources. io... mainly xwkeround. galaxies. with sizes lareer than the PSE from he data tables.," The DAOPHOT parameters sharpness and roundness were used to clip extended sources, i.e., mainly background galaxies, with sizes larger than the PSF from the data tables."
+ The last column in Table 1 (N stars) lists or each. CLIVC'. the number of objects elassed to be stars and detected in V and 1. Our data achieve limiting magnitudes in Vz:25.26. and in 12324.," The last column in Table 1 (N stars) lists for each CHVC, the number of objects classed to be stars and detected in V and I. Our data achieve limiting magnitudes in $\approx$ 25–26, and in $\approx$ 23–24."
+ In Figure 1. we displav the V-L. 1] CMDs of the four. λος Located. at similar. Galactic latitudes (sce Table 1).," In Figure 1, we display the [V-I, I] CMDs of the four CHVCs located at similar Galactic latitudes (see Table 1)."
+ In. Figure 2. we show the V-L. 1] CALD of LIIPASS J1712-64.," In Figure 2, we show the [V-I, I] CMD of HIPASS J1712-64."
+ As an assessment of the data quality we reproduce in Figure 3 the DAOPLIIOT photometric errors in, As an assessment of the data quality we reproduce in Figure 3 the DAOPHOT photometric errors in
+Tere ⋅⋅ =o Gap Faery and σ. are given in Tables 1 aud ? L 2 of |16]..,"Here where and The quantities $m_{B}^{corr}$ , $m_{B}^{eff}$ , $\sigma
+_{m_{B}^{corr}}$, $ \sigma _{m_{B}^{eff}}$ and $\sigma _{z}$ are given in Tables 1 and 2 of \cite{perlmutter}."
+ The results of our analvsis for the CCG Universe are displaved in Fig., The results of our analysis for the GCG Universe are displayed in Fig.
+ 1., $1$ .
+" Iu this figure we show Gs aud 95 coufidence level contours. iu the (o. Q3,)-"," In this figure we show $68$ and $95$ confidence level contours, in the $ (\alpha ,$ $\Omega _{M}^*)$ -plane."
+" We observe that current SNIa data coustrain O3, to the range OL. but do not strougly coustrain the parauicter à in the considered range."," We observe that current SNIa data constrain $\Omega _{M}^*$ to the range $0.15\lesssim \Omega _{M}^*\lesssim 0.4$ , but do not strongly constrain the parameter $\alpha $ in the considered range."
+ Other tests may impose further coustraints., Other tests may impose further constraints.
+ For iustauce. in Ref.," For instance, in Ref."
+ [17] it shown that CAIB alone. imposes Ty=112:0.5 Gyr (10) for the age of the Universe.," \cite{knox} it shown that CMB alone, imposes $T_{0}=14\pm 0.5$ Gyr $1\sigma $ ) for the age of the Universe."
+ If we also assume the HIST. hey Project result. fy=72+ |15].. and that {10 and fy measurements are uncorrelated. we obtain for theproduct ITTo. the following rauee: 0.79<MyTy1.27. at the 20 confidencelevel.," If we also assume the HST Key Project result, $H_{0}=72\pm 8$ \cite{freedman}, and that $H_{0}$ and $T_{0}$ measurements are uncorrelated, we obtain for theproduct $H_{0}T_{0}$, the following range: $0.79<H_{0}\;T_{0}<1.27$, at the $2\sigma $ confidencelevel."
+ Thecentral value occurs at10=1.03., Thecentral value occurs at$H_{0}T_{0}=1.03$.
+ In Fig., In Fig.
+ 1l. we also display the contours IlyTy=0.79 aud HogTy= 1.27.," 1, we also display the contours $H_{0}T_{0}=0.79$ and $H_{0}T_{0}=1.27$ ."
+ As remarked before. we cau see that negative," As remarked before, we can see that negative"
+as it would for an ideal gas of massless constituents.,as it would for an ideal gas of massless constituents.
+ In the simple model of the previous section. such an effect arose due to the bag pressure. and in actual QCD. one can also interpret it in such a fashion [6]..," In the simple model of the previous section, such an effect arose due to the bag pressure, and in actual QCD, one can also interpret it in such a fashion \cite{Asakawa}."
+" It has also been considered in terms of a gradual onset of deconfinement starting from high momenta [16].. and most recently as a possible consequence of coloured ""resonance? states |17].."," It has also been considered in terms of a gradual onset of deconfinement starting from high momenta \cite{interaction}, and most recently as a possible consequence of coloured “resonance” states \cite{Shuryak}."
+ The second point to note is that the thermodynamic observables remain about LO below their Stefan-Doltzmann values (marked “SB” in 4)) even al very. hieh temperatures. where (he interaction measure becomes very sniall.," The second point to note is that the thermodynamic observables remain about 10 below their Stefan-Boltzmann values (marked “SB” in \ref{edens}) ) even at very high temperatures, where the interaction measure becomes very small."
+" such deviations from ideal gas behaviour can be expressed (o a large extent in terms of effective ""thermal masses nig, of quarks and gluons. with img,g(L)T [18]. - [20]."," Such deviations from ideal gas behaviour can be expressed to a large extent in terms of effective `thermal' masses $m_{\rm th}$ of quarks and gluons, with $\m \simeq g(T)~T$ \cite{ERSW} - \cite{Patkos}."
+ Maintaining the next-to-leacding order term in mass in the Stefan-Doltzmann form gives for the pressure P=cTT! |1—-eT[-a gs) and lor the energy density ε-- 3ecT T! y )] =3ecT T! |. where c and e are colour- ancl flavour-dependent positive constants.," Maintaining the next-to-leading order term in mass in the Stefan-Boltzmann form gives for the pressure P = c T^4 ] = c T^4 [1 - a g^2(T)] and for the energy density = 3 c T^4 )^2 - ] =3 c T^4 ], where $c$ and $a$ are colour- and flavour-dependent positive constants."
+ Since {ον1/log1. the deviations of P? and € from the massless Stefan-Bolizmann form vanish as (logT)?. while the interaction measure decreases more rapidly by one power of log.T.," Since $g(T) \sim 1/\log T$, the deviations of $P$ and $\e$ from the massless Stefan-Boltzmann form vanish as $(\log~T)^{-2}$, while the interaction measure decreases more rapidly by one power of $\log~T$."
+ Finally we turn to the value of the transition temperature., Finally we turn to the value of the transition temperature.
+ Since QCD (in the limit of inassless quarks) does not contain anv dimensional parameters. 7;. can only be obtained in physical units by expressing it in ternis of some other known observable which can also be caleulated on the lattice. such as the p-mass. the proton mass. or the string (tension.," Since QCD (in the limit of massless quarks) does not contain any dimensional parameters, $T_c$ can only be obtained in physical units by expressing it in terms of some other known observable which can also be calculated on the lattice, such as the $\rho$ -mass, the proton mass, or the string tension."
+ In the continuum limit. all different. wavs should lead to the same result.," In the continuum limit, all different ways should lead to the same result."
+ Within the present accuracy. (hey define (he uncertainty so lar still inherent in the lattice evaluation of QCD.," Within the present accuracy, they define the uncertainty so far still inherent in the lattice evaluation of QCD."
+" Using the p-mass to fix the scale leads to T;70.15 GeV. while the string tension still allows values as large as 7;20.20 GeV. Very recently. fine structure charmonium calculations (the mass splitting betweenΟν, and 0/75) have been used to fix the dimensional scale. leading to |21]. to T;1902:10 MeV. In any case. energy densities of some | - 2 /Im are needed in order to produce a medium of deconfined quarks and eluons."," Using the $\rho$ -mass to fix the scale leads to $T_c\simeq 0.15$ GeV, while the string tension still allows values as large as $T_c \simeq 0.20$ GeV. Very recently, fine structure charmonium calculations (the mass splitting between, and ) have been used to fix the dimensional scale, leading to \cite{BBC} to $T_c \simeq 190
+\pm 10$ MeV. In any case, energy densities of some 1 - 2 $^3$ are needed in order to produce a medium of deconfined quarks and gluons."
+ In summnmarx. finite temperature lattice QCD at vanishing baryon density shows," In summary, finite temperature lattice QCD at vanishing baryon density shows"
+In the outer 28 '4 bv radius of the Sun. known as the Solar Convection. Zone (SCZ). most of the vertical energy transport is by convection.,"In the outer 28 $\%$ by radius of the Sun, known as the Solar Convection Zone (SCZ), most of the vertical energy transport is by convection."
+ Due to the combination of convection and rotation. the gaseous body exhibits dillerential rotation Le. non-uniform angular velocity.," Due to the combination of convection and rotation, the gaseous body exhibits differential rotation i.e. non-uniform angular velocity."
+ LUcloscismology observations (Scherrer et al 1995. Libbrecht 1989) of the nearly 10 million. acoustic p-modes that leak [rom the interior into the atmosphere. have provided through frequency. splittings. information on the angular velocity cistribution of the solar interior.," Helioseismology observations (Scherrer et al 1995, Libbrecht 1989) of the nearly 10 million acoustic p-modes that leak from the interior into the atmosphere, have provided through frequency splittings, information on the angular velocity distribution of the solar interior."
+ The latest results suggest that the isorotation surfaces in the solar convection zone are cone-like (aligned. racially). in disagreement with most numerical simulations. which tend. to produce cvlindrical contours parallel to the rotation axis.," The latest results suggest that the isorotation surfaces in the solar convection zone are cone-like (aligned radially), in disagreement with most numerical simulations, which tend to produce cylindrical contours parallel to the rotation axis."
+ Specifically. such observations have revealed. that: the angular velocity near the equator first increases then gently decreases with depth. at. mid-latitudes it is almost constant with depth. while at high latituces it increases with depth: on the surface of the Sun. the angular velocity increases from the poles (35 day period) to the equator (25 day period) and meridional circulation. vo. is much. weaker than cilferential rotation (zonal velocity z2km/s.ry2 25m/s)," Specifically, such observations have revealed that: the angular velocity near the equator first increases then gently decreases with depth, at mid-latitudes it is almost constant with depth, while at high latitudes it increases with depth; on the surface of the Sun, the angular velocity increases from the poles (35 day period) to the equator (25 day period) and meridional circulation $v_\theta$, is much weaker than differential rotation (zonal velocity $\approx 2{\rm km/s},
+v_\theta \approx 25 {\rm m/s}$ )."
+ Alodeling the SCZ is a formidable task., Modeling the SCZ is a formidable task.
+ Observations sugeest: that the SCZ is highly stratified. spanning about 19 pressure scale heights in depth. with the Mach. number (square of the ratio of flow velocity to. sound. speed) approaching unity at the top: the Prandtl number (ratio of the timescales of thermal to viscous cdillusion) is extremely small (10. ): the equation of state is complex: and. finally. the motion is highlv turbulent. revealed by the large number," Observations suggest: that the SCZ is highly stratified, spanning about 19 pressure scale heights in depth, with the Mach number (square of the ratio of flow velocity to sound speed) approaching unity at the top; the Prandtl number (ratio of the timescales of thermal to viscous diffusion) is extremely small $10^{-6}$ ); the equation of state is complex; and finally, the motion is highly turbulent, revealed by the large number"
+following the lines O|O4=1.OQconst.,"following the lines $\Omega
++ \Omega_\Lambda = 1 - \Omega_R = \mbox{const}$."
+ Thus our method should be able to constrain well the curvature £j. while the quantity O.—O4 should remain practically unknown. at least when the number .V of galaxies used is not too high.," Thus our method should be able to constrain well the curvature $\Omega_R$, while the quantity $\Omega - \Omega_\Lambda$ should remain practically unknown, at least when the number $N$ of galaxies used is not too high."
+ This approximate invariance may explain the behavior of the quantity (» for finite .., This approximate invariance may explain the behavior of the quantity $\ell_2$ for finite $N$.
+ In fact. if the invariance were exact. f» would be expected to follow a chi-square distribution with (only) one degree of freedom.," In fact, if the invariance were exact, $\ell_2$ would be expected to follow a chi-square distribution with (only) one degree of freedom."
+" This is indeed observed for ANS50000,"," This is indeed observed for $N \lesssim
+50\,000$."
+ Thus in order to draw the confidence regions. one should use the results of Monte Carlo simulations. unless the number .N of galaxies is large.," Thus in order to draw the confidence regions, one should use the results of Monte Carlo simulations, unless the number $N$ of galaxies is large."
+ The contours presented in this paper have been obtained from Monte Carlo simulations., The contours presented in this paper have been obtained from Monte Carlo simulations.
+ The use of the asymptotic limit would have led indeed to confidence regions., The use of the asymptotic limit would have led indeed to confidence regions.
+ In this section we briefly discuss the applicability of our method to observations., In this section we briefly discuss the applicability of our method to observations.
+ As noted earlier. the main difficulty with our method is probably the high number of galaxies (and redshifts) needed to obtain significant constraints on the geometry of the universe.," As noted earlier, the main difficulty with our method is probably the high number of galaxies (and redshifts) needed to obtain significant constraints on the geometry of the universe."
+ In order to show that the examples chosen for our simulations can be considered to be not unrealistic. we first note that the required density of galaxies is between LOO and 500 galaxies per square areminute.," In order to show that the examples chosen for our simulations can be considered to be not unrealistic, we first note that the required density of galaxies is between $400$ and $800$ galaxies per square arcminute."
+ This density ts certainly too high for current instruments., This density is certainly too high for current instruments.
+" In fact. the current achievable density of galaxies (with ""reasonable"" integration times) is of the order of 60 galaxies per square areminute (Cloweαἱ."," In fact, the current achievable density of galaxies (with “reasonable” integration times) is of the order of $60$ galaxies per square arcminute (Clowe."
+ 1998). and this number ts only likely to double with the Advanced Camera on HST.," 1998), and this number is only likely to double with the Advanced Camera on HST."
+ However. current estimates for NGST (Stiavellial.," However, current estimates for NGST (Stiavelli."
+ 1997) give more than 1000 galaxies per square areminute.," 1997) give more than $1\,000$ galaxies per square arcminute."
+ It is also interesting to note that the baseline for the camera currently considered for NGST is about l'. not far from the value used in our simulations (5).," It is also interesting to note that the baseline for the camera currently considered for NGST is about $4\arcmin$, not far from the value used in our simulations $5\arcmin$ )."
+ A second difficulty with the method proposed here is the large number of redshifts required., A second difficulty with the method proposed here is the large number of redshifts required.
+ As suggested in the Introduction. photometric redshifts might be used.," As suggested in the Introduction, photometric redshifts might be used."
+ In this case. however. sizable statistical errors on the measured redshifts are expected.," In this case, however, sizable statistical errors on the measured redshifts are expected."
+ At this stage. for simplicity. such errors have not been taken into account in the simulations. but they obviously should be considered in detail before applications are tried. (," At this stage, for simplicity, such errors have not been taken into account in the simulations, but they obviously should be considered in detail before applications are tried. ("
+"We note also that errors on the measurement of ellipticities may have an impact, especially if small and weak source galaxies are used.)","We note also that errors on the measurement of ellipticities may have an impact, especially if small and weak source galaxies are used.)"
+ Another important ingredient ts the cluster chosen in our simulations., Another important ingredient is the cluster chosen in our simulations.
+ In an Einstein-de Sitter universe with //=(0.5. the cluster would have a mass within à 1Mpe aperture of about 2< M...," In an Einstein-de Sitter universe with $h = 0.5$, the cluster would have a mass within a $1 \, \mbox{Mpc}$ aperture of about $2 \times 10^{15} \, \mbox{M}_\odot$ ."
+" We may compare our cluster with the X-ray cluster MS 105103 located at :,=1,523 (Luppino Kaiser 1997); weak lensing estimates for this cluster give a mass within 1Mpe in the range 1.2«101M.. and 5.5«104 M... depending on the assumed redshift distribution for the source galaxies."," We may compare our cluster with the X-ray cluster MS $1054-03$ located at $z_d = 0.83$ (Luppino Kaiser 1997); weak lensing estimates for this cluster give a mass within $1 \, \mbox{Mpc}$ in the range $1.2 \times 10^{15} \,
+\mbox{M}_\odot$ and $5.5 \times 10^{15} \, \mbox{M}_\odot$ , depending on the assumed redshift distribution for the source galaxies."
+ In conclusion. the proposed method should give interesting results if applied to NGST observations.," In conclusion, the proposed method should give interesting results if applied to NGST observations."
+ Of course. several refinements and improvements would be required for the purpose.," Of course, several refinements and improvements would be required for the purpose."
+ For example. as suggested in the Introduction. the use of several clusters should provide tighter constraints on the cosmological parameters.," For example, as suggested in the Introduction, the use of several clusters should provide tighter constraints on the cosmological parameters."
+ Broadly speaking. the use of Δι similar clusters should make the errors on the cosmological parameters smaller by a factor ν΄να.," Broadly speaking, the use of $N_\mathrm{cl}$ similar clusters should make the errors on the cosmological parameters smaller by a factor $\sqrt{N_\mathrm{cl}}$."
+ However. for this purpose the method should be generalized to be applied to different clusters at different redshifts (actually. we expect that if the redshifts of the clusters are different. more information on the geometry of the universe ts available).," However, for this purpose the method should be generalized to be applied to different clusters at different redshifts (actually, we expect that if the redshifts of the clusters are different, more information on the geometry of the universe is available)."
+ We should also keep in mind that. intuitively. the use of several clusters may introduce additional difficulties. much like the case of a single cluster with pronounced substructure.," We should also keep in mind that, intuitively, the use of several clusters may introduce additional difficulties, much like the case of a single cluster with pronounced substructure."
+ Finally. we should stress that a potential problem for the application of our method could be that of multiple lensing.," Finally, we should stress that a potential problem for the application of our method could be that of multiple lensing."
+ In fact. if weak lensing analysis ts used (in order to," In fact, if weak lensing analysis is used (in order to"
+ion continuity. the ion density perturbation can be written as (\Weilancl 2000):: The electron parallel dynamics vields just the 3oltzmann distribution for the electron number density. and using the quasi-neutrality condition. one then obtains the dispersion equation (o τοΠΠ.,"ion continuity, the ion density perturbation can be written as \citep{w3}: : + The electron parallel dynamics yields just the Boltzmann distribution for the electron number density, and using the quasi-neutrality condition, one then obtains the dispersion equation + _i - + ) + k_y^2 _s^2 =0."
+ Who—whi =Twy. Lt i," Here, $\Omega\equiv \omega/\omega_{*e}$, $\rho_s =c_s/\Omega_i$, $c_s^2= T_{e0}/m_i$, $\eta_b=\omega_{bi}/\omega_{*i}=
+L_n/L_{\sss B}$, and $\omega_{be}=- \tau \omega_{bi}$."
+s seen that. without the magnetic field. inhomogeneity. Iq. (4))," It is seen that, without the magnetic field inhomogeneity, Eq. \ref{e4}) )"
+" vields only the standard. driftmode driven by the density eracicnt. with w~lI/L,."," yields only the standard driftmode driven by the density gradient, with $\omega\sim 1/L_n$."
+ Phe magnetic drift terms o; are responsible for the appearance of the additional plasma mode. while all three gradients together are responsible for the instability.," The magnetic drift terms $\omega_{bi}$ are responsible for the appearance of the additional plasma mode, while all three gradients together are responsible for the instability."
+ As shown by Weiland(2000).. Eq. (4))," As shown by \citet{w3}, Eq. \ref{e4}) )"
+ can be solved analytically vielding the approximate erowth-rate (normalized to ων) Lt is seen that for a;Z7qn there will be an instability., can be solved analytically yielding the approximate growth-rate (normalized to $\omega_{*e}$ ) It is seen that for $\eta_i>\eta_{th}$ there will be an instability.
+ For a fixed 7. the stability. conditions are. completely determined by the three characteristic inhomogencity scale- ἐννι Lo.," For a fixed $\tau$, the stability conditions are completely determined by the three characteristic inhomogeneity scale-lengths $L_n, L_{\sss T}, L_{\sss
+B}$ ."
+ This is demonstrated in Fig., This is demonstrated in Fig.
+" 1 where we give the contour plot of £(L,.Ly.L5)=ipna, Lor Ayps=0.1 and for 7=ο1."," 1 where we give the contour plot of $F(L_n, L_{\sss T}, L_{\sss B})=\eta_i - \eta_{th}$ for $k_y
+\rho_s =0.1$ and for $\tau=T_{e0}/T_{i0}=1$."
+ Lhe positive lines denote the values lor which the eracient-driven instability akes place., The positive lines denote the values for which the gradient-driven instability takes place.
+" In. application to the solar atmosphere with dio=ToLo? E. and assuming £4,=3-107? T. or hydrogen ions we have p,=0.032 m and. hence. the condition Ayp.=0.1 implies the perpendicular wavelength A,=2 nm. The assumption of the nearly perpendicular »erturbed ion motion for such a short perpendicular wave-ength in fact implies a Iute-Llike mode that is very elongated along the magnetic field vector. with a parallel wave-length hat is measured in hundreds of kilometers."," In application to the solar atmosphere with $T_{e0}=T_{i0}=10^6\;$ K, and assuming $B_0=3\cdot 10^{-2}\;$ T, for hydrogen ions we have $\rho_s=0.032\;$ m and, hence, the condition $k_y \rho_s =0.1$ implies the perpendicular wavelength $\lambda_y=2\;$ m. The assumption of the nearly perpendicular perturbed ion motion for such a short perpendicular wave-length in fact implies a flute-like mode that is very elongated along the magnetic field vector, with a parallel wave-length that is measured in hundreds of kilometers."
+ Such strongly elongated. modes kj hy) are easily excited. under aboratory conditions iin a tokamak plasma) in spite of he rather limited scales in the parallel direction.," Such strongly elongated modes $k_\bot\gg
+k_{\|}$ ) are easily excited under laboratory conditions in a tokamak plasma) in spite of the rather limited scales in the parallel direction."
+ In the solar magnetic structures with naturally drastic dillerences in the »rpendieular and. parallel scale-Iengths. their excitation is expected to be even more ellicient.," In the solar magnetic structures with naturally drastic differences in the perpendicular and parallel scale-lengths, their excitation is expected to be even more efficient."
+ The second order dispersion. equation (4)) is solved numerically and some results are. presented in Figs., The second order dispersion equation \ref{e4}) ) is solved numerically and some results are presented in Figs.
+ 2-5., 2-5.
+ In accordance with Fig., In accordance with Fig.
+ 1. in Fig.," 1, in Fig."
+ 2 for ji;=1 we have two real solutions for the wave frequency. one essentially due to the density gracient ancl the other due to the magnetic field eracient.," 2 for $\eta_i=1$ we have two real solutions for the wave frequency, one essentially due to the density gradient and the other due to the magnetic field gradient."
+ The initially positive solution changes the sign for mo=0.4 and then both solutions propagate in the direction of the jon cliamagnetic drift., The initially positive solution changes the sign for $\eta_b\simeq 0.4$ and then both solutions propagate in the direction of the ion diamagnetic drift.
+" Note that for the parameters used. above and for hyp.= 0.1. the normalized. frequency O—1 Hz here implies L,=Ly,~105 m. In Fig."," Note that for the parameters used above and for $k_y \rho_s =0.1$ , the normalized frequency $\Omega\sim 1\;$ Hz here implies $L_n=L_{\sss
+T} \sim 10^3\;$ m. In Fig."
+ 3. the case i;=3 is presented.," 3, the case $\eta_i=3$ is presented."
+" Here. both initial solutions are positive and real for small values of g,."," Here, both initial solutions are positive and real for small values of $\eta_b$."
+" They merge at around i,20.09. vielding a pair of complex-conjugate solutions with the real part becoming negative [or nyz» 0.25."," They merge at around $\eta_b\simeq 0.09$, yielding a pair of complex-conjugate solutions with the real part becoming negative for $\eta_b>0.25$ ."
+ Phe instability. vanishes for gi>2.04 when two negative real solutions appear., The instability vanishes for $\eta_b>2.04$ when two negative real solutions appear.
+" The mode is particularly strongly growing ([2]|2 |.) in the range my,€(0.15. 0.7).Considering g;= 5. in Fig."," The mode is particularly strongly growing $|\gamma| >|\Omega_r|$ ) in the range $\eta_b\in (0.15, \, 0.7)$ .Considering $\eta_i=5$ , in Fig."
+ 4 we present the same mocdebehavior as in Fig., 4 we present the same modebehavior as in Fig.
+ 3., 3.
+" Such a larger jj; value vields a that is larger by about a[actor 2 and the instability range in terms of ap, is widened tom,€(0.05. 3.7)."," Such a larger $\eta_i$ value yields a growth-rate that is larger by about afactor 2 and the instability range in terms of $\eta_b$ is widened to $\eta_b\in (0.05, \, 3.7)$ ."
+to reconcile the results of Πα the shorter trajectories illustrated in our Appendix.,to reconcile the results of \citet{both} with the shorter trajectories illustrated in our Appendix.
+ We also stress. that our results about the radial displacement. of the metals are in agreement with recent chemical models of the Alilky Way aimed to explain the 'hemical gradient along the Galactic disk (7)., We also stress that our results about the radial displacement of the metals are in agreement with recent chemical models of the Milky Way aimed to explain the chemical gradient along the Galactic disk \citep{cemafr07}.
+. In these models the disk is approximated. by several independent rings. 2 kpe wide. without exchange of matter between em.," In these models the disk is approximated by several independent rings, 2 kpc wide, without exchange of matter between them."
+ Because such models are successful in reproducing the observed &radients. one can infer that most of the metals do =100 move racially more than 1 kpe from the place where 10v are created. as found by our simulations.," Because such models are successful in reproducing the observed gradients, one can infer that most of the metals do not move radially more than 1 kpc from the place where they are created, as found by our simulations."
+ We finally make some remarks about the interaction oetween the fountains and the hot gas of the halo., We finally make some remarks about the interaction between the fountains and the hot gas of the halo.
+ The SN LL feedback. infact. is an important. element to understand he energv budget of the halo (e.g.2)..," The SN II feedback, infact, is an important element to understand the energy budget of the halo \citep[e.g.][]{wa05}."
+ Unfortunately. at his stage. a definite conclusion can not be drawn by our models of single fountains.," Unfortunately, at this stage, a definite conclusion can not be drawn by our models of single fountains."
+ Phe amount of thermal energy of the gas of the halo in our Galactic model is Zi107 erg., The amount of thermal energy of the gas of the halo in our Galactic model is $E_{\rm th}\sim 10^{55}$ erg.
+ Although this gas is set initially at rest. some angular momentum is transferred to it by the rotating LSAL of the disk through numerical viscosity.," Although this gas is set initially at rest, some angular momentum is transferred to it by the rotating ISM of the disk through numerical viscosity."
+" While the amount of this angular momentum is small (and the induced. rotation is highlv subsonic). it. introduces some noise at the vertical edges of our computational box. and the thermal energy of the eas reduces of an amount AZ,—351075. about of the total."," While the amount of this angular momentum is small (and the induced rotation is highly subsonic), it introduces some noise at the vertical edges of our computational box, and the thermal energy of the gas reduces of an amount $\Delta E_{\rm th}=3\times 10^{53}$, about of the total."
+ This energy must be compared with £j=5.107EN erg. which is the energy injected in the numerical grid by the SNe LH.," This energy must be compared with $E_{\rm fnt}=5\times 10^{52}$ erg, which is the energy injected in the numerical grid by the SNe II."
+ This energy is ten times lower than the perturbation generated at the boundaries. and any evaluation of the energy transfer from the fountain to the halo is prevented.," This energy is ten times lower than the perturbation generated at the boundaries, and any evaluation of the energy transfer from the fountain to the halo is prevented."
+" This subject will be toroughly investigated in the companion paper on the multiple fountains. when the energy injected by the SNe is ~ LO""? erg and we will be able to give a reasonably accurate evaluation of the fraction of energy. transferred. to the halo gas."," This subject will be toroughly investigated in the companion paper on the multiple fountains, when the energy injected by the SNe is $\sim$ $10^{54}$ erg and we will be able to give a reasonably accurate evaluation of the fraction of energy transferred to the halo gas."
+ We are indebted to an. (anonymous) referee. and to Alex Raga for their extremly profitable. comments and insightful suggestions that have greatly. helped to improve this manuscript., We are indebted to an (anonymous) referee and to Alex Raga for their extremly profitable comments and insightful suggestions that have greatly helped to improve this manuscript.
+ IE.M.CG.D.P. also acknowledges the partial support from erants of the Brazilian Agencies FAPLESP and CNPq. 7.. 2.1)," E.M.G.D.P. also acknowledges the partial support from grants of the Brazilian Agencies FAPESP and CNPq. \citet{breg80}, \ref{subsec:galmod})"
+). Al Al Al , \ref{fig:bal} \ref{fig:bal} \ref{fig:bal} 
+M. ,"M. nuclei of blazar type, these variations are the most dramatic.}"
+"observatioons, most volatile species are found inH5»O-dominated ices.","ons, most volatile species are found in$_2$ O-dominated ices."
+ parameteriizatior of ice desorption., ization of ice desorption.
+"mode inertia 2,=£/3MLB? as a function of the evclic frequeney vo=wf2r for our model of a UMa.","mode inertia $I_{\rm n}=I/3MR^2$ as a function of the cyclic frequency $\nu=\omega/2\pi$ for our model of $\alpha\,$ UMa."
+ Phe choice of normalization is not important here. except that all modes are assumed to have the same surface. amplitude of radial displacement.," The choice of normalization is not important here, except that all modes are assumed to have the same surface amplitude of radial displacement."
+ Phere are two sequences of model results: in the first. sequence (upper plots) we calculated h with 5 obtained. [rom our code: in the second. sequence we suppressed. all nonadiabatic effects. where kr2rq., There are two sequences of model results: in the first sequence (upper plots) we calculated $h$ with $\gamma$ obtained from our code; in the second sequence we suppressed all nonadiabatic effects where $r>r_{\rm f}$.
+ A comparison allows us to assess some of the consequences of the uncertainties of the physics in the convective zone., A comparison allows us to assess some of the consequences of the uncertainties of the physics in the convective zone.
+ Symbols are used to denote the nonradial modes that are most trapped in the acoustic cavity., Symbols are used to denote the nonradial modes that are most trapped in the acoustic cavity.
+" For the (=2 sequence the minima in fy, almost coincide with the radial-mode frequencies. while the minima for 6=1 are located roughly half-way between the /—0 and (=2 minima."," For the $\ell=2$ sequence the minima in $I_{\rm n}$ almost coincide with the radial-mode frequencies, while the minima for $\ell=1$ are located roughly half-way between the $\ell=0$ and $\ell=2$ minima."
+ The positions of these minima resemble the positions of modes in the whole-clise spectra of solar oscillations., The positions of these minima resemble the positions of modes in the whole-disc spectra of solar oscillations.
+ There are many more modes with (=1 and (—2 than are depicted by the svinbols., There are many more modes with $\ell=1$ and $\ell=2$ than are depicted by the symbols.
+ The frequency separation Av between nonracdial modes of consecutive order » is indeed very small., The frequency separation $\Delta\nu$ between nonradial modes of consecutive order $n$ is indeed very small.
+ H0 may be evaluated from theasymptotic formula.. where v is expressed in 112.," It may be evaluated from theasymptotic formula, where $\nu$ is expressed in $\mu$ Hz."
+ Phe numerical constant is specific to the model., The numerical constant is specific to the model.
+ At Ξξ1 and ν=2.syllz the value of Av is about 111. much. less than that found. by ((2000).," At $\ell=1$ and $\nu=2.8\mu{\rm Hz}$ the value of $\Delta\nu$ is about $\,\mu$ Hz, much less than that found by (2000)."
+ There is no substantial dilference between the trapping pattern of the two sequences. except for the differences in the depths of the minima. particularly those of the (=2 modes.," There is no substantial difference between the trapping pattern of the two sequences, except for the differences in the depths of the minima, particularly those of the $\ell=2$ modes."
+ Greater driving in the outer lavers results in deeper münima., Greater driving in the outer layers results in deeper minima.
+ Wo there is net damping in the outer layers. as in the case of solar oscillations. the minima are shallower than in the aciabatic approximation.," If there is net damping in the outer layers, as in the case of solar oscillations, the minima are shallower than in the adiabatic approximation."
+ In Fig., In Fig.
+" 3. we plot the rate of energy. dissipation. 2,= Ms in the asymptotic interior for the same two sequences of modes."," 3, we plot the rate of energy dissipation, $D_{\rm g}\equiv-\omega W_{\rm g}$ , in the asymptotic interior for the same two sequences of modes."
+ In addition. in the upper panel. we show the total energy gain rate.— δρ«M. for racial modes.," In addition, in the upper panel, we show the total energy gain rate, $-D_{\rm p}\equiv-\omega W$, for radial modes."
+ The total riving rate for the nonracdial modes is given approximately wae(DG)LED.71. because significant nonacdiabatic ellects may arise only either near the aand then rey are (independent oor in the deep interior where 10 emode asvmptoties applies Some nonracdial modes with requencies larger then yllz are found to be unstable: if 10 racial modes with v712 plz ave indeed unstable. then rere are also some unstable Iow-degree nonradial modes.," The total driving rate for the nonradial modes is given approximately by $\gamma\simeq(D_{\rm p}(\nu)+D_{\rm g})/2\omega^2I$, because significant nonadiabatic effects may arise only either near the and then they are $\ell$ or in the deep interior where the g-mode asymptotics applies Some nonradial modes with frequencies larger then $\,\mu$ Hz are found to be unstable: if the radial modes with $\nu>12\,\mu$ Hz are indeed unstable, then there are also some unstable low-degree nonradial modes."
+ When (=2 the modes that are most. trapped: are detached. from the remaining modes. except for the one at vex 1112.," When $\ell=2$ the modes that are most trapped are detached from the remaining modes, except for the one at $\nu\simeq10\,\mu$ Hz."
+ Except for this particular mode. all the other modes satisfv d; 1.," Except for this particular mode, all the other modes satisfy $\Phi_{\rm i,f}>1$ ."
+ Thus. the unstable (=2 modes are STU modes. and their growth rates are nearly the same as," Thus, the unstable $\ell=2$ modes are STU modes, and their growth rates are nearly the same as"
+we find iez2 and therefore Dz0.12.,we find $w\approx 2 $ and therefore $\Gamma\approx 0.12$.
+ Εις calculation assumes a Gaussian. photo-z error distribution. which is unlikely to be valid in detail.," This calculation assumes a Gaussian $z$ error distribution, which is unlikely to be valid in detail."
+ Large. tails in the photo-z error distribution would tend to further weaken the observed signal., Large tails in the $z$ error distribution would tend to further weaken the observed signal.
+ The theoretical predictions shown in Figs., The theoretical predictions shown in Figs.
+ 1. and 2 have been multiplied. by this. scale-dependent correction [actor for photometric redshift’ error., \ref{fig:corr_results} and \ref{fig:point_results} have been multiplied by this scale-dependent correction factor for photometric redshift error.
+ Still. the observed signals (especially the pointing angle alignment) are weaker than the theoretical. predictions. even accounting for this contamination.," Still, the observed signals (especially the pointing angle alignment) are weaker than the theoretical predictions, even accounting for this contamination."
+ The correlation alignment is mareinally consistent with the theoretical predictions., The correlation alignment is marginally consistent with the theoretical predictions.
+ For the pointing angle. while the alignments are strongly detected (~ 106). they are considerably weaker than the theoretical predictions.," For the pointing angle, while the alignments are strongly detected $\sim 10\sigma$ ), they are considerably weaker than the theoretical predictions."
+ The ellect of the second. observational error considered here.— cluster centroiding error. is more cillicult to model accurately.," 	 The effect of the second observational error considered here, cluster centroiding error, is more difficult to model accurately."
+ While maxBC'G centroiding errors have been modeled. on mock catalogues (τοι. realistic cluster centroiding errors have only been estimated: for special cluster subsamples such as the very massive ones that have strong X-ray detections (e.g. 2)).," While maxBCG centroiding errors have been modeled on mock catalogues \citep{2007arXiv0709.1159J}, realistic cluster centroiding errors have only been estimated for special cluster subsamples such as the very massive ones that have strong X-ray detections (e.g., \citealt{2009ApJ...697.1358H}) )."
+ Thus. modeling this elfect is bevond the scope of this paper. and we merely state that it should reduce the observed. correlations.," Thus, modeling this effect is beyond the scope of this paper, and we merely state that it should reduce the observed correlations."
+ Third. we estimate the impact of computing the cluster ellipticity and. position angle from only =5 (twpically 56) galaxies.," Third, we estimate the impact of computing the cluster ellipticity and position angle from only $\ge 5$ (typically 5–6) galaxies."
+ In principle. this introduces measurement error in the position angles that is tvpicallv 15 degrees (2).. which will clilute the predicted. signal (it also increases the noise. but we have correctly accounted for this in our estimation of errorbars already).," In principle, this introduces measurement error in the position angles that is typically 15 degrees \citep{2010MNRAS.405.2023N}, which will dilute the predicted signal (it also increases the noise, but we have correctly accounted for this in our estimation of errorbars already)."
+ While the elfect of statistical error in the cluster position angles is in. principle complicatect. we can appeal to simple arguments to roughly estimate its impact.," While the effect of statistical error in the cluster position angles is in principle complicated, we can appeal to simple arguments to roughly estimate its impact."
+" Lt should. result. in the true distribution. of correlation or pointing angles. picos 8,.,,). being convolved with some error distribution. which we assume to be Gaussian."," It should result in the true distribution of correlation or pointing angles, $p(\cos^2{\theta_{c,p}})$ , being convolved with some error distribution, which we assume to be Gaussian."
+ Unfortunately. ? do not report a full distribution of cos?θε. only its mean value.," Unfortunately, \cite{2005ApJ...618....1H} do not report a full distribution of $\cos^2{\theta_{c,p}}$, only its mean value."
+" Thus. to estimate the elfect on our statistic feos""8,μὲν we arbitrarily assume that picos?A.)=A1Bcos? 6..."," Thus, to estimate the effect on our statistic $\langle\cos^2{\theta_{c,p}}\rangle$, we arbitrarily assume that $p(\cos^2{\theta_{c,p}}) = A + B
+ \cos^2{\theta_{c,p}}$ ."
+" Welix the values of Land 2 by imposing two requirements: that the probability distribution be normalised to 1. and that σος”8,μὲ should be consistent with the simulations."," We fix the values of $A$ and $B$ by imposing two requirements: that the probability distribution be normalised to $1$, and that $\langle\cos^2{\theta_{c,p}}\rangle$ should be consistent with the simulations."
+" Then. we convolve that clistribution with a Gaussian distribution with aj...) Gvhere e;=v2, since the former includes two position angles and the latter includes one)."," Then, we convolve that distribution with a Gaussian distribution with $\sigma_{\mathrm{obs},c,p}$ (where $\sigma_{\mathrm{obs},c}
+ = \sqrt{2} \sigma_{\mathrm{obs},p}$ since the former includes two position angles and the latter includes one)."
+ In this simple limit. for small Tobwepe We expect an observed correlation U4 ," In this simple limit, for small $\sigma_{\mathrm{obs},c,p}$, we expect an observed correlation = 0.5 + - ]."
+"This reduces. to. σος65,4,=feos?85. as ω-- approaches zero."," This reduces to $\langle\cos^2\theta\rangle_\mathrm{obs}=\langle\cos^2{\theta}\rangle$ as $\sigma_{\mathrm{obs},c,p}$ approaches zero."
+ Following ?.. if we assume that the position angle errors are typically ~15 deg. then the correlations are reduced by ~15 per cent.," Following \cite{2010MNRAS.405.2023N}, if we assume that the position angle errors are typically $\sim 15$ deg, then the correlations are reduced by $\sim 15$ per cent."
+ Fhis reduction cannot account for the apparent dilference between theory. ancl observation in Fig. 2..," This reduction cannot account for the apparent difference between theory and observation in Fig. \ref{fig:point_results},"
+ which means that reconciling this detection with the theory will require more detailed modeling of observational systematic errors., which means that reconciling this detection with the theory will require more detailed modeling of observational systematic errors.
+ Finally. we consider the impact. of excluding. those clusters that are within lo of being round and therefore have poorly defined. position angles.," Finally, we consider the impact of excluding those clusters that are within $1\sigma$ of being round and therefore have poorly defined position angles."
+ Phere are two options o consider., There are two options to consider.
+ The first is that those clusters truly ave round (tvpicallvy cluster dark matter halos are. triaxial in A’- σον simulations. but could. appear nearly. round. clue to »ojection at certain position angles).," The first is that those clusters truly are round (typically cluster dark matter halos are triaxial in $N$ -body simulations, but could appear nearly round due to projection at certain position angles)."
+ In that case. the worctical predictions for the pointing and correlation angle μαalistics include a contribution of zero from such. round ‘lusters. and our exclusion of them will artificially inflate 1 signal.," In that case, the theoretical predictions for the pointing and correlation angle statistics include a contribution of zero from such round clusters, and our exclusion of them will artificially inflate the signal."
+ “Phe second option is that the clusters are not round. but that they appear so due to noise given that vpically only 56 galaxies are used to define the shape: us option likely represents the majority. of the clusters ju were excluded.," The second option is that the clusters are not round, but that they appear so due to noise given that typically only 5–6 galaxies are used to define the shape; this option likely represents the majority of the clusters that were excluded."
+ In that case. the theoretical predictions include some real alignment signal for these clusters. which cannot be measured in reality since their position angles are too noisy.," In that case, the theoretical predictions include some real alignment signal for these clusters, which cannot be measured in reality since their position angles are too noisy."
+ Lowe were to include them. it would. dilute the measured signal.," If we were to include them, it would dilute the measured signal."
+ EExeluding them is the proper choice in this case. and as lone as those clusters that are excluded are a representative subsample with respect to intrinsic alignment properties. then it should not cause any bias in the signal.," Excluding them is the proper choice in this case, and as long as those clusters that are excluded are a representative subsample with respect to intrinsic alignment properties, then it should not cause any bias in the signal."
+ Nonetheless. future work should include a more careful simulation with such effects. clirecthy incorporated before generating theoretical predictions.," Nonetheless, future work should include a more careful simulation with such effects directly incorporated before generating theoretical predictions."
+ In this paper. we have presented a strong detection. of ealaxy cluster intrinsic alignments to very large scales of 100h tAlpe. representing a tendency of clusters to. point preferentially towards other clusters.," In this paper, we have presented a strong detection of galaxy cluster intrinsic alignments to very large scales of $100h^{-1}$ Mpc, representing a tendency of clusters to point preferentially towards other clusters."
+ Depending on the method used to select clusters ancl assign photometric redshifts. the streneth of the detection (averaged: over 1]«Hιθ0ή IMpe) ranges [rom Ga to 10e.," Depending on the method used to select clusters and assign photometric redshifts, the strength of the detection (averaged over $1<R<100h^{-1}$ Mpc) ranges from $6\sigma$ to $10\sigma$."
+ The alignment of pairs of cluster position angles with each other was only detected at the 2 2.56 level., The alignment of pairs of cluster position angles with each other was only detected at the $2$ $2.5\sigma$ level.
+ This measurement constitutes the first strong detection. of the intrinsic alignments of clusters to 1005. Mpe with a large. statistical cluster sample (6 810° clusters) using a uniform set of whotometric cata.," This measurement constitutes the first strong detection of the intrinsic alignments of clusters to $100h^{-1}$ Mpc with a large, statistical cluster sample $6$ $8\times 10^3$ clusters) using a uniform set of photometric data."
+ The observed correlation angle alignment is consistent. within the large error bars. with the theoretical. AC'DAL wedietion.," The observed correlation angle alignment is consistent, within the large error bars, with the theoretical $\Lambda$ CDM prediction."
+ The pointing angle alignment. while strongly detected to 1005.1 Mpc. is weaker than expected: this is ikelv clue to various systematic observational uncertainties. all of which tend to weaken the observed. signal.," The pointing angle alignment, while strongly detected to $100h^{-1}$ Mpc, is weaker than expected; this is likely due to various systematic observational uncertainties, all of which tend to weaken the observed signal."
+ We ind that photometric redshift error cannot fully account or the dilution of the signal., We find that photometric redshift error cannot fully account for the dilution of the signal.
+ Future work may include an exploration of how the signal varies with the BCC dominance. to test whether merging clusters might be responsible for diluting the alignment signal.," Future work may include an exploration of how the signal varies with the BCG dominance, to test whether merging clusters might be responsible for diluting the alignment signal."
+ Le would also be helpful to carry out this measurement using a sample spanning a broader redshift range. to check for redshift evolution.," It would also be helpful to carry out this measurement using a sample spanning a broader redshift range, to check for redshift evolution."
+ These future investigations can be used to understand. the primary source of the alignments. i.e. whether they are primordial or due to tidal torques from. the large-scale density Ποιά that would [ead to an increase at later times.," These future investigations can be used to understand the primary source of the alignments, i.e. whether they are primordial or due to tidal torques from the large-scale density field that would lead to an increase at later times."
+ A more detailedcomparison of these observations with ACDAL would. require simulations that fully incorporate cluster selection ellects (including False cluster identification. incorrect division of larger clusters into twosmaller clusters. ete).," A more detailedcomparison of these observations with $\Lambda$ CDM would require simulations that fully incorporate cluster selection effects (including false cluster identification, incorrect division of larger clusters into twosmaller clusters, etc.),"
+ centroiding errors. the shape estimation procedure," centroiding errors, the shape estimation procedure"
+"SNR of 3-9 for rj,—510, assuming a distance of 100 Mpc.","SNR of 3-9 for $r_p=5-10$, assuming a distance of 100 Mpc."
+" Scaling the GW emission up to (Mys,Μη)=(2,8)Mg and assuming optimal orientation gives SNR of 8 out to 340 Mpc for rp=7.5."," Scaling the GW emission up to $(M_{\rm NS},M_{\rm BH})=(2,8)M_{\odot}$ and assuming optimal orientation gives SNR of 8 out to 340 Mpc for $r_p=7.5$."
+ Encounters with larger ry could be easier to detect since they will have a number of fly-by GW bursts before the merger., Encounters with larger $r_p$ could be easier to detect since they will have a number of fly-by GW bursts before the merger.
+" Given that some BH-NS capture binaries emit the majority of their GW power at the high frequency end of the LIGO noise curve and could be relatively weak (both in terms of GW emission and extrapolated electromagnetic emission based on disk mass), it may be a worthwhile exercise to revisit the analysis of GRB 070201 (Abbottetal.2008) with burst-templates adapted to capture driven BH-NS encounters."," Given that some BH-NS capture binaries emit the majority of their GW power at the high frequency end of the LIGO noise curve and could be relatively weak (both in terms of GW emission and extrapolated electromagnetic emission based on disk mass), it may be a worthwhile exercise to revisit the analysis of GRB 070201 \citep{Abbott:2007rh}
+ with burst-templates adapted to capture driven BH-NS encounters."
+" We thank Adam Burrows, John Friedmann, Roman Gold, Benjamin Lackey, and Richard O’Shaughnessy for useful conversations."," We thank Adam Burrows, John Friedmann, Roman Gold, Benjamin Lackey, and Richard O'Shaughnessy for useful conversations."
+" This research was supported by the NSF through TeraGrid resources provided by NICS under grant TG-PHY100053, the Bradley Program fellowship the NSF Graduate Research Fellowship under grant (BCS),DGE-0646086 (WE), NSF grants PHY-0745779 (FP) and PHY-1001515 (BCS), and the Alfred P. Sloan Foundation "," This research was supported by the NSF through TeraGrid resources provided by NICS under grant TG-PHY100053, the Bradley Program fellowship (BCS), the NSF Graduate Research Fellowship under grant DGE-0646086 (WE), NSF grants PHY-0745779 (FP) and PHY-1001515 (BCS), and the Alfred P. Sloan Foundation (FP)."
+Simulations were also run on the Woodhencluster at (FP).Princeton University., Simulations were also run on the cluster at Princeton University.
+relative to the Sun. in the LAW? filter than in bolometric Mux.,"relative to the Sun, in the LW2 filter than in bolometric flux."
+ In the brown dwarf models. more power per solar mass is produced by low mass than by higher mass young brown εναν». reflecting the fact that their mass to light ratio changes by [ess than their mass ratio.," In the brown dwarf models, more power per solar mass is produced by low mass than by higher mass young brown dwarfs, reflecting the fact that their mass to light ratio changes by less than their mass ratio."
+ Phe reverse is true for older brown dwarfs., The reverse is true for older brown dwarfs.
+ The integrated. power per unit solar mass for the various brown dwarf mass functions is almost independant of the adopted mass function., The integrated power per unit solar mass for the various brown dwarf mass functions is almost independant of the adopted mass function.
+ This rellects the result. above: for voung brown ciwarfs. the emissivity is dominated by the lowest masses. for older brown chwarls it is dominated by the highest masses.," This reflects the result above: for young brown dwarfs, the emissivity is dominated by the lowest masses, for older brown dwarfs it is dominated by the highest masses."
+ Thus. one has minimal sensitivity to the mass spectrum in the range toΕ," Thus, one has minimal sensitivity to the mass spectrum in the range to."
+ΟΟ Finally. old. white cars are not well suited to direct detection by ISO.," Finally, old white dwarfs are not well suited to direct detection by ISO."
+ Thev are too hot. and have too small a surface area for their mass. for easy detection.," They are too hot, and have too small a surface area for their mass, for easy detection."
+" Given these caleulated central Luminosity densities per unit mass. together with the parameterised central mass densities per unit volume. representing the rotation curve fits in figure οον, it is straightforward. to compare the mocdels with our observations."," Given these calculated central luminosity densities per unit mass, together with the parameterised central mass densities per unit volume, representing the rotation curve fits in figure \ref{rotcurves}, it is straightforward to compare the models with our observations."
+ The graphical comparison of these results with the observations is given in figure ??.., The graphical comparison of these results with the observations is given in figure \ref{results}.
+ The solid lines indicate the hvdrogen-burning star models (ALLMA). dotted lines represent the brown dwarf models. and the dashed line represents the white dwarf model.," The solid lines indicate the hydrogen-burning star models (M1–M4), dotted lines represent the brown dwarf models, and the dashed line represents the white dwarf model."
+ Phe bold lines labelled with the galaxy name represents the results from the ISO observations. with in each case the 95 percent upper limits indicated by the lines labelled with the galaxy. name.," The bold lines labelled with the galaxy name represents the results from the ISO observations, with in each case the 95 percent upper limits indicated by the lines labelled with the galaxy name."
+ Models involving significant contributions to dark halos from hydrogen burning stars with a plausible mass function model AIS are clearly excluded in both passbands., Models involving significant contributions to dark halos from hydrogen burning stars with a plausible mass function – model M5 – are clearly excluded in both passbands.
+ Even the extreme mocel. of a d-function mass function of stars at the hyelrogen burning limit is directly excluded.," Even the extreme model, of a $\delta$ -function mass function of stars at the hydrogen burning limit is directly excluded."
+ Similarly. models involving a substantial population of voung brown cwarfs. with masses near-to. but below. the hyvdrogen-burning limit. are also directly excluded.," Similarly, models involving a substantial population of young brown dwarfs, with masses near-to, but below, the hydrogen-burning limit, are also directly excluded."
+ Interestingly. improvements in sensitivity in the longer wavelength.LAWS data by an order or magnitude — which might be feasible with the present observations. given improved. calibrations and cata reduction algorithms. Phough may well require SIRTE. can directly test any model involving brown cwarfs as a significant contribution to ealactic dark halos.," Interestingly, improvements in sensitivity in the longer-wavelength LW3 data by an order or magnitude – which might be feasible with the present observations, given improved calibrations and data reduction algorithms, Though may well require SIRTF, can directly test any model involving brown dwarfs as a significant contribution to galactic dark halos."
+ We illustrate the present situation below. in figure ??..," We illustrate the present situation below, in figure \ref{bdrange}."
+ Phis shows. with better mass resolution. the emissivity as a function of mass of brown chwarls of a variety of ages. as they would be observed through the ISO LAW2 and LAWS filters.," This shows, with better mass resolution, the emissivity as a function of mass of brown dwarfs of a variety of ages, as they would be observed through the ISO LW2 and LW3 filters."
+ The sensitivity of these ISO observations allows us to limit the baryonic possibilities for halo dark matter. and to exclude hyelrogen-burning stars from consideration.," 	 The sensitivity of these ISO observations allows us to limit the baryonic possibilities for halo dark matter, and to exclude hydrogen-burning stars from consideration."
+ The ISO observations are more sensitive and. have significantly better spatial resolution than was previously possible with observations at these wavelengths., The ISO observations are more sensitive and have significantly better spatial resolution than was previously possible with observations at these wavelengths.
+where the density p aud the sound speed ce; have beeu normalized to the typical values fouud in our simulations and shown in Figure 1.,where the density $\rho$ and the sound speed $c_s$ have been normalized to the typical values found in our simulations and shown in Figure 1.
+ For the plivsical conditions iu our four simmlatious. the Bondi aceretion rate is —10!107 Myr. anch bueer than the Eddington 3.10PCAZpu/M.fnM ifyr (CIKrolik 1999). where 1 is the radiative efficiency. to be discussed further below.," For the physical conditions in our four simulations, the Bondi accretion rate is $\sim 10^{-4}-10^{-3} ~\hbox{M}_{\odot}/\hbox{yr}$ , much larger than the Eddington $~~$ $\sim3\times10^{-9}\,(M_{BH}/M_{\odot})/\eta\,\,\hbox{M}_{\odot}$ /yr (Krolik 1999), where $\eta$ is the radiative efficiency, to be discussed further below."
+ The black hole will therefore accrete gas at the Eddingtou rate., The black hole will therefore accrete gas at the Eddington rate.
+ Under these circumstances. the black hole mass will erow on a Salpeter time scale where Lp~LIS«10ALE/AL. erg/s is the Eddington Lhuuinosity and a value of the mass per electron gio1.15 has been adopted (see c.g. Krolik 1999).," Under these circumstances, the black hole mass will grow on a Salpeter time scale where $L_E \sim 1.45 \times 10^{38} M_{BH}/M_{\odot}$ erg/s is the Eddington luminosity and a value of the mass per electron $\mu_e\sim
+1.15$ has been adopted (see e.g. Krolik 1999)."
+ The value of jj is uncertain and depends on the properties of the accretiou disk aud on the spin of the accretiug black hole., The value of $\eta$ is uncertain and depends on the properties of the accretion disk and on the spin of the accreting black hole.
+ Here we adopt yo~0.05—0.1. a range conunonlv adopted in studies of black hole erowth in galactic uucleij.," Here we adopt $\eta\sim
+0.05-0.1$, a range commonly adopted in studies of black hole growth in galactic nuclei."
+ We uote that values of j larecr than 0.1. appropriate for rapidly rotating black holes. would reduce the erowth of the IAIBIT. while much sinaller values (4j«1l ) typical of thick accretion disks would significantly increase the IMDIT erowth rate.," We note that values of $\eta$ larger than 0.1, appropriate for rapidly rotating black holes, would reduce the growth of the IMBH, while much smaller values $\eta \ll 1$ ) typical of thick accretion disks would significantly increase the IMBH growth rate."
+ For accretion at the Eddingtou rate with gj~0.05.0.1. Ty~200 My.," For accretion at the Eddington rate with $\eta\sim 0.05-0.1$, $T_S\sim 20-40$ Myr."
+ The duration of the SG star formation episode in the models of D'Ercole et al. (, The duration of the SG star formation episode in the models of D'Ercole et al. (
+2008) is AT©109 My.,2008) is $\Delta T \approx 10^2$ Myr.
+ This time scale is constrained by spectroscopic observations showing little variation in the total CNO of FC aud SC stars. nuplviug that only ACB progenitors with MomLoSAL. must have coutributed eas to the SC formation process.," This time scale is constrained by spectroscopic observations showing little variation in the total CNO of FG and SG stars, implying that only AGB progenitors with $M>4-5 M_{\odot}$ must have contributed gas to the SG formation process."
+ Thus. if the seed black hole accretes at close to the Eddington rate during this time. its mass will grow by a factor of ~10100. leacling. bx the eud of the SC. formation phase. to a final black hole mass of Alpi~teTohAL...," Thus, if the seed black hole accretes at close to the Eddington rate during this time, its mass will grow by a factor of $\sim 10-100$ , leading, by the end of the SG formation phase, to a final black hole mass of $M_{BH} \sim 10^3-10^4~M_{\odot}$."
+ Several aspects of the scenario just prescuted merit urther discussion aud exploration., $~~$ Several aspects of the scenario just presented merit further discussion and exploration.
+ We consider first the effect on the cooling flow of the enerev radiated by the accreting black hole., We consider first the effect on the cooling flow of the energy radiated by the accreting black hole.
+ To address lis issue we follow the analysis of black hole outflow oeseuted by Wine (2003. 2005). m which it is shown hat the wind produced by the black hole sweeps up he surrounding gas iuto an expanding shell.," To address this issue we follow the analysis of black hole outflow presented by King (2003, 2005), in which it is shown that the wind produced by the black hole sweeps up the surrounding gas into an expanding shell."
+ Because of effective euergv losses due to inverse Compton scattering. he shell expands in the momeutuni-conservius regine.," Because of effective energy losses due to inverse Compton scattering, the shell expands in the momentum-conserving regime."
+ Under these conditions (sce ine 2003. 2005). in order or the momentum flux from a black hole accreting at the Eddington rate to overcome eravity aud expel the shell roni the cluster. the black hole mass iust be larger than Iere. & is the electron scattering opacity. fy is the fraction of the total cluster mass in the form of gas. and it has been assmmed that the eas is enmibedded imu an isothermal svstem with velocity cliispersion 9.," Under these conditions (see King 2003, 2005), in order for the momentum flux from a black hole accreting at the Eddington rate to overcome gravity and expel the shell from the cluster, the black hole mass must be larger than Here, $\kappa$ is the electron scattering opacity, $f_g$ is the fraction of the total cluster mass in the form of gas, and it has been assumed that the gas is embedded in an isothermal system with velocity dispersion $\sigma$."
+" Asstuning σz20kms (the smallest value for which the AGB ejecta cau be retained) and f,z:0.05 (D'Excole et al.", Assuming $\sigma \gtorder 20~\hbox{km/s}$ (the smallest value for which the AGB ejecta can be retained) and $f_g\approx0.05$ (D'Ercole et al.
+ 2008). we fud Ap;Z101A... uch larger than the assumed nass of our initial seed black hole.," 2008), we find $M_{BH,crit} \gtorder 10^4 M_{\odot}$, much larger than the assumed mass of our initial seed black hole."
+ Thus feedback from the accreting black hole isnot sufficient to alter significantly the elobal dvuaiuices of the cooling flow or the SC star formation process., Thus feedback from the accreting black hole is sufficient to alter significantly the global dynamics of the cooling flow or the SG star formation process.
+ On the other haud. the effects of feedback may sienificautly alter the local gas dynamics in the vicinity of the black hole.," On the other hand, the effects of feedback may significantly alter the local gas dynamics in the vicinity of the black hole."
+ For example. recent detailed 2-D hvcdvodvuamical models of accretion outo a LOO AL... seed black hole in a dense protogalactic cloud (Milosavljevic et al.," For example, recent detailed 2-D hydrodynamical models of accretion onto a 100 $M_{\odot}$ seed black hole in a dense protogalactic cloud (Milosavljevic et al."
+ 2009) show iutermuttent accretion. reducing the uet accretion rate to about 1/3 of Eddineton.," 2009) show intermittent accretion, reducing the net accretion rate to about $1/3$ of Eddington."
+ Such a reduction in the mean accretion rate in our case would reduce the erowth of the seed black hole to a factor 25 for the range of jj values adopted in refsec:bh.., Such a reduction in the mean accretion rate in our case would reduce the growth of the seed black hole to a factor $\sim2-5$ for the range of $\eta$ values adopted in \\ref{sec:bh}.
+ We note here that during the accretion phase the black hole will be very Iuniuous aud might be observed as a (possibly intermittent) ultraluniuous X-ray source (CLA)., We note here that during the accretion phase the black hole will be very luminous and might be observed as a (possibly intermittent) ultraluminous X-ray source (ULX).
+ We caution however that. before linking this scenario to the ULXs observed in nearby. voung massive clusters (see e.g. Portegies Zwart et al.," We caution however that, before linking this scenario to the ULXs observed in nearby young massive clusters (see e.g. Portegies Zwart et al."
+ 2010). one must first verity whether these clusters (1) meet the concitions to form a seed black hole. (2) axe sufficieutlv massive aud (3) have the structural properties required for SG. star formation.," 2010), one must first verify whether these clusters (1) meet the conditions to form a seed black hole, (2) are sufficiently massive and (3) have the structural properties required for SG star formation."
+ A second important issue is the possiblerelation between the IMDITI amass and the curent structural properties of the pareut cluster., A second important issue is the possiblerelation between the IMBH mass and the current structural properties of the parent cluster.
+assume that the gas has a density of 0.020m* and velocity/sound speed of I0kms..,assume that the gas has a density of $0.02 \pcmcu$ and velocity/sound speed of $10^3 \kmps$.
+. With a typical bolometric correction of 10 in the 0.5—7 keV band for sources at low Eddington rate (Vasudevan Fabian 2007). most of our sources have X-ray luminosities close to that expected from Bondi accretion (Fig. 8).," With a typical bolometric correction of 10 in the 0.5–7 keV band for sources at low Eddington rate (Vasudevan Fabian 2007), most of our sources have X-ray luminosities close to that expected from Bondi accretion (Fig. \ref{fig:luminosity}) ),"
+ especially when variations in velocity ete., especially when variations in velocity etc.
+ are considered., are considered.
+ Where there are mini-haloes. and it is plausible that they all have some form of mini-halo due to central stellar mass loss. then the expected," Where there are mini-haloes, and it is plausible that they all have some form of mini-halo due to central stellar mass loss, then the expected"
+multiplying the amplitudes by f). where Var(f) is total variance of the data f.,"multiplying the amplitudes by $f$ ), where $f$ ) is total variance of the data $f$."
+ This ensures the powers follow à X5 distribution when the lighteurve is dominated. by (measurement) white noise eeach point is independent and. randomly taken from. the initial lux distribution)., This ensures the powers follow a $\chi^2_2$ distribution when the lightcurve is dominated by (measurement) white noise each point is independent and randomly taken from the initial flux distribution).
+ The optical power spectra typically showed power at low frequencies ancl white noise at high frequencies., The optical power spectra typically showed power at low frequencies and white noise at high frequencies.
+ In these cases we divided by the average power at the frequencies where measurement. noise dominates to recover the AS statistics., In these cases we divided by the average power at the frequencies where measurement noise dominates to recover the $\chi^2_2$ statistics.
+ The Fourier zunplitudes give the fractional variance per frequency. bin and. after subtraction of the constant white noise component. we plot the power [requenev in (dimensionless) units of rns to highlight the timescales with significant variability ancl ease comparison with the N-rav PDS (Fig.," The Fourier amplitudes give the fractional variance per frequency bin and, after subtraction of the constant white noise component, we plot the power $\times$ frequency in (dimensionless) units of $^2$ to highlight the timescales with significant variability and ease comparison with the X-ray PDS (Fig."
+ 3)., 3).
+ The low duty evele of the data acquisition. scheme introduces large amounts of spurious power in a Fourier ranslorm covering several frames., The low duty cycle of the data acquisition scheme introduces large amounts of spurious power in a Fourier transform covering several frames.
+ To circumvent this »oblem. we separated the calculation of the power spectrum into two parts so that the transformed lighteurves had duty eveles close to10074.," To circumvent this problem, we separated the calculation of the power spectrum into two parts so that the transformed lightcurves had duty cycles close to."
+.. At high frequencies. we computed a rower spectrum of the images in cach frame separately. then averaged the power spectra of all the frames.," At high frequencies, we computed a power spectrum of the images in each frame separately, then averaged the power spectra of all the frames."
+ Since there are »etween 46 to 48 images (photometry points) in cach [rame with a time resolution of 5 ms (white light) to 100 ms )). he resulting frequency range is about 5.100 Lz (white light) or 0.24.8 Hz (Lo). depending upon the observational setup (sce Tab.," Since there are between 46 to 48 images (photometry points) in each frame with a time resolution of 5 ms (white light) to 100 ms ), the resulting frequency range is about 5–100 Hz (white light) or 0.2–4.8 Hz ), depending upon the observational setup (see Tab."
+ 2)., 2).
+ At lower frequencies we took the average Lux of each frame (the sum of all images in a frame) and computed he power spectrum of the resulting lighteurve., At lower frequencies we took the average flux of each frame (the sum of all images in a frame) and computed the power spectrum of the resulting lightcurve.
+ The resulting requeney range is about 105 0.05 Lz (see Tab., The resulting frequency range is about $10^{-3}$ –0.05 Hz (see Tab.
+ 2)., 2).
+ The combined power spectrum is free of power leakage at the expense of a loss in frequency resolution and coverage., The combined power spectrum is free of power leakage at the expense of a loss in frequency resolution and coverage.
+ A uigh frequencies. the averaging of the individual spectra diminish the measured power from a highly coherent signal. cause the contributions from each frame are not summe in-phase.," At high frequencies, the averaging of the individual spectra diminish the measured power from a highly coherent signal, because the contributions from each frame are not summed in-phase."
+ Significant power is present in all of the datasets a very low frequencies (below 310 LLz in white light anc 19 Lz in Lla)) except during the short August 3 oobservation., Significant power is present in all of the datasets at very low frequencies (below $3\times 10^{-3}$ Hz in white light and $10^{-3}$ Hz in ) except during the short August 3 observation.
+ The total rims., The total r.m.s.
+ of the Liekcring is estimatec by integrating over the frequencies where power is detecte above the 3e single trial level., of the flickering is estimated by integrating over the frequencies where power is detected above the $\sigma$ single trial level.
+ We obtain values. of rans. (, We obtain values of r.m.s. (
+see ‘Tab.,see Tab.
+ 2) which are consistent with the standard deviation of the 32 s binned lighteurve., 2) which are consistent with the standard deviation of the 32 s binned lightcurve.
+ Phe optica variability is about an order of magnitude smaller than tha seen in the X-ray lighteurves on the same frequency range., The optical variability is about an order of magnitude smaller than that seen in the X-ray lightcurves on the same frequency range.
+ At higher frequencies. statistically significant power is detected. in the white light observation of Aug. 3. when XN-22 was on the horizontal branch with a veh X-rav flux.," At higher frequencies, statistically significant power is detected in the white light observation of Aug. 3, when 2 was on the horizontal branch with a high X-ray flux."
+ The total r.nis., The total r.m.s.
+ is about (0.02OTF Liz)., is about (0.02--0.07 Hz).
+ Unfortunatelv. this is the only detection. of optical variability at high frequency in our limited. dataset.," Unfortunately, this is the only detection of optical variability at high frequency in our limited dataset."
+ Independent confirmation would be desirable., Independent confirmation would be desirable.
+" Phe very low requency power in both optical and. X-rays is reasonably approximated by D,7&? UD.cp in Fig.", The very low frequency power in both optical and X-rays is reasonably approximated by $P_\nu \sim \nu^{-2}$ $\nu P_\nu \sim \nu^{-1}$ in Fig.
+ 3)., 3).
+ We obtained an upper limit to the lickering power in the cases where none was detected: by calculating the total rons., We obtained an upper limit to the flickering power in the cases where none was detected by calculating the total r.m.s.
+ of power with this particular shape that would have vielded a 3m detection. (tvpically a ewἄν που Tab.," of power with this particular shape that would have yielded a $\sigma$ detection (typically a few, see Tab."
+ 2)., 2).
+ We also computed 3o upper limits to the ros., We also computed $\sigma$ upper limits to the r.m.s.
+ of a sinusoidal signal., of a sinusoidal signal.
+ Both of these upper limits take into account the number of trials in the frequency range., Both of these upper limits take into account the number of trials in the frequency range.
+ ALL of the results may be found in Tab., All of the results may be found in Tab.
+ 2., 2.
+ We cross-correlated each of the nine simultaneous X-ray. and optical datasets (Lab., We cross-correlated each of the nine simultaneous X-ray and optical datasets (Tab.
+ 1)., 1).
+ The X-ray data. which have the highest time resolution. were corrected to the local Palomar time and rebinned in an identical way to the optical data set. vielding two continuous lighteurves with the same resolution (Af= 5. 10 or 100119).," The X-ray data, which have the highest time resolution, were corrected to the local Palomar time and rebinned in an identical way to the optical data set, yielding two continuous lightcurves with the same resolution $\Delta t=$ 5, 10 or 100ms)."
+ The cross-correlation function at a lag AM is delined as (e.g. Exdelson&νο]1988)): where o is the optical lighteurve. w the X-ray. the bar symbolises the average of the dataset and e are the standard deviations.," The cross-correlation function at a lag $k\Delta t$ is defined as (e.g. \citealt{krolik}) ): where $o$ is the optical lightcurve, $x$ the X-ray, the bar symbolises the average of the dataset and $\sigma$ are the standard deviations."
+ C'(AzM) is an average over the AL pairs of points which are AM apart., $C(k\Delta t)$ is an average over the $M$ pairs of points which are $k\Delta t$ apart.
+ νο variance of the AL averaged discrete correlations gives an estimate of the error on C'., The variance of the $M$ averaged discrete correlations gives an estimate of the error on $C$.
+ The above is efficiently computed in Fourier space for continuous lichteurves with the same 2M: where E ijs the (fast) Fourier transform. « ijs the conjugate and the v are window functions (a timeseries with value 1: when data are taken. 0 elsewhere).," The above is efficiently computed in Fourier space for continuous lightcurves with the same $\Delta t$: where $\rm F$ is the (fast) Fourier transform, $\star$ is the conjugate and the $w$ are window functions (a timeseries with value 1 when data are taken, 0 elsewhere)."
+ There were no obvious trends on timescales 210 min where the datasets are not long enough to clearly establish correlations., There were no obvious trends on timescales $\gsim$ 10 min where the datasets are not long enough to clearly establish correlations.
+ For the ceross-correlations below 10 min. we subtracted piecewise linear fits in 600 s increments to the datasets to smooth out anv underlying long timescale variation.," For the cross-correlations below 10 min, we subtracted piecewise linear fits in 600 s increments to the datasets to smooth out any underlying long timescale variation."
+ The X-ray and optical data did not show any significant. correlations down to 5 ms (100 ms)., The X-ray and optical data did not show any significant correlations down to 5 ms (100 ms).
+ Cross-correlating with cdillerent energv. selected N-rav. lighteurves (2.05.7 keV. 5.79.8 keV. 9.8 keV and above) gave identical results.," Cross-correlating with different energy selected X-ray lightcurves (2.0–5.7 keV, 5.7–9.8 keV, 9.8 keV and above) gave identical results."
+ Our dataset is consistent with previous studies reporting variations in the optical continuum of N-22 of ~0.5 mag in davs. —0.1 mag in hours and ~0.05 mag in minutes (Ixristianetal.1967:Ixruszewski1974:WilvachlovBeskinctal. 1979)).," Our dataset is consistent with previous studies reporting variations in the optical continuum of 2 of $\sim$ 0.5 mag in days, $\sim$ 0.1 mag in hours and $\sim$ 0.05 mag in minutes \citealt{kristian,kru74,kilyachkov,beskin}) )."
+ Orbital ellipsoidal modulation aside. there are few. well-established features in the optical variability of X-22.," Orbital ellipsoidal modulation aside, there are few well-established features in the optical variability of 2."
+ There is a trend for the continuum to become bluer in C-D as N-22 becomes brighter (1- Vl staving ~ constant: Lvutvi&Sunvaey1976:Daskoet 1976)).," There is a trend for the continuum to become bluer in $U$ $B$ as 2 becomes brighter $B$ $V$ staying $\sim$ constant; \citealt{lyutyi,basko}) )."
+ Varving lluxes and EWILM on timescales = 1 hr were also noticed early on in Ho. L2. A4686 and the Bowen blend. Johnson&Colson1969:Cowleyctal. 1979)).," Varying fluxes and FWHM on timescales $\gsim$ 1 hr were also noticed early on in $\alpha$, $\beta$, $\lambda$ 4686 and the Bowen blend \citealt{johnson,cowley}) )."
+ The evolution from the horizontal to the Laring branches in Cve N-2 like sources has been proposed to reflect an increasing mass accretion rate in the disc (see discussion in Llomanetal. 2002))., The evolution from the horizontal to the flaring branches in Cyg X-2 like sources has been proposed to reflect an increasing mass accretion rate in the disc (see discussion in \citealt{homan}) ).
+ This could be due to Ductuations in the mass transfer rate from the donor star or to Lluctuations in the angular momentum transport. processes. presumably on the disc viscous timescale.," This could be due to fluctuations in the mass transfer rate from the donor star or to fluctuations in the angular momentum transport processes, presumably on the disc viscous timescale."
+ Interestingly.," Interestingly,"
+on redshift of stellar mass. luminosity. colour. SER. SSER and metallicity of the observed. host galaxies compiled. by Savaglioetal.(2009). over the redshift range 0<2«3.,"on redshift of stellar mass, luminosity, colour, SFR, SSFR and metallicity of the observed host galaxies compiled by \citet{Sav09} over the redshift range $0 < z < 3$."
+ Our main findings are: Our results are. mainly a consequence of the joint requirements to have high star formation activity to ensure observability and to reproduce the current. distribution of stellar masses of observed. host. galaxies., Our main findings are: Our results are mainly a consequence of the joint requirements to have high star formation activity to ensure observability and to reproduce the current distribution of stellar masses of observed host galaxies.
+ However. within the current constrains. galaxies with high masses would be too metal-rich. to. produce LORBs ancl low mass systenis would have low probability of being observed.," However, within the current constrains, galaxies with high masses would be too metal-rich to produce LGRBs and low mass systems would have low probability of being observed."
+ LE the observed sample were modified by the incorporation of other galaxics. [or example clusty hosts. then our model would need to be reacljustecl accordinglv.," If the observed sample were modified by the incorporation of other galaxies, for example dusty hosts, then our model would need to be readjusted accordingly."
+ We thank M.IZ. De Rossi and Gerard. Lemson for helping us to manage theSiulalHion. and Sandra Savaglio for her useful comments and. suggestions.," We thank M.E. De Rossi and Gerard Lemson for helping us to manage the, and Sandra Savaglio for her useful comments and suggestions."
+ This work was partially supported by grants PIC'T 2005-32342. ιο 2006-245 Max Planck. and PIC’P 2006-2015 from Argentine ANPCVT.," This work was partially supported by grants PICT 2005-32342, PICT 2006-245 Max Planck, and PICT 2006-2015 from Argentine ANPCyT."
+will be confused (SNR. supersolt sources. LALXBs). ancl only a very rough separation on the basis of hard. color will be possible.,"will be confused (SNR, supersoft sources, LMXBs), and only a very rough separation on the basis of hard color will be possible."
+ The soft sources identified in Section 4. as supernova remnants have considerably less scalter in their luminosities than do sources with //1>-0.6 — Lie. a higher fraction of the sources in the LMXD. part of the diagram have luminosities >10° erg !., The soft sources identified in Section \ref{sec:id} as supernova remnants have considerably less scatter in their luminosities than do sources with $H1>$ -0.6 – i.e. a higher fraction of the sources in the LMXB part of the diagram have luminosities $> 10^{37}$ erg $^{-1}$.
+ This is demonstrated in Figure 7.., This is demonstrated in Figure \ref{fig:colors_lumin}.
+ The left panel shows the soft X-ray color plotted as a function of luminosity [or MIOI ancl M32., The left panel shows the soft X-ray color plotted as a function of luminosity for M101 and M83.
+" This difference in the distribution of luminosities in LMXD and ""soft"" sources is significant: a AS test gives the probability that they are drawn from the same distribution as 6xLO+.", This difference in the distribution of luminosities in LMXB and “soft” sources is significant; a KS test gives the probability that they are drawn from the same distribution as $6\times10^{-4}$.
+ The larger scatter in the luminosities of sources with Hl —0.5is naturally explained if many of these objects are accreting binaries., The larger scatter in the luminosities of sources with $H1>-0.5$ is naturally explained if many of these objects are accreting binaries.
+" Binaries can reach much higher huminosities (especially in a flare state) than is tvpically observed in evolved SNR. ( 109-107"" erg 4).", Binaries can reach much higher luminosities (especially in a flare state) than is typically observed in evolved SNR ( $10^{36}$ $10^{37}$ erg $^{-1}$ ).
+ The X-rav. luminosities of the soft sources are (vpical of brighter SNR: there are certainly other soft sources below our detection threshold., The X-ray luminosities of the soft sources are typical of brighter SNR; there are certainly other soft sources below our detection threshold.
+ Several of the brightest soft sources have very extreme colors (IL1—-1) and have colors and luminosities characteristic of supersoft sources (Ixaliabka.Pietsch.&Hasinger 1994).., Several of the brightest soft sources have very extreme colors (H1=-1) and have colors and luminosities characteristic of supersoft sources \citep{kah94}. .
+ These are probably accretion powered., These are probably accretion powered.
+ The right panel of Figure 7 shows the Iuminosity plot for X-ray hare color., The right panel of Figure \ref{fig:colors_lumin} shows the luminosity plot for X-ray hard color.
+ This plot suggests that sources with the hardest colors (47/2> 0.2. IIMXD candidates) have sinaller scatter than less extreme sources.," This plot suggests that sources with the hardest colors $H2\ge$ 0.2, HMXB candidates) have smaller scatter than less extreme sources."
+ A [x-9 test shows that the. Iuminositv. distributions of the hardest sources is significantlv different. from those in the LMXD part of the diagram (the probability that thev are drawn [rom the same distribution is 5x10°., A K-S test shows that the luminosity distributions of the hardest sources is significantly different from those in the LMXB part of the diagram (the probability that they are drawn from the same distribution is $5\times10^{-5}$.
+ The huninosity function of LMXD sources in the Alilky Way extends to higher luminosities than the IIMXD huminosityv function (Grimm.Gillanov.&Sunvaev2001).. consistent with what is observed here.," The luminosity function of LMXB sources in the Milky Way extends to higher luminosities than the HMXB luminosity function \citep{grimm01}, consistent with what is observed here."
+ We note. however. (hat extremely Iuminous sources tentatively associated with IAINBs are seen in starburst galaxies (Zezas&Fabbiano2002:οἱal2001:Prestwich2001).," We note, however, that extremely luminous sources tentatively associated with HMXBs are seen in starburst galaxies \citep{zez02,fab01,prestwich01}."
+. If the soft sources discussed in Section 4d. are supernova remnants. thev should show little or no evidence lor variability.," If the soft sources discussed in Section \ref{sec:id} are supernova remnants, they should show little or no evidence for variability."
+ When SNR. are voung (< 1000 vears) several emission mechanisms might contribute to the X-ray. emission. and the (ux may be variable 1995).," When SNR are young $\le$ 1000 years) several emission mechanisms might contribute to the X-ray emission, and the flux may be variable \citep{schl95}."
+. However. once (he remnant enters the adiabatic phase the luminosity should decline eracduallv (Jones.S11ithi.&Straka1981:Laanilton.Chevalier.Sarazin1933).," However, once the remnant enters the adiabatic phase the luminosity should decline gradually \citep{JSS81, ham83}."
+. In contrast. accretion-powered supersoft sources are known to be variable ((xahabka.Pietsch.1994:IXongetal. 2002).. and detection of variability in a large fraction of the soft. sources would support the hypothesisthat thev are accretion-powerecl," In contrast, accretion-powered supersoft sources are known to be variable \citep{kah94,kong02}, and detection of variability in a large fraction of the soft sources would support the hypothesisthat they are accretion-powered."
+ Both MIOLI and M83 have, Both M101 and M83 have
+universe sienilicantlv decrease wilh increasing bhuminositv.,universe significantly decrease with increasing luminosity.
+ Di (his letter we point out that two correction have to be made to the samples to study of the fraction of obscured quasars., In this letter we point out that two correction have to be made to the samples to study of the fraction of obscured quasars.
+ The corrections are: a) radio loud AGNs have to be excluded since their X-ray emission might be dominated by the jet component thus the measured luminosity ancl (he obscuration does not reflect the intrinsic values in the nuclei: b) Compton thick sources have to be excluded {oo since their soft ganna rav. emission are also strongly attenuated by Compton scattering and their intrinsic huminosities are hard to estimate., The corrections are: a) radio loud AGNs have to be excluded since their X-ray emission might be dominated by the jet component thus the measured luminosity and the obscuration does not reflect the intrinsic values in the nuclei; b) Compton thick sources have to be excluded too since their soft gamma ray emission are also strongly attenuated by Compton scattering and their intrinsic luminosities are hard to estimate.
+ After the corrections. we find only marginal decrease in the fraction of obscured. AGN with huninositv in the local universe.," After the corrections, we find only marginal decrease in the fraction of obscured AGN with luminosity in the local universe."
+ Larger samples are required (o reach a more robust conclusion., Larger samples are required to reach a more robust conclusion.
+ Bassani et al. (, Bassani et al. (
+"2006) provides an AGN sample selected in the 20 100 keV band with the IBIS on INTEGRAL,",2006) provides an AGN sample selected in the 20 – 100 keV band with the IBIS on $INTEGRAL$.
+" The sample contains 62 active galactic nuclei (14 of them are unclassified) above a flux limit of ~1.5x10H eres ? ',", The sample contains 62 active galactic nuclei (14 of them are unclassified) above a flux limit of $\sim 1.5\times10^{-11}$ ergs $^{-2}$ $^{-1}$.
+ BOG listed the available column densities (obtained [rom archival X-ray spectra) lor 35 sources will redshifts., B06 listed the available column densities (obtained from archival X-ray spectra) for 35 sources with redshifts.
+ In this letter we provicle an updated list of Ny (and relerences) wilh more recent measurements available in literature (mostly from Chandra or IXMM. observations. Table 1)!.," In this letter we provide an updated list of $N_H$ (and references) with more recent measurements available in literature (mostly from $Chandra$ or $XMM$ observations, Table 1)."
+". Within (his subsample. 10 are radio loud (including five blazars). five are Compton-thick (Nj,>107! 7). the rest 20 are radio quiet and Compton-thin."," Within this subsample, 10 are radio loud (including five blazars), five are Compton-thick $_{H} \geqslant 10^{24}$ $^{-2}$ ), the rest 20 are radio quiet and Compton-thin."
+ The luminosities in the 0. 100 keV band were ealeulated by assuming a Crab-like spectrum., The luminosities in the 20 – 100 keV band were calculated by assuming a Crab-like spectrum.
+" The first 3 months of the SWIFT Durst Alert Telescope (BAT) high Galactic Intitude survev provide a sample of 14. 195 keV band selected sources (M05). with a flux limit of ~ H eres 7s ! and ~2.7 confidence) positional uncertainties for the faintest sources,", The first 3 months of the $SWIFT$ Burst Alert Telescope (BAT) high Galactic latitude survey provide a sample of 14 – 195 keV band selected sources (M05) with a flux limit of $\sim$ $^{-11}$ ergs $^{-2}$ $^{-1}$ and $\sim 2\arcmin.7$ confidence) positional uncertainties for the faintest sources.
+ of the 66 higb-Iatitude sources were identified., of the 66 high-latitude sources were identified.
+ Twelve are Galactic-tvpe sources. and 44 are identified with known AGNs.," Twelve are Galactic-type sources, and 44 are identified with known AGNs."
+ The luminosities in the 14. 195 keV. band. were calculated by: assuming a tvpical power-law spectrum (D~ 1.7)., The luminosities in the 14 – 195 keV band were calculated by assuming a typical power-law spectrum $\Gamma \sim 1.7$ ).
+ MOS listed the absorption column densitv derived from literature or archive N-raw spectra for 39 AGNs., M05 listed the absorption column density derived from literature or archive X-ray spectra for 39 AGNs.
+" We also provide an updated list of Ny, (and references) will more recent measurements available in literature", We also provide an updated list of $N_H$ (and references) with more recent measurements available in literature
+one mass resolution was simulated in this case. corresponding to npa=11325⋅1075∣∙⊥M. and m;=L5−10>hn7M. for. dark matter and gas within the high-resolution region. respectively.,"one mass resolution was simulated in this case, corresponding to $m_{\rm DM}=1.13\times
+10^9\,h^{-1}{\rm M}_\odot$ and $m_{\rm gas}=1.7\times 10^8\,h^{-1}{\rm
+ M}_\odot$ for dark matter and gas within the high–resolution region, respectively."
+ As such. this mass resolution is about 4 times better than the LR runs. and therefore lies intermediate between the MR and the HR runs of#1.," As such, this mass resolution is about 4 times better than the LR runs, and therefore lies intermediate between the MR and the HR runs of."
+. The softening length is set to «14Oh Κρο fixed in physical units below 2=5. while it is kept fixed in comoving units at higher redshift.," The softening length is set to $\epsilon_{\rm Pl}=5.0\, h^{-1}$ kpc, fixed in physical units below $z=5$, while it is kept fixed in comoving units at higher redshift."
+ The reference runs for clusters of assume that 50 per cent of the energy provided by SN G.e. half of that used for the #1)) is carried by winds., The reference runs for clusters of assume that 50 per cent of the energy provided by SN (i.e. half of that used for the ) is carried by winds.
+" This gives a wind speed of e,&340kms.1.", This gives a wind speed of $v_w\simeq 340\vel$.
+ Unless otherwise stated. our analysis presented in this paper is restricted to the CL5 and CL6 simulations out of the 20 clusters of the full#2.," Unless otherwise stated, our analysis presented in this paper is restricted to the CL5 and CL6 simulations out of the 20 clusters of the full."
+. We note that for all of our simulations. we set the smallest allowed value for the SPH smoothing length to ον/4.," We note that for all of our simulations, we set the smallest allowed value for the SPH smoothing length to $\epsilon_{\rm
+ Pl}/4$."
+ In summary. our two sets of simulated. clusters differ in the following aspects: We will use the convention that a certain simulation of each cluster is labeled by the name of the cluster itself. followed. when required. by a label which specifies the resolution used.," In summary, our two sets of simulated clusters differ in the following aspects: We will use the convention that a certain simulation of each cluster is labeled by the name of the cluster itself, followed, when required, by a label which specifies the resolution used."
+ In this way. CLI-MR will indicate the reference run of medium-resolution of the CLI cluster from#1.. while CLS will designate the reference run of this cluster from#2.," In this way, CL1-MR will indicate the reference run of medium–resolution of the CL1 cluster from, while CL5 will designate the reference run of this cluster from."
+. For the latter. we do not specify the resolution. since clusters from are simulated at only one resolution.," For the latter, we do not specify the resolution, since clusters from are simulated at only one resolution."
+ In addition. a further extension of the name of euch run will be provided whenever the run differs from the reference run of the set it belongs to.," In addition, a further extension of the name of each run will be provided whenever the run differs from the reference run of the set it belongs to."
+ The list of such extensions and their description is given in Table 3.., The list of such extensions and their description is given in Table \ref{tab:tests}.
+ For instance. CL4+-HR indicates the high-resolution run of the CL4 cluster. using the standard setup of#1.. while CL4-HR-NW stands for the same simulation. but neglecting the effect of galactic winds.," For instance, CL4-HR indicates the high–resolution run of the CL4 cluster, using the standard setup of, while CL4-HR-NW stands for the same simulation, but neglecting the effect of galactic winds."
+ We show in Figure | the gas density maps of the six simulated clusters within boxes each having a size of 41/4...," We show in Figure \ref{fi:maps} the gas density maps of the six simulated clusters within boxes each having a size of $4R_{\rm
+ vir}$."
+ The four clusters from. the are shown .in their. highest. resolution. version., The four clusters from the are shown in their highest resolution version.
+. In the first step of our simulation analysis we determine suitable cluster centers., In the first step of our simulation analysis we determine suitable cluster centers.
+ These are defined as the position of the most bound particle among those grouped together by a FOF algorithm with linking length 6=0.15 tin units of the mean interparticle separation in the high-resolution region)., These are defined as the position of the most bound particle among those grouped together by a FOF algorithm with linking length $b=0.15$ (in units of the mean interparticle separation in the high–resolution region).
+" Once the center is identified. we apply a spherical overdensity algorithm to determine the virial radius.2,4."," Once the center is identified, we apply a spherical overdensity algorithm to determine the virial radius,."
+".. We here defined this as the radius that encompasses an average density equal to the virial density for the adopted cosmological model. pou(2)=2NGp Gas)=Al(z)Ha]p.ο is the critical density at redshift 2). where the overdensity A,(2) is computed as described by2.. with (0)=100 for the assumed cosmology."," We here defined this as the radius that encompasses an average density equal to the virial density for the adopted cosmological model, $\rho_{\rm vir}(z)=\Delta_c(z)\rho_{c}(z)$ $\rho_c(z)=[H(z)/H_0]^2\rho_{c,0}$ is the critical density at redshift $z$ ), where the overdensity $\Delta_c(z)$ is computed as described by, with $\Delta_c(0)\simeq 100$ for the assumed cosmology."
+" The virial mass.Άλι, is simply the mass contained within the virial radius."," The virial mass, is simply the mass contained within the virial radius."
+ In Table 1.. we provide the typical values ofii. aand mass-weighted temperature.μον. for the six simulated clusters.," In Table \ref{tab:sets}, we provide the typical values of, and mass–weighted temperature, for the six simulated clusters."
+ Profiles of gas-related quantities are computed within 200iOO equispacedüispae| linearinear radialfal |bins.bin outut tovir.. Startingstartingfrom from ia minimum radius which contains 100 gas particles.," Profiles of gas-related quantities are computed within 200 equispaced linear radial bins, out to, starting from a minimum radius which contains 100 gas particles."
+ As shown by?.. numerically stable results can be expected for this choice in radiative simulations of galaxy clusters.," As shown by, numerically stable results can be expected for this choice in non--radiative simulations of galaxy clusters."
+ We identify galaxies in our simulations by applying the SKID to the distribution of star particles., We identify galaxies in our simulations by applying the SKID to the distribution of star particles.
+ In the following. we provide a short description of our algorithm. while a more detailed discussion and presentation of tests is provided elsewhere (Murante et al.," In the following, we provide a short description of our algorithm, while a more detailed discussion and presentation of tests is provided elsewhere (Murante et al."
+ 2005. in preparation).," 2005, in preparation)."
+ Briefly. the SKID algorithm works as follows:," Briefly, the SKID algorithm works as follows:"
+Transiting extrasolar plaucts. like IIDI1180733b offer unique opportunities to scrutinize their atuoxpleric content (e.g. Charbonneau et 22002. Vidal-Madgar et 22003).,"Transiting extrasolar planets, like 189733b offer unique opportunities to scrutinize their atmospheric content (e.g., Charbonneau et 2002, Vidal-Madjar et 2003)."
+ Tn particular. primary transits can reveal tenuous quantities of eas or dust by spectral absorption leading to variations in the apparent planet size as a function of the wavelength.," In particular, primary transits can reveal tenuous quantities of gas or dust by spectral absorption leading to variations in the apparent planet size as a function of the wavelength."
+ 1159732bb is. presently the vost taveet for such studies because of its nearby iud bright host star. its large absorption depth im its transit helt cive (Bouchy ct 22005). aud a short period (IIébbrard Lecavelicr des Etaugs 2006).," b is presently the best target for such studies because of its nearby and bright host star, its large absorption depth in its transit light curve (Bouchy et 2005), and a short period (Hébbrard Lecavelier des Etangs 2006)."
+" Using transit observatious with the ACS camera of the IIubble Space Telescope. the apparent radius of 1159733b las been nieasure from 0.55 to 1.05 nücrous. thus providing the ""ransnmüssion spectrum of the atinosphere (Pout et 22008)."," Using transit observations with the ACS camera of the Hubble Space Telescope, the apparent radius of 189733b has been measured from 0.55 to 1.05 microns, thus providing the “transmission spectrum” of the atmosphere (Pont et 2008)."
+ Iu this wavelength range. strong features of abundant atomic species aud molecules were predicted but rot detected. leacing Pout et ((2008) to the conclusion that a laze of sub-micron particles is present in the upper atinosphliere of the planet.," In this wavelength range, strong features of abundant atomic species and molecules were predicted but not detected, leading Pont et (2008) to the conclusion that a haze of sub-micron particles is present in the upper atmosphere of the planet."
+ Tere we use these measurements of the planet radius as a function of the wavelength from 550 to qui., Here we use these measurements of the planet radius as a function of the wavelength from 550 to nm.
+ Tn Sect., In Sect.
+ 2. we deseribe how a trausimüssiou spectrum can be interpreted to derive basic quantities., \ref{Interpreting} we describe how a transmission spectrum can be interpreted to derive basic quantities.
+ We then estimate the temperature of 1189733b from fit to the ACS spectrum and show that the observed absorption can be caused by Ravleigh scattering (Sect. 3))., We then estimate the temperature of 189733b from fit to the ACS spectrum and show that the observed absorption can be caused by Rayleigh scattering (Sect. \ref{Temperature}) ).
+ Several species that are possibly responsible for the Rayleigh scattering are discussed. aud the correspouding pressure at the absorption level are estimated in Sect. [..," Several species that are possibly responsible for the Rayleigh scattering are discussed, and the corresponding pressure at the absorption level are estimated in Sect. \ref{Rayleigh}."
+ To interpret fransif spectra. one can use detailed atinosphere models that iuclude temperature. pressure. and composition as a function of the altitude and uunercallv iuteerate the radiative transfer equations to obtain a theoretical spectrum to compare to the observations (o... Ehnreurecich ct 22006).," To interpret transit spectra, one can use detailed atmosphere models that include temperature, pressure, and composition as a function of the altitude and numerically integrate the radiative transfer equations to obtain a theoretical spectrum to compare to the observations (e.g., Ehrenreich et 2006)."
+ Tere we propose another approach by deriving first-order equations to obtain a better feeling for the basic quantities hat can be obtained from hese lacasurenents., Here we propose another approach by deriving first-order equations to obtain a better feeling for the basic quantities that can be obtained from these measurements.
+ Following Fortuevs (2005) derivation of the path cheth. the optical depth. r. in a line of sight erazine he planetary Bub at an altitude + is given by 7(À.:)cmG(íAÀ)n(z)VEπΠ ώμο. where Dogs IS the assured lancet radius. Z7 the atinosphere scale height. aud »(:)=DpyycuplLAIT) the volume cusity at the altitude z of the main absorbent with a cross section a(A).," Following Fortney's (2005) derivation of the path length, the optical depth, $\tau$, in a line of sight grazing the planetary limb at an altitude $z$ is given by $\tau(\lambda,z)\approx \sigma(\lambda)n(z)\sqrt{2\pi R_{planet} H}$ , where $R_{planet}$ is the assumed planet radius, $H$ the atmosphere scale height, and $n(z)=n_{(z=0)} exp(-z/H)$ the volume density at the altitude $z$ of the main absorbent with a cross section $\sigma(\lambda)$."
+ For a cluperature T. Π is given by Π=Τμ where pr is he mean uass of atmospheric particles taken to be 2.3 iues the mass of the proton. aud g the eravitv.," For a temperature $T$, $H$ is given by $ H = k T/\mu g$, where $\mu$ is the mean mass of atmospheric particles taken to be 2.3 times the mass of the proton, and $g$ the gravity."
+" Using dlanctary transit measurements. a fit to the light curve xovides the ratio of the effective planetary racius as a πλοίο o| waveleugth to the stellar radius: &,(A)/R"," Using planetary transit measurements, a fit to the light curve provides the ratio of the effective planetary radius as a function of wavelength to the stellar radius: $R_p(\lambda)/R_*$."
+" Ποιο we defue Ty bv the optical depth at altitude τν such that a sharp occultiug disk of radius Figsner|leqxoduces the same absorption depth as the planet with its raushicent atinosphiere: iu other words. 0z4,38 defined by"," Here we define $\tau_{eq}$ by the optical depth at altitude $z_{eq}$ such that a sharp occulting disk of radius $R_{planet}+z_{eq}$produces the same absorption depth as the planet with its translucent atmosphere; in other words, $\tau_{eq}$ is defined by"
+results from the large number of multiple spirals (Elmegreen propagating simultaneously through the disk during the first ~500 Myr.,"results from the large number of multiple spirals \citep{elmegreen92,rix93,mq06} propagating simultaneously through the disk during the first $\sim500$ Myr."
+" Due to the various SS pattern speeds, there are no radii at which the effect of resonance overlap differs distinctly from the rest of the disk (such as the bar’s corotation and OLR)."," Due to the various SS pattern speeds, there are no radii at which the effect of resonance overlap differs distinctly from the rest of the disk (such as the bar's corotation and OLR)."
+" Consequently, we observe a smooth distribution of AL in contrast to the case of bar + SS."," Consequently, we observe a smooth distribution of $\Delta L$ in contrast to the case of bar + SS."
+ Fig., Fig.
+" 3 shows a power spectrum displaying the pattern speeds of the m=2 and m=4 (two-armed and four-armed structure) Fourier components, over the whole simulation, of the gSa (top) and gSb (bottom) models."," \ref{fig:omega} shows a power spectrum displaying the pattern speeds of the $m=2$ and $m=4$ (two-armed and four-armed structure) Fourier components, over the whole simulation, of the gSa (top) and gSb (bottom) models."
+" Contour levels are indicated for each panel in units of Qr, the ratio of the maximum tangential force to the azimuthally averaged radial force at a given radius."," Contour levels are indicated for each panel in units of $Q_T$, the ratio of the maximum tangential force to the azimuthally averaged radial force at a given radius."
+" For both simulations, the SS (rZ5 kpc) spans a range of pattern speeds, always slower than the bar."," For both simulations, the SS $r\ga5$ kpc) spans a range of pattern speeds, always slower than the bar."
+" As expected from the results of MF10, the SS pattern speed has little influence on the location of the maxima in the bimodal distribution of AL."," As expected from the results of MF10, the SS pattern speed has little influence on the location of the maxima in the bimodal distribution of $\Delta L$."
+" Note that the gSa bar evolves (slows down and extends over time) much more, and that its two-armed SS is twice stronger compared to the gSb model (top left)."," Note that the gSa bar evolves (slows down and extends over time) much more, and that its two-armed SS is twice stronger compared to the gSb model (top left)."
+" Despite the slightly stronger gSb bar, the effect on AL for the gSa is much more prominent in the outer disk (Fig.2)) because of the combined effect of the bar and stronger spirals."," Despite the slightly stronger gSb bar, the effect on $\Delta L$ for the gSa is much more prominent in the outer disk \ref{fig:gSall}) ) because of the combined effect of the bar and stronger spirals."
+ Strong spirals are thus needed for this migration mechanism to be efficient., Strong spirals are thus needed for this migration mechanism to be efficient.
+" However, the relation between the strength of the perturbers and the angular momentum redistribution is nonlinear as previously shown by MF10."," However, the relation between the strength of the perturbers and the angular momentum redistribution is nonlinear as previously shown by MF10."
+" To determine whether the strong mixing we described above is caused by a deficiency in the resolution of the GalMer simulations, we ran pure N-body collisionless simulations with 107 and 1.27x10? particles in the disk."," To determine whether the strong mixing we described above is caused by a deficiency in the resolution of the GalMer simulations, we ran pure N-body collisionless simulations with $10^7$ and $1.27\times10^5$ particles in the disk."
+ Full description and details can be found in WozniakandMichel-Dansac(2009)., Full description and details can be found in \cite{wozniak09}.
+". With a scale-length of 3.5 kpc, the initial conditions of the disk are comparable to those of the gSb GalMer model (Fig. 2))"," With a scale-length of 3.5 kpc, the initial conditions of the disk are comparable to those of the gSb GalMer model (Fig. \ref{fig:gSall}) )"
+ but lack a gas component., but lack a gas component.
+" These simulations have no halo, which causes a ~30% drop in the RC at ~10 kpc."," These simulations have no halo, which causes a $\sim30\%$ drop in the RC at $\sim10$ kpc."
+ Fig., Fig.
+" 4 shows the changes in angular momentum at t=1 and 3 Gyr for 1.27x10° (top) and 10"" (bottom) disk particles.", \ref{fig:nbody} shows the changes in angular momentum at t=1 and 3 Gyr for $1.27\times10^5$ (top) and $10^7$ (bottom) disk particles.
+" At the beginning of the simulations, the effect is stronger for the low-resolution case due to the faster bar formation."," At the beginning of the simulations, the effect is stronger for the low-resolution case due to the faster bar formation."
+" However, at f&2 Gyr both the high- and low-resolution runs yield a similar result."," However, at $t\approx2$ Gyr both the high- and low-resolution runs yield a similar result."
+ This suggests that the magnitude of the radial migration induced by the bar-spiral interaction and the timescale over which it is effective are not strongly dependent on the numerical resolution of the simulations., This suggests that the magnitude of the radial migration induced by the bar-spiral interaction and the timescale over which it is effective are not strongly dependent on the numerical resolution of the simulations.
+" We conclude that the properties of the mixing, i.e., its magnitude and timescale, that we infer from the GalMer simulations are only weakly affected by the insufficient resolution."," We conclude that the properties of the mixing, i.e., its magnitude and timescale, that we infer from the GalMer simulations are only weakly affected by the insufficient resolution."
+ We now investigate whether the resonance overlap mechanism is also efficient in low-mass galaxies., We now investigate whether the resonance overlap mechanism is also efficient in low-mass galaxies.
+" Gogartenetal.(2010) showed that although transient spirals can explain extended disks for MW-type galaxies, this mixing mechanism is inefficient for galaxies with RCs of V.~100 km/s, such as NGC 300 and M33."," \citet{gogarten10} showed that although transient spirals can explain extended disks for MW-type galaxies, this mixing mechanism is inefficient for galaxies with RCs of $V_c\sim100$ km/s, such as NGC 300 and M33."
+" Nevertheless, both NGC 300 and M33 are observed to have extended radial profiles of up to ten"," Nevertheless, both NGC 300 and M33 are observed to have extended radial profiles of up to ten"
+Atmospheric dispersion is not accounted for im our pipeline. but we are aware that in certain conditions (presence of emission lines in the spectrum). it can niic asviuuetry along the elevation axis.,"Atmospheric dispersion is not accounted for in our pipeline, but we are aware that in certain conditions (presence of emission lines in the spectrum), it can mimic asymmetry along the elevation axis."
+ Such false detection can be avoided by observation at different parallactic aueles., Such false detection can be avoided by observation at different parallactic angles.
+ The effect of fringe bluriug due to the large baudwith of L' filter can technically be accounted for iu the fitting process., The effect of fringe blurring due to the large bandwith of L' filter can technically be accounted for in the fitting process.
+ It would require including a loss of füuge coherence ou the edge of the diffraction pattern., It would require including a loss of fringe coherence on the edge of the diffraction pattern.
+ However. at this stage. the SAAIP pipeline does not iuclude this correction.," However, at this stage, the SAMP pipeline does not include this correction."
+ It is. however. not seen as critical with the seveu-hole mask. since the diffraction pattern of the 1.2 neter holes are significautly smaller than the coherence cheth (see Fie. 1)).," It is, however, not seen as critical with the seven-hole mask, since the diffraction pattern of the 1.2 meter holes are significantly smaller than the coherence length (see Fig. \ref{Image}) )."
+ It becomes an acute problem with he znall holes of the 18-hole mask. which then justifics using a narrow band filter.," It becomes an acute problem with the small holes of the 18-hole mask, which then justifies using a narrow band filter."
+ The complex visibilities consist of amsplitucles and yhases. but for the purposes of binary detection. we vpically rely on the phases only (of the bispectruni).," The complex visibilities consist of amplitudes and phases, but for the purposes of binary detection, we typically rely on the phases only (of the bispectrum)."
+ This is because of the dependence of absolute visibilities ou the adaptive optics correction., This is because of the dependence of absolute visibilities on the adaptive optics correction.
+ They are thus harder to calibrate. and are subject to different biases between the science and the calibrator object.," They are thus harder to calibrate, and are subject to different biases between the science and the calibrator object."
+" Tlowever. by default. the pipelines also process the power spectra. and estimate fringe visibilities,"," However, by default, the pipelines also process the power spectrum and estimate fringe visibilities."
+ Performance of SAAS visibility estimation will be the subject of a forthcoming paper., Performance of SAM's visibility estimation will be the subject of a forthcoming paper.
+ Aperture masking data vields Fourier pliases aud amplitudes., Aperture masking data yields Fourier phases and amplitudes.
+ They are functionally equivalent to those recovered iu lone-bascline interferometry. aud therefore the same techniques aud software can be applied.," They are functionally equivalent to those recovered in long-baseline interferometry, and therefore the same techniques and software can be applied."
+ Image reconstruction is a possibility. but to fulfill the difficult requirements of faint companion cletection. direct fitting iu the Fourier domain is essential.," Image reconstruction is a possibility, but to fulfill the difficult requirements of faint companion detection, direct fitting in the Fourier domain is essential."
+ Modeling the preseuce of a two point-like sources in the Fourier domain is trivial., Modeling the presence of a two point-like sources in the Fourier domain is trivial.
+ The complex visibility can be written as: where wand e are orthogonal spatial frequency vectors (respectively to the east and to the north). à and à are he right ascension aud declination in radians. aud + is the Hux ratio between the primary aud secoudarv.," The complex visibility can be written as: where $u$ and $v$ are orthogonal spatial frequency vectors (respectively to the east and to the north), $\alpha$ and $\delta$ are the right ascension and declination in radians, and $r$ is the flux ratio between the primary and secondary."
+ Figure 3 presents the interferometer fiuge phase for a erid of ime models with three differeut angular separations aud hree different flux ratios., Figure \ref{Model} presents the interferometer fringe phase for a grid of nine models with three different angular separations and three different flux ratios.
+ A direct relationship betweeu he phases and the basic properties of the binary svsteni can be noted., A direct relationship between the phases and the basic properties of the binary system can be noted.
+ The phases exhibit a sinusoidal curve across the Fourier plane. with the amplitude proportional o the contrast ratio while the period is proportional," The phases exhibit a sinusoidal curve across the Fourier plane, with the amplitude proportional to the contrast ratio while the period is proportional"
+"and GRB 090510, which may include a jet break), and we determine the presence of jet breaks in the X-ray afterglows of the BAT and GBM samples using the criteria from ?,, we can evaluate the jet opening angles and collimation-corrected energetics as a function of these populations, as shown in Figure 9..","and GRB 090510, which may include a jet break), and we determine the presence of jet breaks in the X-ray afterglows of the BAT and GBM samples using the criteria from \cite{racusin09}, we can evaluate the jet opening angles and collimation-corrected energetics as a function of these populations, as shown in Figure \ref{fig:energetics}."
+" For those GRBs with only lower limits on the jet break we use the time of last detection to determine the lower times,limit on 0;, and therefore also E."," For those GRBs with only lower limits on the jet break times, we use the time of last detection to determine the lower limit on $\theta_j$, and therefore also $E_{\gamma}$ ."
+" As demonstrated in ?,, there are several different characteristic times for which one can place limits on jet breaks, and the large error bars on late-time light curve data points can mask jet breaks (see also ?))."," As demonstrated in \cite{racusin09}, there are several different characteristic times for which one can place limits on jet breaks, and the large error bars on late-time light curve data points can mask jet breaks (see also \citealt{curran08}) )."
+" However, for Figure 9,, we simply use the time of last detection."," However, for Figure \ref{fig:energetics}, we simply use the time of last detection."
+" We have showed that there are observational differences between the BAT, GBM, and LAT samples throughout the previous sections."," We have showed that there are observational differences between the BAT, GBM, and LAT samples throughout the previous sections."
+" However, the difficulty lies in out the instrumental selection effects from the physical separatingdifferences between GRBs that produce >100MeV emission, and those that do not."," However, the difficulty lies in separating out the instrumental selection effects from the physical differences between GRBs that produce appreciable $> 100 ~\textrm{MeV}$ emission, and those that do not."
+ The appreciablemedian and standard deviation of the distributions of the many observational parameters discussed in the previous and following sections are presented for each sample in Table 3.., The median and standard deviation of the distributions of the many observational parameters discussed in the previous and following sections are presented for each sample in Table \ref{tab:medsig}.
+" In the following section, we will explore physical explanations for the observable parameter distributions including a calculation of the radiative efficiency, and speculate on the origin of the afterglow luminosity clustering."," In the following section, we will explore physical explanations for the observable parameter distributions including a calculation of the radiative efficiency, and speculate on the origin of the afterglow luminosity clustering."
+ We also will compare and contrast the other recent studies of the broadband observations of the LAT bursts., We also will compare and contrast the other recent studies of the broadband observations of the LAT bursts.
+ Our first attempt to explore the underlying physics is by calculating the radiative efficiency of the GRBs at turning their kinetic energy into radiation during the prompt emission., Our first attempt to explore the underlying physics is by calculating the radiative efficiency of the GRBs at turning their kinetic energy into radiation during the prompt emission.
+" We follow the formulation of ?,, which derives the kinetic energy from the X-ray and by comparing(£;) the 4-ray prompt afterglowemission observations,output, we can estimatea radiative"," We follow the formulation of \cite{zhang07}, which derives the kinetic energy $E_k$ ) from the X-ray afterglow observations, and by comparing the $\gamma$ -ray prompt emission output, we can estimatea radiative"
+"this particular softening both because it is numerically convenient, and because it reduces exactly to the solution for a disk with scale height 8=h/R (with p2X/2h at |z| <A) when h«R.","this particular softening both because it is numerically convenient, and because it reduces exactly to the solution for a disk with scale height $\beta=h/R$ (with $\rho=\Sigma/2\,h$ at $|z|<h$ ) when $h\ll R$."
+ This formulation of the indirect potential follows ?.., This formulation of the indirect potential follows \citet{murraydermott}.
+" The inclusion of the indirectpotential is necessary because the m=1 mode changes the center of mass, leading to motion by the BH about that center of mass."," The inclusion of the indirectpotential is necessary because the $m=1$ mode changes the center of mass, leading to motion by the BH about that center of mass."
+" Recall that our coordinate frame is centered on the BH, thus r is the vector distance to the BH, and the frame is rotating about the BH."," Recall that our coordinate frame is centered on the BH, thus $r$ is the vector distance to the BH, and the frame is rotating about the BH."
+" For a given c, 7, and f, then, the solution to these coupled equations determines the perturbation structure."," For a given $\omega$, $\eta$, and $\beta$, then, the solution to these coupled equations determines the perturbation structure."
+" Because we are interested in the behavior where w/€) may be non-trivial, the equations are non-linear in w."," Because we are interested in the behavior where $\omega/\Omega$ may be non-trivial, the equations are non-linear in $\omega$."
+" However, we can combine the equation in Na, Ua, and να to obtain a single equation [ drr(PG,r)/8r)Xa’)and. 9"" f(8PG,r)/Or)Dar), and vy =OInX/Olnr."," However, we can combine the equation in $\Sigma_{a}$, $u_{a}$, and $v_{a}$ to obtain a single equation Where $\Phi^{\prime} = \int\,{dr^{\prime}}\,r^{\prime}\,(\partial P(r,\,r^{\prime})/\partial r)\,\Sigma_{a}(r^{\prime})$ and $\Phi^{\prime\prime} = 
+\int\,{dr^{\prime}}\,r^{\prime}\,(\partial^{2} P(r,\,r^{\prime})/\partial r^{2})\,\Sigma_{a}(r^{\prime})$ , and $\nu_{X}\equiv \partial \ln{X}/\partial \ln{r}$."
+"If we discretize 2, into some grid in logr and apply some summation rule to evaluate the integrals in , then we can write this as a linear operator acting on the perturbed density, (731) 20 The eigenvalues w and eigenvectors 32, represent the exact homogenous solutions to this equation."," If we discretize $\Sigma_{a}$ into some grid in $\log{r}$ and apply some summation rule to evaluate the integrals in $\Phi$, then we can write this as a linear operator acting on the perturbed density, ) = The eigenvalues $\omega$ and eigenvectors $\bar{\Sigma}_{a}$ represent the exact homogenous solutions to this equation."
+" In practice, we obtain solutions following the method outlined in Appendix B of ?"," In practice, we obtain solutions following the method outlined in Appendix B of \citet{adams89:eccentric.instab.in.keplerian.disks}."
+" If we multiply Equation 39 by A’, then we can eliminate all occurrences of w in the denominator and write the resulting matrix equation as a fifth-order equation in w:"," If we multiply Equation \ref{eqn:mode.eom} by $\Delta^{2}$ , then we can eliminate all occurrences of $\omega$ in the denominator and write the resulting matrix equation as a fifth-order equation in $\omega$ :"
+T~3.10seeto with au uncertainty of a factor 3 either way.,"$\Gamma\sim3\times10^{-12}\,\rm {sec}^{-1}$, with an uncertainty of a factor 3 either way."
+ They also found that.within their wacertainty. DL did not vary significantly with z. Of the many later attempts to determine P aud its possible variation with z. he most recent are due to Cdallougo et al (1996). Lu et αἱ (1996). Savaelio et al (1997) aud Cooke. Epsev Carswell (1997).," They also found that, within their uncertainty, $\Gamma$ did not vary significantly with z. Of the many later attempts to determine $\Gamma$ and its possible variation with z, the most recent are due to Giallongo et al (1996), Lu et al (1996), Savaglio et al (1997) and Cooke, Epsey Carswell (1997)."
+" Despite the considerable remaining uncertainties here is a consensus that D~2<10sce!, with πο sieuificaut z variation out to +=[.5."," Despite the considerable remaining uncertainties there is a consensus that $\Gamma\sim2\times10^{-12}\,\rm{sec}^{ -1}$, with no significant z variation out to $z=4.5$."
+ There has been much discussion iu the literature as o whether the population of QSOs alone cau provide a sufficieut ionising flux tfo accouut for P., There has been much discussion in the literature as to whether the population of QSOs alone can provide a sufficient ionising flux to account for $\Gamma$ .
+ According to IHanudt Madan (1996) QSOToso~10seo bat. —2. mut only 2«1013seelatz~L5.," According to Haardt Madau (1996) $\Gamma_{\rm QSO}\sim10^{-12}\,\rm {sec}^{-1}$ at $z \sim 2$, but only $2\times10^{-13}\,\rm {sec}^{-1}\rm {at}\,z\sim 4.5$."
+ The discrepancy of a actor ~10 at 2.~L5 appears to be significant. even ifhe one at 2~ lies within the uncertainties.," The discrepancy of a factor $\sim 10$ at $z\sim 4.5$ appears to be significant, even ifthe one at $z\sim 2$ lies within the uncertainties."
+ the estimates of. P from ⋅⋅the proximity effect⋅ which.... Lyo rave so far been made have been based on the simple asstunption that the influence of the QSO on the nearby clouds is entirely due to the additional direct ionisatiou which it produces., All the estimates of $\Gamma$ from the proximity effect which have so far been made have been based on the simple assumption that the influence of the QSO on the nearby clouds is entirely due to the additional direct ionisation which it produces.
+ However. as Miraldz-Escudóé Rees (1991) (ATR) pointed out. one should also allow for the additional leat input due to the radiation frou the QSO.," However, as Miralda-Escudé Rees (1994) (MR) pointed out, one should also allow for the additional heat input due to the radiation from the QSO."
+ The resulting expansion of the clouds reduces thea electron density. aud the increase of temperature reduces 10 recombination coefficient. so that the recombination rate of the clouds is reduced.," The resulting expansion of the clouds reduces their electron density, and the increase of temperature reduces the recombination coefficient, so that the recombination rate of the clouds is reduced."
+" The uct additional ionisation xoduced bythe QSO is thus increased. aud sothe iuplied value of the background ionisation rate is also increased,"," The net additional ionisation produced by the QSO is thus increased, and so the implied value of the background ionisation rate is also increased."
+ AIR further pointed out that. since the main contribution o the additional heating comes frou the ionisation of Te II in the clouds. the effect is larger where the QSO (but rot the general UV backeround) is capable of ionising Πο IL," MR further pointed out that, since the main contribution to the additional heating comes from the ionisation of He II in the clouds, the effect is larger where the QSO (but not the general UV background) is capable of ionising He II."
+ This remark has become particularly pertinout now. j)oenmuse it sccus that iouisation breakthrough for He ID iu the iuterealactic iiedium may uot have occured until Do3 (Sonegaila Cowie 1996. Hogan et al 1997. Reimers et al 1997. Bokscubere 1998. Songaila 1998). a situation which was theoretically anticipated by AIR aud by Macau Meissindle (1991).," This remark has become particularly pertinent now, because it seems that ionisation breakthrough for He II in the intergalactic medium may not have occured until $z\sim 3$ (Songaila Cowie 1996, Hogan et al 1997, Reimers et al 1997, Boksenberg 1998, Songaila 1998), a situation which was theoretically anticipated by MR and by Madau Meiksin (1994)."
+10€ While the ποΊσα]. value of. the MB. effect. is. model dependent (preliminary∙∙ attempts to estinate∙ it. haviug. been made by MB themselves) it is clear that it has two consequences for our discussion., While the numerical value of the MR effect is model dependent (preliminary attempts to estimate it having been made by MR themselves) it is clear that it has two consequences for our discussion.
+ It leads to au increase in the (sinall) discrepancy between EP aud Toso at i~2. aud to an iucrease of D with z. in contrast to the rapid decrease in. Tose.," It leads to an increase in the (small) discrepancy between $ \Gamma$ and $\Gamma_{\rm QSO}$ at $z\sim
+2$, and to an increase of $\Gamma$ with z, in contrast to the rapid decrease in $\Gamma_{\rm QSO}$."
+ It has often been suggested that any discrepaney between Τ aud Poso could be resolved by appealing to ionising radiation ciuitted by hot stars in galaxies (ee., It has often been suggested that any discrepancy between $\Gamma$ and $\Gamma_{\rm QSO}$ could be resolved by appealing to ionising radiation emitted by hot stars in galaxies (eg.
+ Miralda-Escudé aud Ostriker 1990)., Miralda-Escudé and Ostriker 1990).
+ Recent discussions of this possibility have been given by Ciroux Shapiro (1996) aud by Madan Shull (1996)., Recent discussions of this possibility have been given by Giroux Shapiro (1996) and by Madau Shull (1996).
+ Metalchrichinent aremuents inply that if ~25{ ofthe Lyman continu photons emitted by lot stars escape from them parent galaxies. then these photons would be responsible for an ionisation rate ~10κος| per atom at olo3.," Metal-enrichment arguments imply that if $\sim 25 \%$ of the Lyman continuum photons emitted by hot stars escape from their parent galaxies, then these photons would be responsible for an ionisation rate $\sim10^{-12}\,\rm
+{sec}^{-1}$ per H atom at $z\sim 3$."
+ The actual escape fraction is uot known at lieh z. but at :~(0 it is less than 1% (Deharveng ct al 1997).," The actual escape fraction is not known at high z, but at $z\sim 0$ it is less than $1 \%$ (Deharveng et al 1997)."
+" This strong constraint was deduced bv relating the 77, Iuuinositv density of star-forming galaxies iu the local universe (Gallego et al 1995) to the 11 observatious of Χους et al (1995) anc Donahue et al (1995) which. as already mentioned. lead to an upper limit on the ivdrosen-donising flux at :~0."," This strong constraint was deduced by relating the $H_\alpha$ luminosity density of star-forming galaxies in the local universe (Gallego et al 1995) to the $H_\alpha$ observations of Vogel et al (1995) and Donahue et al (1995) which, as already mentioned, lead to an upper limit on the hydrogen-ionising flux at $z\sim 0$."
+ Although the escape fraction is not known at ligh z. there exists evidence for considerable dust extinction in galaxies at these redshifts (Aleurer 1997. Cimatti et al 1997).," Although the escape fraction is not known at high z, there exists evidence for considerable dust extinction in galaxies at these redshifts (Meurer 1997, Cimatti et al 1997)."
+ Absorption by atomic livdrogen⊓ iu these ⊓galaxies may also be important., Absorption by atomic hydrogen in these galaxies may also be important.
+ This arguinent has recently been much strenettlence x the observations of Spinrad et al (1998) which failed detect any Lyman coutimmun radiation escaping fron All eee)man-linit Oogalaxies;, This argument has recently been much thened by the observations of Spinrad et al (1998) which failed to detect any Lyman continuum radiation escaping from $z>3$ Lyman-limit galaxies.
+" AccordingC» to these authors it is plausible that ionising radiatiou from voung galaxies can replace QSO ionisation at 21,", According to these authors it is implausible that ionising radiation from young galaxies can replace QSO ionisation at $z>4$.
+ Another muclie-discussed. possibility is that radiation from stars at red shifts much ereater than 5 (the so-called Population- IIT stars) might make au appreciable coutzibution to P aud be respousible for the reiouisation of the universe at z25., Another much-discussed possibility is that radiation from stars at red shifts much greater than 5 (the so-called Population III stars) might make an appreciable contribution to $\Gamma$ and be responsible for the reionisation of the universe at $z>5$.
+ Recent discussious of this possibility have been given bv Taman Loeb (1997). Cuedin Ostriker (1997) and Cuedin (1998).," Recent discussions of this possibility have been given by Haiman Loeb (1997), Gnedin Ostriker (1997) and Gnedin (1998)."
+ An important aspect of this proposal concerns the level of metallicity which would result from the required stellay activity., An important aspect of this proposal concerns the level of metallicity which would result from the required stellar activity.
+ It is not clear whether this level agrees with observation. especially iu view ofthe low metallicity recently derived by Sougaila (1997) for the Lyman a clouds.," It is not clear whether this level agrees with observation, especially in view of the low metallicity recently derived by Songaila (1997) for the Lyman $\alpha$ clouds."
+ We cannot go iuto this iutricate question lere.and since we are seckiug au upper linit ou the intergalactic flax of decay plhotous. it will wow be assuiued that most of the iuissng ionising photons at +72toL5 are produced by the cosmological distribution of decaying ueutrinos and not by hot stars.," We cannot go into this intricate question here,and since we are seeking an upper limit on the intergalactic flux of decay photons, it will now be assumed that most of the missing ionising photons at $z\sim 2\; \rm {to}\; 4.5$ are produced by the cosmological distribution of decaying neutrinos and not by hot stars."
+" This assumption would enable us to uuderstaud why the usual interpretation of the proximity effect leads to a value of P which is approximately incepeudent of z. If. 2 oy.~2.9DE is of uuthe sa atue general order as Tose then. as z Increases, D; would iucrease as (1|2)?3 while Tose decreases. leaving DP approximately coustaut."," This assumption would enable us to understand why the usual interpretation of the proximity effect leads to a value of $\Gamma$ which is approximately independent of z. If, at $z\sim 2$, $\Gamma_{\nu}$ is of the same general order as $\Gamma_{\rm QSO}$ then, as z increases, $\Gamma_{\nu}$ would increase as ${(1+z)}^{3\over 2}$ while $\Gamma_{\rm QSO}$ decreases, leaving $\Gamma$ approximately constant."
+" It would seem that a rough upper Iuuit for D, can be derived by setting P,~22Poso at 2=2."," It would seem that a rough upper limit for $\Gamma_\nu$ can be derived by setting $\Gamma_\nu\sim 2\,\Gamma_{\rm QSO}$ at $z=2$."
+" Then. since at this. red shiftκο Pose~410soeDt according. to ILaaret Madau (1996). D, would ~2.101sec| and P—«10.4?see 1, "," Then, since at this red shift $\Gamma_{\rm QSO}\sim\,10^{-12}\,{\rm sec} ^{-1}$ according to Haardt Madau (1996), $\Gamma_\nu$ would $\sim2\times 10^{-12}\,{\rm sec}^{-1}$ and $\Gamma\sim 3\times 10^{- 12}\,{\rm sec} ^{-1}$ ."
+This excess ofP over the proximity effect value of ~2«10eet could be attributed to the MIR effect.," This excess of $\Gamma$ over the proximity effect value of $\sim2\times 10^{-12}\,
+{\rm sec}^{-1}$ could be attributed to the MR effect."
+" Tudeed. since in this picture D, isthe douiuaut contributor to T. the spectrum of the QSO would differ appreciably from that of the backerouud. aud the AIR effect would then be greater than in the pure QSO case (Rees 1990. Sciama 1995)."," Indeed, since in this picture $\Gamma_\nu$ isthe dominant contributor to $\Gamma$ , the spectrum of the QSO would differ appreciably from that of the background, and the MR effect would then be greater than in the pure QSO case (Rees 1990, Sciama 1995)."
+ Finally. it should be noted that. with our adopted values. P would increase bv a factor," Finally, it should be noted that, with our adopted values, $\Gamma$ would increase by a factor"
+ 5-ray (HIewishetal.1969).. 2006).. (Lorimeretal.2007;Ruf2009:Durke-Spolaor&Bailes2010:Keaneetal.2011).. (Rossietal.2008)..(Drunthaleretal.2009).. (Giannios&Metzger2011:Bower2011).. (Falckeetal.1999).. (Fiedleretal.1937).. elal.2007).. (Nakar&Piran2011)... (Levinsonal.2011:Bellet2011).. (Ixidaetal.2008:Matsumura2009)..," $\gamma$ \citep{1969Natur.224..472H}, \citep{2006Nature...mcl}, \citep{2007Sci...318..777L,2009GeoRL..3613202R,2010MNRAS.402..855B,2011arXiv1104.2727K}. \citep{2008MNRAS.390..675R},\citep{2009A&A...499L..17B}, \citep{2011arXiv1102.1429G,2011arXiv1103.4328B}, \citep{1999ApJ...514L..17F}, \citep{1987Natur.326..675F}, \citep{2007ApJ...663L..25H}, \citep{2011arXiv1102.1020N}. \citep{2002ApJ...576..923L,2006ApJ...639..331G,2007ApJ...666..346B,2010AJ....140..157B,2011ApJ...728L..14B,2011MNRAS.412..634B,2011arXiv1103.0511B}. \citep{2008NewA...13..519K,2009AJ....138..787M}."
+In this work we have examined the structure of accretion flows that includes shocks of more general character than those discussed so lar in the literature (e.g..Chakrabarti1990).,"In this work we have examined the structure of accretion flows that includes shocks of more general character than those discussed so far in the literature \citep[e.g.,][]{Cha90}."
+. In particular we have examined whether il is possible to have flows with shocks a( which part of the mass and/or energy fluxes are lost ancl do not participate in the shock transition (this maybe the case in multidimensional shocks or in shocks that involve acceleration of particles that escape Irom the shock region)., In particular we have examined whether it is possible to have flows with shocks at which part of the mass and/or energy fluxes are lost and do not participate in the shock transition (this maybe the case in multidimensional shocks or in shocks that involve acceleration of particles that escape from the shock region).
+ Such generalized shocks obev jump conditions more general than those of Rankine-IIugoniot and it is notprior? certain if accretion flows can allow for such shocks., Such generalized shocks obey jump conditions more general than those of Rankine-Hugoniot and it is not certain if accretion flows can allow for such shocks.
+ We have also examined the degree of mass and energy loss allowed that are at the same time consistent with the continuation of the downstream flow through another sonic transition onto the black hole., We have also examined the degree of mass and energy loss allowed that are at the same time consistent with the continuation of the downstream flow through another sonic transition onto the black hole.
+ In this respect we have focused our study on shocks in which the energy per particle for those escaping the shock (ransilion is greater (lan that of the local gravitational potential., In this respect we have focused our study on shocks in which the energy per particle for those escaping the shock transition is greater than that of the local gravitational potential.
+ In such a case one can argue (that these particles can escape to infinity producing the jets/winds observed in many accretion-powered svstems., In such a case one can argue that these particles can escape to infinity producing the jets/winds observed in many accretion-powered systems.
+ We, We
+0.,.
+"0158. From the narrow Ha emission line measured in the optical spectra presented below, we measure a value of z=0.0157+0.0013 (the error corresponds to anuncertainty of +400 km 51)."," From the narrow $\alpha$ emission line measured in the optical spectra presented below, we measure a value of $z=0.0157\pm 0.0013$ (the error corresponds to anuncertainty of $\pm 400$ km $^{-1}$ )."
+ Adopting a Hubble constant of Ho = 73.8+2.4 km s! Mpc7! our measured redshift from the spectra of SN 2006V corresponds to a distance of 72.7+5.0 Mpc.," Adopting a Hubble constant of $_0$ = $73.8\pm2.4$ km $^{-1}$ $^{-1}$ \citep{riess11}
+ our measured redshift from the spectra of SN 2006V corresponds to a distance of $72.7\pm5.0$ Mpc."
+" Here we have taken into account peculiar motion corrections (Virgo+GA+Shapley, see Mould et al."," Here we have taken into account peculiar motion corrections (Virgo+GA+Shapley, see Mould et al."
+" 2000) and cosmological parameters of Q,,=0.30 and Qa=0.70.", 2000) and cosmological parameters of $\Omega_{m}=0.30$ and $\Omega_{\Lambda}=0.70$.
+ was discovered on 7.2 March 2006 UT in the Scd galaxyUGC 11057 by the Tenagra Observatory Supernova Search (Trondaletal]2006)., was discovered on 7.2 March 2006 UT in the Scd galaxyUGC 11057 by the Tenagra Observatory Supernova Search \citep{trondal06}.
+" With J2000.0 coordinates a= 17°57™13.56°, 6=12°11/03/2, this object was located 17/00 West and 18""22 North from the center of the host galaxy (see Fig. 2))."," With J2000.0 coordinates $\alpha = 17^{\rm h}57^{\rm m}13.56^{\rm s}$ , $\delta =
+12^\circ11\arcmin03\farcs2$, this object was located 0 West and 2 North from the center of the host galaxy (see Fig. \ref{gal06au}) )."
+" On 13.6 March UT, classified SN 2006au as a Type II SN, and noted a well developed Ha P-Cygni profile."," On 13.6 March UT, \citet{blanc06} classified SN 2006au as a Type II SN, and noted a well developed $\alpha$ P-Cygni profile."
+" NED lists a heliocentric redshift to UGC 11057 of z=0.0099, which is in agreement with the value measured from our spectra of z=0.0098+0.0013."," NED lists a heliocentric redshift to UGC 11057 of $z = 0.0099$, which is in agreement with the value measured from our spectra of $z = 0.0098\pm0.0013$."
+" Our measured redshift to UGC 11057 corresponds to a distance of 46.2+3.2 Mpc, where we again adopted the above-mentioned cosmology and peculiar motions."," Our measured redshift to UGC 11057 corresponds to a distance of $46.2\pm3.2$ Mpc, where we again adopted the above-mentioned cosmology and peculiar motions."
+ The organization of this paper is as follows: Section contains brief details regarding the observations and subsequent data reduction techniques; Section ?? presents the broad-band optical and near-infrared photometry; Section ??contains the spectroscopic observations; Section gives our discussion including a comparison with other 1987A-like objects and SNe from BSGs; and this is followed by our conclusions in Section ??.., The organization of this paper is as follows: Section \ref{sec:data_red} contains brief details regarding the observations and subsequent data reduction techniques; Section \ref{sec:photometry} presents the broad-band optical and near-infrared photometry; Section \ref{sec:spectra} contains the spectroscopic observations; Section \ref{sec:discussion} gives our discussion including a comparison with other 1987A-like objects and SNe from BSGs; and this is followed by our conclusions in Section \ref{sec:conclusions}.
+ Broad-band imaging of SNe 2006V and 2006au was obtained with facilities at the Las Campanas Observatory (LCO)., Broad-band imaging of SNe 2006V and 2006au was obtained with facilities at the Las Campanas Observatory (LCO).
+ Optical (ugriBV) and near-infrared (Y JH) imaging was performed with the Henrietta Swope 1-m and the Irénnéee du Pont 2.5-m telescopes equipped with CSP filters., Optical $ugriBV$ ) and near-infrared $YJH$ ) imaging was performed with the Henrietta Swope 1-m and the Irénnéee du Pont 2.5-m telescopes equipped with CSP filters.
+" The optical images were obtained with the Swope Direct Camera equipped with the CCD named Site 3, and with the du Pont Direct Camera equipped with the CCD named Tek 5."," The optical images were obtained with the Swope Direct Camera equipped with the CCD named Site 3, and with the du Pont Direct Camera equipped with the CCD named Tek 5."
+ In the following we adopt Site 3 and Tek 5 as names to distinguish these cameras., In the following we adopt Site 3 and Tek 5 as names to distinguish these cameras.
+ Near-infrared imaging was obtained with RetroCam on Swope and with the Wide Field IR Camera (WIRC) on the du Pont., Near-infrared imaging was obtained with RetroCam on Swope and with the Wide Field IR Camera (WIRC) on the du Pont.
+ Details regarding these instruments and the bandpasses used are given in etal(2006). ECaT] StritzingerGOTT).," Details regarding these instruments and the bandpasses used are given in \citet{hamuy06}, \citet{contreras10}, and \citet{stritzinger11}."
+" Detailed COTO),descriptionsand of the observing techniques and data reduction methodology can be found etal.in|Contreras(2010);; in what follows we briefly summarize the data reduction process.", Detailed descriptions of the observing techniques and data reduction methodology can be found in \citet{contreras10}; in what follows we briefly summarize the data reduction process.
+" All optical images were reduced in a standard manner including: (i) bias subtraction, (ii) flat-field division, and (iii) the application of a shutter time and linearity correction."," All optical images were reduced in a standard manner including: (i) bias subtraction, (ii) flat-field division, and (iii) the application of a shutter time and linearity correction."
+" The near-infrared images were also reduced following several steps, consisting of (i) dark subtraction, (ii) flat-field division, (iii) sky subtraction, and (iv) geometric alignment and combination of the dithered frames."," The near-infrared images were also reduced following several steps, consisting of (i) dark subtraction, (ii) flat-field division, (iii) sky subtraction, and (iv) geometric alignment and combination of the dithered frames."
+" Months after each SN faded, deep template images of their host galaxy were obtained under excellent seeing conditions."," Months after each SN faded, deep template images of their host galaxy were obtained under excellent seeing conditions."
+ Optical and near-infrared template imaging was performed with the du Pont telescope using Tek 5 and WIRC., Optical and near-infrared template imaging was performed with the du Pont telescope using Tek 5 and WIRC.
+" Following the method highlighted in Contrerasetal](2010),, the template images allow us to subtract away the host background light at the position of the SN in each science image."," Following the method highlighted in \citet{contreras10}, the template images allow us to subtract away the host background light at the position of the SN in each science image."
+ Observed magnitudes of each SN were computed differentially from the science frames with respect to a local sequenceILandoli of stars., Observed magnitudes of each SN were computed differentially from the science frames with respect to a local sequence of stars.
+ The photometric sequences were calibrated using (1992) (BV). (ugri) and [Perssonetal](1998) (YJH) fields observed over a minimum of three photometric nights.," The photometric sequences were calibrated using \citet{landolt92} $BV$ ), \citet{smith02} $ugri$ ) and \citet{persson98} $YJH$ ) fields observed over a minimum of three photometric nights."
+ The local sequences of SNe 2006V and 2006au in the standard system are provided in Table[I]., The local sequences of SNe 2006V and 2006au in the standard system are provided in Table \ref{tabstars}.
+" Nine epochs of optical spectroscopy were obtained for both SNe 2006V and 2006au with telescopes at LCO, and in one instance each, with the New Technology Telescope (NTT)."," Nine epochs of optical spectroscopy were obtained for both SNe 2006V and 2006au with telescopes at LCO, and in one instance each, with the New Technology Telescope (NTT)."
+ A journal of the spectroscopic observations is provided in Table D]., A journal of the spectroscopic observations is provided in Table \ref{tabspectra}.
+" Depending on the exact instrument used, the wavelength interval generally ranges from ~ 3800 tto 9500A."," Depending on the exact instrument used, the wavelength interval generally ranges from $\sim$ 3800 to 9500."
+. Standard reductions of each spectrum were performed as described in (2006)., Standard reductions of each spectrum were performed as described in \citet{hamuy06}.
+". Briefly, this consisted of overscan correction, bias subtraction and flat-fielding."," Briefly, this consisted of overscan correction, bias subtraction and flat-fielding."
+ The 2-D spectra were then optimally extracted and wavelength calibrated with respect to arc lamps., The 2-D spectra were then optimally extracted and wavelength calibrated with respect to arc lamps.
+" The wavelength corrected spectra were corrected for telluric absorption through the division of a telluric standard spectrum, and then flux-calibrated."," The wavelength corrected spectra were corrected for telluric absorption through the division of a telluric standard spectrum, and then flux-calibrated."
+ Multiple exposures of a particular object were then combined to produce a final high signal-to-noise science product., Multiple exposures of a particular object were then combined to produce a final high signal-to-noise science product.
+ This allowed for the removal of cosmic rays., This allowed for the removal of cosmic rays.
+ Optical and near-infrared light curves of SNe 2006V and 2006au are plotted in Figs., Optical and near-infrared light curves of SNe 2006V and 2006au are plotted in Figs.
+ B| and H]. respectively.," \ref{lc06V} and \ref{lc06au}, respectively."
+" The corresponding optical photometry — in the CSP natural system — is listed in Tables 3] and [B], whereas the final near-IR photometry in the Persson et al. ("," The corresponding optical photometry – in the CSP natural system – is listed in Tables \ref{tabphot06V} and \ref{tabphot06au}, whereas the final near-IR photometry in the Persson et al. ("
+1998) system is listed in Tables B] and [@.,1998) system is listed in Tables \ref{tabphotIR06V} and \ref{tabphotIR06au}.
+" The differences in measured magnitudes between the natural and the standard system are insignificant for the purpose of this study (Hamuyetal.|2006)., whereas they would be important for precision cosmology."," The differences in measured magnitudes between the natural and the standard system are insignificant for the purpose of this study \citep{hamuy06}, whereas they would be important for precision cosmology."
+" The transformation equations between the two systems are reported in (2006),, with updated color terms provided by (2011).."," The transformation equations between the two systems are reported in \citet{hamuy06}, , with updated color terms provided by \citet{stritzinger11}. ."
+" The light curves of SN 2006V follow the flux evolution from —50 to +75 days past B- maximum (Bmax, ie. JD= 2453823.7), while those of SN 2006au range from —59 to +33 days past Bnax (JD 2453865.5)."," The light curves of SN 2006V follow the flux evolution from $-50$ to $+75$ days past $B$ -band maximum $B_{\rm max}$ , i.e. $JD=2453823.7$ ), while those of SN 2006au range from $-59$ to $+33$ days past $B_{\rm max}$ $JD=2453865.5$ )."
+" The long rise to maximum, the broad peak and the subsequent decline to the radioactive tail are very similar to"," The long rise to maximum, the broad peak and the subsequent decline to the radioactive tail are very similar to"
+whereas for binaries the combination of initial orbital period and inital mass ratio must by be carefully arranged.,whereas for binaries the combination of initial orbital period and inital mass ratio must by be carefully arranged.
+ The derived rate depends on the adopted initial mass ratio distribution. the assumed binary fraction and on whether or not the companion star results in a core-collapse supernovae. see Table 5..," The derived rate depends on the adopted initial mass ratio distribution, the assumed binary fraction and on whether or not the companion star results in a core-collapse supernovae, see Table \ref{tab:rates}."
+ Even when we adopt the assumptions favoring a high SN IIb rate from binaries binaries. mass ratio distribution skewed to equal-mass systems and no core-collapse supernovae from the companions) the rate derived from our conservative models is still below the observed rate.," Even when we adopt the assumptions favoring a high SN IIb rate from binaries binaries, mass ratio distribution skewed to equal-mass systems and no core-collapse supernovae from the companions) the rate derived from our conservative models is still below the observed rate."
+ Another uncertainty to consider is the efficiency of mass transfer., Another uncertainty to consider is the efficiency of mass transfer.
+ In our models. lower efficiencies shift and widen the parameter space.," In our models, lower efficiencies shift and widen the parameter space."
+ In Table 5 we show that the number of SNe IIb over the number of core-collapse supernovae may increase to over1%.. with the warning that for a proper evaluation of the parameter space one needs a more extensive model grid.," In Table \ref{tab:rates} we show that the number of SNe IIb over the number of core-collapse supernovae may increase to over, with the warning that for a proper evaluation of the parameter space one needs a more extensive model grid."
+ A further uncertainty we have to consider is in the range of hydrogen-envelope masses in the progenitor that give rise to a SN IIb., A further uncertainty we have to consider is in the range of hydrogen-envelope masses in the progenitor that give rise to a SN IIb.
+ The lower limit of 0.1) is determined by the boundary between compact and extended type IIb., The lower limit of 0.1 is determined by the boundary between compact and extended type IIb.
+ The upper limit of about 0.5 is determined by the lack of a plateau in the light curve. which depends on the chemical structure of the stellar envelope and therefore is less certain.," The upper limit of about 0.5 is determined by the lack of a plateau in the light curve, which depends on the chemical structure of the stellar envelope and therefore is less certain."
+ If we change the boundary from 0.5 to 0.6M... test models show that the initial orbital period range will only widen by about 80 days.," If we change the boundary from 0.5 to 0.6, test models show that the initial orbital period range will only widen by about 80 days."
+ This means an increase of the overal rate of type IIb SNe of about10%., This means an increase of the overal rate of type IIb SNe of about.
+. The uncertainties 1 these boundaries will therefore not have a large effect on the rate., The uncertainties in these boundaries will therefore not have a large effect on the rate.
+ As an extreme assumption. we consider that all systems with orbital periods between 1000 and 2000 days produce IIb progenitors.," As an extreme assumption, we consider that all systems with orbital periods between 1000 and 2000 days produce IIb progenitors."
+ This corresponds to of all binary systems with initial mass ratios and orbital periods in the ranges specified earlier., This corresponds to of all binary systems with initial mass ratios and orbital periods in the ranges specified earlier.
+ This would increase the relative rate compared to the rate of core collapse supernova to2.3-3.5%.. depending on whether the companion star produces à core-collapse supernova or not and assuming a binary fraction of50%.," This would increase the relative rate compared to the rate of core collapse supernova to, depending on whether the companion star produces a core-collapse supernova or not and assuming a binary fraction of."
+. This number is consistent with the observed rate., This number is consistent with the observed rate.
+ Table 6 indicates the relative rates for the different characteristics of the companion at the moment of explosion. as described in section 5.1..," Table \ref{tab:rates2} indicates the relative rates for the different characteristics of the companion at the moment of explosion, as described in section \ref{reaction_companion}."
+ The large majority. about90%.. of type IIb supernova resulting from stable mass transfer is expected to have an O-star companion at the moment of explosion.," The large majority, about, of type IIb supernova resulting from stable mass transfer is expected to have an O-star companion at the moment of explosion."
+ The rate of systems with a blue supergiant companion. which applies to SN 1993J and possibly 200112. is predicted to be quite low according to our models. about of the type ΠΟ supernova rate.," The rate of systems with a blue supergiant companion, which applies to SN 1993J and possibly 2001ig, is predicted to be quite low according to our models, about of the type IIb supernova rate."
+ These relative rates reflect the large and very small regions in parameter space for Ostar and B-supergiant companions. respectively. as depicted in Fig. 3..," These relative rates reflect the large and very small regions in parameter space for O-star and B-supergiant companions, respectively, as depicted in Fig. \ref{Area}."
+ However. these percentages do not necessarily reflect the probability of a compantor of a certain type.," However, these percentages do not necessarily reflect the probability of a companion of a certain type."
+ Detecting an O-star companion in post-explosion images of a supernova is more difficult than detecting a B or K supergiant. as we demonstrated in Section 5.1..," Detecting an O-star companion in post-explosion images of a supernova is more difficult than detecting a B or K supergiant, as we demonstrated in Section \ref{reaction_companion}."
+ The probability for the presence of blue companions to type ΠΟ SNe we derive is strikingly lower than predicted by ?.., The probability for the presence of blue companions to type IIb SNe we derive is strikingly lower than predicted by \citet{Podsiadlowski92}.
+" They find that. if the companion starts to acerete after becoming a giant star. it will contract in a similar way to a feature known as “blue loops"" seen in some evolutionary tracks of single stars."," They find that, if the companion starts to accrete after becoming a giant star, it will contract in a similar way to a feature known as “blue loops” seen in some evolutionary tracks of single stars."
+ We do not find this behavior with our stellar evolution code., We do not find this behavior with our stellar evolution code.
+ Whether or not stellar models perform blue loops is very sensitive to the ratio of core mass to total stellar mass and to details in the chemical profile outside the stellar core (?).., Whether or not stellar models perform blue loops is very sensitive to the ratio of core mass to total stellar mass and to details in the chemical profile outside the stellar core \cite[]{Kippenhahn94}.
+ If a significant fraction of the case 3 models would result in blue companions. the relative rate for the presence of blue companions may increase by a factor of two or three (see Table 6)).," If a significant fraction of the case 3 models would result in blue companions, the relative rate for the presence of blue companions may increase by a factor of two or three (see Table \ref{tab:rates2}) )."
+ In the most common case. aceretion during the main sequence. the star rejuvenates. ie. the size of its convective core increases to adapt to the new stellar mass. see section 5.1.," In the most common case, accretion during the main sequence, the star rejuvenates, i.e. the size of its convective core increases to adapt to the new stellar mass, see section \ref{reaction_companion}."
+ In our code. all main sequence stars that accrete rejuvenate.," In our code, all main sequence stars that accrete rejuvenate."
+ However. whether stars rejuvenate or not depends on the assumed efficiency of semi-convection.," However, whether stars rejuvenate or not depends on the assumed efficiency of semi-convection."
+ Stars that have evolved towards the enc of the mail sequence build up a chemical gradient between the helium rich core and the hydrogen rich envelope., Stars that have evolved towards the end of the main sequence build up a chemical gradient between the helium rich core and the hydrogen rich envelope.
+ ? show that. assuming a low efficiency of semi-convective mixing. the growth of the convective core in an accreting main-sequence star is prevented.," \citet{Braun95} show that, assuming a low efficiency of semi-convective mixing, the growth of the convective core in an accreting main-sequence star is prevented."
+ These stars may also appear as blue supergiants at the moment of explosion., These stars may also appear as blue supergiants at the moment of explosion.
+ Based on their models. we estimate that in our case of conservative mass transfer the companion must have reached a central helium mass fraction of about 0.75 or more to prevent rejuvenation.," Based on their models, we estimate that in our case of conservative mass transfer the companion must have reached a central helium mass fraction of about 0.75 or more to prevent rejuvenation."
+ In our models this occurs for initial mass ratios larger than about 0.87. and we estimate that this may increase the relative rate of the presence of blue companions up to about50%.," In our models this occurs for initial mass ratios larger than about 0.87, and we estimate that this may increase the relative rate of the presence of blue companions up to about."
+. For lower accretion efficiencies it becomes even easier to prevent rejuvenation: a central helium mass fraction for the acereting star of just 0.6 and higher are required if B0., For lower accretion efficiencies it becomes even easier to prevent rejuvenation: a central helium mass fraction for the accreting star of just 0.6 and higher are required if $\beta=0$.
+ However. the companion must accrete enough to finish its main sequence evolution before the explosion of the primary star. to be observed as a blue supergiant.," However, the companion must accrete enough to finish its main sequence evolution before the explosion of the primary star, to be observed as a blue supergiant."
+ More detailed estimates are beyond the scope of this study. but we do emphasize this as an interesting possibility to use the companions of supernovae to gain insight in internal mixing processes.," More detailed estimates are beyond the scope of this study, but we do emphasize this as an interesting possibility to use the companions of supernovae to gain insight in internal mixing processes."
+ We identified binary progenitor models for extended type IIb supernovae., We identified binary progenitor models for extended type IIb supernovae.
+ In these models the most massive star is stripped from its envelope by interaction with its companion. such that only several tenths of solar masses of hydrogen remain at the moment of explosion.," In these models the most massive star is stripped from its envelope by interaction with its companion, such that only several tenths of solar masses of hydrogen remain at the moment of explosion."
+ We find that the most massive star must fill its Roche lobe during central helium burning 1n order to achieve this., We find that the most massive star must fill its Roche lobe during central helium burning in order to achieve this.
+ We derive for which range of initial orbital periods and mitial mass ratios binary systems are expected to produce type IIb supernova progenitors., We derive for which range of initial orbital periods and initial mass ratios binary systems are expected to produce type IIb supernova progenitors.
+ We have discussed in detail the properties of the companion star at the moment of explosion. motivated by the detection of a companion for SN 19951. a possible companion for 200115 and the non-detection of a companion in the supernova remnant," We have discussed in detail the properties of the companion star at the moment of explosion, motivated by the detection of a companion for SN 1993J, a possible companion for 2001ig and the non-detection of a companion in the supernova remnant"
+and terminal-age IIB (ZAIID and TAUB. respectively) from Aloehleretal.(2003).. for canonical (¥=0.23) and Ile-enriched (Y=0.33) models.,"and terminal-age HB (ZAHB and TAHB, respectively) from \citet{Moehler03}, for canonical (Y=0.23) and He-enriched (Y=0.33) models."
+ In the same figure we include 73 stars [rom Moehlleretal.(2011)... who measured the atmospheric parameters with (he same procedure as in the present work. but based on mecdium-resolution FLAMES spectra. and a different set of model spectra for stars above 200000 Ix. The two sets of data behave very similarly in (he Tey-loe (gy) plane.," In the same figure we include 78 stars from \citet{Moehler11}, who measured the atmospheric parameters with the same procedure as in the present work, but based on medium-resolution FLAMES spectra, and a different set of model spectra for stars above 000 K. The two sets of data behave very similarly in the $_\mathrm{eff}$ $\log{\mathrm (g)}$ plane."
+ The comparison of the eieht stars in common confirms the good agreement between the two works: the mean difference in gravity is null (<0.01 dex). while the difference in temperature is small (155 Ix). ancl becomes negligible (25 Ix) alter the exclusion of (wo stars with verv large errors (1500 Ix).," The comparison of the eight stars in common confirms the good agreement between the two works: the mean difference in gravity is null $\leq0.01$ dex), while the difference in temperature is small (155 K), and becomes negligible (25 K) after the exclusion of two stars with very large errors $\geq$ 1500 K)."
+ Our mass estimates are on average higher by 0.035 MN... an offset accounted for by the fainter magnitudes (0.09 mag ihe mean difference) Irom the Castellaniοἱal.(2007) catalog used bv Moehler et al.," Our mass estimates are on average higher by 0.035 $_\sun$, an offset accounted for by the fainter magnitudes (0.09 mag the mean difference) from the \citet{Castellani07} catalog used by Moehler et al."
+ Given the excellent agreement. we will merge (he two datasets together.," Given the excellent agreement, we will merge the two datasets together."
+ The surface gravity of stars cooler than ~18 0000 Ix. is svstematically lower by about 0.2-0.3 dex with respect to canonical models. while (μον closely follow the trend. of models at all temperatures.," The surface gravity of stars cooler than $\sim$ 000 K is systematically lower by about 0.2-0.3 dex with respect to canonical models, while they closely follow the trend of He-enhanced models at all temperatures."
+ In the lower panel of Figure 1. we compare these results with similar measurements obtained in three GCs. namely 66752 (Monietal. 2007).. M80 and 55986 (MoniBiclinetal.2009).," In the lower panel of Figure \ref{f_tg} we compare these results with similar measurements obtained in three GCs, namely 6752 \citep{Moni07}, M80 and 5986 \citep{Moni09}."
+. We adopt as vertical coordinate ihe difference between the measured log(g) and the value of the canonical ZAIID at the corresponding temperature., We adopt as vertical coordinate the difference between the measured $\log{\mathrm (g)}$ and the value of the canonical ZAHB at the corresponding temperature.
+ The comparison reveals that the IIB stars in oí Cen and in the other GCs behave very differently., The comparison reveals that the HB stars in $\omega$ Cen and in the other GCs behave very differently.
+ We note that the stars of the comparison clusters do nol completely agree with canonical models. whose expectation is to find the majority of the objects next to the ZAIID. and not to the TAIID as observed.," We note that the stars of the comparison clusters do not completely agree with canonical models, whose expectation is to find the majority of the objects next to the ZAHB, and not to the TAHB as observed."
+ Even so. ω CCen stars clearly show lower eravilies al a given effective temperature will respect to stus in other GCs.," Even so, $\omega$ Cen stars clearly show lower gravities – at a given effective temperature – with respect to stars in other GCs."
+ Observational errors tend (o mask the general trend. but there is an offset of 0.15 dex at the cooler end. which decreases at higher temperatures ancl fades out around. 0000 Ix. (o reappear even larger (20.2 dex) among hot stars al 0000. Ix. However. we find a problem in the estimate of stellar masses. summarized in (he upper panel of Figure 2: while the results in the comparison clusters roughly agree wilh expectations (seeMoniDidin2007.2009.[oracomplete discussion)... the masses of ο Cen stars are constantly underestimated at all temperatures.," Observational errors tend to mask the general trend, but there is an offset of 0.15 dex at the cooler end, which decreases at higher temperatures and fades out around 000 K, to reappear even larger $\geq$ 0.2 dex) among hot stars at 000 K. However, we find a problem in the estimate of stellar masses, summarized in the upper panel of Figure \ref{f_masabs}: while the results in the comparison clusters roughly agree with expectations \citep[see][for a complete discussion]{Moni07,Moni09}, the masses of $\omega$ Cen stars are constantly underestimated at all temperatures."
+ Interestingly enough. Moehler&Sweigart(2006). found very simular results for WB stus in 66338.," Interestingly enough, \citet{Moehler06} found very similar results for HB stars in 6388."
+ This is a very. peculiar cluster. as it has an extended. blue HB (Richetal.1997). despite its high metallicity —0.6).," This is a very peculiar cluster, as it has an extended, blue HB \citep{Rich97} despite its high metallicity $-$ 0.6)."
+ This has been interpreted in terms of an extreme IIe-enrichment 200T7).. as in (he case of its twin cluster. 66441.," This has been interpreted in terms of an extreme He-enrichment \citep[up to Y=0.40,][]{Caloi07}, as in the case of its twin cluster, 6441."
+ While Moehler&Sweigart(2006) cast doubts on their results due to the large uncertainties caused by stellar crowcling in this compact. bulge cluster. these results are very similar to what we find now in v CCen.," While \citet{Moehler06} cast doubts on their results due to the large uncertainties caused by stellar crowding in this compact, bulge cluster, these results are very similar to what we find now in $\omega$ Cen."
+scattering terii in the reflection: the iucideut radiatiou heats up the planet and is always re-cinitted locally.,scattering term in the reflection: the incident radiation heats up the planet and is always re-emitted locally.
+" Scattered light is implicitly accounted for by allowing the global temperature T, to be de-coupled from the dav-de teniperature.", Scattered light is implicitly accounted for by allowing the global temperature $T_{p}$ to be de-coupled from the day-side temperature.
+ The ELC inodoel provides a good fit to the observations., The ELC model provides a good fit to the observations.
+ The chi square is 57.225 for 19.016 degrees of freedom (reduced 2-167).," The chi square is 57,225 for 49,016 degrees of freedom (reduced $\chi^{2}_{\nu}$ =1.167)."
+ Uncertainties ou the parameters were boosted by a factor of VA )) to account for the formally. ligh. 9χω. which. we attribute. to systematic. nou-CGatssian noise.," Uncertainties on the parameters were boosted by a factor of $\sqrt{\chi^{2}_{\nu}}$ ) to account for the formally high $\chi^{2}_{\nu}$, which we attribute to systematic non-Gaussian noise."
+ The values of the parameters are listed in Table 4, The values of the parameters are listed in Table 1.
+" The L8 Alig, planet orbiting only L1 stellar radii from its host star induces a tidal distortion on the star. changing its shape from oblate (not spherical. due to its rotation) to a more triaxial shape. with the longest axis along the direction toward the planet aud the shortest axis perpendicular to the orbital This shape causes the well-known ""ellipsoidal variation” effect secu in binary stars: a modulation in light at half the orbital period. with maxima at phases near 0.25 and 0.75. aud nuequal minima at phases 0.0 and 0.5."," The 1.8 $\Mjup$ planet orbiting only 4.1 stellar radii from its host star induces a tidal distortion on the star, changing its shape from oblate (not spherical, due to its rotation) to a more triaxial shape, with the longest axis along the direction toward the planet and the shortest axis perpendicular to the orbital This shape causes the well-known “ellipsoidal variation” effect seen in binary stars: a modulation in light at half the orbital period, with maxima at phases near 0.25 and 0.75, and unequal minima at phases 0.0 and 0.5."
+ The ellipsoidal varlation is primarily a ecometrical (projected surface area) effect whose relative amplitude depends ou the lnass ratio and the inclination of the binary svsteiu., The ellipsoidal variation is primarily a geometrical (projected surface area) effect whose relative amplitude depends on the mass ratio and the inclination of the binary system.
+ Given the tight constraint on the inclination from the eclipses. they imav provide some limits on the mass ratio. and hence the mass of the planet.," Given the tight constraint on the inclination from the eclipses, they may provide some limits on the mass ratio, and hence the mass of the planet."
+ In the optical. where the phase-depeudeut planct-to-star flux ratio is sxnall. the presence of the stella ellipsoidal variation beconies significant.," In the optical, where the phase-dependent planet-to-star flux ratio is small, the presence of the stellar ellipsoidal variation becomes significant."
+ Neglecting its contribution can lead to a coufused interpretation of the phase curve aud thus au incorrect measurement of the albedo aud phliase-dependent scatteredthermal enission., Neglecting its contribution can lead to a confused interpretation of the phase curve and thus an incorrect measurement of the albedo and phase-dependent scattered/thermal emission.
+ The presence of ellipsoidal variations im cxoplanctary systems was anticipated by Loeb&Candi(2003). and Drake(2003).. and Pfaliletal.(2008) preseut a detailed theoretical investigation of the tidal force on the star by the planet.," The presence of ellipsoidal variations in exoplanetary systems was anticipated by \citet{Loeb03} and \citet{Drake03}, and \citet{Pfahl08} present a detailed theoretical investigation of the tidal force on the star by the planet."
+ Figs., Figs.
+ 1 and 3 show the first detection of the effect in an exoplanct svstem., 1 and 3 show the first detection of the effect in an exoplanet system.
+ In Fig., In Fig.
+ 3 we show the light curve binned to 30 min. along with the best-fit model decomposition.," 3 we show the light curve binned to 30 min, along with the best-fit model decomposition."
+ The dotted curve shows the stellaz-ouly elpsoidal light curve while the dashed curve i9 the Xauet-oulv light curve (offset vertically)., The dotted curve shows the stellar-only ellipsoidal light curve while the dashed curve is the planet-only light curve (offset vertically).
+ The solid curve is the stun of the two and equals our best-fit model., The solid curve is the sum of the two and equals our best-fit model.
+ The Xauet-oulv model can match the Πο curve ouly very rear phase 0.5: if is a very poor fit to the data at other ghases. indicating the need for the ellipsoidal variation conrponeut.," The planet-only model can match the light curve only very near phase 0.5; it is a very poor fit to the data at other phases, indicating the need for the ellipsoidal variation component."
+ The amplitude of the ellipsoidal component is 37.3 ppiu. detectable only because ofs hieh-wecision photometric capability.," The amplitude of the ellipsoidal component is 37.3 ppm, detectable only because of high-precision photometric capability."
+ Ellipsoidal variations arise as a consequence of eravity on a luminous fluid body., Ellipsoidal variations arise as a consequence of gravity on a luminous fluid body.
+ Within the Roche framework. its amplitude is exactly known if the mass ratio. inclination. and stellar radius are known.," Within the Roche framework, its amplitude is exactly known if the mass ratio, inclination, and stellar radius are known."
+ However. since iu exoplanet systems the star is not expected to be ij svuchronous rotation. the Roche poteutial is au approximation. albeit a good one since the maxima idal distortion. AR/R is ouly 10.1.," However, since in exoplanet systems the star is not expected to be in synchronous rotation, the Roche potential is an approximation, albeit a good one since the maximum tidal distortion $\Delta R/R$ is only $10^{-4}$."
+ However. we stress hat the model presented here is only a starting point. sed on the well-developed aud successtul Roche model.," However, we stress that the model presented here is only a starting point, based on the well-developed and successful Roche model."
+ Effects such as those discussed. in Pfahletal.(2008) based ou the equilibrium tide approximation can ο present., Effects such as those discussed in \citet{Pfahl08} based on the equilibrium tide approximation can be present.
+ As moreKepler data become available it will be interesting to try to distinguish between these approximations. as they can lead to measuring internal xoperties of the star.," As more data become available it will be interesting to try to distinguish between these approximations, as they can lead to measuring internal properties of the star."
+ For example. in the frame of he rotating binary. the star is spimuine retrograde if <Por. Which is expected to be true for most Pepinrot Jupiter planets.," For example, in the frame of the rotating binary, the star is spinning retrograde if $P_{spin} < P_{orb}$, which is expected to be true for most hot Jupiter planets."
+" This could induce a phase shift iu he ellipsoidal variations. as the stars spin ""drags the ide away from the line of centers."," This could induce a phase shift in the ellipsoidal variations, as the star's spin “drags” the tide away from the line of centers."
+ If a phasc-lag not associated with the planct light can be unambiguously neasured. it may be possible to put constraints ou he tidal OQ factor.," If a phase-lag not associated with the planet light can be unambiguously measured, it may be possible to put constraints on the tidal Q factor."
+ This asvnunetry could also lead to interesting orbital dynamics effects., This asymmetry could also lead to interesting orbital dynamics effects.
+ The detailed shape and amplitude of the ellipsoidal variation also depeucds ou the stellar cuvelope (convective vs. radiative. ίαetal.(2008))). thus potentially offering a probe of stellar interiors.," The detailed shape and amplitude of the ellipsoidal variation also depends on the stellar envelope (convective vs. radiative – \citet{Pfahl08}) ), thus potentially offering a probe of stellar interiors."
+ Towever. for ILAT-P-7. the situation is coniplicated by the fact that the spin axis is not aligned with the orbit axis (Winnetal.2009).. so these optimistic statements must be tempered with caution.," However, for HAT-P-7, the situation is complicated by the fact that the spin axis is not aligned with the orbit axis \citep{Winn09}, so these optimistic statements must be tempered with caution."
+ Fie., Fig.
+ 2 shows model light curves that illustrate the contributions of the star to the optical light., 2 shows model light curves that illustrate the contributions of the star to the optical light.
+ The lower curve is the stellar light only. showine the ellipsoidal variations.," The lower curve is the stellar light only, showing the ellipsoidal variations."
+ If is not coustant outside of transit. and its auplitude can be significant compared to the depth of the occultation.," It is not constant outside of transit, and its amplitude can be significant compared to the depth of the occultation."
+ Above this ave four curves that include the planet contribution., Above this are four curves that include the planet contribution.
+ The lowest is for au isothermal isotropic planet with no plase-dependent scattered or re-racliated enassiou., The lowest is for an isothermal isotropic planet with no phase-dependent scattered or re-radiated emission.
+ To first order. it is essentially au additive offset to the stars liebt. the amount depending ou the relative radius aud temperature of the planet.," To first order, it is essentially an additive offset to the star's light, the amount depending on the relative radius and temperature of the planet."
+ If the absorbed incident cucrev from the star is nof vortectly unifonulv re-distributed around the planet. the day side of the planet will be hotter than the might side. resulting in an orbit-phase dependent modulation »eakiug at the sub-stellay point (phase 0.5).," If the absorbed incident energy from the star is not perfectly uniformly re-distributed around the planet, the day side of the planet will be hotter than the night side, resulting in an orbit-phase dependent modulation peaking at the sub-stellar point (phase 0.5)."
+ Or. if the atmosphere scatters the iucident radiatiou. again the dav side will be brighter than the might side (though not recessarily hotter).," Or, if the atmosphere scatters the incident radiation, again the day side will be brighter than the night side (though not necessarily hotter)."
+ The upper three curves show this shase-cependcut effect for different values of the heat albedo. which is a proxy for scattered and re-radiate ulssjon on the dav side.," The upper three curves show this phase-dependent effect for different values of the heat albedo, which is a proxy for scattered and re-radiated emission on the day side."
+ Note that advection can move i6 peak dowmwind aud produce a phase shift iu the lanet's enmüssiou. as seen in the infrared. phase curve of ΠΟΣΟτο (I&nutsonetal.2007): this is not included in ese mnodels.," Note that advection can move the peak downwind and produce a phase shift in the planet's emission, as seen in the infrared phase curve of HD189733 \citep{Knutson07}; this is not included in these models."
+ Given a stellar tempcrature estimate. the ELC mode can vield dav and might hemisphere tempcratures.," Given a stellar temperature estimate, the ELC model can yield day and night hemisphere temperatures."
+ Because the planet is uot black it contributes light at al yhases (other than occultation) and its fíux coutribution at each orbital phase. relative to ellipsoidal aud other effects. allows its temperature to be estimated.," Because the planet is not black it contributes light at all phases (other than occultation) and its flux contribution at each orbital phase, relative to ellipsoidal and other effects, allows its temperature to be estimated."
+ Tn articular. the ability to measure relative teniperatures arises from the requirciment of ectting the occultation and transit depths correct given the tight geometric constraints (c.g. a 2600 Is pluiet reduces the 6000 transito depth by 29 ppm compared to a zero temperature plauct).," In particular, the ability to measure relative temperatures arises from the requirement of getting the occultation and transit depths correct given the tight geometric constraints (e.g. a 2600 K planet reduces the 6000 transit depth by 29 ppm compared to a zero temperature planet)."
+ We fud the peak temperature on the planet (at the sub-stellar poiut) is 3160 Ik. but a," We find the peak temperature on the planet (at the sub-stellar point) is 3160 K, but a"
+it iuto halves.,it into halves.
+ Decorrelated CoBoT-1 data aud the best-fit eclipse are shown in the middle panel of Figure 1. aud are binned (to 100 bins) in the bottom panel of Figure 1.," Decorrelated CoRoT-1 data and the best-fit eclipse are shown in the middle panel of Figure 1, and are binned (to 100 bins) in the bottom panel of Figure 1."
+ We use a nearly identical procedure to fit the 4 eclipse of CoRoT-2. shown in Figure 3. except that we do not break the decorrelation at nud-eclipse.," We use a nearly identical procedure to fit the $\mu$ m eclipse of CoRoT-2, shown in Figure 3, except that we do not break the decorrelation at mid-eclipse."
+ The first 30 uunutes of these data (not illustrated in Figure 3) exhibit a transient decrease m fiux. simular to the raz effect seen at longer wavelength. but decreasing instead of increasing. and uot correlated with the inge position ou the detector.," The first $\sim 30$ minutes of these data (not illustrated in Figure 3) exhibit a transient decrease in flux, similar to the ramp effect seen at longer wavelength, but decreasing instead of increasing, and not correlated with the image position on the detector."
+ Transient effects at this wavelength are not well understood. so we simply omit the 19 data cubes prior to orbital phase 0.11.," Transient effects at this wavelength are not well understood, so we simply omit the 19 data cubes prior to orbital phase 0.41."
+ Another difference for CoRoT-2 is that a correction is needed for diffracted πο from an Aechwart wine {απο sec distaut (Cullonetal.2010)., Another difference for CoRoT-2 is that a correction is needed for diffracted light from an M-dwarf lying 4 arc-sec distant \citep{gillon10}.
+". Since we also re-analyze archival data at 1.5 and Sgan for CoRoT-2 (see below). we need to estimate the diffracted light contributed by the A-dwarf in the CoRoT-2 aperture at 3.6. 15. aud μια, We calculated the flux ratio (M-dwarf to CoRoT-2) in the IRAC lauds. using the the flux estimation tool (STAR-PET) ou the Spitzer website. aud the 2421ASS K-naenitudes aud J-I& colors of the two stars."," Since we also re-analyze archival data at 4.5 and $\mu$ m for CoRoT-2 (see below), we need to estimate the diffracted light contributed by the M-dwarf in the CoRoT-2 aperture at 3.6, 4.5, and $\mu$ m. We calculated the flux ratio (M-dwarf to CoRoT-2) in the IRAC bands, using the the flux estimation tool (STAR-PET) on the Spitzer website, and the 2MASS K-magnitudes and J-K colors of the two stars."
+ ποπας their relative brightuess. we also need to know the fraction of he Me-dwarf fux that is diffracted iuto the photometiy aperture for CoRoT-2.," Knowing their relative brightness, we also need to know the fraction of the M-dwarf flux that is diffracted into the photometry aperture for CoRoT-2."
+ We estinated this by placing he aperture at a svuuuetiic location on the other side of CoRoT-2. where the diffracted Leht is coutributed aliiost exclusively by CoRoT-2 itself.," We estimated this by placing the aperture at a symmetric location on the other side of CoRoT-2, where the diffracted light is contributed almost exclusively by CoRoT-2 itself."
+ Using that diffracted yaction together with the flux ratio of M-dwarf to CoRoT-2. we inter that the diffracted light. from the. A\L-dwart contributes5.0%.. aud to CoRoT-2 at 3.6. 1.5 aud 8.0422. respectively.," Using that diffracted fraction together with the flux ratio of M-dwarf to CoRoT-2, we infer that the diffracted light from the M-dwarf contributes, and to CoRoT-2 at 3.6, 4.5 and $\mu$ m, respectively."
+ The eclipse photometry aud derived parameters forCoRoT-2 in our Figures and Tables have all been corrected for this diffracted light., The eclipse photometry and derived parameters forCoRoT-2 in our Figures and Tables have all been corrected for this diffracted light.
+" Cüllonal.(2010). interred aud at L5 and 8402. respectively, but he used aperture radii of [O0 and 3.5 pixels. respectively. 3.0 pixels in our case."," \citet{gillon10} inferred and at 4.5 and $\mu$ m, respectively, but he used aperture radii of 4.0 and 3.5 pixels, respectively, 3.0 pixels in our case."
+ As a check. we repeated our diffracted light correction using apertures having the same size as Cüllonetal. (2010).," As a check, we repeated our diffracted light correction using apertures having the same size as \citet{gillon10}."
+.. Because the diffracted Πο is not unifon. the values do not simply scale as the area of the aperture.," Because the diffracted light is not uniform, the values do not simply scale as the area of the aperture."
+ For the same apertures as Callonetal.(201 we obtain corrections of and at L5 and jan. respectively. iu reasonably good agreenmeut with the independent determination of (οιetal.(2010).," For the same apertures as \citet{gillon10}, , we obtain corrections of and at 4.5 and $\mu$ m, respectively, in reasonably good agreement with the independent determination of \citet{gillon10}."
+. Uncertaüntv in the diffracted Πο correction is nof included in our eclipse amplitude error estimates., Uncertainty in the diffracted light correction is not included in our eclipse amplitude error estimates.
+ Cave our eood agreement with the diffracted Πο corrections of Callonetal.(2010).. aud. given that we use smaller photometric apertures than (οι.etal.(2010).. we conchide that uncertainty in the diffracted light correction does not contribute siguificauth to the errors on our measured eclipse depths.," Given our good agreement with the diffracted light corrections of \citet{gillon10}, and given that we use smaller photometric apertures than \citet{gillon10}, we conclude that uncertainty in the diffracted light correction does not contribute significantly to the errors on our measured eclipse depths."
+ The best-fit eclipse depths and errors are listed in Table 1. aud the ceutral phases aud errors are listed iu Table 2.," The best-fit eclipse depths and errors are listed in Table 1, and the central phases and errors are listed in Table 2."
+ The ideal method to calculate crrors would be to repeat all of the observations and analysis. and compare he results from analvziug many independent sets of observations.," The ideal method to calculate errors would be to repeat all of the observations and analysis, and compare the results from analyzing many independent sets of observations."
+ This is obviously impractical. so we unnuic sole kev aspects of that ideal procedure.," This is obviously impractical, so we mimic some key aspects of that ideal procedure."
+ We generate ake photomoetnc datasets having the same properties as he real photometry. aud we repeat the cutive iterative fitting process - includiug iutra-pixel corrections aud ran fitting - on cach fake dataset.," We generate fake photometric datasets having the same properties as the real photometry, and we repeat the entire iterative fitting process - including intra-pixel corrections and ramp fitting - on each fake dataset."
+ We calculate the standard deviation of the collection of eclipse depths and central phases resulting from the repetitions of the analysis ou he fake data., We calculate the standard deviation of the collection of eclipse depths and central phases resulting from the repetitions of the analysis on the fake data.
+ To egcuerate cach fake dataset. we subtract the best-fit eclipse curve (plus baseline aud. iutra-pixel decorrelation Wuction) from the original photometry to produce a set of plotometric residuals.," To generate each fake dataset, we subtract the best-fit eclipse curve (plus baseline and intra-pixel decorrelation function) from the original photometry to produce a set of photometric residuals."
+ We likewise produce a set of image position residuals bv subtracting a iulti-point ruunius average of the XN aud. Y-pixel positions from cach individual (X.Y) position measurement.," We likewise produce a set of image position residuals by subtracting a multi-point running average of the X and Y-pixel positions from each individual (X,Y) position measurement."
+ We permite all of he residuals and add them back to the best-fit fiction (photometry) or running average coordinate (position) to nake an individual fake dataset., We permute all of the residuals and add them back to the best-fit function (photometry) or running average coordinate (position) to make an individual fake dataset.
+ We permite the residuals using two methods. to make two distinct collections of ake data.," We permute the residuals using two methods, to make two distinct collections of fake data."
+" The first permutation method scrambles the residuals raudomly, which is equivalent to the conventional )otstrap Afoute Carlo technique (Pressctal.1992)."," The first permutation method scrambles the residuals randomly, which is equivalent to the conventional bootstrap Monte Carlo technique \citep{press}."
+ We generate 101 bootstrap datasets (trials) using this uethod. aud calculate the standard deviation of eclipse depth and central phase frou the distributions of these xuümnmeters over the 104 trials.," We generate $10^4$ bootstrap datasets (trials) using this method, and calculate the standard deviation of eclipse depth and central phase from the distributions of these parameters over the $10^4$ trials."
+ These distributions are close to Cassia., These distributions are close to Gaussian.
+" A second method to permute the residuals preserves hei relative order but shifts their initial phase: this is sometimes called the ""praver-bead! method (Cillonetal.2009b).", A second method to permute the residuals preserves their relative order but shifts their initial phase; this is sometimes called the `prayer-bead' method \citep{gillon09b}.
+. Iu this case. the number of trials equals he number of originalft photometry poiuts.," In this case, the number of trials equals the number of original photometry points."
+ This is Sse or CoBRoT-1. and 13.515 for version 2 of the CoRoT-2 subarrav photometry.," This is 888 for CoRoT-1, and 13,545 for version 2 of the CoRoT-2 subarray photometry."
+ These are adequate to define the distributions of eclipse depth aud phase., These are adequate to define the distributions of eclipse depth and phase.
+ The praver-bead ucthod is more scusitive to the presence of red noise in the data., The prayer-bead method is more sensitive to the presence of red noise in the data.
+ Nevertheless we find that the distribution of eclipse depth remains consistent with a Gaussian. but for CoRoT- the distribution of eclipse ράσο shows about 7% of the central phases lie below the 30 poiut in the distribution.," Nevertheless we find that the distribution of eclipse depth remains consistent with a Gaussian, but for CoRoT-1 the distribution of eclipse phase shows about $7\%$ of the central phases lie below the $3\sigma$ point in the distribution."
+ We attribute this to the presence of some red noise before wic-eclipse. visible iu the bottom panel of Figure 1.," We attribute this to the presence of some red noise before mid-eclipse, visible in the bottom panel of Figure 1."
+ For CoRoT-2. the distributions of eclipse depth aud ghase were close to Gaussian. so errors from) the prayer-)'ad method were quite close to the values from the vootstrap method.," For CoRoT-2, the distributions of eclipse depth and phase were close to Gaussian, so errors from the prayer-bead method were quite close to the values from the bootstrap method."
+ This indicates relatively little red noise iu the CoRoT-2 data (after we omitted the first 19 data yoluts. as noted above)," This indicates relatively little red noise in the CoRoT-2 data (after we omitted the first 19 data points, as noted above)."
+ For both CoRoT-1 aud -2. we adopted the greater of the bootstrap and praver-bead errors for each pariüneter.," For both CoRoT-1 and -2, we adopted the greater of the bootstrap and prayer-bead errors for each parameter."
+ CoRoT-1 errors are uniformly arecr than for CoRoT-2 because the star is fainter aud the red noise is sliglitlv exeater., CoRoT-1 errors are uniformly larger than for CoRoT-2 because the star is fainter and the red noise is slightly greater.
+ Tables 1 2 list the errors on eclipse depth aud. central phase for all three eclipses. plus our results from re-analysis of CoRoT-2 at 5 aud aun (sce below).," Tables 1 2 list the errors on eclipse depth and central phase for all three eclipses, plus our results from re-analysis of CoRoT-2 at 4.5 and $\mu$ m (see below)."
+ We check our methodology by analyzing archivalSpitzer data for CoRoT-2 at L5- and µια. for comparison to the results of Cillonetal. (2010)..," We check our methodology by analyzing archivalSpitzer data for CoRoT-2 at 4.5- and $\mu$ m, for comparison to the results of \citet{gillon10}. ."
+ Our analysis at pau, Our analysis at $\mu$ m
+survey having a unique combination of depth and area coverage.,survey having a unique combination of depth and area coverage.
+ The additional V-banud tages. which altogether overlap the £ images over ~ of the surveyed area in vatches A and D of the whole EIS-wicde). are also useful in the search of galaxy clusters Olsen 1998. paper V).," The additional $V$ -band images, which altogether overlap the $I$ images over $\sim$ of the surveyed area in patches A and B of the whole EIS-wide), are also useful in the search of galaxy clusters Olsen 1998, paper V)."
+ Finally. it is iuportaut to enipliasize that he preseut data offer an excellent opportunity to assess he conmplexitv of cficiently handling large. volumes of nulticolor data anc extracting useful target lists.," Finally, it is important to emphasize that the present data offer an excellent opportunity to assess the complexity of efficiently handling large volumes of multicolor data and extracting useful target lists."
+ This is a sev Clete for the exploration of the full range of science offered by the iuulticolor surveys envisioned for the new wide-field camera (NFETo2.211) at the ESO/AIPIA 2.21 clescope a La Silla., This is a key element for the exploration of the full range of science offered by the multicolor surveys envisioned for the new wide-field camera (WFI2.2m) at the ESO/MPIA 2.2m telescope at La Silla.
+ Iu section 2. he observations are described aud the characteristics of he unulticolor data are presented.," In section 2, the observations are described and the characteristics of the multicolor data are presented."
+ This section also prescuts the extracted object catalogs and discusses thei completeness and reliability., This section also presents the extracted object catalogs and discusses their completeness and reliability.
+ Iu section 3 the catalogs are evaluated by comparison with models aud other data., In section 3 the catalogs are evaluated by comparison with models and other data.
+ Concluding remarks are preseuted im section L., Concluding remarks are presented in section 4.
+ The observations of patch D were carried out over several mouths in the period July 1997 to Deceiiber 1997. using the red channel of the EMMI camera on the 2.510. New Technology Telescope (NTT) at La Silla.," The observations of patch B were carried out over several months in the period July 1997 to December 1997, using the red channel of the EMMI camera on the 3.5m New Technology Telescope (NTT) at La Silla."
+ The red chanel of EMMZT is equipped with a Tektronix 2016 « 2016 chip with a pixel size of 0.266 arcsec aud a useful field-ofview of about 9«8.5., The red channel of EMMI is equipped with a Tektronix 2046 $\times$ 2046 chip with a pixel size of 0.266 arcsec and a useful field-of-view of about $9' \times 8.5'$.
+ EIS uses a special set of BVI filters aud the response function of the system cau be found iu oper I. As described in paper I. the observations were carried out bv a sequence of overlapping exposures (hereafter refer ο as ονομαο frames) 150 sec cach. with cach position on ie sky being sampled at least twice.," EIS uses a special set of BVI filters and the response function of the system can be found in paper I. As described in paper I, the observations were carried out by a sequence of overlapping exposures (hereafter refer to as even/odd frames) 150 sec each, with each position on the sky being sampled at least twice."
+ A total of 701 frames were obtained in the area with 200 in D. 282 in V and 219 oe1i I bands.," A total of 701 frames were obtained in the area with 200 in $B$, 282 in V and 219 in I bands."
+ Ouly 150 frames in cach band were required ο cover the field mt poor weather conditions required several frauues to ο re-observed., Only 150 frames in each band were required to cover the field but poor weather conditions required several frames to be re-observed.
+ The stroug variations in the observing conditious can be seen in Figure 1 which shows. or cach band. the seeing distribution of all observed frames.," The strong variations in the observing conditions can be seen in Figure \ref{fig:seeinghistall}
+ which shows, for each band, the seeing distribution of all observed frames."
+ For comparison he shaded histograms show the seciug distribution of the frames finally accepted. with the solid vertical line in cach panel indicating the necdian secine for each xud.," For comparison the shaded histograms show the seeing distribution of the frames finally accepted, with the solid vertical line in each panel indicating the median seeing for each band."
+ The B-baud is he worst overall with a median seem of 2 arcsec and with a ew frames extending to very huge seeing (~2.5 arcsec)., The $B$ -band is the worst overall with a median seeing of 1.2 arcsec and with a few frames extending to very large seeing $\sim 2.5$ arcsec).
+ Iu the analysis below 9 frames with secing 21.8 arcsec were discarded because of their incompleteness at faint uaeuitudes., In the analysis below 9 frames with seeing $\gsim 1.8$ arcsec were discarded because of their incompleteness at faint magnitudes.
+ Fieure 2. shows. again for cach baud. the lo initiug isophote within 1 arcsec.," Figure \ref{fig:mlim} shows, again for each band, the $1\sigma$ limiting isophote within 1 arcsec."
+ The transparency of the Helits also showed significant variations especially for the I-haud nuages. with one frame reaching Qs21 /arcsecr (not shown iu the Ποιο).," The transparency of the nights also showed significant variations especially for the $I$ -band images, with one frame reaching 21 $^2$ (not shown in the figure)."
+C» For this frame. which was removed from analysis. the depth reached is considerably shallower than the remaining frames leadiug not only to a bright lnuitine magnitude but also to the detection of a significant umuber of spurious objects.," For this frame, which was removed from analysis, the depth reached is considerably shallower than the remaining frames leading not only to a bright limiting magnitude but also to the detection of a significant number of spurious objects."
+ Other frames with bright limiting isophotes have no significant impact in the analysis preseuted iu section 3.., Other frames with bright limiting isophotes have no significant impact in the analysis presented in section \ref{analysis}.
+ Figures 3.Pa and show. for each band. the two-dimensional distribution of the ποιο and lanitine isophote as determined from the even frames.," Figures \ref{fig:2dseeing} and \ref{fig:2dmlim} show, for each band, the two-dimensional distribution of the seeing and limiting isophote as determined from the even frames."
+ Similar results are obtained for the odd frames which altcrnate with the even ones., Similar results are obtained for the odd frames which alternate with the even ones.
+ Such maps allow the poteutial user of the derived catalogs to evaluate their reliability., Such maps allow the potential user of the derived catalogs to evaluate their reliability.
+ The final data is reasonably homogeneous with the median secing in all bands <1.2 arcsec., The final data is reasonably homogeneous with the median seeing in all bands $< 1.2$ arcsec.
+ However. some poor images do exist aud for some applications nist be removed. as discussed above.," However, some poor images do exist and for some applications must be removed, as discussed above."
+ The area that is affected i8. <0.1 square degree., The area that is affected is $\lsim 0.1$ square degree.
+ Finally. it ds worth inentioudug that since the conipletiou of paper L the V inages for patch A over au area l.l square degrees have also been reduced aud are being made available together with the catalogs extracted from them which are used below.," Finally, it is worth mentioning that since the completion of paper I, the $V$ images for patch A over an area $\sim 1.1$ square degrees have also been reduced and are being made available together with the catalogs extracted from them which are used below."
+ The data Were processec »* the EIS pipeline being developed to haudle large inagiug programs aud described in detail iu paper L The software development is still In progress with new functionalities being constantlv added to the pipeline as well as euhauced features in Skveat driven by the survey needs. in particular to facilitate the visual iuspection of the target lists," The data were processed by the EIS pipeline being developed to handle large imaging programs and described in detail in paper I. The software development is still in progress with new functionalities being constantly added to the pipeline as well as enhanced features in Skycat driven by the survey needs, in particular to facilitate the visual inspection of the target lists"
+the verification of the theories of dvnamical evolution.,the verification of the theories of dynamical evolution.
+ Tam grateful to all the members of the Group of HWieh Anenlar Resolution Methods in Astronomy of the SAO RAS for their help in securing the observations made. and personally to Yu.," I am grateful to all the members of the Group of High Angular Resolution Methods in Astronomy of the SAO RAS for their help in securing the observations made, and personally to Yu."
+ Dalega. A. Tokovinin and an anonvinous referee for valuable comments that led to a notable improvement of this article.," Balega, A. Tokovinin and an anonymous referee for valuable comments that led to a notable improvement of this article."
+ I (hank D. Latham. who provided the data on spectroscopic multiplicity of objects from the CLLA catalog along with the corrected periods of some binary stars. V. Dvachenko for finding the components of the stellar streams members in the literature and A. Zvazeva for the help with translation.," I thank D. Latham, who provided the data on spectroscopic multiplicity of objects from the CLLA catalog along with the corrected periods of some binary stars, V. Dyachenko for finding the components of the stellar streams members in the literature and A. Zyazeva for the help with translation."
+ I acknowledge the continuous support provided by the Russian Foundation for Basic Research (project no., I acknowledge the continuous support provided by the Russian Foundation for Basic Research (project no.
+ 04-02-17563) and the Programme of Physical Sciences of the Russian Academy of Sciences.).. To be able to estimate (he number of stus in a given volume by an extrapolation of the calculated. quantity of stars in a smaller volume. we have to lake into account the structure ol the Galaxy.," 04-02-17563) and the Programme of Physical Sciences of the Russian Academy of Sciences.. To be able to estimate the number of stars in a given volume by an extrapolation of the calculated quantity of stars in a smaller volume, we have to take into account the structure of the Galaxy."
+ A null hypothesis is that in the volume of 250 pe from the Sun. the amount of thick disc stars is 1000 times bigger than that within 25 pe.," A null hypothesis is that in the volume of 250 pc from the Sun, the amount of thick disc stars is $1\ 000$ times bigger than that within 25 pc."
+ We assume a uniform —istvibution of halo stars on the scale of 100 pe’ni in the solar vicinity., We assume a uniform distribution of halo stars on the scale of 100 $^3$ in the solar vicinity.
+ Let us consider how 1 exponential decrease in the number density wilh an increase in (he distance from the ealactic plane in the sample of thick dise stars affects the null hypothesis., Let us consider how an exponential decrease in the number density with an increase in the distance from the galactic plane in the sample of thick disc stars affects the null hypothesis.
+" Let's introduce a &=N4/N, coefficient. where"," Let's introduce a $k=N_{1}/N_{0}$ coefficient, where"
+an increase in (he number of objects detected on the inverted image.,an increase in the number of objects detected on the inverted image.
+ This suggests that we are not reliably adding new real clusters to our sample., This suggests that we are not reliably adding new real clusters to our sample.
+ In addition to these false clusters. we also expect to have a contribution lrom faint ealaxies behind M87.," In addition to these false clusters, we also expect to have a contribution from faint galaxies behind M87."
+ This contribution can be estimated by measuring how many objects we find with our settings on very deep HST images like those of the Hubble Deep Field (Williamsetal.1996)., This contribution can be estimated by measuring how many objects we find with our settings on very deep HST images like those of the Hubble Deep Field \citep{HDF}.
+. The detection ellicieney [or faint extended objects is very. sensitive to (he signal to noise ratio of the image., The detection efficiency for faint extended objects is very sensitive to the signal to noise ratio of the image.
+ This required us to degrade the quality of the IDF image to match the average variance of our images., This required us to degrade the quality of the HDF image to match the average variance of our images.
+ Το help remove as many contaaninating objects as possible. we clemand (hat (wo requirements be met before an object was determined to be a valid globular cluster.," To help remove as many contaminating objects as possible, we demand that two requirements be met before an object was determined to be a valid globular cluster."
+ The first requirement is that all globular cluster candidates must be detected in both filters., The first requirement is that all globular cluster candidates must be detected in both filters.
+ This removes a significant portion of the false clusters. as well as a few of the IIDF objects.," This removes a significant portion of the false clusters, as well as a few of the HDF objects."
+ By further requiring that all candidates have a final V—7 color between 0.5 aud 1.7. we exclude almost. all of the remaining background sources and false clusters.," By further requiring that all candidates have a final $V -
+I$ color between 0.5 and 1.7, we exclude almost all of the remaining background sources and false clusters."
+ As shown in Table 1.. each of these cuts removes approximately the same number of clusters from both the observed and false sets which suggests that these cuts preserve the real clusters. and only remove objects that. are not elobular clusters.," As shown in Table \ref{tab:detection}, each of these cuts removes approximately the same number of clusters from both the observed and false sets which suggests that these cuts preserve the real clusters, and only remove objects that are not globular clusters."
+ With the detection. parameters set. we next used model clusters of known flux to determine the completeness of our sample.," With the detection parameters set, we next used model clusters of known flux to determine the completeness of our sample."
+ All of the model clusters were generated [rom King model profiles with a fixed concentration of 1.26 at eight tidal radius steps evenly spaced between 2 and 16 pixels., All of the model clusters were generated from King model profiles with a fixed concentration of 1.26 at eight tidal radius steps evenly spaced between 2 and 16 pixels.
+ These templates were (hen scaled (o form a grid in magnitude and size. Wilh ES14W. and F606W apparent magnitudes spaced every 0.5 magnitude between 20 and 28 magnitude.," These templates were then scaled to form a grid in magnitude and size, with F814W and F606W apparent magnitudes spaced every 0.5 magnitude between 20 and 28 magnitude."
+ This range of radii ancl magnitudes was chosen (o cover (he properties of the real clusters in our data sample., This range of radii and magnitudes was chosen to cover the properties of the real clusters in our data sample.
+ For each combination of magnitude and (tidal radius. 150 copies of the template were randomly placed on the background subtracted images.," For each combination of magnitude and tidal radius, 150 copies of the template were randomly placed on the background subtracted images."
+ Source Extractor was (hen used {ο lind these objects with the same detection parameters as were used on the data image., Source Extractor was then used to find these objects with the same detection parameters as were used on the data image.
+ By measuring the rate al which the templates are recovered. we can determine how complete our data sample is.," By measuring the rate at which the templates are recovered, we can determine how complete our data sample is."
+ Because the change in the completeness did not vary much lor template clusters of different sizes. the final completeness is found by simply averaging the results for all radii together.," Because the change in the completeness did not vary much for template clusters of different sizes, the final completeness is found by simply averaging the results for all radii together."
+ As the detection process is sensitive to the background uncertainty. the completeness limit becomes fainter as the distance [rom the center of the galaxy increases.," As the detection process is sensitive to the background uncertainty, the completeness limit becomes fainter as the distance from the center of the galaxy increases."
+ To take advantage of this increased. sensitivitv. we divided our sample into (wo radial bins with equal numbers of clusters in each.," To take advantage of this increased sensitivity, we divided our sample into two radial bins with equal numbers of clusters in each."
+ This gives bins that have rages of projected galactocentric, This gives bins that have rages of projected galactocentric
+o2.092.12.,$\alpha \sim 2.09-2.12$.
+ These results demonstrate that CBM is able to see siguificaut emission above 300 keV at a level consistent with the canonical hard spectra of the Crab., These results demonstrate that GBM is able to see significant emission above 300 keV at a level consistent with the canonical hard spectrum of the Crab.
+ Future analysis with the finer spectral resolution of the GBAL CSPEC data will allow a better determination of the the break energy aud spectral index above the break., Future analysis with the finer spectral resolution of the GBM CSPEC data will allow a better determination of the the break energy and spectral index above the break.
+ Cyveuus N-1l is a highaunass X-ray binary and Was Olle of the first svstenis determined contain a black hole (Bolton1972:Paczvusk. 1971).," Cygnus X-1 is a high-mass X-ray binary and was one of the first systems determined to contain a black hole \citep{Bolton1972,Paczynski1974}."
+ The N-ray ciuission is bimodal. with the 10 keV enussou anticorrelated witi the <10 keV cussion (Dolanetal.1977).," The X-ray emission is bimodal, with the $>10$ keV emission anticorrelated with the $<10$ keV emission \citep{Dolan1977}."
+. I has CCL observed to enüt sienificaut ewission above 100 keV inchiding a power law tail exteudiug out ereater than 1 MeV (MeConnelletal.2000:Line&Wheaton 2005b).," It has been observed to emit significant emission above 100 keV including a power law tail extending out to greater than 1 MeV \citep{McConnell2000,Ling2005b}."
+. Based on BATSE occultation analysis. Lii5&Wheaton(20)b) have show hat in the high eamuna-rav iucusity (lard) state. he spectrum consists of a Comptonized shape yvclow 200300 keV with a soft (D 3) power-aw tail extending to at least 1 MeV. Iu the ow-iutensitv (soft) state. however. the spectra. akes on a cifferent shape: In this case. the cutive Spectruu from 30 keV to 1 MeV is characterized Noa suele power law with a larder index D-227.," Based on BATSE occultation analysis, \citet{Ling2005b} have shown that in the high gamma-ray intensity (hard) state, the spectrum consists of a Comptonized shape below 200–300 keV with a soft $\Gamma > 3$ ) power-law tail extending to at least 1 MeV. In the low-intensity (soft) state, however, the spectrum takes on a different shape: In this case, the entire spectrum from 30 keV to 1 MeV is characterized by a single power law with a harder index $\Gamma \sim 2 - 2.7$."
+ Shuilar behavior has been reported for two low-mass N-ray binaries (LMXNDs) that also contain black holes. the trausicut θααν sources GRO J0122|32 (Ling&Wheaton2003a) and GRO J1719-21 (Ling&Wheaton2005a).," Similar behavior has been reported for two low-mass X-ray binaries (LMXBs) that also contain black holes, the transient gamma-ray sources GRO J0422+32 \citep{Ling2003a} and GRO J1719-24 \citep{Ling2005a}."
+. Figure 3 shows the CBM light curves., Figure \ref{CygX1} shows the GBM light curves.
+ The helt curves show significant variability with emission above 300 keV up until about NLJD 55355., The light curves show significant variability with emission above 300 keV up until about MJD 55355.
+ Starting at about MJD 55355.mma the 100300 keV baud enissiou began to decrease (Wilson-IHodgoe&Case 2010).. dropping from an average level of about 1200 iiCvab down to nearly uudetectable levels on MJD 5510555106.," Starting at about MJD 55355, the 100–300 keV band emission began to decrease \citep{WilsonCase2010}, dropping from an average level of about 1200 mCrab down to nearly undetectable levels on MJD 55405–55406."
+ On MJD 55371 MANI detected a rapid rise in the soft 2| keV baud (Negoroetal.20105).. which combined with the decrease in the low energv gamma-ray filx indicated a transition to a thermally-cdouunatecd soft state.," On MJD 55374 MAXI detected a rapid rise in the soft 2–4 keV band \citep{Negoro2010b}, which combined with the decrease in the low energy gamma-ray flux indicated a transition to a thermally-dominated soft state."
+ As of NLID 55119 the one-day. average CGDM light curves show that the 1250 keV f has begun to rise. while the 100300 keV fn renmiadns at a low level of =150 miCrab. consiste with the συ state of Lingetal.(1997).," As of MJD 55419 the one-day average GBM light curves show that the 12–50 keV flux has begun to rise, while the 100–300 keV flux remains at a low level of $\approx 150$ mCrab, consistent with the $\gamma_0$ state of \citet{Ling1997}."
+. We will continue to monitor Cve N-1 during the thermalv-dominated state and follow its transition back O the low/hard state., We will continue to monitor Cyg X-1 during the thermally-dominated state and follow its transition back to the low/hard state.
+ The GBAL light curves (Fig. 3)), The GBM light curves (Fig. \ref{CygX1}) )
+ reveal siguificant cnussion above 300 keV. consistent with the power law tail observed when (νο N-1 is in its low/haud state.," reveal significant emission above 300 keV, consistent with the power law tail observed when Cyg X-1 is in its low/hard state."
+ The 50100 keV flix level observec by GBA, The 50–100 keV flux level observed by GBM
+ The 50100 keV flix level observec by GBAL, The 50–100 keV flux level observed by GBM
+center. and b ancl | are the Galactic coordinates of the nebulae.,"center, and b and l are the Galactic coordinates of the nebulae."
+ While the statistical distance scale can not provide errors. SSV showed that uncertainties are well within 30% ol the calibration.," While the statistical distance scale can not provide errors, SSV showed that uncertainties are well within $\%$ of the calibration."
+ We assign a rather conservative uncertainty of 420% to the heliocentric distances. and we then propagate the error to the galactocentric distances. assuming that the Galactic coordinates | and b have much lower relative uncertainties than the heliocentric distances.," We assign a rather conservative uncertainty of $\pm$ $\%$ to the heliocentric distances, and we then propagate the error to the galactocentric distances, assuming that the Galactic coordinates l and b have much lower relative uncertainties than the heliocentric distances."
+ The galactocentric distances and their (formal) uncertainties are given in Table ]. eolumn (3).," The galactocentric distances and their (formal) uncertainties are given in Table 1, column (3)."
+ There are 1143 PNe in Acker’s catalog. whose diameters are measured. (either directlv or as an upper limit) in 1053 cases.," There are 1143 PNe in Acker's catalog, whose diameters are measured (either directly or as an upper limit) in 1053 cases."
+ We were able to determine the PN distances for ~730 Ne., We were able to determine the PN distances for $\sim$ 730 PNe.
+ Perinotto οἱ al. (, Perinotto et al. (
+2004) have carefully reviewed all elemental abundance analysis of Galactic Όλο published since the seventies. and selected a sample that is homogeneous in abundance determination methods and high quality of the observations (their sample A).,"2004) have carefully reviewed all elemental abundance analysis of Galactic PNe published since the seventies, and selected a sample that is homogeneous in abundance determination methods and high quality of the observations (their sample A)."
+ We use the A sample from PO4 in our study. augmented with the abundances published later than the PO4 sample was analvzed. selecting the sets with the same criteria. used by POL," We use the A sample from P04 in our study, augmented with the abundances published later than the P04 sample was analyzed, selecting the sets with the same criteria used by P04."
+ Our updated sample includes abundances from (1) PO4: (2) Stanehellini et al. (, Our updated sample includes abundances from (1) P04; (2) Stanghellini et al. (
+2006): (3) the series of papers by Pottasch aud collaborators. summarized by Pottasch Dernard-Salas 2006: (4) IxXvitter Henry (2001). Milingo et al. (,"2006); (3) the series of papers by Pottasch and collaborators, summarized by Pottasch Bernard-Salas 2006; (4) Kwitter Henry (2001), Milingo et al. ("
+2002). and INwitter et al. (,"2002), and Kwitter et al. ("
+2003): and (5) and Costa et al. (,2003); and (5) and Costa et al. (
+2004).,2004).
+ It is worth noting that the series of papers by Ix. Kwitter and collaborators are formally dated before the PO4 review. but actually the DPO4 work was completed before (he former abundauces have became available. (hus not included in POS.," It is worth noting that the series of papers by K. Kwitter and collaborators are formally dated before the P04 review, but actually the P04 work was completed before the former abundances have became available, thus not included in P04."
+ There are several other papers on Galactic PN abundances published after. 2004. but with the exception of the ones listed above thev are reanalvsis of earlier databases.," There are several other papers on Galactic PN abundances published after 2004, but with the exception of the ones listed above they are reanalysis of earlier databases,"
+at ULIRGs in a narrow slice around ((Lip fL. 2212.5 to be consistent with the selection function. displaved. however the samples are too small and. hence we plot all 12x (bin /L 212.5 sources.,"at ULIRGs in a narrow slice around $_{\rm IR}$ $_{\odot}$ =12.5 to be consistent with the selection function displayed, however the samples are too small and hence we plot all $\le$ $_{\rm IR}$ $_{\odot}$ $\le$ 12.5 sources."
+ In addition. all subsequent estimates of the average peak wavelength and temperature refer to these log((Lin/L. 12.5 sources only. in order to achieve a fair comparison with local ULIRGs whose luminosities do not extend bevond ((Lin /L.)212.5.," In addition, all subsequent estimates of the average peak wavelength and temperature refer to these $_{\rm IR}$ $_{\odot}$ $\le$ 12.5 sources only, in order to achieve a fair comparison with local ULIRGs whose luminosities do not extend beyond $_{\rm IR}$ $_{\odot}$ )=12.5."
+ The MIPS iini survey can potentially detect objects up to z=2. but with severe incompleteness above z=1 it is unbiased up to z~0.7 and sensitive to warmer SEDs above that redshift.," The MIPS $\mu$ m survey can potentially detect objects up to z=2, but with severe incompleteness above z=1 — it is unbiased up to $\sim$ 0.7 and sensitive to warmer SEDs above that redshift."
+ Our predictions are consistent with results from Casey ct al. (2009), Our predictions are consistent with results from Casey et al. )
+) who find. all zz1. 4m- ULIRGs to be hot.," who find all $>$ 1, $\mu$ m-detected ULIRGs to be hot."
+ Surveys at. μι have a similar selection function to the ones at 705m. but with a temperature range which varies less strongly with recshilt. apart [from the very hot (Apos i440 jn) SED region.," Surveys at $\mu$ m have a similar selection function to the ones at $\mu$ m, but with a temperature range which varies less strongly with redshift, apart from the very hot $\lambda_{\rm peak}$ $<$ $\mu$ m) SED region."
+ On the other hand for the longer wavelength surveys the likelihood. of detection is significantly more dependent on SED type., On the other hand for the longer wavelength surveys the likelihood of detection is significantly more dependent on SED type.
+ DLAST is highly sensitive to colder SEDs. but sources with Apenk= pmi can still be detected. at. 250g up to z=l.," BLAST is highly sensitive to colder SEDs, but sources with $\lambda_{\rm peak}$ $\mu$ m can still be detected at $\mu$ m up to z=1."
+ Down to 2PmmJv SCUBA should be able to detect ULIRGs over the entire Apcsua-z parameter space. primarilv because of the negative A--correction which olfers a large advantage above z 11.," Down to mJy SCUBA should be able to detect ULIRGs over the entire $\lambda_{\rm peak}$ -z parameter space, primarily because of the negative -correction which offers a large advantage above $\sim$ 1."
+ However above z=0.5. there is à strong predisposition towards colder CA:277 pm) objects.," However above z=0.5, there is a strong predisposition towards colder $\lambda_{\rm peak}$$>$ $\mu$ m) objects."
+ In addition. the large area below A4: 077m covered. by the blue regions in the BLAST and SCUBA panels suggests that warm SEDs can only be recovered if they have sullicient submum Iux. te. an additional cold temperature Component.," In addition, the large area below $\lambda_{\rm peak}$ $\mu$ m covered by the blue regions in the BLAST and SCUBA panels suggests that warm SEDs can only be recovered if they have sufficient submm flux, i.e. an additional cold temperature component."
+ Llow does the dust temperature distribution. of the BLAST. SCUBA and samples compare in the context of their selection characteristics?," How does the dust temperature distribution of the BLAST, SCUBA and samples compare in the context of their selection characteristics?"
+ The 809 sample's Apes distribution shifts [rom longer to shorter values with increasing redshift. following the shape of the selection function. which shows increased sensitivity towards warm SEDs above z=0.7.," The S09 sample's $\lambda_{\rm peak}$ distribution shifts from longer to shorter values with increasing redshift, following the shape of the selection function, which shows increased sensitivity towards warm SEDs above z=0.7."
+ Although above z=0.8. where the survey becomes more sensitive tO Ape i70 yan. most ULIRGs are warm. it is interesting to note that in the unbiased region (20.7) most ULLItCis are cold.," Although above z=0.8, where the survey becomes more sensitive to $\lambda_{\rm peak}$ $<$ $\mu$ m, most ULIRGs are warm, it is interesting to note that in the unbiased region $<$ 0.7) most ULIRGs are cold."
+ We calculate, We calculate
+group luminosity function.,group luminosity function.
+" On a similar argument, the same effect neatly stems, as well, from a study of the po nSérrsicparameters, likein Fig. 4.."," On a similar argument, the same effect neatly stems, as well, from a study of the $\mu_0$ $n$ Sérrsic parameters, like in Fig. \ref{f4}."
+" The plot shows, in addition, that early-type profiles actually tend to smoothly match the standard de Vaucouleurs case (n— 0.25) as far as galaxy luminosity (and accordingly uo) brightens up reaching the distinctive range of ""standard"" ellipticals, as confirmed by the straight L—n observed relationship."," The plot shows, in addition, that early-type profiles actually tend to smoothly match the standard de Vaucouleurs case $n \to 0.25$ ) as far as galaxy luminosity (and accordingly $\mu_0$ ) brightens up reaching the distinctive range of “standard” ellipticals, as confirmed by the straight $L-n$ observed relationship."
+" At the faint end, dSph's display a wide range in n, while their central surface brightnesses remain almost constant at Lo(g)©25 magaarcsec""?."," At the faint end, dSph's display a wide range in $n$, while their central surface brightnesses remain almost constant at $\mu_0(g) \approx 25$ $^{-2}$."
+ The Gunn griz magnitudes and the Johnson BV photometry provide a minimal but effective set of measures to probe galaxy SEDs along the 4400-9000 wwavelength range., The Gunn $griz$ magnitudes and the Johnson $BV$ photometry provide a minimal but effective set of measures to probe galaxy SEDs along the 4400-9000 wavelength range.
+" In addition to the morphological piece of information, in fact, a multicolour study of integrated luminosity of our targets, as well as of their surface brightness distributions could provide us with important clues to tackle the distinctive evolutionary properties of the galaxy population in the NGC 5044 group."," In addition to the morphological piece of information, in fact, a multicolour study of integrated luminosity of our targets, as well as of their surface brightness distributions could provide us with important clues to tackle the distinctive evolutionary properties of the galaxy population in the NGC 5044 group."
+" Again, our final results are collected in Table A1 and A2,, respectively for the and In both tables, the total apparent magnitude (encircled within the u=27 magaarcsec? isophote in the g and V bands, respectively) is reported in col."," Again, our final results are collected in Table \ref{a2} and \ref{a3}, respectively for the and In both tables, the total apparent magnitude (encircled within the $\mu = 27$ $^{-2}$ isophote in the $g$ and $V$ bands, respectively) is reported in col."
+" 11, together with the mean surface brightness within the same isophotal radius, and within one effective radius (cols."," 11, together with the mean surface brightness within the same isophotal radius, and within one effective radius (cols."
+" 12 and 13, respectively) according to the corresponding values of p27 and pe."," 12 and 13, respectively) according to the corresponding values of $\rho_{27}$ and $\rho_\mathrm{e}$."
+ Our output is finally completed with the griz colours (cols., Our output is finally completed with the $griz$ colours (cols.
+ 14 to 19 in Table A1)) and (B—V) (cols., 14 to 19 in Table \ref{a2}) ) and $(B-V)$ (cols.
+ 14 and 15 inTable A2)) across the same relevant apertures., 14 and 15 inTable \ref{a3}) ) across the same relevant apertures.
+" Unless explicitely stated, note that no correction for Galactic reddening has been introduced."," Unless explicitely stated, note that no correction for Galactic reddening has been introduced."
+" According to Burstein&Heiles (1982), the colour excess in the sky region around NGC 5044 amounts to E(g—r)~E(BV)0.03 mag, a figure that may raise to ~0.07 mag according to the Schlegel,Finkbeiner,&Davis(1998) reddening map."," According to \citet{burstein82}, , the colour excess in the sky region around NGC 5044 amounts to $E(g-r) \simeq E(B-V) \sim 0.03$ mag, a figure that may raise to $\sim 0.07$ mag according to the \citet{schlegel98}
+ reddening map."
+" The colour distribution in the plane of integrated (g—r) for the whole sample of 59 “likely member"" (membership code m.<2 inTables Al and A2) galaxies with available colours is displayed in the upper panel of Fig. 5..", The colour distribution in the plane of integrated $(g-r)$ for the whole sample of 59 “likely member” (membership code $m_c \le 2$ inTables \ref{a2} and ) galaxies with available colours is displayed in the upper panel of Fig. \ref{f5}. .
+ Inthe, Inthe
+photometric redshift space: the algorithm might miss some members in the outer regions more likely to be star-forming galaxies.,photometric redshift space; the algorithm might miss some members in the outer regions more likely to be star-forming galaxies.
+ Note also that due to the near-infrared selection the MUNICS group sample may be biased against high-SFR galaxies with low dust attenuation., Note also that due to the near-infrared selection the MUNICS group sample may be biased against high-SFR galaxies with low dust attenuation.
+ For easier analysis. the approximate upper boundaries. to the SSFR S for the different samples in Fig.," For easier analysis, the approximate upper boundaries to the SSFR ${\cal S}$ for the different samples in Fig."
+ | can be described by the following functional form (similar to the Schechter parametrisation of the luminosity function. ?)): The free parameters of this function. describe. the normalisation (So). the location of the break (Ato). and the slope at lower stellar masses (a).," \ref{fig:ssfr} can be described by the following functional form (similar to the Schechter parametrisation of the luminosity function, \citealt{Schechter76}) ): The free parameters of this function describe the normalisation ${\cal S}_{\,0}$ ), the location of the break ${\cal M}_{\,0}$ ), and the slope at lower stellar masses $\alpha$ )."
+ Their approximate values are (logSo/Gyr.logΛίαΛίο.a)=(5.25.10.5.-1.5) for the field galaxies. (5.25.11.7.—1.5) for the groups. and (5.25.13.0.—1.5) for the clusters.," Their approximate values are $(\log {\cal S}_{\,0} / \mathrm{Gyr}^{-1}, \: \log {\cal M}_{\,0}
+/ {\cal M}_{\,\odot}, \: \alpha) \simeq (5.25, \: 10.5, \: -1.5)$ for the field galaxies, $(5.25, \: 11.7, \: -1.5)$ for the groups, and $(5.25, \: 13.0, \: -1.5)$ for the clusters."
+ The slope « can be derived from the field-galaxy sample only but seems to apply also to the other samples., The slope $\alpha$ can be derived from the field-galaxy sample only but seems to apply also to the other samples.
+ The curves corresponding to these values are plotted in Fig. 1.., The curves corresponding to these values are plotted in Fig. \ref{fig:ssfr}.
+ Note. again. that there is a smooth transition from groups to clusters. so the upper mass limit for groups should be taken with a grain of salt.," Note, again, that there is a smooth transition from groups to clusters, so the upper mass limit for groups should be taken with a grain of salt."
+ In this Letter we have for the first time presented the integrated SSFR for groups and clusters of galaxies and compared it to the SSFR of the field galaxy population.," \nocite{fdfssfr, munics8, Finn2005}
+ In this Letter we have for the first time presented the integrated SSFR for groups and clusters of galaxies and compared it to the SSFR of the field galaxy population."
+ Moreover. we tentatively find a continuous upper limit for galaxies. groups. and clusters in the SSFR-stellar mass plane over seven orders of magnitude in stellar mass.," Moreover, we tentatively find a continuous upper limit for galaxies, groups, and clusters in the SSFR-stellar mass plane over seven orders of magnitude in stellar mass."
+ This might indicate that the processes which control star formation in dark matter haloes of different mass have the same scaling with mass over a wide range of masses from dwarf galaxies to massive clusters of galaxies., This might indicate that the processes which control star formation in dark matter haloes of different mass have the same scaling with mass over a wide range of masses from dwarf galaxies to massive clusters of galaxies.
+" The physical processes responsible for the ""down-sizing"" phenomenon witnessed in galaxies are not yet well understood.", The physical processes responsible for the “down-sizing” phenomenon witnessed in galaxies are not yet well understood.
+" An early formation epoch for massive galaxies or ""dry merging"" (22). of lower-mass galaxies as well as quenching of star formation in. more massive haloes by feedback mechanisms (e.g.?) are among the discussed possibilities."," An early formation epoch for massive galaxies or “dry merging” \citep{Faber2006, Bell2006} of lower-mass galaxies as well as quenching of star formation in more massive haloes by feedback mechanisms \citep[e.g.][]{Scannapieco2005} are among the discussed possibilities."
+ Of course. we could also see the result of a combination of these processes or be faced with different evolutionary paths leading to the population of massive galaxies with old stellar populations.," Of course, we could also see the result of a combination of these processes or be faced with different evolutionary paths leading to the population of massive galaxies with old stellar populations."
+" The fact that the SSFRs of groups and clusters of galaxies seem to continue the trend displayed by the field galaxy population towards higher masses is intrigui5,", The fact that the SSFRs of groups and clusters of galaxies seem to continue the trend displayed by the field galaxy population towards higher masses is intriguing.
+ For the galaxies. within these structures. one ca naturally expect a similar general behaviour às for their counterparts in the field. modified by environmental effects.," For the galaxies within these structures, one can naturally expect a similar general behaviour as for their counterparts in the field, modified by environmental effects."
+ It has been known for a long time that higher density enviroments are occupied by galaxies with morphologically earlier types (e.g.POD) and with overall redder colours (and thus lower star-formation activity: ?)).," It has been known for a long time that higher density environments are occupied by galaxies with morphologically earlier types \citep[e.g.][]{Dressler1980,
+Postman1984, Dressler1997} and with overall redder colours (and thus lower star-formation activity; \citealt{BO78a}) )."
+ Moreover. ellipticals i higher-density environments are on average older than their low-density counterparts (?).. and star formation activity in groups seems to be lower than in the field (?)..," Moreover, ellipticals in higher-density environments are on average older than their low-density counterparts \citep{Thomas2005}, and star formation activity in groups seems to be lower than in the field \citep{Wilman2005}."
+ But the fact that the upper limit of the SSERs of all these objects. from dwarf galaxies to rich clusters. seems to follow a continuous sequence in the SSFR-stellar mass plane seems to suggest that there could be a smooth transition from the field to the clusters. which in turn might imply that the physical processes responsible for the lower integrated star formation activity in higher mass haloes are the same over this wide range of stellar masses. or at least have the same scaling with stellar mass.," But the fact that the upper limit of the SSFRs of all these objects, from dwarf galaxies to rich clusters, seems to follow a continuous sequence in the SSFR–stellar mass plane seems to suggest that there could be a smooth transition from the field to the clusters, which in turn might imply that the physical processes responsible for the lower integrated star formation activity in higher mass haloes are the same over this wide range of stellar masses, or at least have the same scaling with stellar mass."
+ The analysis presented here is made possible by the availability of large samples of field galaxies with well studied properties. and by the advent of group and cluster catalogues with photometric and spectroscopic data for large number of members.," The analysis presented here is made possible by the availability of large samples of field galaxies with well studied properties, and by the advent of group and cluster catalogues with photometric and spectroscopic data for large number of members."
+ However. statistics is still rather poor for groups and clusters. and we could not derive SFRs and stellar masses using the same methods in all samples.," However, statistics is still rather poor for groups and clusters, and we could not derive SFRs and stellar masses using the same methods in all samples."
+ Future studies of large and homogeneous samples of groups. clusters and their member galaxies will result in progress in the study of galaxy evolution as a function of local density. and allow us to better constrain the physical processes responsible for controlling star formation in different environments and thus haloes of different masses.," Future studies of large and homogeneous samples of groups, clusters and their member galaxies will result in progress in the study of galaxy evolution as a function of local density, and allow us to better constrain the physical processes responsible for controlling star formation in different environments and thus haloes of different masses."
+Ganma-hayv Bursts (GRBs) are indeed. “gamma-ray bursts since (heir spectrum is dominated by 5gamma radiation.,Gamma-Ray Bursts (GRBs) are indeed 'gamma-ray' bursts since their spectrum is dominated by gamma radiation.
+" E,,,peak is the characteristic photon energve measured. [rom", $E_{peak}$ is the characteristic photon energy measured from
+"With the £15 filter covering the restframe 8j/m. we simply convert the observed flux to ὅμπι monochromatic luminosity L5, ) using a standard cosmology.","With the $L15$ filter covering the restframe $\mu$ m, we simply convert the observed flux to $\mu$ m monochromatic luminosity $L_{8\mu m}$ ) using a standard cosmology."
+ Completeness was measured by distributing artificial point sources with varying flux within the field and by examining what fraction of them was recovered as a function of input flux., Completeness was measured by distributing artificial point sources with varying flux within the field and by examining what fraction of them was recovered as a function of input flux.
+ Since we have deeper coverage at the center of the cluster. the completeness was measured separately in the central deep region and the outer regions of the field.," Since we have deeper coverage at the center of the cluster, the completeness was measured separately in the central deep region and the outer regions of the field."
+ More detail of the method is deseribed in Wadaetal.(2008)., More detail of the method is described in \citet{2008PASJ...60S.517W}.
+ Once the flux is converted to luminosity and completeness is taken into account. it is straight forward to construct ζω LFs. which we show in the squares in Fig.1..," Once the flux is converted to luminosity and completeness is taken into account, it is straight forward to construct $L_{8\mu m}$ LFs, which we show in the squares in \ref{fig:8umlf}."
+ Errors of the LFs are assumed to follow Poisson distribution., Errors of the LFs are assumed to follow Poisson distribution.
+ Here. we take an angular distance of the most distant source from the cluster center as a cluster radius (2)...= 6.2Mpe).," Here, we take an angular distance of the most distant source from the cluster center as a cluster radius $R_{max}=6.2$ Mpc)."
+ We assumed STR}a as the volume of the cluster to obtain galaxy density (0)., We assumed $\frac{4}{3}\pi R_{\max}^3$ as the volume of the cluster to obtain galaxy density $\phi$ ).
+ This ts only one of many ways to define a cluster volume. and thus. a caution must be taken to compare absolute values of our LFs to other work such as Shimetal.(2010).," This is only one of many ways to define a cluster volume, and thus, a caution must be taken to compare $absolute$ values of our LFs to other work such as \citet{Shim2010}."
+. This cluster is elongated in angular direction (Koyamaetal..2007) . dend thus. the volume might not be spherical.," This cluster is elongated in angular direction \citep{2007MNRAS.382.1719K} , and thus, the volume might not be spherical."
+ Yet. comparison of the shape of the LFs is valid.," Yet, comparison of the $shape$ of the LFs is valid."
+ Our field LFs are based on the AKARI NEP Deep field data., Our field LFs are based on the AKARI NEP Deep field data.
+ The AKARI performed deep imaging in the North Ecliptic Pole Region (NEP) from 2-24jm. with 4 pointings in each field over 0.4 deg? (Matsuharaetal.2006.2007;Wadaetαἱ... 2008)..," The AKARI performed deep imaging in the North Ecliptic Pole Region (NEP) from $\mu$ m, with 4 pointings in each field over 0.4 $^2$ \citep{2006PASJ...58..673M,2007PASJ...59S.543M,2008PASJ...60S.517W}."
+ The 5 σ sensitivity in the AKARI IR filters CN2.N3..NLOST.SOW.S11.L15.L1SW. and £21) are 14.2. 11.0. 5.0. 48. 58. 71. 117. 121 and 2754/Jy (Wadaetal..2008)..," The 5 $\sigma$ sensitivity in the AKARI IR filters $N2,N3,N4,S7,S9W,S11,L15,L18W$ and $L24$ ) are 14.2, 11.0, 8.0, 48, 58, 71, 117, 121 and $\mu$ Jy \citep{2008PASJ...60S.517W}."
+ Flux is measured m 3 pix radius aperture (277). then corrected to total flux.," Flux is measured in 3 pix radius aperture (=7”), then corrected to total flux."
+ A subregion of the NEP-Deep field (0.25 deg?) has ancillary data from Subaru BVΠο πιαetal..2007;Wadaetal. 2008). CFHT «(Serjeant et al.," A subregion of the NEP-Deep field (0.25 $^2$ ) has ancillary data from Subaru $BVRi'z'$ \citep{2007AJ....133.2418I,2008PASJ...60S.517W}, CFHT $u'$ (Serjeant et al."
+ in prep.).," in prep.),"
+ KPNO2n/FLAMINGOs ο and As(lmaietal.2007). GALEX FUV and NUV (Malkan et al.," KPNO2m/FLAMINGOs $J$ and $Ks$ \citep{2007AJ....133.2418I}, GALEX $FUV$ and $NUV$ (Malkan et al."
+ in prep.)., in prep.).
+ For the optical identification of MIR sources. we adopt the likelihood ratio method (Sutherland&Saunders.1992)..," For the optical identification of MIR sources, we adopt the likelihood ratio method \citep{1992MNRAS.259..413S}."
+ Using these data. we estimate photometric redshift of £15 detected sources in the region with the (Ibertetal..2006:Arnoutsetal.. 2007)..," Using these data, we estimate photometric redshift of $L15$ detected sources in the region with the \citep{2006A&A...457..841I,2007A&A...476..137A}."
+ The measured errors on the photo-: against 293 spec-z galaxies from Keck/DEIMOS (Takagi et al., The measured errors on the $z$ against 293 spec-z galaxies from Keck/DEIMOS (Takagi et al.
+ in prep.), in prep.)
+ are =0.036 at:<0.8., are $\frac{\Delta z}{1+z}$ =0.036 at $z\leq0.8$.
+ We have excluded those sources better fit P=with QSO templates from the LFs., We have excluded those sources better fit with QSO templates from the LFs.
+ To construct fieldLFs. we have selected £15 sources at 0.65“photos0.9.," To construct fieldLFs, we have selected $L15$ sources at $0.65<z_{photoz}<0.9$."
+ There remained 289 IR. galaxies with a median redshift of 0.76., There remained 289 IR galaxies with a median redshift of 0.76.
+" L15 flux is converted to Lj, using the photometric redshift of each galaxy.", $L15$ flux is converted to $L_{8\mu m}$ using the photometric redshift of each galaxy.
+ LPs are computed using the [Έως method., LFs are computed using the $V_{\max}$ method.
+ We used the SED templates (Lagache.Dole.&Puget.2003) for À-corrections to obtain the maximum observable redshift from the flux limit., We used the SED templates \citep{2003MNRAS.338..555L} for $k$ -corrections to obtain the maximum observable redshift from the flux limit.
+ Completeness of the £15 detection is corrected using(2009b)., Completeness of the $L15$ detection is corrected using.
+. This correction is at maximum. since we only use the sample where the completeness is greater than80%.," This correction is at maximum, since we only use the sample where the completeness is greater than."
+ The resulting field LFs are shown in the dotted line and triangles in Fig.1.., The resulting field LFs are shown in the dotted line and triangles in \ref{fig:8umlf}.
+ Errors of the LFs are computed using a 1000 Monte Caro simulation with varying + and flux within their errors., Errors of the LFs are computed using a 1000 Monte Caro simulation with varying $z$ and flux within their errors.
+ These estimated errors are added to the Poisson errors in each LF bin in quadrature., These estimated errors are added to the Poisson errors in each LF bin in quadrature.
+ We performed a detaild comparison of restframe 8jm LFs to those in the literature in Gotoetal.(2010).., We performed a detaild comparison of restframe $\mu$ m LFs to those in the literature in \citet{GOTO_NEP_LF}.
+" Briefly. there is an oder of difference between Caputietal.(2007) and Babbedgeetal.(2006).. reflecting difficulty in estimating Ls, dominated by PAH emissions using Spitzer 245m flux."," Briefly, there is an oder of difference between \citet{2007ApJ...660...97C} and \citet{2006MNRAS.370.1159B}, reflecting difficulty in estimating $L_{8\mu m}$ dominated by PAH emissions using Spitzer $\mu$ m flux."
+ Our field 8jjm LF lies between Caputietal.(2007) and Babbedgeetal.(2006)., Our field $\mu$ m LF lies between \citet{2007ApJ...660...97C} and \citet{2006MNRAS.370.1159B}.
+ Compared with these work. we have directly observed restframe 8j/m using the AKARI £15 filter. eliminating the uncertaintly in flux conversion based on SED models.," Compared with these work, we have directly observed restframe $\mu$ m using the AKARI $L15$ filter, eliminating the uncertaintly in flux conversion based on SED models."
+ More details and evolution of field IR LFs are described in Gotoetal.(2010).., More details and evolution of field IR LFs are described in \citet{GOTO_NEP_LF}.
+ In. Fig.]l.. we show restframe Sim LFs of cluster RXJI716.4|6708 in the squares. and LFs of the field region in the triangles.," In \ref{fig:8umlf}, we show restframe $\mu$ m LFs of cluster $+$ 6708 in the squares, and LFs of the field region in the triangles."
+ First of all. cluster LFs have by a factor of ~700 higher density than the field LFs. reflecting the fact the galaxy clusters is indeed high density regions in terms of infrared sources.," First of all, cluster LFs have by a factor of $\sim$ 700 higher density than the field LFs, reflecting the fact the galaxy clusters is indeed high density regions in terms of infrared sources."
+ Next. to compare the shape of the LFs. we normalized the cluster LF to match the field LFs at the faintest end. and show in the dash-dotted line.," Next, to compare the shape of the LFs, we normalized the cluster LF to match the field LFs at the faintest end, and show in the dash-dotted line."
+" In contrast to the field LFs. which show flattening of the slope at logLs,,«10.8... the cluster LF maintains the steep slope in the range of 10.01.logLa,<10.6L ..."," In contrast to the field LFs, which show flattening of the slope at $L_{8\mu m}<10.8L_{\odot}$, the cluster LF maintains the steep slope in the range of $10.0L_{\odot}<\log L_{8\mu m}<10.6L_{\odot}$ ."
+ The difference is significant. considering the size of errors on each LF.," The difference is significant, considering the size of errors on each LF."
+ We fit a double-power law to both cluster and field LFs using the following Free parameters are: £ (characteristic luminosity. Zo). 0 (normalization. à and 3 (faint and bright end slopes). respectively.," We fit a double-power law to both cluster and field LFs using the following Free parameters are: $L^*$ (characteristic luminosity, $L_{\odot}$ ), $\phi^*$ (normalization, $^{-3}$ ), $\alpha$ and $\beta$ (faint and bright end slopes), respectively."
+ The best fit values for field and cluster LFs are summarised in Table! and shown in the dotted lines in Fig.1..," The best fit values for field and cluster LFs are summarised in \ref{tab:fit_parameters}
+ and shown in the dotted lines in \ref{fig:8umlf}."
+ The bright-end slopes are not very different. but L ofthe cluster LF is smaller than the field by a factor of 2.4. and the faint-end tail of cluster LF is steeper than that of field LF.," The bright-end slopes are not very different, but $L^*$ ofthe cluster LF is smaller than the field by a factor of 2.4, and the faint-end tail of cluster LF is steeper than that of field LF."
+ To further examine the difference at the faint end of the LFs. we divide the cluster LF using the local galaxy density (X55) measured by Koyamaetal. (2008)..," To further examine the difference at the faint end of the LFs, we divide the cluster LF using the local galaxy density $\Sigma_{5th}$ ) measured by \citet{2008MNRAS.391.1758K}. ."
+ This density is based on the distance to the 5th nearest neighbor in the transverse, This density is based on the distance to the 5th nearest neighbor in the transverse
+eqn. (5)),eqn. \ref{eq:merge}) )
+ has been derived. on the basis of a number of assumptions. for example that the galaxy may be treated as à point mass.," has been derived on the basis of a number of assumptions, for example that the galaxy may be treated as a point mass."
+ Thus. whilst they recommenced a default value of fap=1. they also allow themselves the freedom to choose a cilferent value if required. to. produce a realistic model.," Thus, whilst they recommend a default value of $f_{\rm df}=1$, they also allow themselves the freedom to choose a different value if required to produce a realistic model."
+ Reeenth. Colpi..Mayer.&Covernato(1999) rave examined the validity of eqn.," Recently, \scite{colpi99}
+ have examined the validity of eqn."
+ 5 in detail using roth analytic ancl numerical techniques., \ref{eq:merge} in detail using both analytic and numerical techniques.
+" On the basis of heir investigations. they suggest small modifications to the ormula. in particular a slightly cilferent dependence of 6,4 on the orbital eccentricity."," On the basis of their investigations, they suggest small modifications to the formula, in particular a slightly different dependence of $\theta_{\rm orb}$ on the orbital eccentricity."
+ They. also find that the ellects of tidal stripping produce fap%2.7 (lor a specifie model of he satellites dark matter halo)., They also find that the effects of tidal stripping produce $f_{\rm df}\approx 2.7$ (for a specific model of the satellite's dark matter halo).
+ Both of these changes act ο INCLCASC Ty. and so reduce the merger rate.," Both of these changes act to increase $\tau_{\rm mrg}$, and so reduce the merger rate."
+ Since the nunocdel already contains too few high mass galaxies compared to the SDPII simulations. a slower merger rate would merely increase the discrepancy.," Since the model already contains too few high mass galaxies compared to the SPH simulations, a slower merger rate would merely increase the discrepancy."
+ At the resolution of our SPL simulations. an infalling satellite galaxy will lose nearly all of its original dark matter alo shortly after entering the larger halo.," At the resolution of our SPH simulations, an infalling satellite galaxy will lose nearly all of its original dark matter halo shortly after entering the larger halo."
+ On the other mand. in the SPILL simulations. merging with the central object is driven not only by dynamical friction. but also by drag due to viscous clleets as the satellite moves through the ιοί halo of the cluster (Erenketal.1996:Tittley.Couchman&Pearce. 1999).," On the other hand, in the SPH simulations, merging with the central object is driven not only by dynamical friction, but also by drag due to viscous effects as the satellite moves through the hot halo of the cluster \cite{fews,tittley99}."
+. Phese processes may drive the elfective fa ο à value less than unity.," These processes may drive the effective $f_{\rm
+df}$ to a value less than unity."
+ The smaller the value of fay. the faster a galaxy will sink to the centre of its host halo.," The smaller the value of $f_{\mathrm df}$, the faster a galaxy will sink to the centre of its host halo."
+ In Fig., In Fig.
+ 19 we show the ollect of reducing the merging rate from the default. value. far=αν το fay20.4 on the e&ealaxv mass function in the ACDAL model.," \ref{fig:cg_altmrg} we show the effect of reducing the merging rate from the default value, $f_{\mathrm df}=1$, to $f_{\mathrm df} = 0.4$ on the galaxy mass function in the $\Lambda$ CDM model."
+ This change improves the match between the aand SPLIT models at z=0. but only slightly. and mostly at high masses.," This change improves the match between the and SPH models at $z=0$, but only slightly, and mostly at high masses."
+ Below 10175.1M... the two models still differ significantly.," Below $10^{12}h^{-1}{\rm M}_\odot$, the two models still differ significantly."
+ At higher redshift the improvement is slightly better. but the nimocel still contains fewer high-mass objects than the SPILL simulation.," At higher redshift the improvement is slightly better, but the model still contains fewer high-mass objects than the SPH simulation."
+ Thus. a change in the global merger timescale does not help to reconcile the two models.," Thus, a change in the global merger timescale does not help to reconcile the two models."
+ This is perhaps not too surprising since this simple alteration cannot mimic any redshift dependence of the merger timescales in the SPL simulation., This is perhaps not too surprising since this simple alteration cannot mimic any redshift dependence of the merger timescales in the SPH simulation.
+ Decreasing fa further produces more high mass ealaxies. at the expense of depleting the number of lower mass galaxies. thereby worsening the agreement between the aand SPI galaxy mass functions at the low mass end.," Decreasing $f_{\rm df}$ further produces more high mass galaxies, at the expense of depleting the number of lower mass galaxies, thereby worsening the agreement between the and SPH galaxy mass functions at the low mass end."
+ A second. possible explanation for the cliserepaney tween the galaxy. mass functions is that the relative gas cooling rates in the SPII and mmoclels are cillerent., A second possible explanation for the discrepancy between the galaxy mass functions is that the relative gas cooling rates in the SPH and models are different.
+ We investigate the effects of changing he rate at which gas cools in the numodel. by keeping the core radius of the hot gas density wofile equal to a fixed fraction of the NEW scale radius in he halo.," We investigate the effects of changing the rate at which gas cools in the model, by keeping the core radius of the hot gas density profile equal to a fixed fraction of the NFW scale radius in the halo."
+ This has the effect. of allowing gas to cool more rapidlv in any given halo., This has the effect of allowing gas to cool more rapidly in any given halo.
+ The galaxy mass function of the nminiodel for this case is shown by the dot-dashed line in Fig. 10.., The galaxy mass function of the model for this case is shown by the dot-dashed line in Fig. \ref{fig:cg_altmrg}.
+ The agreement with the SPILL simulation is now much iniproved., The agreement with the SPH simulation is now much improved.
+ In the next section we will explore how changing merger and cooling rates allects the number of galaxies found in halos of à given mass and thus identifv the underlving cause of the dillerences between the mass functions in the SDSA and SPLHI models., In the next section we will explore how changing merger and cooling rates affects the number of galaxies found in halos of a given mass and thus identify the underlying cause of the differences between the mass functions in the SDSA and SPH models.
+ The redshift’ evolution of the mean number of progenitor galaxies per halo in the ACDAL cosmology is shown in Fig. LL.," The redshift evolution of the mean number of progenitor galaxies per halo in the $\Lambda$ CDM cosmology is shown in Fig. \ref{fig:haloNgal},"
+ for two ranges of present-day halo masses: 1077 lothIN (upper panel) and 10575 10Η7M. (lower panel). (, for two ranges of present-day halo masses: $10^{13}$ $10^{14}h^{-1}{\rm M}_\odot$ (upper panel) and $10^{14}$ $10^{15}h^{-1}{\rm M}_\odot$ (lower panel). (
+Qualitatively similar results are obtained for the SCDAL cosmology.),Qualitatively similar results are obtained for the SCDM cosmology.)
+ In the SPL simulation there are. 165 halos in the lower mass range and 12 in the higher mass range., In the SPH simulation there are 165 halos in the lower mass range and 12 in the higher mass range.
+ Since the galaxy. abundance is à steep function. of galaxy mass (c.f, Since the galaxy abundance is a steep function of galaxy mass (c.f.
+ Fig. 7)), Fig. \ref{fig:cgdiff}) )
+ the number of galaxies per halo above some particular mass cut is always dominated by ealaxies close to the cutoll, the number of galaxies per halo above some particular mass cut is always dominated by galaxies close to the cutoff.
+ As a result. the number of galaxies per halo more massive than 64 SPII particles may be allected by differences in the way in which the resolution limit works in SPII and ΠΠ," As a result, the number of galaxies per halo more massive than 64 SPH particles may be affected by differences in the way in which the resolution limit works in SPH and models."
+ Pherefore. we plot. results for galaxies more massive than G4 SPILL gas particles (thin lines). but also for those more massive than 128 SPL gas particles (heavy lines) which are much less alfected by resolution cllects.," Therefore, we plot results for galaxies more massive than 64 SPH gas particles (thin lines), but also for those more massive than 128 SPH gas particles (heavy lines) which are much less affected by resolution effects."
+ The shape of the abundance curves in Fig., The shape of the abundance curves in Fig.
+ Ll ids determined. by the relative rates of galaxy. Formation. aud merging and has the same basic form in the aand SPLE models., \ref{fig:haloNgal} is determined by the relative rates of galaxy formation and merging and has the same basic form in the and SPH models.
+ Phe number of galaxies more massive than 64 SPLIT eas particles predicted by the, The number of galaxies more massive than 64 SPH gas particles predicted by the
+oceurs quickly in GMCs after their formation.,occurs quickly in GMCs after their formation.
+ The fact that clusters with ages greater than S Myr are seldom associated with molecular gas. suggests that clouds disperse quickly (2)..," The fact that clusters with ages greater than 5 Myr are seldom associated with molecular gas, suggests that clouds disperse quickly \citep*{Leisawizetal1989}."
+ In the original cloud lifetime proposition. it was assumed that all of the CO observed in the galaxy was associated with molecular ivdrogen involved in star formation.," In the original cloud lifetime proposition, it was assumed that all of the CO observed in the galaxy was associated with molecular hydrogen involved in star formation."
+ We now realise that the vast majority of the gus that comprises a GMC is never involved in he star formation process., We now realise that the vast majority of the gas that comprises a GMC is never involved in the star formation process.
+ In fact the star formation efficiency in GMCs is only a few percent., In fact the star formation efficiency in GMCs is only a few percent.
+ The reason behind this lies with the ‘act that little of the cloud is actually dense and bound enough to urn into stars in the cloud's lifetime (2?2).," The reason behind this lies with the fact that little of the cloud is actually dense and bound enough to turn into stars in the cloud's lifetime \citep{Padoan1995, Hartmann1998, Zinnecker2002}."
+. Also. if GMCs are short-lived features then there is little time for the more tenuous oar!s of the cloud to get involved in the star formation via accretion.," Also, if GMCs are short-lived features then there is little time for the more tenuous parts of the cloud to get involved in the star formation via accretion."
+ The old GMC model also required the cloud to be supported and in virial equilibrium. since that would permit them to remain as coherent structures for as long as was necessary.," The old GMC model also required the cloud to be supported and in virial equilibrium, since that would permit them to remain as coherent structures for as long as was necessary."
+ This support pressure had to be in the form of non-thermal kinetic energy. such as turbulence (2).. since the thermal energy component of these clouds is typically very small.," This support pressure had to be in the form of non-thermal kinetic energy, such as turbulence \citep{Larson1981}, since the thermal energy component of these clouds is typically very small."
+ To counteraet the gravity on the large scales however requires motions which are supersonic. and it is known that these quickly damp in shocks €?:: 2)). even in the presence of magnetic tields.," To counteract the gravity on the large scales however requires motions which are supersonic, and it is known that these quickly damp in shocks \citealt{Maclowetal1998}; \citealt*{Stoneetal1998}) ), even in the presence of magnetic fields."
+ Thus the bound GMC model requires some method of continually driving the turbulence on the large scale., Thus the bound GMC model requires some method of continually driving the turbulence on the large scale.
+ These driving mechanisms are not necessary in the short cloud lifetime model. and there also is no need to assume that the clouds are in virial equilibrium.," These driving mechanisms are not necessary in the short cloud lifetime model, and there also is no need to assume that the clouds are in virial equilibrium."
+ GMCs can therefore exist in a variety of dynamical states., GMCs can therefore exist in a variety of dynamical states.
+ ? have examined the stability of molecular clouds in the outer galaxy and come to the conclusion that most clouds are indeed globally unbound by their internal motions., \citet*{Heyeretal2001} have examined the stability of molecular clouds in the outer galaxy and come to the conclusion that most clouds are indeed globally unbound by their internal motions.
+ They also point out the difficulty in producing mass estimates (which a great number of papers on the subject of GMCs pass over) and note that even mass estimates determined via CO measurements (both CO and 0) assumes at some stage the cloud is bound., They also point out the difficulty in producing mass estimates (which a great number of papers on the subject of GMCs pass over) and note that even mass estimates determined via CO measurements (both $^{12}$ CO and $^{13}$ CO) assumes at some stage the cloud is bound.
+ Although they do find the clouds approach dynamical stability at large masses €»107M η. here is still considerable scatter in the data. suggesting that there is no typical dynamic state for GMCs.," Although they do find the clouds approach dynamical stability at large masses $>10^{5}$ ), there is still considerable scatter in the data, suggesting that there is no typical dynamic state for GMCs."
+ ? have shown that it might also be possible to build GMCs by accumulating very low density hydrogen gas. already in a molecular state.," \citet*{Pringleetal2001} have shown that it might also be possible to build GMCs by accumulating very low density hydrogen gas, already in a molecular state."
+ Their study came in response to the ideas presented by ?. in an attempt to provide a new mechanism for GMC ormation that can occur quickly.," Their study came in response to the ideas presented by \citet{Elmegreen2000}, in an attempt to provide a new mechanism for GMC formation that can occur quickly."
+ They point out that it is quite yossible that a large fraction of the interstellar medium may be in molecular form. but either simply too low a density to be detectable by current methods or too far away from illuminating sources.," They point out that it is quite possible that a large fraction of the interstellar medium may be in molecular form, but either simply too low a density to be detectable by current methods or too far away from illuminating sources."
+ The GMCs are then formed from large scale shocks. from spiral arm passage or feedback from high mass stars. such as winds and supernovae.," The GMCs are then formed from large scale shocks, from spiral arm passage or feedback from high mass stars, such as winds and supernovae."
+ This cloud formation can occur within a few million years., This cloud formation can occur within a few million years.
+ ?. also point out that GMCs are probably not in. virial equilibrium. and note that their wind-swept appearance suggests that they are anything but.," \citet{Pringleetal2001} also point out that GMCs are probably not in virial equilibrium, and note that their wind-swept appearance suggests that they are anything but."
+ The simulation that we present here draws on the above studies for motivation., The simulation that we present here draws on the above studies for motivation.
+ We assume that large scale flows are able to create an unbound GMC in a few millions years., We assume that large scale flows are able to create an unbound GMC in a few millions years.
+ Instead of being contained by external forces (e.g. 2)). we assume that the cloud is free to expand into the ISM.," Instead of being contained by external forces (e.g. \citealt{Heyeretal2001}) ), we assume that the cloud is free to expand into the ISM."
+ Thus the flows that created the cloud are assumed to have been used up in its formation., Thus the flows that created the cloud are assumed to have been used up in its formation.
+ Since the cloud is assumed to be short lived and not quasi-static. there is no need for the internal turbulent energy. which will dissipate on the crossing time. to be replenished (2?)," Since the cloud is assumed to be short lived and not quasi-static, there is no need for the internal turbulent energy, which will dissipate on the crossing time, to be replenished \citep*{Paredesetal1999,
+Elmegreen2000}."
+ OB associations are historically identified simply as extended groups of OB stars. having diameters of tens of parsecs (?)}.," OB associations are historically identified simply as extended groups of OB stars, having diameters of tens of parsecs \citep{Ambart1955}."
+. Furthermore they are rather more diffuse than open clusters. with the mass density of OB type stars at 0.1 pe ¢2222)..," Furthermore they are rather more diffuse than open clusters, with the mass density of OB type stars at $\sim 0.1$ $^{-3}$ \citep{Blaauw1964, Ambart1955,
+Garmany1994, Ladas2003}."
+" It was found that these associations contain considerable substructure which are referred to as ""OB subgoups’ (?)..", It was found that these associations contain considerable substructure which are referred to as `OB subgoups' \citep{Blaauw1964}.
+ These subgroups are unbound from one another as was deduced from their expansion about the centre of the region (2)., These subgroups are unbound from one another as was deduced from their expansion about the centre of the region \citep{Blaauw1952}.
+". Some regions or ""subgroups? are shown to be associated with molecular gas.", Some regions or `subgroups' are shown to be associated with molecular gas.
+ In general these regions are not coeval but can exhibit a spread of ages between the subgroup population as large as 10 Myr (2).., In general these regions are not coeval but can exhibit a spread of ages between the subgroup population as large as 10 Myr \citep{Blaauw1964}.
+ The fact that OB associations are very young. with some of the subgroups possessing ages of the order of a millions years. suggests that unbound nature of the subgroups from one another is primordial.," The fact that OB associations are very young, with some of the subgroups possessing ages of the order of a millions years, suggests that unbound nature of the subgroups from one another is primordial."
+ The relationship between OB associations and other types of clusters found in the galactic disc. such as open clusters and embedded clusters. is still rather unclear.," The relationship between OB associations and other types of clusters found in the galactic disc, such as open clusters and embedded clusters, is still rather unclear."
+ The OB associations do however have a classic theory regarding their formation., The OB associations do however have a classic theory regarding their formation.
+ ? ooposed that OB associations form via triggering. prompted by he ionised regions produced by previous generations of OB stars.," \citet{ElmegreenLada1977} proposed that OB associations form via triggering, prompted by the ionised regions produced by previous generations of OB stars."
+ In this manner. the star formation is self propagating. with one generations of OB stars triggering the formation of the next.," In this manner, the star formation is self propagating, with one generations of OB stars triggering the formation of the next."
+ Since he shocked layer in which the new group of OB stars forms is moving away from the older OB stars. at a few . the new group is unbound from its parent group.," Since the shocked layer in which the new group of OB stars forms is moving away from the older OB stars, at a few $^{-1}$, the new group is unbound from its parent group."
+ The region then naturally ws the dynamics of the observed OB groups., The region then naturally has the dynamics of the observed OB groups.
+ Motivation came Tom observations of stars forming at the boundaries of molecular clouds and HIT. such as NGC7538. MI7 and M8 (2:: 25).," Motivation came from observations of stars forming at the boundaries of molecular clouds and HII, such as NGC7538, M17 and M8 \citealt*{Habingetal1972}; \citealt{Ladaetal1976}) )."
+ The issue is complicated however when one considers the detailed stellar population of OB associations (22)..," The issue is complicated however when one considers the detailed stellar population of OB associations \citep{Garmany1994, Brown2001}."
+ In the self propagating model. OB type stars form in the shocked layers where conditions are naturally more suited to forming high mass stars.," In the self propagating model, OB type stars form in the shocked layers where conditions are naturally more suited to forming high mass stars."
+ Low mass stars form spontaneously in the rest of the cloud., Low mass stars form spontaneously in the rest of the cloud.
+ Thus the model assumes a two step formation process whereby low mass stars and high mass stars are formed by different mechanisms and in physically separated locations., Thus the model assumes a two step formation process whereby low mass stars and high mass stars are formed by different mechanisms and in physically separated locations.
+ The IMF of the OB associations however do not exhibit this feature and generally possess the standard field star IMF. at least within the Salpeter range (e.g. Sco OB2. ?: ?)).," The IMF of the OB associations however do not exhibit this feature and generally possess the standard field star IMF, at least within the Salpeter range (e.g. Sco OB2, \citealt{deGeus1992}; \citealt{PreibischZinnecker1999}) )."
+ Since up to nearly of star formation is thought to occur in embedded clusters. with a field star IMF and primordial mass segregation (for a discussion see 21). it may be that the formation of OB associations has more in common with standard clustered star formation.," Since up to nearly of star formation is thought to occur in embedded clusters, with a field star IMF and primordial mass segregation (for a discussion see \citealt{Ladas2003}) ), it may be that the formation of OB associations has more in common with standard clustered star formation."
+ ? and ? have modelled cluster formation in a turbulently supported cloud., \citet*{BBV2003} and \citet*{BVB2004} have modelled cluster formation in a turbulently supported cloud.
+ They modelled a molecular cloud that was initially supported against collapse by a turbulent velocity field., They modelled a molecular cloud that was initially supported against collapse by a turbulent velocity field.
+ It was found that the dissipation of the large scale supersonic flows produced a number of distinct subclusters., It was found that the dissipation of the large scale supersonic flows produced a number of distinct subclusters.
+ Each subcluster contains at the core a massive star., Each subcluster contains at the core a massive star.
+ The subclusters were mass segregated and each had a protostellar population consistent with that of the observed field star IMF. both of which are the result of competitive accretion.," The subclusters were mass segregated and each had a protostellar population consistent with that of the observed field star IMF, both of which are the result of competitive accretion."
+ Since the cloud was initially bound. even more so after the dissipation of the turbulent energy. the whole system of subclusters are themselves bound to one another.," Since the cloud was initially bound, even more so after the dissipation of the turbulent energy, the whole system of subclusters are themselves bound to one another."
+ They quickly merge within roughly 0.5 Myr troughly twice the free-fall time for the original cloud)., They quickly merge within roughly 0.5 Myr (roughly twice the free-fall time for the original cloud).
+ Tf this merging process was to occur on large scales. such as a whole GMC. one would never be able to form OB associations.," If this merging process was to occur on large scales, such as a whole GMC, one would never be able to form OB associations."
+ The massive stars at the centres of the subclusters would tind themselves in one large cluster., The massive stars at the centres of the subclusters would find themselves in one large cluster.
+ Our proposal in this paper is that OB associations are just a, Our proposal in this paper is that OB associations are just a
+"in the total image (dotted), arm (solid) and interarm (dashed) regions.","in the total image (dotted), arm (solid) and interarm (dashed) regions."
+" Once again, we find that there is little difference between the arm and interarm region for NGC 5194 and only slight difference for NGC 628."," Once again, we find that there is little difference between the arm and interarm region for NGC 5194 and only a slight difference for NGC 628."
+" However, NGC 6946 showsa an excess of higher SFE(H2) pixels in the arm region (i.e., there are more pixels with a SFE value higher than 6 x 1071? yr-! in the arm regions versus the interarm regions)."," However, NGC 6946 shows an excess of higher $_{2}$ ) pixels in the arm region (i.e., there are more pixels with a SFE value higher than 6 $\times$ $^{-10}$ $^{-1}$ in the arm regions versus the interarm regions)."
+" One would expect that the grand design spirals would show the highest SFE in the arms as opposed to flocculent galaxies, as here the spiral shocks should be strongest."," One would expect that the grand design spirals would show the highest SFE in the arms as opposed to flocculent galaxies, as here the spiral shocks should be strongest."
+" It is interesting then that NGC 6946, the most flocculent galaxy in this study, seems to show a higher SFE in the arms."," It is interesting then that NGC 6946, the most flocculent galaxy in this study, seems to show a higher SFE in the arms."
+" While the source of this is not clear, one possible explanation is that our spiral arm definition is not probing the underlying density enhancement for this galaxy, since it is very weak."," While the source of this is not clear, one possible explanation is that our spiral arm definition is not probing the underlying density enhancement for this galaxy, since it is very weak."
+ Instead the spiral arm mask has isolated regions of high star formation., Instead the spiral arm mask has isolated regions of high star formation.
+" Looking at the spiral arm masks Figure it is clear that the spiral arm structure is (seemuch more 1)),complex than the grand design structures of NGC 628 and NGC 5194."," Looking at the spiral arm masks (see Figure \ref{mask}) ), it is clear that the spiral arm structure is much more complex than the grand design structures of NGC 628 and NGC 5194."
+" If the arms were defined based on young, recent star forming regions, then it would be biased towards high SFRs and hence show seemingly higher SFE in the arm"," If the arms were defined based on young, recent star forming regions, then it would be biased towards high SFRs and hence show seemingly higher SFE in the arm"
+particles will scale as «Νο SeVi.,particles will scale as $N_t$ $N_t$.
+ Thus. there can be a substantial variation in tle computational ti required. between different trees. aud evolving the largest tree cai comprise a significant [racti of the total computational load.," Thus, there can be a substantial variation in the computational time required between different trees, and evolving the largest tree can comprise a significant fraction of the total computational load."
+ Au efficient. parallel code must iancddle this situation well wl dividiug work up amoug NCPU processors., An efficient parallel code must handle this situation well when dividing work up among NCPU processors.
+ Load balancing of trees is achieved iu three ways., Load balancing of trees is achieved in three ways.
+" First. trees a'e sorted by NV, aud then divicex iuto “copses” of roughly equal amounts of weyk using the op[n]eedy algoritli1."," First, trees are sorted by $N_t$ and then divided into “copses” of roughly equal amounts of work using the “greedy” algorithm."
+ That is. starting witu the largest tree. each one in turn is given to the copse which up o that point has been assigiec the least auuou uo ‘total work.," That is, starting with the largest tree, each one in turn is given to the copse which up to that point has been assigned the least amount of total work."
+ Usually there is oue copse per CPU. but tleye Call be two or nore per CPU if required by space coustrains.," Usually there is one copse per CPU, but there can be two or more per CPU if required by space constraints."
+ There is a comunulcation step. wren all data associatec with a copse is sen to one CPU.," There is a communication step, when all data associated with a copse is sent to one CPU."
+ A CPU then evolves each tree in its loca copse in turu. starting with the most massive.," A CPU then evolves each tree in its local copse in turn, starting with the most massive."
+ Additionally. tie largest trees can ye doue in paralel.," Additionally, the largest trees can be done in parallel."
+ If a ew large trees dominate the work load. then it is imj»ossible for all copses to contain eqal amouts of work— ideally. one would want NCPU copses. each comprising 1/NCPL of the total amout of work. but this clearly can’t happen if the largest t‘ee takes a larger (ractiou jus by itsel.," If a few large trees dominate the work load, then it is impossible for all copses to contain equal amounts of work— ideally, one would want NCPU copses, each comprising 1/NCPU of the total amount of work, but this clearly can't happen if the largest tree takes a larger fraction just by itself."
+ In this case. the force calculation for these large trees is dome in parallel by a small uumber of CPUs.," In this case, the force calculation for these large trees is done in parallel by a small number of CPUs."
+ Tis is currently doue quite crudely. with only the tree walk auc force calculation acually ca‘ried out in parallel: iu theory a fully parallel tree code could be used for every tree. allocaling ukre processors to those copses containing the most work.," This is currently done quite crudely, with only the tree walk and force calculation actually carried out in parallel; in theory a fully parallel tree code could be used for every tree, allocating more processors to those copses containing the most work."
+ Ouly a few nodes are usually. recuired to reduce the time spent on the largest tree to the level required for oad balancing., Only a few nodes are usually required to reduce the time spent on the largest tree to the level required for load balancing.
+ As a final means of balancing the load. when a CPU finishes evolving all the trees iu its local copse. it teu sends a sigual to the other CPUs.," As a final means of balancing the load, when a CPU finishes evolving all the trees in its local copse, it then sends a signal to the other CPUs."
+ A CPU with work still eft will seu«c an unevolvect t‘ee to the idle €‘PL for it to carry out the evolution., A CPU with work still left will send an unevolved tree to the idle CPU for it to carry out the evolution.
+ The scaling of the current implementation of TPM with nuuber of processors is shown iu Fig. velfie:scale.., The scaling of the current implementation of TPM with number of processors is shown in $.$ \\ref{fig:scale}.
+" The test case lor these timine rus is a standard LCDM. inodel 1i redshilt 220.16 witl N=256"" particles in a 3204. 'Mpe cube.", The test case for these timing runs is a standard LCDM model at redshift $z$ =0.16 with $N$ $^3$ particles in a $h^{-1}$ Mpc cube.
+ Tus particular set of runs was carried out on an LÀ-61 Linux cluster at NCSA named. “Titan”., This particular set of runs was carried out on an IA-64 Linux cluster at NCSA named “Titan”.
+ TIis inachine consists of 128 nodes. each with dual Inte SOOMHAz Itauium processors. and a Myrine uetwork intercounect.," This machine consists of 128 nodes, each with dual Intel 800MHz Itanium processors, and a Myrinet network interconnect."
+ While the total time require depeuds ou processor performance. similar scaling with NCPU ias been found on a number of uachines with various types of processors axl interconnects.," While the total time required depends on processor performance, similar scaling with NCPU has been found on a number of machines with various types of processors and interconnects."
+ The topiost line in Fig. is total time. calculated as the utuber of secouds walleock lune per PM step multipliec w NCPU: perect scaling would be a horizoital line.," The topmost line in $.$ \\ref{fig:scale}
+ is total time, calculated as the number of seconds wallclock time per PM step multiplied by NCPU; perfect scaling would be a horizontal line."
+ TPM perforiIs (t Hewell— at NCPU=61 the efficiency 15 stl] as compared to NCPU-I., TPM performs quite well— at NCPU=64 the efficiency is still as compared to NCPU=4.
+ Bevyoud this poiu sCaing begius to deerace. with he efficieucy. droppite to for NCPU=128.," Beyond this point scaling begins to degrade, with the efficiency dropping to for NCPU=128."
+ The 'OasO1 TPA] scales well can be seen in the secoud line {ror1 the top. which shows the otal time speit in t‘ee evolution.," The reason TPM scales well can be seen in the second line from the top, which shows the total time spent in tree evolution."
+ This pat of the code takes up most of the CPU time. but it 'equires no communication aid there are eiough trees so that just coarse-erainecl parallelization works reasonably wel.," This part of the code takes up most of the CPU time, but it requires no communication and there are enough trees so that just coarse–grained parallelization works reasonably well."
+ The nest two curves shown indicate the amount of time iu the PM portion of the code aud overlead related to trees (ideutifviug tree regions aud particles. etc.).," The next two curves shown indicate the amount of time in the PM portion of the code and overhead related to trees (identifying tree regions and particles, etc.)."
+ These take a siuall fraction of the total time aud scale well since the grid is distributed across all processors., These take a small fraction of the total time and scale well since the grid is distributed across all processors.
+are the mean field expectation values of the vector mesons ω and ὁ in quark matter. p; is lepton pressure. po is the vacuum pressure and 5 is an effective bag constant.,"are the mean field expectation values of the vector mesons $\omega$ and $\phi$ in quark matter, $p_l$ is lepton pressure, $p_0$ is the vacuum pressure and $B^*$ is an effective bag constant."
+ The quark chemical potentials are modified by the vector fields as follow fe=diag(μμ-Wo.LtWo.flyGo).," The quark chemical potentials are modified by the vector fields as follow $\hat\mu^*={\rm diag}_f(\mu_u-\omega_0,\mu_d-\omega_0,\mu_s-\phi_0)$."
+" The numerical values of the parameters of the Lagrangian are 77,4,=5.5 MeV. m,=140.7 MeV. A=602.3 MeV. G.A?=1.835. KA?=12.36. and Gp/Gs=I."," The numerical values of the parameters of the Lagrangian are $m_{u,d}
+= 5.5$ MeV, $m_s = 140.7$ MeV, $\Lambda = 602.3$ MeV, $G_S\Lambda^2 =
+1.835$, $K\Lambda^5 =12.36$, and $G_D/G_S = 1$."
+ The surface tension between the (hyperinuclear and quark matter is not well-known. therefore we shall adopt the working hypothesis that this tension is high enough to prevent the formation of mixed phases.," The surface tension between the (hyper)nuclear and quark matter is not well-known, therefore we shall adopt the working hypothesis that this tension is high enough to prevent the formation of mixed phases."
+ Then. the transition. from (hyper)nuclear matter to quark matter occurs at a certam baryo-chemical potential at which the pressures of these phases are equal.," Then, the transition from (hyper)nuclear matter to quark matter occurs at a certain baryo-chemical potential at which the pressures of these phases are equal."
+ This is equivalent to the condition that pressure. P. vs. chemical potential. µ. curves for these phases cross (Maxwell construction).," This is equivalent to the condition that pressure, $P$, vs. chemical potential, $\mu$, curves for these phases cross (Maxwell construction)."
+ Thus. according to the Aaxwell construction of the deconfinement phase transition. there is a jump in the density at constant pressure.," Thus, according to the Maxwell construction of the deconfinement phase transition, there is a jump in the density at constant pressure."
+ However. the transition density itself cannot be fixed. because the current NJL model does not allow us to fix the low-density normalization of the pressure: (this is the consequence of the fact that this class of models does not capture the confinement feature of the QCD).," However, the transition density itself cannot be fixed, because the current NJL model does not allow us to fix the low-density normalization of the pressure; (this is the consequence of the fact that this class of models does not capture the confinement feature of the QCD)."
+" For this reason we introduced above an additional ""bag"" parameter B. which allows us to vary the density at which the quark phase sets-in. thus fixing the density of deconfinement f."," For this reason we introduced above an additional “bag” parameter $B^*$, which allows us to vary the density at which the quark phase sets-in, thus fixing the density of deconfinement $\rho_{\rm tr}$."
+ The transition density increases with 5 (as well as with the vector coupling Gy)., The transition density increases with $B^*$ (as well as with the vector coupling $G_V$ ).
+ For example. varying B in the range -40 MeV fm? to 50 MeV fm we find variations in the transition density in the range 2.4po«pu<dpo. where po is the nuclear saturation density.," For example, varying $B^*$ in the range -40 MeV $^{-3}$ to 50 MeV $^{-3}$ we find variations in the transition density in the range $2.4\rho_0<\rho_{\rm tr}< 4 \rho_0$, where $\rho_0$ is the nuclear saturation density."
+ Before turning to the results of calculations. we briefly summarize the physical input of the underlying models.," Before turning to the results of calculations, we briefly summarize the physical input of the underlying models."
+ The masses and coupling in the Lagrangian (1)) are constrained by the phenomenology of nuclear and hyper-nuclear matter; we will use the NL3 parameterization throughout., The masses and coupling in the Lagrangian \ref{eq:L_RMF}) ) are constrained by the phenomenology of nuclear and hyper-nuclear matter; we will use the NL3 parameterization throughout.
+" The parameters of the Lagrangian (2)) are fixed to reproduce the observables of vacuum QCD: the pion mass ii,=135 MeV. the kaon mass my=397.7 MeV. the eta-prime mass 7;=957.8 MeV and the pion-decay constant f£;292.4 MeV. There remain two free parameters Gy and B7."," The parameters of the Lagrangian \ref{eq:NJL_Lagrangian}) ) are fixed to reproduce the observables of vacuum QCD: the pion mass $m_{\pi} = 135$ MeV, the kaon mass $m_K = 497.7$ MeV, the eta-prime mass $m_{\eta'} = 957.8$ MeV and the pion-decay constant $f_{\pi} = 92.4$ MeV. There remain two free parameters $G_V$ and $B^*$."
+ The first one can be used to regulate the stiffness of the quark matter equation of state (Gy=0 gives the softest equation of state). while B* can be varied to change the density of the phase transition to the quark phase (deconfinement).," The first one can be used to regulate the stiffness of the quark matter equation of state $G_V=0$ gives the softest equation of state), while $B^*$ can be varied to change the density of the phase transition to the quark phase (deconfinement)."
+ The contributions to the pressure from the condensates arises from two phases: the two-flavor 2SC phase. where A;=Ax0 and A;#O and the three-flavor CFL phase. where A;#0. A»z0. and Axz0.," The contributions to the pressure from the condensates arises from two phases: the two-flavor 2SC phase, where $\Delta_1 = \Delta_2 = 0$ and $\Delta_3 \neq 0$ and the three-flavor CFL phase, where $\Delta_1 \neq
+0$, $\Delta_2 \neq 0$, and $\Delta_3 \neq 0$."
+" The strength of the pairing field is fixed by setting Gp=Gs. which implies stronger pairing interaction than the one deduced from the Fierz transformation Gy,=0.75Gs."," The strength of the pairing field is fixed by setting $G_D = G_S$, which implies stronger pairing interaction than the one deduced from the Fierz transformation $G_D = 0.75 G_S$."
+ Figure 1. shows the equations of state of (hyper)nuclear and quark phases and their matehing via Maxwell construction for several values of the parameters Gy and py., Figure \ref{eos} shows the equations of state of (hyper)nuclear and quark phases and their matching via Maxwell construction for several values of the parameters $G_V$ and $\rho_{\rm tr}$.
+ For low values of the Gy parameter there are two sequential transitions from (hyper)nuclear matter to the 2SC phase and from 2SC phase to the CFL phase., For low values of the $G_V$ parameter there are two sequential transitions from (hyper)nuclear matter to the 2SC phase and from 2SC phase to the CFL phase.
+ As the onset density of the quark matter is shifted to higher densities. the CFL phase is expelled from the density range relevant for the phenomenology.," As the onset density of the quark matter is shifted to higher densities, the CFL phase is expelled from the density range relevant for the phenomenology."
+ For higher values of the Gy parameter the density jump(s) are smaller. therefore the CFL phase can nucleate at lower densities.," For higher values of the $G_V$ parameter the density jump(s) are smaller, therefore the CFL phase can nucleate at lower densities."
+ This tendency eventually leads to a smooth transition between the phases for Gy=0.6G;. if the transition density 1s not too high.," This tendency eventually leads to a smooth transition between the phases for $G_V = 0.6 G_S$, if the transition density is not too high."
+ The € Hooft term makes the transition from the 2SC to the CFL phase “smoother”., The t' Hooft term makes the transition from the 2SC to the CFL phase “smoother”.
+ The equations of state described above were supplemented by an equation of state of the low-density crust (Baymetal.1971).., The equations of state described above were supplemented by an equation of state of the low-density crust \citep{1971ApJ...170..299B}.
+ These were used as an input into the Oppenheimer-Volkoff equations to obtain sequences of stellar configurations., These were used as an input into the Oppenheimer-Volkoff equations to obtain sequences of stellar configurations.
+ The resulting mass-radius relationship for massive stars is shown in Fig., The resulting mass-radius relationship for massive stars is shown in Fig.
+ 2. together with the largest mass measurement to date M=1.97+0.04Moe (Demorestetal.2010).., \ref{fig:2} together with the largest mass measurement to date $M= 1.97\pm 0.04 M_{\sun}$ \citep{2010Natur.467.1081D}.
+ Masses above the lower bound on the maximum mass are obtained for purely hadronic stars: this feature is prerequisite for finding similar stars with quark phases., Masses above the lower bound on the maximum mass are obtained for purely hadronic stars; this feature is prerequisite for finding similar stars with quark phases.
+ Evidently only for high values of vector coupling Gy one finds stable stars that contain (at the bifurcation from the hadronic sequence) the 2SC phase. which are followed by stars that additionally contain the CFL phase (for higher central densities).," Evidently only for high values of vector coupling $G_V$ one finds stable stars that contain (at the bifurcation from the hadronic sequence) the 2SC phase, which are followed by stars that additionally contain the CFL phase (for higher central densities)."
+ Thus we find that the stable branch of the sequence contains stars with quark matter in the 2SC and CFL phases., Thus we find that the stable branch of the sequence contains stars with quark matter in the 2SC and CFL phases.
+ Similar results were obtained from the GM3 parameterization of nuclear matter 1f. however. the hyperons were excluded from the consideration.," Similar results were obtained from the GM3 parameterization of nuclear matter if, however, the hyperons were excluded from the consideration."
+ Therefore. a less stiff GM3 parameterization is qualitatively equivalent to the NL3 parametrization. which features the softening of the equation of state due to the hyperons.," Therefore, a less stiff GM3 parameterization is qualitatively equivalent to the NL3 parametrization, which features the softening of the equation of state due to the hyperons."
+ Figure 3. summarizes several results for the masses and composition of compact stars as a function of the parameters of the model Gy and py., Figure \ref{fig:3} summarizes several results for the masses and composition of compact stars as a function of the parameters of the model $G_V$ and $\rho_{\rm tr}$.
+ First. it shows the tracks of constant maximum mass compact stars within the parameter space.," First, it shows the tracks of constant maximum mass compact stars within the parameter space."
+ The decrease of these lines with increasing vector coupling reflects the fact that non-zero vector coupling stiffens the equation state., The decrease of these lines with increasing vector coupling reflects the fact that non-zero vector coupling stiffens the equation state.
+ It other words. to obtain a given maximum mass one can admit à small amount of soft quark matter with vanishing vector coupling by choosing a high transition. density: the same result is obtained with a low transition density. but strong vector coupling.{οι a stiffer quark equation of state.," In other words, to obtain a given maximum mass one can admit a small amount of soft quark matter with vanishing vector coupling by choosing a high transition density; the same result is obtained with a low transition density, but strong vector coupling, a stiffer quark equation of state."
+ For low transition densities one finds 2SC matter in. stars. which means that weaker vector couplings slightly disfavor 2SC matter.," For low transition densities one finds 2SC matter in stars, which means that weaker vector couplings slightly disfavor 2SC matter."
+ Substantial CFL cores appear in configurations for strong vector coupling and almost independent of the transition density (nearly vertical dashed lines with ὁ~0.1 in Fig. 3))., Substantial CFL cores appear in configurations for strong vector coupling and almost independent of the transition density (nearly vertical dashed lines with $\delta\sim 0.1$ in Fig. \ref{fig:3}) ).
+ Note that for a high transition density there ts a direct transition from hyper-nuclear to the CFL phase., Note that for a high transition density there is a direct transition from hyper-nuclear to the CFL phase.
+ For transition densities blow 3.5pg a 2SC layer emerges that separates these phases., For transition densities blow $3.5\rho_0$ a 2SC layer emerges that separates these phases.
+ On the other hand. weak vector couplings and low transition densities produce stars with a 2SC phase only.," On the other hand, weak vector couplings and low transition densities produce stars with a 2SC phase only."
+ In Fig., In Fig.
+ 4. we show the baryon density profiles of the maximum-mass stars with a transition density of py= 2.506.," \ref{fig:4} we show the baryon density profiles of the maximum-mass stars with a transition density of $\rho_{\rm tr}=2.5
+\rho_0$ ."
+"the central sources (an accreting black hole). consisting of silicate and eraphite grains. confined in a toroidal shape. as opposed. το ""chlumpy tori models([SCC CLE. Nenkovaet 2)).","the central sources (an accreting black hole), consisting of silicate and graphite grains, confined in a toroidal shape, as opposed to “clumpy” tori models (see e.g. \citealt{nenkova02}) )."
+ For a detailed description of the model we refer to Fritzetal.(2).., For a detailed description of the model we refer to \cite{fritz06}.
+ Here. however. we summarize some of the major characteristies of the mocel.," Here, however, we summarize some of the major characteristics of the model."
+ The central source is assumed to be point-like and its emission isotropic., The central source is assumed to be point-like and its emission isotropic.
+ Its spectral energy distribution is define by means of à composition of power laws (ic. AL(A)x A‘) with cillerent values for the spectral index ;.., Its spectral energy distribution is defined by means of a composition of power laws (i.e. $\lambda L(\lambda) \propto \lambda^i$ ) with different values for the spectral index $i$.
+ We adopte values found in the literature (e.g. Granato&Danese1994 or Schartmannetal.2005)). namely: ¢=1.2 for 0.001< 0.03.)=O for 0.08<A« 0.125. /—OS andi=1.0 or 0.125«A10. and a Iavleigh-Jeans decline (;= 3.0) is assumed. longward of A=10. where A is given in.," We adopted values found in the literature (e.g. \citealt{granato94} or \citealt{schartmann05})), namely: $i=1.2$ for $0.001 < \lambda < 0.03$ , $i=0$ for $0.03 < \lambda < 0.125$ , $i=-0.5$ and $i=-1.0$ for $0.125 < \lambda < 10$, and a Rayleigh-Jeans decline $i=-3.0$ ) is assumed longward of $\lambda =10$, where $\lambda$ is given in."
+. The torus geometry adopted to describe the shape anc he spatial cistribution of dust is the so-called “Harect cisk (sce e.g. Efstathiou&Rowan-Robinson1995... Manskeeal.1998. and vanDemmel&Dullemoned 2003)). that. is a sphere with the polar cones removed.," The torus geometry adopted to describe the shape and the spatial distribution of dust is the so-called “flared disk” (see e.g. \citealt{efstathiou95}, , \citealt{manske98} and \citealt{vanbemmel03}) ), that is a sphere with the polar cones removed."
+ Is size is define w its outer radius. Row. and the opening angle. QO. of the orus itself.," Its size is defined by its outer radius, $_{out}$, and the opening angle, $\Theta$, of the torus itself."
+ The dust components that dominate both the absorption and the emission of radiation are graphite anc silicate., The dust components that dominate both the absorption and the emission of radiation are graphite and silicate.
+" The location of the inner racius. τωclepencls Mon thesublimation temperature of the dust. grains ""a(150t ancl 1000 IW. for graphite. (Darvainis1987) and"," The location of the inner radius, $_{in}$ , depends bothon thesublimation temperature of the dust grains (1500 and 1000 K, for graphite \citep{barvainis87} and silicate"
+We selected the old systems LP. 100-22. (kawkactal.2006:INilicetal.2009:Veunes 2000)..NLTT 11718 (Ikiwka&Veunes2009:Steiufadtetal.2010:hawkaetal.2010). and a sample of low-mass white dwarfs iu the SDSS survey (ντκανά&vanIkerkwijk2010:Brownetal2010:ilieal. 2011a.b)..,"We selected the old systems LP 400-22 \citep{kaw2006,kil2009,ven2009}, NLTT 11748 \citep{kaw2009,ste2010,kaw2010} and a sample of low-mass white dwarfs in the SDSS survey \citep{kul2010,bro2010,kil2011a,kil2011b}. ."
+ The ELM precursor TID 188112 Ueberetal.2003) is also shown with a pre-flash model describing low-1uass subdwarf stars (ποσοetal.1998)., The ELM precursor HD 188112 \citep{heb2003} is also shown with a pre-flash model describing low-mass subdwarf stars \citep{dri1998}.
+". The evolutionary tracks cover the population with masses rauging from ΟΛΗ, Πο z0.15..", The evolutionary tracks cover the population with masses ranging from $\ga 0.15$ to $\approx0.45$.
+. The ELM white dwarfs (M= 027.3) occupy a relatively hieb-Tuninosity region of the diagram characterized by residual burning., The ELM white dwarfs $M\la0.2$ ) occupy a relatively high-luminosity region of the diagram characterized by residual burning.
+" Althausetal.(2001) estimates that. taking iuto account the effect of diffusion. objects with masses Z01514, wwill experience ao thermonuclear fash that delays their cooling."," \citet{alt2001} estimates that, taking into account the effect of diffusion, objects with masses $\ga0.18$ will experience a thermonuclear flash that delays their cooling."
+ The effect is also metallicity dependent. with ligher-inetallicity objects expericucing a flash at a comparatively lower mass.," The effect is also metallicity dependent, with higher-metallicity objects experiencing a flash at a comparatively lower mass."
+ Independent age aud mass estimates are possible iu binaries comprising au ELM white dwarf and a pulsar (seeAlthausctal.2001) but are more difficult to obtain in double white dwarf systems., Independent age and mass estimates are possible in binaries comprising an ELM white dwarf and a pulsar \citep[see][]{alt2001} but are more difficult to obtain in double white dwarf systems.
+ The ELM GALEN J1717|6757 is one of the vounecst known class iienibers. but the estimated age of the system remains mocdol-«depenudeut with the metallicity of the progenitors still uucoustrained.," The ELM GALEX J1717+6757 is one of the youngest known class members, but the estimated age of the system remains model-dependent with the metallicity of the progenitors still unconstrained."
+" The ELM. primary in GALEN J1717|6757. with au estimated radius Ry0.1 falls in oulv about of its Roche volume (R4, --0.13run... resulting in. low-auplitude ellipsoidal variations compounded with a dominant relativistic beamine effect."," The ELM primary in GALEX J1717+6757, with an estimated radius $R_1=0.1$ fills in only about of its Roche volume \citep[$R_{1,L} =, resulting in low-amplitude ellipsoidal variations compounded with a dominant relativistic beaming effect."
+ The system is expected to imnerge (Z4) In 7.0 Car assunmnius a secondary mnass Mo=0.9 citepritl986.. and the product of the merger will likely be an ultraanassive white dwarf with Mz1.1nusuli..," The system is expected to merge $t_{\rm merge}$ ) in 7.0 Gyr assuming a secondary mass $M_2=0.9$ \\citep{rit1986}, and the product of the merger will likely be an ultra-massive white dwarf with $M\approx 1.1$."
+ Using a distauco d=195+7 pe (Section 3.1). and the velocity. and proper motion nueasurenients we computed the (VT) velocity vector following Johuson&Soderblom(1987): The velocity componcuts are characteristic of the thin Galactic disc and similar to the white chwart population in general (Sionoetal.1988:Pauliet 2006).," Using a distance $d=195\pm7$ pc (Section 3.1), and the velocity and proper motion measurements we computed the $U,V,W$ ) velocity vector following \citet{joh1987}: The velocity components are characteristic of the thin Galactic disc and similar to the white dwarf population in general \citep{sio1988,pau2006}."
+. We conclude that uulike the old svsteius LP. £00- and NLTT 117185 that show peculiar kinematics. the svsteii CALEX JI7TI7|6757 is kinematically voung.," We conclude that unlike the old systems LP 400-22 and NLTT 11748 that show peculiar kinematics, the system GALEX J1717+6757 is kinematically young."
+ The calculation of additional evolutionary trackscovering the observed rauge of paraicters. particularly between 0.17 and 0.19 AF... would benefit studies of ELM white dwats," The calculation of additional evolutionary trackscovering the observed range of parameters, particularly between 0.17 and 0.19 would benefit studies of ELM white dwarfs"
+"For each real shift. Ox, (and Ovi. We will simplify the notation by referring to both axis directions with . when possible). we calculate the mean and standard deviation. cr. of the measured shifts. Ovyoasured. by use of the outlier-robust statistics based on Tukey's Biweight as implemented in the [DL Astronomy User's Library (Landsman1993)..","For each real shift, $\delta x_\text{real}$ (and $\delta
+y_\text{real}$, we will simplify the notation by referring to both axis directions with $x$ when possible), we calculate the mean and standard deviation, $\sigma$, of the measured shifts, $\delta
+x_\text{measured}$, by use of the outlier-robust statistics based on Tukey's Biweight as implemented in the IDL Astronomy User's Library \citep{1993ASPC...52..246L}."
+ After removing dc outliers. we fit the data to the relationship where ¢=27/1 pixel. by use of the robust nonlinear LS curve fitting procedure MPFIT (Moré1978:Markwardt2009).," After removing $4\sigma$ outliers, we fit the data to the relationship where $a=2\pi/1$ pixel, by use of the robust nonlinear LS curve fitting procedure MPFIT \citep{more78levenberg,2009ASPC..411..251M}."
+ The linear coefficient. pj. of these fits is listed in the tables below.," The linear coefficient, $p_1$, of these fits is listed in the tables below."
+" In 2.. the SDF. ADF"". CFL. and CFF CAs are compared for the case of no noise and no bias mismatch. using the 2QI LA and 24x24-pixel subfields."," In \ref{fig:perfect-scatter-four-methods}, the SDF, $^2$, CFI, and CFF CAs are compared for the case of no noise and no bias mismatch, using the 2QI IA and $\times$ 24-pixel subfields."
+ The errors are a mix of systematic and random errors., The errors are a mix of systematic and random errors.
+" The SDF and ADF"" algorithms give a tighter correlation between the true and measured image shifts than the CFI and CFF algorithms.", The SDF and $^2$ algorithms give a tighter correlation between the true and measured image shifts than the CFI and CFF algorithms.
+ November&Simon(1988) found similar undulating effect of small errors near whole and half pixel shifts calculated by CFF and larger errors in between., \citet{1988ApJ...333..427N} found similar undulating effect of small errors near whole and half pixel shifts calculated by CFF and larger errors in between.
+ Ballesterosetal.(1996) make a similar observation with ADF., \citet{ballesteros96two-dimensional} make a similar observation with ADF.
+ The ADF and CFF methods produce many outliers. which is the main reason we need the robust statistics mentioned 1 Sect. 3.2..," The ADF and CFF methods produce many outliers, which is the main reason we need the robust statistics mentioned in Sect. \ref{sec:perfect-processing}."
+ A small fraction of these outliers. on the order 107? of all cases for CFF. zero for all other CAs. are caused by the correlation matrix minimum falling on an outer row or column.," A small fraction of these outliers, on the order $10^{-5}$ of all cases for CFF, zero for all other CAs, are caused by the correlation matrix minimum falling on an outer row or column."
+ The rest are caused by false minima in the correlation matrix. which happen to be deeper than the real minimum.," The rest are caused by false minima in the correlation matrix, which happen to be deeper than the real minimum."
+ This can happen occasionally if à secondary minimum ts located oi or near a grid point. while the real minimum its between eric points.," This can happen occasionally if a secondary minimum is located on or near a grid point, while the real minimum is between grid points."
+ The effect is more severe for ADF because of the more pointy shape of its minimum., The effect is more severe for ADF because of the more pointy shape of its minimum.
+ For CFF. minima corresponding to large shifts are attenuated by two effects: apodization lowers the intensity away from the center of the subimage and the digital Fourier transform wraps in mismatching structure from the other side of the subimage.," For CFF, minima corresponding to large shifts are attenuated by two effects: apodization lowers the intensity away from the center of the subimage and the digital Fourier transform wraps in mismatching structure from the other side of the subimage."
+ This can lead to detection of false minima corresponding to small shifts., This can lead to detection of false minima corresponding to small shifts.
+ ((1990)) used GO? for their thorough: kinematical analvsis of the optical euissiou lines.,\cite{mgoetz}) ) used $60\degr$ for their thorough kinematical analysis of the optical emission lines.
+ Other angles are also used. AAchtermann Lacy (1995)) use 707 in the analysis of their data.," Other angles are also used, Achtermann Lacy \cite{achtermann}) ) use $70\degr$ in the analysis of their data."
+ We re-determined the position of the kinematical major axis of the molecular gas frou our data cube together with the position of the dynamical ceutre (see Fie. 5))., We re-determined the position of the kinematical major axis of the molecular gas from our data cube together with the position of the dynamical centre (see Fig. \ref{fig:velfi}) ).
+" The of this axis is 757. 10"" differeut from that of the outer cüsk."," The of this axis is $75\degr$, $10\degr$ different from that of the outer disk."
+" The ανασα ceutre is found to coincide with the AZ jan nucleus fa=0951""EzL. à=69*55'0070 DB1950.0. Tov et citejov))."," The dynamical centre is found to coincide with the $\ \lambda 2\,\mu$ m nucleus $\alpha = 
+09^h51^m43\fs4$, $\delta = 69\degr55'00\farcs0$ B1950.0, Joy et \\cite{joy}) )."
+ Fie., Fig.
+ 6 shows the positiou-velocity (p-v) diagrauuue taken along the major axis (PA= 757)., \ref{fig:posvel} shows the position-velocity (p-v) diagramme taken along the major axis $PA=75\degr$ ).
+ Emission cau be traced frou 190 + to |130 !., Emission can be traced from $-190$ $^{-1}$ to $+130$ $^{-1}$.
+ This implies a rotation speed —- of 165/sin(/) +. exactly the same figure as derived bv Shen Lo (1995)).," This implies a rotation speed – – of $165/\sin(i)$ $^{-1}$, exactly the same figure as derived by Shen Lo \cite{shen}) )."
+" We would then obtain the same indicative mass for the central ~800 pe reeion of as Shen Lo (7«105 ML, ).", We would then obtain the same indicative mass for the central $\sim800$ pc region of as Shen Lo $7\times10^8$ $_{\odot}$ ).
+ Although the PCO brightness distribution is very asvuuuetric. the main pattern of the p-v diaeramuuc and the ridge of teriniual frequencies are rather svuuetric with respect to the dynamical centre.," Although the $^{13}$ CO brightness distribution is very asymmetric, the main pattern of the p-v diagramme and the ridge of terminal frequencies are rather symmetric with respect to the dynamical centre."
+ We note that chussion from gas at lieh velocities is detected on both sides of the uucleus. although with a higher iuteusitv for SU<0.," We note that emission from gas at high velocities is detected on both sides of the nucleus, although with a higher intensity for $x<0$."
+ The rotation curve derived from the terminal niethod appliedto p-v diagramuue 1s very steep: we estimate ¢yor~LlOkkinss + at +~ Thppe., The rotation curve derived from the terminal-velocities method appliedto p-v diagramme is very steep: we estimate $v_{rot}\sim 140$ $^{-1}$ at $r\sim75$ pc.
+" From r= TMppe tor= 300ppc the estimated rotation curve stays flat (6,4,—LlOkkiiss 5) and reaches a relative", From $r=75$ pc to $r=300$ pc the estimated rotation curve stays flat $v_{rot}=140$ $^{-1}$ ) and reaches a relative
+channels which are contoured in Fig.,channels which are contoured in Fig.
+ 4(d) are the same as in Fig., \ref{fig:obs_M31} are the same as in Fig.
+ 4(b) and Fig. 4(c).., \ref{fig:synth_blur1_M31} and Fig. \ref{fig:synth_blur3_M31}.
+ A good general agreement is seen between the synthetic and real data whereby the galaxy's main outer disc is traced as the velocity channels increase., A good general agreement is seen between the synthetic and real data whereby the galaxy's main outer disc is traced as the velocity channels increase.
+ The synthetic data with the least blurring (i.e. the highest spatial resolution) show structure associated with the spiral shock in the disc., The synthetic data with the least blurring (i.e. the highest spatial resolution) show structure associated with the spiral shock in the disc.
+ This structure is not particularly evident in the observed data for M31., This structure is not particularly evident in the observed data for M31.
+ This may be because of lack of resolution (it is also more difficult to distinguish the spiral structure in our, This may be because of lack of resolution (it is also more difficult to distinguish the spiral structure in our
+"In this section, we analyze when the first major merger that produces the black hole in the galaxy is predicted to occur.","In this section, we analyze when the first major merger that produces the black hole in the galaxy is predicted to occur."
+" In Fig. 6,,"," In Fig. \ref{fig:firstmerge},"
+ we plot the distribution of the times of the first major merging events for galaxies with different present-day stellar masses., we plot the distribution of the times of the first major merging events for galaxies with different present-day stellar masses.
+ The vertical dashed lines indicate the median values of the distributions., The vertical dashed lines indicate the median values of the distributions.
+ The dotted bar in each panel indicates the fraction of galaxies in each stellar mass bin that have not experienced a major merger and are hence not included in the distribution of merging times., The dotted bar in each panel indicates the fraction of galaxies in each stellar mass bin that have not experienced a major merger and are hence not included in the distribution of merging times.
+" As can be seen, massive galaxies experience their first major merging event at earlier epochs than less massive galaxies."," As can be seen, massive galaxies experience their first major merging event at earlier epochs than less massive galaxies."
+ Almost no new black holes form in massive galaxies at the present day., Almost no new black holes form in massive galaxies at the present day.
+ The distribution of black hole formation times in low mass galaxies is much flatter., The distribution of black hole formation times in low mass galaxies is much flatter.
+" If the bulk of the black hole mass is built up in a short period following the first major merger, this would explain why present-day massive black holes have stopped growing, while low mass black holes are still growing at a significant rate (7).."," If the bulk of the black hole mass is built up in a short period following the first major merger, this would explain why present-day massive black holes have stopped growing, while low mass black holes are still growing at a significant rate \citep{heckman2004}."
+ We now assume that the black holes formed from the first major mergers of galaxies can shine and be observed for 107 years., We now assume that the black holes formed from the first major mergers of galaxies can shine and be observed for $10^7$ years.
+" By counting the numbers of such events at different redshifts, we can compute the evolution in the number density of newly formed black holes."," By counting the numbers of such events at different redshifts, we can compute the evolution in the number density of newly formed black holes."
+ The result is plotted as diamonds in Fig. 7.., The result is plotted as diamonds in Fig. \ref{fig:BHrate}.
+" The comoving number density of such events peaks at redshift of z ~2—3, consistent with the observed peak in the number density of bright quasars (?).."," The comoving number density of such events peaks at redshift of z $\sim2-3$, consistent with the observed peak in the number density of bright quasars \citep{richards2006}."
+" The decrease in the number density of newly formed black holes to high redshifts is less pronounced than that found by ?,,"," The decrease in the number density of newly formed black holes to high redshifts is less pronounced than that found by \citet{fan2001},"
+where the Zi are the ionization energies of the poto absorbing species in temperature units.,"where the $T_A$ are the ionization energies of the photo absorbing species in temperature units, and,"
+temperature anisotropy.,temperature anisotropy.
+ In only one of the cases exeanined. (hat of (A=107.5»40°) shown in Figure 5. did both Ryfa=0.3) and. ρα=0.5) exceed unity by an amount that was more (han twice the adopted error.," In only one of the cases examined, that of $(\lambda = 10^{\circ}, \beta = 40^{\circ})$ shown in Figure 5, did both $R_T(a=0.3)$ and $R_T(a=0.5)$ exceed unity by an amount that was more than twice the adopted error."
+ For this direction. we caleulate τρία=0.3)1.4520.08 and Rp(a=0.5)1.23+0.06.," For this direction, we calculate $R_T(a=0.3) = 1.45 \pm 0.08$ and $R_T(a=0.5) = 1.23 \pm 0.06$."
+ The errors were calculated [rom the dispersion in the measurements of 7 given in RedfieldandLinsky(2004).. in à manner similar to that used in section 4.1.1.," The errors were calculated from the dispersion in the measurements of $T$ given in \cite{Redfield04}, in a manner similar to that used in Section 4.1.1."
+ These values for ἐς would correspond (o anisotropy parameters e=0.34—0.44., These values for $R_T$ would correspond to anisotropy parameters $\epsilon = 0.34 - 0.44$.
+ We do not claim this direction as a case for temperature anisotropy because the data shown in Figure 5 do not show clear adherence to Equation (3)., We do not claim this direction as a case for temperature anisotropy because the data shown in Figure 5 do not show clear adherence to Equation (3).
+ For the other directions. the data indicate Hj factors closer to unity. aud. anisotropy [actors closer to zero.," For the other directions, the data indicate $R_T$ factors closer to unity, and anisotropy factors closer to zero."
+ We thus find no evidence in the data for an anisotropy in the sense T>7). although (as illustrated by the discussion in the previous paragraph) we cannot exclude the possibility that €<0.40 could be present. but hidden by random variations in T [rom one line of sight to another.," We thus find no evidence in the data for an anisotropy in the sense $T_{\perp} > T_{\parallel}$, although (as illustrated by the discussion in the previous paragraph) we cannot exclude the possibility that $\epsilon \leq 0.40$ could be present, but hidden by random variations in $T$ from one line of sight to another."
+ An upper limit to (he temperature anisotropy e<0.40 corresponds to TXL6T. which is considerably less (han that reported for the solar corona. and also less than many cases reported in the solar wind 2009).," An upper limit to the temperature anisotropy $\epsilon \leq 0.40$ corresponds to $\frac{T_{\perp}}{T_{\parallel}} \leq 1.67$, which is considerably less than that reported for the solar corona, and also less than many cases reported in the solar wind \citep[see Figure 1 of ][]{Kasper09}."
+. The analvsis of Section 4.1 was done in an unbiased fashion. ie. with no a-priori estimate of the local direction of the interstellar magnetic field.," The analysis of Section 4.1 was done in an unbiased fashion, i.e. with no a-priori estimate of the local direction of the interstellar magnetic field."
+ No direction examined had a compelling case for anisotropy of the turbulent amplitude € or (he ion temperature 7., No direction examined had a compelling case for anisotropy of the turbulent amplitude $\xi$ or the ion temperature $T$.
+" With (his analvsis completed. we then re-examined the data for ""prelerred candidate directions b advocated bv Lallementetal(2005) and Ophleretal(2009)... as cliscussed in Section 4.1 above."," With this analysis completed, we then re-examined the data for “preferred” candidate directions $\hat{b}$ advocated by \cite{Lallement05} and \cite{Opher09}, as discussed in Section 4.1 above."
+" Lallementetal(2005) propose a direction of the local interstellar magnetic field οἱ 2057<A<2407.—60""«€3x—38""."," \cite{Lallement05} propose a direction of the local interstellar magnetic field of $205^{\circ} \leq \lambda \leq 240^{\circ}, -60^{\circ} \leq \beta \leq -38^{\circ}$."
+ Taking the means for each coordinates. we have a candidate field direction for Lallementetal(2005) of (A=222.55.9—49*).," Taking the means for each coordinates, we have a candidate field direction for \cite{Lallement05} of $(\lambda=222.5^{\circ},\beta=-49^{\circ})$."
+ Since —b serves equally well as a direction for anisotropy. we have (A=42.5°.949°) as a candidate direction lor the local field.," Since $-\hat{b}$ serves equally well as a direction for anisotropy, we have $(\lambda=42.5^{\circ},\beta=49^{\circ})$ as a candidate direction for the local field."
+ It is interesting that this direction is verv close to the sel. of directions chosen as the best candidates [rom the unbiased. analvsis of Section 4.1. and illustrated in Figure 2.," It is interesting that this direction is very close to the set of directions chosen as the best candidates from the unbiased analysis of Section 4.1, and illustrated in Figure 2."
+ Although the weak. if not nonexistent. evidence of anisotropy in Figure 2 precludes any futher claims. the coincidence of (he set of directions chosen [ου closer examination aud the proposed direction of the local field of Lallementοἱal(2005) could motivate future investigations with more lines of sight.," Although the weak, if not nonexistent, evidence of anisotropy in Figure 2 precludes any further claims, the coincidence of the set of directions chosen for closer examination and the proposed direction of the local field of \cite{Lallement05} could motivate future investigations with more lines of sight."
+" For comparison, the Dyer-Roeder approximation with a=0.99 is also presented (when a=1 the Dyer-Roeder approximation reduces to the standard Friedmann relation for the distance)."," For comparison, the Dyer–Roeder approximation with $\alpha = 0.99$ is also presented (when $\alpha=1$ the Dyer–Roeder approximation reduces to the standard Friedmann relation for the distance)."
+ As seen 6p is of negligible amplitude ~10”, As seen $\delta_D$ is of negligible amplitude $\sim 10^{-3}$.
+" Thus, for this configuration, and under considered assumptions,3. there is no need to take into account inhomogeneities."," Thus, for this configuration, and under considered assumptions, there is no need to take into account inhomogeneities."
+" Also, another important fact is that dp in the weak lensing approximation is negative, while the Dyer-Roeder approximation implies positive dp."," Also, another important fact is that $\delta_D$ in the weak lensing approximation is negative, while the Dyer–Roeder approximation implies positive $\delta_D$."
+ This is a consequence of positive density fluctuations along the line of sight., This is a consequence of positive density fluctuations along the line of sight.
+" is (6)ip=6x107? and the standard deviation 2x 107?, which implies dp=—5x107!€3107°."," is $\av{\delta}_{1D} = 6 \times 10^{-3}$ and the standard deviation $2 \times 10^{-2}$ , which implies $\delta_D = - 5 \times 10^{-4} \pm 3 \times 10^{-3}$."
+" Thus, as pointed out by Rássánnen (2009) in order toproperly"," Thus, as pointed out by Rässännen (2009) in order toproperly"
+The results discussed in this paper may be summarized as follows: a) From the Fe Ίνα edge à value of clpo=l.7 in number with respect to solar (using the set of values given by Anders Cirevesse 1989) has been found. assuming the primary illaminating continuum measured by BeppoSAX (Matt οἱ al.,"The results discussed in this paper may be summarized as follows: a) From the Fe $\alpha$ edge a value of $A_{\rm Fe}$ =1.7 in number with respect to solar (using the set of values given by Anders Grevesse 1989) has been found, assuming the primary illuminating continuum measured by BeppoSAX (Matt et al."
+ 1999)., 1999).
+ L£the slope of the primary power law ts left free to vary. a steeper spectrum and a lower iron abundance (about 1.2) is found.," If the slope of the primary power law is left free to vary, a steeper spectrum and a lower iron abundance (about 1.2) is found."
+ b) From the flux ratio of the Ni to Fe Ίνα lines. a nickel-to-iron. abundance ratio of 0.055-0.075 is found. ie. a factor of 1.5-2 Ni relative overabundance compared. to the Anders Crevesse (19890) cosmic values. ancl instead roughly consistent with the Anders Ebihara (1982) set.," b) From the flux ratio of the Ni to Fe $\alpha$ lines, a nickel-to-iron abundance ratio of 0.055-0.075 is found, i.e. a factor of 1.5-2 Ni relative overabundance compared to the Anders Grevesse (1989) cosmic values, and instead roughly consistent with the Anders Ebihara (1982) set."
+ To the best of our knowledge this is the first unambiguous Xray detection of a Ni line in an AGN., To the best of our knowledge this is the first unambiguous X–ray detection of a Ni line in an AGN.
+ Previous cletections. like in the ASCA spectrum of Ark 564 (Purner ct al.," Previous detections, like in the ASCA spectrum of Ark 564 (Turner et al."
+ 2001) and in the BeppoSAX spectrum of Circinus itself (Cuainazzi et al., 2001) and in the BeppoSAX spectrum of Circinus itself (Guainazzi et al.
+ 1999) were allected by confusion with the Fe kes line., 1999) were affected by confusion with the Fe $\beta$ line.
+ c) The Fe Ix;/Ixo [lux ratio is 0.14. somewhat lower than expected.," c) The Fe $\beta$ $\alpha$ flux ratio is $\sim$ 0.14, somewhat lower than expected."
+ A possible explanation is that the iron is mildly ionized. but not much more than FeX. otherwise the Fe line emission would be much fainter then observed (and of course absent at all for Fo or more).," A possible explanation is that the iron is mildly ionized, but not much more than Fe, otherwise the Fe $\beta$ line emission would be much fainter then observed (and of course absent at all for Fe or more)."
+ d) The oesence of the Fe Ίνα Compton Shoulder. already discovered. bv Bianchi οἱ al. (," d) The presence of the Fe $\alpha$ Compton Shoulder, already discovered by Bianchi et al. ("
+2002) in the Chandra ALETG spectrum. is confirmed.,"2002) in the $Chandra$ /HETG spectrum, is confirmed."
+ The relative Dux of the Compton Shoulder. is 20434.. [ullv consistent with the Chandra finding. ancl implving Comptonthick matter (Matt 2002b).," The relative flux of the Compton Shoulder is $\pm$, fully consistent with the $Chandra$ finding, and implying Compton–thick matter (Matt 2002b)."
+ Phis provides further support to the identification of the rellecting region with the absorber and suggests that the torus is fail: homogeneous., This provides further support to the identification of the reflecting region with the absorber and suggests that the torus is fairly homogeneous.
+ GAL and SD. acknowledges ASL and. MIUIH (under. erant COFIN-00-02-36) for financial support., GM and SB acknowledges ASI and MIUR (under grant ) for financial support.
+Repzd kn (Zepka. et al.,"$R_{\rm BB}\simeq 1$ km (Zepka, et al."
+ 1996)., 1996).
+ Again. the source was unfortunately shadowed by the detector supporting ribs. mnakineg the obtaiued spectral parameters rather uncertain.," Again, the source was unfortunately shadowed by the detector supporting ribs, making the obtained spectral parameters rather uncertain."
+ Tuterestinely. it was found that this source wieght be pulsating iu the 2-rav cucrey baud.," Interestingly, it was found that this source might be pulsating in the $\gamma$ -ray energy band."
+ Zepka. ct al. (," Zepka, et al. ("
+1996) folded the arrival times of 267 counts fromCGRO EGRET at the expected pulse period.,1996) folded the arrival times of 267 counts from EGRET at the expected pulse period.
+ They found that the folded lisht curve was sienificautly displaced from wniform distribution at more than 99 confidence., They found that the folded light curve was significantly displaced from uniform distribution at more than 99 confidence.
+ We have proposed and carried out the oobservation of dduring 1998. October 16-18. (," We have proposed and carried out the observation of during 1998, October 16-18. ("
+"(Tanaka. Inoue. Πο, 1991) carries two kinds of X-ray detectors at the foci of four identical N-rav telescopes.","Tanaka, Inoue, Holt, 1994) carries two kinds of X-ray detectors at the foci of four identical X-ray telescopes."
+ The ταν CCD camera. SIS (Solid-state Πασάς Spectrometer: Burke. et al.," The X-ray CCD camera, SIS (Solid-state Imaging Spectrometer; Burke, et al."
+ 1991). has the lugher energy resolution and relatively high detection efficiency in the soft energy band.," 1994), has the higher energy resolution and relatively high detection efficiency in the soft energy band."
+ The time resolution of SIS is 1-16 sin the standard mode which is not suitable for tiniug observation of fast pulsars., The time resolution of SIS is 4-16 s in the standard mode which is not suitable for timing observation of fast pulsars.
+ The imagine easscintillation proportional counter. CUS (Cas Luagine Spectrometer: Olashi. et al.," The imaging gas-scintillation proportional counter, GIS (Gas Imaging Spectrometer; Ohashi, et al."
+ 1996: Malishima. al.," 1996; Makishima, et al."
+ 1996). has relatively high detection efficiency in the hard euergy band aud high time resolution.," 1996), has relatively high detection efficiency in the hard energy band and high time resolution."
+ We operated the SIS in 1-CCD faint mode with Ls time resolution., We operated the SIS in 1-CCD faint mode with 4-s time resolution.
+ Therefore. only pulse phase averaged spectroscopy could be made with the SIS.," Therefore, only pulse phase averaged spectroscopy could be made with the SIS."
+ We operated the CIS in the PIT mode aud assigned a part of +ιο telemetry bit to increase the time resolution reducing the spectral information., We operated the GIS in the PH mode and assigned a part of the telemetry bit to increase the time resolution reducing the spectral information.
+ The resultant time resolution was 3.9 ms or better depeuding on the telemetry rate., The resultant time resolution was 3.9 ms or better depending on the telemetry rate.
+ We used screened event data according to +1ο standard REV2 processing (Pier 1997)., We used screened event data according to the standard REV2 processing (Pier 1997).
+" The effective exposure time was 69.6 aud 76.1 ks for each SIS and GIS. respectively, and the net time span of observation was T=160 ks."," The effective exposure time was 69.6 and 76.1 ks for each SIS and GIS, respectively, and the net time span of observation was $T=160$ ks."
+ We analyzed the CIS data to search for pulsations in the N-ray. baud., We analyzed the GIS data to search for pulsations in the X-ray band.
+ The CIS observation began at 51103.0185 ALTD and cuded at 511018260 ALTD., The GIS observation began at 51103.0185 MJD and ended at 51104.8260 MJD.
+" The data from GIS2 and CIS3 were co-added. viclding ~1100 events within 3’ radius circular regionOo in the total enerevOo, baud of 0.7-7 keV. includingC» backerouud."," The data from GIS2 and GIS3 were co-added, yielding $\simeq 1100$ events within $'$ radius circular region in the total energy band of 0.7-7 keV, including background."
+" The data from GIS2 and CIS3 were co-added. viclding ~1100 events within 3’ radius circular regionOo in the total enerevOo, baud of 0.7-7 keV. includingC» backerouud.C"," The data from GIS2 and GIS3 were co-added, yielding $\simeq 1100$ events within $'$ radius circular region in the total energy band of 0.7-7 keV, including background."
+" The data from GIS2 and CIS3 were co-added. viclding ~1100 events within 3’ radius circular regionOo in the total enerevOo, baud of 0.7-7 keV. includingC» backerouud.C»"," The data from GIS2 and GIS3 were co-added, yielding $\simeq 1100$ events within $'$ radius circular region in the total energy band of 0.7-7 keV, including background."
+calibration of ?..,calibration of \cite{Mamajek:2008fk}.
+ The isochrones of ? indicated that HD 5388 is at the end of the core-H burning phase just past the turn-off point., The isochrones of \cite{2008A&A...482..883M} indicated that HD 5388 is at the end of the core-H burning phase just past the turn-off point.
+" Therefore, HD 5388 may be about to leave the main sequence or even be slightly evolved."," Therefore, HD 5388 may be about to leave the main sequence or even be slightly evolved."
+" This would introduce additional uncertainties and possibly biases into our estimates of logRig: Prot, the stellar radius, and eventually οι, and may therefore alter the above findings."," This would introduce additional uncertainties and possibly biases into our estimates of $\log R'_{H,K}$, $P_\mathrm{rot}$, the stellar radius, and eventually $i_\mathrm{rot}$, and may therefore alter the above findings."
+" In addition to HD 5388, we examined five other stars that satisfied our selection criteria in Sect. ??.."," In addition to HD 5388, we examined five other stars that satisfied our selection criteria in Sect. \ref{sec:targets}. ."
+" They are listed in Table 4 and although their astrometric motions were previously studied in the literature, we describe our findings below: Among the potential planets around the six stars selected in Sect. ??,,"," They are listed in Table \ref{tab:4} and although their astrometric motions were previously studied in the literature, we describe our findings below: Among the potential planets around the six stars selected in Sect. \ref{sec:targets},"
+" one turned out to be a star (HD 195019 b), two are likely planets (55 Cnc d and y Cep b), and one is a brown- companion (HD 5388)."," one turned out to be a star (HD 195019 b), two are likely planets (55 Cnc d and $\gamma$ Cep b), and one is a brown-dwarf companion (HD 5388)."
+ The astrometric orbits of two stars were not detected with Hipparcos and the companions remain planetary candidates., The astrometric orbits of two stars were not detected with Hipparcos and the companions remain planetary candidates.
+ The significant fraction of one third of non-planetary companions found in this small sample is a selection effect and does not reflect the properties of the planetary population., The significant fraction of one third of non-planetary companions found in this small sample is a selection effect and does not reflect the properties of the planetary population.
+" That a large population of non-planetary companions is observed with low orbital inclinations, hence mistaken for planets, is statistically unlikely and would have been detected with Hipparcos (?).."," That a large population of non-planetary companions is observed with low orbital inclinations, hence mistaken for planets, is statistically unlikely and would have been detected with Hipparcos \citep{Zucker:2001ve}."
+" At the individual level, however, a planet detected by RV remains a planet until its mass range is confined by complementary constraints derived for instance from a transit lightcurve, dynamical interactions, by direct imaging, or from astrometric data, as in the case of HD 5388."," At the individual level, however, a planet detected by RV remains a planet until its mass range is confined by complementary constraints derived for instance from a transit lightcurve, dynamical interactions, by direct imaging, or from astrometric data, as in the case of HD 5388."
+" In the future, as the data of the GAIA astrometry satellite becomes available, we will be able to measure the astrometric orbits of many planetary systems, thus obtaining an accurate census of the planetary mass distribution."," In the future, as the data of the GAIA astrometry satellite becomes available, we will be able to measure the astrometric orbits of many planetary systems, thus obtaining an accurate census of the planetary mass distribution."
+ A few more planet candidates may then turn out to be brown-dwarf companions., A few more planet candidates may then turn out to be brown-dwarf companions.
+" We have detected the astrometric orbit of HD 5388, which harboured a planet candidate of minimum mass M)sini=1.96 M;."," We have detected the astrometric orbit of HD 5388, which harboured a planet candidate of minimum mass $M_2 \sin i = 1.96\,M_\mathrm{J}$ ."
+ We found an almost face-on orbit with an inclination Of igit= ," We found an almost face-on orbit with an inclination of $i_\mathrm{orbit}= 178.3_{- 0.7}^{+ 0.4\,\circ}$."
+"Consequently, we determined the mass of the companion HD 178.3*07*,5388 b to be M;=69+20Mj."," Consequently, we determined the mass of the companion HD 5388 b to be $M_2 = 69\pm20\,M_\mathrm{J}$."
+ A mass lower than 13M; was excluded at the 3c confidence level.," A mass lower than $13\,M_\mathrm{J}$ was excluded at the $3\,\sigma$ confidence level."
+" Thus, HD 5388 b can no longer be considered a planet."," Thus, HD 5388 b can no longer be considered a planet."
+" It is instead a close brown-dwarf companion, which is remarkable because at most 0.6% of Sun-like stars have close brown-dwarf companions (?).."," It is instead a close brown-dwarf companion, which is remarkable because at most $0.6~\%$ of Sun-like stars have close brown-dwarf companions \citep{Sahlmann:2011fk}."
+" In particular, it is one of the very few brown-dwarf companions for which the astrometric orbit could be measuredand whose orbital motion is therefore fully characterised, such as (?).."," In particular, it is one of the very few brown-dwarf companions for which the astrometric orbit could be measuredand whose orbital motion is therefore fully characterised, such as \citep{Benedict:2010ph}."
+" The measurement of the spin-orbit alignment in planetary and binary systems, e.g., ? and ?, gives insight into the formation of these systems."," The measurement of the spin-orbit alignment in planetary and binary systems, e.g., \cite{Triaud:2010qy} and \cite{Albrecht:2011uq}, gives insight into the formation of these systems."
+ The most frequently applied technique uses spectroscopic measurements during the transit of eclipsing systems and only determines the projected spin-orbit angle., The most frequently applied technique uses spectroscopic measurements during the transit of eclipsing systems and only determines the projected spin-orbit angle.
+" Because it is restricted to eclipsing systems, it favours small-separation, short-period systems (another method relies on interferometric imaging of the stellar surface, see ?))."," Because it is restricted to eclipsing systems, it favours small-separation, short-period systems (another method relies on interferometric imaging of the stellar surface, see \citealt{2011ApJ...726..104H}) )."
+" We determined the stellar spin axis inclination on the basis of the spectroscopic estimate of vsinig,and the rotation period calibration using logR7, ,."," We determined the stellar spin axis inclination on the basis of the spectroscopic estimate of $\upsilon \sin i_\mathrm{rot}$and the rotation period calibration using $\log R'_{H,K}$ ."
+ We found evidence of non- spin-orbit angles with numerical values of y=23*[9419 ," We found evidence of non-zero spin-orbit angles with numerical values of $\psi= 23_{- 8}^{+10\,[+3.4]\,\circ}$ "
+Given these internal tests of our cross sections and our comparisons to experimental measurements refpl.vp)). PICrOSs section stobeungandartis ground —50%intheenergyregionneartheweestimarteourcalculateddirect stateionizationthreshold.,"Given these internal tests of our cross sections and our comparisons to experimental measurements \\ref{pi_exp}) ), we estimate our calculated direct PI cross sections to be uncertain by in the energy region near the ground state ionization threshold."
+ To benchmark and test the accuracy of our results. we have compared our (medium CI expansion) calculations to experimental absolute PI cross-section measurements performed at the ALS synchrotron radiation facility. Se. Se--. Se*-.," To benchmark and test the accuracy of our results, we have compared our (medium CI expansion) calculations to experimental absolute PI cross-section measurements performed at the ALS synchrotron radiation facility. $^+$, $^{2+}$, $^{3+}$,"
+ and Se cross were measured from the threshold of the highest-energy metastable state in the ground configuration up to 10 eV beyond the ground state ionization threshold (2).., and $^{5+}$ cross were measured from the threshold of the highest-energy metastable state in the ground configuration up to 10 eV beyond the ground state ionization threshold \citep{esteves10}.
+ Experimental results for Se (?) and Se (?) have been published. and we compare our theoretical results to experiment for those tons in this section (a similar comparison to the as-yet unpublished Se7 and Se cross sections was also performed).," Experimental results for $^+$ \citep{sterling11} and $^{3+}$ \citep{esteves09} have been published, and we compare our theoretical results to experiment for those ions in this section (a similar comparison to the as-yet unpublished $^{2+}$ and $^{5+}$ cross sections was also performed)."
+ The experimental measurements were described by ? ?.. and we refer the reader to those papers for full al," The experimental measurements were described by \citet{esteves09} and \citet{sterling11}, and we refer the reader to those papers for full details."
+ Utm. the PI cross sections were measured via the merged beam method (?)..," Briefly, the PI cross sections were measured via the merged beam method \citep{lyon86}."
+ The photoionization spectra were measured by scanning across photon energies. and normalized to an absolute scale with absolute cross-section measurements performed at discrete energies.," The photoionization spectra were measured by scanning across photon energies, and normalized to an absolute scale with absolute cross-section measurements performed at discrete energies."
+ Overall. uncertainties in the experimental cross sections were estimated to be «306€ for Se and ~20% for Se (?2)..," Overall, uncertainties in the experimental cross sections were estimated to be $\sim$ for $^+$ and $\sim$ for $^{3+}$ \citep{sterling11, esteves09}."
+ The experimental data are complicated by the fact that the primary ion beams were composed of an unknown admixture of ground and metastable states., The experimental data are complicated by the fact that the primary ion beams were composed of an unknown admixture of ground and metastable states.
+ We therefore linearly combined the computed cross sections of the ground configuration states to best reproduce the direct experimental cross sections and (when possible) families of resonances., We therefore linearly combined the computed cross sections of the ground configuration states to best reproduce the direct experimental cross sections and (when possible) families of resonances.
+ refselexp and 5 compare our calculated Se and ὃς) PI cross sections (solid lines) with experimental measurements (dotted lines)., \\ref{se1exp} and \ref{se3exp} compare our calculated $^+$ and $^{3+}$ PI cross sections (solid lines) with experimental measurements (dotted lines).
+ The computed Se cross section was convolved with a Gaussian of 28 meV FWHM to reproduce the experimental resolution., The computed $^+$ cross section was convolved with a Gaussian of 28 meV FWHM to reproduce the experimental resolution.
+ By visual comparison to the experimental data. we found best agreement by weighting the contribution of the," By visual comparison to the experimental data, we found best agreement by weighting the contribution of the"
+The distributions of the field changes are shown at the top of Figure 7..,The distributions of the field changes are shown at the top of Figure \ref{fchangehist}.
+ We will describe ihe remainder of Figure 7 in Section 5.., We will describe the remainder of Figure \ref{fchangehist} in Section \ref{fluxchanges}.
+ The significance of a field change is defined as the ratio of the amplitude of the field change to the root-mean-square scatter of the data with respect to the fit before the field change occurred., The significance of a field change is defined as the ratio of the amplitude of the field change to the root-mean-square scatter of the data with respect to the fit before the field change occurred.
+ The top right plot of Figure 7 shows a histogram of this quantity., The top right plot of Figure \ref{fchangehist} shows a histogram of this quantity.
+ The minimum in the significance plot is due to the fact that we reject field changes with amplitudes less than 1.4 times (he noise level. which is arouncl 20 G for strong active fields aud around 3 G lor quiet fields.," The minimum in the significance plot is due to the fact that we reject field changes with amplitudes less than 1.4 times the noise level, which is around 20 G for strong active fields and around 3 G for quiet fields."
+ Di practice we calculated the noise level lor each pixel under consideration., In practice we calculated the noise level for each pixel under consideration.
+ The histograms suggest power laws., The histograms suggest power laws.
+exceed ~10°? ere if some of the bursts with uukuown redshifts are located at 2>L. as suggested by their fainthost galaxies (Rz23 mag: Dergeretal.2006)).,"exceed $\sim 10^{52}$ erg if some of the bursts with unknown redshifts are located at $z\gtrsim 1$, as suggested by their fainthost galaxies $R\gtrsim 23$ mag; \citealt{bfp+06}) )."
+ The full range of 23— [orders of magnitude is similar to the spread in £.νο observed for long GRDs (Frailetal.2001:Bloometal. 2003).," The full range of $3-4$ orders of magnitude is similar to the spread in $E_{\gamma,{\rm iso}}$ observed for long GRBs \citep{fks+01,bfk03}."
+. The data preseuted in this paper couclusively show that the majority (80%) of all short CRBs exhibit a tight correlation between their isotropic-equivaleut prompt 5- ray and blast wave kinetic energies (see also Nakar2007))., The data presented in this paper conclusively show that the majority $80\%$ ) of all short GRBs exhibit a tight correlation between their isotropic-equivalent prompt $\gamma$ -ray and blast wave kinetic energies (see also \citealt{nak07}) ).
+ The inferred 5-rav cficicucy is ~85'A with a narrow spread of about 0.114 dex., The inferred $\gamma$ -ray efficiency is $\sim 85\%$ with a narrow spread of about $0.14$ dex.
+ This result is iudeed verified bv the three bursts for which detailed afterglow observatious are available. with derived ε. values of 0.8 0050709: Foxetal. 2005)). 0.2 0050721: Bergeretal. 2005))). and 0.65 0051221a: Soderbergetal. 2006a)).," This result is indeed verified by the three bursts for which detailed afterglow observations are available, with derived $\epsilon_\gamma$ values of 0.8 050709; \citealt{ffp+05}) ), 0.2 050724; \citealt{bpc+05}) ), and 0.65 051221a; \citealt{sbk+06}) )."
+ The remaining 20% of short bursts lave X-ray fluxes that are suppressed by about three orders of magnitude conparedto their x-ray fluence., The remaining $20\%$ of short bursts have X-ray fluxes that are suppressed by about three orders of magnitude comparedto their $\gamma$ -ray fluence.
+ I stress that if this is indeed due to a low cireiuniburst density refsec:outlicrs)). which leads to my«ÓÁ.. then observations of the afterelow at higher N-ray cuereies (above Ji.) should recover the same narrow distribution of e. seen for the bulls of the population.," I stress that if this is indeed due to a low circumburst density \\ref{sec:outliers}) ), which leads to $\nu_X<\nu_c$, then observations of the afterglow at higher X-ray energies (above $h\nu_c$ ) should recover the same narrow distribution of $\epsilon_\gamma$ seen for the bulk of the population."
+ The observed correlation and the interred narrow distribution of have several crucial implications for short CRB progenitor e.models. the energy extraction 1iechanisui. and burst properties such as the circumburst density aud shock unicroplysics.," The observed correlation and the inferred narrow distribution of $\epsilon_\gamma$ have several crucial implications for short GRB progenitor models, the energy extraction mechanism, and burst properties such as the circumburst density and shock microphysics."
+ The overall narrow spread ἐ.. aud the median value. are similar to those observed in long GRBs (Panaitescu&Iunar2002).. sugeestiug that the properties of the relativistic outflow are similar for long iid short bursts and are generally independent of the ientity of the progenitor or thecireunburst environment.," The overall narrow spread in $\epsilon_\gamma$, and the median value, are similar to those observed in long GRBs \citep{pk02}, suggesting that the properties of the relativistic outflow are similar for long and short bursts and are generally independent of the identity of the progenitor or thecircumburst environment."
+ Iu addition. the tight correlation between στα aud afterelow energies indicates that for most bursts the nuderling assumption that vy2νι is indeed correct.," In addition, the tight correlation between $\gamma$ -ray and afterglow energies indicates that for most bursts the underlying assumption that $\nu_X>\nu_c$ is indeed correct."
+ Therefore the required cireunburst deusities are (Carnot&Savi2002) n>0.05(1|:)Eey(Ex 7. tvpical of mterstellu environinents.," Therefore the required circumburst densities are \citep{gs02} $n\gtrsim 0.05\,(1+z)^{-1/2}
+\epsilon_{B,-1}^{-3/2} E_{50}^{1/2}$ $^{-3}$, typical of interstellar environments."
+ Iu the context of binary conrpact object progenitors (NS-NS. NS-DII). the inferred densities indicate that the majority of the progenitors do not experience kick velocities that are large enough for ejection into the intergalactic medium and/or that the merger timescales are short enough that the distance traveled is <10 kpe.," In the context of binary compact object progenitors (NS-NS, NS-BH), the inferred densities indicate that the majority of the progenitors do not experience kick velocities that are large enough for ejection into the intergalactic medium and/or that the merger timescales are short enough that the distance traveled is $\lesssim 10$ kpc."
+" Finally. the fraction of energy in relativistic electrons. e,zz(1MEAE... ust also have a narrow distribution. since both e. aud νιο. have a narrow spread."," Finally, the fraction of energy in relativistic electrons, $\epsilon_e\approx(1-\epsilon_\gamma)E_K/E_\gamma$, must also have a narrow distribution, since both $\epsilon_\gamma$ and $E_K/E_\gamma$ have a narrow spread."
+ While there is a clear correlation between £.ο aud gis the overall spread in isotropic-equivaleut energies appears to be wider than for the boeanuüueg-corrected energies of long GRBs.collimated.," While there is a clear correlation between $E_{\gamma,{\rm iso}}$ and $E_{K,{\rm iso}}$, the overall spread in isotropic-equivalent energies appears to be wider than for the beaming-corrected energies of long GRBs,."
+" To date. ouly 0051221a exhibits evidence for siguificaut beaming. with fj=|Iσωστο«1052006a).. while 0050709 appears to have a wide jet with f,z0.06 (Foxct20053."," To date, only 051221a exhibits evidence for significant beaming, with $f_b\equiv [1-{\rm
+cos}(\theta_j)]\approx 7.5\times 10^{-3}$, while 050709 appears to have a wide jet with $f_b\approx 0.06$ \citep{ffp+05}."
+". In general short CRB jets appear to be wider than those of long GRBs with 0;zLO"" compared to (ο25 (Ρούσσοςetal.2006a).", In general short GRB jets appear to be wider than those of long GRBs with $\theta_j \gtrsim 10^\circ$ compared to $\langle\theta_j\rangle\approx 5^\circ$ \citep{sbk+06}.
+. Tn the larger sample preseuted here two additional bursts with kuowu redshifts (061006 aud 061210) exhibit steep decavs. ay~2. at late time. reminiscent of a post jet break evolution (for which the expected value is ay=pwith p>2: Sarictal.1999)).," In the larger sample presented here two additional bursts with known redshifts (061006 and 061210) exhibit steep decays, $\alpha_X\sim -2$, at late time, reminiscent of a post jet break evolution (for which the expected value is $\alpha_X=-p$ with $p\gtrsim 2$; \citealt{sph99}) )."
+ Tn addition to CGRD0051221a these bursts also have the largest secure values of Ee3.4.," In addition to 051221a these bursts also have the largest secure values of $E_{\gamma,{\rm iso}}$."
+" Iu the case of 0061006 there is a clear break at fz1&100, while for 0061210 the helt curve is already in the rapid decay phase at the time of the first observations. f2«10? s, Using the conversion frou jet break time to opening angle (Sarictal.1999): I fud 0;xO.12 for 0061006. aud 0;<O16 Or 0061210 (ay=10? an?)"," In the case of 061006 there is a clear break at $t\approx 1\times 10^5$ s, while for 061210 the light curve is already in the rapid decay phase at the time of the first observations, $t\approx 2\times 10^5$ s. Using the conversion from jet break time to opening angle \citep{sph99}: I find $\theta_j\approx 0.12$ for 061006 and $\theta_j\lesssim
+0.16$ for 061210 $n_0=10^{-2}$ $^{-3}$ )."
+ Thus. in the coutext of jet breaks. the beamine-corrected chereies are E.zxjsLO erg and z6&107 erg. respectively.," Thus, in the context of jet breaks, the beaming-corrected energies are $E_\gamma\approx
+5\times 10^{48}$ erg and $\lesssim 6\times 10^{48}$ erg, respectively."
+ These values are within a factor of few of E.~(]2j«Lol? cre interred for 0051221a (Soderbereetal.2006a).. aud E.22.1«10 interred for for 0050709 (Fox2005).," These values are within a factor of few of $E_\gamma\approx (1-2)
+\times 10^{49}$ erg inferred for 051221a \citep{sbk+06}, and $E_\gamma\approx 2.1\times 10^{48}$ inferred for for 050709 \citep{ffp+05}."
+". If the observed steep decays are indeed due to jets. this leads to a siuilar inverse correlation between £i, and beaming correction that is observed in long GRBs (Frailctal.2001:Dergeret2003a:Bloomoet 2003).. and thus to a narrow distribution of Ly.)~—107 ere."," If the observed steep decays are indeed due to jets, this leads to a similar inverse correlation between $E_{\rm iso}$ and beaming correction that is observed in long GRBs \citep{fks+01,bkf03,bfk03}, and thus to a narrow distribution of $E_{\rm rel}\sim 10^{48}-10^{49}$ erg."
+ Models of the relativistic outflows from black hole accretion systems relevant for short GRBs indicate E444~Lol107 with opening angles of 515° (Jankaetal.2006).. in good agreciment with the observations presented here.," Models of the relativistic outflows from black hole accretion systems relevant for short GRBs indicate $E_{\rm rel,iso}\sim
+10^{49}-10^{52}$ with opening angles of $\sim 5-15^\circ$ \citep{jam+06}, in good agreement with the observations presented here."
+ It is thus possible that short GRBs also exhibit a nearly staudard energy release. with an overall scale that is two orders of magnitude simaller than that of long GRBs.," It is thus possible that short GRBs also exhibit a nearly standard energy release, with an overall scale that is two orders of magnitude smaller than that of long GRBs."
+ The proposed beaming corrections also impact the nature of the energy extraction mechanisin., The proposed beaming corrections also impact the nature of the energy extraction mechanism.
+ In particular. extraction via 17 aunihilation has been argued to provide at most Ey~few«1005 erg (Rosswoe&Ramirez-Ruiz 2002).. which may be sufficient if the beaming corrections are valid.," In particular, extraction via $\nu\bar{\nu}$ annihilation has been argued to provide at most $E_{\rm rel}\sim {\rm
+few}\times 10^{48}$ erg \citep{rr02}, which may be sufficient if the beaming corrections are valid."
+ On the other haud. if the observed steep decavs are not due to jets. then the required large cuereics likely poiut to MIID. processes such as the Blaucdtord-Zuajek mechanigu (Blandford&Zuajek1977:Rosswoe 2006a).. which cau. produce minosities in excess of 1077 cre 2.," On the other hand, if the observed steep decays are not due to jets, then the required large energies likely point to MHD processes such as the Blandford-Znajek mechanism \citep{bz77,rrd03,sbk+06}, , which can produce luminosities in excess of $10^{52}$ erg $^{-2}$ ."
+ These possibilities have to be assessed with a larger sample of bursts. with particular attention to multinvaveleneth observations that can differentiate between bona-fide jet breaks aud other scenarios (6.9.. energy or density variations).," These possibilities have to be assessed with a larger sample of bursts, with particular attention to multi-wavelength observations that can differentiate between bona-fide jet breaks and other scenarios (e.g., energy or density variations)."
+ Ifurther investigate whether there is a correlation between the enerev release and burst duration (159): Figure , Ifurther investigate whether there is a correlation between the energy release and burst duration $T_{90}$ ); Figure \ref{fig:fxfgamma}. .
+For the full sample I find a Speariuau's rank. correlation coefficient. p=(0.58. or a null hypothesis probability of ouly 7.5«10. 7.," For the full sample I find a Spearman's rank correlation coefficient, $\rho=0.58$ , or a null hypothesis probability of only $7.5\times 10^{-3}$ ."
+ Removing 0050509).," Removing 050509b,"
+in Figs.,in Figs.
+ + and 7. for y and ή respectively., \ref{fig:chi.eps} and \ref{fig:lh.eps} for $\chi$ and $\lambda/h$ respectively.
+ In particular. at v=0. we have which results in izo5x0.690x0.2539 (this value would be obtained from Eq.," In particular, at $x=0$, we have which results in $\frac{\lambda}{h} \approx e^{-1} \times 0.690 \approx 0.2539$ (this value would be obtained from Eq."
+ 45. for g=0.849. which ts close to the parabolic case).," \ref{eq:sl_x=0_vertprofile} for $q \approx 0.849$, which is close to the parabolic case)."
+ In this paper. we have reported the first reliable prescription for the softening length t to be used in the numerical determination of the gravitational potential of geometrically thin dises within the framework of softened gravity.," In this paper, we have reported the first reliable prescription for the softening length $\lambda$ to be used in the numerical determination of the gravitational potential of geometrically thin discs within the framework of softened gravity."
+ This expression has been found by rigorously comparing the Newtonian potential of a geometrically thin dise (of finite thickness) with the softened potential of a flat disc., This expression has been found by rigorously comparing the Newtonian potential of a geometrically thin disc (of finite thickness) with the softened potential of a flat disc.
+ This ts a function of the radius and disc local aspect ratio. and also depends on vertical stratification.," This is a function of the radius and disc local aspect ratio, and also depends on vertical stratification."
+ It is accurate at the singularity and around (typically a few disc thicknesses in radius)., It is accurate at the singularity and around (typically a few disc thicknesses in radius).
+ Although this formula is valid only locally. and obtained in the axisymmetric limit. it should help to improve the quality (realism. accuracy. and computing time) of 2D- and3D-simulations!.," Although this formula is valid only locally, and obtained in the axisymmetric limit, it should help to improve the quality (realism, accuracy, and computing time) of 2D- and."
+. If necessary. the accuracy of the prescription can be improved by considering higher order terms in the expanded complete elliptic integrals of the first kind.," If necessary, the accuracy of the prescription can be improved by considering higher order terms in the expanded complete elliptic integrals of the first kind."
+ Finally. it would then be interesting to generalize our results 1) to mass density profiles that extend to infinity. such as the Gaussian profile (which. corresponds to vertically isothermal discs). and m) to the entire disc. since this prescription is expected to be inaccurate far from the singularity.," Finally, it would then be interesting to generalize our results i) to mass density profiles that extend to infinity, such as the Gaussian profile (which corresponds to vertically isothermal discs), and ii) to the entire disc since this prescription is expected to be inaccurate far from the singularity."
+Turbulence is a fundamental open problem im classical plvsics aud is of broad importance i science ancl technology.,Turbulence is a fundamental open problem in classical physics and is of broad importance in science and technology.
+ While non-relaIvistic turbulence has beeu studied iu ereat detail (6.8.. Davidson 2001). relatively little attention has been devoted to the properties of turbulence iu relativisic gas.," While non-relativistic turbulence has been studied in great detail (e.g., Davidson 2004), relatively little attention has been devoted to the properties of turbulence in relativistic gas."
+ Yet easOS flows iu astroplivsics are known to )o relativistic and turbuleu In a wide varicty of systems under active investigation., Yet gas flows in astrophysics are known to be relativistic and turbulent in a wide variety of systems under active investigation.
+ Extensive evidence of relatIvistic gas flow now exists. nuost dramatically in asrononical observations ¢ outflows from θαΤαν bursts (CRBs) (seereceureviewsbv?T?T) and active ealactic nuclei (e.g. 222740," Extensive evidence of relativistic gas flow now exists, most dramatically in astronomical observations of outflows from gamma-ray bursts (GRBs) \citep[see recent reviews
+by][]{Piran:2004p4605, Fox:2006p4649, Woosley:2006p4650,
+ Zhang:2004p4732, Gehrels:2009p4728} and active galactic nuclei \citep[e.g.][]{Begelman:1984p4745, Antonucci:1993p4756,
+ Krolik:1999p4787, Frank:2002p4792}."
+ Nearer by. relativistic outflows are observed from a variety of sources iLichiding müero-quasars. sof Sauna repeaters (SGT). and some superuovae (?)..," Nearer by, relativistic outflows are observed from a variety of sources including micro-quasars, soft gamma repeaters (SGRs), and some supernovae \citep{Soderberg:2010p4890}."
+ Theoretical models for the| central engines powering abroad range of astrophysical phenomena invoke eas accretion onto black holes aud neutroji stars., Theoretical models for the central engines powering a broad range of astrophysical phenomena invoke gas accretion onto black holes and neutron stars.
+ In addition. the eravitational wave sources targeted by LIGO involve the hydrodyaiuuical mereiug of neutrou starneutron star or neuticn starblack hole binaries.," In addition, the gravitational wave sources targeted by LIGO involve the hydrodynamical merging of neutron star–neutron star or neutron star–black hole binaries."
+ The τοσο of neutrou star binaries generates shear as the stars touch aud ieroneC, The merging of neutron star binaries generates shear as the stars touch and merge.
+", Shear is also generated as neutron stars are shred whei Πιοας with a black hole binary companion.", Shear is also generated as neutron stars are shred when merging with a black hole binary companion.
+ Nelviu-IHehuhloltz iustability iu these shear flows gencrate turbuleuce which can lead to large amplifications of magnetic field (72).., Kelvin-Helmholtz instability in these shear flows generate turbulence which can lead to large amplifications of magnetic field \citep{Zhang:2009p3418}.
+ Such field wuplification Lav be Cruüucial for creating conditions capable of exracting observed ποος from the CRB ceutral οiciue (?).., Such field amplification may be crucial for creating conditions capable of extracting observed luminosities from the GRB central engine \citep{Gehrels:2009p4728}.
+" Many astroplivsical flows, including the aforementioned. απο at least partially relativistic (bulk aud/or thermal Loraitz factor = 1) and all are hiehlv susceptible to turbulence due to the extremely hieh Revnolds unubers cjuracterizing astroplivsical gas."," Many astrophysical flows, including the aforementioned, are at least partially relativistic (bulk and/or thermal Lorentz factor $\gtrsim
+1$ ) and all are highly susceptible to turbulence due to the extremely high Reynolds numbers characterizing astrophysical gas."
+ Inuowledee of the properties of relativistic turbulence is also of general mnüportaunce in plysics. with direct applications to carly ujiverse cosmnologv. heavy don colliders. aud high euergv density plivsies aud laboratory plasmas.," Knowledge of the properties of relativistic turbulence is also of general importance in physics, with direct applications to early universe cosmology, heavy ion colliders, and high energy density physics and laboratory plasmas."
+ Research clucidating the basic properties of relativistic firbuleuce is thus broadly motivated., Research elucidating the basic properties of relativistic turbulence is thus broadly motivated.
+ Tn order to avance knowledge of turbulence iu the lareclv ujiexplored relativistic case. we lave beeun a SCTIOS «M4 numerical simmlations of turbulent flows iu maeueized gas with relativistic οποιον. deusity or veocitv.," In order to advance knowledge of turbulence in the largely unexplored relativistic case, we have begun a series of numerical simulations of turbulent flows in magnetized gas with relativistic energy density or velocity."
+ Iu tus first in a series of papers on this topic we presen simulations aud analysis of driven turbulence iu a frans-sondc. super-Alfvénnic mediuu for which the raio of iuterwal to rest mass energv density is of order nmitv.," In this first in a series of papers on this topic we present simulations and analysis of driven turbulence in a trans-sonic, super-Alfvénnic medium for which the ratio of internal to rest mass energy density is of order unity."
+ Diffforenees frou non-relativistic turbulence may he expected since pressure and magnetic fluctuations COΠαςated by acoustic aud ΑΠΟ waves inodifv the fluid iuertial term iu the relativistic case., Differences from non-relativistic turbulence may be expected since pressure and magnetic fluctuations communicated by acoustic and MHD waves modify the fluid inertial term in the relativistic case.
+ This is cistinct from the non-relativistie case where fid inertia 15 xoportional to rest nass density but docs uot depend on internal cuerey. xessure or maguetie field streneth.," This is distinct from the non-relativistic case where fluid inertia is proportional to rest mass density but does not depend on internal energy, pressure or magnetic field strength."
+ We have studied he same model (SATSI) at a range of resolutions up to 1021? iu order to demoustrate nunerical convergeicc., We have studied the same model (SATS1) at a range of resolutions up to $1024^3$ in order to demonstrate numerical convergence.
+ We iake comparisons with the extensive body of iterature available for compressible non-relativistic lydrodvnamic and ΑΠΟ turbulence (e.g.TTTPUDDDPPTT).," We make comparisons with the extensive body of literature available for compressible non-relativistic hydrodynamic and MHD turbulence \citep[e.g.,][]{Padoan:1997p2749, VazquezSemadeni:2000p2760,
+ Cho:2003p3040, Beresnyak:2005p4444, Kritsuk:2007p3858,
+ Kritsuk:2009p2270, Schmidt:2008p4878, Federrath:2010p4876,
+ Lemaster:2009p2711, Burkhart:2009p4462, Kowal:2010p3546}."
+" Wo also drew from studies of Alfvénnic turbulence in the nowrelativistico (6.8.222277) aud relativistic force-free (2) τοσο»,"," We also draw from studies of Alfvénnic turbulence in the non-relativistic \citep[e.g.,][]{Politano:1995p4542, Cho:2000p4507, Maron:2001p3039,
+ Beresnyak:2005p4444, Beresnyak:2009p4471, Beresnyak:2010p4470} and relativistic force-free \citep{Cho:2005p4453} regimes."
+ By applying the same analyses utilized iu these stucies to our own models. οἱ roa is to elucidate the similarities ancl clifferences exisine between the relativistic aud," By applying the same analyses utilized in these studies to our own models, our aim is to elucidate the similarities and differences existing between the relativistic and"
+(with a correlation coefficient r=0.99 and a rms scatter of only «=0.086 dex from 17 clusters) is therefore strictly only applicable up to the metallicity of NGC 104 (47 Tuc). [Fe/H]=-0.77 dex.,"(with a correlation coefficient $r=0.99$ and a $rms$ scatter of only $\sigma=0.086$ dex from 17 clusters) is therefore strictly only applicable up to the metallicity of NGC 104 (47 Tuc), [Fe/H]=-0.77 dex."
+ However. from the abundance analysis of two very metal rich bulge clusters. Carretta et al. (," However, from the abundance analysis of two very metal rich bulge clusters, Carretta et al. ("
+2001) already noted that the index W might deviate from linearity at high metallicity. whereas in their work Rutledge et al. (,"2001) already noted that the index $W'$ might deviate from linearity at high metallicity, whereas in their work Rutledge et al. ("
+1997) claimed that a linear relation existed between W and ]ccGos. simply because results for metal-rich clusters were not available at the time.,"1997) claimed that a linear relation existed between W' and $_{CG97}$, simply because results for metal-rich clusters were not available at the time."
+ With the present sample we cannot extend the calibration further. but in the Appendix we add two near solar metallicity bulge clusters to obtain a more extended calibration of this index. to be inserted in a final compilation of GC metallicities.," With the present sample we cannot extend the calibration further, but in the Appendix we add two near solar metallicity bulge clusters to obtain a more extended calibration of this index, to be inserted in a final compilation of GC metallicities."
+ Within our survey of 19 Galactic GCs with FLAMES (UVES and GIRAFFE). we have measured the iron abundance of about 2000 RGB stars in a accurate and homogeneous way (determination of atmospheric parameters. measure of EWs. spectroscopic analysis. ete).," Within our survey of 19 Galactic GCs with FLAMES (UVES and GIRAFFE), we have measured the iron abundance of about 2000 RGB stars in a accurate and homogeneous way (determination of atmospheric parameters, measure of $EW$ s, spectroscopic analysis, etc)."
+ In the present paper we exploit this huge. homogeneous data set. covering (almost) the entire range of GC metallicity in our Galaxy and a summary of our findings follows:," In the present paper we exploit this huge, homogeneous data set, covering (almost) the entire range of GC metallicity in our Galaxy and a summary of our findings follows:"
+values within stamps and of the distribution of stamps themselves.,values within stamps and of the distribution of stamps themselves.
+" However, for the M31 bulge, where relatively high fraction of pixels change due to astrophysical variation,a sigma clipping can lead to a significant fraction of the image area being excluded, particularly around the M31 core."," However, for the M31 bulge, where a relatively high fraction of pixels change due to astrophysical variation, sigma clipping can lead to a significant fraction of the image area being excluded, particularly around the M31 core."
+ To test the potential impact of the high sky density of variables we perform a double-pass iterative minimization of Equation (1)) that corrects the target image for the presence of variable sources., To test the potential impact of the high sky density of variables we perform a double-pass iterative minimization of Equation \ref{dia}) ) that corrects the target image for the presence of variable sources.
+ On the first pass ISIS is run in the normal way., On the first pass ISIS is run in the normal way.
+" Then, all pixel fluxes on the resulting difference image above a specified absolute flux threshold are added back on to the target frame to form a modified target image in which intrinsically variable sources are largely removed (i.e. matched to their flux values on the reference image)."," Then, all pixel fluxes on the resulting difference image above a specified absolute flux threshold are added back on to the target frame to form a modified target image in which intrinsically variable sources are largely removed (i.e. matched to their flux values on the reference image)."
+ A second pass of ISIS is then performed on this modified target frame., A second pass of ISIS is then performed on this modified target frame.
+ The resulting kernel solution for the modified frame is then applied to the reference image and the convolved reference is subtracted from the original (unmodified) target frame to create the final difference image., The resulting kernel solution for the modified frame is then applied to the reference image and the convolved reference is subtracted from the original (unmodified) target frame to create the final difference image.
+ This method ensures that the final difference image is largely insensitive to the presence of variable stars and that as much of the image area as possible is used to compute the kernel., This method ensures that the final difference image is largely insensitive to the presence of variable stars and that as much of the image area as possible is used to compute the kernel.
+" Reassuringly, we find virtually identical results are recovered using either a standard ISIS run or the double-pass method described above, showing that ISIS is capable of finding an optimal kernel solution even for high crowding densities of variable stars."," Reassuringly, we find virtually identical results are recovered using either a standard ISIS run or the double-pass method described above, showing that ISIS is capable of finding an optimal kernel solution even for high crowding densities of variable stars."
+ In Figures 8 and 9 we show two periodic variable star light-curves from the Angstrom Project data processed with and without the photometric alignment stage described in Section 3.1.., In Figures \ref{lc-comp} and \ref{lc-comp2} we show two periodic variable star light-curves from the Angstrom Project data processed with and without the photometric alignment stage described in Section \ref{photscale}.
+ Figure 8 shows the light-curve of a variable with a period of 501 days located 2/4 from the M31 core monitored over five observing seasons., Figure \ref{lc-comp} shows the light-curve of a variable with a period of 501 days located $2 \farcm 4$ from the M31 core monitored over five observing seasons.
+" The right-hand and left-hand columns show the results with and without photometric alignment, respectively."," The right-hand and left-hand columns show the results with and without photometric alignment, respectively."
+" Panel (a) shows the light-curve itself, panel (b) is the computed periodogram, and panel (c) is the result of folding the light-curve with the period corresponding to the peak in the periodogram."," Panel (a) shows the light-curve itself, panel (b) is the computed periodogram, and panel (c) is the result of folding the light-curve with the period corresponding to the peak in the periodogram."
+ Looking at panel (a) in the right-hand column of Figure 8 we see that the light-curve derived from the photometrically aligned images appears much more coherent and less scattered than that derived from unaligned data in the left-hand panel., Looking at panel (a) in the right-hand column of Figure \ref{lc-comp} we see that the light-curve derived from the photometrically aligned images appears much more coherent and less scattered than that derived from unaligned data in the left-hand panel.
+ Panel (b) shows the computed periodogram for periods from 5 to 5000 days calculated using the Generalised Lomb Scargle method of Zechmeister&Kürster(2009).., Panel (b) shows the computed periodogram for periods from 5 to 5000 days calculated using the Generalised Lomb Scargle method of \cite{zech09}.
+ This generalised method extends the standard Lomb Scargle approach by taking account of errors, This generalised method extends the standard Lomb Scargle approach by taking account of errors
+a slow increase above a few TeVs.,a slow increase above a few TeVs.
+ This is due to the comparatively slower diffusion of CRs in this model and the dominance by a single source (Vela) over a wide range of energy spectrum., This is due to the comparatively slower diffusion of CRs in this model and the dominance by a single source (Vela) over a wide range of energy spectrum.
+" The total CR anisotropy A expected under our model can be calculated using the following equation (Thoudam 2007), where the summation is over the nearby SNRs, 7; denotes the direction of the i SNR giving an intensity J; at the Earth, £5, denotes the direction of maximum intensity, [7 represents the total observed CR intensity and δι denotes the anisotropy amplitude due to a single SNR."," The total CR anisotropy $\Delta$ expected under our model can be calculated using the following equation (Thoudam 2007), where the summation is over the nearby SNRs, $\hat{r}_i$ denotes the direction of the $i^{th}$ SNR giving an intensity $I_i$ at the Earth, $\hat{r}_m$ denotes the direction of maximum intensity, $I_T$ represents the total observed CR intensity and $\delta_i$ denotes the anisotropy amplitude due to a single SNR."
+" 6; under the diffusion approximation is given by (Mao & Shen 1972), where N; (given by Eq."," $\delta_i$ under the diffusion approximation is given by (Mao $\&$ Shen 1972), where $N_i$ (given by Eq."
+ 7) denotes the CR density from an SNR with distance r; and age t;., 7) denotes the CR density from an SNR with distance $r_i$ and age $t_i$.
+" For our best fit proton results, we get Az(1.7x10?—0.12) and (1—4)x10? for Model A and B respectively in the energy range of (1—100) TeV. Our estimates are larger than the measured anisotropies of ~(0.5—1)x107? in the same energy range."," For our best fit proton results, we get $\Delta\approx (1.7\times 10^{-2}-0.12)$ and $(1-4)\times10^{-2}$ for Model A and B respectively in the energy range of $(1-100)$ TeV. Our estimates are larger than the measured anisotropies of $\sim (0.5-1)\times 10^{-3}$ in the same energy range."
+" But, compared to Model A, Model B looks more closer to the measured values (see also Ptuskin et al."," But, compared to Model A, Model B looks more closer to the measured values (see also Ptuskin et al."
+ 2006)., 2006).
+" We show that for both Models A and B, the nearby SNRs contribute mostly above ~(0.5—1) TeV/n and they may account for the observed spectral hardening at high energies."," We show that for both Models A and B, the nearby SNRs contribute mostly above $\sim(0.5-1)$ TeV/n and they may account for the observed spectral hardening at high energies."
+ We show this for a wide range of CR injection efficiencies and CR escape parameters from the SNRs., We show this for a wide range of CR injection efficiencies and CR escape parameters from the SNRs.
+" Looking into the averaged spectra in Figs. 1&2,"," Looking into the averaged spectra in Figs. $1\&2$,"
+ we find that both the models predict that the hardening of the helium spectrum should occur at lower energies as compared to the protons., we find that both the models predict that the hardening of the helium spectrum should occur at lower energies as compared to the protons.
+" We also find that the averaged result of Model A seems to explain the overall data better than that of Model B. However, the wide range of parameters values considered in our study allow both the propagation models to successfully explain the observed data with a careful choice of model parameters."," We also find that the averaged result of Model A seems to explain the overall data better than that of Model B. However, the wide range of parameters values considered in our study allow both the propagation models to successfully explain the observed data with a careful choice of model parameters."
+ We show this with our best fit results in Fig., We show this with our best fit results in Fig.
+ 3., 3.
+" But, the high CR injection efficiency of ερ=20% required in Model B is around a factor of 2 larger than the normally considered value of ~10% for CR studies in the Galaxy."," But, the high CR injection efficiency of $\epsilon_p=20\%$ required in Model B is around a factor of $2$ larger than the normally considered value of $\sim 10\%$ for CR studies in the Galaxy."
+" Moreover, the steep source index of required in this model is also hard to reconcile with the results of diffusive shock acceleration theory which predict an index of y= 2."," Moreover, the steep source index of $\gamma\sim 2.4$ required in this model is also hard to reconcile with the results of diffusive shock acceleration theory which predict an index of $\gamma= 2$ ."
+" Model A, on the other hand, looks favorable considering its relatively more reasonable values of the source index (y—2.15) and the proton injection efficiency {ερ=9%) required to explain the observed hardening."," Model A, on the other hand, looks favorable considering its relatively more reasonable values of the source index $(\gamma= 2.15)$ and the proton injection efficiency $(\epsilon_p=9\%)$ required to explain the observed hardening."
+" In addition, Model A also better explain the apparent observed property that the spectral hardening does not persist above a few TeVs."," In addition, Model A also better explain the apparent observed property that the spectral hardening does not persist above a few TeVs."
+" However, the measured anisotropy seem to favor Model B which assumes a weaker energy dependence of CR diffusion in the Galaxy."," However, the measured anisotropy seem to favor Model B which assumes a weaker energy dependence of CR diffusion in the Galaxy."
+ Our results look different from the predictions of other models., Our results look different from the predictions of other models.
+ Models based on constant diffusion coefficient at high energies or spectral dispersion in the source spectrum are expected to produce a high energy spectrum which remains hard up to the maximum energy (Ave et al., Models based on constant diffusion coefficient at high energies or spectral dispersion in the source spectrum are expected to produce a high energy spectrum which remains hard up to the maximum energy (Ave et al.
+ 2009)., 2009).
+" But, the data indicates that the spectral hardening happens only up to ~(20—30) TeV for protons and 10 TeV/n for helium which in general agrees well with our predictions (especially with Model A)."," But, the data indicates that the spectral hardening happens only up to $\sim (20-30)$ TeV for protons and $\sim 10$ TeV/n for helium which in general agrees well with our predictions (especially with Model A)."
+" It should be mentioned that our results may not be significantly different from others if the CR spectrum has a break or a cut-off (normally assumed to be exponential) at energies <0.1 PeV. In such a case, the spectrum will start rolling over before it starts showing noticeable differences."," It should be mentioned that our results may not be significantly different from others if the CR spectrum has a break or a cut-off (normally assumed to be exponential) at energies $\lesssim 0.1$ PeV. In such a case, the spectrum will start rolling over before it starts showing noticeable differences."
+" But, note that a cut-off somewhere between ~(3—5) PeV is preferred, irrespective of the nature and the origin of the cut-off, in order to explain the observed knee in the energy spectrum of CRs (Hórrandel 2003)."," But, note that a cut-off somewhere between $\sim (3-5)$ PeV is preferred, irrespective of the nature and the origin of the cut-off, in order to explain the observed knee in the energy spectrum of CRs (Hörrandel 2003)."
+ The secondary CR spectrum under our model can be even more different from other models., The secondary CR spectrum under our model can be even more different from other models.
+ Secondaries are those which are considered to be produced by the spallation of the primaries only during the propagation in the Galaxy., Secondaries are those which are considered to be produced by the spallation of the primaries only during the propagation in the Galaxy.
+" Their spectrum N, in the Galaxy is related to their primary spectrum N, as N&(E)οςN,(E)/ ", Their spectrum $N_s$ in the Galaxy is related to their primary spectrum $N_p$ as $N_s(E)\propto N_p(E)/D(E)$ .
+"Thus, for N,(E)οςET, the secondary spectrum follows N,(E)D(E).οςΕ0*9) which is steeper than their primaries by the diffusion index a."," Thus, for $N_p(E)\propto E^{-\Gamma}$, the secondary spectrum follows $N_s(E)\propto E^{-(\Gamma+a)}$ which is steeper than their primaries by the diffusion index $a$."
+" Therefore, once D(E) is fixed, the shape of the secondary spectrum depends on the shape of their primary spectrum."," Therefore, once $D(E)$ is fixed, the shape of the secondary spectrum depends on the shape of their primary spectrum."
+" This means that models which assume the same D(E) but different N,(E) will produce different N.(E).", This means that models which assume the same $D(E)$ but different $N_p(E)$ will produce different $N_s(E)$.
+" Under our model, if we neglect the production of secondaries inside the SNRs, we can assume that all the secondaries are produced by the background CRs."," Under our model, if we neglect the production of secondaries inside the SNRs, we can assume that all the secondaries are produced by the background CRs."
+" As our background primary spectrum is steeper than the spectrum used in other models to explain the spectral hardening, e.g. Yuan et al."," As our background primary spectrum is steeper than the spectrum used in other models to explain the spectral hardening, e.g. Yuan et al."
+" 2011, we expect a steeper secondaryspectrum in our case."," 2011, we expect a steeper secondaryspectrum in our case."
+ This difference would be even more significant when compared to propagation models which, This difference would be even more significant when compared to propagation models which
+Observations in the last decade revealed. the existence of eroups of galaxies containing extended: N-rav-mitting hot gas with properties expected for poor clusters such as the Virgo cluster. but in the optical light completely dominated bv a single. luminous. giant elliptical. galaxy (??)..,"Observations in the last decade revealed the existence of groups of galaxies containing extended X-ray-emitting hot gas with properties expected for poor clusters such as the Virgo cluster, but in the optical light completely dominated by a single luminous, giant elliptical galaxy \cite[][]{Ponman1994a,
+ Vikhlinin1999a}."
+ The SCCOLL brightest galaxv in these svstems is more than a factor five less luminous than the dominant elliptical., The second brightest galaxy in these systems is more than a factor five less luminous than the dominant elliptical.
+ n| be specific. these svstems are defined as spatially extened. N-rav sources with luminosities Lx0baHUVES . do," To be specific, these systems are defined as spatially extended X-ray sources with luminosities $L_{\rm X,bol} \ge 10^{42} h_{50}\;{\rm erg}\; {\rm s}^{-1}$ ."
+cThe optical1. counterpartsI are 8galaxy NgerOUps with Ans»2 mag. where Amys is the absolute R-xuad magniuce-gap between the brightest ancl second-brigWest galaxies.," The optical counterparts are galaxy groups with $\Delta m_{12} \ge 2 $ mag, where $\Delta m_{12}$ is the absolute R-band magnitude-gap between the brightest and second-brightest galaxies."
+" ""jese systems are extremely interesting for several DOO8O0LIS: Although they have X-ray temperaues comparable to the Virgo cluster. these. syst«nus show a ealaxy luminosity function with a deficit of brieht ealaxies bevond the characteristic magnitude of the Schechter unction AL* or of visible galaxies as compared to the xedietions of cosmological simulations (?).. whereas Virgo contains six AL* galaxies. (?).."," These systems are extremely interesting for several reasons: Although they have X-ray temperatures comparable to the Virgo cluster these systems show a galaxy luminosity function with a deficit of bright galaxies beyond the characteristic magnitude of the Schechter function M*, or of visible galaxies as compared to the predictions of cosmological simulations \cite[][]{DOnghia2004a}, whereas Virgo contains six M* galaxies \cite[][]{Jones2000a}."
+ Therefore they have been interpreted as the final outcome of galaxv-galaxy mergers., Therefore they have been interpreted as the final outcome of galaxy-galaxy mergers.
+ umerical simulations suggest that the luminous galaxies in a group will eventually merge form a single giant elliptical ealaxy (e.g., Numerical simulations suggest that the luminous galaxies in a group will eventually merge to form a single giant elliptical galaxy \cite[e.g.][]{Barnes1989a}.
+ The merging timescales for the xiehtest. group members (with magnitudes Al ~ ME or xiehter) in conipact eroups are typically a [ew tenths of a Llubble time., The merging timescales for the brightest group members (with magnitudes M $\sim$ M* or brighter) in compact groups are typically a few tenths of a Hubble time.
+ Therefore. by the present clay several eroup galaxies have likely merged into the giant. elliptical.," Therefore, by the present day several group galaxies have likely merged into the giant elliptical."
+ Outside of the high-density. core. the cooing time for the intra-group medium is larger than a llubble time: thus. while the luminous galaxies in some DInFOUps have had enough time to merge into a single object. the [large-scale X-ray halo of the original groups should. remain intact.," Outside of the high-density core, the cooling time for the intra-group medium is larger than a Hubble time; thus, while the luminous galaxies in some groups have had enough time to merge into a single object, the large-scale X-ray halo of the original groups should remain intact."
+ This means that a merged group might apyear todav as an isolated elliptical galaxy with a eroup-like X-ray halo (?).., This means that a merged group might appear today as an isolated elliptical galaxy with a group-like X-ray halo \cite[][]{Ponman1993a}.
+ llence these svstems have been termed “fossil” groups., Hence these systems have been termed “fossil” groups.
+ Using the ROSAT All-Skv Survey data. ? found. the first “fossil” eroup canclidate.," Using the ROSAT All-Sky Survey data, \citet[][]{Ponman1994a} found the first “fossil” group candidate."
+ The RNJ1840.6|401SN svsten has an X-ray luminosity comparable to a group. but of the optical light comes from a single elliptical. galaxy," The RXJ1340.6+4018 system has an X-ray luminosity comparable to a group, but of the optical light comes from a single elliptical galaxy"
+Numerical simulations were performed with OSIRIS. a fully relativistic. electromagnetic. and massively parallel PIC! code (Fonsecaetal.2003.2008).. which has been widely used in other astrophysical problems (Silvaοἱal.2003:Medvedev2009 ).,"Numerical simulations were performed with OSIRIS, a fully relativistic, electromagnetic, and massively parallel PIC code \citep{fonseca03,fonseca08}, which has been widely used in other astrophysical problems \citep{silva03,medvedev05,medvedev06,martins09}."
+". We simulate 3D svstems of shearing slabs of cold (ay29μμ. where cy, is the thermal velocity) unmagnetized electron-proton plasmas with a realistic mass ralio mg/m,=1836 (my is the proton mass). and evolve it until the electromagnetic energy saturales on the electron time-scale."," We simulate 3D systems of shearing slabs of cold $v_{0} \gg v_{th}$, where $v_{th}$ is the thermal velocity) unmagnetized electron-proton plasmas with a realistic mass ratio $m_p/m_e=1836$ $m_p$ is the proton mass), and evolve it until the electromagnetic energy saturates on the electron time-scale."
+ In this Letter. we explore a subrelativistic (so=1.02) and a relativistic (sy= 3) WIL scenario. in (he regime of two plasma slabs of equal density. which is directly relevant to internal shocks where 1«5510 with comparable density shells (Piran2005).," In this Letter, we explore a subrelativistic $\gamma_0 = 1.02$ ) and a relativistic $\gamma_0 = 3$ ) KHI scenario, in the regime of two plasma slabs of equal density, which is directly relevant to internal shocks where $1 < \gamma_0 < 10$ with comparable density shells \citep{Piran05}."
+. The munerical simulations are prepared as follows., The numerical simulations are prepared as follows.
+" The shear flow initial condition is set by a velocity field with ey pointing in the positiver, 1 direction. in the upper and lower quarters of the simulation box. and a svuumetric velocity field with ο pointing in (he negative wry direction. in the midcdle-half of the box."," The shear flow initial condition is set by a velocity field with $v_0$ pointing in the positive $x_1$ direction, in the upper and lower quarters of the simulation box, and a symmetric velocity field with $-v_0$ pointing in the negative $x_1$ direction, in the middle-half of the box."
+ Initially. the svstems are charge and current neutral.," Initially, the systems are charge and current neutral."
+" In the subrelativistic case. the simulation box dimensions are 20xο where wy=(dane?/m,)'? is the plasma frequency. and we use 20 cells per electron skin depth (c/u)."," In the subrelativistic case, the simulation box dimensions are $20 \times 20 \times 20 ~ ( c/\omega_p )^3$, where $\omega_p = (4\pi n e^2/m_e)^{1/2}$ is the plasma frequency, and we use $20$ cells per electron skin depth $c/\omega_p$ )."
+ The simulation box dimensions for the relativistic scenario are 250x80(ορ) with a resolution of 4 cells per cfc.," The simulation box dimensions for the relativistic scenario are $250 \times 80 \times 80 ~ ( c/\omega_p )^3$, with a resolution of $4$ cells per $c/\omega_p$."
+ Periodic boundary conditions are imposed in every direction., Periodic boundary conditions are imposed in every direction.
+ The magnetic field structures generated in the subrelativistic and relativistic scenarios are displaved in Figures D. and 2.. respectively. ancl (he magnetic fiekl line topology. is illustrated in Figure 3..," The magnetic field structures generated in the subrelativistic and relativistic scenarios are displayed in Figures \ref{fig:3dclas} and \ref{fig:3drel}, respectively, and the magnetic field line topology is illustrated in Figure \ref{fig:3dlines}."
+ In the linear regime. the onset of the fluid NIL. discussed in Section 2.. occurs in (he ccc plane around (he shear surfaces. generating a magnetic field normal to this plane. D».," In the linear regime, the onset of the fluid KHI, discussed in Section \ref{sec:theory}, , occurs in the $x_1x_3$ plane around the shear surfaces, generating a magnetic field normal to this plane, $B_2$."
+ This growing magnetic field component D» is responsible [for the electrons to form the signature III vortices (insets a2 and b2 of Figure 4))., This growing magnetic field component $B_2$ is responsible for rolling-up the electrons to form the signature KH vortices (insets a2 and b2 of Figure \ref{fig:bevol}) ).
+" The tvpical length-scale of the NIL modulations. observed in Figures 1. and 2. f agrees with (he wavelength of the most unstable mode predicted bv the {hud theory (Anas=2cf, and Aga&50οως in the subrelativistic and relativistic cases. respectively)."," The typical length-scale of the KHI modulations, observed in Figures \ref{fig:3dclas} f and \ref{fig:3drel} f, agrees with the wavelength of the most unstable mode predicted by the fluid theory $\lambda_\mathrm{max} = 2~c/\omega_p$ and $\lambda_\mathrm{max} \simeq 50~c/\omega_p$ in the subrelativistic and relativistic cases, respectively)."
+ The III modulations. however. are less noticeable in the relativistic regime because (hese are maskecl bv a strong DC component with a magnitude higher than the AC component. which is negligible in the subrelativistic regime.," The KHI modulations, however, are less noticeable in the relativistic regime because these are masked by a strong DC component with a magnitude higher than the AC component, which is negligible in the subrelativistic regime."
+ This DC magnetic field (&= 0). is not unstable according to the fIuid model as can be seen in equation (1)).," This DC magnetic field $k'=0$ ), is not unstable according to the fluid model as can be seen in equation \ref{eq1}) )."
+ As the amplitude of the III mocdulations erow. the electrons [rom one How cross the shear surfaces aud enter the flow.," As the amplitude of the KHI modulations grow, the electrons from one flow cross the shear surfaces and enter the counter-streaming flow."
+ Since (he protons are unperturbed. due to their inertia. (he current neutrality around (he shear surfaces is unbalanced. forming DC currentsheets. which point in the," Since the protons are unperturbed, due to their inertia, the current neutrality around the shear surfaces is unbalanced, forming DC currentsheets, which point in the"
+snapshotsunits duringare inn its last stages. Perhaps the most interesting result from Figure 8 is the normalization as a function of redshift.,units are comoving.} n its last stages.} Perhaps the most interesting result from Figure \ref{fig:lum_mass_64} is the normalization as a function of redshift.
+" At z=0, a 101M mass halo emits~ while a similar mass halo at z=2 emits 10?'erg/s/Hz,~107erg/s/Hz."," At $z=0$, a $10^{14}M_\odot$ mass halo emits$\sim10^{20} erg/s/Hz$ , while a similar mass halo at $z=2$ emits $\sim10^{22} erg/s/Hz$."
+ This amplifies our hypothesis that the merger state of the halo is very important., This amplifies our hypothesis that the merger state of the halo is very important.
+" The 10!M halo at z=2 is one of the most massive objects at that time and has likely recently formed, whereas the same mass halo at z—0 is a fairly common object that was likely formed some time ago."," The $10^{14}M_\odot$ halo at $z=2$ is one of the most massive objects at that time and has likely recently formed, whereas the same mass halo at $z=0$ is a fairly common object that was likely formed some time ago."
+" Additionally, the probability of a merger with a 1:1 mass ratio is very low for the largest halos at a given time."," Additionally, the probability of a merger with a $1:1$ mass ratio is very low for the largest halos at a given time."
+" For both simulations at multiple redshifts, the radio power is correlated with the mass of the halo."," For both simulations at multiple redshifts, the radio power is correlated with the mass of the halo."
+" 'This dependence is expected since the radio emission is a function of the temperature, density and magnetic field strength of the halo, all of which scale positively with the mass."," This dependence is expected since the radio emission is a function of the temperature, density and magnetic field strength of the halo, all of which scale positively with the mass."
+ The second result of these radio luminosity-mass relationships is the large scatter around the best fit., The second result of these radio luminosity-mass relationships is the large scatter around the best fit.
+ We see that there can be scatter of up to 2—3 orders of magnitude for a given mass cluster., We see that there can be scatter of up to $2-3$ orders of magnitude for a given mass cluster.
+" This suggests that while the radio emission is correlated with mass, the merger state of the halo plays a major role in determining the radio power."," This suggests that while the radio emission is correlated with mass, the merger state of the halo plays a major role in determining the radio power."
+" As can be seen from the projections of these halos in Figure 7,, themost radio luminous objects have very disturbed morphology and are undergoing"," As can be seen from the projections of these halos in Figure \ref{fig:halos}, , themost radio luminous objects have very disturbed morphology and are undergoing"
+ab these wavelengths.,at these wavelengths.
+ One therefore. draws the conclusion that ellipticals and. more generally. most spheroids mus have formed in clust-shroucecl starbursts.," One therefore draws the conclusion that ellipticals and, more generally, most spheroids must have formed in dust-shrouded starbursts."
+ In this paper. we show how separate tracking of disk and. spheroid star formation within a simple bu explicit cosmological framework enables one to substantiate this view and show that current data is. consisten with spheroic formation clominating the submillimetre background.," In this paper, we show how separate tracking of disk and spheroid star formation within a simple but explicit cosmological framework enables one to substantiate this view and show that current data is consistent with spheroid formation dominating the submillimetre background."
+" La section 2. we brielly describe how we connec the various pieces of our semi-analvtie model together. to differentiate between spheroids and. disk galaxies ancl use the model galaxy spectra (ονος, Guiderdoni Sadat 1999) to make the link with observations."," In section 2, we briefly describe how we connect the various pieces of our semi-analytic model together, to differentiate between spheroids and disk galaxies and use the model galaxy spectra (Devriendt, Guiderdoni Sadat 1999) to make the link with observations."
+ The approach is very similar in spirit to that of Devricnelt CGuiderdoni (2000) - DCOO hereafter. so we point out the dilferences/improvements with respect to this previous work.," The approach is very similar in spirit to that of Devriendt Guiderdoni (2000) - DG00 hereafter, so we point out the differences/improvements with respect to this previous work."
+ Section 3 presents the results obtained. with this new implementation of collision-incluced galaxy formation., Section 3 presents the results obtained with this new implementation of collision-induced galaxy formation.
+ Finally. we put our work in context with other studies and draw conclusions in Section 4.," Finally, we put our work in context with other studies and draw conclusions in Section 4."
+ To match. the combined observations of the cosmic microwave backeround anisotropies on sub-degree scales by BOOAMERanG (de Dernardis et al., To match the combined observations of the cosmic microwave background anisotropies on sub-degree scales by BOOMERanG (de Bernardis et al.
+ 2000) and ALANIALA (Balbi et al., 2000) and MAXIMA (Balbi et al.
+ 2000). and of high-redshift supernovae (Riess et al.," 2000), and of high-redshift supernovae (Riess et al."
+ 1998: Schmidt et αἱ., 1998; Schmidt et al.
+ 1998: Perlmutter ct αἱ., 1998; Perlmutter et al.
+ 1999). we have chosen a flat. universe model with a cosmological constant.," 1999), we have chosen a flat universe model with a cosmological constant."
+ More specifically. the ratio of total matter density to critical density and the reduced. cosmological constant are taken to beO4=0.3 and Ay=0.7 respectively.," More specifically, the ratio of total matter density to critical density and the reduced cosmological constant are taken to be$\Omega_0 = 0.3$ and $\lambda_0 = 0.7$ respectively."
+ The IIubble constant ZZ was chosen to lie within the error bar quoted. by the Hubble Space Telescope. key project. value (Freedman ct al., The Hubble constant $H_0$ was chosen to lie within the error bar quoted by the Hubble Space Telescope key project value (Freedman et al.
+ 2001)., 2001).
+ The value of Og=0.00257 is the one currently favored by various. cosmological probes of Big Bane nucleosvnthesis (Dania et al.," The value of $\Omega_B = 0.02 \, h^{-2}$ is the one currently favored by various cosmological probes of Big Bang nucleosynthesis (Bania et al."
+ 2002)., 2002).
+ Finally the normalization of our non-tilted primordial. Llarrison-Zeldovich power spectrum ox=0.86 ensures fair agreement both with the amplitude of the cosmic background radiation measured by CODE and the local cluster abundance., Finally the normalization of our non-tilted primordial Harrison-Zeldovich power spectrum $\sigma_8 = 0.86$ ensures fair agreement both with the amplitude of the cosmic background radiation measured by COBE and the local cluster abundance.
+ We note that recent cluster surveys seem to favor slightly lower values around as0.73 (e.g. Lahav ct al., We note that recent cluster surveys seem to favor slightly lower values around $\sigma_8 \sim 0.73$ (e.g. Lahav et al.
+ 2002) but our conclusions are quite insensitive to a change of 20% in any of the parameters presented in table 1L.., 2002) but our conclusions are quite insensitive to a change of 20 in any of the parameters presented in table \ref{cosmoparam}.
+ In the semi-analvtical approach that we outline in the next section. galaxies form from Gaussian. randoni density (uetuations in the primordial matter cistribution. dominated by cold dark matter (CDAL).," In the semi-analytical approach that we outline in the next section, galaxies form from Gaussian random density fluctuations in the primordial matter distribution, dominated by cold dark matter (CDM)."
+ To follow a bound perturbation. we use the well known spherical top-hat collapse model: a homogeneous spherical perturbation grows along with the expanding universe. until self.gravity results in turn around and (nondissipative) collapse.," To follow a bound perturbation, we use the well known spherical top-hat collapse model: a homogeneous spherical perturbation grows along with the expanding universe, until self–gravity results in turn around and (non–dissipative) collapse."
+ We then compute the number of perturbations ofa given mass which collapse at à given redshilt using the standard. Press Schechter (1974) prescription and starting with a primorclial power spectrum for the density. field. given by Bardeen οἱ al. (, We then compute the number of perturbations of a given mass which collapse at a given redshift using the standard Press Schechter (1974) prescription and starting with a primordial power spectrum for the density field given by Bardeen et al. (
+1986). corrected for the presence of barvons (Sugivama 1995).,"1986), corrected for the presence of baryons (Sugiyama 1995)."
+ To take into account mass build up of halos through mergers. we employ the extended PressSchechter formalism described in e.g. Lacey Cole (1993). in à similar fashion as in Dalland. Silk Schaelfer (1998) - DSS98 hereafter.," To take into account mass build up of halos through mergers, we employ the extended Press–Schechter formalism described in e.g. Lacey Cole (1993), in a similar fashion as in Balland, Silk Schaeffer (1998) - BSS98 hereafter."
+ Note that this is slightly different from Devricndt Ciuiderdoni (2000). where the prescription to compute the number of collapsed haloes was given by the peaks formalism developed in Dardeen et al. (," Note that this is slightly different from Devriendt Guiderdoni (2000), where the prescription to compute the number of collapsed haloes was given by the peaks formalism developed in Bardeen et al. ("
+1986) and where halo merging histories were not considered.,1986) and where halo merging histories were not considered.
+ To make contact with observations. star formation aud spectral evolution must. be incorporated.," To make contact with observations, star formation and spectral evolution must be incorporated."
+ We follow the eencral framework described by White Frenk (1991). that is to say. within virializecl cark matter halos. we let gas cool radiativelv. settle into ai disk. and form stars.," We follow the general framework described by White Frenk (1991), that is to say, within virialized dark matter halos, we let gas cool radiatively, settle into a disk and form stars."
+ Technical details about the semi-empirical recipes used to model the different astrophysical processes are given in. DCGOO., Technical details about the semi-empirical recipes used to model the different astrophysical processes are given in DG00.
+. Chemical ancl spectral evolution are then computed. with the stellar population svnthesis model from. Devriendt et al. (, Chemical and spectral evolution are then computed with the stellar population synthesis model from Devriendt et al. (
+L999).,1999).
+ Detailed references can be looked up in this paper. but. as our results in the far-IIt/submim crucially depend upon it. we summarize the basic assumptions/ingredients concerning dust that eo into this modelling.," Detailed references can be looked up in this paper, but, as our results in the far-IR/submm crucially depend upon it, we summarize the basic assumptions/ingredients concerning dust that go into this modelling."
+ The first. of these is that metals ancl cust are homogeneously mixed. with stars in. galaxies of oblate ellipsoidal shape., The first of these is that metals and dust are homogeneously mixed with stars in galaxies of oblate ellipsoidal shape.
+ The extinction. curve of each ealaxy is then assumed to scale. with its gas and meta contents. in a manner that is calibrated on the Milky Way and nearby spirals.," The extinction curve of each galaxy is then assumed to scale with its gas and metal contents, in a manner that is calibrated on the Milky Way and nearby spirals."
+ From such a model. the radiative transfer can be computed analytically. ancl vields the tota amount of stellar light absorbed at UVoptical wavelength.," From such a model, the radiative transfer can be computed analytically, and yields the total amount of stellar light absorbed at UV/optical wavelength."
+ Following energv conservation. this energy must be re-emitted at HU wavelengths.," Following energy conservation, this energy must be re-emitted at IR wavelengths."
+ Implementing the 3-componen absorption/emission model of Déssert. Boulanger Puge (1990). (à mixture of. PALIs. very small grains and laree erains) as inSTAHRDUST. Lc. in combination with the empirical correlation of the bolometric Ht luminosity with colours. we finally link the optical and the Ilt/submnm windows in a selfconsistent way.," Implementing the 3-component absorption/emission model of Déssert, Boulanger Puget (1990), (a mixture of PAHs, very small grains and large grains) as in, i.e. in combination with the empirical correlation of the bolometric IR luminosity with colours, we finally link the optical and the far-IR/submm windows in a self–consistent way."
+ We emphasize that. as discussed. in detail by 600. rere are three Κον parameters in such models. the star ormation elliciency 37. the feedback ellieiencv csy. and the sxtent of the gaseous disks. fi...," We emphasize that, as discussed in detail by DG00, there are three key parameters in such models, the star formation efficiency $\beta^{-1}$, the feedback efficiency $\epsilon_{SN}$ and the extent of the gaseous disks $f_c$ ."
+ We take parameter values at are fairly similar to the ones used by these authors.," We take parameter values that are fairly similar to the ones used by these authors,"
+to microlensing comes out the same for both images.,to microlensing comes out the same for both images.
+ We leave this for future investigation., We leave this for future investigation.
+ We noted. in Section 2.4 that when a map generated [rom | smooth matter is subtracted from the corresponding map generated from stars| small masses. the residual magnitudes have a roughly Gaussian distribution.," We noted in Section \ref{procedure} that when a map generated from $+$ smooth matter is subtracted from the corresponding map generated from $+$ small masses, the residual magnitudes have a roughly Gaussian distribution."
+ The are variabiliVN. suggesting residualthat Gaussian magnitudesnoise in themicrolensing nanolensingobservations will be indistinguishable from nanolensing.," The residual magnitudes are the nanolensing variability, suggesting that Gaussian noise in microlensing observations will be indistinguishable from nanolensing."
+ However. of subtractionthe maps in Figure1. resolves only the caustic bands. which produce the variability seen in Figures 5 and 6...," However, subtraction of the maps in Figure \ref{maps strip} resolves only the caustic bands, which produce the variability seen in Figures \ref{wiggles_by_mass} and \ref{wiggles_by_size}."
+ LE variability due to nanolensing caustics could be resolved. we expect the magnitude distribution to assume a shape more tvpical of microlensing (Lewis&Irwin1995:Schechter.Wambsganss.&Lewis 2004).," If variability due to nanolensing caustics could be resolved, we expect the magnitude distribution to assume a shape more typical of microlensing \citep{lewis1995b,schechter2004}."
+. This would make it possible to distinguish nanolensing from instrument noise., This would make it possible to distinguish nanolensing from instrument noise.
+ Another method. of. distinguishing noise be may be to examine the power spectrum of the lighteurves in. Figures 5 and 6., Another method of distinguishing noise be may be to examine the power spectrum of the lightcurves in Figures \ref{wiggles_by_mass} and \ref{wiggles_by_size}.
+ By inspection. the caustics bands produce a periodicity in the lighteurve that should appear as a peak in a power spectrum: such a spectrum. will. distinguish nanolensing [rom white and other “colours” of noise.," By inspection, the caustics bands produce a periodicity in the lightcurve that should appear as a peak in a power spectrum; such a spectrum will distinguish nanolensing from white and other “colours” of noise."
+ A detailed investigation of noise issues will be set. aside for future work., A detailed investigation of noise issues will be set aside for future work.
+ As seen in Figures 5. and 6... we find that events have a ime scale on the order of a vear or less.," As seen in Figures \ref{wiggles_by_mass} and \ref{wiggles_by_size}, we find that events have a time scale on the order of a year or less."
+ Figure δ shows a lO-vear long segment of the light curve for the case of he 2.5.10 M. masses. source size 0.01 ER. for image M (top row) and image S (bottom row): that is. corresponding o the micelle row in Figure 6..," Figure \ref{big_cadence} shows a 10-year long segment of the light curve for the case of the $2.5\times 10^{-5}$ $_\odot$ masses, source size 0.01 ER, for image M (top row) and image S (bottom row); that is, corresponding to the middle row in Figure \ref{wiggles_by_size}."
+ Phe light curves have been sampled: at regular intervals of 90 days. and the resultant »oints. plotted in crosses on the light curve.," The light curves have been sampled at regular intervals of 90 days, and the resultant points plotted in crosses on the light curve."
+ This indicates hat an observing cadence of 90 days is enough to see the nanolensing variability produced by the caustic bands., This indicates that an observing cadence of 90 days is enough to see the nanolensing variability produced by the caustic bands.
+ To find out how often. observations need to be made, To find out how often observations need to be made
+"ο(0), equation (1)). varies with respect to the clilferent sky reconstructions.","$C(\vartheta)$, equation \ref{Eq:C_theta}) ), varies with respect to the different sky reconstructions."
+ At first one has to determine the minimal value of the reconstruction parameter fis which is needed for the reconstructed. map in order to give a correlation function C'()) in agreement. with the one obtained. fron the initial map with full resolution., At first one has to determine the minimal value of the reconstruction parameter $l_{\hbox{\scriptsize max}}$ which is needed for the reconstructed map in order to give a correlation function $C(\vartheta)$ in agreement with the one obtained from the initial map with full resolution.
+ Therefore. the 2-point correlation functions computed. from the multipoles up to fons=ο. 10 and 15 are compared with the correlation function obtained from the full LLC map with maximal resolution in figure 16..," Therefore, the 2-point correlation functions computed from the multipoles up to $l_{\hbox{\scriptsize max}}=5$, 10 and 15 are compared with the correlation function obtained from the full ILC map with maximal resolution in figure \ref{Fig:correlation_function_ilc_full_lmax_dependence}."
+ From this figure one can read olf that at least a reconstruction parameter asLO is needed in order to have a good estimation of the correlation function ab large scales., From this figure one can read off that at least a reconstruction parameter $l_{\hbox{\scriptsize max}} \gtrsim 10$ is needed in order to have a good estimation of the correlation function at large scales.
+ The investigations of the previous section lead to the conclusion that the KOQT5S mask is already too laree in order to allow a reconstruction for such [large values Obs., The investigations of the previous section lead to the conclusion that the KQ75 mask is already too large in order to allow a reconstruction for such large values of $l_{\hbox{\scriptsize max}}$.
+ Por this reason the following computations are based all on the KOSS mask with a mask threshold aj=0.5 at a resolution of Nuus=16., For this reason the following computations are based all on the KQ85 mask with a mask threshold $x_{\hbox{\scriptsize th}}=0.5$ at a resolution of $N_{\hbox{\scriptsize side}} = 16$.
+ Furthermore. the reconstruction algorithm using the covariance matrix is applied. in the following. since it gives for Aliis10 almost the same results as the direct inversion methoc.," Furthermore, the reconstruction algorithm using the covariance matrix is applied in the following, since it gives for $l_{\hbox{\scriptsize max}} \gtrsim 10$ almost the same results as the direct inversion method."
+ In figure 17. three correlation functions C'(7) are shown which are calculated from the reconstructed ILC (Tr) map. from the full ILC map and. from the cata outside the IxXQs5 mask.," In figure \ref{Fig:correlation_function_ilc_KQ85_FWHM_600arcmin_nside_16}
+ three correlation functions $C(\vartheta)$ are shown which are calculated from the reconstructed ILC (7yr) map, from the full ILC map and from the data outside the KQ85 mask."
+ The three panels show the results for the reconstruction parameters fas=10. 12 and 15.," The three panels show the results for the reconstruction parameters $l_{\hbox{\scriptsize max}}=10$, 12 and 15."
+ Here the ILC (rvr) map with an additional smoothing of 600 arcmin is used., Here the ILC (7yr) map with an additional smoothing of 600 arcmin is used.
+ The reconstructed correlation function is given with lo and 2e deviations of the reconstruction errors which are estimated by where the index & counts the No=1000 CAIB simulations of the ACDAL model., The reconstructed correlation function is given with $\sigma$ and $\sigma$ deviations of the reconstruction errors which are estimated by where the index $k$ counts the $N=1000$ CMB simulations of the $\Lambda$ CDM model.
+ This quantity takes into account only the error of the reconstruction of pure CALB maps., This quantity takes into account only the error of the reconstruction of pure CMB maps.
+ Additional contributions due to residual foreground and detector noise in the ILC map propagate through the reconstruction aleorithm and are not considered here., Additional contributions due to residual foreground and detector noise in the ILC map propagate through the reconstruction algorithm and are not considered here.
+" In all three eases (4,44=10. 12 and 15) the correlation function of the reconstructed ILC map agrees well within the 20 errors with the correlation function of the full LLC map and less with the correlation function computed outside the mask."," In all three cases $l_{\hbox{\scriptsize max}}=10$, 12 and 15) the correlation function of the reconstructed ILC map agrees well within the $2\sigma$ errors with the correlation function of the full ILC map and less with the correlation function computed outside the mask."
+ This observation is used by ? as an argument. in favour of the statement that the struc” correlation function is the one obtained. from the full LLC map.," This observation is used by \cite{Efstathiou_Ma_Hanson_2009}
+ as an argument in favour of the statement that the “true” correlation function is the one obtained from the full ILC map."
+ However. as already emphasised. the large additional smoothing of GOO aremin transfers information from the masked region into the cata," However, as already emphasised, the large additional smoothing of 600 arcmin transfers information from the masked region into the data"
+with an eccentricity e=0.617.,with an eccentricity $e = 0.617$.
+ The angle between the NS1 spin axis and the post supernova orbital angular momentum axis is 18?+6° or 162?+6°., The angle between the NS1 spin axis and the post supernova orbital angular momentum axis is $18^\circ \pm 6^\circ$ or $162^\circ \pm 6^\circ$.
+" At the current time, this binary is about 300 away the galactic plane, at /=50.0? and b=2.1°, and pc8.3 kpc away from us."," At the current time, this binary is about 300 pc away the galactic plane, at $l = 50.0^\circ$ and $b = 2.1^\circ$, and 8.3 kpc away from us."
+" The observed proper motion in right ascension and declination is jg4,2—3.27 mas yr! and 44.=—1.04 mas yr!.", The observed proper motion in right ascension and declination is $\mu_{R.A.} = -3.27$ mas $^{-1}$ and $\mu_{dec.} = -1.04$ mas $^{-1}$ .
+" At the measured distance, this implies Vg4.2-129-25 km s! and Vj,241£18 km sl, in the reference frame of the Sun."," At the measured distance, this implies $V_{R.A.} = -129 \pm 25$ km $^{-1}$ and $V_{dec.} = -41 \pm 18$ km $^{-1}$, in the reference frame of the Sun."
+" In addition, instead of the characteristic age, we use the spin down age, which is 80 Myr (Arzoumanian et al."," In addition, instead of the characteristic age, we use the spin down age, which is 80 Myr (Arzoumanian et al."
+" 1999), as the upper limit on PSR B1913416 's age."," 1999), as the upper limit on PSR B1913+16 's age."
+ The results are displayed in Figures 3 and 4., The results are displayed in Figures 3 and 4.
+" As shown in the 2D V;-M;; joint probability distribution in figure 3, the kick to NS2 at birth is between 190 and 450 km s, while the impartedprogenitor mass M»; is between 1.4 and 5.0 Mo, both at confidence."," As shown in the 2D $V_k$ $M_{2i}$ joint probability distribution in figure 3, the kick imparted to NS2 at birth is between 190 and 450 km $^{-1}$, while the progenitor mass $M_{2i}$ is between 1.4 and 5.0 $_\sun$, both at confidence."
+" As seen from figure 4, both polar and planar kick directions are allowed."," As seen from figure 4, both polar and planar kick directions are allowed."
+ The current radial velocity is between —255 and 315 km s! at confidence., The current radial velocity is between $-255$ and 315 km $^{-1}$ at confidence.
+" Within the observationally constrained range, the V;,;,,, PDF is flat, with a slight for transverse velocities roughlyof ~120 km/s. Furthermore,preference although observations show a degeneracy in 0, at ~18? and ~162°, our analysis shows that 0; is much more likely to be ~18°."," Within the observationally constrained range, the $V_{trans}$ PDF is roughly flat, with a slight preference for transverse velocities of $\sim$ 120 km/s. Furthermore, although observations show a degeneracy in $\theta_t$ at $\sim\,18^\circ$ and $\sim\,162^\circ$, our analysis shows that $\theta_t$ is much more likely to be $\sim\,18^\circ$."
+" In addition, although our confidence intervals of V; and Mo; are consistent with Willems et al. ("," In addition, although our confidence intervals of $V_k$ and $M_{2i}$ are consistent with Willems et al. ("
+2004) and Wang et.,2004) and Wang et.
+ al. (, al. (
+"2006), our reported intervals are tighter, because Willems et al. (","2006), our reported intervals are tighter, because Willems et al. ("
+2004) and Wang et al. (,2004) and Wang et al. (
+"2006) reported the allowed ranges of V, and M5j instead of confidence intervals.",2006) reported the allowed ranges of $V_k$ and $M_{2i}$ instead of confidence intervals.
+" The masses of the two neutron stars in this are 1.34 and 1.25 Mo, where the more massive onebinary is NS1."," The masses of the two neutron stars in this binary are 1.34 and 1.25 $_\sun$, where the more massive one is NS1."
+" They are orbiting each other with a period of 0.102 days, and an eccentricity of 0.0878."," They are orbiting each other with a period of 0.102 days, and an eccentricity of 0.0878."
+" At the current time, the binary is close to the galactic plane, at /=245.2? and b=—4.5?."," At the current time, the binary is close to the galactic plane, at $l = 245.2^\circ$ and $b = -4.5^\circ$."
+" Through direct measurement of geometric parallax, Deller et al. ("," Through direct measurement of geometric parallax, Deller et al. ("
+"2009) measured the distance to PSR J0737-3039 to be 1.15 kpc, and the proper motion in R.A. and declination to be —3.82 and 2.13 mas yr“! respectively.","2009) measured the distance to PSR J0737-3039 to be 1.15 kpc, and the proper motion in R.A. and declination to be $-$ 3.82 and 2.13 mas $^{-1}$ respectively."
+" At the measured distance, this implies Vg4,=—20.8 km s! and Vyec.=11.6 km s!, in the reference frame of the Sun."," At the measured distance, this implies $V_{R.A.} = -20.8$ km $^{-1}$ and $V_{dec.} = 11.6$ km $^{-1}$, in the reference frame of the Sun."
+" In addition, Lorimer et al. ("," In addition, Lorimer et al. ("
+2007) constrained the possible age range to be either 70—90 or 170—190 Myr.,2007) constrained the possible age range to be either $70 - 90$ or $170 - 190$ Myr.
+ We adopt these age ranges in our analysis of PSR J0737-3039., We adopt these age ranges in our analysis of PSR J0737-3039.
+" As the radio pulse profile of PSR JO737-3039A (i.e., NS1) has not changed six years of observation, the spin tilt angle of significantlyJO737-3039A throughwith respect to the post supernova orbital angular momentum axis is believed to be small."," As the radio pulse profile of PSR J0737-3039A (i.e., NS1) has not changed significantly through six years of observation, the spin tilt angle of J0737-3039A with respect to the post supernova orbital angular momentum axis is believed to be small."
+" However, there is no well constrained value for this angle."," However, there is no well constrained value for this angle."
+ Ferdman et al. (, Ferdman et al. (
+2008) and Breton (2008) argue that there are two possible models to explain the observed pulse profile: the single-cone and two-cone emission model.,2008) and Breton (2008) argue that there are two possible models to explain the observed pulse profile: the single-cone and two-cone emission model.
+" In this context, the number of emission cones means the number of emitting cones that intersect with our line of sight."," In this context, the number of emission cones means the number of emitting cones that intersect with our line of sight."
+ Each model gives a different spin tilt angle estimate., Each model gives a different spin tilt angle estimate.
+ Ferdman et al. (, Ferdman et al. (
+"2008) estimate this angle to be « 15? for the single- model and « 6.1? for thetwo-cone model, both at 68.3 confidence.","2008) estimate this angle to be < $^\circ$ for the single-cone model and < $^\circ$ for thetwo-cone model, both at 68.3 confidence."
+" Due to this uncertainty, we take the spin tilt angle 0,to be « 15? in our work, which corresponds to 68.3 confidence value in the single-cone model and 95.4"," Due to this uncertainty, we take the spin tilt angle $\theta_t$ to be < $^\circ$ in our work, which corresponds to 68.3 confidence value in the single-cone model and 95.4"
+matter with the SMDITI is far from clear.,matter with the SMBH is far from clear.
+ Some of the outstanding questions are: This article describes how future X-ray observations may attempt to disentangle these phenomena., Some of the outstanding questions are: This article describes how future X-ray observations may attempt to disentangle these phenomena.
+ The accretion disk model is highly successfu at explaining the X-ray reprocessing s2οςστα observed m many ACN., The accretion disk model is highly successful at explaining the X-ray reprocessing spectrum observed in many AGN.
+ A small iiuber of ACN (AICC6-30-15. Tanaka e al.," A small number of AGN (MCG–6-30-15, Tanaka et al."
+ 1995: NGC 3516. Nandra e al.," 1995; NGC 3516, Nandra et al."
+ 1999: NGC 1151. Wang et al.," 1999; NGC 4151, Wang et al."
+ 1999) have been the subject of very ong inteeyatious Wwith vieldiug hie1 quality iron line profiles which match he predictious of re aceretion disk mode well (Fabian ( al., 1999) have been the subject of very long integrations with yielding high quality iron line profiles which match the predictions of the accretion disk model well (Fabian et al.
+ 1995)., 1995).
+ However. there are still ambienles preseut.," However, there are still ambiguities present."
+ Firstly. a tiuc-averaeged iron line profile οςfais no information about the mass Q tje ceutral back hole.," Firstly, a time-averaged iron line profile contains no information about the mass of the central black hole."
+ A] paxauueters relevait to determining the line profile scale witi the eravitational radius., All parameters relevant to determining the line profile scale with the gravitational radius.
+ Secoidlv. aud nore 1teresting from au astroplivsics poiut of view. t1e line profi eds sensitive to the Nouv source ecolmetry. accretion disk. stcture (inchdiic the region inke the immernost stable orbi). aud he spin of the SMDII.," Secondly, and more interesting from an astrophysics point of view, the line profile is sensitive to the X-ray source geometry, accretion disk structure (including the region inside the innermost stable orbit), and the spin of the SMBH."
+ DegeneLac desexists in the sense that differeu astrophysical asstuuptions aud space-tinie geouetres can produce very similar ion ine profiles., Degeneracies exists in the sense that different astrophysical assumptions and space-time geometries can produce very similar iron line profiles.
+ The best studied exauple of this (egeneracyv is the case of the verv-broad stae of the irou liue iu AICMCG6-30-15 nd by Iwasawa et al. (, The best studied example of this degeneracy is the case of the very-broad state of the iron line in MCG–6-30-15 found by Iwasawa et al. (
+1996).,1996).
+ Masing the sandard assumptions tha the iue endssion is axisviuiaetrie. aud there is oilv cluission from outside of the ius of uareilal staülitv. Twasawa et a. (," Making the standard assumptions that the line emission is axisymmetric, and there is only emission from outside of the radius of marginal stability, Iwasawa et al. ("
+1996) πιogested that the SMDIT in this ject unst be rapidly rotating to prodiος α liue as broad axd redshitted as tiat nu.,1996) suggested that the SMBH in this object must be rapidly rotating to produce a line as broad and redshifted as that seen.
+ Da)owski et al. (, Dabowski et al. (
+1996) computed erids of iron ine profiles for various values he SNBI sun wih the same asstations and set a formal limit of «-0.91 the sq)in of tus SMDIT.,1996) computed grids of iron line profiles for various values of the SMBH spin with the same assumptions and set a formal limit of $a>0.94$ on the spin of this SMBH.
+ However. Revrolds Beechnuan (1997) showed that he sonne mon line profile «uai result from à non-rotating SMDITI if a hiel-latitude N-rav source πιατος disk material within tfje radius of mareinal stability.," However, Reynolds Begelman (1997) showed that the same iron line profile can result from a non-rotating SMBH if a high-latitude X-ray source illuminates disk material within the radius of marginal stability."
+ This is ali explicit demonstration of how unucertaiuies in the assumed astroplivsics (6.8. he, This is an explicit demonstration of how uncertainties in the assumed astrophysics (e.g. the
+IL aud VLT data.,II and VLT data.
+ Ser &* was highly active curing this period showing several UR flares., Sgr A* was highly active during this period showing several IR flares.
+ Figure la, Figure 1a
+mereine events) iu common DAL haloes are also included. as described iu Mencei et al. (,merging events) in common DM haloes are also included as described in Menci et al. (
+2001).,2004).
+ ILowwever. several chanecs have been made iu both the DAL and the barvonic sector of our model with respect to our previous papers.," However, several changes have been made in both the DM and the baryonic sector of our model with respect to our previous papers."
+ We describe such changes in turn., We describe such changes in turn.
+ To derive the whole color distribution we run a Monte Carlo version of our SAM. where many realizations of the mereing histories of present-day DAL haloes (with differcut final masses) are drawn by extracting mereine probabilities according to the Extended Press Schechter Theory (Bond ct al.," To derive the whole color distribution we run a Monte Carlo version of our SAM, where many realizations of the merging histories of present-day DM haloes (with different final masses) are drawn by extracting merging probabilities according to the Extended Press Schechter Theory (Bond et al."
+ 1991: Lacey Cole 1993)., 1991; Lacey Cole 1993).
+ To generate the merging trees we adopt the same algoritlin described in Cole et al. (, To generate the merging trees we adopt the same algorithm described in Cole et al. (
+2000). with a erid of 25 final DM masses ranging from 3.16-10°AL. to La107AL. with equally spaced logaritlunic values. aud a ass resolution of 5.10*AL: a total of 100 realizations are drawn for cach final DAL mass.,"2000), with a grid of 25 final DM masses ranging from $3.16\cdot 10^9\,M_{\odot}$ to $1.8\cdot10^{15}\,M_{\odot}$ with equally spaced logarithmic values, and a mass resolution of $5\cdot 10^7\,M_{\odot}$; a total of 100 realizations are drawn for each final DM mass."
+ The evolution of DAL haloes described above is connected to the processes involving the barvous as described below., The evolution of DM haloes described above is connected to the processes involving the baryons as described below.
+ The mass 52. of cooled gas m a disk with radius à aud rotation velocity 04 is derived (Moenuci et al., The mass $m_c$ of cooled gas in a disk with radius $r_d$ and rotation velocity $v_d$ is derived (Menci et al.
+" 2002. and Somerville Primack 1999) by computing. for cach time-step. the increment in the cooling radius Ar, of the central galaxy in all he DM haloes."," 2002, and Somerville Primack 1999) by computing, for each time-step, the increment in the cooling radius $\Delta r_c$ of the central galaxy in all the DM haloes."
+" The corresponding increment of mass of the cooled eas is Aim,=LarpglhJACO asstunine a simple isothermal gas density profile py(r)xr?."," The corresponding increment of mass of the cooled gas is $\Delta m_c=4\, \pi\,r_c^2\rho_g(r_c)\,\Delta r_c$ , assuming a simple isothermal gas density profile $\rho_g(r)\propto r^{-2}$."
+ The normalization is set as to recover the hot eas mass when the density profile is integrated over the volume., The normalization is set as to recover the hot gas mass when the density profile is integrated over the volume.
+ After mereing events we determine whether the mass of the largest progenitor my Comprises ess than a fraction 1/2 of the post-merecr mass (as in Somerville Primack 1999): if this is the case (.c.. the mereing wartuers have comparable masses) we reset the cooling time aud cooling radius to zero (ax sugeested by recent. aimed wdrodvuamucal simulation of major galaxy mereiug. see Cox et al.," After merging events we determine whether the mass of the largest progenitor $m_1$ comprises less than a fraction $1/2$ of the post-merger mass (as in Somerville Primack 1999); if this is the case (i.e., the merging partners have comparable masses) we reset the cooling time and cooling radius to zero (as suggested by recent, aimed hydrodynamical simulation of major galaxy merging, see Cox et al."
+ 2000. and the eas is reheated to the virial temperature of the DM halo.," 2004), and the gas is reheated to the virial temperature of the DM halo."
+ The introduction of such a process suppresses the cooling iu the massive haloes. which undergo a larger iuuber of such major mereine eveuts compared to low-uass haloes.," The introduction of such a process suppresses the cooling in the massive haloes, which undergo a larger number of such major merging events compared to low-mass haloes."
+ Iu addition. the steepuess of tle faiut-cud slope of the iunuinositv function (LE hereafter. see Sect.," In addition, the steepness of the faint-end slope of the luminosity function (LF hereafter, see Sect."
+ 2.5 and fig., 2.5 and fig.
+ 1) may contribute to account for the nonexistence of extremely xieht (monster) galaxies (see Deuson ct al., 1) may contribute to account for the nonexistence of extremely bright (”monster”) galaxies (see Benson et al.
+ 2003)., 2003).
+ See also Sect., See also Sect.
+ 1 for a discussion ou this point., 4 for a discussion on this point.
+ As for the star formation. we assume the canonical Sceliniidt. form i.=if(yr). where ty—ο aud q is left as a free parameter.," As for the star formation, we assume the canonical Schmidt form $\dot 
+m_*=m_c/(q\,\tau_d)$, where $\tau_d\equiv r_d/v_d$ and $q$ is left as a free parameter."
+" At cach time step. the mass Arn), veturued from the cold gas content of the disk to the hot gas phase due to Supernovae (SNe) activity is estimated from canonical cnerey balance arguments (I&auffinan 1996. Ikaufftinaun Charlot 1998: sce also Dekel Biruboin 2005) as Ay,=EsxeytyAin.fe? where Av, is the mass of stars formed in the timestep.©35-10OAL. is the imber of SNe per unit solar mass (depending ou the asstaned IMF). Γον=10°!eres is the energy of ejecta of cach SN. and e. is the circular velocity of the galactic halo: ey=0.010.5 is the efitcicucy of the cnerey transfer to the cold iuterstellaz eas."," At each time step, the mass $\Delta m_h$ returned from the cold gas content of the disk to the hot gas phase due to Supernovae (SNe) activity is estimated from canonical energy balance arguments (Kauffman 1996, Kauffmann Charlot 1998; see also Dekel Birnboin 2005) as $\Delta 
+m_h=E_{SN}\,\epsilon_0\,\eta_0\,\Delta m_*/v_c^2$ where $\Delta m_*$ is the mass of stars formed in the timestep, $\eta\approx 3-5\cdot 10^{-3}/M_{\odot}$ is the number of SNe per unit solar mass (depending on the assumed IMF), $E_{SN}=10^{51}\,{\rm ergs}$ is the energy of ejecta of each SN, and $v_c$ is the circular velocity of the galactic halo; $\epsilon_0=0.01-0.5$ is the efficiency of the energy transfer to the cold interstellar gas."
+" The above mass An), is made available for cooling at the next. timestep.", The above mass $\Delta m_h$ is made available for cooling at the next timestep.
+ The model free parameters ¢=30 aud ej=0.1 are chosen as to match the local D-baud LF aud the Tih-Fisher relation adopting a Salpeter IAIF., The model free parameters $q=30$ and $\epsilon_0=0.1$ are chosen as to match the local B-band LF and the Tully-Fisher relation adopting a Salpeter IMF.
+" The above star formation law depeuds strouglv on ry cuterime the timescale τη,", The above star formation law depends strongly on $r_d$ entering the timescale $\tau_d$.
+" The disk radius ry=reg(e) is related to the DAL virial radius kc, by the function gle.) accounting for the properties of the disk: we shall adopt as our fiducial choice the model by Mo. Mao White (1998) which vields gle.)2(AV2)Qusfm)fe.)127Foley. whore we take for the DAL spin parameter the average value A=0.05. while the ratio of the eas to DM angular momentum is kept to dd=0.05 as in Mo. Mao White (1998)."," The disk radius $r_d=r_v\,g(v_c)$ is related to the DM virial radius $r_v$ by the function $g(v_c)$ accounting for the properties of the disk; we shall adopt as our fiducial choice the model by Mo, Mao White (1998) which yields $g(v_c)\approx (\lambda/\sqrt{2})[j_d/(m_c/m)]f_c(v_c)^{-
+1/2}\,f_R(v)$, where we take for the DM spin parameter the average value $\lambda=0.05$, while the ratio of the gas to DM angular momentum is kept to $j_d=0.05$ as in Mo, Mao White (1998)."
+" The slowly varving fictions f. aud fg account for the eas density profile aud for the self-gravitv of the disk: the former is assumed to have the form given by Navarro. Freeuk White (1996): the concentration parameter cutering their profile depends on the mass, aud is computed following the procedure given in the appendix of the above paper."," The slowly varying functions $f_c$ and $f_R$ account for the gas density profile and for the self-gravity of the disk; the former is assumed to have the form given by Navarro, Freenk White (1996); the concentration parameter entering their profile depends on the mass, and is computed following the procedure given in the appendix of the above paper."
+ Note that our form of gfe) vields a star formation timescale 7;Xm/m. and hence appreciably affects the star formation in the most massive haloe where Óma is not effectively balanced by Supernovac feedback.," Note that our form of $g(v)$ yields a star formation timescale $\tau_d\propto m/m_c$, and hence appreciably affects the star formation in the most massive haloe where $\dot m_*$ is not effectively balanced by Supernovae feedback."
+ Tn particulary. such haloes are caracterized by a larger i. at high redshifts (where rapid cooling vields larger m0) ratios). which miplies faster coustunption of cold gas: at low-:. conversely. the longer star formation timescales due to the low afi ratio inhibits the formation of stars in galaxies within massive DM haloes.," In particular, such haloes are caracterized by a larger $\dot m_*$ at high redshifts (where rapid cooling yields larger $m_c/m$ ratios), which implies faster consumption of cold gas; at $z$, conversely, the longer star formation timescales due to the low $m_c/m$ ratio inhibits the formation of stars in galaxies within massive DM haloes."
+ We stummarize in fe., We summarize in fig.
+" 1 some results of the model obtained for our fiducial choice of cosinological parameters: O9=0.3. Q4=0.7. OQ,=0.05 and My=ταιςtAMpe+: the iietalliitv aud the dust extinction are computed as in Meuci et al. ("," 1 some results of the model obtained for our fiducial choice of cosmological parameters: $\Omega_0=0.3$, $\Omega_{\Lambda}=0.7$, $\Omega_b=0.05$ and $H_0=70\,{\rm km\,s^{-1}\,Mpc^{-1}}$; the metallicity and the dust extinction are computed as in Menci et al. ("
+2002). ax well as the evolution of the stellay populations. with cussion derived frou svuthetic spectral energy distributions (Druzual Charlot 1993). adopting a Salpeter IME.,"2002), as well as the evolution of the stellar populations, with emission derived from synthetic spectral energy distributions (Bruzual Charlot 1993), adopting a Salpeter IMF."
+ To test the cousisteucy of our set of cooling. star formation aud feedback laws with the available observations. we first compare with the local observed distribution ofcold gas. stars and of disk sizes uneastred at low-: (pauels à. b. c). aud with the Tully-Fisher relation iu the LEbaud shown iu pancl d).," To test the consistency of our set of cooling, star formation and feedback laws with the available observations, we first compare with the local observed distribution ofcold gas, stars and of disk sizes measured at $z$ (panels a, b, c), and with the Tully-Fisher relation in the I-band shown in panel d)."
+ Then we compare with the B-baud (at low :) aud the, Then we compare with the B-band (at low $z$ ) and the
+dust grains in the disc.,dust grains in the disc.
+ In this case. forward scattering on big dust grains at the outer edge of the disc is most effective in producing the observed radiation.," In this case, forward scattering on big dust grains at the outer edge of the disc is most effective in producing the observed radiation."
+ However. the expected polarization degree is small. if scattering occurs at small angels.," However, the expected polarization degree is small, if scattering occurs at small angels."
+ Radiative transfer modelling of scattering 1n a dise seen at iclination angles 80°<ix ddone in Bastien&Menard(1990) predicts à maximum polarization degree of8-11%., Radiative transfer modelling of scattering in a disc seen at inclination angles $\leq i \leq$ done in \citet{bastien} predicts a maximum polarization degree of.
+. Gomez's Hamburger (IRAS 18059-3211) is à good example of a high-inclination disc scattering. light from the central star (Woodetal.2008)., Gomez's Hamburger (IRAS 18059-3211) is a good example of a high-inclination disc scattering light from the central star \citep{wood08}.
+. This object shows a maximum polarization of in outer faint regions of the observed image (Ruizetal.1987).. Le.. regions produced by seattering at large angles.," This object shows a maximum polarization of in outer faint regions of the observed image \citep{ruiz}, i.e., regions produced by scattering at large angles."
+ In the central. bright regions. produced by forward scattering. the observed polarization is much lower.," In the central, bright regions, produced by forward scattering, the observed polarization is much lower."
+ If the image of Gomez's Hamburger were not resolved. as for V4332 Ser. the net polarization would have been observed at ~ or less.," If the image of Gomez's Hamburger were not resolved, as for V4332 Sgr, the net polarization would have been observed at $\sim$ or less."
+ Kaminskietal.(2010) also postulate that there Is an outflowing matter in V4332 Sgr. where the observed emission features of atoms and molecules are formed by absorption and re-emission of the radiation from the central object.," \citet{kst} also postulate that there is an outflowing matter in V4332 Sgr, where the observed emission features of atoms and molecules are formed by absorption and re-emission of the radiation from the central object."
+ This matter is cold. as indicated by the profiles of the emission molecular bands analysed in Kaminskietal.(2010).. and most probably contains dust.," This matter is cold, as indicated by the profiles of the emission molecular bands analysed in \citet{kst}, and most probably contains dust."
+ Thus the observed optical continuum can also be partly produced in the outflow scattering the central star radiation., Thus the observed optical continuum can also be partly produced in the outflow scattering the central star radiation.
+ If the outflow is concentrated along the dise axis. scattering would occur at angles close to907.. leading to a considerable polarization of the observed spectrum.," If the outflow is concentrated along the disc axis, scattering would occur at angles close to, leading to a considerable polarization of the observed spectrum."
+ Bastien&Menard(1990) (seealsoSole1989) found that for bipolar outflows from dises with inclination angles of tthe maximum polarization can reach60-94%., \citet{bastien} \citep[see also][]{solc} found that for bipolar outflows from discs with inclination angles of the maximum polarization can reach.
+. Therefore the high polarization of the observed spectrum of V4332 Sgr may indicate that this spectrum is produced by dust scattering not only in the disc but also in the matter outflowing along the dise axis., Therefore the high polarization of the observed spectrum of V4332 Sgr may indicate that this spectrum is produced by dust scattering not only in the disc but also in the matter outflowing along the disc axis.
+ In either case. the present polarization. measurements confirm the conclusion of Kaminskietal.(2010). that the main object in. V4332 Ser is now hidden in a dusty dise seen almost edge-on.," In either case, the present polarization measurements confirm the conclusion of \citet{kst} that the main object in V4332 Sgr is now hidden in a dusty disc seen almost edge-on."
+ This implies that the image of the object should display a dark lane across the image. às it observed e.g. in Gomez's Humburger.," This implies that the image of the object should display a dark lane across the image, as it observed e.g. in Gomez's Humburger."
+ In November 1997 an image of V4332 Ser was taken by HST (brieflydiscussedinKimeswenger 2006)., In November 1997 an image of V4332 Sgr was taken by HST \citep[briefly discussed in][]{kimes}.
+. No dark lane is seen but the image profile 15 not significantly different from that of the point source., No dark lane is seen but the image profile is not significantly different from that of the point source.
+ With the lower limit distance of 1 kpe. as derived in Kammiskietal. (2010).. the angular resolution of HST of 0.1 aresee implies that the linear dimensions of the V4332 Ser disc are <100 AU.," With the lower limit distance of 1 kpc, as derived in \citet{kst}, the angular resolution of HST of 0.1 arcsec implies that the linear dimensions of the V4332 Sgr disc are $\la$ 100 AU."
+ We do not know the nature of the disc. nor when it was formed.," We do not know the nature of the disc, nor when it was formed."
+ If it resulted from the 1994 eruption. e.g. as a remnant of a binary merger. its dimensions are expected to be much smaller that the above upper limit.," If it resulted from the 1994 eruption, e.g. as a remnant of a binary merger, its dimensions are expected to be much smaller that the above upper limit."
+expected to be only minor.,expected to be only minor.
+ In this case the inclination should be quite large and the system would be viewed nearly edge-on at all times., In this case the inclination should be quite large and the system would be viewed nearly edge-on at all times.
+ A strong flare-like event with 100 MK plasma is not expected to be produced by wind shocks even for magnetically channelled winds. thus these X-rays must be generated by magnetic reconnection.," A strong flare-like event with 100 MK plasma is not expected to be produced by wind shocks even for magnetically channelled winds, thus these X-rays must be generated by magnetic reconnection."
+ X-ray emitting structures on the surface of IQ Aur appear rather unlikely given the result from the X-ray diagnostics: additionally the magnetic field of Ap/Bp stars is overall a large scale one and structurally differs from coronae where most X-ray plasma is confined in small-scale structures., X-ray emitting structures on the surface of IQ Aur appear rather unlikely given the result from the X-ray diagnostics; additionally the magnetic field of Ap/Bp stars is overall a large scale one and structurally differs from coronae where most X-ray plasma is confined in small-scale structures.
+ Therefore in the MCWS scenario the flare requires the presence of a rigidly rotating equatorial disk as predicted by the above outlined MHD simulations (?).., Therefore in the MCWS scenario the flare requires the presence of a rigidly rotating equatorial disk as predicted by the above outlined MHD simulations \citep{dou08}.
+ In this scenario magnetic reconnection occurs. leading to the production of hot X-ray emitting plasma and further. episodic centrifugal breakouts of confined plasma produce fare-like events. originally proposed to explain X-ray flares on the Bp star c Ori (see discussion below) and (?)..," In this scenario magnetic reconnection occurs, leading to the production of hot X-ray emitting plasma and further, episodic centrifugal breakouts of confined plasma produce flare-like events, originally proposed to explain X-ray flares on the Bp star $\sigma$ Ori (see discussion below) and \citep{dou06}."
+ In contrast to coronal flares. these magnetic reconnection events are driven by centrifugal mass ejections that disrupt outer disk regions.," In contrast to coronal flares, these magnetic reconnection events are driven by centrifugal mass ejections that disrupt outer disk regions."
+ During breakout. the centrifugal forces on the accumulated disk matter stretch the field lines more and more outwards and finally reconnection occurs.," During breakout, the centrifugal forces on the accumulated disk matter stretch the field lines more and more outwards and finally reconnection occurs."
+ The event is accompanied by the ejection of outer disk mass. an inward propagating rebound and releases sufficient energy to produce hard X-ray flares.," The event is accompanied by the ejection of outer disk mass, an inward propagating rebound and releases sufficient energy to produce hard X-ray flares."
+ The finding of metallicity changes during the flare implies that the accumulated material in the circumstellar disk. which ts heated during the disruption. is more metal rich compared to the plasma that generates the quasi-quiescent X-ray emission.," The finding of metallicity changes during the flare implies that the accumulated material in the circumstellar disk, which is heated during the disruption, is more metal rich compared to the plasma that generates the quasi-quiescent X-ray emission."
+ Our rough estimations of the size of the flaring structure would imply that the outbreak-flare from IQ Aur was a more local event in the disk. involving a large clump or segment rather than the entire disk.," Our rough estimations of the size of the flaring structure would imply that the outbreak-flare from IQ Aur was a more local event in the disk, involving a large clump or segment rather than the entire disk."
+ Subsequent energy release or a change in the mner disk structure might be responsible for the flattening of the flare decay and the continuously enhanced level of X-ray emission after the event., Subsequent energy release or a change in the inner disk structure might be responsible for the flattening of the flare decay and the continuously enhanced level of X-ray emission after the event.
+ Overall this scenario appears suitable to also explain the flare event on IQ Aur., Overall this scenario appears suitable to also explain the flare event on IQ Aur.
+ In this case. we have likely observed the first X-ray flare from an AOp star.," In this case, we have likely observed the first X-ray flare from an A0p star."
+ As an alternative explanation it has been postulated that ray detected intermediate mass stars possess unknown late-type companions that actually produce the X-ray emission., As an alternative explanation it has been postulated that X-ray detected intermediate mass stars possess unknown late-type companions that actually produce the X-ray emission.
+ Considering the estimated age of IQ Aur. re. several ten up to hundred Myr and comparing its X-ray luminosity with studies of the 70 Myr old Pletades (?).. an active G or early-K dwarf as well as an M dwarf binary would be suitable candidate X-ray sources.," Considering the estimated age of IQ Aur, i.e. several ten up to hundred Myr and comparing its X-ray luminosity with studies of the 70 Myr old Pleiades \citep{mic96}, an active G or early-K dwarf as well as an M dwarf binary would be suitable candidate X-ray sources."
+ However. young active stars have typically significantly higher plasma temperatures than observed for IQ Aur (seee.g.?) and their X-ray light curves show quite frequent X-ray flares of factor up to a few. whereas IQ Aur is dominated by softer emission refcoolex)) and exhibits except for the large flare only minor variability.," However, young active stars have typically significantly higher plasma temperatures than observed for IQ Aur \citep[see e.g.][]{gue04} and their X-ray light curves show quite frequent X-ray flares of factor up to a few, whereas IQ Aur is dominated by softer emission \\ref{coolex}) ) and exhibits except for the large flare only minor variability."
+ Other properties like the high f/;-ratio in the triplet could be easily explained with a companion. typical densities are usually low and already a distance of 0.1. AU to IQ Aur would be in sufficient to completely erase the influence of the UV-field.," Other properties like the high $f/i$ -ratio in the triplet could be easily explained with a companion, typical densities are usually low and already a distance of 0.1 AU to IQ Aur would be in sufficient to completely erase the influence of the UV-field."
+ Also the occurrence of a large flare are no problem for an active late-type companion and the associated abundance changes are reminiscent of the chromospherie evaporation scenario describing the flare evolution., Also the occurrence of a large flare are no problem for an active late-type companion and the associated abundance changes are reminiscent of the chromospheric evaporation scenario describing the flare evolution.
+ Another criterion to consider is the fact that IQ Aur is also a bright radio source at ccm with logLe=16 ergss7!' HHz!. whereas a CCVn with logLa<14.9 ss7! HHz7! is radio dark (radiopropertiesfrom?)..," Another criterion to consider is the fact that IQ Aur is also a bright radio source at cm with $\log L_{6} \approx 16$ $^{-1}$ $^{-1}$, whereas $\alpha^2$ CVn with $\log L_{6} < 14.9$ $^{-1}$ $^{-1}$ is radio dark \citep[radio properties from][]{lin92}."
+ Since a rather tight correlation between X-ray and radio brightness measured at vz5 GHz is present for coronal sources. these measurements allow to test the companion scenario by checking the well established X-ray/Microwavetradio) correlation (also known as Giiddel-Benz relation). which predicts Lx/Lg=102-9? HHz for magnetically active stars (2)..," Since a rather tight correlation between X-ray and radio brightness measured at $\nu \gtrsim 5$ GHz is present for coronal sources, these measurements allow to test the companion scenario by checking the well established X-ray/Microwave(radio) correlation (also known as G{\""u}ddel-Benz relation), which predicts $L_{\rm X}/L_{\rm R}=10^{15.5 \pm 0.5}$ Hz for magnetically active stars \citep{gue93}."
+ With a quasi-quiescent X-ray luminosity of logZx=29.6 ergss!. IQ Aur is radio-overluminous by roughly two orders of magnitudes and thus significantly violates this correlation.," With a quasi-quiescent X-ray luminosity of $\log L_{\rm X}\approx29.6$ $^{-1}$, IQ Aur is radio-overluminous by roughly two orders of magnitudes and thus significantly violates this correlation."
+ The deviation is well in excess of the typical observed scatter that is. typically of a few., The deviation is well in excess of the typical observed scatter that is typically of a few.
+ Note that similar X-ray and radio fluxes were repeatedly observed for IQ Aur. making intrinsic variability quite unlikely to be the reason for the discrepancy.," Note that similar X-ray and radio fluxes were repeatedly observed for IQ Aur, making intrinsic variability quite unlikely to be the reason for the discrepancy."
+ Also the B8p He-w/Si star HD 144334 9929 Sco) with logLao.x eress! HHz! and logLx.=30.3 erg/s violates the Güddel-Benz relation by one to two orders of magnitude. similar to the B8V He-w star HD 146001 with a radio detection atlogZgx16.0 ergss! HHz! (plus one ULL. at logLa<15.9 eress! HHz! ).," Also the B8p He-w/Si star HD 144334 929 Sco) with $\log L_{6} \approx 16.0$ $^{-1}$ $^{-1}$ and $\log L_{\rm X} \approx 30.3$ erg/s violates the Güddel-Benz relation by one to two orders of magnitude, similar to the B8V He-w star HD 146001 with a radio detection at $\log L_{6} \approx 16.0$ $^{-1}$ $^{-1}$ (plus one U.L. at $\log L_{6} \lesssim 15.9$ $^{-1}$ $^{-1}$ )."
+ An even stronger violation is observed for the B8IIIp He-w/Si star HD 142301 9927 Sco). with logLy-17.1 eress! HHz! and logLx.=28.1 ss. this star is radio overluminous by more than four orders of magnitude.," An even stronger violation is observed for the B8IIIp He-w/Si star HD 142301 927 Sco), with $\log L_{6} \approx 17.1$ $^{-1}$ $^{-1}$ and $\log L_{\rm X} \approx 28.1$ $^{-1}$, this star is radio overluminous by more than four orders of magnitude."
+ On the other hand. all non X-ray detected comparison stars studied by ?.. i.e. HD 122532. HD 147010. remained. also undetected at radio wavelengths at a limit of logLa€15.7 eress! Hz!.," On the other hand, all non X-ray detected comparison stars studied by \cite{lin92}, i.e. HD 122532, HD 147010, remained also undetected at radio wavelengths at a limit of $\log L_{6} \lesssim 15.7$ $^{-1}$ $^{-1}$."
+ This points to a joint occurrence of strong X-ray and radio emission for the sample stars., This points to a joint occurrence of strong X-ray and radio emission for the sample stars.
+ The fact that IQ Aur and all other X-ray detected Ap/Bp stars violate the Güddel-Benz relation by orders of magnitudes suggests. that their emission does not originate from a late-type dwarfs. but a contributing companion in individual objects is not ruled out.," The fact that IQ Aur and all other X-ray detected Ap/Bp stars violate the Güddel-Benz relation by orders of magnitudes suggests, that their emission does not originate from a late-type dwarfs, but a contributing companion in individual objects is not ruled out."
+ The MCWS model has also been applied to a variety of hotter stars and it is illustrative to compare them with IQ Aur., The MCWS model has also been applied to a variety of hotter stars and it is illustrative to compare them with IQ Aur.
+ A well studied example is the magnetic OS star 4! CC. where ? find that the dynamic MCWS model describes the X-ray properties of the star very well.," A well studied example is the magnetic O5 star $\theta ^1$ C, where \cite{gag05}
+ find that the dynamic MCWS model describes the X-ray properties of the star very well."
+" It has an average X-ray luminosity of Ly=|10? eress7! and utilizing among other diagnostics f/i-line ratios. they conclude that the X-ray emitting plasma is located at distances of only 0.2...0.5R, rather close to the surface of the star."," It has an average X-ray luminosity of $L_{\rm X} =1 \times 10^{33}$ $^{-1}$ and utilizing among other diagnostics $f/i$ -line ratios, they conclude that the X-ray emitting plasma is located at distances of only $0.2 \dots 0.5\,R_{*}$ rather close to the surface of the star."
+ 6! CC exhibits rotational modulated X-ray emission. since It ds à tilted magnetic rotator with /=45° and B= 42°. leading to," $\theta ^1$ C exhibits rotational modulated X-ray emission, since it is a tilted magnetic rotator with $i=45^{\circ}$ and $\beta =42^{\circ}$ , leading to"
+environment. and does not depeud ou the mass of the galaxy or on the iustauitaneons star formation Thus. feedback in our model is specified by three parameters: €rchear:€lialos€aw,"environment, and does not depend on the mass of the galaxy or on the instantaneous star formation Thus, feedback in our model is specified by three parameters: $\erh, \eho, \esw$."
+ If is useful to note that the total energv available froii supernovae is of order LOere for cach Solar mass of stars that is formed. for a standard IME.," It is useful to note that the total energy available from supernovae is of order $10^{49} {\hbox{erg}}$ for each Solar mass of stars that is formed, for a standard IMF."
+Alost stars in. the field⋅ are observed to be in. multiple. (e.g.227).,"Most stars in the field are observed to be in multiple \citep[e.g.][]{Duquennoy91,Fischer92}."
+ The binary fraction of voung stars appears to be even higher: than that of the field⋅ (e.g.n2???.andreferences therein).," The binary fraction of young stars appears to be even higher than that of the field \citep[e.g.][and references therein]{Mathieu94, Patience02, Goodwin05,
+Goodwin07}."
+ ? show that in dense clusters it is very dillicult to make significant numbers of binaries from initially single stars and that it is impossible to widen an initially. narrow separation cistribution to the observed wide ciüstribution., \citet{Kroupa01b} show that in dense clusters it is very difficult to make significant numbers of binaries from initially single stars and that it is impossible to widen an initially narrow separation distribution to the observed wide distribution.
+ Therefore. the field binary fractions and properties must in some way mirror the outcome of star formation.," Therefore, the field binary fractions and properties must in some way mirror the outcome of star formation."
+ Thus. one of the most. significant constraints on models of star formation is that these models correctly predict the numbers and properties of binary and multiple systems (2).," Thus, one of the most significant constraints on models of star formation is that these models correctly predict the numbers and properties of binary and multiple systems \citep{Goodwin07}."
+ However. at least 50 per cent (probably τὸ 90 per cent) of stars are thought to form in clusters (e.g.2)..," However, at least 50 per cent (probably 70 – 90 per cent) of stars are thought to form in clusters \citep[e.g.][]{Lada03}."
+ H is known hat internal dynamical processes in clusters can rapidly and significant∢⊀⋅ allect the properties of binaries in clusters (c.g.227777). U," It is known that internal dynamical processes in clusters can rapidly and significantly affect the properties of binaries in clusters \citep*[e.g.][]{Heggie75, Hills75a, Hills75b, Kroupa95a, Kroupa95b, Kroupa99}."
+AThe degree to which binary properties will xf altered. depends! also on the density and. lifetime of the cluster in which the binaries are born., The degree to which binary properties will be altered depends also on the density and lifetime of the cluster in which the binaries are born.
+ Phe field binary population is thus a mixture of binaries rom cilferent density environments (which may or may not » The same initial population) which have been dynamically ocessed in different ways., The field binary population is thus a mixture of binaries from different density environments (which may or may not be the same initial population) which have been dynamically processed in different ways.
+ Therefore. the outcome οἱ (hvelrodvnamical) star. formation simulations should not. and cannot. be directly compared with the field population.," Therefore, the outcome of (hydrodynamical) star formation simulations should not, and cannot, be directly compared with the field population."
+ This paper is the first in a series in which we attempt to reconstruct the initial binary population resulting from the carly stages (Class 0/1) of star formation., This paper is the first in a series in which we attempt to reconstruct the initial binary population resulting from the early stages (Class 0/I) of star formation.
+ This paper will investigate the dvnamical processing of binarics in clusters of different densities and examine the extent to which binarics of dillerent. separations are. processed., This paper will investigate the dynamical processing of binaries in clusters of different densities and examine the extent to which binaries of different separations are processed.
+" This will allow us to discuss the main aspects of attempting to model the origin of⋅ the field.⋅ and reconstruct⋅. the initial""EE binary⋠⋅⇁ population."," This will allow us to discuss the main aspects of attempting to model the origin of the field, and reconstruct the initial binary population."
+R In this paper. we investigate how a Licled-like initial binary population is modified in clusters.," In this paper, we investigate how a field-like initial binary population is modified in clusters."
+ In Section 2. we," In Section \ref{method}
+ we"
+13mm (Paper 11). and by. iing a detailed: model of the racliative transfer. we investigate the physical parameters of this core.,"1.3mm (Paper II), and by using a detailed model of the radiative transfer, we investigate the physical parameters of this core."
+ The remainder of t1ο paper is laicl out as follows: Section 2 presents the C1°O οoservations of LIGSOB. Section 3 presents a racliative transfer model of LI689DB. In Section 4 the variaion of density and. temperature of the eas is explored. ancl we introduce he L3mm data to show how that leads o further constraints. including the variation of the abundance of CO realive to the dust.," The remainder of the paper is laid out as follows: Section 2 presents the $ \rm ^{18}$ O observations of L1689B. Section 3 presents a radiative transfer model of L1689B. In Section 4 the variation of density and temperature of the gas is explored, and we introduce the 1.3mm data to show how that leads to further constraints, including the variation of the abundance of CO relative to the dust."
+ Section. 5r suminarises the main conclusions., Section 5 summarises the main conclusions.
+ The observations were carried out at the James Clerk Alaxwell Telescope. located. on Mauna Kea. Llawaii. on 1995 July Isth. 22nd and 25th 17:3001:30 = 0:3:30.11:30). on 1996 August 31st. ancl September Ist 17:3001:30 T= 03:30.11:30).," The observations were carried out at the James Clerk Maxwell Telescope, located on Mauna Kea, Hawaii, on 1995 July 18th, 22nd and 25th 17:30–01:30 = 03:30–11:30), on 1996 August 31st and September 1st 17:30–01:30 = 03:30–11:30)."
+ The HC) :2) transition. with a rest. frequency of 329.33 CGllz. and he CO 11) transition. with a rest [requeney. of 29.56 CGllz. were observed using the common user heterodyne receivers RxBSi (Cunningham οἱ al.," The $ \rm ^{18}$ O $ \rm \rightarrow$ 2) transition, with a rest frequency of 329.33 GHz, and the $ \rm ^{18}$ O $ \rm \rightarrow$ 1) transition, with a rest frequency of 219.56 GHz, were observed using the common user heterodyne receivers RxB3i (Cunningham et al."
+ 1992) and RxAP (Davies et. al., 1992) and RxA2 (Davies et al.
+ 1992)., 1992).
+ The JCALLT hall-power beam-wicth (LIDA) is 19 arcsec at 220 GlIz and 14 aresece ad 329 Gllz., The JCMT half-power beam-width (HPBW) is 19 arcsec at 220 GHz and 14 arcsec at 329 GHz.
+ Double-sideband: system. teniperatures were 200012000 Ix for receiver D31 ancl 480Ix. [or receiver A2 observations., Double-sideband system temperatures were 2000–12000 K for receiver B3i and 480K for receiver A2 observations.
+ The backend used was a digital auto-correlation spectrometer (DAS). with a resolution of 37SkI1z and 95 klz per channel. for the 2) and +1) transitions respectively. corres»ding to 0.34 + at 3):uad 0.13 tat 0D.," The backend used was a digital auto-correlation spectrometer (DAS), with a resolution of 378kHz and 95 kHz per channel, for the $ \rm \rightarrow$ 2) and $ \rm \rightarrow$ 1) transitions respectively, corresponding to 0.34 $ \rm ^{-1}$ at $ \rm \rightarrow$ 2) and 0.13 $ \rm ^{-1}$ at $ \rm \rightarrow$ 1)."
+ Data reduction was carried ou using theSPECN package (Packman 1990)., Data reduction was carried out using the package (Padman 1990).
+ The weather during the observations was in general eood for millimetre (RxA2) observations and. short. periods. proved adequate or observations of the submillimetre 2) Uransition., The weather during the observations was in general good for millimetre (RxA2) observations and short periods proved adequate for observations of the submillimetre $ \rm \rightarrow$ 2) transition.
+ Regular pointing checks on bright sources. observed with the continuum backend system. were carried. out," Regular pointing checks on bright sources, observed with the continuum backend system, were carried out"
+alc| A straight forward caleulation gives The C* algebra A is defined as the algebra as the free stun of (non necessarily connected) 4-valent spin networks over SO(4) as in (191).,and A straight forward calculation gives The $C^*$ algebra $\cal A$ is defined as the algebra as the free sum of (non necessarily connected) 4-valent spin networks over $SO(4)$ as in \ref{eq:v}) ).
+ The « operation is simply defined by complex conjugation of the components of .1., The $*$ operation is simply defined by complex conjugation of the components of $A$.
+ We define the functional Wo over the algebra bv means of the corresponding Wightnan functions of our spin foam model., We define the functional $W$ over the algebra by means of the corresponding Wightman functions of our spin foam model.
+ The [act (hat the positivity condition holds for W (namely. HWCA)> 0) can be easily seen from the form o [the functional measure.," The fact that the positivity condition holds for $W$ (namely, $W({
+A}^*{A}) \ge 0$ ) can be easily seen from the form of the functional measure."
+" We can explicitly construct an orthonormal basis i YH», as follows.", We can explicitly construct an orthonormal basis in ${\cal H}_{ph}$ as follows.
+ There ave two kind of situations: spin network which do not interfere with the vacuum |0) (the empty spin network). and those which do.," There are two kind of situations: spin network which do not interfere with the vacuum $\left|0 \right>$ (the empty spin network), and those which do."
+" In the first case the projection is trivial and the elements of the physical Hilbert space Hp, are the simply the original spin network states.", In the first case the projection is trivial and the elements of the physical Hilbert space ${\cal H}_{ph}$ are the simply the original spin network states.
+ some examples are the lollowing, Some examples are the following
+Obsenration by dust. affects seriously anv measurement of galaxy luminosities at the optical aud. UV. wavelengths.,Obscuration by dust affects seriously any measurement of galaxy luminosities at the optical and UV wavelengths.
+ Correcting for (hiis effect is a difficult task as measurements of parameters sensitive to obscuration are not always easy {ο obtain., Correcting for this effect is a difficult task as measurements of parameters sensitive to obscuration are not always easy to obtain.
+ In the case of obscuration corrections lor emission lines. i( is common to assume an average value of 1 in and the extinction curve of ?. to apply Che value to other emission lines.," In the case of obscuration corrections for emission lines, it is common to assume an average value of $^m$ in and the extinction curve of \citet{1989ApJ...345..245C} to apply the value to other emission lines."
+" This value of 4,5, has been found as (he average of the obscuration of several samples of local emiütüng galaxies (e.g.?2?77).."," This value of $A_{H\alpha}$ has been found as the average of the obscuration of several samples of local emitting galaxies \citep[e.g.][]{1983ApJ...272...54K,1995ApJ...455L...1G,2003ApJ...599..971H}."
+ This procedure has the advantage that is very simple to apply to a luminosity Fancetion or huminosity density. since it is a scaling factor.," This procedure has the advantage that is very simple to apply to a luminosity function or luminosity density, since it is a scaling factor."
+" It is worth to note here that the average obscuration applied in UV studies is about ely=1"" (?. ee). "," It is worth to note here that the average obscuration applied in UV studies is about $A_V=1^m$ \citep[e.g.]{1996MNRAS.281..847T,2000MNRAS.312..442S}. ."
+In the GALEN far-ultraviolet (FUV). ? found a mean obseuration of App=LAU for galaxies selected in the GALEN near-ultraviolet band.," In the GALEX far-ultraviolet (FUV), \citet{2005ApJ...619L..51B}
+ found a mean obscuration of $A_{FUV}=1.1^m$ for galaxies selected in the GALEX near-ultraviolet band."
+ This value is roughly Aec90.4., This value is roughly $A_V \simeq 0.4^m$.
+" Galaxies in the GALEN UV Imaging Atlas of Nearby Galaxies (??) have a nean obscuration. considering ouly spirals and irregulars. of yyy=1.8"". corresponding {ο Ayc07 (Gil de Paz. private communication) as this sample ean be considered as selected in (he optical."," Galaxies in the GALEX UV Imaging Atlas of Nearby Galaxies \citep{2004AAS...205.4201G,2006astro.ph..6440G}
+ have a mean obscuration, considering only spirals and irregulars, of $A_{FUV}=1.8^m$, corresponding to $A_V \simeq 0.7^m$ (Gil de Paz, private communication) as this sample can be considered as selected in the optical."
+" The average obscuration of ly,=1” corresponds to a value of about id,c1.22"" as ioled by ?..", The average obscuration of $A_{H\alpha}=1^m$ corresponds to a value of about $A_V\simeq 1.22^m$ as noted by \citet{2004ApJ...615..209H}.
+ This disagreement in the observation is a selection effect of the different (FUY. UV and }) samples.," This disagreement in the obscuration is a selection effect of the different (FUV, UV and ) samples."
+" Galaxies selected at redder wavelengths can have greater internal ohseurations than those selected at the UV. where heavily absorbed svstems will not be selected,"," Galaxies selected at redder wavelengths can have greater internal obscurations than those selected at the UV, where heavily absorbed systems will not be selected."
+ The value of A at the wavelength: of the emission line is obtained applving a given extinction curve., The value of $A$ at the wavelength of the emission line is obtained applying a given extinction curve.
+ We are assuming that the continuum and the emission line are affected by (he same obscuration. so the EW does not need obscuration correction.," We are assuming that the continuum and the emission line are affected by the same obscuration, so the EW does not need obscuration correction."
+ This last assumption is nol true if we use (he obscuration recipe of 2.. that is applied lor strong starbursts.," This last assumption is not true if we use the obscuration recipe of \citet{1996ApJ...458..132C}, that is applied for strong starbursts."
+ llowever. normal galaxies and moderate starbursts can be assumed as showing no obscuration correction in EW.," However, normal galaxies and moderate starbursts can be assumed as showing no obscuration correction in EW."
+ In (his Section. we study the underlying statistics of the magnitude-color diagram used to select ELGs in the most simple scenario. a broad-band filler aud a narrow-band. filter.," In this Section, we study the underlying statistics of the magnitude-color diagram used to select ELGs in the most simple scenario, a broad-band filter and a narrow-band filter."
+ A sample selected in broad-narrow color is directly selected in EW (Eq. 22))., A sample selected in broad-narrow color is directly selected in EW (Eq. \ref{eq:def:ew}) ).
+ Thus. ideally. a sample can be created. selecting the objects with color inexcess ofa mininum color," Thus, ideally, a sample can be created selecting the objects with color inexcess ofa minimum color"
+to simulate statistically equivalent data sets.,to simulate statistically equivalent data sets.
+ By statistically equivalent we mean datasets that contain the same variance and noise properties as the actual data., By statistically equivalent we mean datasets that contain the same variance and noise properties as the actual data.
+ It is worth noting. that unlike the direct. bootstrapping approach. where the actual datapoints are picked in random order in order to simulate synthetic light curves. the AL approach retains the correlations in between the neighbouring data points. thus retaining the red noise properties of the dataset.," It is worth noting, that unlike the direct bootstrapping approach, where the actual datapoints are picked in random order in order to simulate synthetic light curves, the AR approach retains the correlations in between the neighbouring data points, thus retaining the red noise properties of the dataset."
+ In order to choose the right order for the AR. mocelling. we fitted. the light curves with several AR) models with an increasing order (up to 6th order)," In order to choose the right order for the AR modelling, we fitted the light curves with several AR models with an increasing order (up to 6th order)."
+ The fitting itself is carried out. by estimating the Ah cocllicients from the autocorrelations., The fitting itself is carried out by estimating the AR coefficients from the autocorrelations.
+ We then examined the resulting correlations (al clillerent time lags) to see when the correlations become insignificant OL level olf (Chatheld 1989)., We then examined the resulting correlations (at different time lags) to see when the correlations become insignificant or level off (Chatfield 1989).
+ As a result we find that typically a second: order autoreeressive model ie. where NX; ave the time series values and N(o) is the white noise component. is sullicient to. describe the correlation properties in the data.," As a result we find that typically a second order autoregressive model i.e. where $_{i}$ are the time series values and $\sigma$ ) is the white noise component, is sufficient to describe the correlation properties in the data."
+ Thus a second order autoregressive model (i.c. A11(2)) is used as a basis for the following significance estimates for the peaks seen in power spectra., Thus a second order autoregressive model (i.e. AR(2)) is used as a basis for the following significance estimates for the peaks seen in power spectra.
+ The procedure for obtaining these estimates is as follows: As a result. we will get 00.000 frequency. vs. maximum relative power pairs for cach of our real datasets.," The procedure for obtaining these estimates is as follows: As a result, we will get 100,000 frequency vs. maximum relative power pairs for each of our real datasets."
+ Vhese are then plotted in the frequency vs. power space together with the confidence contours tha encloseΟπή... and of the cases (Fig 2).," These are then plotted in the frequency vs. power space together with the confidence contours that enclose, and of the cases (Fig 2)."
+ Finally. we express the periodogram power in terms of relative »ower in the same Wavy as Wels done above in step 4 for the svnthetic data. and. overplot it (in red. Fig 2).," Finally, we express the periodogram power in terms of relative power in the same way as was done above in step 4 for the synthetic data, and overplot it (in red, Fig 2)."
+ Now. if the highest. peak in the observed periodogram falls outside the contours. the peak represents à," Now, if the highest peak in the observed periodogram falls outside the contours, the peak represents a"
+Hierarchical models of structure formation predict that galaxies form after a series of accretion events involving lower mass systems (??)..,"Hierarchical models of structure formation predict that galaxies form after a series of accretion events involving lower mass systems \citep{2003ApJ...591..499A,2007MNRAS.374.1479G}."
+ Within this scenario. à massive spiral galaxy such as the Milky Way should have accreted 100-200 satellites during its evolutionary history (?)..," Within this scenario, a massive spiral galaxy such as the Milky Way should have accreted 100-200 satellites during its evolutionary history \citep{2005ApJ...635..931B}."
+ This process would leave its imprint in a diffuse stellar halo extending out to 10-100 times =he optical radius of the galactic disc., This process would leave its imprint in a diffuse stellar halo extending out to 10-100 times the optical radius of the galactic disc.
+ The density. luminosity. and metallicity of the stellar halo can thus provide direct nformation about the history and evolution of a galaxy. and enable us to test hierarchical formation scenarios.," The density, luminosity, and metallicity of the stellar halo can thus provide direct information about the history and evolution of a galaxy, and enable us to test hierarchical formation scenarios."
+ The discovery of the Sagittarius stream in the Milky Way AW) revealed for the first time the role of satellite accretion n the build-up of a galaxy mass (?).., The discovery of the Sagittarius stream in the Milky Way (MW) revealed for the first time the role of satellite accretion in the build-up of a galaxy mass \citep{1994Natur.370..194I}.
+ Studies of the MW halo csing data from the Sloan Digitised Sky Survey (SDSS:?) revealed other tidal streams around our galaxy. proving that these features are relatively common and can survive for a few Gyrs (2222).," Studies of the MW halo using data from the Sloan Digitised Sky Survey \citep[SDSS;][]{2000AJ....120.1579Y} revealed other tidal streams around our galaxy, proving that these features are relatively common and can survive for a few Gyrs \citep{2006ApJ...642L.137B,2006ApJ...645L..37G,2007ApJ...658..337B,2009ApJ...693.1118G}."
+ Similar features were discovered around M31. the other massive spiral of the Local Group (LG). that implied this galaxy has experienced a more violent merger history than the MW (??)..," Similar features were discovered around M31, the other massive spiral of the Local Group (LG), that implied this galaxy has experienced a more violent merger history than the MW \citep{2006AJ....131.1436F,2008ApJ...682L..33F}."
+ A giant stellar stream with a neutral hydrogen (Ht) counterpart was found in this galaxy (????)..," A giant stellar stream with a neutral hydrogen ) counterpart was found in this galaxy \citep{1977MNRAS.181..573N,2001Natur.412...49I,2002AJ....124.1452F,2004ApJ...601L..39T}."
+ Much effort has been devoted to obtaining deeper optical images of the halo of M31 with the Pan-Andromeda Archaeological Survey (PAndAS;?).. which revealed a large-scale metal-poor stellar distribution (stretching out to 150 kpe). and uncovered additional extended substructures in the halo similar to the giant southern stream.," Much effort has been devoted to obtaining deeper optical images of the halo of M31 with the Pan-Andromeda Archaeological Survey \citep[PAndAS;][]{2009Natur.461...66M}, which revealed a large-scale metal-poor stellar distribution (stretching out to 150 kpc), and uncovered additional extended substructures in the halo similar to the giant southern stream."
+" While M31 clearly shows signs of disturbance in its spheroid and outer disc. M33. its brightest satellite and the third most massive member of the LG. has long been considered an example of a ""pure disc"" system evolving in relative isolation. because of the lack of a prominent bulge. a seemingly unperturbed stellar disc. and the absence of known nearby satellites."," While M31 clearly shows signs of disturbance in its spheroid and outer disc, M33, its brightest satellite and the third most massive member of the LG, has long been considered an example of a “pure disc” system evolving in relative isolation, because of the lack of a prominent bulge, a seemingly unperturbed stellar disc, and the absence of known nearby satellites."
+ The origin of bulgeless galaxies such as M33 is difficult to explain within the hierarchical structure formation scenario (?).. because mergers of galaxies are expected to form à bulge component.," The origin of bulgeless galaxies such as M33 is difficult to explain within the hierarchical structure formation scenario \citep{2001MNRAS.327.1334V}, because mergers of galaxies are expected to form a bulge component."
+ The inflow of gas at later times (z>1) from the local environment is thought to play a major role in the growth of the discs of these systems (??)..," The inflow of gas at later times $z > 1$ ) from the local environment is thought to play a major role in the growth of the discs of these systems \citep{2006MNRAS.368....2D,2009ApJ...694..396B}."
+ Nonetheless. the presence of a diffuse faint stellar halo even in disc galaxies seems to be expected within the framework of current galaxy formation scenarios. predicting that a diffuse stellar halo rather than a centrally-concentrated bulge would result after the late major-merger event experienced by these galaxies (?)..," Nonetheless, the presence of a diffuse faint stellar halo even in disc galaxies seems to be expected within the framework of current galaxy formation scenarios, predicting that a diffuse stellar halo rather than a centrally-concentrated bulge would result after the late major-merger event experienced by these galaxies \citep{2010arXiv1010.1004B}."
+ Studies of the outskirts of M33 have started providing evidence of an extended stellar distribution beyond the optical disc. although it remains unclear whether this component corresponds to a stellar halo or an extended disc.," Studies of the outskirts of M33 have started providing evidence of an extended stellar distribution beyond the optical disc, although it remains unclear whether this component corresponds to a stellar halo or an extended disc."
+" ?) derived a radial density stellar profile out to 1 degree from the centre of M33 finding a clear dropoff in the star counts as a function of radius (xΚΙ ΙΟ], that they interpreted as a evidence of a stellar spheroid component."," \citet{2004AJ....128..237B} derived a radial density stellar profile out to 1 degree from the centre of M33 finding a clear dropoff in the star counts as a function of radius $\propto R^{-1.46}$ ), that they interpreted as a evidence of a stellar spheroid component."
+ The metallicity peak of this sparse stellar population. [Fe/H] = -1.24 + 0.04. was lower than the abundance in the disc ([Fe/H] ~ -0.9). and comparable to the mean metallicity of a sample of halo globular clusters (2)..," The metallicity peak of this sparse stellar population, [Fe/H] = -1.24 $\pm$ 0.04, was lower than the abundance in the disc ([Fe/H] $\sim$ -0.9), and comparable to the mean metallicity of a sample of halo globular clusters \citep{2000AJ....120.2437S}."
+ In a spectroscopic survey of red giant branch (RGB) stars in M33. ?) found that the stellar velocity distributions could be described by three different components: the galaxy disc. à halo (with [Fe/H] = -1.5). and a third distribution possibly associated with a stellar stream.," In a spectroscopic survey of red giant branch (RGB) stars in M33, \citet{2006ApJ...647L..25M} found that the stellar velocity distributions could be described by three different components: the galaxy disc, a halo (with [Fe/H] = -1.5), and a third distribution possibly associated with a stellar stream."
+ Finally. the PAndAS survey revealed an extended stellar distribution around M33 well beyond the main optical disc out to 3 degrees (or a projected radius of ~ 45 kpe) from its centre (22).," Finally, the PAndAS survey revealed an extended stellar distribution around M33 well beyond the main optical disc out to 3 degrees (or a projected radius of $\sim$ 45 kpc) from its centre \citep{2009Natur.461...66M,2010ApJ...723.1038M}."
+ This structure could either be the stellar debris of an accreted satellite. or the result of a tidal interaction with M31.," This structure could either be the stellar debris of an accreted satellite, or the result of a tidal interaction with M31."
+ This latter interpretation would be supported by the similar spatial orientation of the stellar features and the M33 warp., This latter interpretation would be supported by the similar spatial orientation of the stellar features and the M33 warp.
+ The origin of the warp remains disputed and previous studies have proposed that it 15 the result of an interaction with its massive neighbour (?2)..," The origin of the warp remains disputed and previous studies have proposed that it is the result of an interaction with its massive neighbour \citep{1997ApJ...479..244C,2008MNRAS.390L..24B}."
+ Simulations of the relative motion of the M31/M33 system (???).. converge on the possibility of a close encounter between the two galaxies occurring a few Gyr ago with a pericenter distance around or greater than 40 kpe. which could strip material from the M33 outer dise without severely distorting or disrupting the bulk of the stellar dise.," Simulations of the relative motion of the M31/M33 system \citep{2005ApJ...633..894L,2009ApJ...703.1486P,2009Natur.461...66M}, converge on the possibility of a close encounter between the two galaxies occurring a few Gyr ago with a pericenter distance around or greater than 40 kpc, which could strip material from the M33 outer disc without severely distorting or disrupting the bulk of the stellar disc."
+Pulsating DA (LH-rich atmospheres) white chwarks. commonly known as ZZ Ceti or DAW variable stars. comprise the most numerous Class of compact pulsators.,"Pulsating DA (H-rich atmospheres) white dwarfs, commonly known as ZZ Ceti or DAV variable stars, comprise the most numerous class of compact pulsators."
+ They are located in a narrow ancl probably pure instability strip with cllective temperatures between 10500. Ix. and. 12500. Wk (e.g. Winect Wepler 2008: Fontaine Brassard 2008: Althaus et al.," They are located in a narrow and probably pure instability strip with effective temperatures between $10\,500$ K and $12\,500$ K (e.g., Winget Kepler 2008; Fontaine Brassard 2008; Althaus et al."
+ 2010a)., 2010a).
+ ZZ Ceti stars are characterized. |N multimode photometric variations of up to 0.30 mag causec bv nonracdial g-mode pulsations of low degree (£< 2) aux periods between TO and 1500 s. The driving mechanism thought to excite the pulsations is a sort. of combination of the &~ mechanism acting in the hyerogen partia ionization zone (Dolez Vauelair 1981: Dziembowski ]|xoester. 1981: Winget et al., ZZ Ceti stars are characterized by multimode photometric variations of up to $0.30$ mag caused by nonradial $g$ -mode pulsations of low degree $\ell \leq 2$ ) and periods between 70 and 1500 s. The driving mechanism thought to excite the pulsations is a sort of combination of the $\kappa-\gamma$ mechanism acting in the hydrogen partial ionization zone (Dolez Vauclair 1981; Dziembowski Koester 1981; Winget et al.
+" 1959) anc the “convective driving"" mechanism proposed first by Brickhill (1991) ane later re-examineck by Coldreich Wu (1999).", 1982) and the “convective driving” mechanism proposed first by Brickhill (1991) and later re-examined by Goldreich Wu (1999).
+ The later mechanism is supposed to be dominant once a thick convection zone has developed at the stellar surface., The later mechanism is supposed to be dominant once a thick convection zone has developed at the stellar surface.
+ White-dwarl asteroseismologv fully exploits the comparison between the observed. pulsation periods in white dwarfs and the periods computed. for appropriate theoretical models., White-dwarf asteroseismology fully exploits the comparison between the observed pulsation periods in white dwarfs and the periods computed for appropriate theoretical models.
+ H allows us to infer details of the origin. internal structure ancl evolution of white cwarks (Winget Wepler 2008: Fontaine Brassard 2008: Althaus οἱ al.," It allows us to infer details of the origin, internal structure and evolution of white dwarfs (Winget Kepler 2008; Fontaine Brassard 2008; Althaus et al."
+ 2010a)., 2010a).
+ In. particular. constraints on the stellar mass. the thickness of the outer envelopes. the core chemical composition. weak magnetic fields and slow rotation rates can be inferred. from the observed: period patterns of ZZ Ceti stars.," In particular, constraints on the stellar mass, the thickness of the outer envelopes, the core chemical composition, weak magnetic fields and slow rotation rates can be inferred from the observed period patterns of ZZ Ceti stars."
+ In addition. asteroscismology of ZZ Ceti stars is à valuable tool for studying axions (Isern et al.," In addition, asteroseismology of ZZ Ceti stars is a valuable tool for studying axions (Isern et al."
+ 1992: Corrsico et al., 1992; Córrsico et al.
+ 2001: Discholl-Ixim ct al., 2001; Bischoff-Kim et al.
+ 2008b: Isern ct al., 2008b; Isern et al.
+ 2010. Córrsico et al.," 2010, Córrsico et al."
+ 2011). crystallization (Montgomery Wineet 1999: C'órrsico et al.," 2011), crystallization (Montgomery Winget 1999; Córrsico et al."
+ 2004. 2005: Metcalfe ct al.," 2004, 2005; Metcalfe et al."
+ 2004: Ilxanaan et al., 2004; Kanaan et al.
+ 2005). and important properties of the outer convection. zones (Montgomery. 2005ab.. 2007).," 2005), and important properties of the outer convection zones (Montgomery 2005ab, 2007)."
+ Finally. the temporal changes in the observed stable periods allow the measurement of the white dwarf. evolutionary timescale and the detection possible planets orbiting white chwarls (Alullally et al.," Finally, the temporal changes in the observed stable periods allow the measurement of the white dwarf evolutionary timescale and the detection possible planets orbiting white dwarfs (Mullally et al."
+ 2008)., 2008).
+ Among the numerous ZZ Ceti stars currently. known (148 stars: Castanheira ct al., Among the numerous ZZ Ceti stars currently known (148 stars; Castanheira et al.
+ 2010). in this paper we will analyze 44 bright ZZ Ceti stars which are listed in Table 1 of Fontaine Brassarel (2008).," 2010), in this paper we will analyze 44 bright ZZ Ceti stars which are listed in Table 1 of Fontaine Brassard (2008)."
+ We defer to a future work the study of the fainter ZZ Ceti stars discovered within the Sloan Digital Sky Survey (SDSS) (Mukacdam et al., We defer to a future work the study of the fainter ZZ Ceti stars discovered within the Sloan Digital Sky Survey (SDSS) (Mukadam et al.
+ 2004: Mullally et al., 2004; Mullally et al.
+ 2005: Kepler et al., 2005; Kepler et al.
+ 2005b: Castanheira et al., 2005b; Castanheira et al.
+ 2006. 2007. 2010).," 2006, 2007, 2010)."
+ The first target star of our seismological survey is the most studied member of the class. the paracigniatic," The first target star of our seismological survey is the most studied member of the class, the paradigmatic"
+"the HB luminosity, set by the value of the He-core mass at the He-flash (see Salaris et al.","the HB luminosity, set by the value of the He-core mass at the He-flash (see Salaris et al."
+" 2002, for a detailed discussion)."," 2002, for a detailed discussion)."
+ The first detailed comparison between predicted and observed measurements of in galactic GCs was performed by Fusi Pecci et al. (, The first detailed comparison between predicted and observed measurements of in galactic GCs was performed by Fusi Pecci et al. (
+1990).,1990).
+" By employing the theoretical models available at that time, these authors found that the observed dependence of on the cluster metallicity was in good agreement with theory, but the zero point of the theoretical values was too small by ~0.4 mag."," By employing the theoretical models available at that time, these authors found that the observed dependence of on the cluster metallicity was in good agreement with theory, but the zero point of the theoretical values was too small by $\sim 0.4$ mag."
+" For some time, this result has been considered a clear drawback of standard theoretical models of low-mass RGB stars."," For some time, this result has been considered a clear drawback of standard theoretical models of low-mass RGB stars."
+" This result has been reanalyzed by Cassisi Salaris (1997), who concluded that their models provided a good match to the available observational data."," This result has been reanalyzed by Cassisi Salaris (1997), who concluded that their models provided a good match to the available observational data."
+" More recently, this topic has been reviewed by several authors (e.g., Zoccali et al."," More recently, this topic has been reviewed by several authors (e.g., Zoccali et al."
+" 1999, Ferraro et al."," 1999, Ferraro et al."
+" 1999, Di Cecco et al."," 1999, Di Cecco et al."
+" 2010), and from their comparisons between theory and observations one can draw the conclusion that lingering uncertainties on the HB theoretical brightness and the GC metallicity scale leave open the possibility that a discrepancy at the level of ~0.20 mag between theory and observations may exist."," 2010), and from their comparisons between theory and observations one can draw the conclusion that lingering uncertainties on the HB theoretical brightness and the GC metallicity scale leave open the possibility that a discrepancy at the level of $\sim$ 0.20 mag between theory and observations may exist."
+" In the following we investigate how the new rate for the VN(p,y)?O reaction affects the evolutionary lifetimes along the RGB and the brightness of the RGB bump, compared to calculations with the NACRE rate."," In the following we investigate how the new rate for the $^{14}N(p,\gamma)^{15}O$ reaction affects the evolutionary lifetimes along the RGB and the brightness of the RGB bump, compared to calculations with the NACRE rate."
+" Figure | displays the bolometric luminosity as a function of time from the end of the core H-burning stage to the TRGB, for a post-MS stellar model of a typical GC star (M=0.8Mo, Y=0.248, Z=0.002)."," Figure \ref{ltime} displays the bolometric luminosity as a function of time from the end of the core H-burning stage to the TRGB, for a post-MS stellar model of a typical GC star $M_{\odot}$, $Y$ =0.248, $Z$ =0.002)."
+ As expected from previous investigations (Imbriani et al., As expected from previous investigations (Imbriani et al.
+" 2004, Weiss et al."," 2004, Weiss et al."
+ 2005) the model computed with the new reaction rate reaches the end of the H-burning stage with an age larger by about 120 Myr (~1% age difference) compared to results with the NACRE rate., 2005) the model computed with the new reaction rate reaches the end of the H-burning stage with an age larger by about 120 Myr $\sim$ age difference) compared to results with the NACRE rate.
+ The impact of this age difference on the M.TO_age relation in GC isochrones can be estimated from Fig. 2.., The impact of this age difference on the $M_V^{TO}$ -age relation in GC isochrones can be estimated from Fig. \ref{mvage}.
+" The Mτο, relation obtained using the LUNA rate is systematically brighter compared to results with the NACRE rate.", The $M_V^{TO}$ -age relation obtained using the LUNA rate is systematically brighter compared to results with the NACRE rate.
+" As a consequence, for a fixed observed value of M.LO. LUNA isochrones give ages higher by ~0.7 Gyr, compared to calculations with the NACRE rate."," As a consequence, for a fixed observed value of $M_V^{TO}$, LUNA isochrones give ages higher by $\sim$ 0.7 $Gyr$, compared to calculations with the NACRE rate."
+" When the models displayed in Fig.1 reach the He-flash, the age difference between calculations with the NACRE or the LUNA rate is reduced from ~120 Myr to ~20 Myr, implying a slightly faster RGB evolution when the LUNA reaction rate is employed."," When the models displayed in \ref{ltime} reach the He-flash, the age difference between calculations with the NACRE or the LUNA rate is reduced from $\sim$ 120 Myr to $\sim$ 20 Myr, implying a slightly faster RGB evolution when the LUNA reaction rate is employed."
+ Figure 3 shows the evolution of the luminosity with time in the region of the RGB bump., Figure \ref{bump1} shows the evolution of the luminosity with time in the region of the RGB bump.
+" The LUNA rate causes a brighter bump (by ~0.06 mag), due to a slightly shallower convective envelope at the first dredge up, as shown by Fig. 4.."," The LUNA rate causes a brighter bump (by $\sim 0.06$ mag), due to a slightly shallower convective envelope at the first dredge up, as shown by Fig. \ref{bump2}."
+ This is due to the the fact that the LUNA based model has an hotter H-burning shell for a fixed He core mass., This is due to the the fact that the LUNA based model has an hotter H-burning shell for a fixed He core mass.
+" Finally, Fig."," Finally, Fig."
+ 5 shows the trend of the absolute visual magnitude of the RGB bump as a function of the metallicity for a 12 Gyr old stellar population., \ref{bump3} shows the trend of the absolute visual magnitude of the RGB bump as a function of the metallicity for a 12 Gyr old stellar population.
+" The LUNA rate causes a systematically brighter RGB bump in the models, for the whole Galactic GC metallicity range."," The LUNA rate causes a systematically brighter RGB bump in the models, for the whole Galactic GC metallicity range."
+" The difference with respect to results with the NACRE rate is of the order of 0.07 mag at the lowest metallicities, becoming almost negligible in the metal rich regime."," The difference with respect to results with the NACRE rate is of the order of 0.07 mag at the lowest metallicities, becoming almost negligible in the metal rich regime."
+sub-1u (Deckwithetal.2000: (Woodenetal.2007:Poutoppidan&Brearley2010..,"$\mu$ \citep{BE00,LD01,KE04,HG10}. \citealt{WO07,PB10},"
+ As time passes. the small dust responsible for the infrared (IR) excess observed around. voung stars ds subjected to different processes that affect. and will eventually determine how this progression will cud.," As time passes, the small dust responsible for the infrared (IR) excess observed around young stars is subjected to different processes that affect, and will eventually determine how this progression will end."
+ Planets and planetary svstenis have been observed around hundreds of stars other than the Sun. showing hat this result is rather common (Udrv&Sautos2007).," Planets and planetary systems have been observed around hundreds of stars other than the Sun, showing that this result is rather common \citep{US07}."
+. IR observations have revealed a ereat nuuber ofdisks. composed of large planctesimal rocks aud smaller )odies. around a varicty of stars spamming a larec ranec in spectral types and ages (Rickeetal.2005:Brydeneuteretal. 2009).," IR observations have revealed a great number of, composed of large planetesimal rocks and smaller bodies, around a variety of stars spanning a large range in spectral types and ages \citep{RI05,BR06,SU06,GA07,CA09}."
+. A few debris disks are kuownu to iarbor planets (e.g...) Pictoris aud Fomalhaut. Lagrangeetal.2010:alas 20051). although it is still nuclear whether this is often true (xdspalotal.2009).," A few debris disks are known to harbor planets (e.g., $\beta$ Pictoris and Fomalhaut, \citealt{LG10,KA08}) ), although it is still unclear whether this is often true \citep{KO09}."
+. The uajoritv of niüiu-sequenuce stars show no sigus of plancts or debris within the current observational Liuitatious. rowever. Indicating that the disks around such stars at he time of their formation have dissipated completely. caving no dust behind to tell the story.," The majority of main-sequence stars show no signs of planets or debris within the current observational limitations, however, indicating that the disks around such stars at the time of their formation have dissipated completely, leaving no dust behind to tell the story."
+ Which processes are imiportaut aud determinant for the aftermath of disk evolution are still uuder debate. and this topic is the subject of many theoretical auc observational studies over the last decade. stimulated in large by receut IR and (sub-anillineter facilities.," Which processes are important and determinant for the aftermath of disk evolution are still under debate, and this topic is the subject of many theoretical and observational studies over the last decade, stimulated in large by recent IR and (sub-)millimeter facilities."
+ Specifically on the subject of the muneralogical conrpositiou. spectra from the eround aud the," Specifically on the subject of the mineralogical composition, spectra from the ground and the"
+that beaming does not affect the relations between two conrponeuts discussed above (sce Fig.,that beaming does not affect the relations between two components discussed above (see Fig.
+ 5 in Chiahberee et al., 5 in Chiaberge et al.
+ 2000)., 2000).
+ Based on the TeV variability. it was estimated that Mrk 121 and ALK 501 have 0<G (Padovani 1999).," Based on the TeV variability, it was estimated that Mrk 421 and Mrk 501 have $\theta\lesssim6\arcdeg$ (Padovani 1999)."
+ Tf they were observed at 0=607. according to the Eq. (," If they were observed at $\theta=60\arcdeg$, according to the Eq. ("
+"5) aud assundue typical values for P aud p. ie. D—5 and p~3. the observed flux would decrease by a factor of 1.34101, which is not detectable even during the largest outburst.","5) and assuming typical values for $\Gamma$ and $p$ , i.e., $\Gamma\sim$ 5 and $p\sim$ 3, the observed flux would decrease by a factor of $1.3\times10^{-4}$, which is not detectable even during the largest outburst."
+ Iowever. if they were at the distance of Cen A. z-0.00085. the observed flux would increase by a factor of about (eP," However, if they were at the distance of Cen A, z=0.0008, the observed flux would increase by a factor of about $(\frac{0.034}{0.0008})^2$."
+ Therefore. if Mak 121 and Alek 501 were moved to the position of Cen A and observed at a viewing anele of GU. or moved to the position of AIST (z=0.0013) and observed at 0=357. they would still be detectable at TeV energies duriug outburst.," Therefore, if Mrk 421 and Mrk 501 were moved to the position of Cen A and observed at a viewing angle of $60\arcdeg$, or moved to the position of M87 (z=0.0043) and observed at $\theta=35\arcdeg$, they would still be detectable at TeV energies during outburst."
+ As mentioned in Section 1. BL Lac objects are divided into IIBLs and LBLs according to the peak locations iu the SEDs.," As mentioned in Section 1, BL Lac objects are divided into HBLs and LBLs according to the peak locations in the SEDs."
+ Further studies show that for LBLs the N-ravs are inverse-Couptou emission. while N-ravs in IIBEs are an extension of svuchrotron ciission at lower energies (Padovani Cüonuui 1996: Fossati et al.," Further studies show that for LBLs the X-rays are inverse-Compton emission, while X-rays in HBLs are an extension of synchrotron emission at lower energies (Padovani Giommi 1996; Fossati et al."
+ 1998)., 1998).
+ In. the following subsections we will show that N-ravs in Cou A ancl Ms cousist of svuchrotron enmission aud thus they are TIBLs., In the following subsections we will show that X-rays in Cen A and M87 consist of synchrotron emission and thus they are HBLs.
+ Figure 1l shows spectral cnerey distributions of Cen A. Ms? and Ak 501 based on several observations in the literature.," Figure 1 shows spectral energy distributions of Cen A, M87 and Mrk 501 based on several observations in the literature."
+ The N-vav spectral iudex of Cen À nucleus in 2-10 keV was found to be less than 1 (Turner ct al., The X-ray spectral index of Cen A nucleus in 2-10 keV was found to be less than 1 (Turner et al.
+ 1997: Mivazaki ct al., 1997; Miyazaki et al.
+" 1996). aud. heuce log(zf,) is a monotonously rising function of logip."," 1996), and hence $\log(\nu f_{\nu})$ is a monotonously rising function of $\log\nu$."
+ Between LO and 100 keV. the spectral index is 0.7 0.8 (Miyazaki et al.," Between 10 and 100 keV, the spectral index is 0.7 $-$ 0.8 (Miyazaki et al."
+ 1996: Dond et al., 1996; Bond et al.
+ 1996: Wheaton et al., 1996; Wheaton et al.
+ 1996: Iiuzer et al., 1996; Kinzer et al.
+ 1995). also less than 1. with a mouotonously rising SED of Cen A in this energv range.," 1995), also less than 1, with a monotonously rising SED of Cen A in this energy range."
+ The SED of Cen A continuously rises below ~ 150 keV where the spectiuu breaks. then the SED οσο» down coutiuuouslv to the MeV. aud. even GeV enerev range (see Fie.," The SED of Cen A continuously rises below $\sim$ 150 keV where the spectrum breaks, then the SED goes down continuously to the MeV and even GeV energy range (see Fig."
+ 1)., 1).
+ As in the case of blazars. this mnnup represents one of the two radiation commpoucuts. aud he radiation mechanisia of this component is the same as hat of N-ravs.," As in the case of blazars, this hump represents one of the two radiation components, and the radiation mechanism of this component is the same as that of X-rays."
+ Steiule et al. (, Steinle et al. (
+1998) aud Iiuzer et al. (,1998) and Kinzer et al. (
+1995) pointed out hat the spectra of Con A show interesting similarities to hose of jet-aligued blazars (MeNarou-Brown et al.,1995) pointed out that the spectra of Cen A show interesting similarities to those of jet-aligned blazars (McNaron-Brown et al.
+ 1995). and iu particilar to those of he well-stucied quasar 3€ 273 (Jolhuson et a.," 1995), and in particular to those of the well-studied quasar 3C 273 (Johnson et al."
+ 1995: Lichtiet al., 1995; Lichti et al.
+" 1995): both show spectral weaks in the soft οταν τοσο, aud both have iuteusity-indepenudenut yower-law shaves below the break."," 1995); both show spectral breaks in the soft $\gamma$ -ray regime, and both have intensity-independent power-law shapes below the break."
+ For 3€ 273 the peak in the soft 2-rav regine is a Compton peak. so Steiule ct al. (," For 3C 273 the peak in the soft $\gamma$ -ray regime is a Compton peak, so Steinle et al. ("
+1998) and Winger et al. (,1998) and Kinzer et al. (
+1995) regarded he peak of Cen A at —150 keV as a Compton peak.,1995) regarded the peak of Cen A at $\sim$ 150 keV as a Compton peak.
+ However. based on the observations byROSAT audASCA. Turner et al. (," However, based on the observations by and, Turner et al. ("
+1997) found that the spectrum of he jet of Cen A is actually consistent with the spectra of the nuclear coutimmun. aud that the spectrum of the jet is flatter than that expected as a result of an iuverse-Compton scattering of radio photons but is consistent with hat predicted by a simple svuchrotrou model.,"1997) found that the spectrum of the jet of Cen A is actually consistent with the spectrum of the nuclear continuum, and that the spectrum of the jet is flatter than that expected as a result of an inverse-Compton scattering of radio photons but is consistent with that predicted by a simple synchrotron model."
+ Recent Chandra observations also indicated that X-ravs from he jet of Cen A comprise synchrotron cussion (Birk Lesch 2000)., Recent Chandra observations also indicated that X-rays from the jet of Cen A comprise synchrotron emission (Birk Lesch 2000).
+ Therefore. the peak at ~150 keV is not he Compton peak but the svuchrotron peal. as the case of Ak 501 (see Fig.," Therefore, the peak at $\sim$ 150 keV is not the Compton peak but the synchrotron peak, as the case of Mrk 501 (see Fig."
+ 1)., 1).
+ The peak frequency of Con A is variable. which is also similar to the behavior in TeV dazirs.," The peak frequency of Cen A is variable, which is also similar to the behavior in TeV blazars."
+ Observations byCGRO OSSE in 1991. 1991 revealed peak frequencies of ~150 keV (Iiuzer et al., Observations by OSSE in 1991 $-$ 1994 revealed peak frequencies of $\sim$ 150 keV (Kinzer et al.
+ 1995). while Mivazaki et al. (," 1995), while Miyazaki et al. ("
+1996) reported a peak frequency of 150 keV. based on Wolcomoe-1 observations.,"1996) reported a peak frequency of 180 keV, based on Welcome-1 observations."
+ The peak of the Compton componeut of Mk 501 is at ~l TeV. so that of Cen A is probably also at —1 TeV. During the outburst state. the flux density at 100 keV is =11&107 photon 7s Κον | (see Fie.," The peak of the Compton component of Mrk 501 is at $\sim$ 1 TeV, so that of Cen A is probably also at $\sim$ 1 TeV. During the outburst state, the flux density at 100 keV is $\approx 11\times10^{-5}$ photon $^{-2}$ $^{-1}$ $^{-1}$ (see Fig."
+ 5 in Bound ct al., 5 in Bond et al.
+ 1996). and the svuchrotrou peak flux density is Εθν=2.o10Were 7s fz +. assuming a spectral iudex of 0.7.," 1996), and the synchrotron peak flux density is $f_{150keV} =
+2.1\times10^{-29}$ erg $^{-2}$ $^{-1}$ $^{-1}$, assuming a spectral index of 0.7."
+ According to Eq. (, According to Eq. (
+"CI) with a,=0.7 and a»-—1.3. the estimated TeV dus in 0.25 30 TeV for Con A is which is about LO times brighter than that of Mrk 501 in the 1997 outburst. sugeestingOO that Cen A can be a strong TeV οταν source curving outburst.","4) with $\alpha_1$ =0.7 and $\alpha_2$ =1.3, the estimated TeV flux in 0.25 $-$ 30 TeV for Cen A is which is about 10 times brighter than that of Mrk 501 in the 1997 outburst, suggesting that Cen A can be a strong TeV $\gamma$ -ray source during outburst."
+ The radio to X-ray spectral euergv distribution of M87 is very smooth like that of a typical BL Lac object CIsvetanov ct al., The radio to X-ray spectral energy distribution of M87 is very smooth like that of a typical BL Lac object (Tsvetanov et al.
+ 1998: see Fie., 1998; see Fig.
+ 1)., 1).
+" If it is really a iisalieued BL Lac object. the radio to ταν hup should be the svuchrotron component. and N-ravs are also svuchrotron οσο, sugeestine that Ms? ds an —IBL."," If it is really a misaligned BL Lac object, the radio to X-ray hump should be the synchrotron component, and X-rays are also synchrotron emission, suggesting that M87 is an HBL."
+EINSTEIN audROSAT observations show that the ταν enüssion from the jet of M87 is probably the —[umigh frequency tail of the radio to optical svuclirotrou PAectrun. although other origins (thermal or iuverse-Compton) for the N-ravs cannot be ruled out completely Diretta. Stern Tlarris 1991: Nowunaun et al.," and observations show that the X-ray emission from the jet of M87 is probably the high frequency tail of the radio to optical synchrotron spectrum, although other origins (thermal or inverse-Compton) for the X-rays cannot be ruled out completely (Biretta, Stern Harris 1991; Neumann et al."
+ 1997)., 1997).
+ The EUVE observations of MIST are consistent with the spectral cutoff iu the spectrum of the jet im AIS? as sugeested by Aleisenheiuer et al. (, The $EUVE$ observations of M87 are consistent with the spectral cutoff in the spectrum of the jet in M87 as suggested by Meisenheimer et al. (
+1996). which further supports the idea that the EUV and A-ray cussion of the jet is svuchrotron radiation (Berehotter et al.,"1996), which further supports the idea that the EUV and X-ray emission of the jet is synchrotron radiation (Berghöffer et al."
+ 2000)., 2000).
+ Revuolds et al. (, Reynolds et al. (
+1999) also icdeutify M87. with au IIBL. by combining hard N-ray (4 15keV) observations withROSAT data aud inferring that the N-rav spectra of the AIST core and jet must be steep like those of IBLs.,"1999) also identify M87 with an HBL, by combining hard X-ray (4 $-$ 15keV) observations with data and inferring that the X-ray spectra of the M87 core and jet must be steep like those of HBLs."
+ The svuchrotron peak of MS7 lies between LOY101 TIz (sce Fig., The synchrotron peak of M87 lies between $10^{15} - 10^{16}$ Hz (see Fig.
+ 1)., 1).
+ According to Eq.(1). the Compton compoucut of ALIS? may peak at ~0.1 TeV. The flux of the core aud jet of MIST iu the EUVE DS band pass is 16.25 jiJy. Cuius a core/jet flux ratio of —1.5) with a spectral iudex of the UV to N-vay power-law spectrum of az1.1 (Berehotteret al.," According to Eq.(1), the Compton component of M87 may peak at $\sim$ 0.1 TeV. The flux of the core and jet of M87 in the $EUVE$ DS band pass is 16.25 $\mu$ Jy (using a core/jet flux ratio of $\sim$ 1.5) with a spectral index of the UV to X-ray power-law spectrum of $\alpha \approx 1.4$ (Berghöfferet al."
+ 2000: Πας et al., 2000; Harris et al.
+ 1997)., 1997).
+ According to Eq. C, According to Eq. (
+L). the estimated TeV flux (0.25. 30 TeV) for MST is which is comparable to that of Alek 501 in 1995. sugecsting that M87 may be a TeV +-rav-loud FRI radio ealaxy.,"4), the estimated TeV flux (0.25 $-$ 30 TeV) for M87 is which is comparable to that of Mrk 501 in 1995, suggesting that M87 may be a TeV $\gamma$ -ray-loud FRI radio galaxy."
+Quasars are variable in different wavebands. and their variability was been extensively explored for several decades (Smith&Hof-Heit1963:Matthews&Sandage 1963).,"Quasars are variable in different wavebands, and their variability has been extensively explored for several decades \citep{s1963,m1963}."
+. It has been accepted that quasars are powered by black hole accretion. and their variability jerefore provides an important clue to understand the physics at work in accretion discs.," It has been accepted that quasars are powered by black hole accretion, and their variability therefore provides an important clue to understand the physics at work in accretion discs."
+ The timescales of optical variability for quasars vary from hours to decades., The timescales of optical variability for quasars vary from hours to decades.
+ The variability is found to be correlated with other quasar properties. such as redshift and uminosity.," The variability is found to be correlated with other quasar properties, such as redshift and luminosity."
+" There is a strong negative correlation of variability with quasar luminosity. while a signiticant positive correlation between a""uriability and redshift is also present (Cristianietal.1996:Van-denBerketal. 2004)."," There is a strong negative correlation of variability with quasar luminosity, while a significant positive correlation between variability and redshift is also present \citep{c1996,vb2004}."
+. An anti-correlation between variability amplitude and rest-frame wavelength has also been observed. so the correlation between variability and redshift was explained by the selection effect. because quasars at higher redshifts are more eusier to be probed in shorter rest-frame wavelength (CidFernan-desetal.1996:Cristiani 1996).," An anti-correlation between variability amplitude and rest-frame wavelength has also been observed, so the correlation between variability and redshift was explained by the selection effect, because quasars at higher redshifts are more easier to be probed in shorter rest-frame wavelength \citep{cf1996,c1996}."
+. However. VandenBerketal.2004). claimed that the evolution of quasar population or the change of quasar variability mechanism can cause quasars to be more variable at higher redshifts based on a sample of over 25.000 quasars from the Sloan Digital Sky Survey. which is different from the selection effect scenario.," However, \citet{vb2004} claimed that the evolution of quasar population or the change of quasar variability mechanism can cause quasars to be more variable at higher redshifts based on a sample of over 25,000 quasars from the Sloan Digital Sky Survey, which is different from the selection effect scenario."
+ Recently. Woldetal.(2007) matched quasars from the the Quasar Equatorial Survey Team Phase | (QUESTI) variability survey with broad-line objects from the Sloan Digital Sky Survey. which leads to ~1060 quasars.," Recently, \citet{wold2007} matched quasars from the the Quasar Equatorial Survey Team Phase 1 (QUEST1) variability survey with broad-line objects from the Sloan Digital Sky Survey, which leads to $\sim100$ quasars."
+ The black hole masses of these quasars are estimated from the widths of broad-line Ly., The black hole masses of these quasars are estimated from the widths of broad-line ${\rm H}_\beta$.
+ They measured the variability amplitudes from the QUEST| light curves. and found a correlation between quasar variability amplitude and black hole mass.," They measured the variability amplitudes from the QUEST1 light curves, and found a correlation between quasar variability amplitude and black hole mass."
+ They found that the correlation seems to be unlikely caused by selection effects or the contamination coming from host galaxy or other well-known correlations., They found that the correlation seems to be unlikely caused by selection effects or the contamination coming from host galaxy or other well-known correlations.
+ This correlation between black hole mass and variability amplitude was contirmed by Wilhiteetal.(2007) by using a sample over 2500 quasars., This correlation between black hole mass and variability amplitude was confirmed by \citet{w2007} by using a sample over 2500 quasars.
+ There were many efforts devoted to reveal the physies behind the correlations between quasar optical-UV variability and other quasar properties., There were many efforts devoted to reveal the physics behind the correlations between quasar optical-UV variability and other quasar properties.
+ Many different scenarios were suggested. such as. variation of accretion rates. dise instabilities. stellar collisions. and microlensing effect (e.g..Pereyraetal.2006:Mineshige 2006).," Many different scenarios were suggested, such as, variation of accretion rates, disc instabilities, stellar collisions, and microlensing effect \citep*[e.g.,][]{p2006,m1994,k1998,t1992,t2000,cf2000,m2006}."
+. It is widely believed that quasars are powered by black hole accretion., It is widely believed that quasars are powered by black hole accretion.
+ However. the spectral shape in optical-UV bands for a standard thin accretion dise can be approximated as fi.xvt5 while the observed quasar continua usually exhibit fx1," However, the spectral shape in optical-UV bands for a standard thin accretion disc can be approximated as $f_\nu\propto\nu^{1/3}$, while the observed quasar continua usually exhibit $f_\nu\propto\nu^{-1}$."
+ Gaskell(2007) suggested that the observed quasar spectra can be reproduced by accretion discs with a temperature gradient of TOUR)xRU instead of TOUR)xRO as predicted by standard thin dise model., \citet{g2007} suggested that the observed quasar spectra can be reproduced by accretion discs with a temperature gradient of $T(R)\propto{R^{-0.57}}$ instead of $T(R)\propto{R^{-0.75}}$ as predicted by standard thin disc model.
+ This means that there is some flow of heat outwards in the disk., This means that there is some flow of heat outwards in the disk.
+ They use the optical variability timescales of quasars to argue that the variations must propagate at close to the speed of light. rather than on viscous timescales (Gaskell2007).," They use the optical variability timescales of quasars to argue that the variations must propagate at close to the speed of light, rather than on viscous timescales \citep{g2007}."
+. Pereyraetal.(2006) found that the multi-epoch quasar spectra in the rest-frame wavelength range 192006000 can be fitted by the standard thin dise model. provided their accretion rates vary from," \citet{p2006} found that the multi-epoch quasar spectra in the rest-frame wavelength range $1300-6000 {\rm \AA}$ can be fitted by the standard thin disc model, provided their accretion rates vary from"
+"iin theMS-II, and simulations at z=0.","in the, and simulations at $z=0$."
+" The first two samples comprise the same 12 haloes at different resolutions, while the third contains 1676 haloes."," The first two samples comprise the same 12 haloes at different resolutions, while the third contains 1676 haloes."
+ The mean cumulative mass functions of subhaloes for all three samples are plotted in Fig. 1.., The mean cumulative mass functions of subhaloes for all three samples are plotted in Fig. \ref{fig:conv}.
+" Here, and below, we show cumulative mass functions multiplied by Maur/M20o in order to take out the dominant mass dependence and make the differences between curves more apparent."," Here, and below, we show cumulative mass functions multiplied by $M_{\rm sub}/M_{200}$ in order to take out the dominant mass dependence and make the differences between curves more apparent."
+" Following previous work (e.g.,Stoehretal.2003;Gao2004a;Diemandetal.2007;Springel 2008a),, we restrict the samples to include only subhaloes with 100 particles or more in order to avoid numerical effects."," Following previous work \citep[e.g.,][]{stoehr03, g04a,
+ diemand07, springel08a}, we restrict the samples to include only subhaloes with 100 particles or more in order to avoid numerical effects."
+" There is good convergence between the results from the and the and, within the errors, there is good agreement with the cumulative mass function of theMS."," There is good convergence between the results from the and the and, within the errors, there is good agreement with the cumulative mass function of the."
+ The error bars show the rms scatter about the mean for the MS and MS-II haloes. (, The error bars show the rms scatter about the mean for the MS and MS-II haloes. (
+They are plotted at a selection of well-separated points to reduce the effect of correlations between adjacent bins.),They are plotted at a selection of well-separated points to reduce the effect of correlations between adjacent bins.)
+" This scatter, which reflects the varied formation histories of haloes, is substantial, particularly for large values of Miu/M200."," This scatter, which reflects the varied formation histories of haloes, is substantial, particularly for large values of $M_{\rm sub}/M_{200}$."
+" In a recent study, Boylan-Kolchinetal.(2010) demonstrated that the subhalo occupation distribution is well described by a negative binomial, independently of host halo mass."," In a recent study, \cite{mike} demonstrated that the subhalo occupation distribution is well described by a negative binomial, independently of host halo mass."
+" In this case, the scatter at the high mass end of the subhalo mass function is well described by a Poisson distribution, but at the low mass end, i104, intrinsic scatter of fractional amplitude e~0.15 dominates."," In this case, the scatter at the high mass end of the subhalo mass function is well described by a Poisson distribution, but at the low mass end, $<10^{-4}$, intrinsic scatter of fractional amplitude $\sigma \sim 0.15$ dominates."
+ The authors further found that the cumulative subhalo mass function of galactic size haloes can be well fitted with the following form: The dashed line in Fig., The authors further found that the cumulative subhalo mass function of galactic size haloes can be well fitted with the following form: The dashed line in Fig.
+ 1 is a fit to the mean subhalo cumulative mass function for the cluster size haloes from the using eqn. (1)), \ref{fig:conv} is a fit to the mean subhalo cumulative mass function for the cluster size haloes from the using eqn. \ref{eq:fits}) )
+" with parameters jii=0.01, µειι=0.1, b=1.2 and power-law index a=—0.94."," with parameters $\muone=0.01$, ${\mu_{\rm cut} =0.1}$, $b=1.2$ and power-law index $a=-0.94$."
+" Clearly, the fit is very good."," Clearly, the fit is very good."
+ The best-fit parameters differ slightly from the values given by Boylan-Kolchinetal.(2010) because a different definition of virial radius has been used here and the range of halo masses is different.," The best-fit parameters differ slightly from the values given by \cite{mike}
+ because a different definition of virial radius has been used here and the range of halo masses is different."
+ Gaoetal.(2004a) noted that the abundance of relatively massive subhaloes (Msub/M200> 0.001) increases systematically with host halo mass., \cite{g04a} noted that the abundance of relatively massive subhaloes ${\rm M_{sub}/M_{200}} > 0.001$ ) increases systematically with host halo mass.
+" This trend reflects the fact that in the CDM cosmology more massive haloes are, in the mean, both less centrally concentrated and younger than less massive haloes."," This trend reflects the fact that in the CDM cosmology more massive haloes are, in the mean, both less centrally concentrated and younger than less massive haloes."
+" Thus, they exert weaker tidal forces and have had less time to disrupt their substructure."," Thus, they exert weaker tidal forces and have had less time to disrupt their substructure."
+ This result has been confirmed by subsequent numerical and semi-analytical studies (Zentneretal.2005;vandenBoschetal.2005;Shaw2006;Giocoli 2008b).," This result has been confirmed by subsequent numerical and semi-analytical studies \citep{zentner05,van05,shaw06,giocoli08b}."
+". The much larger samples of well resolved haloes in our set of simulations allows us to explore the trend with host halo mass to much smaller mass ratios than have been resolved in earlier simulations which, for the most part, have focused on cluster size haloes2005)."," The much larger samples of well resolved haloes in our set of simulations allows us to explore the trend with host halo mass to much smaller mass ratios than have been resolved in earlier simulations which, for the most part, have focused on cluster size haloes."
+". Combining data from the Aquarius andMS-II simulations, we have a sample of haloes spanning the range ~10?—3x10'4h~'Mo,, each resolved with at least 10? particles within the virial radius, and often with 10 to 1000 times more."," Combining data from the Aquarius and simulations, we have a sample of haloes spanning the range $\sim 10^{12}- 3 \times 10^{14}$, each resolved with at least $10^5$ particles within the virial radius, and often with 10 to 1000 times more."
+ This represents an improvement by a factor of 100-1000 over the simulations analyzed by Gaoetal., This represents an improvement by a factor of 100-1000 over the simulations analyzed by \cite{g04a}.
+ (200, Fig.
+4a).. Fig. 2 shows cumulative subhalo mass functions for host haloes of different mass., \ref{fig:massdep} shows cumulative subhalo mass functions for host haloes of different mass.
+" The magenta squares and the solid lines respectively represent galactic size haloes ([1,3]x1013 ΤΜο)) from Aquarius andMS-II, the blue lines represent group size haloes ([1,3]x 10'3h~'Mo)) from theMS-II and the black lines represent cluster size haloes ([1,3]x1045! Ma)) from the and theMS."," The magenta squares and the solid lines respectively represent galactic size haloes $[1, 3]
+\times 10^{12}$ ) from Aquarius and, the blue lines represent group size haloes $[1, 3]
+\times 10^{13}$ ) from the and the black lines represent cluster size haloes $[1, 3] \times 10^{14}$ ) from the and the."
+ The corresponding best fits of the form of eqn. (1)), The corresponding best fits of the form of eqn. \ref{eq:fits}) )
+ are overplotted as dotted lines., are overplotted as dotted lines.
+" Clearly, eqn. (1))"," Clearly, eqn. \ref{eq:fits}) )"
+" provides an excellent fit to the cumulative subhalo mass function for all three ranges of host halo mass, albeit with varying best-fit parameters."," provides an excellent fit to the cumulative subhalo mass function for all three ranges of host halo mass, albeit with varying best-fit parameters."
+ We fixed the subhalo function slope index to a=—0.94 when fitting all three curves., We fixed the subhalo function slope index to $a=-0.94$ when fitting all three curves.
+" Of the remaining three free parameters in eqn. (1)),"," Of the remaining three free parameters in eqn. \ref{eq:fits}) ),"
+ 41 determines the overall amplitude of the subhalo abundance function and is largely independent of the parameters µειι and 6 which combine to determine the shape of the cut-off at the high mass end., $\muone$ determines the overall amplitude of the subhalo abundance function and is largely independent of the parameters $\mu_{\rm cut}$ and $b$ which combine to determine the shape of the cut-off at the high mass end.
+ We, We
+sshowed that repeated impacts lead to the growth of a conical structure at the bottom of the target (Fig.,showed that repeated impacts lead to the growth of a conical structure at the bottom of the target (Fig.
+" Ba, "," \ref{fig.VG12_Abbrechen}a a, b)."
+The growth continues until the structure becomes too b).heavy and breaks off from the target surface (Fig., The growth continues until the structure becomes too heavy and breaks off from the target surface (Fig.
+ Bkc)., \ref{fig.VG12_Abbrechen}c c).
+ This happens usually after a few hundred impacts and is often triggered by an impact., This happens usually after a few hundred impacts and is often triggered by an impact.
+" To compare our results with those by ?,, we used the high-speed images to calculate the maximum angles of the cones as illustrated in Fig."," To compare our results with those by \citet{TeiserWurm:2009b}, we used the high-speed images to calculate the maximum angles of the cones as illustrated in Fig."
+" μα. Ordered by increasing velocities, these angles are 75°, 59°, 58°, and 61°."," \ref{fig.VG12_Abbrechen}a a. Ordered by increasing velocities, these angles are $75^\circ$, $59^\circ$, $58^\circ$, and $61^\circ$."
+" It should be noted that due to the breaking off of the cones after a few hundred impacts, these are not necessarily the maximum angles which could have been achieved under microgravity conditions or in top-down collisions as studied by ?.."," It should be noted that due to the breaking off of the cones after a few hundred impacts, these are not necessarily the maximum angles which could have been achieved under microgravity conditions or in top-down collisions as studied by \citet{TeiserWurm:2009a}."
+ The mass determinations after each series of 25 shots showed a steady increase of the accreted mass on the target (Fig. , The mass determinations after each series of 25 shots showed a steady increase of the accreted mass on the target (Fig. \ref{fig.eac_all}) ).
+"After some impacts, this efficiency decreases, probably[4)). owing to the increasing impact angle onto the cone."," After some impacts, this efficiency decreases, probably owing to the increasing impact angle onto the cone."
+" To define the value of the accretion efficiency for impacts on a flat target, or with"," To define the value of the accretion efficiency for impacts on a flat target, or with"
+vorthern (ELODIE) aud southern (CORALIE) datasets as a free parameter.,northern (ELODIE) and southern (CORALIE) datasets as a free parameter.
+ The two systems were found to ve entirely consistent. and this parameter was thus hek fixed to zero for the final solution.," The two systems were found to be entirely consistent, and this parameter was thus held fixed to zero for the final solution."
+ Figure 1 shows the individual radial velocity measurements as a function of orbital phase (the 16 orbital periods elapsed since the firs neasurement make unpractical a display as a function of iue: we however have essentially coutinuous coverage of one period in June and July 1998. excluding any possible spectral alias).," Figure 1 shows the individual radial velocity measurements as a function of orbital phase (the 16 orbital periods elapsed since the first measurement make unpractical a display as a function of time; we however have essentially continuous coverage of one period in June and July 1998, excluding any possible spectral alias)."
+ The orbital period js two mouths anne he velocity senmi-auuplitude is ~250 E. over 10 times he standard error of one radial velocity measurement.," The orbital period is two months and the velocity semi-amplitude is $\sim$ 250 $^{-1}$, over 10 times the standard error of one radial velocity measurement."
+ The radial velocity curve miplies a moderate but highly sienificant eccentricity of 60.31., The radial velocity curve implies a moderate but highly significant eccentricity of $e$ =0.34.
+ The large amplitude aud moderate period of the radial velocity variation argue strouslv for orbital motion as its case., The large amplitude and moderate period of the radial velocity variation argue strongly for orbital motion as its cause.
+ An integration of the radial velocity curve iuples a ΜΗπα. plysical motion of ~0.5R..., An integration of the radial velocity curve implies a minimum physical motion of $\sim$ 0.5.
+ This varlation ds about twice the radius of an ALL dwarf (Baratte Chabrier. 1996: Chabrier Baraffe. 1997). excluding oilsation as a possible explanation.," This variation is about twice the radius of an M4 dwarf (Baraffe Chabrier, 1996; Chabrier Baraffe, 1997), excluding pulsation as a possible explanation."
+" Gl 876 in addition oulv has low level plotometric variability. aud indecd happens to be oue of he standard stars of the original UDV system (Johuson Παπάς, 1951)."," Gl 876 in addition only has low level photometric variability, and indeed happens to be one of the standard stars of the original UBV system (Johnson Harris, 1954)."
+ It is actually variable. but with low rus amplitudes of 13 munae at V. 9 ummae at R and 6 nuuag at I (Weis. 1991).," It is actually variable, but with low rms amplitudes of 13 mmag at V, 9 mmag at R and 6 mmag at I (Weis, 1994)."
+ The photometric variations dou't phase well at 61 days and appear consistent with a BY Dra type variability (Sasselov aud Cody. private communication).," The photometric variations don't phase well at 61 days and appear consistent with a BY Dra type variability (Sasselov and Cody, private communication)."
+ Since CL 8576 js a very slow rotator (Vsiu/<Qhaus +. Delfosse et al.," Since Gl 876 is a very slow rotator $\sin{i}<$ $^{-1}$, Delfosse et al.,"
+ 1998a). rotational modulation of such low level surface inhomogeneities cannot explain its large radial velocity variations.," 1998a), rotational modulation of such low level surface inhomogeneities cannot explain its large radial velocity variations."
+ Deuscly sampled plotometry would ou the other haud be of considerable iuterest to establish the stellar rotation period., Densely sampled photometry would on the other hand be of considerable interest to establish the stellar rotation period.
+ The mass of Cl 876 unfortunately coutributes some uncertaintv to the minium mass of its companion., The mass of Gl 876 unfortunately contributes some uncertainty to the minimum mass of its companion.
+ As a consequence of IT recombination iu the plotosphere aud the deepening couvection for lower mass stars (Ixyoupa ct al., As a consequence of $_2$ recombination in the photosphere and the deepening convection for lower mass stars (Kroupa et al.
+ 1990). the 1unuinositv does not drop nearlv as quickly 201 uut mass or mid-M durs as it does for both higher acl lower mass stars (Henry Me Cartlia. 1993). aud it las a stronecr metallicity depeudeuce.," 1990), the luminosity does not drop nearly as quickly per unit mass for mid-M dwarfs as it does for both higher and lower mass stars (Henry Mc Carthy, 1993), and it has a stronger metallicity dependence."
+ Between 70.50. aud 0.15AL... Mass-Luninuositv relatious therefore have oth shallow slopes aud large intrinsic dispersious (Παν Ale Carthy. 1993).," Between $\sim$ and $\sim$, Mass-Luminosity relations therefore have both shallow slopes and large intrinsic dispersions (Henry Mc Carthy, 1993)."
+ Cd 876 thus belongs to a spectral ype range where the mass of a sinele star is poorly constrained by its observable characteristics., Gl 876 thus belongs to a spectral type range where the mass of a single star is poorly constrained by its observable characteristics.
+ Takius at ace value either the solar neigliborlioo observational nass-Dunuinositv relation of Teury Me Carthy (1993) or the solar metallicity mocels of Baraffe et al. (, Taking at face value either the solar neighborhood observational mass-luminosity relation of Henry Mc Carthy (1993) or the solar metallicity models of Baraffe et al. (
+"1998). he absolute magnitudes of GI 876 (ατα σεν AL) =7.56. AT, =6.96. Mj =6.70. Legeett (1992) and ESA (1997)) Πρίν a aiuass of 0.304:0.05 for Gl στο,","1998), the absolute magnitudes of Gl 876 $_V$ =11.81, $_J$ =7.56, $_H$ =6.96, $_K$ =6.70, Leggett (1992) and ESA (1997)) imply a mass of $\pm$ for Gl 876."
+ As au illustration of possible uncertainties however. Delfosse et al (199860) measure M = 0.132. 40.001 and Al = 11.740.2 (AL =11.81 for Cd 876) for the brighter star in CJ 2069A. an ΑΟΡΝ eclipsing binary which is probably superauetabrich (MJII|A 10.5).," As an illustration of possible uncertainties however, Delfosse et al (1998c) measure M = 0.432 $\pm$ and $_V$ = $\pm$ 0.2 $_V$ =11.81 for Gl 876) for the brighter star in GJ 2069A, an M3.5V eclipsing binary which is probably super-metal-rich $\sim$ +0.5)."
+ From its vosition in colour-colour diagrams (Leeectt.ao 1992). alca youn the relative depth of its cross-correlation dips with several binary templates (Delfosse et al..," From its position in colour-colour diagrams (Leggett, 1992), and from the relative depth of its cross-correlation dips with several binary templates (Delfosse et al.,"
+ in preparation). C] στο is probably more metallic than the suu. thoug[um rot as michi as CJ 2069À. We adopt a imass ofAL... for he rest of the discussion but warn that it is uncertain * perhaps30:4.," in preparation), Gl 876 is probably more metallic than the sun, though not as much as GJ 2069A. We adopt a mass of for the rest of the discussion but warn that it is uncertain by perhaps."
+. The ΓΕΕΕ semi-major axis and anetarv mass which result from the orbital solutio- aro then asini = 0.20 AU and Mosin; = 2.0M., The minimum semi-major axis and planetary mass which result from the orbital solution are then $\sin{i}$ = 0.20 AU and $_{2}\sin{i}$ = $_J$.
+ Thev respectivelyH scale as M51 aud M7.2, They respectively scale as $^{\frac{1}{3}}$ and $^{\frac{2}{3}}$.
+ The orbital eleineuts of Cdiese 876b axe worth notius., The orbital elements of Gliese 876b are worth noting.
+" Iu spite of the low mass of C1 876. the ice-coucdensation radius at the time of planet formation around this star is ~ 1 AU, only lower than around a solar type star (Boss 1995)."," In spite of the low mass of Gl 876, the ice-condensation radius at the time of planet formation around this star is $\sim$ 4 AU, only lower than around a solar type star (Boss 1995)."
+ Once again the orbital seiuinajor axis (a = 0.20 AU) is thus muuch smaller than the expected ΙΙΙ racius for eiur planet formation. and some orbital miüeration must have occeured.," Once again the orbital semi-major axis (a = 0.20 AU) is thus much smaller than the expected minimum radius for giant planet formation, and some orbital migration must have occcured."
+ However the observed orbital separation is also l| times larecr than the measured semi-major axes of 51 Pee. τ Boo and (e Aud (0.01-0.05 AU).," However the observed orbital separation is also 4 times larger than the measured semi-major axes of 51 Peg, $\tau$ Boo and $\upsilon$ And (0.04-0.05 AU)."
+ The excess of planets with such small semi-major axes is bolieved to result from outward torques which counteract at shor distances the imward torque induced by the imteraction of the planct and the protoplanctary disk (Lin ct al., The excess of planets with such small semi-major axes is believed to result from outward torques which counteract at short distances the inward torque induced by the interaction of the planet and the protoplanetary disk (Lin et al.
+ 1996. Trilline et al.," 1996, Trilling et al."
+ 1998)., 1998).
+ These torques only become effective at separations significantly smaller than 0.20 AU. cannot have plaved a siguficaut role for the Gl 876 system," These torques only become effective at separations significantly smaller than 0.20 AU, and cannot have played a signficant role for the Gl 876 system."
+ I is also interesting to note iat the orbit of GL S76) Is eccentric (¢=0.35). while interaction with an accretion disk 1s expectec to damp any sienificaut orbital eccentricity of a planet (Goldreich and Tremaine. 1980).," It is also interesting to note that the orbit of GL 876b is eccentric $e$ =0.35), while interaction with an accretion disk is expected to damp any significant orbital eccentricity of a planet (Goldreich and Tremaine, 1980)."
+ Several mechanisius nav be called in to explain, Several mechanisms may be called in to explain
+"The study of disks of dust and eas surroundiug voue stabl xovides insight inte tιο environment in which lanets are likely to ]o formu19,",The study of disks of dust and gas surrounding young stars provides insight into the environment in which planets are likely to be forming.
+ Moleenlar abindances in the diskcan be used f) understand the chemical evolution of the gas. (lisk pxcesses stich as radial and vertical mixing. and f1ο effeοs of stellar irradiation on he disk chemistry (Beremi e al.," Molecular abundances in the disk can be used to understand the chemical evolution of the gas, disk processes such as radial and vertical mixing, and the effects of stellar irradiation on the disk chemistry (Bergin et al."
+ 2007)., 2007).
+ While models xediet disk chemistry in tliο earth-like planet forming ‘eeions (ne. MarkwiekX ct al., While models predict disk chemistry in the earth-like planet forming regions (i.e. Markwick et al.
+ 2002: AÁennnudez ct al., 2002; Agúnndez et al.
+ 2008: Classeokl et al., 2008; Glassgold et al.
+ 2009: Woods Willacy 2009). supporting observatiois of the chemical tracers have ven scarce to date.," 2009; Woods Willacy 2009), supporting observations of the chemical tracers have been scarce to date."
+ The Spitzer Infrared Spectrograph (RS) made observations that probed the chenustry ofthe inner disk. detecting 1iolecular absorion (Lahuis ct al.," The Spitzer Infrared Spectrograph (IRS) made observations that probed the chemistry of the inner disk, detecting molecular absorption (Lahuis et al."
+ 2006) and enission (Carr Najita 20tδι SalwkX ct , 2006) and emission (Carr Najita 2008; Salyk et al.
+2008)., 2008).
+ The resolution of IRS was ποκited for detecting molecules with large equivalent widtli frou ro-vibrational Q-branches., The resolution of IRS was well-suited for detecting molecules with large equivalent widths from ro-vibrational Q-branches.
+ Lahmis et al. (, Lahuis et al. (
+2006) detected warn. Colb. IICN. and CO» in absoption toward IRS 16. likely originating m the imucr xEous of a ciretuustellar disk seen near edec-on.,"2006) detected warm $_{2}$ $_{2}$, HCN, and $_{2}$ in absorption toward IRS 46, likely originating in the inner regions of a circumstellar disk seen near edge-on."
+" Untotuately. svstenas displaying, molecular absorption are ALC with. for example. IRS 16 being the ouly one of roughv 100 voung stellar objects (YSO) observed in the eX Spitzer Legacy program (Evans ct al."," Unfortunately, systems displaying molecular absorption are rare with, for example, IRS 46 being the only one of roughly 100 young stellar objects (YSO) observed in the c2d Spitzer Legacy program (Evans et al."
+ 2003)., 2003).
+ Molecular CoILIo and HCN absorption have alsobeen detected toward CY Tan N (Doppnuizun ct al., Molecular $_{2}$ $_{2}$ and HCN absorption have also been detected toward GV Tau N (Doppmann et al.
+ 2008: cf., 2008; cf.
+ Gibb et al., Gibb et al.
+ 2007). a classical -- eonipeuuion IRC: Ixoresko et al.," 2007), a classical infrared companion (IRC; Koresko et al."
+ 1997)., 1997).
+ IRCs are binary companions to voune sars that are faint im the visible but dominate he system flux in the infrared., IRCs are binary companions to young stars that are faint in the visible but dominate the system flux in the infrared.
+ Correia et al. (, Correia et al. (
+2007) ‘ound with MIDIVLTI data that three IRCs. includ GV Tan. are likely To Tani starsbi (TTS) with tf (isk seen aluost edge-on.,"2007) found with MIDI/VLTI data that three IRCs, including GV Tau, are likely T Tauri stars (TTS) with the disk seen almost edge-on."
+ Doppiann ct al. (, Doppmann et al. (
+200 ‘ound the wari temperature and the relative veloci shifts of the molecalar absorption further supported t edge-on disk model.,2008) found the warm temperature and the relative velocity shifts of the molecular absorption further supported the edge-on disk model.
+ Cibb et al. (, Gibb et al. (
+2007) compared t nolecular absorption abundances to predictions made he chemical αἱsk nodel of Markwick et al. (,2007) compared the molecular absorption abundances to predictions made by the chemical disk model of Markwick et al. (
+2002) and ound them to IC consistent with the imucr region of a xotoplanetary disk.,2002) and found them to be consistent with the inner region of a protoplanetary disk.
+ Because IRS 16 and CV TanN show the rare features hat are valuable for probing the inner disk chemistry. it is advantageous to study ore disks with simular attributes.," Because IRS 46 and GV Tau N show the rare features that are valuable for probing the inner disk chemistry, it is advantageous to study more disks with similar attributes."
+ We were granted time with Spitzer IRS to ostain hieh signal-to-noise (8.N) spectra of seven YSOs., We were granted time with Spitzer IRS to obtain high signal-to-noise (S/N) spectra of seven YSOs.
+ Our target selection was limited to systems thought to be simiku to IRS 16 and GV Tau N with a laree cohuun of eas between the observer and the cutting source necessary for absorption studies: three circmustellar disks seen near edgc-on and four classical IRCs., Our target selection was limited to systems thought to be similar to IRS 46 and GV Tau N with a large column of gas between the observer and the emitting source necessary for absorption studies: three circumstellar disks seen near edge-on and four classical IRCs.
+ DC Tau Band VV CrA are two of the objects observed in tliis project., DG Tau B and VV CrA are two of the objects observed in this project.
+ Abundt Fried (1983) first designated DG Tau D as the, Mundt Fried (1983) first designated DG Tau B as the
+In addition to the sparsity condition. the other important ingredient of our technique relies on the election of the sensing matrix.,"In addition to the sparsity condition, the other important ingredient of our technique relies on the election of the sensing matrix."
+ This matrix is the one that relates the underlving spectrum with (he measurements we use in (he reconstruction., This matrix is the one that relates the underlying spectrum with the measurements we use in the reconstruction.
+ Our ain is to test whether photometric data can be used to reconstruct spectra. so that (he sensing matrix is nol a choice. but. given by the weighting functions of the SDSS filler set.," Our aim is to test whether photometric data can be used to reconstruct spectra, so that the sensing matrix is not a choice, but given by the weighting functions of the SDSS filter set."
+ Figure 1. shows the ugriz liller sel and an example of an observed spectrum., Figure \ref{fig:sensing_matrix} shows the $ugriz$ filter set and an example of an observed spectrum.
+ Note that fillers 4 and 2 have important contributions outside the observed spectral range., Note that filters $u$ and $z$ have important contributions outside the observed spectral range.
+ Consequently. the information they contain cannot be easily utilized under the scheme presented in (lis paper ancl we carry out the reconstructions using only fillers g. kr ancl j.," Consequently, the information they contain cannot be easily utilized under the scheme presented in this paper and we carry out the reconstructions using only filters $g$, $r$ and $i$ ."
+ Obtaining the magnitude in the filler 4: of the SDSS svstem [rom spectroscopic data reduces to the calculation of the following quantity (Fukugitaοἱal.1996): where (ris the (lus distribution of the star., Obtaining the magnitude in the filter $k$ of the SDSS system from spectroscopic data reduces to the calculation of the following quantity \citep{fukugita96}: where $f(\nu)$ is the flux distribution of the star.
+ Therefore. S(v) is the effective (ansmissivily of the filter. including the filler response. the CCD quantun efficiency and the typical (vansmission of the sky [or a point source (in our cases adopted for an airmass of 1.35)).," Therefore, $S(\nu)$ is the effective transmissivity of the filter, including the filter response, the CCD quantum efficiency and the typical transmission of the sky for a point source (in our cases adopted for an airmass of )."
+ The flux measured [or each filler can be estimated using the Riemann integral as (note (hat more precise quadrature rules can be used without a significant change in the following discussion): where we have made (he integration in the wavelength axis and usecl /(A) instead of {εν} since the principal components of MeGurketal.(2010) are given in terms of f(A)., The flux measured for each filter can be estimated using the Riemann integral as (note that more precise quadrature rules can be used without a significant change in the following discussion): where we have made the integration in the wavelength axis and used $f(\lambda)$ instead of $f(\nu)$ since the principal components of \cite{mcgurk10} are given in terms of $f(\lambda)$.
+ The normalization constant for each filter is obtained likewise: Note that one is able toisolate the [αν from the observed magnitudes as:, The normalization constant for each filter is obtained likewise: Note that one is able toisolate the flux from the observed magnitudes as:
+Iu this work. we propose that black hole selfreeulation nay be mediated bv 1-100 GeV protons. or iu other wordsrugs. produced iu the jet core duriug xhases of radio-loud quasar activity.,"In this work, we propose that black hole self-regulation may be mediated by 1-100 GeV protons, or in other words, produced in the jet core during phases of radio-loud quasar activity."
+ We show that enough cosnüc ravs escape the radio core iuto the host ealaxy to power a cosmic rav-driven wind that ejects the interstellar gas. thus removing the fuel for further star ormation and black hole accretion.," We show that enough cosmic rays escape the radio core into the host galaxy to power a cosmic ray-driven wind that ejects the interstellar gas, thus removing the fuel for further star formation and black hole accretion."
+ Our arguinuenuts are organized as follows., Our arguments are organized as follows.
+ In refs:basic woe lav out the basic idea behind our jet-powered cosmde ray feedback scenario by define and recognizing the miportauce of the Eddington Πιτ I cosnic ravs., In \\ref{s: basic} we lay out the basic idea behind our jet-powered cosmic ray feedback scenario by defining and recognizing the importance of the Eddington limit in cosmic rays.
+ refs:model wemsbriefly summarize the plysical features of the standard model of radio-loud ACN NSjets that are Huportant for our Feedback mechaisin., In \\ref{s: model} we briefly summarize the physical features of the standard model of radio-loud AGN jets that are important for our feedback mechanism.
+ ith the jet modelin haud. iu vets: effiainercedeter rininetheinterste Harcosinicrayliin itypreténsyyg Toudquasar acticityinterinsofobsercablequantitics.," With the jet model in hand, in \\ref{s: eff_lum} we determine the interstellar cosmic ray luminosity resulting from radio-loud quasar activity in terms of observable quantities."
+Forreadesaporubsiestadioet ios icsthastdridtoonre aswellasit: :anodelands fram," For readers uninterested in the details of our jet and radio core model, as well as its phenomenological underpinnings, we suggest skipping \\ref{s: model} and \\ref{s: eff_lum}."
+ Dns fecdbaekweeveamiunctheconscquencesofourcosimicreasgdapllbeesRsc dst tdt TO ack hol regulationmodelprescntedherewithothermodelsofquasar fccdbhagdo, In \\ref{sec:feedback} we examine the consequences of our cosmic ray feedback scenario and make a critical comparison of the black hole self-regulation model presented here with other models of quasar feedback.
+dlemitimmosminirqusdlingsint test... wherewealsopresentabricfrecicwo fwcll knownobservationsthatlendsupporttoour pet∖CHR poweredsclf regulationmnecchanisin," We summarize our findings in \\ref{s: test}, where we also present a brief review of well-known observations that lend support to our jet-powered self-regulation mechanism."
+" Inaddition. ο... listinital i5,"," In addition, we propose an observational test."
+ If supermassive black hole growth results primarily from racdiatively-cficient accretion. as the arguneut sugecsts.yooputu then the ratio of the total euergv released during the accretion process to the binding encrev of the ealaxv« gaseous phase is given by in case the black hole ealaxy system follows the AM.AL. relation in eq. (03).," If supermassive black hole growth results primarily from radiatively-efficient accretion, as the argument suggests, then the ratio of the total energy released during the accretion process to the binding energy of the galaxy's gaseous phase is given by in case the black hole – galaxy system follows the $M_\bullet-M_\star$ relation in eq. \ref{e: Magg}) )."
+" Tere. 6,=Oley is the radiative efficiency accro flow."," Here, $\epsilon_\ditto{rad}=0.1\,\epsilon_\ditto{rad,\min}$ is the radiative efficiency of the accretion flow."
+" The biudius energy of the eas LByisothermalthe f,M.a?js fionappropriate for a galaxy οποιο au sphere whose eravitational force is mainly provided by stars aud dark mnatter: f,OLfa is the galaxw’s gas fraction aud σι=3000,39)kins+ is the stellar velocity dispersion."," The binding energy of the gas $E_g\simeq f_gM_{\star} \sigma^2_{\star}$is appropriate for a galaxy resembling an isothermal sphere whose gravitational force is mainly provided by stars and dark matter; $f_g=0.1\,f_{g,\min}$ is the galaxy's gas fraction and $\sigma_\star=300\,\sigma_{\star,300} \unit{km\,s^{-1}}$ is the stellar velocity dispersion."
+ Clearly. only a small amount of the photon energv released during optically-bright quasar epochs is needed o couple to the galactic gas iu order to eject it from he galaxv gravitational potential. thus self-Imitiug the colmbined growths of theblack hole aud the stellar bulec.," Clearly, only a small amount of the photon energy released during optically-bright quasar epochs is needed to couple to the galactic gas in order to eject it from the galaxy's gravitational potential, thus self-limiting the combined growth of the black hole and the stellar bulge."
+ Again. duriug radio-Ioud phases a significant amount of he eucrey resulting from black hole accretion is released in mechanical form as powerful relativistic jets.," Again, during radio-loud phases a significant amount of the energy resulting from black hole accretion is released in mechanical form as powerful relativistic jets."
+" Even if he time-integrated kinetic output of the jet is AF, ΔΙ. eq. (3)} "," Even if the time-integrated kinetic output of the jet is $\Delta E_\ditto{J} \ll 
+\Delta E_{\bullet}$ , eq. \ref{e: DeltaE}) )"
+indicates that there may be ample euerev o unbind the galactic gas. provided that the fraction of jet kinetic energwv deposited in the radio core can efficieutly. couple to the gaseous component of the host ealaxy.," indicates that there may be ample energy to unbind the galactic gas, provided that the fraction of jet kinetic energy deposited in the radio core can efficiently couple to the gaseous component of the host galaxy."
+ Iu other words. radio-loud kincetically-domunated epochs may represent another mode of ACN activity. other than optically-lhinuneus radiativelv-efficieut quasar phases. that could potentially lead to efficient black hole selt-xeeulation.," In other words, radio-loud kinetically-dominated epochs may represent another mode of AGN activity, other than optically-luminous radiatively-efficient quasar phases, that could potentially lead to efficient black hole self-regulation."
+ The photou lhuuinositv of the sub-pc scale jet core results from the cooling of relativistic electrons through a combination of svuchnrotron aud inverse Compton cluission., The photon luminosity of the sub-pc scale jet core results from the cooling of relativistic electrons through a combination of synchrotron and inverse Compton emission.
+ By modeling the spectral energy.distribution of powerful radio-loud quasars. infer that the cluitting electrons iust be accelerated. presumably by sole magnetized Fermi mechanism. up to liehh-relativistic energies. iu excess of a few teus of GeV (in the jet comoving frame).," By modeling the spectral energy distribution of powerful radio-loud quasars, infer that the emitting electrons must be accelerated, presumably by some magnetized Fermi mechanism, up to highly-relativistic energies, in excess of a few tens of GeV (in the jet comoving frame)."
+ The details of the acceleration mechanisin aside. it is reasonable to couclude as well that protons are raudomized with cucreics upto 1.10Gey.," The details of the acceleration mechanism aside, it is reasonable to conclude as well that protons are randomized with energies up to $\sim1-10\unit{GeV}$."
+" If a significant fraction of these randomized relativistic ions. or in other wordsrays. escape the site of acceleration aud diffuse iuto the iuterstellar medina of the host galaxy. they may provide he coupling between jet power and galactic gas required for black hole selt-regulation duriug radio-loud ράσο»,"," If a significant fraction of these randomized relativistic ions, or in other words, escape the site of acceleration and diffuse into the interstellar medium of the host galaxy, they may provide the coupling between jet power and galactic gas required for black hole self-regulation during radio-loud phases."
+ For example. iu actively star-forming ealaxics the eeneration and subsequent diffusion of COSLUC Ταν freanwalbbnown o by-product of core-collapse ΕΙΡ niadak act as an ageut of sclfreeulation. luting the rate of star-formation aud therefore the galaxy huninositv.," For example, in actively star-forming galaxies the generation and subsequent diffusion of cosmic ray protons – a well-known by-product of core-collapse supernovae, and thus massive star-formation – may act as an agent of self-regulation, limiting the rate of star-formation and therefore the galaxy luminosity."
+" The tornawgelesstaac jet alghanse where £,4,, is the Thomson Eddington lauit for a ealaxy with mass M,: we have pinued the ratio between the cosmic rav mean free path and the Thomson meau free path to the value A./A,~105. as in the case of the Milkv Way."," The simplest way to understand this is by defining the where $L_{_{\rm Edd, \star}}$ is the Thomson Eddington limit for a galaxy with mass $M_\star$; we have pinned the ratio between the cosmic ray mean free path and the Thomson mean free path to the value $\lambda_\ditto{CR}/\lambda_\ditto{T} \sim
+10^{-6}$, as in the case of the Milky Way."
+" Au interstellar cosmic rav. Iuninosity of LCh~Jüleygs| equal to the corresponding cosunic rav Eddinetou lait 1,Toddht for the most massive ealaxies may result from the act of forming stars atf a rate of ~105ALvr| corresponding to the brightest star-forming galaxies."," An interstellar cosmic ray luminosity of $L_\ditto{CR}\sim 10^{44}\,{\rm
+erg\,s^{-1}}$ – equal to the corresponding cosmic ray Eddington limit $L_{_{\rm Edd,
+CR}}$ for the most massive galaxies – may result from the act of forming stars at a rate of $\sim 10^3\,{M}_{\odot}\,{\rm yr^{-1}}$ – corresponding to the brightest star-forming galaxies."
+ The formation of stars at a ligher rate leads to the breaking of the cosmic rav Eddington hmit aud a cosmic rav-divenu wind develops. removing the gaseous phase of the galaxy thus choline off star-formation.," The formation of stars at a higher rate leads to the breaking of the cosmic ray Eddington limit and a cosmic ray-driven wind develops, removing the gaseous phase of the galaxy thus choking off star-formation."
+ Frou this straightforward aremueut. it is quite possible that the Iuninosities of star-forming ealaxies are capped at their observed values because they are Eddineton-limited iu cosmic rays. as proposed byο," From this straightforward argument, it is quite possible that the luminosities of star-forming galaxies are capped at their observed values because they are Eddington-limited in cosmic rays, as proposed by."
+"ν With respect to ACN. radio-loud objects display powerful jets with kinetic huuinositv upwards of L,~10Mores3 2). comparable to the Thomson Eddington lait £,,,,, for the most massive black holes with Af,~10?AL..."," With respect to AGN, radio-loud objects display powerful jets with kinetic luminosity upwards of $L_\ditto{J}\sim 10^{47}\, {\rm erg\,s^{-1}}$ , comparable to the Thomson Eddington limit $L_\ditto{Edd,\bullet}$ for the most massive black holes with $M_{\bullet} \sim 10^9\, M_{\odot}$."
+" Tfthe jot kinetic power is a fraction Av,~ Lof the black hole Thomson Eddington linüt. we have where Af,=LOYAL AL..."," If the jet kinetic power is a fraction $\Lambda_{_{
+\rm Edd}}\sim1$ of the black hole Thomson Eddington limit, we have where $M_{\bullet}=10^9M_{\bullet,9}\,M_{\sun}$ ."
+" Let ἐωμ be the efficieney for converting the jet kinetic power £, iutointerstellar cosmic rav power ", Let $\epsilon_\ditto{CR}$ be the efficiency for converting the jet kinetic power $L_\ditto{J}$ intointerstellar cosmic ray power $L_\ditto{CR}$ .
+"A statement of moment conservation and livcrostatic£,.,. balance. the Eddinegtou lit m cosnüc ravs defined in eq. (1))"," A statement of momentum conservation and hydrostatic balance, the Eddington limit in cosmic rays defined in eq. \ref{e: LeddCR}) )"
+ mdicates that a ealactic cosmic rav-driiven wind develops if, indicates that a galactic cosmic ray-driven wind develops if
+"entire $,—a area observed. leaving the impression that flux and spectral index are uncorrelated.","entire $S_{\nu}-\alpha$ area observed, leaving the impression that flux and spectral index are uncorrelated."
+ The 3-mm (1.3-mm) power spectrum globally follows a red noise powerlaw with index B«0.7 (B«0.5)., The 3-mm (1.3-mm) power spectrum globally follows a red noise powerlaw with index $\beta\approx0.7$ $\beta\approx0.5$ ).
+ However. for frequencies f<5x107 days! the power spectrum seem to be an upscaled — by a factor «3 (x2) — version of the average powerlaw.," However, for frequencies $f\lesssim5\times10^{-3}$ $^{-1}$ the power spectrum seem to be an upscaled – by a factor $\approx$ 3 $\approx$ 2) – version of the average powerlaw."
+ The discrete correlation function. points to time delays τ significantly different from zero in terms of the 3c confidence threshold. located in the range tr=—20..—110 days.," The discrete correlation function points to time delays $\tau$ significantly different from zero in terms of the $3\sigma$ confidence threshold, located in the range $\tau \approx -20 ... -110$ days."
+ The negative delay indicates that the. 1.3-mm emission. preceeds the 3-mm emission., The negative delay indicates that the 1.3-mm emission preceeds the 3-mm emission.
+" However. as 3C 454.3 shows a dramatic transition from ""quiescent"" to “flaring” around epoch 2005.2. the curve presented in Fig."," However, as 3C 454.3 shows a dramatic transition from “quiescent” to “flaring” around epoch 2005.2, the curve presented in Fig."
+ 8 (bottom left) might be misleading as the timing of the emission might be different in the two states., \ref{fig_timelag} (bottom left) might be misleading as the timing of the emission might be different in the two states.
+ We therefore computed in addition one DCF for each of the two activity phases. one covering all data from 1996.8 to 2005.1. one including all measurements obtained between 2005.2 and 2010.6.," We therefore computed in addition one DCF for each of the two activity phases, one covering all data from 1996.8 to 2005.1, one including all measurements obtained between 2005.2 and 2010.6."
+ The results are summarized in Fig. 9.., The results are summarized in Fig. \ref{fig_dcf3C454.3}.
+ We note that for the “quiescent” phase (before 2005.2) the most probable time delay — on a low absolute level. DCF(0)«0.6 — is consistent with zero. whereas we find for the “flaring” phase (after 2005.2) 7=—15...—80 days. with the maximum of the DCF located at r=—55 days.," We note that for the “quiescent” phase (before 2005.2) the most probable time delay – on a low absolute level, $DCF(0)\approx0.6$ – is consistent with zero, whereas we find for the “flaring” phase (after 2005.2) $\tau\approx -15 ... -80$ days, with the maximum of the DCF located at $\tau\approx-55$ days."
+ Again. the negative delay indicates that the. 1.3-mm emission preceeds the 3-mm emission.," Again, the negative delay indicates that the 1.3-mm emission preceeds the 3-mm emission."
+ Finally. we probed the impact of the spectral time delay on the calculation of the spectral index a.," Finally, we probed the impact of the spectral time delay on the calculation of the spectral index $\alpha$."
+ For this. we repeated the calculation outlined in Sect.," For this, we repeated the calculation outlined in Sect."
+ 3.3 using 1.3-mm and 3-mm flux density data only., 3.3 using 1.3-mm and 3-mm flux density data only.
+ We compensated the spectral time delay by artificially retarding the 1.3-mm lighteurve by 55 days., We compensated the spectral time delay by artificially retarding the 1.3-mm lightcurve by 55 days.
+ We present the resulting modified spectral index evolution in Fig., We present the resulting modified spectral index evolution in Fig.
+ 11 (left panel)., \ref{fig_specindex_timeshift} (left panel).
+ Time bins are the same as in Fig., Time bins are the same as in Fig.
+ + (bottom left panel) which should be compared to Fig., \ref{fig_specindex} (bottom left panel) which should be compared to Fig.
+ 11. (left panel)., \ref{fig_specindex_timeshift} (left panel).
+ We note that retarding the 1.3-mm data removes the inverted values for « that are visible in Fig., We note that retarding the 1.3-mm data removes the inverted values for $\alpha$ that are visible in Fig.
+ 4. (bottom left panel) during the outburst 2005/2006., \ref{fig_specindex} (bottom left panel) during the outburst 2005/2006.
+ The 3-mm (1.3-mm) lightcurve of this Seyfert 2 galaxy varies smoothly during the entire monitoring time between 23.5 Jy (x2 Jy) in 2002 and «11 Jy (=7 Jy) in 2010., The 3-mm (1.3-mm) lightcurve of this Seyfert 2 galaxy varies smoothly during the entire monitoring time between $\approx$ 3.5 Jy $\approx$ 2 Jy) in 2002 and $\approx$ 11 Jy $\approx$ 7 Jy) in 2010.
+ Since about 2004 the fluxes raise continuously with exponential rise time scales of x6 years (x4 years)., Since about 2004 the fluxes raise continuously with exponential rise time scales of $\approx$ 6 years $\approx$ 4 years).
+ The flux distributions show somewhat complicated double-peak profiles that can be approximated às superpositions of three Gaussians peaking at z4 Jy. 25.5 Jy. and x9 Jy (=2 Jy. =4 Jy. and =6 Jy).," The flux distributions show somewhat complicated double-peak profiles that can be approximated as superpositions of three Gaussians peaking at $\approx$ 4 Jy, $\approx$ 5.5 Jy, and $\approx$ 9 Jy $\approx$ 2 Jy, $\approx$ 4 Jy, and $\approx$ 6 Jy)."
+ The spectral index is steep. with a=0.9 (within errors) until about 2005 and then drops smoothly ίσα=0.5 after 2007.," The spectral index is steep, with $\alpha\approx0.9$ (within errors) until about 2005 and then drops smoothly to $\alpha\approx0.5$ after 2007."
+ We find the flattest spectral index for the highest flux value: however. as this is based on a single data point that occupies a large range in flux. this observation does not establish a significant trend.," We find the flattest spectral index for the highest flux value; however, as this is based on a single data point that occupies a large range in flux, this observation does not establish a significant trend."
+ The power spectra can be described as simple red-noise spectra. with powerlaw indices 8=0.7 for the 3-mm and Ps0.4 for the 1.3-mm lighteurve. respectively.," The power spectra can be described as simple red-noise spectra, with powerlaw indices $\beta\approx0.7$ for the 3-mm and $\beta\approx0.4$ for the 1.3-mm lightcurve, respectively."
+ The discrete correlation function computed from. the 1.3-mm and 3-mm lighteurves finds a spectral time delay significantly different from zero., The discrete correlation function computed from the 1.3-mm and 3-mm lightcurves finds a spectral time delay significantly different from zero.
+ The most probable time delay is located at 7.x—270...—330days. meaning that the 1.3-mm emission preceeds the 3-mm emission by almost one year.," The most probable time delay is located at $\tau\approx -270 ... -330$days, meaning that the 1.3-mm emission preceeds the 3-mm emission by almost one year."
+ We note the high absolute value of the maximum correlation. DCF(r=-277d)1.04+0.05. which is in agreement with the theoretical maximum value of 1.," We note the high absolute value of the maximum correlation, $DCF(\tau=-277{\rm d})=1.04\pm0.05$, which is in agreement with the theoretical maximum value of 1."
+ However. as already DCF(O)«0.7. one might wonder how large the difference between τ30 and r«—300 actually is 1n terms.," However, as already $DCF(0)\approx0.7$, one might wonder how large the difference between $\tau\approx0$ and $\tau\approx-300$ actually is in terms."
+ We therefore tested the impact of the time offset graphically., We therefore tested the impact of the time offset graphically.
+ For this we normalize the 1.3-mm and 3-mm lighteurves to zero mean and unity standard deviation. in analogy to the computation of the unbinned discrete correlations according to Eq. 7..," For this we normalize the 1.3-mm and 3-mm lightcurves to zero mean and unity standard deviation, in analogy to the computation of the unbinned discrete correlations according to Eq. \ref{eq_udcf}."
+ In Fig., In Fig.
+" 10 we accordingly present the quantity (S)«o for each of the two frequency bands: here S, is the flux density. (S,) denotes the time average of the flux density. and σ ndicates the standard deviation of the flux density."," \ref{fig_3C84shift} we accordingly present the quantity $(S_{\nu}-\langle S_{\nu}\rangle)/\sigma$ for each of the two frequency bands; here $S_{\nu}$ is the flux density, $\langle S_{\nu}\rangle$ denotes the time average of the flux density, and $\sigma$ indicates the standard deviation of the flux density."
+ In Fig., In Fig.
+ 10 we present the normalized data separately for the two cases r=0 days (left panel) and r=—300 days days (right panel)., \ref{fig_3C84shift} we present the normalized data separately for the two cases $\tau=0$ days (left panel) and $\tau=-300$ days days (right panel).
+ Indeed the case r=-300 days shows a better agreement between the 1.3-mm and 3-mm lightcurves than the case r=0 days., Indeed the case $\tau=-300$ days shows a better agreement between the 1.3-mm and 3-mm lightcurves than the case $\tau=0$ days.
+ This agrees with the result obtained from the correlation analysis., This agrees with the result obtained from the correlation analysis.
+ However. we note that the lightcurves show clear trends that are only partially covered by our observations.," However, we note that the lightcurves show clear trends that are only partially covered by our observations."
+ This can mislead a correlation analysis. meaning that the time delay we find has to be taken with care.," This can mislead a correlation analysis, meaning that the time delay we find has to be taken with care."
+ Like for 3C 454.3. we probed the impact of the spectral time delay on the calculation of the spectral index.," Like for 3C 454.3, we probed the impact of the spectral time delay on the calculation of the spectral index."
+ We compensated the spectral time delay by artificially retarding the 1.3-mm lightcurve by 300 days., We compensated the spectral time delay by artificially retarding the 1.3-mm lightcurve by 300 days.
+ We present the resulting modified spectral index evolution in Fig., We present the resulting modified spectral index evolution in Fig.
+ 11. (right panel)., \ref{fig_specindex_timeshift} (right panel).
+ Time bins are the same as in Fig., Time bins are the same as in Fig.
+ 4. (bottom right panel) which should be compared to Fig., \ref{fig_specindex} (bottom right panel) which should be compared to Fig.
+ 11. (right panel)., \ref{fig_specindex_timeshift} (right panel).
+ Comparison of the two diagrams indicates that removing the spectral time delay systematically shifts a towards steeper values by =0.1., Comparison of the two diagrams indicates that removing the spectral time delay systematically shifts $\alpha$ towards steeper values by $\approx0.1$.
+ Our analysis reveals substantial — by factors «2-8) — flux variability on timescales of several years in all sources of our sample., Our analysis reveals substantial – by factors $\approx$ 2–8 – flux variability on timescales of several years in all sources of our sample.
+ Even though all lighteurves (see Fig. 1)), Even though all lightcurves (see Fig. \ref{fig_lightcurve}) )
+ show strong variations. the characteristic (exponential rise/decay) timescales differ by more than one order of magnitude. ranging from x0.5 years (3C454.3) to =7 years (09234392).," show strong variations, the characteristic (exponential rise/decay) timescales differ by more than one order of magnitude, ranging from $\lesssim$ 0.5 years (3C454.3) to $\approx$ 7 years (0923+392)."
+" Accordingly. we find lighteurves with smooth morphologies during the entire z14-year monitoring time (09234392, 3C 84) as well as ltghteurves that are “spiky” (3C ΤΙ. 3C 454.3)."," Accordingly, we find lightcurves with smooth morphologies during the entire $\approx$ 14-year monitoring time (0923+392, 3C 84) as well as lightcurves that are “spiky” (3C 111, 3C 454.3)."
+ The flux distributions (Figs. 2.. 3) ," The flux distributions (Figs. \ref{fig_fluxdist1}, , \ref{fig_fluxdist2}) )"
+of five out of six (the exception being 3C 111) AGN show ΡΙ- or even multimodality., of five out of six (the exception being 3C 111) AGN show bi- or even multimodality.
+ This indicates that the flux variability of the concerned sources cannot be deseribed by uniform stochastic processes:, This indicates that the flux variability of the concerned sources cannot be described by uniform stochastic processes:
+technique in Deshpande Raulsin. which proceeds under the assumption that the depolarized samples contain equal PPM aud SPM levels of power.,"technique in Deshpande Rankin, which proceeds under the assumption that the depolarized samples contain equal PPM and SPM levels of power."
+" ""Though pulsar 2020]28's imodal behavior appears more complex(e.y.. Cordes 11978: MS98). we see may of the same features‘or dustanee. that the UP power approaches the otal power ou the extreme; edgesge of profile"," Though pulsar 2020+28's modal behavior appears more complex, Cordes 1978; MS98), we see many of the same features---for instance, that the UP power approaches the total power on the extreme edges of profile."
+" Tudeed. MS98*s analysis based ou ""superposed nodes” sugeests simular couchisions."," Indeed, MS98's analysis based on “superposed modes” suggests similar conclusions."
+ The well Ποστος profile) demonstrates that Πω πο componeits cach have a good deal of structureSOOL aS “breaks” in the total-power curvesbut he PPA. SPAL and UP curves demonstrate. in adclition. that auch of t16 colplesity is modal iu origin.," The well measured profile demonstrates that its “two” components each have a good deal of structure---seen as “breaks” in the total-power curves—but the PPM, SPM and UP curves demonstrate, in addition, that much of the complexity is modal in origin."
+ T16 complex modal behavior of this pulsar deserves much fuller stidy. and a well measured polariuctiic pulse sequence in the 100200-MITz range would add much to our knowledge.," The complex modal behavior of this pulsar deserves much fuller study, and a well measured polarimetric pulse sequence in the 100–200-MHz range would add much to our knowledge."
+ Overall. we see that the conal compoucut pairs depicted in Figure 2. all have moderate to hie evels of fractional linear polarizationthat is. vpically some though most have narrow. interior regions of longitude where the Linear »olarization is higher.," Overall, we see that the conal component pairs depicted in Figure \ref{fig2} all have moderate to high levels of fractional linear polarization—that is, typically some —though most have narrow, interior regions of longitude where the linear polarization is higher."
+ We shall see that this stands in sharp contrast to the pulsars considered iu 1ο following section., We shall see that this stands in sharp contrast to the pulsars considered in the following section.
+ Our point is that when |./|/p is relatively sinallproducing well resolved coual conrponeut pairsthe mode mixing depoluizes 1e outer edges. but not the profile interior.," Our point is that when $|\beta|/\rho$ is relatively small—producing well resolved conal component pairs—the mode mixing depolarizes the outer edges, but not the profile interior."
+ This ien reflects propertiesboth dynamic iu terius of modulation phenomena and polavizationalof ie conal enidssion beam. aud we nmst reflect on just how this is possible.," This then reflects properties—both dynamic in terms of modulation phenomena and polarizational—of the conal emission beam, and we must reflect on just how this is possible."
+ We have just considered a group of stars in which our sighthue makes a fairly ceutral traverse through their' emissionDOM cone(s). and we now tim fo members of the conal single eroup. all of which are configure by a tanecutial traverse along he average[m] enüsson conce.," We have just considered a group of stars in which our sightline makes a fairly central traverse through their emission cone(s), and we now turn to members of the conal single group, all of which are configured by a tangential traverse along the average emission cone."
+ Tere we have the oyportumity both to explore the conal depolarizaion phenomena i τα verv differeut ecolctrical couext and then to investigate how the modulation and depolarizajon phenomena are connected., Here we have the opportunity both to explore the conal depolarization phenomena in a very different geometrical context and then to investigate how the modulation and depolarization phenomena are connected.
+ Figure 3 elves uode-segreesated polarization plos (Similar to those in Fie. 2)), Figure \ref{fig3} gives mode-segregated polarization plots (similar to those in Fig. \ref{fig2}) )
+ for six stars., for six stars.
+" ere. it is niportaut to keep in mind that each ofthese pulsars has prominent ""diiftiug"" subpulses. so that the profiles eive ouly a static average of t1e subpulse polarization."," Here, it is important to keep in mind that each of these pulsars has prominent “drifting” subpulses, so that the profiles give only a static average of the subpulse polarization."
+ The displays ofFig., The displays of Fig.
+ 3 show that the UP (perhaps. mode-mixed) yower is fvXeallv of the total. so that he overall inodal contributions are comparable aud he aggregate lear polarization is often small.," \ref{fig3} show that the UP (perhaps, mode-mixed) power is typically of the total, so that the overall modal contributions are comparable and the aggregate linear polarization is often small."
+ While all of the toal-power profiles are roughly unimiodal (only BoOs20|02 is really svuuuetrical). he modal profies are more coniplex: the two peaks. at cüffereut onugituces in DUSO9|71: the SPAL las a double form 1 D2016|28: anc we have already roted the pecular “triple” form. of the agerceateOO cay n pulsar 08532yy02.," While all of the total-power profiles are roughly unimodal (only B0820+02 is really symmetrical), the modal profiles are more complex; the two peak at different longitudes in B0809+74; the SPM has a double form in B2016+28; and we have already noted the peculiar “triple” form of the aggregate linear in pulsar 0820+02."
+ As a class. t16 stars exhibit couspieuously depolarized profiles at metre wavelengths.," As a class, the stars exhibit conspicuously depolarized profiles at metre wavelengths."
+ Iudeed. his has been one of the ereat obstacles to understaudius thei characteristics. because. for nany(Ae. ΕΟΟΤΟ. the modal complexity and low fractional linear polarization male it difteult to accurately deteziumne even such a siuple paranieter as the PA sweep rate(ey. Iunachandzan 22002).," Indeed, this has been one of the great obstacles to understanding their characteristics, because, for many, 0809+74), the modal complexity and low fractional linear polarization make it difficult to accurately determine even such a simple parameter as the PA sweep rate, Ramachandran 2002)."
+ Puradoxicallv. some also have nearly complete linear polarization at certain lougitudes and frequencies(50. as docs 080917Es leading edee at higher frequencies) suggesting that ulxing is hot always operativo.," Paradoxically, some also have nearly complete linear polarization at certain longitudes and frequencies, as does 0809+74's leading edge at higher frequencies) suggesting that mode-mixing is not always operative."
+ The pulsars are also the profile class uost closely associated with the problematic phenomenon of “absorption”., The pulsars are also the profile class most closely associated with the problematic phenomenon of “absorption”.
+" It was in 0809|71 hat the effect was first identified (Bartel 11981: Bartel 1981)that is. evicence that parts of the profiles were ""niüssiusand strong evidence o this effect through subbeam-napping methods rave also been adduced for 0913|10 (Deshpande Rankin 20Hn)."," It was in 0809+74 that the effect was first identified (Bartel 1981; Bartel 1981)—that is, evidence that parts of the profiles were “missing”—and strong evidence to this effect through subbeam-mapping methods have also been adduced for 0943+10 (Deshpande Rankin 2001)."
+ Surely oue coud nuaegine from 3913]108. asvuunetric profile that a part of its trailine-cdee cussion ds “absorbed”. though WOO|7Vs more svinmuetric leading edee at meter waveleueths gives little clue that cussion appears o be missing here as well.," Surely one could imagine from 0943+10's asymmetric profile that a part of its trailing-edge emission is “absorbed”, though 0809+74's more symmetric leading edge at meter wavelengths gives little clue that emission appears to be missing here as well."
+" In short. the eireiuustauc""X defining the profile edges appear to be more colmplicated for conal single stars than for the other species. aud their modal polarization characteristics are all aspect €of this complexity."," In short, the circumstances defining the profile edges appear to be more complicated for conal single stars than for the other species, and their modal polarization characteristics are an aspect of this complexity."
+,blazars.
+example ofa cdiffereut blazar categorization that may, With more statistics we may be able to divide the class into subsets that better clarify the processes behind the dramatic variability.
+ vield optimally con, We thank S. G. Djorgovski and his group for helpful discussions.
+trasting. aud therefore ΠεΠΟ ΠΕvariability results. Suchrebinnine, This work is supported by the Office of Science of the Department of Energy and the National Science Foundation.
+ We haveused thePalomar-QUEST Survey to examine the," This research has made use of data from the MOJAVE (Lister and Homan, 2005, AJ, 130, 1389), 2cm Survey (Kellermann et al.,"
+ ensemble optical variability," 2004, ApJ, 609, 539), and CRATES (Healey et al.,"
+ of276 FSRQs aud86 BLLacs.these," 2007, ApJS, 171, 61) programs."
+ of276 FSRQs aud86 BLLacs.theseo," 2007, ApJS, 171, 61) programs."
+ of276 FSRQs aud86 BLLacs.theseof," 2007, ApJS, 171, 61) programs."
+annulus is given by where NyGc.0) denotes the number of observable modes in the annulus (only modes whose line-of-sight components fit within the observed. bandpass are included).,"annulus is given by where $N_{\rm m}(k,\theta)$ denotes the number of observable modes in the annulus (only modes whose line-of-sight components fit within the observed bandpass are included)."
+ In terms of the A-vector components & and 8. the number of independent Fourier modes within an annulus of radial width αλ and angular width d6 is Αμ.)=2rhTVsin(6)dkd6/ (2x). where V.τεDAD(OCLA) is the. observed. volume.," In terms of the $k$ -vector components $k$ and $\theta$, the number of independent Fourier modes within an annulus of radial width $dk$ and angular width $d\theta$ is $N_{\rm m}(k,\theta) = 2\pi k^2 V \sin(\theta) dk\,d\theta/(2\pi)^3$ , where $V = D^2\Delta D(\lambda^2/A_{\rm e})$ is the observed volume."
+ Averaging ap(A.@) over 6 gives the spherically averaged sensitivity to the 21-em power spectrum ap(fr). In order to find the sensitivity in terms ofAS); we use," Averaging $\sigma_P(k,\theta)$ over $\theta$ gives the spherically averaged sensitivity to the 21-cm power spectrum $\sigma_P(k)$, In order to find the sensitivity in terms of$\Delta^2_{21}$ we use"
+As the gas density in these star [ormüng disces is enormous compared to “norm” molecular clouds. the 3oncli-Llovle formula frequently vields super-IEXddington accretion rates (Navakshin.2006).,"As the gas density in these star forming discs is enormous compared to “normal” molecular clouds, the Bondi-Hoyle formula frequently yields super-Eddington accretion rates \citep{Nayakshin06a}."
+" Assuming| that gas liberates CAL.fA, of energy. per unit mass. where M; and R, ave the mass and radius of the accretor. one estimates that the accretion luminosity is where Al is the accretion rate."," Assuming that gas liberates $G M_*/R_*$ of energy per unit mass, where $M_*$ and $R_*$ are the mass and radius of the accretor, one estimates that the accretion luminosity is where $\dot M$ is the accretion rate."
+" Setting this equal to the Edeington luminosity. Lear=4286mycfar13810(AL,/M.) erg s1 we obtain the EddingtonM accretion rate: Note that this expression depends only on the size of the object. &,."," Setting this equal to the Eddington luminosity, $\ledd = 4\pi G M_* m_p c/\sigma_T = 1.3 \times 10^{38}
+(M_*/\msun)$ erg $^{-1}$, we obtain the Eddington accretion rate: Note that this expression depends only on the size of the object, $R_*$."
+ For first cores. Ry=Rooye. which vields avery high aceretion rate limit of almost a Solar mass per vear for Rowe=1073 em.," For first cores, $R_* = R_{\rm core}$, which yields a very high accretion rate limit of almost a Solar mass per year for $R_{\rm core} = 10^{14}$ cm."
+ For stars. we use the observational results of Demircan&Wahraman(1991):GordaSvech-nikoy (1998): For R=R.. this vields an Exldingtong accretion rate limit of 510M.," For stars, we use the observational results of \cite{1991Ap&SS.181..313D,
+1998ARep...42..793G}: For $R=\rsun$, this yields an Eddington accretion rate limit of $\sim 5 \times 10^{-4} \msun$ $^{-1}$."
+" We use the tests with cooling parameter 3?=2 and 3. iellv described in Section οὃν, to point out some general rends seen in our simulations."," We use the tests with cooling parameter $\beta=2$ and $3$, briefly described in Section \ref{sec:fragm}, to point out some general trends seen in our simulations."
+ “Phese runs were continued zu longer than was needed to form the first gravitationally »»und. clumps., These runs were continued far longer than was needed to form the first gravitationally bound clumps.
+ Sink particles were introduced as described in Sections 77. and ?7.., Sink particles were introduced as described in Sections \ref{sec:fragm} and \ref{sec:sink}.
+" The collapse parameter 2h.) was set to 30 (see equation 3)). and the gravitational focusing xwameter is f,=3."," The collapse parameter $A_{\rm col}$ was set to 30 (see equation \ref{rhosink}) ), and the gravitational focusing parameter is $f_m = 3$."
+ Figure 2. shows a column density map of a snapshot mace at dimensionless time /=τὸ for run 82 (3= 2)., Figure \ref{fig:fig1} shows a column density map of a snapshot made at dimensionless time $t=75$ for run S2 $\beta=2$ ).
+ The eft panel of the figure shows the whole disc. whereas the right panel zooms in on a smaller region centred on we=1.5.," The left panel of the figure shows the whole disc, whereas the right panel zooms in on a smaller region centred on $x=1.8$."
+ ted asterisks show stars more massive that 3M..., Red asterisks show stars more massive that $\msun$.
+ We do not show lighter stars and first cores for clarity in the figure., We do not show lighter stars and first cores for clarity in the figure.
+ Since gas dynamical ancl cooling times are shorter for smaller radii. high density regions develop faster at smaller radii.," Since gas dynamical and cooling times are shorter for smaller radii, high density regions develop faster at smaller radii."
+ The sequence of events is as follows., The sequence of events is as follows.
+ “Phe innermost οσο of the disc forms collapsed haloes in which sink particles are introcuced., The innermost edge of the disc forms collapsed haloes in which sink particles are introduced.
+ These sink particles grow by mergers with other sink particles and gas accretion., These sink particles grow by mergers with other sink particles and gas accretion.
+ With time. the gaseous disc is depleted while the stellar cise grows.," With time, the gaseous disc is depleted while the stellar disc grows."
+ Looking at larger /? in Fig. 2.," Looking at larger $R$ in Fig. \ref{fig:fig1},"
+ we observe an intermediate epoch in star formation. where bound. high density clumps have formed but there are no stars more massive than M. vet.," we observe an intermediate epoch in star formation, where bound high density clumps have formed but there are no stars more massive than $\msun$ yet."
+ This is because some of these clumps have not vet satisfied the fragmentation criterion. and others satisfied it only recently. thus containing only first cores and low-mass stars.," This is because some of these clumps have not yet satisfied the fragmentation criterion, and others satisfied it only recently, thus containing only first cores and low-mass stars."
+ Finally. at the largest radii. we see dense filaments (spiral structure) but no well defined bound clumps.," Finally, at the largest radii, we see dense filaments (spiral structure) but no well defined bound clumps."
+ Figure 3. is identical to Fig. 2..," Figure \ref{fig:fig2} is identical to Fig. \ref{fig:fig1},"
+ but shows test 83 (1.6. j—3 instead of ο=2)., but shows test S3 (i.e. $\beta=3$ instead of $\beta=2$ ).
+ Comparing the two runs. itis ipparent that fragmentation of the more gradually. cooling disc (83) is naturally much slower than in the test S2.," Comparing the two runs, it is apparent that fragmentation of the more gradually cooling disc (S3) is naturally much slower than in the test S2."
+ There are fewer stars in the disc. and there are fewer high density clumps.," There are fewer stars in the disc, and there are fewer high density clumps."
+ As we shall see later. within the given model. this allows the stars to grow to larger masses. resulting in a more top-heavy mass function.," As we shall see later, within the given model, this allows the stars to grow to larger masses, resulting in a more top-heavy mass function."
+ ligure 4. shows the column density maps for the run 83 at two later times. /=175 and /=275. in the left and right panels. respectively.," Figure \ref{fig:fig3} shows the column density maps for the run S3 at two later times, $t=175$ and $t=275$, in the left and right panels, respectively."
+ Star formation spreads to larger radii as time progresses., Star formation spreads to larger radii as time progresses.
+ In the right panel of Fig. 4..," In the right panel of Fig. \ref{fig:fig3},"
+ the innermost region of the dise is almost completely. eleared of eas bv continuing accretion onto existing stars (rather than further disc fragmentation. see Section 22)).," the innermost region of the disc is almost completely cleared of gas by continuing accretion onto existing stars (rather than further disc fragmentation, see Section \ref{sec:quenching}) )."
+ At the end of the simulations that co fragment. Le. SL83. very little gas remains.," At the end of the simulations that do fragment, i.e. S1–S3, very little gas remains."
+ Jo characterise the rate at which star formation consumes the gas in the disc. we define the disc half-time. fia as the time from the start of the simulation to the point at which the gaseous mass in the disc is equal to one half of the initial mass.," To characterise the rate at which star formation consumes the gas in the disc, we define the disc half-time, $t_{\rm half}$, as the time from the start of the simulation to the point at which the gaseous mass in the disc is equal to one half of the initial mass."
+ Table 1 lists fis for each simulation., Table 1 lists $t_{\rm half}$ for each simulation.
+ Longer cooling times allow the dise to survive longer., Longer cooling times allow the disc to survive longer.
+ Nevertheless. even for the run 8S3. where fiai2ανα the hall-time of the disc is only around. 12.000 vears in physical units.," Nevertheless, even for the run S3, where $t_{\rm half} \gg
+t_{\rm dyn}$ , the half-time of the disc is only around 12,000 years in physical units."
+ One notices that some of the stars are followed and/or preceeded: by. low column depth regions., One notices that some of the stars are followed and/or preceeded by low column depth regions.
+ These are caused by angular momentum transfer. [rom the stars to. the surrounding gas (Goldreich&Tremaine.1980)., These are caused by angular momentum transfer from the stars to the surrounding gas \citep{GoldreichTremaine80}.
+ Lo stars are massive enough. they could open up gaps in accretion cdises (Sveretal.1991:Navakshin.2004) in the way planets open gaps in proto-planetary dises.," If stars are massive enough, they could open up gaps in accretion discs \citep{Syer91,Nayakshin04} in the way planets open gaps in proto-planetary discs."
+ We do not observe such “clean” gaps around stars in our simulations., We do not observe such “clean” gaps around stars in our simulations.
+ We believe the main reason for this is the large number of stars that we have here., We believe the main reason for this is the large number of stars that we have here.
+ “Phe interactions between the gravitational potentials of neighboring stars. destroy. the resonant charachter of the gas-star interactions (Goldreich&“Tremaine.1980)., The interactions between the gravitational potentials of neighboring stars destroy the resonant charachter of the gas-star interactions \citep{GoldreichTremaine80}.
+. 1n addition. selt-interactions between stars force stars on slightly. elliptical and slightly inclined orbits (NavakshinCuaclra. 2005)... which also makes gap opening harder.," In addition, self-interactions between stars force stars on slightly elliptical and slightly inclined orbits \citep{NC05}, which also makes gap opening harder."
+ The disc starts olf as a razor-thin gaseous disc and evolves into a thicker stellar disc., The disc starts off as a razor-thin gaseous disc and evolves into a thicker stellar disc.
+ This behaviour is to be expected., This behaviour is to be expected.
+ LInitiallv. gas cooling is sullicicnthy strong. and the disc temperature is regulated to maintain a geometrical thickness of order of 4/12MaiifAdan. which is equivalent to having Q=| (eg.Commie.2001).," Initially, gas cooling is sufficiently strong, and the disc temperature is regulated to maintain a geometrical thickness of order of $H/R \simeq M_{\rm disc}/\mbh$, which is equivalent to having $Q\approx 1$ \citep[e.g.,][]{Gammie01}."
+. When a good fraction of the gas is turned. into stars. however. the rate at which the disc can cool decreases.," When a good fraction of the gas is turned into stars, however, the rate at which the disc can cool decreases."
+ At the same time. a stellar disc. considered on its own. would thicken with time since stellar orbital energy is transferred into stellar random motions bv the normal N-hocly relaxation processes.," At the same time, a stellar disc, considered on its own, would thicken with time since stellar orbital energy is transferred into stellar random motions by the normal N-body relaxation processes."
+ In the intermediate state. when both gas and stars are present in important fractions.disc evolution is not trivial.," In the intermediate state, when both gas and stars are present in important fractions,disc evolution is not trivial."
+ In the runs described in this section. stars transfertheir random energy to the gas via Chandrasekhar's dynamical," In the runs described in this section, stars transfertheir random energy to the gas via Chandrasekhar's dynamical"
+(2003) with only minor modifications.,(2003) with only minor modifications.
+ Because the most significant limitation to this analysis is inevitably the assumptions that must be enforced upon our model. we summarise here the most important aspects of this model.," Because the most significant limitation to this analysis is inevitably the assumptions that must be enforced upon our model, we summarise here the most important aspects of this model."
+ In order to derive a model to fit the observed £(a.3) grid. we need three main ingredients.," In order to derive a model to fit the observed $\xi(\sigma,\pi)$ grid, we need three main ingredients."
+ The first is to assume some form for the real-space correlation function 0)., The first is to assume some form for the real-space correlation function $\xi(r)$.
+ Because we are only going to be concerned with relatively small-scale separations between galaxies (<<20f ΤΜΡο). we shall assume a power-law form of this function. In converting from real space to redshift space the next step is to account for the distortions in the correlation function which are caused by the linear coherent in-fall of galaxies into cluster over-densities (Kaiser 1987: Hamilton 1992). combined with non-linear velocity dispersion (e.g. Peacock et al.," Because we are only going to be concerned with relatively small-scale separations between galaxies $\le 20\;h^{-1}$ Mpc), we shall assume a power-law form of this function, In converting from real space to redshift space the next step is to account for the distortions in the correlation function which are caused by the linear coherent in-fall of galaxies into cluster over-densities (Kaiser 1987; Hamilton 1992), combined with non-linear velocity dispersion (e.g. Peacock et al."
+ 2001)., 2001).
+" The linear-theory fall distortion can be written as (Hamilton 1992): where Ῥ,(7) are Legendre polynomials. j/=cos(8) and 6 is the angle between r and π."," The linear-theory in-fall distortion can be written as (Hamilton 1992): where $_{\ell}(\mu)$ are Legendre polynomials, $\ct = \cos(\theta)$ and $\theta$ is the angle between $r$ and $\pi$."
+" The relations between £,. £(r) and .7 for a simple power-law £(r)=(r/ro) are We use these relations to create a model £'(o.7) which we then convolve with the distribution function of random pairwise motions. (0). to give the tinal model (Peebles 1980): and we choose to represent the random motions by an exponential form. where e is the pairwise peculiar velocity dispersion (often known as σι»)."," The relations between $\xi_\ell$, $\xir$ and $\beta$ for a simple power-law $\xir=(r/r_0)^{-\gamma}$ are We use these relations to create a model $\xi'(\sigma, \pi)$ which we then convolve with the distribution function of random pairwise motions, $f(v)$, to give the final model (Peebles 1980): and we choose to represent the random motions by an exponential form, where $a$ is the pairwise peculiar velocity dispersion (often known as $\sigma_{12}$ )."
+ An exponential form for the random motions has been found to fit the observed data better than other functional forms RRateliffe 1998: Landy 2002)., An exponential form for the random motions has been found to fit the observed data better than other functional forms Ratcliffe 1998; Landy 2002).
+" The factor 3=—OT""771/b arises from the growth rate in linear theory. which is almost independent of the cosmological constant (Lahav 11991). combined with the scale-independent biasing parameter b."," The factor $\beta \equiv \Omega_{\rm m}^{0.6}/b$ arises from the growth rate in linear theory, which is almost independent of the cosmological constant (Lahav 1991), combined with the scale-independent biasing parameter $b$."
+ A simple consequence of this model is that the redshift-space power spectrum will also appear to be amplified compared to its real-spacec, A simple consequence of this model is that the redshift-space power spectrum will also appear to be amplified compared to its real-space.
+ounterpart! It is worth explicitly re-stating that all of these derivations are based upon the assumptions of the linear theory of perturbations and linear bias. and assume the far-tield approximation (although this is not a significant issue for the 2dFGRS. for which zidua=0.10.," It is worth explicitly re-stating that all of these derivations are based upon the assumptions of the linear theory of perturbations and linear bias, and assume the far-field approximation (although this is not a significant issue for the 2dFGRS, for which $z_{\rm median} = 0.1$ )."
+ It is also interesting to note that there is another cosmologica effect that can result in the flattening of the observed £(m.3) contours.," It is also interesting to note that there is another cosmological effect that can result in the flattening of the observed $\xi(\sigma,\pi)$ contours."
+ It was first noted by Alcock Paczynski (1979) that the presence of a significant cosmological constant. A. would resul in geometric distortions of the inferred clustering if an incorrec geometry was assumed.," It was first noted by Alcock Paczynski (1979) that the presence of a significant cosmological constant, $\Lambda$, would result in geometric distortions of the inferred clustering if an incorrect geometry was assumed."
+" However. Ballinger. Peacock Heavens (1996) have shown that this is likely to be negligible for the low redshift data-set being considered here. for our assumed model of O,—1.04-03."," However, Ballinger, Peacock Heavens (1996) have shown that this is likely to be negligible for the low redshift data-set being considered here, for our assumed model of $\Omega_{\rm m}= 1-\Omega_\Lambda=0.3$."
+ To summarise. the parameters of our model are therefore the real-space correlation function. €(r)=(rfro). 3. and flr). parameterised in terms of the velocity dispersion @.," To summarise, the parameters of our model are therefore the real-space correlation function, $\xi(r)=(r/r_0)^{-\gamma}$, $\beta$, and $f(v)$ , parameterised in terms of the velocity dispersion $a$."
+ The best-fitting parameters of this model can now be determined from the model correlation function that best matches the observed £(a.7).," The best-fitting parameters of this model can now be determined from the model correlation function that best matches the observed $\xi(\sigma,\pi)$."
+ We have calculated the real-space correlation function. €(1). independently using the non-parametric method of Saunders. Rowan-Robinson Lawrence (1992). in order to confirm the range over which it is a power-law (see Fig. 3).," We have calculated the real-space correlation function, $\xi(r)$, independently using the non-parametric method of Saunders, Rowan-Robinson Lawrence (1992), in order to confirm the range over which it is a power-law (see Fig. \ref{fig:xx}) )."
+ For both samples of galaxies. £(17) is adequately tit by a power-law to separations of + Mpc.," For both samples of galaxies, $\xi(r)$ is adequately fit by a power-law to separations of $r<20$ $h^{-1}$ Mpc."
+ This limit provides the upper bound to which we can compare the observed £(o.7) with our assumed model.," This limit provides the upper bound to which we can compare the observed $\xi(\sigma,\pi)$ with our assumed model."
+ In addition. because we have assumed the linear theory of perturbations in deriving our model we must impose a lower limit to the separations that we will use in our fit.," In addition, because we have assumed the linear theory of perturbations in deriving our model we must impose a lower limit to the separations that we will use in our fit."
+ To ensure that we have a sufficiently large fitting range we set this lower-limit to s=8f+ Mpc., To ensure that we have a sufficiently large fitting range we set this lower-limit to $s=8\;h^{-1}$ Mpc.
+ In fact it is quite plausible that the assumptions of linear theory are no longer valid at this separation., In fact it is quite plausible that the assumptions of linear theory are no longer valid at this separation.
+ However. as will be shown. the ability of our model to recover the observed £(o.7) at these scales is reassuring.," However, as will be shown, the ability of our model to recover the observed $\xi(\sigma,\pi)$ at these scales is reassuring."
+ On large scales it is known that the correlation function must deviate from a pure power law form. and there is some evidence to support this in Section 3.2.. on scales ~205.+ Mpc.," On large scales it is known that the correlation function must deviate from a pure power law form, and there is some evidence to support this in Section \ref{section:xir}, on scales $\sim20\;h^{-1}$ Mpc."
+ A number of methods to account for this expected curvature in £(r) were investigated in Hawkins (2003)., A number of methods to account for this expected curvature in $\xi(r)$ were investigated in Hawkins (2003).
+ However. the analysis presented there suggests that so long as we our restrictfitting range tor«20f+ Mpe the curvature has a negligible impact upon our parameter estimation.," However, the analysis presented there suggests that so long as we restrict our fitting range to $r<20\;h^{-1}$ Mpc the curvature has a negligible impact upon our parameter estimation."
+ For this reason we neglect the possibility of curvature in the present analysis and restrict ourselves to using the simple power-law form for £(i)., For this reason we neglect the possibility of curvature in the present analysis and restrict ourselves to using the simple power-law form for $\xi(r)$.
+ The other major assumption we have made is that the peculiar velocity distribution. (0). has an exponential form.," The other major assumption we have made is that the peculiar velocity distribution, $f(v)$, has an exponential form."
+ This can be tested using the method outlined by Landy. Szalay Broadhurst (1998).," This can be tested using the method outlined by Landy, Szalay Broadhurst (1998)."
+ This method makes a non-parametric estimate of the velocity distribution. using the Fourier decompositions of the observed £(o.z) grid along theAz=O and ἂν=O axes.," This method makes a non-parametric estimate of the velocity distribution, using the Fourier decompositions of the observed $\xi(\sigma,\pi)$ grid along the$k_\sigma=0$ and $k_\pi=0$ axes."
+ Unfortunately this method ignores the effects of coherent in-fall. which ean substantially change the resulting estimate of @ (see," Unfortunately this method ignores the effects of coherent in-fall, which can substantially change the resulting estimate of $a$ (see"
+where x is a unit vector in the wv direction.,where $\hat{\mathbf{x}}$ is a unit vector in the $x$ direction.
+ We consider only the Olimic dissipation in equation (4)) because influence of ambipolar diffusion is negligible at the density level that is expected in the inner (S100AU) midplane of a protoplanetary disk (Jin1996:Sano1999:Chiang&Murrav-Clay 2007).," We consider only the Ohmic dissipation in equation \ref{induction}) ) because influence of ambipolar diffusion is negligible at the density level that is expected in the inner $\lesssim 100\rm AU$ ) midplane of a protoplanetary disk \citep{jin96, sam99, chi07}."
+. The last terms in rls., The last terms in r.h.s.
+ of equation (2)) allows the elect of the global pressure gradient to be taken into account in our local model., of equation \ref{motion}) ) allows the effect of the global pressure gradient to be taken into account in our local model.
+" The parameter 3 is related to the global pressure gradient. as where we assume 2?x rr"". with constant a and an isothermal disk e;=Πω. where Lf is disk scale height."," The parameter $\beta$ is related to the global pressure gradient as where we assume $P \propto r^{\alpha}$ , with constant $\alpha$ and an isothermal disk $c_{s}=H\Omega$, where $H$ is disk scale height."
+ The value of ο(=all/r=0c./r£2) is treated as a constant. parameter in our local model., The value of $\beta \left(=\alpha H/r =\alpha c_s/r\Omega \right)$ is treated as a constant parameter in our local model.
+ We solve equation (4)) with MOCCT method (Stone&Norman1992) and the advection terms in equation (2)) with CIP method (Yabe&Aoki1991)., We solve equation \ref{induction}) ) with MOCCT method \citep{sto92} and the advection terms in equation \ref{motion}) ) with CIP method \citep{yabe91}.
+. The clust oSeralns are (reatecl as test »articles., The dust grains are treated as test particles.
+ The equation of motion of the i-th »article ls given by where w; is (he gas velocity at the location of the /-ih particle. which is interpolated with the three-dimensional first-order interpolation scheme by values of u at the eight erid points surrounding the particle.," The equation of motion of the $i$ -th particle is given by where $\mathbf{w}_{i}$ is the gas velocity at the location of the $i$ -th particle, which is interpolated with the three-dimensional first-order interpolation scheme by values of $\mathbf{u}$ at the eight grid points surrounding the particle."
+ The friction force. the last term in equation (9)). is proportional to the velocity difference between (he dust ancl gas since the dust sizes (hat we are concerned with are in Epstein and Stokes drag-law regimes.," The friction force, the last term in equation \ref{particle}) ), is proportional to the velocity difference between the dust and gas since the dust sizes that we are concerned with are in Epstein and Stokes drag-law regimes."
+" The characteristic friction (stopping) ümescale τι Is given by where dg. p, and ὃς are the dust particle radius. dust particle density and (he sound speed. respectively."," The characteristic friction (stopping) timescale $\tau_{f}$ is given by where $a_{\bullet}$ , $\rho_{\bullet}$ and $c_{s}$ are the dust particle radius, dust particle density and the sound speed, respectively."
+" v is the molecular viscosity and equal (o 0,À/2. where A is the mean free path of the gas molecules."," $\nu$ is the molecular viscosity and equal to $c_{s}\lambda/2$, where $\lambda$ is the mean free path of the gas molecules."
+" Depending on (the particle size relative to (he mean Iree path (a,>or<(9/4)↜∕ A) the drage law changese from the Epstein to the Stokes regime.", Depending on the particle size relative to the mean free path $a_{\bullet} > {\rm or} <\left(9/4\right)\lambda$ ) the drag law changes from the Epstein to the Stokes regime.
+"e The scaled friction times of το{=0.1 and 1.0 correspondto the 5grain sizes of approximately 0.1 and 1 meter. respectively. atthe radial location of Jupiter in the minimum mass solar nebula model,"," The scaled friction times of $\tau_{f}\Omega=0.1$ and $1.0$ correspondto the grain sizes of approximately 0.1 and 1 meter, respectively, atthe radial location of Jupiter in the minimum mass solar nebula model."
+estimation. including an accounting for uncertainties in. the limb-darkening law and due to time-correlated noise. as well as linkage between the light curve parameters and stellar-evolutionary models. that were not always applied by previous authors.,"estimation, including an accounting for uncertainties in the limb-darkening law and due to time-correlated noise, as well as linkage between the light curve parameters and stellar-evolutionary models, that were not always applied by previous authors."
+ In the past these efforts have occasionally led to significant revisions of the planetary dimensions Winnetal.2007b.2008)..," In the past these efforts have occasionally led to significant revisions of the planetary dimensions \citep[see, e.g.][]{2007AJ....134.1707W,2008ApJ...683.1076W}."
+ In the present case our results are in agreement with the previously reported results., In the present case our results are in agreement with the previously reported results.
+ All of the TrES-4 parameters agree to within 2σ with the results reported by Torresetal.(2008).., All of the TrES-4 parameters agree to within $\sigma$ with the results reported by \citet{2008ApJ...677.1324T}.
+ Two of the most important parameters. the mass and radius of the planet. agree to within ασ.," Two of the most important parameters, the mass and radius of the planet, agree to within $\sigma$."
+ Our HAT-P-3 parameters agree to within 20. with those reported by Gibson and the planetary mass and radius agree to within 1.37.," Our HAT-P-3 parameters agree to within $\sigma$ with those reported by \citet{2010MNRAS.401.1917G}, and the planetary mass and radius agree to within $\sigma$."
+" Our WASP-12 parameters agree to within 2o with those reported by Hebbetal.(2009). and the mass and radius of the planet agree to within I,"," Our WASP-12 parameters agree to within $\sigma$ with those reported by \citet{2009ApJ...693.1920H}, and the mass and radius of the planet agree to within $\sigma$."
+ Our transit ephemeris for TrES-4+ agrees with that of Mandushevetal.(2007).. and the refined orbital period ts about 20 times more precise.," Our transit ephemeris for TrES-4 agrees with that of \cite{2007ApJ...667L.195M}, and the refined orbital period is about 20 times more precise."
+ Our ephemeris for HAT-P-3 agrees with that of Gibsonetal.(2010).. and is 2-3 times more precise.," Our ephemeris for HAT-P-3 agrees with that of \citet{2010MNRAS.401.1917G}, and is 2-3 times more precise."
+ Our ephemeris for WASP-12 agrees with that of Hebbetal.(2009).. and is of comparable precision. despite being based on only a few well-documented data points.," Our ephemeris for WASP-12 agrees with that of \citet{2009ApJ...693.1920H}, and is of comparable precision, despite being based on only a few well-documented data points."
+ The reason that TrES-4. HAT-P-3. and WASP-12 are of particular interest is because their measured dimensions do not agree with standard models of gas giant planets.," The reason that TrES-4, HAT-P-3, and WASP-12 are of particular interest is because their measured dimensions do not agree with standard models of gas giant planets."
+ TrES-4 and WASP-12 are heavily bloated. with radit too large for their masses. while HAT-P-3 is too small.," TrES-4 and WASP-12 are heavily bloated, with radii too large for their masses, while HAT-P-3 is too small."
+ With reference to the tables of Fortneyetal.(2007).. TrES-4+ and WASP-12 are incompatible with pure hydrogen-helium giant planets at the 100 and 5o levels. respectively.," With reference to the tables of \citet{2007ApJ...659.1661F}, TrES-4 and WASP-12 are incompatible with pure hydrogen-helium giant planets at the $\sigma$ and $\sigma$ levels, respectively."
+ Enhancing these planets with metals to the degree of Jupiter or Saturn would only make the problem worse., Enhancing these planets with metals to the degree of Jupiter or Saturn would only make the problem worse.
+" HAT-P-3. in contrast. 1s compatible with the tabulated models if it is endowed with approximately 100 M,. of heavy elements."," HAT-P-3, in contrast, is compatible with the tabulated models if it is endowed with approximately 100 $M_\earth$ of heavy elements."
+ Our results do not change these interpretations of the three systems we have studied., Our results do not change these interpretations of the three systems we have studied.
+ Rather. our results lend more confidence to the claims that the dimensions of the planets are anomalous. and merit attention by theoreticians who seek to solve the radius anomaly problem.," Rather, our results lend more confidence to the claims that the dimensions of the planets are anomalous, and merit attention by theoreticians who seek to solve the radius anomaly problem."
+ We thank Gerald Nordley for checking some of the entries in Tables 5-7., We thank Gerald Nordley for checking some of the entries in Tables 5-7.
+ We gratefully acknowledge support from the NASA Origins program through award NNX09AB33G and from the Research Science Institute. a program of the Center for Excellence in Education.," We gratefully acknowledge support from the NASA Origins program through award NNX09AB33G and from the Research Science Institute, a program of the Center for Excellence in Education."
+ Some of the data presented herein were obtained with the Nordic Optical Telescope (NOT). operated on the island of La Palma jointly by Denmark. Finland. [eeland. Norway. and Sweden. in the Spanish Observatorio del Roque de los Muchachos Instituto Astrofisica de Canarias. and ALFOSC. which is owned by the Instituto Astrofisica de Andalucia (AA) and operated at the NOT under agreement between [AA and NBIfAFG of the Astronomical Observatory of Copenhagen.," Some of the data presented herein were obtained with the Nordic Optical Telescope (NOT), operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway, and Sweden, in the Spanish Observatorio del Roque de los Muchachos Instituto Astrofisica de Canarias, and ALFOSC, which is owned by the Instituto Astrofisica de Andalucia (IAA) and operated at the NOT under agreement between IAA and NBlfAFG of the Astronomical Observatory of Copenhagen."
+question.,question.
+ Instead. we shall first give a brief. description of NINOO and its main results. and then re-derive it theory using a combinatorial approach which would enable us to compare it. with the LD67 theory. and. point to the reasons of why they diller.," Instead, we shall first give a brief description of NK00 and its main results, and then re-derive it theory using a combinatorial approach which would enable us to compare it with the LB67 theory, and point to the reasons of why they differ."
+ Finally we will analyse the two-levels configuration which goes into the water-bag configuration in a limiting case., Finally we will analyse the two-levels configuration which goes into the water-bag configuration in a limiting case.
+ The result of this analysis will be used in the next section when we examine the transitivity of the NIXOO theory in a double relaxation experiment., The result of this analysis will be used in the next section when we examine the transitivity of the NK00 theory in a double relaxation experiment.
+ In the NIXOO theory we adopt the probabilistic description of the phase-space density f(£.r.v).Let. fo(r. be the initial phase-space densitv of the system. ancl definev) the initial probability distribution po(r.v) for finding a (single) test point at f=0 by Then we let the test. point move under gravity just ike any phase-space element. and we define the probability distribution. p(r.v./) as the probability distribution. of inding the test point at time /20.," In the NK00 theory we adopt the probabilistic description of the phase-space density $f(t, \Br,\Bv)$.Let $f_0(\Br,\Bv)$ be the initial phase-space density of the system, and define the initial probability distribution $p_0(\Br,\Bv)$ for finding a (single) test point at $t=0$ by Then we let the test point move under gravity just like any phase-space element, and we define the probability distribution $p(\Br,\Bv,t)$ as the probability distribution of finding the test point at time $t>0$."
+ The conservation of ohase-space volume guarantees that p(r.v./)=f(r.v.ΟΛ or all />0.," The conservation of phase-space volume guarantees that $p(\Br,\Bv,t)=f(\Br,\Bv,t)/M$ for all $t>0$."
+ Next. we divide phase-space into macro-cclls¢=1.2.3.... of volume ac. and celine the coarse-erainecl oobabilitv p; as the probability of finding the point in he Zth macro-cell when the svstem reaches an equilibrium.," Next, we divide phase-space into macro-cells$i=1,2,3,\ldots$ of volume $\tw$, and define the coarse-grained probability $\bar{p}_i$ as the probability of finding the point in the $i$ 'th macro-cell when the system reaches an equilibrium."
+ From the above discussion it is clear that p; is equal to f;/M with f; being the coarse-grained. DE in the macro-cell 7 at ecquilibrium., From the above discussion it is clear that $\bar{p}_i$ is equal to $\bar{f}_i/M$ with $\bar{f}_i$ being the coarse-grained DF in the macro-cell $i$ at equilibrium.
+ To calculate p; using the information-theory approach. we define the p;(v.v) which measures the probability of initially finding the test point at the (r.v) and later at the macro-cell 7.," To calculate $\bar{p}_i$ using the information-theory approach, we define the $p_i(\Br,\Bv)$ which measures the probability of initially finding the test point at the $(\Br,\Bv)$ and later at the macro-cell $i$."
+ Phen we maximise the Shanon entropy subject to constraints of energy. conservation. phase-space volume conservation and. initial conditions.," Then we maximise the Shanon entropy subject to constraints of energy conservation, phase-space volume conservation and initial conditions."
+ The resultant. distribution can be best written in terms of the Ai(v.v)=pi(r.v)/po(v.v): with e; being the energy per unit mass of the macro-cell foamed 3.0(x.v).A; ave the Lagrange multipliers. to be found from the energy conservation constraint. from the initial conditions and from the phase-space volume conservation constraint From Eqs. (35--87))," The resultant distribution can be best written in terms of the $K_i(\Br,\Bv) = p_i(\Br,\Bv)/p_0(\Br,\Bv)$: with $\epsilon_i$ being the energy per unit mass of the macro-cell $i$, and $\beta, \delta(\Bx,\Bv), \lambda_i$ are the Lagrange multipliers, to be found from the energy conservation constraint, from the initial conditions and from the phase-space volume conservation constraint From Eqs. \ref{eq:NK-K}- \ref{eq:NK-con2}) )"
+ it is evident that the dependence of A(rev) on the indices 7.Y.v is only via e; ancl po(r.v). the latter can be trivially be replaced by οιv).," it is evident that the dependence of $K_i(\Br,\Bv)$ on the indices $i, \Br, \Bv$ is only via $\epsilon_i$ and $p_0(\Br,\Bv)$, the latter can be trivially be replaced by $f_0(\Br,\Bv)$."
+ Fherefore. the above equations can be re-written using the A(e.7) Function. together with the τί) and gfe) functions which were defined in Sec. 3.1:," Therefore, the above equations can be re-written using the $K(\epsilon, \eta)$ function, together with the $\tau(\eta)$ and $g(\epsilon)$ functions which were defined in Sec. \ref{sec:LB-res}:"
+ The coarse-grained equilibrium DE is then given by As noted by Nakamura. a prominent difference between his result ancl the LD67 results is that in the non-degenerate limit his expression reduces to a single. Maxwellian istribution. whereas the LD67 expression is a superposition X Maxwellian clistributions with cdillerent dispersions.," The coarse-grained equilibrium DF is then given by As noted by Nakamura, a prominent difference between his result and the LB67 results is that in the non-degenerate limit his expression reduces to a single Maxwellian distribution, whereas the LB67 expression is a superposition of Maxwellian distributions with different dispersions."
+ This ilference can be attributed to the fact that in LBGT we iscretise phase-space using phase-space elements of equal 7'obume anclmasses.. while. as we shall see below. —16 NWOO theory can be derived by cliscretising phase-space using elements. of equal mass. which arc associated with ilferent phase-space volumes.," This difference can be attributed to the fact that in LB67 we discretise phase-space using phase-space elements of equal volume and, while, as we shall see below, the NK00 theory can be derived by discretising phase-space using elements of equal mass, which are associated with different phase-space volumes."
+ Another evident difference comes from the phase-space volume conservation constraint(36)., Another evident difference comes from the phase-space volume conservation constraint.
+. This constraint guarantees that the total phase-space volume of all phase-space patches that ended up in macro-cell ¢ will be equal to 1e macro-cell volume., This constraint guarantees that the total phase-space volume of all phase-space patches that ended up in macro-cell $i$ will be equal to the macro-cell volume.
+ Consequently the total phase-space volume of the initial system au)dr must be equal to the total phase-space volume of the non-vanishing phase-space density in the equilibrium. configuration.," Consequently the total phase-space volume of the initial system $\int_0^\infty\!\! \tau(\eta)\, d\eta$ must be equal to the total phase-space volume of the non-vanishing phase-space density in the equilibrium configuration."
+ This constraint does not exist in the LDB67 theory where the macro-cells can be only partly Full - as is the case. for example. inthe non-degenerate equilibrium of a svstem which is initially in 1ο water-bag configuration.," This constraint does not exist in the LB67 theory where the macro-cells can be only partly full - as is the case, for example, inthe non-degenerate equilibrium of a system which is initially in the water-bag configuration."
+ Furthermore. this can not be trivially changed. by adding a volume of zero phase-space ensitv to the initial condition. because setting p(r.v)=0 would leacl to divergences in(35).," Furthermore, this can not be trivially changed by adding a volume of zero phase-space density to the initial condition, because setting $p(\Br,\Bv)=0$ would lead to divergences in."
+. In Sec., In Sec.
+ 4.3. we shall see how this problem can be overcome by. using a limiting procedure., \ref{sec:NK-2-levels} we shall see how this problem can be overcome by using a limiting procedure.
+ ‘Lo derive the NINO0 theory in a combinatorial approach. we realise the phase-space density distribution using IN.x»1 elements of m.," To derive the NK00 theory in a combinatorial approach, we realise the phase-space density distribution using $N\gg 1$ elements of $m$."
+ As in Sec. 3.1...," As in Sec. \ref{sec:LB-res},"
+ we assume that initially the system is mace of a ciscrete set of density Levels Wet... occupying phase-space volumes V1.15.....," we assume that initially the system is made of a discrete set of density levels $\eta_1, \eta_2, \ldots$ occupying phase-space volumes $V_1, V_2,
+\ldots$."
+ Then the overall number of elements that realise a phase-space density 3; is Nip=η fm., Then the overall number of elements that realise a phase-space density $\eta_J$ is $N_J = V_J\eta_J/m$ .
+ Next. we let the system reach an equilibrium through the process of violent relaxation. ancl divide phase-space into macro-cells of equal volume a. which are label by the index /— 1.2.....," Next, we let the system reach an equilibrium through the process of violent relaxation, and divide phase-space into macro-cells of equal volume $\tw$ , which are label by the index $i=1,2,\ldots$ ."
+ We define a by specifying the macro-cell in which every clement ended up., We define a by specifying the macro-cell in which every element ended up.
+" A is then defined by the matrix ο which counts how many elements that initially realised the density level ap,ended up in the macro-cell ¢.", A is then defined by the matrix $\{ n_{iJ}\}$ which counts how many elements that initially realised the density level $\eta_J$ended up in the macro-cell $i$ .
+ Using (n;;]. the coarse-grained DE. at," Using $\{n_{iJ}\}$ , the coarse-grained DF at"
+(Wolffetal.2002).,\citep{Wolff02}.
+. When complemented with the BYRIJIUN photometry of Bergeronetal.(2001) and the © magnitude (MeCook&Sion1999) the complete spectral energy distribution (SED) of this cool white dwarl is obtained.," When complemented with the BVRIJHK photometry of \citet{BLR01} and the U magnitude \citep{MS99}
+ the complete spectral energy distribution (SED) of this cool white dwarf is obtained."
+ The UV spectrum extends well into the wing of the Lya line and thus provides an excellent measure of the strength anc prolile of the line.," The UV spectrum extends well into the wing of the $\rm Ly \, \alpha$ line and thus provides an excellent measure of the strength and profile of the line."
+ Our fit of a pure I model (Tuy=5820 Kk. loggy= 8.30) to the entire SED of DPM 4729 is excellent (Fig.," Our fit of a pure H model $T_{\rm eff}=5820 \,$ K, $\log
+g=8.30$ ) to the entire SED of BPM 4729 is excellent (Fig."
+ 2)., 2).
+ These valuesof Tug ancl y agree with the values of at the lo level., These valuesof $T_{\rm eff}$ and $g$ agree with the values of \citet{BLR01} at the $1\sigma$ level.
+ The high quality of the fit of the UV spectrum demonstrates the validitv of our model lor the far red wing of the Lva line under the conditions encountered in DPM 4729.," The high quality of the fit of the UV spectrum demonstrates the validity of our model for the far red wing of the $\rm Ly \, \alpha$ line under the conditions encountered in BPM 4729."
+ At the photosphere. the composition is log(1I)=19.3 and log(IIs)=19.0. where » is the number density in em..," At the photosphere, the composition is $\log n({\rm H})=19.8$ and $\log
+n({\rm H_2})=19.0$, where $n$ is the number density in $\rm cm^{-3}$."
+ Despite the lower abundance of Is. it is an important contributor to the Lya line opacity in the far red wing (Fig.," Despite the lower abundance of $\rm H_2$, it is an important contributor to the $\rm Ly \, \alpha$ line opacity in the far red wing (Fig."
+ 2)., 2).
+ Wolffetal.(2002) obtained a good fit of the UV and blue spectrum with a IHe-rich model with IL/Ile—3x10 ., \citet{Wolff02} obtained a good fit of the UV and blue spectrum with a He-rich model with $\rm H/He=3 \times 10^{-5}$ .
+ They, They
+of the PSS sample: Rhee (1996). Itoscoe (1999): ο) the finding of a very tight RYE in. PSS and other three different samples Williek (1997. see below). Yegorova et al. (,"of the PSS sample: Rhee (1996), Roscoe (1999); c) the finding of a very tight RTF in PSS and other three different samples Willick (1997, see below), Yegorova et al. ("
+2007).,2007).
+ The claim has been also tested by comparing the RCs of two samples of spirals (Courteau. 1998. 131 objects: Verheijen. 1997. 30 objects) with the circular velocities predicted by the URC). once that the values of galaxy Luminosity and disk are inserted. in it.," The claim has been also tested by comparing the RCs of two samples of spirals (Courteau, 1998, 131 objects; Verheijen, 1997, 30 objects) with the circular velocities predicted by the $_0$, once that the values of galaxy luminosity and disk length-scale are inserted in it."
+ “Phe face-value result of the test: 2/3 of the RCs are in pretty good agreement with the universal curve. while 1/3 show some disagreement. indicates that the URC is a useful tool to investigate the svstematies of the mass," The face-value result of the test: 2/3 of the RCs are in pretty good agreement with the universal curve, while 1/3 show some disagreement, indicates that the URC is a useful tool to investigate the systematics of the mass"
+since this information simply is not in the data.,since this information simply is not in the data.
+ However. when the angular resolution is converted to a linear scale one finds that to resolve objects out in these cases they must be relatively large.," However, when the angular resolution is converted to a linear scale one finds that to resolve objects out in these cases they must be relatively large."
+" At redshifts of0.1. 0.5 and 1. a resolution of aarcsec corresponds to kkpc. kkpc. and kkpe. respectively (using Ho71.0. On,= 0.27. Que= 0.73)."," At redshifts of 0.1, 0.5 and 1, a resolution of arcsec corresponds to kpc, kpc, and kpc, respectively (using $H_0=71.0$ $\Omega_{\rm M}=0.27$ , $\Omega_{\rm vac}=0.73$ )."
+ Only jets and lobes of radio galaxies are so large. and this would have been noticed in the GGHz images. as is the case in ESOL.," Only jets and lobes of radio galaxies are so large, and this would have been noticed in the GHz images, as is the case in ES011."
+ However. the majority of the sources suffering from high-frequency drop-outs are not visibly resolved. and given these size constraints we conclude that resolution is not the dominant effect.," However, the majority of the sources suffering from high-frequency drop-outs are not visibly resolved, and given these size constraints we conclude that resolution is not the dominant effect."
+ Furthermore. resolution can only explain spectra where both the GGHz and GGHz sensitivity limits are lower than what would be expected from the lower-frequency data. (," Furthermore, resolution can only explain spectra where both the GHz and GHz sensitivity limits are lower than what would be expected from the lower-frequency data. ("
+1) In cases where the GGHz limit is below the extrapolation (as in. e.g.. ES798 and ES973) resolution effects cannot account for the observed discrepancy because we carefully selected baselines such as to obtain matched-resolution images.,"ii) In cases where the GHz limit is below the extrapolation (as in, e.g., ES798 and ES973) resolution effects cannot account for the observed discrepancy because we carefully selected baselines such as to obtain matched-resolution images."
+ In these cases. the spectral steepening must be related to the distribution of the particle energies in the source.," In these cases, the spectral steepening must be related to the distribution of the particle energies in the source."
+ In general. a population of particles with a power-law distribution of energies (with index g) in a uniform magnetic field will result in synchrotronemission with a constant radio spectral index with @=(1—q)/2.," In general, a population of particles with a power-law distribution of energies (with index $q$ ) in a uniform magnetic field will result in synchrotronemission with a constant radio spectral index with $\alpha=(1-q)/2$."
+ Unless fresh particles are continuously injected into the source. the particles. with higher energies will lose their energy faster than lower-energy particles. and the spectrum will steepen.," Unless fresh particles are continuously injected into the source, the particles with higher energies will lose their energy faster than lower-energy particles, and the spectrum will steepen."
+ Hence steeper spectra at higher frequencies could indicate that the AGN has recently been inactive., Hence steeper spectra at higher frequencies could indicate that the AGN has recently been inactive.
+ CSS/GPS sometimes show a change in spectral slope at frequencies of GGHz to 10GGHz (e.g.. Readheadetal.1996.. Μυσία.2003). which is commonly attributed. to synchrotron losses (Readheadetal.1996)).," CSS/GPS sometimes show a change in spectral slope at frequencies of GHz to GHz (e.g., \citealt{Readhead1996}, \citealt{Murgia2003}) ), which is commonly attributed to synchrotron losses \citealt{Readhead1996}) )."
+ In our subsample of 10 sources where independent measurements were available for ary and ats. the median spectral index changes from —].14 to a median of —1.71.," In our subsample of 10 sources where independent measurements were available for $\alpha_{1.4}^{2.4}$ and $\alpha_{4.8}^{8.6}$ , the median spectral index changes from $-1.14$ to a median of $-1.71$."
+ This difference of Ae.=0.57 ts comparable to the change in « reported by Readheadetal.(1996) in the prototypical CSS source JJ2355+4950. which was found to be 0.6.," This difference of $\Delta\alpha=0.57$ is comparable to the change in $\alpha$ reported by \cite{Readhead1996}
+ in the prototypical CSS source J2355+4950, which was found to be 0.6."
+ Whilst this similarity is prone to coincidence and. small-number statistics we note that Middelbergetal.(2008a) already pointed out the similarities between CSS sources and a VLBI-detected IFRS., Whilst this similarity is prone to coincidence and small-number statistics we note that \cite{Middelberg2008a} already pointed out the similarities between CSS sources and a VLBI-detected IFRS.
+ We also searched for polarised emission in our targets., We also searched for polarised emission in our targets.
+ We are in the process of making a careful analysis of the polarisation levels in the ATLAS radio data. accounting correctly for the positivity bias in images of polarised intensity (Hales et al.," We are in the process of making a careful analysis of the polarisation levels in the ATLAS radio data, accounting correctly for the positivity bias in images of polarised intensity (Hales et al."
+ 2010. in prep).," 2010, in prep)."
+ A preliminary result is that three sources in our sample are significantly polarised at GGHz.. whereas upper limits were obtained for the remaining 14 sources.," A preliminary result is that three sources in our sample are significantly polarised at GHz, whereas upper limits were obtained for the remaining 14 sources."
+ The polarised sources are ES509 (P=2.7040.03 mmJy). ES973 (P=0.8340.06 mmJy). and CS703 (P=1.90+0.05 mmJy).," The polarised sources are ES509 $\pm$ mJy), ES973 $\pm$ mJy), and CS703 $\pm$ mJy)."
+ All other sources are unpolarised at the various levels indicated in Table 1., All other sources are unpolarised at the various levels indicated in Table 1.
+ The limits are given at >99% credibility (Bayesian confidence) that sources are unpolarised (Vaillancourt 2006))., The limits are given at $>$ credibility (Bayesian confidence) that sources are unpolarised \citealt{Vaillancourt2006}) ).
+ It is well-established that there is a tight correlation. extending over five orders of magnitude. between the radio and far-infrared (FIR) luminosity. or flux density. of. star-forming galaxies (vanderKruit1973.. Condonetal. 1982..Dickey 1984.. deJongetal. 1985)).," It is well-established that there is a tight correlation, extending over five orders of magnitude, between the radio and far-infrared (FIR) luminosity, or flux density, of star-forming galaxies \citealt{vanderKruit1973}, \citealt{Condon1982}, ,\citealt{Dickey1984}, , \citealt{deJong1985}) )."
+ It is attributedto massive star. formation. which generates far-infrared emission by heating dust. and generates radio emission by accelerating cosmic rays that then generate radio synchrotron," It is attributedto massive star formation, which generates far-infrared emission by heating dust, and generates radio emission by accelerating cosmic rays that then generate radio synchrotron"
+show a bump around 11.3 that is not detected in the low resolution spectra.,show a bump around 11.3 $\mu$ m that is not detected in the low resolution spectra.
+ wmThis feature originates from aromatic emission and is very common in HII regions., This feature originates from aromatic emission and is very common in HII regions.
+ The reason we detect it at high resolution and do not detect it at low resolution can be explained by the difference in the background subtraction in the two observing modes., The reason we detect it at high resolution and do not detect it at low resolution can be explained by the difference in the background subtraction in the two observing modes.
+ In the short-low (SL) mode we subtract the two nod positions from each other to eliminate the background while in high resolution mode a different spectrum taken at a nearby location was used for the same purpose., In the short-low (SL) mode we subtract the two nod positions from each other to eliminate the background while in high resolution mode a different spectrum taken at a nearby location was used for the same purpose.
+ We see the 11.3 um line in the latter case as a result of imperfect background subtraction due to spatial variation of the line flux., We see the 11.3 $\mu$ m line in the latter case as a result of imperfect background subtraction due to spatial variation of the line flux.
+ To support this explanation we show the relevant part of the low resolution spectrum of the source in NGC 2264 (Fig., To support this explanation we show the relevant part of the low resolution spectrum of the source in NGC 2264 (Fig.
+ 5 left panel) and the same part of a low resolution spectrum taken at the nearby background location where the high resolution background spectra were taken (Fig., \ref{fig:compPAH} left panel) and the same part of a low resolution spectrum taken at the nearby background location where the high resolution background spectra were taken (Fig.
+ 5 right panel)., \ref{fig:compPAH} right panel).
+ The two sharp vertical lines are the spectra of our target and a nearby star that happened to fall on the slit while the smeared out horizontal line in the middle of the image is the 11.3 wm aromatic feature., The two sharp vertical lines are the spectra of our target and a nearby star that happened to fall on the slit while the smeared out horizontal line in the middle of the image is the 11.3 $\mu$ m aromatic feature.
+ The facts that the aromatic feature is visible through the entire slit width and that it is in both spectra suggest that it comes from the background rather than from the source itself., The facts that the aromatic feature is visible through the entire slit width and that it is in both spectra suggest that it comes from the background rather than from the source itself.
+ It is also clear that the feature is much weaker at the background location (right panel) than at the source position (left panel)., It is also clear that the feature is much weaker at the background location (right panel) than at the source position (left panel).
+ We downloaded all the available 3.6-8 zm data from theSpitzer archive to see if new observations have become available since the appearance of the paper reporting the discovery of these objects (Balog 2006)., We downloaded all the available 3.6-8 $\mu$ m data from the archive to see if new observations have become available since the appearance of the paper reporting the discovery of these objects \citep{Balo06}.
+". Only in the case of NGC 2244 did we find new deeper IRAC observations (Bouwman et al.,"," Only in the case of NGC 2244 did we find new deeper IRAC observations (Bouwman et al.,"
+ Prog ID: 30726)., Prog ID: 30726).
+ For the rest of the sources we used the data from the Prog ID: 37 (PI: Fazio) and Prog ID: 58 (PI: Rieke) GTO programs., For the rest of the sources we used the data from the Prog ID: 37 (PI: Fazio) and Prog ID: 58 (PI: Rieke) GTO programs.
+ We re-mosaiced all the IRAC observations using custom IDL scripts and extracted the flux at the position of the “head” of each source using 2 pixel apertures., We re-mosaiced all the IRAC observations using custom IDL scripts and extracted the flux at the position of the “head” of each source using 2 pixel apertures.
+ We estimated 30 upper limits for the flux of the “tail” in the case of NGC 2264 and IC 1396., We estimated $\sigma$ upper limits for the flux of the “tail” in the case of NGC 2264 and IC 1396.
+ The situation for NGC 2244 was more complicated since the nearby O star affects the photometry at shorter wavelengths., The situation for NGC 2244 was more complicated since the nearby O star affects the photometry at shorter wavelengths.
+" As a result, although we detect the source at 3.6 and 4.5 jum we cannot extract its flux."," As a result, although we detect the source at 3.6 and 4.5 $\mu$ m we cannot extract its flux."
+ At 5.8 and 8.0 jum both the “head” and the “tail” are measurable in the new mosaics., At 5.8 and 8.0 $\mu$ m both the “head” and the “tail” are measurable in the new mosaics.
+ IK Tau. also known as NML Tau. is an extremely red. oxygen-rich. Mira-type variable. with a period of about 470 days (?)..," IK Tau, also known as NML Tau, is an extremely red, oxygen-rich, Mira-type variable, with a period of about 470 days \citep{Wing1973ApJ...184..873W}."
+ Its dust-driven wind produces a cool circumstellar envelope (CSE). which fosters a rich gas-phase chemistry (e.g..?)..," Its dust-driven wind produces a cool circumstellar envelope (CSE), which fosters a rich gas-phase chemistry \citep[e.g.,][]{Duari1999AandA...341L..47D}."
+ IK Tau is relatively nearby. at a distance of ~265 ppe (?)..," IK Tau is relatively nearby, at a distance of $\sim$ pc \citep{Hale1997ApJ...490..407H}."
+ Estimates of its mass-loss rates range from 3.8x1079 0) to 3x107 Msun//vr (?).., Estimates of its mass-loss rates range from $3.8 \times 10^{-6}$ \citep{Neri1998AandAS..130....1N} to $3 \times 10^{-5}$ /yr \citep{GonzalesDelgado2003AandA...411..123G}.
+" IK Tau’s proximity and relatively high mass-loss rate facilitate the observation of molecular emission lines,", IK Tau's proximity and relatively high mass-loss rate facilitate the observation of molecular emission lines.
+ Currently. a dozen different molecules and some of their isotopologs have been discovered in IK Tau. including CO. HCN. 1Ο. SiS. SO. SO». and NaCl (e.g. ??).. ," Currently, a dozen different molecules and some of their isotopologs have been discovered in IK Tau, including CO, HCN, SiO, SiS, SO, $_2$, and NaCl \citep[e.g.][]{Milam2007ApJ...668L.131M,Decin2010}. ."
+In this Letter. we report on the detection of thermal emission of water (H3O) in the envelope around IK Tau.," In this Letter, we report on the detection of thermal emission of water $_2$ O) in the envelope around IK Tau."
+ The main isotopolog (H!?O) as well as the rarer isotopologs (H!’O and H!SO) are detected for both the ortho- and para-states., The main isotopolog $_2^{16}$ O) as well as the rarer isotopologs $_2^{17}$ O and $_2^{18}$ O) are detected for both the ortho- and para-states.
+" We also present observations of high-excitation rotational transitions of ""CO. CO. SiO. ""SiO. ""SiO. HCN. and SO and demonstrate that the observed line widths characterize the acceleration region in the inner wind zone."," We also present observations of high-excitation rotational transitions of $^{12}$ CO, $^{13}$ CO, $^{28}$ SiO, $^{29}$ SiO, $^{30}$ SiO, HCN, and SO and demonstrate that the observed line widths characterize the acceleration region in the inner wind zone."
+ The HIFI instrument (?) onboard the Herschel satellite (2) offers the possibility to observe molecular fingerprints m the frequency ranges of GGHz and GGHz at à spectral resolution up to kkHz., The HIFI instrument \citep{DeGraauw2010} onboard the Herschel satellite \citep{Pilbratt2010} offers the possibility to observe molecular fingerprints in the frequency ranges of GHz and GHz at a spectral resolution up to kHz.
+ Single-point observations towards TTau were carried out with the Herschel/HIFI instrument in the dual beam switch (DBS)mode with a 3’ chop throw., Single-point observations towards Tau were carried out with the Herschel/HIFI instrument in the dual beam switch (DBS)mode with a $3^{\prime}$ chop throw.
+ The observation strategy and data-reduction are discussed in the online Appendix A and in ?.., The observation strategy and data-reduction are discussed in the online Appendix A and in \citet{Bujarrabal2010}. .
+therein).,.
+. Iu these cases the distances go well bevoud the virial radius ranging frou few hundred kye to several Moc., In these cases the distances go well beyond the virial radius ranging from few hundred kpc to several Mpc.
+ Individual field galaxies or clusters produce a sinall distortion of the backeround eal:uxdes that allows us to measure the surface mass doeusity profile of the dark matter (o5.Moeller1999:Dartelinauun&Schneider 2001).," Individual field galaxies or clusters produce a small distortion of the background galaxies that allows us to measure the surface mass density profile of the dark matter \citep[e.g.,][]{Mellier,BS01}."
+. It Is customary in the weak lensing analysis to specity a dark matter halo density profil ΣΤΟ 1nodol he projected lass profie nieasured Wih this technique., It is customary in the weak lensing analysis to specify a dark matter halo density profile to model the projected mass profile measured with this technique.
+ The NEW analytical foiiuula is often adopted aud extrapolated at large disances. bevoud Ba with poxr?.," The NFW analytical formula is often adopted and extrapolated at large distances, beyond $\Rvir$ with $\rho \propto r^{-3}$."
+ This density moc eliuayv not be accurate enoueli., This density model may not be accurate enough.
+" The motion of satellie galaxies as a test for dark iater distribution a laree radii {οOO.oZarit-al.2003:Brainerd2001:Conroyet2001) is another observatioua method. which requires cletailed heoretical predictions for outer deusity profies and iufall velocitic""S"," The motion of satellite galaxies as a test for dark matter distribution at large radii \citep[e.g.,][]{ZaritskyWhite,Zaritsky97,
+Prada2003,Brainerd04,DEEP} is another observational method, which requires detailed theoretical predictions for outer density profiles and infall velocities."
+ Iu fact. we really do not know how far the halos extend.," In fact, we really do not know how far the halos extend."
+ Indeed. the NEW. fitting forma was proposed and extesively tested to describe dark latter halos within R44.," Indeed, the NFW fitting formula was proposed and extensively tested to describe dark matter halos within $\Rvir$."
+ This is why the NEW density. fits are alwavs done within the vira radius or even well below the virial radius (< O0.5BRG4. where the halos ire expected to be viraized iu order to avoid such non-equilibrimm Huctuatious.," This is why the NFW density fits are always done within the virial radius or even well below the virial radius $<0.5\Rvir$ ), where the halos are expected to be virialized in order to avoid such non-equilibrium fluctuations."
+ In this context. it is also surpris18o hat the xedictiou of the spherical collapse model or the mean profile. which can be reasonably expected to eive good predictious at sufBicieutlv arge distances (> Ray). have not been used.," In this context, it is also surprising that the prediction of the spherical collapse model for the mean profile, which can be reasonably expected to give good predictions at sufficiently large distances $>2\Rvir$ ), have not been used."
+ Iu fact. the predictions of this models ive oulv recently been worked οι (seeBarkana2001). mt they have uot heen ested agaiust miuerical simulatious.," In fact, the predictions of this models have only recently been worked out \citep[see][]{Barkana}, but they have not been tested against numerical simulations."
+ Qur goal is to carry οἱ a detailed stidv of the density profiles im aud around collapse objects., Our goal is to carry out a detailed study of the density profiles in and around collapsed objects.
+ The paper is organized as ollows., The paper is organized as follows.
+ Iu Sectlon 2.. we eive the details of tle imnerical simulaions aud the fits of the deusity profile of clistine ealaxv-size halos.," In Section \ref{sec:num}, we give the details of the numerical simulations and the fits of the density profile of distinct galaxy-size halos."
+ Iu Section ?3.. we have stidied the shape of the deusitv aud πα] velocity profiles of isolated halos up to 2νι for differeu ILASSCN.," In Section \ref{sec:infall}, we have studied the shape of the density and infall velocity profiles of isolated halos up to $2-3\Rvir$ for different masses."
+ We show in Section 1 that for large radi the niea density profile around dark mater halos is in exccllent agreement with the predictions we have obtained via the spherical collapse model. which are soinewlhat different from those fouud by Barkana(200L).," We show in Section \ref{sec:sph} that for large radii the mean density profile around dark matter halos is in excellent agreement with the predictions we have obtained via the spherical collapse model, which are somewhat different from those found by \citet{Barkana}."
+.. Finally. in Section 5. discussions aud couclisiMis are elven.," Finally, in Section \ref{sec:fin} discussions and conclusions are given."
+ Throughout this paper the formal virial radius is the racius within which the mean matter density is equal to :MO thes the average mean matter deusitv ο οf the Universe at 2=0., Throughout this paper the formal virial radius is the radius within which the mean matter density is equal to 340 times the average mean matter density $\rho_m$ of the Universe at $z=0$.
+ The simuatious used in this paper are done using the Acaptive Refinement Tree (ART) code (ART.RravSOV.etal.1997).," The simulations used in this paper are done using the Adaptive Refinement Tree (ART) code \citep[ART,][]{Kravtsov97}."
+ The simulations were done for the standard ACDAL coszinological model with Qj=O35. Q4=T. aud fi=O.7.," The simulations were done for the standard $\LCDM$ cosmological model with $\Omega_{0}=0.3$, $\Omega_{\Lambda} = 0.7$, and $h = 0.7$ ."
+ We stuvo halos selected from four differeut simulations. which parameters are listed in Table 1..," We study halos selected from four different simulations, which parameters are listed in Table \ref{tab:tab1}."
+ The siniations cover a wide range of scales ancl have liftcrent mass and force resolutious., The simulations cover a wide range of scales and have different mass and force resolutions.
+ The simulation DBox20 has the highest resolution. but it has ouly wo galaxv-ze halos.," The simulation Box20 has the highest resolution, but it has only two galaxy-size halos."
+ We use thei. as examples for the structure of halos simulated with verv lieh resolution., We use them as examples for the structure of halos simulated with very high resolution.
+ In most of the cases we Init the analysis to well resolved halos: those should have more than 20.000 particles inside virial radius.," In most of the cases we limit the analysis to well resolved halos: those should have more than $\approx$ 20,000 particles inside virial radius."
+ The simulation Box120 jas the lareer vohuue. mut its mass resolution allows us to use ouly halos with masses lareer than 1015iW.," The simulation Box120 has the larger volume, but its mass resolution allows us to use only halos with masses larger than $10^{13}\Msunh$."
+ Most of he analysis of galaxy-size iios is done using sinmlatious Boxs0S aud BoxsKC Tn the case of the sinulation BoxsIG. the whole 5001Mpc vohuue was resolved with equalbaiuass particles., Most of the analysis of galaxy-size halos is done using simulations Box80S and Box80G. In the case of the simulation Box80G the whole $80\Mpch$ volume was resolved with equal-mass particles.
+ There were about 150.000 halos in he simulation.," There were about 180,000 halos in the simulation."
+ The sunuulation Boxs0S was done using particles with different. masses., The simulation Box80S was done using particles with different masses.
+" Oilv a sinall fraction a 10hiN peradius region ofthe box was resolved with snalinass particles,", Only a small fraction – a $10\Mpch$ radius region – of the box was resolved with small-mass particles.
+ The high resolution reglon has au average deitv about equal to the uen deitv of the Universe., The high resolution region has an average density about equal to the mean density of the Universe.
+ It was chosen iu such acwav that the halo mass function iu the region is representative for the typical region of this size., It was chosen in such a way that the halo mass function in the region is representative for the typical region of this size.
+ For example. t1e region does not have a massive cluster.," For example, the region does not have a massive cluster."
+ The uost massive halo in the region has lass 2.6vlEphTALL., The most massive halo in the region has mass $2.6\times 10^{13}\Msunh$.
+ There are 5 halos with mass laveer than Lote) TAL..., There are 5 halos with mass larger than $10^{13}\Msunh$ .
+ Altevecther. there are about GU.JOO halos in this sinuation.," Altogether, there are about 60,000 halos in this simulation."
+ The density profile of cac1i lado. which we study. is fitby the ao)proxination eiven iu eq.( 2)).," The density profile of each halo, which we study, is fitby the approximation given in eq.( \ref{eq:Sersic}) )."
+ Each fit provides tie concentration C aid the Sérrsic, Each fit provides the concentration $C$ and the Sérrsic
+"star-disk system at t=1.2 hours (upper panel; the flaring loop is marked in red), and a schematic view of the system during the evolution of the flaring loop (lower panel).","star-disk system at $t=1.2$ hours (upper panel; the flaring loop is marked in red), and a schematic view of the system during the evolution of the flaring loop (lower panel)."
+" Due to the efficient thermal conduction and radiation, the plasma begins to cool immediately after the heat pulse is over, and the maximum temperature rapidly decreases to ~50 MK in =1 hour."," Due to the efficient thermal conduction and radiation, the plasma begins to cool immediately after the heat pulse is over, and the maximum temperature rapidly decreases to $\sim 50$ MK in $\approx 1$ hour."
+" The disk evaporation is fed by thermal conduction from the outer hot plasma, even after the end of the heat pulse."," The disk evaporation is fed by thermal conduction from the outer hot plasma, even after the end of the heat pulse."
+ Figure 2 also shows that the overheating of the disk surface at the loop footpoint makes a significant amount of material expand and be ejected in the magnetosphere., Figure \ref{flare_evol} also shows that the overheating of the disk surface at the loop footpoint makes a significant amount of material expand and be ejected in the magnetosphere.
+" A small fraction of this evaporated disk material streams along the loop accelerated toward the star by its gravity, fills the whole tube in ~10 hours, and the density keeps increasing throughout the loop up to values ranging between 10° and 1010 cm~? at its apex."," A small fraction of this evaporated disk material streams along the loop accelerated toward the star by its gravity, fills the whole tube in $\approx 10$ hours, and the density keeps increasing throughout the loop up to values ranging between $10^{9}$ and $10^{10}$ $^{-3}$ at its apex."
+" On the other hand, most of the evaporated disk material is not efficiently confined by the magnetic field and channeled into the hot loop but is rather ejected away from the central star in the outer stellar corona, carrying away mass and angular momentum (see Fig. 2))."," On the other hand, most of the evaporated disk material is not efficiently confined by the magnetic field and channeled into the hot loop but is rather ejected away from the central star in the outer stellar corona, carrying away mass and angular momentum (see Fig. \ref{flare_evol}) )."
+" Due to the high values of 8 there, the magnetic field lines are dragged away."," Due to the high values of $\beta$ there, the magnetic field lines are dragged away."
+ The outflowing plasma is supersonic and its speed is of the order of the local Alfvénn speed km s!., The outflowing plasma is supersonic and its speed is of the order of the local Alfvénn speed $u\rs{A} \approx 300$ km $^{-1}$.
+ We note that our results on the dynamics of the outflowing plasma are similar to those found with an MHD model proposed by ? to describe hard X-ray flares in protostars observed by the ASCA satellite., We note that our results on the dynamics of the outflowing plasma are similar to those found with an MHD model proposed by \cite{1996ApJ...468L..37H} to describe hard X-ray flares in protostars observed by the ASCA satellite.
+" From the model results, we synthesized the X-ray emission originating from the flare, applying a methodology analogous to that described in the literature in the context of the study of novae and supernova remnants (?7))."," From the model results, we synthesized the X-ray emission originating from the flare, applying a methodology analogous to that described in the literature in the context of the study of novae and supernova remnants \citealt{orlando2, 2009A&A...493.1049O}) )."
+" In particular, we first calculate the emission measure in the j-th domain cell as em;=ni (where nj is the hydrogen number density in the cell,Vj V; is the cell volume, and we assume fully ionized plasma)."," In particular, we first calculate the emission measure in the $j$ -th domain cell as ${\rm em}\rs{j} = n\rs{Hj}^2 V\rs{j}$ (where $n\rs{Hj}^2$ is the hydrogen number density in the cell, $V\rs{j}$ is the cell volume, and we assume fully ionized plasma)."
+" From the values of emission measure and temperature in the cell, we synthesize the corresponding X-ray spectrum, using the PINTofALE spectral code (?)) with the APED V1.3 atomic line database, and assuming the same metal abundances of the simulation, namely 0.5 of the solar values, as deduced from X-ray observations of CTTSs (7))."," From the values of emission measure and temperature in the cell, we synthesize the corresponding X-ray spectrum, using the PINTofALE spectral code \citealt{Kashyap2000BASI}) ) with the APED V1.3 atomic line database, and assuming the same metal abundances of the simulation, namely 0.5 of the solar values, as deduced from X-ray observations of CTTSs \citealt{Telleschi2007A&Ab}) )."
+" We integrate the X-ray spectra from the cells in the whole spatial domain and, then, derive the X-ray luminosity by integrating the spectrum in the [0.6—12] keV band."," We integrate the X-ray spectra from the cells in the whole spatial domain and, then, derive the X-ray luminosity by integrating the spectrum in the $[0.6-12]$ keV band."
+" Figure 4 shows the evolution of the maximum temperature and emission measure of the flaring loop, and its X-ray lightcurve."," Figure \ref{fig2_sup} shows the evolution of the maximum temperature and emission measure of the flaring loop, and its X-ray lightcurve."
+" The peak temperature of the flare is reached very soon at t£α100 s, the emission measure peaks later at tzz1000 s. The X-ray luminosity evolves as the emission measure, reaching a peak value of Lx& ergs s! at tzz500 s, namely after the end of the heat pulse."," The peak temperature of the flare is reached very soon at $t\approx 100$ s, the emission measure peaks later at $t\approx 1000$ s. The X-ray luminosity evolves as the emission measure, reaching a peak value of $L\rs{X} \approx 6.5\times 10^{32}$ ergs $^{-1}$ at $t\approx
+500$ s, namely after the end of the heat pulse."
+ Then the luminosity (and the emission, Then the luminosity (and the emission
+We stress again that we use a mix of dust amounts in our standard analysis to mateh the UV-continuum slope distribution that is measured from the data.,We stress again that we use a mix of dust amounts in our standard analysis to match the UV-continuum slope distribution that is measured from the data.
+ Especially for the a- and g-dropouts this puts strong constraints on the combination of the age and the amount of dust in the model template. so that we can reduce the systematic error to à minimum.," Especially for the $u$ -, and $g$ -dropouts this puts strong constraints on the combination of the age and the amount of dust in the model template, so that we can reduce the systematic error to a minimum."
+ To justify the assumptions we make regarding the shapes of our simulated sources we also estimated the systematic error on the LF due to this component., To justify the assumptions we make regarding the shapes of our simulated sources we also estimated the systematic error on the LF due to this component.
+ Because we expect similar results in the three dropout samples. we only run these simulations for the e-dropouts.," Because we expect similar results in the three dropout samples, we only run these simulations for the $g$ -dropouts."
+ We inject sources that are 0.057 larger than the measured position-dependent PSF. and compare values of Moffat parameter 8=3.0 and 6=5.0 with our fiducial 6=4.0 parameter.," We inject sources that are 0.05” larger than the measured position-dependent PSF, and compare values of Moffat parameter $\beta=3.0$ and $\beta=5.0$ with our fiducial $\beta=4.0$ parameter."
+ We find the following: Furthermore we find that. based on the different template spectra. the estimated UV luminosity density varies.," We find the following: Furthermore we find that, based on the different template spectra, the estimated UV luminosity density varies."
+" The ranges are. upon integration down to L< 0.3L°,. in units |1079eres""!Hz! Mpe7]. 1.30.«puy<1.98 for the e. 0.92«puy<1.15 for the g-. and 0.80<puy0.92 for the r-dropouts."," The ranges are, upon integration down to $L<0.3L^{*}_{z=3}$ , in units $\rm{[10^{26} erg\,s^{-1}\,Hz^{-1}\,Mpc^{-3}}]$ , $1.30< \rho_{\rm{UV}} <1.98$ for the $u$ -, $0.92< \rho_{\rm{UV}} <1.15$ for the $g$ -, and $0.80< \rho_{\rm{UV}} <0.92$ for the $r$ -dropouts."
+ In order to illustrate the effect that the Eddington bias can have on LF estimations we repeat our analysis with slightly changed pn.z)," In order to illustrate the effect that the Eddington bias can have on LF estimations we repeat our analysis with slightly changed $p(m,z)$."
+ In the completeness simulations we bin the recovered sources by their intrinsic magnitude instead of their recovered magnitude., In the completeness simulations we bin the recovered sources by their intrinsic magnitude instead of their recovered magnitude.
+ Doing so leaves the Eddington bias uncorrected., Doing so leaves the Eddington bias uncorrected.
+" We find that the Schechter fit to the resulting LFs are not as good as for our standard analysis. especially for the g-. and r-dropouts. where we find that Vij,/dof would be 1.0 and 3.0. respectively."," We find that the Schechter fit to the resulting LFs are not as good as for our standard analysis, especially for the $g$ -, and $r$ -dropouts, where we find that $\chi_{\rm{min}}^{2}/\rm{dof}$ would be 1.0 and 3.0, respectively."
+ For the u-dropouts the best-fitting parameters are not changed significantly. but for the g-. and r- the faint-end slope steepens while the normalisation of the Schechter function decreases.," For the $u$ -dropouts the best-fitting parameters are not changed significantly, but for the $g$ -, and $r$ -dropouts the faint-end slope steepens while the normalisation of the Schechter function decreases."
+ Note that we did not account for the presence of Lyman- a emission in our simulations. although this line clearly contributes to broadband fluxes.," Note that we did not account for the presence of Lyman- $ \alpha $ emission in our simulations, although this line clearly contributes to broadband fluxes."
+ ? measured the contribution of this line in the spectra of z~3 LBGs and concluded that only ~25% of the sample showed significant Ly a emission such that EW(Lya) >20A. see also ?..," \citet{shapley2003} measured the contribution of this line in the spectra of $ z \sim 3 $ LBGs and concluded that only $ \sim 25 \% $ of the sample showed significant Ly $ \alpha $ emission such that $\alpha$ ) $> 20 \AA$, see also \citet{reddy2008}."
+ Although the relative contribution of this line is thought to increase towards higher redshift and fainter continuum luminosity. this has not been quantified yet.," Although the relative contribution of this line is thought to increase towards higher redshift and fainter continuum luminosity, this has not been quantified yet."
+ We will therefore describe possible biasses due to the presence of this line in the measurement of the LF only qualitatively., We will therefore describe possible biasses due to the presence of this line in the measurement of the LF only qualitatively.
+ If the line appears in emission. it contributes to the flux in the CFHT g-band for redshifts of about 2.5<z3.5. in the r-band for redshifts of about 3.7<z4.6. in the i-band for redshifts of about 4.8<z5.9. and in the z-band for redshifts of about z>5.8.," If the line appears in emission, it contributes to the flux in the CFHT $g$ -band for redshifts of about $2.5<z<3.5$, in the $r$ -band for redshifts of about $3.7<z<4.6$, in the $i$ -band for redshifts of about $4.8<z<5.9$, and in the $z$ -band for redshifts of about $z>5.8$."
+ Following the redshift distributions from ?.. we expect that the line predominantly contributes to the middle band in a two colour selection for the w-. and r- dropout samples. moving the source to the upper left in Fig. 10..," Following the redshift distributions from \citet{hildebrandt09a}, we expect that the line predominantly contributes to the middle band in a two colour selection for the $u$ -, and $r$ - dropout samples, moving the source to the upper left in Fig. \ref{fig:colortrackall}."
+ This effect causes the effective volumes to rise. thereby lowering the LF measurement.," This effect causes the effective volumes to rise, thereby lowering the LF measurement."
+ If the line indeed gets stronger at faint magnitudes. this would bias the LF measurement and results in a shallower alpha.," If the line indeed gets stronger at faint magnitudes, this would bias the LF measurement and results in a shallower alpha."
+ In our g-dropout sample the line is expected to contribute to the flux in both the ;- and /-band. depending on the redshift of the particular source.," In our $g$ -dropout sample the line is expected to contribute to the flux in both the $r$ - and $i$ -band, depending on the redshift of the particular source."
+ In the case where the line falls in the r-band. the effective volumes rise and the LF we measured is too high.," In the case where the line falls in the $r$ -band, the effective volumes rise and the LF we measured is too high."
+ For the redshifts where the line falls in the i-band. the sources would move to the right in Fig.," For the redshifts where the line falls in the $i$ -band, the sources would move to the right in Fig."
+ 10. so that the effective volumes are decreased and the LF increases., \ref{fig:colortrackall} so that the effective volumes are decreased and the LF increases.
+ To estimate which effect prevails. the redshift distribution. needs to be measured spectroscopically.," To estimate which effect prevails, the redshift distribution needs to be measured spectroscopically."
+ Some objects show Ly-c in absorption rather than emission. which would give the opposite effect.," Some objects show $\alpha$ in absorption rather than emission, which would give the opposite effect."
+ Note that ? model the contribution of Lye in the measurement of their LF. by using a simple model where of the z~4—5 LBGs have EW(Ly « )= 50 Á. independent of the continuum luminosity.," Note that \citet{bouwens07} model the contribution of $\alpha$ in the measurement of their LF, by using a simple model where of the $z \sim 4-5$ LBGs have EW(Ly $\alpha$ )= 50 $\AA$, independent of the continuum luminosity."
+ They find that this affects the normalisation of the LF by only ~10%., They find that this affects the normalisation of the LF by only $\sim 10\%$.
+ However. we stress again that the measure of a might be biased if the strength of the line depends on the continuum luminosity.," However, we stress again that the measure of $\alpha$ might be biased if the strength of the line depends on the continuum luminosity."
+ Although we can measure the UV luminosity density with great accuracy. an estimate of the SFR Density depends sensitively on the dust extinction. correction.," Although we can measure the UV luminosity density with great accuracy, an estimate of the SFR Density depends sensitively on the dust extinction correction."
+ Using. the prescription from ?.. we find that the SFR Density shows a significant increase by a factor of 4-5 between z~5 and z~ο ," Using the prescription from \citet{bouwens09}, we find that the SFR Density shows a significant increase by a factor of 4-5 between $z\sim 5$ and $z\sim 3$."
+We find a strong increase in & between z~5 and z~3 by a factor of 2.5. which is a robust result.," We find a strong increase in $\phi^{*}$ between $z\sim 5$ and $z\sim 3$ by a factor of 2.5, which is a robust result."
+ The characteristic magnitude. M. however is not very well constrained by our data set.," The characteristic magnitude, $M_{\rm{UV}}^{*}$, however is not very well constrained by our data set."
+ This is due to uncertainties in the redshift distributions of the source galaxies. as there are no spectroscopic data available.," This is due to uncertainties in the redshift distributions of the source galaxies, as there are no spectroscopic data available."
+" Our data are consistent with a non-evolving Mj, ( —20.9) between z~5 and =~3.", Our data are consistent with a non-evolving $M_{\rm{UV}}^{*}$ $\sim -20.9$ ) between $z\sim 5$ and $z\sim 3$.
+ Our data do not show an evolution of the faint-end slope. and indicate w~—1.6 in this epoch.," Our data do not show an evolution of the faint-end slope, and indicate $\alpha \sim -1.6$ in this epoch."
+ Before wecompare ourresults from each of the samples with Schechter parameters reported from previous determinations in the literature. there are a few things that should be noted.," Before wecompare ourresults from each of the samples with Schechter parameters reported from previous determinations in the literature, there are a few things that should be noted."
+ We will compare results at redshifts of around 3. 4 and 5.," We will compare results at redshifts of $around$ 3, 4 and 5."
+ The LBGs are generally selected from different surveys and filter sets. and therefore intrinsically slightly different galaxies may be selected. studied and compared.," The LBGs are generally selected from different surveys and filter sets, and therefore intrinsically slightly different galaxies may be selected, studied and compared."
+"cm, the density becomes higher and reaches the value of Nchm=10 particles/cm? (the complete profilesof n,j,, and R as functions of the radial coordinate are reported in Fig. 4)).","cm, the density becomes higher and reaches the value of $n_{cbm}=10$ $^3$ (the complete profilesof $n_{cbm}$ and $\mathcal{R}$ as functions of the radial coordinate are reported in Fig. \ref{f4}) )."
+ This is compatible with the observations., This is compatible with the observations.
+" The observed X-ray and optical afterglow of GRB 071227 is indeed superimposed on the plane of the host galaxy, at (15.0+2.2) kpc from its center (D'Avanzoetal.2009)."," The observed X-ray and optical afterglow of GRB 071227 is indeed superimposed on the plane of the host galaxy, at $(15.0\pm2.2)$ kpc from its center \citep{D09}."
+". An interesting possibility observed by D. Arnett (private communication) is that this very low density ""cavity"" could be formed in the coalescing phase of a binary formed by a neutron star and a white dwarf.", An interesting possibility observed by D. Arnett (private communication) is that this very low density “cavity” could be formed in the coalescing phase of a binary formed by a neutron star and a white dwarf.
+ Accurate studies on compact object mergers have shown the distribution of merger locations for different host galaxies (Belczynskietal.2006;Berger 2010).," Accurate studies on compact object mergers have shown the distribution of merger locations for different host galaxies \citep{ba06,b10}."
+". In starburst galaxies, most of the mergers are expected to be found within hosts, while in elliptical galaxies a substantial fraction of mergers take place outside hosts."," In starburst galaxies, most of the mergers are expected to be found within hosts, while in elliptical galaxies a substantial fraction of mergers take place outside hosts."
+" Spiral galaxies, hosting both young and old stellar populations, represent the intermediate case between the preceding two."," Spiral galaxies, hosting both young and old stellar populations, represent the intermediate case between the preceding two."
+ This result is therefore compatible with our hypothesis about the binary nature of the progenitor of GRB 071227., This result is therefore compatible with our hypothesis about the binary nature of the progenitor of GRB 071227.
+" Although they did not consider the case of binary systems formed by a neutron star and a white dwarf, the progenitor of GRB 071227 would fit into the tight binary scenario described by Belczynskietal. (2006)."," Although they did not consider the case of binary systems formed by a neutron star and a white dwarf, the progenitor of GRB 071227 would fit into the tight binary scenario described by \citet{ba06}."
+. These results clearly imply that GRB 071227 is another example of a disguised burst., These results clearly imply that GRB 071227 is another example of a disguised burst.
+" One of the most effective tools for discriminating between ""short"" and ""long"" bursts, and possibly clarifying the interpretation of these different classes of events, is the Amati relation (Amatietal.2002;Amati2006;2007,2009)."," One of the most effective tools for discriminating between “short” and “long” bursts, and possibly clarifying the interpretation of these different classes of events, is the Amati relation \citep{A02,A06b,a07,A09}."
+". This empirical spectrum-energy correlation states that the isotropic-equivalent radiated energy of the prompt emission Ej, is correlated with the cosmological rest-frame vF, spectrum peak energy Ej; Ey;c(Ειω), with a~0.5 (Amatietal.2002)."," This empirical spectrum-energy correlation states that the isotropic-equivalent radiated energy of the prompt emission $E_{iso}$ is correlated with the cosmological rest-frame $\nu F_{\nu}$ spectrum peak energy $E_{p,i}$: $E_{p,i}\propto (E_{iso})^{a}$, with $a \approx 0.5$ \citep{A02}."
+". The existence of the Amati relation, discovered when analyzing the long duration bursts, has been confirmed by studying a sample of GRBs observed bySwift,, intense or soft, with available measurements of redshift and spectral parameters, and also by andKonus/WIND.."," The existence of the Amati relation, discovered when analyzing the long duration bursts, has been confirmed by studying a sample of GRBs observed by, intense or soft, with available measurements of redshift and spectral parameters, and also by and."
+" When the “afterglow revolution” allowed the redshift estimation of also some short GRBs, it was found that these bursts are inconsistent with the Amati relation, holding instead for long GRBs (Amati2006,2010;Antonellietal.2009)."," When the “afterglow revolution” allowed the redshift estimation of also some short GRBs, it was found that these bursts are inconsistent with the Amati relation, holding instead for long GRBs \citep{A06b,a10,aa09}."
+". The most recent updating of the correlation (95 GRBs with the data available up to April 2009) also includes two high-energetic events detected by (GRB 090323 and GRB 080916C), which are fully consistent with the E,;—E;s; relation (Amatietal.2009)."," The most recent updating of the correlation (95 GRBs with the data available up to April 2009) also includes two high-energetic events detected by (GRB 090323 and GRB 080916C), which are fully consistent with the $E_{p,i}$ $E_{iso}$ relation \citep{A09}."
+". We also note that the short event GRB 090510, observed by the same satellite, is an outlier of the relation, as expected for short bursts."," We also note that the short event GRB 090510, observed by the same satellite, is an outlier of the relation, as expected for short bursts."
+ This dichotomy finds a natural explanation within the fireshell model., This dichotomy finds a natural explanation within the fireshell model.
+" As recalled in Sect. ??,,"," As recalled in Sect. \ref{model},"
+" within this theoretical framework the prompt emission of long GRBs is dominated by the peak of the extended afterglow, while the one of the short GRBs is dominated by the P-GRB."," within this theoretical framework the prompt emission of long GRBs is dominated by the peak of the extended afterglow, while the one of the short GRBs is dominated by the P-GRB."
+" Only the extended afterglow emission follows the Amati relation (seeGuidaetal.2008,for details)..", Only the extended afterglow emission follows the Amati relation \citep[see][for details]{Guida}.
+" Therefore, all GRBs in which the P-GRB provides a negligible contribution to the prompt emission (namely the long ones, where the P-GRB is at most a small precursor) fulfill the Amati relation, while all GRBs in which the extended afterglow provides a negligible contribution to the prompt emission (namely the short ones) do not."," Therefore, all GRBs in which the P-GRB provides a negligible contribution to the prompt emission (namely the long ones, where the P-GRB is at most a small precursor) fulfill the Amati relation, while all GRBs in which the extended afterglow provides a negligible contribution to the prompt emission (namely the short ones) do not."
+" Apart from this general feature, there are peculiar cases."," Apart from this general feature, there are peculiar cases."
+" One of these is GRB 050724, a short burst followed by a long, softer tail in the y-ray energy band."," One of these is GRB 050724, a short burst followed by a long, softer tail in the $\gamma$ -ray energy band."
+" While the first short emission is inconsistent with the Amati relation, the soft tail is again consistent with it."," While the first short emission is inconsistent with the Amati relation, the soft tail is again consistent with it."
+" Another one, the intriguing case of GRB 060614, a short burst within the fireshell scenario (Caitoetal.2009),, shows similar behavior: the first hard episode does not fulfill the Amati relation, while the whole event is fully consistent with it (Amatietal.2007)."," Another one, the intriguing case of GRB 060614, a short burst within the fireshell scenario \citep{c09}, shows similar behavior: the first hard episode does not fulfill the Amati relation, while the whole event is fully consistent with it \citep{a07}."
+. This is again fully consistent with the predictions of the fireshell model for the disguised GRB class., This is again fully consistent with the predictions of the fireshell model for the disguised GRB class.
+" Since the first short and hard episode is the P-GRB and the prolonged softer tail is the peak of the extended afterglow, the Amati relation must be fulfilled only when the softer tail is considered alone, or considered together with the first episode if this episode provides a negligible contribution (seeBernardinietal.2008,for details). "," Since the first short and hard episode is the P-GRB and the prolonged softer tail is the peak of the extended afterglow, the Amati relation must be fulfilled only when the softer tail is considered alone, or considered together with the first episode if this episode provides a negligible contribution \citep[see][for details]{GraziaIC4}. ."
+"Thus, to discriminate its nature, we studied the position of GRB 071227 in the E,;-Ejso plane."," Thus, to discriminate its nature, we studied the position of GRB 071227 in the $E_{p,i}$ $E_{iso}$ plane."
+" At the observed redshift, we assumed a “flat A-CDM model” with Hp=70Km/s/Mpc and"," At the observed redshift, we assumed a “flat $\Lambda$ -CDM model” with $H_{0}=70 Km/s/Mpc$ and"
+the lower eccentricity aud the orbital modulation of the uupulsed flux is ~1054.,the lower eccentricity and the orbital modulation of the unpulsed flux is $\sim 10\%$.
+" The wind Iuniuositv of pulsar D cau be estimated as Although the ratio of the wind huninositics E,Ep Is somewhat smaller than that inferred iuJ07237-3039.. it is still EpxxE, and the enerev deusitv of the wind of A exceeds that of D by a factor 10 at D lieht evliuder."," The wind luminosity of pulsar B can be estimated as Although the ratio of the wind luminosities $\dot E_{A}/\dot E_{B}$ is somewhat smaller than that inferred in, it is still $\dot E_{B}\ll\dot E_{A}$ and the energy density of the wind of A exceeds that of B by a factor $\sim 10$ at B light cylinder."
+ This ensures consistency with our starting assuuiption about the shock forming as the result of the iuteraction of the wind of A with the magnetosphere of D. The scenario outlined in the previous section shows that if one chooses for the model parameters values close to those introduced above. the main observational features of ccau be successfully reproduced.," This ensures consistency with our starting assumption about the shock forming as the result of the interaction of the wind of A with the magnetosphere of B. The scenario outlined in the previous section shows that if one chooses for the model parameters values close to those introduced above, the main observational features of can be successfully reproduced."
+ Although the general picture we propose resenibles that ofJ0737-3039.. the two binary svstenis differ iu more than oue respect.," Although the general picture we propose resembles that of, the two binary systems differ in more than one respect."
+ Iu particular. the shock radiative efficiency jieeds to be higher in 5-3009.. ~107 as conrpared to ~10 baa (TTOL.," In particular, the shock radiative efficiency needs to be higher in , $\sim 10^{-3}$ as compared to $\sim 10^{-4}$ in (TT04)."
+" Lower values of ©. in fact. result iu sinaller periods for pulsar D: this xoduces a steep increase of Ej which soon would become E, at Ds light cvlinder radius."," Lower values of $\xi$, in fact, result in smaller periods for pulsar B: this produces a steep increase of $\dot E_{B}$ which soon would become $\sim \dot E_{A}$ at B's light cylinder radius."
+ Albeit coliereuce seems indeed to be required iu both svsteimis. no detailed mocel or coherent radio emission has been put forward so far.," Albeit coherence seems indeed to be required in both systems, no detailed model for coherent radio emission has been put forward so far."
+ In TTOL some possibilities were suggested but the very 110dcl of coliereut radio emission should be revisited possibly iu he elt of most recent studies on this topic (c.g. Iujpes 2001: Zhaug&Loeb 2001))., In TT04 some possibilities were suggested but the very model of coherent radio emission should be revisited possibly in the light of most recent studies on this topic (e.g. \citealt{fuku04}; \citealt{zl04}) ).
+ Ihuuanetal.(2005). reported that wwas not detected in two ~6 hr images obtained J1L7iu 19961998 aud iu ~1 hr observations iu 20022005., \cite{hym05} reported that was not detected in two $\sim 6$ hr images obtained in 1996--1998 and in $\sim 1$ hr observations in 2002–2003.
+ While i0 concluxive evidence can be drawn frou the shorter observations which have duration less than the orbital oxjod. the lack of activity in 19961998 might pose a cifieultv to the NS| binary model. as the same authors state.," While no conclusive evidence can be drawn from the shorter observations which have duration less than the orbital period, the lack of activity in 1996–1998 might pose a difficulty to the NS+NS binary model, as the same authors state."
+" Tere we only mention that an ""on""off state nav be achieved if star À precesses.", Here we only mention that an “on”–“off” state may be achieved if star A precesses.
+" For yy—01. he wind beam is ~TO"" wide and even a moderate secession anele can prevent the beam from A from intercepting the maguetosphere of D. which subteuds au augle =107."," For $\chi_A\sim 0.1$, the wind beam is $\sim 70^\circ$ wide and even a moderate precession angle can prevent the beam from A from intercepting the magnetosphere of B which subtends an angle $\lesssim 10^\circ$."
+ Tuterestingly. the characteristic time 1/Q~Πορδα20c?)s (Damour&Ruffini 1971)) for changing the system ecometry due to ecodetic precession of the spin axis of pulsar A is ~3 yr. assunuines equal mnasses (AZ= 1.237.) for the two stars.," Interestingly, the characteristic time $1/\Omega\sim 116 
+P^{5/3}M^{-2/3}(1-e^2)\, {\rm s}$ \citealt{daruf74}) ) for changing the system geometry due to geodetic precession of the spin axis of pulsar A is $\sim 3$ yr, assuming equal masses $M=1.3 M_\odot$ ) for the two stars."
+ The range of Pp aud Bp(Rp) for which our model lolds implies that neutron star D max be an active radio-pulsar., The range of $P_{B}$ and $B_{B}(R_B)$ for which our model holds implies that neutron star B may be an active radio-pulsar.
+ The required spindown luminosity of neutron star A (see eq. [5]]), The required spindown luminosity of neutron star A (see eq. \ref{edota}] ])
+ is also compatible with it being a canonical radio-pulsar with By~10172.ote? Gand Py~0.1 sora nuldly recveled pulsar with By~5<10?«1029 C aud Py~1020 mx.," is also compatible with it being a canonical radio-pulsar with $B_{A}\sim 10^{12}-10^{13}$ G and $P_{A}\sim 0.1-1$ s or a mildly recycled pulsar with $B_A\sim 5\times
+10^{9}-5\times 10^{10}$ G and $P_{A}\sim 10-20$ ms."
+" The latter possibility fits better with the standard scenario for the formation of a double ucutrou star binary (c.g, vandenTeuvel&deLoore1973)). which predicts that the svstem should iuclude a relatively voung ueutron star (the one resulting from the second superuova explosion) and a reevcled pulsar (the first born ucutrou star) spun up to tens of millisecond period by the short phase of mass transfer from the evolving conrpauion."," The latter possibility fits better with the standard scenario for the formation of a double neutron star binary (e.g. \citealt{vdhdl73}) ), which predicts that the system should include a relatively young neutron star (the one resulting from the second supernova explosion) and a recycled pulsar (the first born neutron star), spun up to tens of millisecond period by the short phase of mass transfer from the evolving companion."
+ Detectability of the pulsed signal from these two sources depends on the radio luminosity aud favorable orieutatiou of their beams., Detectability of the pulsed signal from these two sources depends on the radio luminosity and favorable orientation of their beams.
+" Caven the position in galactic coordinates (=358.97, hb=0.5%) and for a distance in the range 6 kpe. the available models for the distribution of the ionized couponcuts in the iuterstellar medium (Taylor&CordesL993: Cordes&Lazio2001)) inclicate that scattering would prevent detection of any pulsating radio signal with P shorter than few secouds at the frequency of 330 MIIz."," Given the position in galactic coordinates $l=358.9^{\circ}$, $b=-0.5^{\circ}$ ) and for a distance in the range 6--12 kpc, the available models for the distribution of the ionized components in the interstellar medium \citealt{tc93}; \citealt{lc01}) ) indicate that scattering would prevent detection of any pulsating radio signal with $P$ shorter than few seconds at the frequency of 330 MHz."
+ Towever. observations at frequencies higher than 1 CIIz would preserve the possibility of discovering a c0.3 8 pulsar. while observing at frequency Z2 GIIz is required for detectiug a LO 1s pulsar.," However, observations at frequencies higher than 1 GHz would preserve the possibility of discovering a $\sim 0.3$ s pulsar, while observing at frequency $\gtrsim 2$ GHz is required for detecting a $\sim 10$ ms pulsar."
+ An Nav Iuuinosity coincident with the shock of Lx~--107ores| shouk be present. assuming that ~1054 of the wind huuinosiv intercepted by the imagnetosphere is converted into N-ravs.," An X-ray luminosity coincident with the shock of $L_{X}\approx 10^{32}\
+{\rm erg\, s}^{-1}$ should be present, assuming that $\sim 10\%$ of the wind luminosity intercepted by the magnetosphere is converted into X-rays."
+ However. the expected flix is extienielv low. =10!ctss| for XAMALCEDIC and Chandra-ACIS.. asstning a power-law spectrum with index 2 and Ny~L072cm7.," However, the expected flux is extremely low, $\approx 10^{-4}\ \rm{cts\, s}^{-1}$ for -EPIC and -ACIS, assuming a power-law spectrum with index 2 and $N_{H}\sim 10^{22}\, \rm{cm}^{-2}$."
+ We do not expect substantial enüssion iu other euergv buds. as in the case ofJOT37-3039.. although the vouusger NS may be associated with a supernova remmnant.," We do not expect substantial emission in other energy bands, as in the case of, although the younger NS may be associated with a supernova remnant."
+ ludirect support iu favor of a double NS system for ccould come from further radio observations ofJ0737-3039., Indirect support in favor of a double NS system for could come from further radio observations of.
+. Iu fact. evidence that the continuous source js unresolved at the arcsec level would rule out the transparent scenario CETOD). leaving cohercut cluission from the shock as the only viable option.," In fact, evidence that the continuous source is unresolved at the arcsec level would rule out the transparent scenario (TT04), leaving coherent emission from the shock as the only viable option."
+ If the picture of pproposed here will be coufirmmed. one should consider that other simular systenis nay exist iu theGalaxy.," If the picture of proposed here will be confirmed, one should consider that other similar systems may exist in theGalaxy."
+ Their search among radio transicuts with programs analogous to that of Iun (2005).. may be of the paramount mirportance for estimating the coalescence rate ou NS| systems. a vital issue iu connection with the response of now generation eravitational wave detectors like VIRGO and LIGO.," Their search among radio transients with programs analogous to that of \cite{hym05}, , may be of the paramount importance for estimating the coalescence rate on NS+NS systems, a vital issue in connection with the response of new generation gravitational wave detectors like VIRGO and LIGO."
+estimate.,estimate.
+ Thus. (he spacings will help us refine fature stellar interior models for this star.," Thus, the spacings will help us refine future stellar interior models for this star."
+ The organization of (he paper goes as follows: In §2.1 we discuss the specilics of our models and the physics implemented for (heir ecaleulations., The organization of the paper goes as follows: In 2.1 we discuss the specifics of our models and the physics implemented for their calculations.
+ Then in 82.2 we outline how we eo on lo use observational data to help constraàn our models through our HRD analysis., Then in 2.2 we outline how we go on to use observational data to help constrain our models through our HRD analysis.
+ 83 describes the calculation of both huge and small p-mocle spacings aud how thev can assist in consiraining the parameters of 51 Peg., 3 describes the calculation of both large and small -mode spacings and how they can assist in constraining the parameters of 51 Peg.
+ Lastly. in 84. we summarize our analvsis. which is necessarilv exploratory in (he absence of seismic observations.," Lastly, in 4, we summarize our analysis, which is necessarily exploratory in the absence of seismic observations."
+ Some shortcomings of our models which will have to be addressed when precise frequency measurements become available are discussed., Some shortcomings of our models which will have to be addressed when precise frequency measurements become available are discussed.
+ Finally. we point out the need to devise more sophisticated approaches to analyze the hieh quality data expected from the space missions.," Finally, we point out the need to devise more sophisticated approaches to analyze the high quality data expected from the space missions."
+ All evolutionary (racks were computed using (he Yale stellar evolution eode (YREC) (Guenther et al., All evolutionary tracks were computed using the Yale stellar evolution code (YREC) (Guenther et al.
+ 1992)., 1992).
+ The initial zero age main sequence (ZAMS) model used for 51 Peg was created from pre-main sequence evolution calculations., The initial zero age main sequence (ZAMS) model used for 51 Peg was created from pre-main sequence evolution calculations.
+ Post-main sequence models of various compositions were then constructed by first rescaling the composition of the ZAMS models., Post-main sequence models of various compositions were then constructed by first rescaling the composition of the ZAMS models.
+ In constructing the models we made use of the enrrent OPAL equation of state tables (Rogers οἱ al.1996)., In constructing the models we made use of the current OPAL equation of state tables (Rogers et al.1996).
+ OPAL tables (Iglesias Rogers 1996) were also used to caleulate interior opacities. while the surface ancl atmosphere opacities were taken from the work of Alexander and Ferguson (1994).," OPAL tables (Iglesias Rogers 1996) were also used to calculate interior opacities, while the surface and atmosphere opacities were taken from the work of Alexander and Ferguson (1994)."
+ Due to the lack of better information. we assume a solar mixture ol heavy elements.," Due to the lack of better information, we assume a solar mixture of heavy elements."
+ We accounted [or the diffusion of both the helium and heavy element abundances by. weight (Y and Z. respectively). using the diffusion coefficients of Thoul et al. (," We accounted for the diffusion of both the helium and heavy element abundances by weight (Y and Z, respectively), using the diffusion coefficients of Thoul et al. ("
+1994). in the way described bx Guenther Demarque (1997) in their models of (he Sun.,"1994), in the way described by Guenther Demarque (1997) in their models of the Sun."
+ As expected. all the stellar models in this study. exhibit a convective envelope.," As expected, all the stellar models in this study exhibit a convective envelope."
+ In such stellar models. the ealeulated pulsational frequencies are known (o be sensitive to the treatment of the atmosphere and outer convective lavers (Guenther οἱ al.," In such stellar models, the calculated pulsational frequencies are known to be sensitive to the treatment of the atmosphere and outer convective layers (Guenther et al."
+ 1992: Robinson et al., 1992; Robinson et al.
+ 2003)., 2003).
+ In the atmosphere. we used the Eddineton T-7 relation for a grev atmosphere. ancl {he mixing length theory to describe convectively unstable lavers (Bóhhim-Vitense 1953).," In the atmosphere, we used the Eddington $\tau$ relation for a grey atmosphere, and the mixing length theory to describe convectively unstable layers (Böhhm-Vitense 1958)."
+ The caleulated radius also depends on the choice of the mixing length parameter a (mixing length to scale-height ratio)., The calculated radius also depends on the choice of the mixing length parameter $\alpha$ (mixing length to scale-height ratio).
+ We set a = 1.1 for all models. close to the value required to reproduce the solar radius under the same physical assumptions and stellar evolution code.," We set $\alpha$ = 1.7 for all models, close to the value required to reproduce the solar radius under the same physical assumptions and stellar evolution code."
+ A fuller description of the treatment of the outer lavers aud of the constitutive physics incorporated, A fuller description of the treatment of the outer layers and of the constitutive physics incorporated
+sequence (22)...,"sequence \citep{Baumgarte99e,Shibata02}."
+ The probability of an IMDIT formine from such a massive Pop III star certainly depeuds ou the initial stellar mass fiction (IME). as well as highly uncertain details of zero-metallicity stellar evolution.," The probability of an IMBH forming from such a massive Pop III star certainly depends on the initial stellar mass function (IMF), as well as highly uncertain details of zero-metallicity stellar evolution."
+ However. there has beeu some sugeestions that tle IME in the carly universe is quite top-leavy (27)... aud that stellar nmiass loss is neeheible (?7)..," However, there has been some suggestions that the IMF in the early universe is quite top-heavy \citep{Schneider:02topheavy,Abel:02firststar}, and that stellar mass loss is negligible \citep{Fryer:01spin,Heger:03sn}."
+ These sugsestions nuply that DIIS ina proto globular cluster euvirouimoeut are more massive than those formed in more recent times (?).., These suggestions imply that BHs in a proto globular cluster environment are more massive than those formed in more recent times \citep{vandermarel:2004ib}.
+ Ou the other hand. Pop IHE star formation is thought to take place in dark matter overdensitics at ;~12.20 (?) and only the most massive elobular clusters are thought to be cmbedded in a significant οποιο]. dark matter overdensity to allow for Pop III formation.," On the other hand, Pop III star formation is thought to take place in dark matter overdensities at $z\sim 12-20$ \citep{Madau:01Pop3} and only the most massive globular clusters are thought to be embedded in a significant enough dark matter overdensity to allow for Pop III formation."
+ However the IMDBII forms. for the first ~0.5Cr after formation it is dvuamically active (2)..," However the IMBH forms, for the first $\sim 0.5~{\mathrm{Gyr}}$ after formation it is dynamically active \citep{Spitzer:1987gcbook}."
+ Due to mass scerceation. the iitial environment around an IMDBIT ina elobular cluster is especially rich im BITS (2)..," Due to mass segregation, the initial environment around an IMBH in a globular cluster is especially rich in BHs \citep{Fregeau:02msgc}."
+ While may of these DIIs are quickly ejected by few body interactions with the INDIT. enough remain to subject the INIBIT to LOs100 ιασος with DIIs iu the primordial globular cluster svstem (??7)..," While many of these BHs are quickly ejected by few body interactions with the IMBH, enough remain to subject the IMBH to $10~{\mathrm{s}}-100~{\mathrm{s}}$ mergers with BHs in the primordial globular cluster system \citep{Portegies:00bhmerge,Oleary:2005bm,Gultekin:2004gi}."
+ Iu Πο of ecueral relativistic black hole mereer simulations. surviving this IMDBIT-BIT merecr epoch may be difficult.," In light of general relativistic black hole merger simulations, surviving this IMBH-BH merger epoch may be difficult."
+ Receut advances in nunuerical relativity have at last pinned down the dynamics of black hole merecrs VAiuulatius the coalescence. merecr. and rinedown of equalinass circular nonu-spiunuiue binary black holes iu full general relativity —.(277)..," Recent advances in numerical relativity have at last pinned down the dynamics of black hole mergers – simulating the coalescence, merger, and ringdown of equal-mass circular non-spinning binary black holes in full general relativity \citep{Pretorius:2005gq,Campanelli05a,Baker:2006yw}."
+ Iu. the past veur. more astroplivsically relevant numerical simulatious. iucludiug spins aud unequal masses. have been published by several eroups.," In the past year, more astrophysically relevant numerical simulations, including spins and unequal masses, have been published by several groups."
+ Oue of the most exciting results of ecucral relativity for structure formation is that binary black hole svstenus stronely radiate linear momentum im the forma of eravitational waves during the plunge phase of the inspiral., One of the most exciting results of general relativity for structure formation is that binary black hole systems strongly radiate linear momentum in the form of gravitational waves during the plunge phase of the inspiral.
+ This radiatiou directly results from an asvunuetiy in the orbital configuration and can eoncrically vield a gravitational wave “kick” velocity as ast as L000Iuns3 (2921.," This radiation directly results from an asymmetry in the orbital configuration and can generically yield a gravitational wave “kick” velocity as fast as $4000~\KMS$ \citep{Gonzalez:2007hi,Campanelli:2007cg}."
+ Even typical kick velocities (~200kms. +) are interestingly large when compared o the escape velocity of au average elobular cluster (~FO)s 1) (777).," Even typical kick velocities $\sim 200~\KMS$ ) are interestingly large when compared to the escape velocity of an average globular cluster $\sim 50 \KMS$ ) \citep{Webbink:85vesc,Favata:2004wz,Merritt:2004xa}."
+ Deuce. regardless of how au IMDII formes. the biggest challenge may be to determine row a elobular cluster them in the face of a repeated ouslaught of eravitational wave kicks from nergers with other black holes.," Hence, regardless of how an IMBH forms, the biggest challenge may be to determine how a globular cluster them in the face of a repeated onslaught of gravitational wave kicks from mergers with other black holes."
+ Of course if INIBIIS are ormed through mergers of stellaz-mass black holes. the ormation mechanis itself needs to be able to account or these large kicks as well.," Of course if IMBHs are formed through mergers of stellar-mass black holes, the formation mechanism itself needs to be able to account for these large kicks as well."
+ Tn this paper we explore IMDITI retention within voung globular clusters after collisious with DIIs., In this paper we explore IMBH retention within young globular clusters after collisions with BHs.
+ We calculate the retention probability under a variety of asstuuptious for the initial IAIBIT seed amass. the DII nass distribution. aud the iuitial spin distributions.," We calculate the retention probability under a variety of assumptions for the initial IMBH seed mass, the BH mass distribution, and the initial spin distributions."
+ In addition. for the mass range enconrpassed by the IMDII ormation Channels listed above (stellar runaway. stellar-nass DII collisious. Pop III stars). we simulate the IAIBI survival probability after mereiue with the available DIIs over a collision timescale.," In addition, for the mass range encompassed by the IMBH formation channels listed above (stellar runaway, stellar-mass BH collisions, Pop III stars), we simulate the IMBH survival probability after merging with the available BHs over a collision timescale."
+ Caven the known current structure of the Milkv Way elobular cluster svsteun. we can then estimate the fraction of globular clusters hat may have allowed a particular IMIBID formation channel.," Given the known current structure of the Milky Way globular cluster system, we can then estimate the fraction of globular clusters that may have allowed a particular IMBH formation channel."
+ Conmbiuiug these two key estimates. we can determine the maxima expected number of Milkv Wav elobular clusters that could have retained their INIT.," Combining these two key estimates, we can determine the maximum expected number of Milky Way globular clusters that could have retained their IMBHs."
+ We outline the approach in Sec., We outline the approach in Sec.
+ 2. preseut the results in Sec.," 2, present the results in Sec."
+ 3 and discuss the caveats. inplicatious. aud. future directions in Sec.," 3 and discuss the caveats, implications, and future directions in Sec."
+ [., 4.
+ Cravitational recoil estimates of binary black hole iiergers have been addressed using both seud-analvtic iethods in the unequalauass. non-spiuuiug case (272777) and nunerical methods im more general. spiuniug. tunequal mass scenarios (22272???7)..," Gravitational recoil estimates of binary black hole mergers have been addressed using both semi-analytic methods in the unequal-mass, non-spinning case \citep{Fitchett:1983fc,Favata:2004wz,Damour:2006tr,Blanchet:2005rj,2006PhRvD..74l4010S}
+ and numerical methods in more general, spinning, unequal mass scenarios \citep{Herrmann:2006ks,Baker:2006vn,Gonzalez:2006md,Herrmann07,Koppitz:07kick,gonzalez-2007,campanelli-2007-659,campanelli-2007,tichy-2007}."
+" For non-spiuniug binaries. the most comprehensive numerical study ofkicks is that of ον, which was fouud in agreement with the SCni-analytic estimates of (7).. where a miaximunu kick velocity of ~175+410kins was obtained for mass ratio gq=ALΕυ0.02)."," For non-spinning binaries, the most comprehensive numerical study ofkicks is that of \citet{Gonzalez:2006md}, , which was found in agreement with the semi-analytic estimates of \citep{2006PhRvD..74l4010S}, where a maximum kick velocity of $\sim 175\ \pm 10 \; \KMS$ was obtained for mass ratio $q=M_1/M_2 \sim 0.36 \pm 0.02$ )."
+ The eravitational recoil is expected Το increase with increasing spin (??).. aud this behavior las been confirmed by several umuerical relativity eroups.," The gravitational recoil is expected to increase with increasing spin \citep{Redmount:1989rr,WhitbeckPhD}, and this behavior has been confirmed by several numerical relativity groups."
+ Aunuuerical snuauulations lave demonstrated that the radiative linear momentum loss predicted by Post-Newtonianun studies (??)— can be used to fit numerical results. vielding a generalized formula for the recoil velocity as a function| of the idividual black holes spin. initial orieutation. phase at merger aud mass ratio.," Numerical simulations have demonstrated that the radiative linear momentum loss predicted by Post-Newtonian studies \citep{RevModPhys.52.299,PhysRevD.52.821} can be used to fit numerical results, yielding a generalized formula for the recoil velocity as a function of the individual black hole's spin, initial orientation, phase at merger and mass ratio."
+ We here adopt the parameterized fit of ? adding the expected (1|6) contribution for ecceutric orbits (7) to vield the following formula: where aud where the fitting coustauts are d=1.2.101kins B=093. Π=(73403)ς107lans.J| and NK—(G.02:0.1)«104Jans 1 while the subscripts laud 2 refer to the first and secoud.DIT respectively.," We here adopt the parameterized fit of \citet{campanelli-2007-659} adding the expected $(1+e)$ contribution for eccentric orbits \citep{2006astro.ph.11110S}
+ to yield the following formula: where and where the fitting constants are $A=1.2 \times 10^{4} \; \KMS$, $B =
+-0.93$, $H=(7.3 \pm 0.3) \times 10^{3} \; \KMS$, and $K=(6.0 \pm 0.1)
+\times 10^{4} \; \KMS$ , while the subscripts $1$ and $2$ refer to the first and secondBH respectively."
+ The unit vectors Grey) are orthogonal to cach other and span the initial orbital plane. while L aud | stands for perpendicular and parallel to the orbital augular mioimentuni," The unit vectors $(\hat{x},\hat{y})$ are orthogonal to each other and span the initial orbital plane, while $\perp$ and $\parallel$ stands for perpendicular and parallel to the orbital angular momentum."
+" There are | fitting parameters: the mass ratio g=Ab/AL: the reduced spi parameter a;=$;AL}. where ο is the spin aneular momentum of BIT /: aud the eccentricity ο,"," There are $4$ fitting parameters: the mass ratio $q\equiv M_2/M_1$; the reduced spin parameter $\alpha_i \equiv
+S_i/M_i^2$, where $S_i$ is the spin angular momentum of BH $i$; and the eccentricity $e$."
+" Iu addition. there are 3 angles to specity the orientation of the merger: O. the anele between the “in-plane” component of8=(AL,|M3)(50Mo5iAA) and the infall dixection at Ἱποσοι: Oy. theangle between οἱ and the initial direction of motion: aud ©. the angle between the unequal mass and spin contribution to the recoil iu the orbital plane."," In addition, there are 3 angles to specify the orientation of the merger: $\Theta$, the angle between the “in-plane” component of $\delta^i \equiv (M_1 + M_2) \left(S_2^i/M_2
+ - S_1^i/M_1\right)$ and the infall direction at merger; $\Theta_0$ , theangle between $\delta^i$ and the initial direction of motion; and $\xi$ , the angle between the unequal mass and spin contribution to the recoil in the orbital plane."
+ The imiergeer phase withiu the orbital plane may also play a role. but it is not included explicitly," The merger phase within the orbital plane may also play a role, but it is not included explicitly"
+a discussion of future work in $4..,a discussion of future work in \ref{sec:conclusion}.
+ We begin by summarizing the main features of our simmering models. (, We begin by summarizing the main features of our simmering models. (
+For further details. the interested reader should refer to Woosley et al.,"For further details, the interested reader should refer to Woosley et al."
+ 2004: Lesaffre et al., 2004; Lesaffre et al.
+ 2006: Piro Bildsten 2007: Piro 2008.), 2006; Piro Bildsten 2007; Piro 2008.)
+ Unstable ignition of C occurs when the heating from carbon fusion beats neutrino cooling., Unstable ignition of $^{12}$ C occurs when the heating from carbon fusion beats neutrino cooling.
+ The central temperature then rises and à convective zone grows outward. eventually encompassing ~1M.. of the WD after ~10?yrs.," The central temperature then rises and a convective zone grows outward, eventually encompassing $\sim1 M_\odot$ of the WD after $\sim10^3\ {\rm yrs}$."
+ As the central temperature. 7.. increases. carbon burning becomes more vigorous and the heating timescale. tj—(danT./dry!. gets shorter.," As the central temperature, $T_c$, increases, carbon burning becomes more vigorous and the heating timescale, $t_h\equiv (d\ln T_c/dt)^{-1}$, gets shorter."
+ This timescale in general depends on the size of the region responding to the rising central temperature at the core., This timescale in general depends on the size of the region responding to the rising central temperature at the core.
+" Simmering ends and a burning wave commences once rj,= fous. Where fi Is eddy overturn timescale."," Simmering ends and a burning wave commences once $t_h\lesssim t_{\rm conv}$ , where $t_{\rm conv}$ is eddy overturn timescale."
+ At these late times. individual eddies may experience significant heating during their transit (Garcfa-Senz&Woosley1995).. so that the temperature profile is no longer an adiabat and the entire convective core does not respond to the increasing central temperature. (," At these late times, individual eddies may experience significant heating during their transit \citep{gw95}, so that the temperature profile is no longer an adiabat and the entire convective core does not respond to the increasing central temperature. ("
+"In contrast. we show below that during the majority of the time f, depends on the heat capacity of the entire convective mass.)","In contrast, we show below that during the majority of the time $t_h$ depends on the heat capacity of the entire convective mass.)"
+ This makes it difficult to exactly calculate the precise moment when simmering ends., This makes it difficult to exactly calculate the moment when simmering ends.
+ For this reason Lesaffreetal.(2006) explore t=0fj. Wherea=| parameterizes this uncertainty.," For this reason \citet{les06} explore $t_{\rm conv}=\alpha t_h$, where $\alpha~\lesssim~1$ parameterizes this uncertainty."
+ Since weare only roughly concerned with resolving the end of the simmering phase. we take fiom7fmcuffe. where cy is the specific heat capacity at constant pressure. e is the heating from carbon burning. and all these quantities are evaluated at the WD center.," Since weare only roughly concerned with resolving the end of the simmering phase, we take $t_{\rm conv}\approx t_h\approx c_p T_c/\epsilon$, where $c_p$ is the specific heat capacity at constant pressure, $\epsilon$ is the heating from carbon burning, and all these quantities are evaluated at the WD center."
+" This estimates that simmering should end when fj~7s at a central temperature and density of 7.z:7.8«I0K and p,z2.6«10?gem. which is roughly in agreement with the results presented by Woosleyetal.(2004) using the Kepler stellar evolution code (Weaveretal.1978)."," This estimates that simmering should end when $t_h\approx7\ {\rm s}$ at a central temperature and density of $T_c\approx7.8\times10^8\ {\rm K}$ and $\rho_c\approx2.6\times10^9\ {\rm g\ cm^{-3}}$, which is roughly in agreement with the results presented by \citet{woo04} using the Kepler stellar evolution code \citep{wea78}."
+ We follow the simmering phase by calculating a series of hydrostatic WD models. each with a different central temperature. but at a fixed mass (seePiro&Bildsten2007;Piro 2008).," We follow the simmering phase by calculating a series of hydrostatic WD models, each with a different central temperature, but at a fixed mass \citep[see][]{pb07,pir08}."
+. For simplicity we ignore the convective Urea process since our focus is on the last ~10°s., For simplicity we ignore the convective Urca process since our focus is on the last $\sim10^5\ {\rm s}$.
+" The timescale for thermal conduction across the WD is fy,=K,/R10° yrs. where K, is the conductivity and A is the radius. which is much longer the timescale over which heating is occurring."," The timescale for thermal conduction across the WD is $t_{\rm th}\equiv K_c/R^2\sim10^6\ {\rm yrs}$ , where $K_c$ is the conductivity and $R$ is the radius, which is much longer the timescale over which heating is occurring."
+ Therefore the convection efficiently mixes entropy and the convective region nearly follows an adiabat out from the WD center., Therefore the convection efficiently mixes entropy and the convective region nearly follows an adiabat out from the WD center.
+ Outside the convective zone. we assume the WD is isothermal with a temperature 7}.," Outside the convective zone, we assume the WD is isothermal with a temperature $T_i$."
+" To understand how time-dependent convection is different than from that normally found in steady-state convection. we focus on the time-dependent entropy equation where L,. is the convective luminosity."," To understand how time-dependent convection is different than from that normally found in steady-state convection, we focus on the time-dependent entropy equation = - where $L_c$ is the convective luminosity."
+ This equation omits the work required to expand the WD as the heating takes place. which is a significant amount of energy and thus requires some discussion.," This equation omits the work required to expand the WD as the heating takes place, which is a significant amount of energy and thus requires some discussion."
+ For each convective model we compared the total change in WD binding energy to the total change in internal energy of the electrons (which. are primarily degenerate and relativistic). including the electron Coulomb interaction energy (accordingtoChabrier&Potekhin1998).," For each convective model we compared the total change in WD binding energy to the total change in internal energy of the electrons (which are primarily degenerate and relativistic), including the ion-electron Coulomb interaction energy \citep[according to][]{cp98}."
+. These two quantities are equal to the numerical accuracy of our integrations. which demonstrates that all of the work required to expand the WD comes from changes in the internal energy of the electrons.," These two quantities are equal to the numerical accuracy of our integrations, which demonstrates that all of the work required to expand the WD comes from changes in the internal energy of the electrons."
+ Thus. the entropy created from nuclear burning all goes into convective motions or the internal thermal energy. and we are justified in omitting the binding energy and electron internal energy terms from equation (1).," Thus, the entropy created from nuclear burning all goes into convective motions or the internal thermal energy, and we are justified in omitting the binding energy and electron internal energy terms from equation \ref{eq:entropy}) )."
+ The temperature profile in the convective zone follows an adiabat with a power law index n=(OInT/OlnP)44., The temperature profile in the convective zone follows an adiabat with a power law index $n\equiv(\partial\ln T/\partial\ln P)_{\rm ad}$.
+ The time derivative of the temperature at a given pressure can be eXpressed as [üt =(P/P. is set by the central temperature change TIE Therefore there is a well-defined. global heating timescale. ο," The time derivative of the temperature at a given pressure can be expressed as t = is set by the central temperature change = Therefore there is a well-defined, global heating timescale, $t_h$."
+" To account for the changing central pressure in à more rigorous calculation. we must take partial derivatives at constant mass coordinate. M@,."," To account for the changing central pressure in a more rigorous calculation, we must take partial derivatives at constant mass coordinate, $M_r$."
+ This can be performed by inverting the empirically found 4.(7;.T.) relation presented in Piro (2008: or see eq. [14]]," This can be performed by inverting the empirically found $M_c(T_i,T_c)$ relation presented in Piro (2008; or see eq. \ref{eq:mct}] ]"
+ below). where 7; is the nearly isothermal temperature of the non-convective. conductive region.," below), where $T_i$ is the nearly isothermal temperature of the non-convective, conductive region."
+" The result is TC.Mj4 e I where j/, 1s the mean molecular weight per electron."," The result is 	T(T_c,M_r) = ], where $\mu_e$ is the mean molecular weight per electron."
+" This can be used to find (4111/dr)s;,2dInT.dt. which confirms our conclusion that f, is the same at any/ depth within the convective zone."," This can be used to find $(d\ln T/dt)_{M_r}=d\ln T_c/dt$, which confirms our conclusion that $t_h$ is the same at any depth within the convective zone."
+" Multiplying equation (1). by dM,=Amrpdr and integrating. = eM, - LAM;YL0)."," Multiplying equation \ref{eq:entropy}) ) by $dM_r=4\pi r^2\rho dr$ and integrating, = dM_r 		- L_c(M_r)+L_c(0)."
+ We pull f; outside of the left-hand integral to find Cl=, We pull $t_h$ outside of the left-hand integral to find 	=
+ We pull f; outside of the left-hand integral to find Cl=F, We pull $t_h$ outside of the left-hand integral to find 	=
+Changes in the line profile could be explained if some part of the background continuum Lux arises from a region of such small angular size that it gives rise to inter-stellar scintillation (LSS) (Dennett-Thorpe&deBruyn2000:Rick-ett1990:Wedziora-Chuclezeretal.LOOT) on passing through the interstellar medium of the Milkv. Was.,"Changes in the line profile could be explained if some part of the background continuum flux arises from a region of such small angular size that it gives rise to inter-stellar scintillation (ISS) \cite{ger,rickett,lucyna} on passing through the interstellar medium of the Milky Way."
+ The spectrum associated with this region will vary. while that associated with the non-scintillating Ες will not: this will result in changes in the net optical depth.," The spectrum associated with this region will vary, while that associated with the non-scintillating flux will not; this will result in changes in the net optical depth."
+ LSS will also. produce variations in the continuum flux: further. these variations would be related to the changes in the optical depth.," ISS will also produce variations in the continuum flux; further, these variations would be related to the changes in the optical depth."
+ Since we measure both the background Lux and τοι. one can test whether the observations are consistent with LSS models.," Since we measure both the background flux and $\tau_{21}$, one can test whether the observations are consistent with ISS models."
+ As is discussed in more detail below. the lack of correlation between the variations in flux and opacity rule out the simplest such models.," As is discussed in more detail below, the lack of correlation between the variations in flux and opacity rule out the simplest such models."
+ 5 Cllz VLBI observations of (Ivellermannet.αἱ.1998) have shown that it has a simple core-jet structure., 15 GHz VLBI observations of \cite{keller} have shown that it has a simple core-jet structure.
+ The core is unresolved: at. sub-milliaresecond resolution., The core is unresolved at sub-milliarcsecond resolution.
+ “Phe simplest. scintillation moclel is that of a single compact scintillating component in the core. with the rest of the flux not giving rise to scintillation.," The simplest scintillation model is that of a single compact scintillating component in the core, with the rest of the flux not giving rise to scintillation."
+ In this case. however. two epochs with the same observed continuum Hux should have the same observed. spectrum: 'urther. the line depths should be deeper at epochs of higher lux and vice versa.," In this case, however, two epochs with the same observed continuum flux should have the same observed spectrum; further, the line depths should be deeper at epochs of higher flux and vice versa."
+ Table (1) shows that the Hux values agree within the error bars on the 22nd of September and he Sth of October: however. the spectra at these epochs are considerably. dillerent.," Table (1) shows that the flux values agree within the error bars on the 22nd of September and the 8th of October; however, the spectra at these epochs are considerably different."
+ Similarly. the background Lux is he same on the 10th of May and the 40th of October while he spectra are again quite dilferent.," Similarly, the background flux is the same on the 10th of May and the 10th of October while the spectra are again quite different."
+ Given our error. bars. he Hux change required to produce the observed cillerences in the 21em spectra would. have been easily detectable. unless the linc-of-5sight to the scintillating component. also coincidentally happened to have an optical depth. of order unity.," Given our error bars, the flux change required to produce the observed differences in the 21cm spectra would have been easily detectable, unless the line-of-sight to the scintillating component also coincidentally happened to have an optical depth of order unity."
+ The latter situation is also. however. not viable since he epoch with the highest background flux (S=7.9+1.2 Jv. on the 6th of May) has the weakest absorption.," The latter situation is also, however, not viable since the epoch with the highest background flux $S = 7.9 \pm 1.2$ Jy, on the 6th of May) has the weakest absorption."
+ Thus. he simplest ISS model appears to be ruled out.," Thus, the simplest ISS model appears to be ruled out."
+ The next simplest [SS model is one in which there are wo scintillating components and a steady component., The next simplest ISS model is one in which there are two scintillating components and a steady component.
+ The second. scintillating component could. for example. arise in he bright unresolved knot in the jet of which is separated. from the core by ~2.5 milliareseconds.," The second scintillating component could, for example, arise in the bright unresolved knot in the jet of, which is separated from the core by $\sim 2.5$ milliarcseconds."
+ If the optical depth towards cach of the. scintillating components is comparable to the average optical depth over the entire source (Le. τοι 0.1) then. as discussed below. this model is also ruled. out by the data. of 22nd September and Sth October.," If the optical depth towards each of the scintillating components is comparable to the average optical depth over the entire source (i.e. $\tau_{21} \sim 0.1$ ) then, as discussed below, this model is also ruled out by the data of 22nd September and 8th October."
+ The dilference spectrum on sth October consists of two components which are ~ 50 mJv deeper than the average spectrum and three which are 30 Ην deeper than the average., The difference spectrum on 8th October consists of two components which are $\sim$ 50 mJy deeper than the average spectrum and three which are $\sim 30$ mJy deeper than the average.
+ On the 22nc of September the first two components are weaker (hy ~3040 mJy) than the average. while the remaining three agree (within the noise) with the average spectrum.," On the 22nd of September the first two components are weaker (by $\sim 30 - 40$ mJy) than the average, while the remaining three agree (within the noise) with the average spectrum."
+" If one assumes that the absorbing clouds are distributed such that the changes in the first two line components arise from. one of the scintillating source components (total change in line depth. AS~90100 niv. per component) and the changes in the remaining three stem from the other (total change in line depth. AS—30040 mJs. per component). the changes in line depth would correspond to a flux change of at least 1.3 Jv (form, £20.1). which is not seen (other distributions of the clouds would exacerbate the problem)."," If one assumes that the absorbing clouds are distributed such that the changes in the first two line components arise from one of the scintillating source components (total change in line depth, $\Delta S \sim 90 - 100$ mJy, per component) and the changes in the remaining three stem from the other (total change in line depth, $\Delta S \sim 30 - 
+40$ mJy, per component), the changes in line depth would correspond to a flux change of at least 1.3 Jy (for $\tau_{21} \la 0.1$ ), which is not seen (other distributions of the clouds would exacerbate the problem)."
+ Since these two compact components contain far more than 104 of the total Iux. it is rather unlikely that the optical depth in front of the scintillating components is 1. while the average optical depth is ~0.1.," Since these two compact components contain far more than $10\%$ of the total flux, it is rather unlikely that the optical depth in front of the scintillating components is $\sim 1$, while the average optical depth is $\sim 0.1$."
+ One could construct more complicated mocdels. wherein each of the compact VLBI components contain more than one scintillating component.," One could construct more complicated models, wherein each of the compact VLBI components contain more than one scintillating component."
+ The relation between the observed: spectrum and the measured Dux. becomes. more and more complex with an inereasing number of scintillating components and it would hence be possible το contrive situations where both the variations might be explainec purely as a result of ISS in the Galaxy., The relation between the observed spectrum and the measured flux becomes more and more complex with an increasing number of scintillating components and it would hence be possible to contrive situations where both the variations might be explained purely as a result of ISS in the Galaxy.
+" However. from the ""avlor-C'ordes (1993) model of the distribution of tonizec eas in the Galaxy. the scatter broadening expected towards is 0.4 milliareseconds."," However, from the Taylor-Cordes (1993) model of the distribution of ionized gas in the Galaxy, the scatter broadening expected towards is 0.4 milliarcseconds."
+ I thus seems unlikely. that uncorrelated LSS would result. even if the core or kno has multiple compact components.," It thus seems unlikely that uncorrelated ISS would result, even if the core or knot has multiple compact components."
+ We note. however. tha ISS might also occur in the galaxv(ies) giving rise to the absorption at z=0.3127. while all the above mocdels only consider the case of it occurring in the Galaxy.," We note, however, that ISS might also occur in the galaxy(ies) giving rise to the absorption at $z=0.3127$, while all the above models only consider the case of it occurring in the Galaxy."
+ Scintillation could prove to be a viable model if significant LSS does occur in the interstellar medium of the absorber., Scintillation could prove to be a viable model if significant ISS does occur in the interstellar medium of the absorber.
+ Motion of a source component across the sky migh explain the changes in absorption profile. if the absorbing system has opacity variations transverse to the line-of-sight.," Motion of a source component across the sky might explain the changes in absorption profile, if the absorbing system has opacity variations transverse to the line-of-sight."
+ The linc-of-sight to the moving component samples dilferen paths through the absorber at different epochs. resulting in changes in the profile.," The line-of-sight to the moving component samples different paths through the absorber at different epochs, resulting in changes in the profile."
+ Super-Iuminal motion on sub-VLBL scales has often been invoked to explain rapid. variability of QSO Iluxes (Rees1966:Woltjer1966).," Super-luminal motion on sub-VLBI scales has often been invoked to explain rapid variability of QSO fluxes \cite{rees,woltjer}."
+.. LE these variations are intrinsic. they ave dillieult to explain in the stanclare context of incoherent svnchrotron. emission. as causality arguments. would. imply brightness temperatures far in excess of the limit of 107. ΑΝ set. by the inverse Compton catastrophe.," If these variations are intrinsic, they are difficult to explain in the standard context of incoherent synchrotron emission as causality arguments would imply brightness temperatures far in excess of the limit of $10^{12}$ K set by the inverse Compton catastrophe."
+ Super-Iuminal motion of a source component produces a Doppler. boosting of the observed. brightness temperature by a factor 9 over the true source brightness temperature. where 6 is the Doppler factor.," Super-luminal motion of a source component produces a Doppler boosting of the observed brightness temperature by a factor $\delta^3$ over the true source brightness temperature, where $\delta$ is the Doppler factor."
+" ""Πιν. if 9 is sulliciently large. the true brightness temperature would be lower than the inverse Compton limit."," Thus, if $\delta$ is sufficiently large, the true brightness temperature would be lower than the inverse Compton limit."
+ In these models. the timescale of variability can be used to place a lower limit on 9.," In these models, the timescale of variability can be used to place a lower limit on $\delta$."
+ Doppler factors of order LO are sullicient to account for the variability seen in the majority of compact sources: similar values have also been measured. [rom VLBI observations of super-Iuminal motion., Doppler factors of order 10 are sufficient to account for the variability seen in the majority of compact sources; similar values have also been measured from VLBI observations of super-luminal motion.
+ However. observed. variations in the 1420 MIIz flux of AO 0235|164 on timescales of days," However, observed variations in the 1420 MHz flux of AO 0235+164 on timescales of days"
+"in the system of equations as derivatives with respect to the scale factor « for introducing the common parameterisation for the dark energy equation ofstate (22). With solutions for p, and pasto. the Hubble function Ata) is can be obtained by integrating. which in addition gives the detinition of the critical density pane)=32Η”Απ). with Newton's constant G.","in the system of equations as derivatives with respect to the scale factor $a$ for introducing the common parameterisation for the dark energy equation ofstate \citep{1997PhRvD..56.4439T, 2003MNRAS.346..573L}, With solutions for $\rho_m(a)$ and $\rho_\phi(a)$ , the Hubble function $H(a)$ is can be obtained by integrating, which in addition gives the definition of the critical density $\rho_\mathrm{crit}(a)=3H^2(a)/(8\pi G)$, with Newton's constant $G$."
+" Initial conditionsare defined at the present epoch. Πα=1)Ap. Ota=1)O, and Ott=1)|— O4."," Initial conditionsare defined at the present epoch, $H(a=1)=H_0$, $\Omega_m(a=1)=\Omega_m$ and $\Omega_\phi(a=1)=1-\Omega_m$ ."
+" The comoving distance y follows from the solution for Ata). c denoting the speed of light. and the density parameters ean be defined by normalising O,,(@)=pyGI/piuta) and Outlay=PhO pauta."," The comoving distance $\chi$ follows from the solution for $H(a)$, $c$ denoting the speed of light, and the density parameters can be defined by normalising $\Omega_m(a)=\rho_m(a)/\rho_\mathrm{crit}(a)$ and $\Omega_\phi(a)=\rho_\phi(a)/\rho_\mathrm{crit}(a)$ ."
+" Table | gives an overview over the dark energy models considered in this paper. in terms of their dark energy properties wo. w, and their CDM decay rate T."," Table \ref{table_models} gives an overview over the dark energy models considered in this paper, in terms of their dark energy properties $w_0$, $w_a$ and their CDM decay rate $\Gamma$."
+" In the following I focus on spatially flat models with O,,+0,=I. which isconserved in the time evolution of eqn. CI)."," In the following I focus on spatially flat models with $\Omega_m+\Omega_\phi=1$, which isconserved in the time evolution of eqn. \ref{eqn_cosmology}) )."
+" The homogeneous growth of the overdensity field. ó(x.a)=Dtayótx.1) is described by the growth function D(a). which is a solution to the differential equation (22).. In the standard cold dark matter (SCDM) cosmology with Q,=l and 34dlnH/dine=3. the solution for Dyta) is simply the scale factor. D,(a)=a."," The homogeneous growth of the overdensity field, $\delta(\bmath{x},a)=D_+(a)\delta(\bmath{x},1)$ is described by the growth function $D_+(a)$, which is a solution to the differential equation \citep{1998ApJ...508..483W, 2003MNRAS.346..573L}, In the standard cold dark matter (SCDM) cosmology with $\Omega_m=1$ and $3+\dd\ln H/\dd\ln a=\frac{3}{2}$, the solution for $D_+(a)$ is simply the scale factor, $D_+(a)=a$."
+ Eqn. (59) , Eqn. \ref{eqn_growth}) )
+can be applied for describing structure growth in coupled models as well. because the overdensity field. being the ratio between the density perturbation and the mean density. is not atfected by e.g. CDM decay. as long as the dark energy fluid can be considered to be homogeneous and not to be influencing local structure formation.," can be applied for describing structure growth in coupled models as well, because the overdensity field, being the ratio between the density perturbation and the mean density, is not affected by e.g. CDM decay, as long as the dark energy fluid can be considered to be homogeneous and not to be influencing local structure formation."
+" Nevertheless. coupled models influence the growth equation by the scaling of the Hubble function (2) and the time evolution of the density parameter £2,(4)."," Nevertheless, coupled models influence the growth equation by the scaling of the Hubble function $H(a)$ and the time evolution of the density parameter $\Omega_m(a)$."
+ Additionally. the initial conditions of the growth equation are different. as the cosmology does not have an asymptotic SCDM phase with Hxa77 at early times.," Additionally, the initial conditions of the growth equation are different, as the cosmology does not have an asymptotic SCDM phase with $H\propto a^{-3/2}$ at early times."
+ The inital conditions can be specitied by assuming an asymptotic power-law growth ο ια>0.," The inital conditions can be specified by assuming an asymptotic power-law growth $D_+(a)=a^\alpha$ , $\alpha>0$."
+ Substitution into eqn. (5)), Substitution into eqn. \ref{eqn_growth}) )
+ yields a quadratic equation for c. whose positive solution specifies the initial conditions a and dD)(anda=! at some early time a; (c.f.2)..," yields a quadratic equation for $\alpha$, whose positive solution specifies the initial conditions $D_+(a_i)=a_i^\alpha$ and $\dd D_+(a_i)/\dd a=\alpha a_i^{\alpha-1}$ at some early time $a_i$ \citep[c.f.][]{lensing_paper}."
+ Fig., Fig.
+ |. shows au;growth functions for the four exemplary cosmologies., \ref{fig_growth} shows growth functions for the four exemplary cosmologies.
+ Evolving dark energy has the property of suppressing structure formation at an earlier time. which can be partially compensated by CDM decay (because the gravitational fields generated by the overdensity ó(x) are stronger if Ομ) has a higher value). causing degeneracies between the equation of state parameters and the CDM decay rate: The ACDM and ójCDM- for instance. are almost indistinguishable.," Evolving dark energy has the property of suppressing structure formation at an earlier time, which can be partially compensated by CDM decay (because the gravitational fields generated by the overdensity $\delta(\bmath{x})$ are stronger if $\Omega_m(a)$ has a higher value), causing degeneracies between the equation of state parameters and the CDM decay rate: The $\Lambda$ CDM and $\phi_\Gamma$ CDM-models, for instance, are almost indistinguishable."
+ It is worth noting that models with decaying CDM are genuinely different from dark energy models concerning structure growth and probes which use gravitational interaction. such as gravitational lensing or the ISW-etfect., It is worth noting that models with decaying CDM are genuinely different from dark energy models concerning structure growth and probes which use gravitational interaction such as gravitational lensing or the iSW-effect.
+" While it is always possible to construct an equation of state wa) for a dark energy model with stable CDM that gives the identical Hubble function as a model with decaying CDM. this degeneracy is broken in the evolution of the density parameter Q,,(@) and the growth function D,(a) which both would be ditterent inthe two cases."," While it is always possible to construct an equation of state $w(a)$ for a dark energy model with stable CDM that gives the identical Hubble function as a model with decaying CDM, this degeneracy is broken in the evolution of the density parameter $\Omega_m(a)$ and the growth function $D_+(a)$ which both would be different inthe two cases."
+ Theoretically. combining observations of cosmic structure growth and the expansion history should make it to distinguish between the two families of models. which geometrical probes such as supernova observations alone would not be able to do.," Theoretically, combining observations of cosmic structure growth and the expansion history should make it to distinguish between the two families of models, which geometrical probes such as supernova observations alone would not be able to do."
+ The iSW-ettect is caused by gravitational interaction of a CMB photon with a time-evolving potential d., The iSW-effect is caused by gravitational interaction of a CMB photon with a time-evolving potential $\Phi$.
+ The fractional perturbation 7 of the CMB temperature Τον is given by (2) where 5 denotes the conformaltime.," The fractional perturbation $\tau$ of the CMB temperature $T_\mathrm{CMB}$ is given by \citep{1967ApJ...147...73S}
+ where $\eta$ denotes the conformaltime."
+ In the last step. I have replaced the integration variable by thecomoving distance y. which is related tothe conformal time by dy=—cdij —cd//a. and the time derivative of the growth function has been rewritten in terms of the scale factor o. using the definition of the Hubble function de/di= aHta). with the cosmic time /.," In the last step, I have replaced the integration variable by thecomoving distance $\chi$ , which is related tothe conformal time by $\dd\chi = -c\dd\eta = -c\dd t/a$ , and the time derivative of the growth function has been rewritten in terms of the scale factor $a$ , using the definition of the Hubble function $\dd a/\dd t = aH(a)$ , with the cosmic time $t$ ."
+ The gravitational potential ® follows from the Poisson equation in the comoving frame.," The gravitational potential $\Phi$ follows from the Poisson equation in the comoving frame,"
+while interpolation within more general Zi; colourVe/LI] relations for Sun-like stars tends to vield a bluer 0.62 (Sekiguchi Fukugita 2000: irez Melénndez 2005b).,while interpolation within more general $T_{\rm eff}$ –colour–[Fe/H] relations for Sun-like stars tends to yield a bluer $(B-V)_{\odot} \sim 0.62$ (Sekiguchi Fukugita 2000; rez Melénndez 2005b).
+ Phe latter result also implies that the solar 2V predicted by theoretical model atmospheres: (DVy.0.650.67 (Dessell. Castelli Plez 1998: Casagrande et 22005. in prep.)," The latter result also implies that the solar $B-V$ predicted by theoretical model atmospheres: $(B-V)_{\odot} \sim 0.65-0.67$ (Bessell, Castelli Plez 1998; Casagrande et 2005, in prep.)"
+ is too red., is too red.
+ Sekiguchi Fukugita ascribe the discrepaney to systematies in the Zipp scale. sspectroscopic Hafralted Flux Method (HUEM) though within their quoted: uncertainty 0.626+ 0.018. values 70.64 remain plausible.," Sekiguchi Fukugita ascribe the discrepancy to systematics in the $\Teff$ scale, spectroscopic InfraRed Flux Method (IRFM) — though within their quoted uncertainty $(B-V)_{\odot} = 0.626 \pm 0.018$ , values $\sim$ 0.64 remain plausible."
+ We here analyse ‘Sun-like stars’. chosen in a range ol cllective temperatures which brackets the elective temperature of t1c Sun.," We here analyse `Sun-like stars', chosen in a range of effective temperatures which brackets the effective temperature of the Sun."
+ Vhe colours of the stars. in a number of passh:uxds. are Littec [or dependence on Lip and metallicity FeII].," The colours of the stars, in a number of passbands, are fitted for dependence on $\Teff$ and metallicity [Fe/H]."
+ Phe Sun's colours are then indirectly determined. from hese relations or adopted solar vatues of Lipp.—DYYY Wand Fel] =0.0.," The Sun's colours are then indirectly determined from these relations for adopted solar values of $T_{\mathrm eff,\odot} = 5777$ K and [Fe/H] $=
+0.0$."
+ Crucial το this. procedure. is. that the effective temperatures of the stars and he Sun are on the same scale: historically. this has been the basic impasse to estimating the solar colours indirectly from Sun-like stars.," Crucial to this procedure is that the effective temperatures of the stars and the Sun are on the same scale; historically, this has been the basic impasse to estimating the solar colours indirectly from Sun-like stars."
+ During the last two vears. this situation has fundamentally. altered. as interferometric measurements of stellar ciamietcrs have started to become routine al such instruments as the ESO ΝΕΤ (European Southern Observatorys Very Large. ‘Telescope Laterferometer). PPL (Palomar Testbed Interferometer) and NPOL (Navy Prototype Optical Loterferometer)," During the last two years, this situation has fundamentally altered, as interferometric measurements of stellar diameters have started to become routine at such instruments as the ESO VLTI (European Southern Observatory's Very Large Telescope Interferometer), PTI (Palomar Testbed Interferometer) and NPOI (Navy Prototype Optical Interferometer)."
+ These instruments have produced. high accuracy stellar diameters for about 20 main sequence stars rom zz19 lo zz0.7 ((Ixervella. et al.," These instruments have produced high accuracy stellar diameters for about 20 main sequence stars from $\approx
+1.3$ to $\approx 0.7$ (Kervella et al."
+ 2004)., 2004).
+ As a consequence. the elfective emperatures of the sars can be determined in the same manner as for the Sun. from the physical definition of T; rr. in terms of the bolometric luminosity at the Earth. fi... and he angular cliamoeter. 6. via where e is the Stefan-Doltzmann constant.," As a consequence, the effective temperatures of the stars can be determined in the same manner as for the Sun, from the physical definition of $\Teff$ , in terms of the bolometric luminosity at the Earth, $f_{\mathrm bol}$, and the angular diameter, $\theta$, via = ^4 where $\sigma$ is the Stefan-Boltzmann constant."
+" Using 13 stars with directly measured stellar diameters and bolometric Ηχος. jrez Moelénndez (20052) have determined direct. elfective temperatures. ryp ""ον"," Using 13 stars with directly measured stellar diameters and bolometric fluxes, rez Melénndez (2005a) have determined direct effective temperatures, $T_{\mathrm eff}^{\mathrm dir}$."
+ show that the widely used IREAL temperature scale (Alonso. Arribas Alartionez-Roger 1900) is a very good match to T. but also derive an iniproved. version of Alonso et als. IHREM. temperature scale. showing that the systematic dillerence between the IREM and the directly measured temperature scales is not significantly dillerent. from zero (formally. they obtain a difference of LS Ix. with a dispersion around. the mean of 62 Ix. and an error in the mean of 21 [x for their 10 highest quality stars. in the sense that the HUEM scale is the hotter one).," They show that the widely used IRFM temperature scale (Alonso, Arribas Martínnez-Roger 1996) is a very good match to $T_{\mathrm eff}^{\mathrm dir}$, but also derive an improved version of Alonso et al's IRFM temperature scale, showing that the systematic difference between the IRFM and the directly measured temperature scales is not significantly different from zero (formally, they obtain a difference of 18 K, with a dispersion around the mean of 62 K, and an error in the mean of 21 K for their 10 highest quality stars, in the sense that the IRFM scale is the hotter one)."
+ In rez and Moeléónndez (2005b).r the authors derive a number of colour-tomperature relations for main sequence stars (ancl giants) from which estimates of the Sun’s colours are presented in a range of photometric svstenis.," In rez and Melénndez (2005b), the authors derive a number of colour-temperature relations for main sequence stars (and giants), from which estimates of the Sun's colours are presented in a range of photometric systems."
+ We came aware of this study while undertaking our own. and eventually mace use of their extended: set o£ IUEM emperatures in order to supplement material we had collected. as described in detail below.," We became aware of this study while undertaking our own, and eventually made use of their extended set of IRFM temperatures in order to supplement material we had collected, as described in detail below."
+ Our method is to it the colours of Sun-like stars as à function of their ει and metallicities from Nordstrómm οἱ al (2004). and then use these fits to solve for the colours of the Sun. for which we adopt Lipp=5777 lx and. ΕΟΓ) =0.," Our method is to fit the colours of Sun-like stars as a function of their $T_{\mathrm
+eff}$ and metallicities from Nordströmm et al (2004), and then use these fits to solve for the colours of the Sun, for which we adopt $\Teff = 5777$ K and [Fe/H] $ = 0$."
+ Stars similar to the Sun have been selected [rom the sample of rez and Melénndez (2005a) in the range 560Tony 6000. in order to bracket the solar cllective eniperature Tapp=OT7T7 dx. resulting in 115 main sequene stars.," Stars similar to the Sun have been selected from the sample of rez and Melénndez (2005a) in the range $5600 < \Teff < 6000$ , in order to bracket the solar effective temperature $\Teff = 5777$ K, resulting in 115 main sequence stars."
+ For 61 of these. accurate metallicities on a homogentSOUS svstenm are taken from: Nordstrómm οἱ als (2004) study. of the kinematics and chemistry of some 40000 nearlw E.G and ]x stars.," For 67 of these, accurate metallicities on a homogeneous system are taken from Nordströmm et al's (2004) study of the kinematics and chemistry of some 000 nearby F, G and K stars."
+ The photometric metallic‘ities of Nordstrómam οἱ al. (, The photometric metallicities of Nordströmm et al. (
+2004) ha| been shown to be in excellent. agreement with spectrosccipic ones.,2004) have been shown to be in excellent agreement with spectroscopic ones.
+ When compared to e.g. Exbvardsson eta. (, When compared to e.g. Edvardsson et al. (
+1993). 110 total dispersion between photometric and spectroscopie FesLE is 1.08 dex.,"1993), the total dispersion between photometric and spectroscopic [Fe/H] is 0.08 dex."
+ From this sample we select stars with a metallicity bracketing the solar onc. in the range O40<« Fejl] |0.0 resulting in 52 stars which form our basic samxe of miün sequence. Sun-like stars.," From this sample we select stars with a metallicity bracketing the solar one, in the range $-0.40 <$ [Fe/H] $<$ +0.40 resulting in 52 stars which form our basic sample of main sequence, Sun-like stars."
+ πο distribution in Zip anc Fe/L] is shown in Figure 1.., Their distribution in $\Teff$ and [Fe/H] is shown in Figure \ref{tefeh}.
+ We collect. from the Literature wide-band Johnson UÜDV. Cousins Z4. Pwo Micron All Sky Survey (2\LASS) JI. Vyeho iy. Dy and. medium-band. Strómmgren ubeg photometry for our sample stars. so as to deermine indirectly. the solar. colours from the ultraviolet. to. the infrared.," We collect from the literature wide-band Johnson $UBV$, Cousins $RI$, Two Micron All Sky Survey (2MASS) $JHK_s$, Tycho $V_{\rm T}$, $B_{\rm T}$ and medium-band Strömmgren $ubvy$ photometry for our sample stars, so as to determine indirectly the solar colours from the ultraviolet to the infrared."
+ For 9 stars of the basic sample. homogeneously measured DVRL photometry is available from Bessel (1990) these are used to interpolate the solar colours involving DVBHlblters.," For 9 stars of the basic sample, homogeneously measured $BVRI$ photometry is available from Bessell (1990) — these are used to interpolate the solar colours involving $BVRI$filters."
+ For 12 stars. CDV. photometry is taken from Johnson et al.(1966) these are used to infer the solar pD colour.," For 12 stars, $UBV$ photometry is taken from Johnson et al.(1966) — these are used to infer the solar $U-B$ colour."
+" For 31 stars. J4LI, data have been collected"," For 31 stars, $JHK_s$ data have been collected"
+temperatures with advancing pulses. along with an increase in the radial extension of the photosphere. which rellect the asvamptotie evolution towards the redder tip of the AGL.,"temperatures with advancing pulses, along with an increase in the radial extension of the photosphere, which reflect the asymptotic evolution towards the redder tip of the AGB."
+ Also evident is the fact that the bulk of mass-loss takes place precisely during this short phase. with (he mass dropping from 1.90 to 0.63Af. - almost its final value.," Also evident is the fact that the bulk of mass-loss takes place precisely during this short phase, with the mass dropping from $1.90$ to $0.63\ \Msun$ - almost its final value."
+ The mass Οἱ the II-depleted core increases during (lis phase from 0.58 to 0.62M.: the mass of the inner lle-depleted. core increases from 0.43 (o 0.54AL.., The mass of the H-depleted core increases during this phase from $0.58$ to $0.62\ \Msun$; the mass of the inner He-depleted core increases from $0.48$ to $0.54\ \Msun$.
+ Since the He profile is not as steep as the II profile. the mass of the Ie-depleted core is a matter of definition: here IIe-depleted means Y«10°.," Since the He profile is not as steep as the H profile, the mass of the He-depleted core is a matter of definition: here 'He-depleted' means $Y<10^{-6}\,$."
+ Taking the core boundary at the mid-point of the He profile vields a final lle-depleted core mass of 0.6047..., Taking the core boundary at the mid-point of the He profile yields a final He-depleted core mass of $0.60\Msun$.
+ We should note that the total number of pulses in each evolutionary sequence is largely determined by (he mass-loss law adopted., We should note that the total number of pulses in each evolutionary sequence is largely determined by the mass-loss law adopted.
+ Using thecomplete evolutionary tracks for the mass range of 0.8 to 9AL.. for both populations Z—0.01 and Z=0.001. as listed in Table 1. we obtain a theoretical IEMB. displaved in Fie.," Using thecomplete evolutionary tracks for the mass range of $0.8$ to $9\ \Msun$, for both populations $Z=0.01$ and $Z=0.001$, as listed in Table 1, we obtain a theoretical IFMR, displayed in Fig."
+ 7 (solid aud dashed black lines)., \ref{fig:ifmr} (solid and dashed black lines).
+ We increased (he number of points by adding results for masses of 2.5 and 3.5M... and for Z=0.02 and Z=0.005 and masses of 1.3 and 5AL. (anarked in Fig.," We increased the number of points by adding results for masses of $2.5$ and $3.5\ \Msun$, and for $Z=0.02$ and $Z=0.005$ and masses of $1,3$ and $5\ \Msun$ (marked in Fig."
+ 7 bv different svinbols)., \ref{fig:ifmr} by different symbols).
+ Similar relationships have been recently computed by ? and by ?.. the latter including earlier results obtained by ?..," Similar relationships have been recently computed by \citet{2007arXiv0710.2397M} and by \citet{2008A&A...477..213C}, the latter including earlier results obtained by \citet{1999ApJ...524..226D}."
+ A dilferent and independent source for such a relationship is provided by observations (e.g.. ?)). mainly of star clusters. which lead to empirical or semi-empirical linear relations. such as ?. (based on open-cluster data for the range 2.5—6.5 AL.) and others that will be mentioned below.," A different and independent source for such a relationship is provided by observations (e.g., \citealt{2000A&A...363..647W}) ), mainly of star clusters, which lead to empirical or semi-empirical linear relations, such as \citet{2005MNRAS.361.1131F} (based on open-cluster data for the range $2.5-6.5\ \Msun$ ) and others that will be mentioned below."
+ The curves obtained here show that the IFAIR may be divided into three regions with different slopes: 1., The curves obtained here show that the IFMR may be divided into three regions with different slopes: 1.
+ A moderate slope for M;3M... which coincides with the tabulated results of ? plotted in Fig. 7..," A moderate slope for $M_i\lesssim 3\ \Msun$, which coincides with the tabulated results of \citet{2000A&A...363..647W} plotted in Fig. \ref{fig:ifmr}. ."
+ 2., 2.
+ A steeper slope for 3M;<4.M.., A steeper slope for $3\lesssim M_i \lesssim 4\ \Msun$.
+ 3., 3.
+ Again. a more eradual increase until the top end.," Again, a more gradual increase until the top end."
+ We note that the new relation. as displaved in Fig., We note that the 'new relation' as displayed in Fig.
+ 2 ol ?.. meant to fit only the best determined stars of the Hyades ancl Pleiades clusters. has a very similar shape to our curves. only shifted upwiuds from our Pop.," 2 of \citet{1995LIACo..32..441H}, meant to fit only the best determined stars of the Hyades and Pleiades clusters, has a very similar shape to our curves, only shifted upwards from our Pop."
+l curve by about ,I curve by about $0.05\ \Msun$.
+The dependence on metallicity is apparent [rom the divergence of the (wo curves in Fie. 7..," The dependence on metallicity is apparent from the divergence of the two curves in Fig. \ref{fig:ifmr},"
+ in agreement with the conclusions of e.g. ? or the ? curves as plotted in Fig., in agreement with the conclusions of e.g. \citet{2007arXiv0710.2397M} or the \citet{1999ApJ...524..226D} curves as plotted in Fig.
+ 5 of ?.. , 5 of \citet{2008A&A...477..213C}. .
+The effect of metallicity is negligible for M;<2AZ. (in agreement with e.g. 2)). but 1 increases towards higher initial masses: the eurves diverge by 20.13M.at the top end," The effect of metallicity is negligible for $M_i\lesssim2\ \Msun$ (in agreement with e.g. \citet{2008A&A...477..213C}) ), but it increases towards higher initial masses: the curves diverge by $\gtrsim0.13\ \Msun$at the top end"
+ealaxy is traveling face-on through the LGAL and they do no take into account anv of the galaxys orbital history.,galaxy is traveling face-on through the IGM and they do not take into account any of the galaxy's orbital history.
+ Compared. to the observed. deficiency the amoun of ram. pressure stripped is very small. indicating tha under the assumed. criteria rani pressure alone would. no contribute to gas loss in a significant wav.," Compared to the observed deficiency the amount of ram pressure stripped is very small, indicating that under the assumed criteria ram pressure alone would not contribute to gas loss in a significant way."
+ Llowever. tida interactions have been observed to remove large amounts of gas from galaxies in groups (Freelanceetal.2009). ane can increase the elfectiveness of ram. pressure stripping.," However, tidal interactions have been observed to remove large amounts of gas from galaxies in groups \citep{2009AJ....138..295F} and can increase the effectiveness of ram pressure stripping."
+ In his scenario. tidal interactions may disturb the gas disk. »ull eas out. and reduce the column density enabling even weak ram pressure to strip olf low column density gas in arge quantities.," In this scenario, tidal interactions may disturb the gas disk, pull gas out, and reduce the column density enabling even weak ram pressure to strip off low column density gas in large quantities."
+ LO is quite possible that UCC 07049 has undergone such stripping in this eroup given the richness of the svstem. a hostile X-ray emitting dense LGAL. lower han normal gas content. a truncated ancl asvmmetric disk. and the disturbed stellar disk.," It is quite possible that UGC 07049 has undergone such stripping in this group given the richness of the system, a hostile X-ray emitting dense IGM, lower than normal gas content, a truncated and asymmetric disk, and the disturbed stellar disk."
+ In addition to these. the »ossibilitv of this galaxy to have passed through the group core and thereby experience a higher IGM density cannot be ruled out.," In addition to these, the possibility of this galaxy to have passed through the group core and thereby experience a higher IGM density cannot be ruled out."
+ A trip through the core of this group. experiencing intergalactic gas densities nearly ten times higher. can ram pressure strip enough gas to explain the deficiency only if the galaxy velocity is as high as ~SOOkims," A trip through the core of this group, experiencing intergalactic gas densities nearly ten times higher, can ram pressure strip enough gas to explain the deficiency only if the galaxy velocity is as high as $\sim 800 \kms$."
+ COC 07049 has a 1420. Mllz radio continuum Lux censity of 10.9 my., UGC 07049 has a $1420$ MHz radio continuum flux density of 10.9 mJy.
+ At a distance of 100.7. Alpe this corresponds to an upper limit on the star formation rate of 7M.ve|! using the simple relation in Yunetal.(2001).., At a distance of 100.7 Mpc this corresponds to an upper limit on the star formation rate of $7\ \msol\ \mr{yr}^{-1}$ using the simple relation in \citet{2001ApJ...554..803Y}.
+ This assumes a Salpeter initial mass function integrated over stars with masses ranging from 0.1M. to LOOAL.., This assumes a Salpeter initial mass function integrated over stars with masses ranging from $0.1 \msol$ to $100 \msol$.
+ With a measured mass of 2.10*M. and a star formation rate of 7AL.ve+. UGC 07049 will exhaust its gas in 300 Myr.," With a measured mass of $2\times 10^9\ \msol$ and a star formation rate of $7\ \msol\ \mr{yr}^{-1}$, UGC 07049 will exhaust its gas in $\sim 300$ Myr."
+ For comparison. the crossing time for this group. assuming a cliameter of| Alpe. is ~2 Civr.," For comparison, the crossing time for this group, assuming a diameter of 1 Mpc, is $\sim 2$ Gyr."
+ Thus. the deficiency seen in UGC 07049 may also be a result of an elevated star formation rate.," Thus, the deficiency seen in UGC 07049 may also be a result of an elevated star formation rate."
+ We examine two objects in the NGC 4065 group of galaxies: a bent-double radio source and an deficient spiral galaxy., We examine two objects in the NGC 4065 group of galaxies: a bent-double radio source and an deficient spiral galaxy.
+ We use X-ray observations and the bent-double racio source to probe the densitv of intergalactic gas in this group., We use X-ray observations and the bent-double radio source to probe the density of intergalactic gas in this group.
+ For the deficient spiral. UGC 07049. we caleulate the cllectiveness of ram pressure stripping as a gas removal mechanism and find that it alone is not strong cnough to produce the amount of deficiency that is observed.," For the deficient spiral, UGC 07049, we calculate the effectiveness of ram pressure stripping as a gas removal mechanism and find that it alone is not strong enough to produce the amount of deficiency that is observed."
+ An elevated star-formation rate is also observed in this galaxy., An elevated star-formation rate is also observed in this galaxy.
+ A combination of tidal and. ram. pressure stripping. with help from the elevated star-formation rate. are likely strong enough to produce the observed. deficiency.," A combination of tidal and ram pressure stripping, with help from the elevated star-formation rate, are likely strong enough to produce the observed deficiency."
+ We thank the stall of the CMICE who have made. these observations possible., We thank the staff of the GMRT who have made these observations possible.
+ CGMICE is run by the National Centre [or Raclio AXstrophysies of the Tata Institute of Fundamental Research., GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research.
+ This research. has made use of the NASA/LPAC extragalactic database (NED) which is operated by the Jet Propulsion Laboratory. Caltech. under. contract with the National Acronautics and Space Acdministration.," This research has made use of the NASA/IPAC extragalactic database (NED) which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration."
+"The main camera available lor this project was the UCLA Twin Channel Lntrarecl Camera (known locally as Gemini) (McLeanefaf,1993) which mounts onto the Lick 3-1 telescope.",The main camera available for this project was the UCLA Twin Channel Infrared Camera (known locally as Gemini) \citep{McL93} which mounts onto the Lick 3-m telescope.
+ Ou the 3-13. Gemini’s two 2256 arrays have identical {fields of view. one array covering the wavelength range [rom 1 to 2.5 wwith a HeCdTe array. and the other array covering the wavelength range [rom 1 to 2 wwith au InSb array.," On the 3-m, Gemini's two 256 arrays have identical fields of view, one array covering the wavelength range from 1 to 2.5 with a HgCdTe array, and the other array covering the wavelength range from 1 to 5 with an InSb array."
+" The camera employs the subarray techuique for readout when imagine at waveleugthis lounger than 3pan.. outputting 61<661 pixels giving it a field of view of ouly H5""."," The camera employs the subarray technique for readout when imaging at wavelengths longer than 3, outputting 64 pixels giving it a field of view of only."
+"..With a field of view this small it would be impossible to mosaic the required area of the ""dark spot which covers the equivalent area of 20 [ull moous."," With a field of view this small it would be impossible to mosaic the required area of the “dark spot"" which covers the equivalent area of 20 full moons."
+ The solution that was employed was to image the area at Ix with a different camera which due to the lower sky background at Ix. could lave a much larger field of view. then choose the brightest stars [rom that sample to re-image at 3.5 wwith Gemini ou the 3-11. The camera available for this approach was the Lick Πήγασος Camera 2 (LIRC2).," The solution that was employed was to image the area at K with a different camera which due to the lower sky background at K, could have a much larger field of view, then choose the brightest stars from that sample to re-image at 3.5 with Gemini on the 3-m. The camera available for this approach was the Lick InfraRed Camera 2 (LIRC2)."
+ LIRC? could mount onto the 1-12 telescope at Lick aud have a field of view ofxx71.29/.. easily allowing the [ square degrees to be mapped in a short period of time.," LIRC2 could mount onto the 1-m telescope at Lick and have a field of view of, easily allowing the 4 square degrees to be mapped in a short period of time."
+ This was doue over the nights of 1997. March 28 aud 29.," This was done over the nights of 1997, March 28 and 29."
+ The 302 stars catalogued in the LIRC2 survey are shown in Figure 3.., The 302 stars catalogued in the LIRC2 survey are shown in Figure \ref{fig:cold-spot-list}.
+ The process was complicated curing the reduction of the LIRC2 data when it was discovered that LIRC2 had photometric problems which precluded highly accurate photometry., The process was complicated during the reduction of the LIRC2 data when it was discovered that LIRC2 had photometric problems which precluded highly accurate photometry.
+ The accuracy was still sufficient to estimate relative brightnesses. which was the main result that was required ol the data.," The accuracy was still sufficient to estimate relative brightnesses, which was the main result that was required of the data."
+" Having obtaiued the list of stars Grom the brightest to the diminest in the ""dark spot” we could uow proceed with imaging them iu the L ban.", Having obtained the list of stars from the brightest to the dimmest in the “dark spot” we could now proceed with imaging them in the L band.
+ We had three clear aid. photometric uights at the Lick 3-11 with Gemini ou the nights of 1999. March 26 to 28.," We had three clear and photometric nights at the Lick 3-m with Gemini on the nights of 1999, March 26 to 28."
+ Using Cemini’s twin arrays we imagecl our list of stars in both the Ix baud aud the L band simultaneously., Using Gemini's twin arrays we imaged our list of stars in both the K band and the L band simultaneously.
+ The data were reduced using standard data reduction routines with the software package., The data were reduced using standard data reduction routines with the software package.
+ Since the sky brightuess varies dramatically over very short. periods of time at 3.5jun... dilfereucius of sequential images was used in the data reduction to remove detector signature aud sky background.," Since the sky brightness varies dramatically over very short periods of time at 3.5, differencing of sequential images was used in the data reduction to remove detector signature and sky background."
+ Each star was imaged at 1 different locatious on the array. while the LO brightest stars were imaged a total of 12 titmes.," Each star was imaged at 4 different locations on the array, while the 10 brightest stars were imaged a total of 12 times."
+of the sky excluded from the survey: the final mask was the union of the following areas: The mask is specified as a list of exeluded Ldeg? dune-bins’ (ES Ap.,of the sky excluded from the survey; the final mask was the union of the following areas: The mask is specified as a list of excluded $1 \degs$ `lune-bins' (ES Ap.
+X.1). so these values are necessary averages over 1deg.,"X.1), so these values are necessary averages over $1\degs$."
+ The overall area outside the mask is of the sky., The overall area outside the mask is of the sky.
+" For statistical studies of the LRAS galaxy. population and its distribution. where uniformity is more important. than sky coverage. we made a ‘high |b| mask. as above but including all areas with wie 1""'. and leaving οἱ the sky."," For statistical studies of the IRAS galaxy population and its distribution, where uniformity is more important than sky coverage, we made a `high $|b|$ ' mask, as above but including all areas with $A_B>1^m$, and leaving of the sky."
+ In practice. this criterion. is almost identical to Hio212.5Mvster," In practice, this criterion is almost identical to $I_{100} > 12.5 \MJy\pster$."
+" llenceforth. when we say ""high-latitucle’. we simply mean outside this mask."," Henceforth, when we say `high-latitude', we simply mean outside this mask."
+ Both masks are shown in Figure 1., Both masks are shown in Figure 1.
+ Our aim was to relax the eriteria of the QLGC sullicientby to pick up virtually all galaxies. even at low latitude. purely from their LRAS properties: simultaneously. we wanted to keep contamination by Galactic sources to a reasonable level.," Our aim was to relax the criteria of the QIGC sufficiently to pick up virtually all galaxies, even at low latitude, purely from their IRAS properties; simultaneously, we wanted to keep contamination by Galactic sources to a reasonable level."
+ Our actual colour selection criteria were almost the same as the QUOC: Llowever. unlike the OIGC' upper limits were used only where they guaranteed. inclusion or exclusion.," Our actual colour selection criteria were almost the same as the QIGC: However, unlike the QIGC, upper limits were used only where they guaranteed inclusion or exclusion."
+ We did. not at this stage exclude any source. solely on the basis of an identification with a Galactic source., We did not at this stage exclude any source solely on the basis of an identification with a Galactic source.
+ Because the PSCz evolved. from the QICGC. there remain in the catalogue 3 galaxies which actually fail the new selection criteria.," Because the PSCz evolved from the QIGC, there remain in the catalogue 3 galaxies which actually fail the new selection criteria."
+ In total. we selected 16422 sources from the PSC.," In total, we selected 16422 sources from the PSC."
+ Alany very nearby galaxies have multiple PSC sources. associated. with individual starformüng regions.," Many very nearby galaxies have multiple PSC sources, associated with individual starforming regions."
+ TO such sources were pruned to leave for each such galaxy. the single. brightest PSC source.," 70 such sources were pruned to leave for each such galaxy, the single, brightest PSC source."
+ Also at this stage. Local Croup galaxies were excised from. the catalogue. ancl a separate catalogue of [ar-infrared. properties for Local Group galaxies macle.," Also at this stage, Local Group galaxies were excised from the catalogue, and a separate catalogue of far-infrared properties for Local Group galaxies made."
+ Alb URAS surveys are bedevilled by the question of how to deal with galaxies which are multiple or extended. with respect to the LRAS 607422. beam., All IRAS surveys are bedevilled by the question of how to deal with galaxies which are multiple or extended with respect to the IRAS $60\mum$ beam.
+" The size of most. 60r detectors was 1.5/ in-scan by 4.75"" eross-scan.", The size of most $60\mum$ detectors was $1.5'$ in-scan by $4.75'$ cross-scan.
+ The raw data was taken every 0.5’. and for the PSC this was then filtered with an S-point zero-sum linear filter. of the form LdL| .," The raw data was taken every $0.5'$, and for the PSC this was then filtered with an 8-point zero-sum linear filter, of the form $--++++--$ ."
+ Hence galaxies with in-scan [ar-infrared sizes larger than about 1.5%. will have their Düxes underestimated in the PSC.," Hence galaxies with in-scan far-infrared sizes larger than about $1.5'$, will have their fluxes underestimated in the PSC."
+ There is a lesser sensitivity. to cross-scan diameter. caused by some of the scans only. partially crossing the full width exteneded sources.," There is a lesser sensitivity to cross-scan diameter, caused by some of the scans only partially crossing the full width exteneded sources."
+" The approach we have settled on is to preferentially use PSC Ηχος, except for sources identified with individual ealaxics whose large cüameters are likely to lead to significant Hux underestimation in the PSC."," The approach we have settled on is to preferentially use PSC fluxes, except for sources identified with individual galaxies whose large diameters are likely to lead to significant flux underestimation in the PSC."
+ The PSC tux is based. on a least chi-squared template fit to the point-source-Liltered cata-stream. so the [ux is correctly measured for slightly. extended sources. with diameters smaller. than about 1.5/5.," The PSC flux is based on a least chi-squared template fit to the point-source-filtered data-stream, so the flux is correctly measured for slightly extended sources, with diameters smaller than about $1.5'$."
+ As far-infrared diameters are typically half optical. (e.g. Ice 1988). and optical Dos diameters are typically several times the FWILAL for spirals. we can expect that isolated galaxies must have extinetion-correctec diameters larger than several aremin for their Huxes to be badly underestimated by the PSC.," As far-infrared diameters are typically half optical, (e.g. Rice 1988), and optical $D_{25}$ diameters are typically several times the FWHM for spirals, we can expect that isolated galaxies must have extinction-corrected diameters larger than several arcmin for their fluxes to be badly underestimated by the PSC."
+ Fluxes for extended. sources can be derived. using the ADDSCAN (or SCANDPLID software provided at LPAC. which coadds all detector scans passing over a given position.," Fluxes for extended sources can be derived using the ADDSCAN (or SCANPI) software provided at IPAC, which coadds all detector scans passing over a given position."
+" We have tested how the ratio of PSC-to-extended [ux depencds on Des. and. as expected. we find that the ratio is almost constant for galaxies up to 2.5"" (Figure 2). where the adcdscan Dux is typically larecr."," We have tested how the ratio of PSC-to-extended flux depends on $D_{25}$, and, as expected, we find that the ratio is almost constant for galaxies up to $2.5'$ (Figure 2), where the addscan flux is typically larger."
+ However. we also find. that aclelscan LHuxes are svstematically larger than PSC lluxes. even for much smaller galaxies.," However, we also find that addscan fluxes are systematically larger than PSC fluxes, even for much smaller galaxies."
+ This discrepancy is not due to any difference in calibration (we have tested this by extracting à point-source-filtered Dux from the addscan data and substituting this for the PSC llux. but the discrepancy in Figure 2 remains).," This discrepancy is not due to any difference in calibration (we have tested this by extracting a point-source-filtered flux from the addscan data and substituting this for the PSC flux, but the discrepancy in Figure 2 remains)."
+ Phe reason is the large number of multiple and/or interacting ealaxics seen by HUS. for which the addscan picks up the entire combined flux.," The reason is the large number of multiple and/or interacting galaxies seen by IRAS, for which the addscan picks up the entire combined flux."
+ Using such, Using such
+survey!'.,.
+. As a result the average flux of the observed sample of clusters 1s biased high compared to the mean of the parent population (in particular near the flux limit of the survey)., As a result the average flux of the observed sample of clusters is biased high compared to the mean of the parent population (in particular near the flux limit of the survey).
+ To account for Malmquist bias we follow the procedure outlined in Appendix A.2 of ? and correct the X-ray luminosities before fitting the Maso)—Ly. relation., To account for Malmquist bias we follow the procedure outlined in Appendix A.2 of \cite{Vikhlinin09a} and correct the X-ray luminosities before fitting the $M_{2500}-L_X$ relation.
+ We find that the correction for Malmquist bias is relatively modest. increasing the normalisation My by ~5% and reducing the slope a by ~5%!7.," We find that the correction for Malmquist bias is relatively modest, increasing the normalisation $M_X$ by $\sim
+5\%$ and reducing the slope $\alpha$ by $\sim 5\%$."
+. We find an intrinsic scatter of Moorjar=0.23045 (or a relative error of~ 70%). in good agreement with other studies.," We find an intrinsic scatter of $\sigma_{\log
+ L_X|M}=0.23^{+0.10}_{-0.04}$ (or a relative error of$\sim 70\%$ ), in good agreement with other studies."
+" ? lista somewhat larger scatter of 9j,;,|a;=0.29 for the HIFLUGCS sample. whereas ? and ?. find σα,=0.173:0.02 and 0.18+0.02. respectively."," \cite{Reiprich02} list a somewhat larger scatter of $\sigma_{\log L_X|M}=0.29$ for the HIFLUGCS sample, whereas \cite{Vikhlinin09a} and \cite{Mantz10b} find $\sigma_{\log L_X|M}=0.17\pm0.02$ and $0.18\pm0.02$, respectively."
+ Our results also agree well with ο who find ση=0.190.03. which is in good agreement with our value of ," Our results also agree well with \cite{Stanek06} who find $\sigma_{\log
+ M|L_X}=0.19\pm0.03$, which is in good agreement with our value of $0.17^{+0.04}_{-0.03}$."
+The likelihood contours for the power law 0.17*56:.slope o of the Meso)—Ly. relation and the mass Mesoo of a cluster with a luminosity of Ly=2«I0 ASere/s are shown in Figure δ.., The likelihood contours for the power law slope $\alpha$ of the $M_{2500}-L_X$ relation and the mass $M_{2500}$ of a cluster with a luminosity of $L_X=2\times 10^{44} h_{70}^{-2}$ erg/s are shown in Figure \ref{slope}.
+ For the combined sample of clusters we find a best fit slope of à=0.68£0.07 and a normalization My=(1.2£0.12)x105M ..., For the combined sample of clusters we find a best fit slope of $\alpha=0.68\pm0.07$ and a normalization $M_X=(1.2\pm0.12)\times 10^{14} h_{70}^{-1}\msun$ .
+ The inferred slope is consistent with the value of 3 expected for self-similar evolution (?).., The inferred slope is consistent with the value of $\frac{3}{4}$ expected for self-similar evolution \citep{Kaiser86}.
+ It 15 also interesting to compare the constraints from both sub-samples separately., It is also interesting to compare the constraints from both sub-samples separately.
+ The best fit slope for the clusters studied in? is à20.557040 with a normalization of (indicated by the pink, The best fit slope for the clusters studied in \cite{Hoekstra07} is $\alpha=0.55^{+0.10}_{-0.09}$ with a normalization of (indicated by the pink.
+ The blue contours in ..Figure 8 indicate the constraints for the HST sample studied here.," The blue contours in Figure \ref{slope}
+ indicate the constraints for the HST sample studied here."
+ The parameters are. not well constrained. with best fit values a=0.630.24 and a normalization My=(1.040.24«I0/5IM...," The parameters are not well constrained, with best fit values $\alpha=0.63\pm 0.24$ and a normalization $M_X=(1.0\pm0.24\times 10^{14}h_{70}^{-1}\msun$."
+ The difference in the constraints from the (extended) ? sample and the 160SD systems may hint at a deviation from a single power law relation., The difference in the constraints from the (extended) \cite{Hoekstra07} sample and the 160SD systems may hint at a deviation from a single power law relation.
+ As discussed in 822.5. however. the uncertainty in the position of the cluster center leads to an underestimate of the cluster mass. as does the presence of substructure.," As discussed in 2.5, however, the uncertainty in the position of the cluster center leads to an underestimate of the cluster mass, as does the presence of substructure."
+ The ? results are much less sensitive to these problems., The \cite{Hoekstra07} results are much less sensitive to these problems.
+ To quantify this. we combine the ? results with the 160SD clusters with QgcG=2 (12 systems) and the ones with Ὄμου<2 (13 systems).," To quantify this, we combine the \cite{Hoekstra07}
+ results with the 160SD clusters with $Q_{\rm BCG}=2$ (12 systems) and the ones with $Q_{\rm BCG}<2$ (13 systems)."
+ For the former we find My2(1.3030.15)«ΙΟ24M.. and α=0.6340.08. whereas requiring Osco«2 yields My=(1.190.14)«10745]M.. and a=0.68+0.08.," For the former we find $M_X=(1.30\pm0.15)\times 10^{14}h_{70}^{-1}\msun$ and $\alpha=0.63\pm0.08$, whereas requiring $Q_{\rm BCG}<2$ yields $M_X=(1.19\pm0.14)\times 10^{14}h_{70}^{-1}\msun$ and $\alpha=0.68\pm0.08$ ."
+ This comparison suggests that our normalisation may be biased low by as much as ~10%., This comparison suggests that our normalisation may be biased low by as much as $\sim 10\%$ .
+ The relation between the X-ray luminosity and cluster mass has been studied extensively., The relation between the X-ray luminosity and cluster mass has been studied extensively.
+ In this section we compare our measurements to a number of recent results. which are shown in Figure 9..," In this section we compare our measurements to a number of recent results, which are shown in Figure \ref{zoom}."
+ Where needed. the X-ray luminosities are converted to the 0.1—2.4 keV band and Eqn.," Where needed, the X-ray luminosities are converted to the $0.1-2.4$ keV band and Eqn."
+ 10 has been used toconvert masses to M»soo and adjust the slopes (because the relation between M»soo and M00 (or Msoo) i$ a power law with a slope less than 1.)., \ref{mcrel} has been used toconvert masses to $M_{2500}$ and adjust the slopes (because the relation between $M_{2500}$ and $M_{200}$ (or $M_{500}$ ) is a power law with a slope less than 1.).
+ ? studied a sample of 63 clusters. with masses derived under the assumption of hydrostatic equilibrium.," \cite{Reiprich02} studied a sample of 63 clusters, with masses derived under the assumption of hydrostatic equilibrium."
+" Their BCES bisector results for the flux-limited sample yields à20.60+0.05 and M,2(1.32:0.09)«10.. are indicated by the blue triangle in Fig. 9.. 2.. "," Their BCES bisector results for the flux-limited sample yields $\alpha=0.60\pm0.05$ and $M_x=(1.3\pm0.09)\times10^{14} \msun$ are indicated by the blue triangle in Fig. \ref{zoom}. \cite{Stanek06}, ,"
+however. haveM argued that these results suffer from Malmquist bias.," however, have argued that these results suffer from Malmquist bias."
+ Instead they compared the X-ray number counts to the mass function in ACDM cosmologies and derived à=0.54+0.02 and a high normalization of My=(2.140.1)ς10M..., Instead they compared the X-ray number counts to the mass function in $\Lambda$ CDM cosmologies and derived $\alpha=0.54\pm0.02$ and a high normalization of $M_X=(2.1\pm0.1)\times 10^{14}\msun$.
+ This result. however depends strongly on the adopted value for os. and combination with the WMAP3 data (?) lowers the normalization to My=(1.56+0.08)«10M... (open pink triangle in Fig. 9)).," This result, however depends strongly on the adopted value for $\sigma_8$, and combination with the WMAP3 data \citep{Spergel07} lowers the normalization to $M_X=(1.56\pm0.08)\times10^{14}\msun$ (open pink triangle in Fig. \ref{zoom}) )."
+ ? studied a sample of ς~0.05 and ς~0.5 clusters that were observed with Chandra.," \cite{Vikhlinin09a} studied a sample of $z\sim
+0.05$ and $z\sim 0.5$ clusters that were observed with Chandra."
+" Their results with &=0.55+0.05 and M,=(1.4340.08)«10M are indicated by the open orange square.", Their results with $\alpha=0.55\pm0.05$ and $M_x=(1.43\pm0.08)\times10^{14}\msun$ are indicated by the open orange square.
+" Another recent study.. was presented by ? who found a fairly steep slope of a20.70x:0.04 and (1.42+0.28)«I0"".. (purple square).", Another recent study was presented by \cite{Mantz10b} who found a fairly steep slope of $\alpha=0.70\pm0.04$ and $M_x=(1.42\pm0.28)\times10^{14}\msun$ (purple square).
+ The constraints fromM these studies are largely driven by X- luminous clusters (this is particularly true for the results of ? and ?)), The constraints from these studies are largely driven by X-ray luminous clusters (this is particularly true for the results of \cite{Vikhlinin09a} and \cite{Mantz10b}) ).
+ If the slope of the M»so&—Ly varies with Ly. then it is more appropriate to compare these studies to the results from the samplestudied in ?.. for which the agreement is indeed very good.," If the slope of the $M_{2500}-L_X$ varies with $L_X$ , then it is more appropriate to compare these studies to the results from the samplestudied in\cite{Hoekstra07}, , for which the agreement is indeed very good."
+ When combined with the results from the 160SD survey. the resulting scaling relation has a lower normalisation and. steeper slope.," When combined with the results from the 160SD survey, the resulting scaling relation has a lower normalisation and steeper slope."
+ To examine whether this, To examine whether this
+"synchrotron spectrum (equation 2)) and over the electron distribution (equation 4)). we obtain that Therefore. the typical electron. Lorentz factor in the shocked. [uid is where c;—6/10 (the notation convention Q,,=Q/10"" will be use hereafter).","synchrotron spectrum (equation \ref{syspek}) ) and over the electron distribution (equation \ref{dndg}) ), we obtain that Therefore, the typical electron Lorentz factor in the shocked fluid is where $x_1=x/10$ (the notation convention $Q_n = Q/10^n$ will be use hereafter)."
+ The observed peak-energy of the svnchrotron spectrum. is where the [actor 4/3 accounts for the [lux-weighted average frequency (from equation 7... Doppler boost decreases [rom D=2L ο 06— 0:10 —Da6-pt! and goes asvmptoticallv to zero for 6» x). z is the burst redshift and is the comoving-[rame peak-energy. of the svnchrotron spectrum. B being the magnetic field in the shocked. uid and ερ) a factor which a weak dependence on the electron index p.," The observed peak-energy of the synchrotron spectrum is where the factor 4/3 accounts for the flux-weighted average frequency (from equation \ref{D}, , Doppler boost decreases from ${\cal D} = 2\Gamma$ at $\theta=0$ to ${\cal D} = \Gamma$ at $\theta = \Gamma^{-1}$ and goes asymptotically to zero for $\theta \rightarrow \pi$ ), $z$ is the burst redshift and is the comoving-frame peak-energy of the synchrotron spectrum, $B$ being the magnetic field in the shocked fluid and $\psi(p)$ a factor which a weak dependence on the electron index $p$."
+" By integrating the svnchrotron emissivity per electron over the power-law electron distribution. Wijers CGalama (1999) fined that 0(5.2)=0.34. therefore Using equation (16)) and the observed c. this leads lo The received. Εαν at the peak of the svnchrotron spectrum. ts where D;(2) is the burst luminosity distance. the factor | accounts for relativistic beaming of the burst. emission (over the region @< Lot. the specific flux is beamed by a factor D. but that. regions has an area which is a fraction LE of the entire emitting surface. assuming spherical symmetry) and is the comoving-frame luminosity at ος, and IN is the number of radiating electrons (for a spherically symmetric outllow)."," By integrating the synchrotron emissivity per electron over the power-law electron distribution, Wijers Galama (1999) find that $\psi(5.2) = 0.34$, therefore Using equation \ref{gp}) ) and the observed $\epsilon_{p,sy}$, this leads to The received flux at the peak of the synchrotron spectrum is where $D_L(z)$ is the burst luminosity distance, the factor $\Gamma$ accounts for relativistic beaming of the burst emission (over the region $\theta < \Gamma^{-1}$ , the specific flux is beamed by a factor $\Gamma^3$, but that regions has an area which is a fraction $\Gamma^{-2}$ of the entire emitting surface, assuming spherical symmetry) and is the comoving-frame luminosity at $\epsilon'_{p,sy}$ and $N$ is the number of radiating electrons (for a spherically symmetric outflow)."
+ The factor ó(p) caleulated by Wijers Calama (1999) is 6(5.2)=0.10., The factor $\phi(p)$ calculated by Wijers Galama (1999) is $\phi(5.2)=0.70$.
+" The observed. GIU and optical Dux at any time may be the superposition of many (η) emission episodes. (sub-pulses resulting from internal collisions in the outflow) in which a smaller number (CN, ) of electrons radiate: Ν—iN, "," The observed GRB and optical flux at any time may be the superposition of many $\eta$ ) emission episodes (sub-pulses resulting from internal collisions in the outflow) in which a smaller number $N_e$ ) of electrons radiate: $N = \eta\, N_e$."
+The superposition. of these sub-pulses. leads to intrinsic luctuations in the burst light-curve. of relative anyplitucle y0.," The superposition of these sub-pulses leads to intrinsic fluctuations in the burst light-curve, of relative amplitude $\eta^{-1/2}$."
+ M the observing time resolution were shorter than he sub-pulse duration then the amplitude of the GRB light-curve fluctuations (which includes Poisson noise in addition ο the source I[uctuations) would represent an upper limit or the amplitude of the intrinsic source [lluctuations., If the observing time resolution were shorter than the sub-pulse duration then the amplitude of the GRB light-curve fluctuations (which includes Poisson noise in addition to the source fluctuations) would represent an upper limit for the amplitude of the intrinsic source fluctuations.
+ In hat case. the less than Hluctuations of GRB 990123 displaved in figure 1 of Briggs (1999) indicate that ηx100.," In that case, the less than fluctuations of GRB 990123 displayed in figure 1 of Briggs (1999) indicate that $\eta \simg 100$."
+" However. the sub-pulse duration which we obtain xlow. of about of=20 ms. is a factor. 10 less than the vpical /,,,:=256 ms resolution of BATSE light-curves."," However, the sub-pulse duration which we obtain below, of about $\delta t = 20$ ms, is a factor 10 less than the typical $t_{res} = 256$ ms resolution of BATSE light-curves."
+" In this case. there would be INS,=ywfOl)2105 sub-oulses in a GARB pulse."," In this case, there would be $N_{sp} = \eta (t_{res}/\delta t) \simeq 10\, \eta$ sub-pulses in a GRB pulse."
+ Phe resulting amplitude of source luctuations. iV.“is upper bound by the observed. &10/4 luctuations of GRB 990123. hence y10 sullices to accommodate the observed. Ductuations.," The resulting amplitude of source fluctuations, $N_{sp}^{-1/2}$, is upper bound by the observed $\siml 10\%$ fluctuations of GRB 990123, hence $\eta \simg 10$ suffices to accommodate the observed fluctuations."
+ Normalizing 7 to 10. the svnchrotron peak flux is The observed. synchrotron. peak Dux. ον and BU. from equation (20)) lead to Because the svnehrotron ancl inverse-C'ompton spectra are similar. the ratio Y of the inverse-Compton and svnchrotron radiating powers is using equation (15)).," Normalizing $\eta$ to 10, the synchrotron peak flux is The observed synchrotron peak flux $F_{p,sy}$ and $B\Gamma$ from equation \ref{BG}) ) lead to Because the synchrotron and inverse-Compton spectra are similar, the ratio $Y$ of the inverse-Compton and synchrotron radiating powers is using equation \ref{eicsy}) )."
+" At the same time. the Compton parameter is the integral over the electron distribution of the average increase in the photon energy through scattering: where 7, is the sub-pulse optical thickness to electron scattering and ταςxαλαν5U."," At the same time, the Compton parameter is the integral over the electron distribution of the average increase in the photon energy through scattering: where $\tau_e$ is the sub-pulse optical thickness to electron scattering and $d\tau_e/d\gamma
+\propto dN/d\gamma \propto \gamma^{-p}$."
+ The above two equations Iead to using the svnchrotron and peak Duxes of GRB 990123., The above two equations lead to using the synchrotron and peak fluxes of GRB 990123.
+ The electron optical thiekness is set by the electron column density: where er is the Thomson cross-section., The electron optical thickness is set by the electron column density: where $\sigma_T$ is the Thomson cross-section.
+ The last. two equations allowthe determination of the radius a which the burst and. prompt optical emissions are released:The last observational constraint το be used is equation (5))., The last two equations allowthe determination of the radius a which the burst and prompt optical emissions are released:The last observational constraint to be used is equation \ref{gc}) ).
+" Phe radiative cooling timescale of the 5,-electrons is where /£(5,), is. the comoving-frame.⋅ cooling. timescale.", The radiative cooling timescale of the $\gamma_p$ -electrons is where $t'_c (\gamma_p)$ is the comoving-frame cooling timescale.
+. Substituting £ [rom equation (20)) and Y from. equation (25). lends to," Substituting $B$ from equation \ref{BG}) ) and $Y$ from equation \ref{Y}) ), leads to"
+the complicated morphology of the maps suggests a combination of shocks with different geometry aud seeu [rou different angles.,the complicated morphology of the maps suggests a combination of shocks with different geometry and seen from different angles.
+ For HHT. the well-defined bow-shaped structure probably represents a simplified situation.," For HH7, the well-defined bow-shaped structure probably represents a simplified situation."
+ Smith et ((2003) proposed that a bow shock moving at an angle of ~30° to the line-of-sight is consistent with the He liue profiles observed at ciffereut positions in HH7., Smith et (2003) proposed that a bow shock moving at an angle of $\sim 30 ^{\circ}$ to the line-of-sight is consistent with the $_2$ line profiles observed at different positions in HH7.
+ 'Though tle physic:il sizes of the shock structures mapped by IRS [or HH7 anc HH51 are <.0.1 times sinaller han those « “SNRs. the best-fi| density. iudex 6 aud ο column clenity do uot sienilficantlv dilfer betweel hese two classes of source.," Though the physical sizes of the shock structures mapped by IRS for HH7 and HH54 are $\lesssim 0.1$ times smaller than those for SNRs, the best-fit density, index $b$ and $_2$ column density do not significantly differ between these two classes of source."
+ We noted. however. that the OPR for 101 gas (T> 1000 ]x) in HHT atdà HHSS is lower tlan the LTE value —D+) while for a| fou SNRs e departure of tle OPR [roi1 equilibrium is negligible at that hieli temperature.," We noted, however, that the OPR for hot gas $T>$ 1000 K) in HH7 and HH54 is lower than the LTE value $\sim 3$, while for all four SNRs the departure of the OPR from equilibrium is negligible at that high temperature."
+ Tus yhenomelli is also rellectec in the Ha 'otational diagrams. where the zigzag patteru is ore appzwent for the two Herbig-HaὉ objects all is apparent even or the highest rotational levels.," This phenomenon is also reflected in the $_2$ rotational diagrams, where the zigzag pattern is more apparent for the two Herbig-Haro objects and is apparent even for the highest rotational levels."
+ hi OUL moclel. the equilibrium OPR of hot gas in the lour SNRs requires a best-fit n(H)xr that is 0.6 — 3+) orders «of maeuituce larger thau that iiferred for HH? aud HH51.," In our model, the equilibrium OPR of hot gas in the four SNRs requires a best-fit $n$ $\times$$\tau$ that is 0.6 – 3 orders of magnitude larger than that inferred for HH7 and HH54."
+ This difTereice may be caused by a cliflere1 atomic hydrogeu density »(H) within gas around SNRs aud Herbie-Haro objects., This difference may be caused by a different atomic hydrogen density $n$ (H) within gas around SNRs and Herbig-Haro objects.
+ Fast SNR bla1 waves driving interstellar shocks with V;>100 kin 1 can prodice strong ultraviolet eiissioS that are responsible for the dissociation of H» in surrounding regions. While for Herbig-Haro objects the typical shock speeds are observed to be smaller (Herbig Jores 1981:Cohen Fuller 1985).," Fast SNR blast waves driving interstellar shocks with $V_s \geq 100 $ km $^{-1}$ can produce strong ultraviolet emissions that are responsible for the dissociation of $_2$ in surrounding regions, while for Herbig-Haro objects the typical shock speeds are observed to be smaller (Herbig Jones 1981;Cohen Fuller 1985)."
+ Many hieh excitation Lije structure lines which have been detectec previously toward mauy SNRs are faint or absent toward Herbig-Haro objects., Many high excitation fine structure lines which have been detected previously toward many SNRs are faint or absent toward Herbig-Haro objects.
+ In addition to the UV ield produced by nearby [ast shocks. tlie N-ray enission from au SNR interior will also induce clissociation of p'e-sliock gas.," In addition to the UV field produced by nearby fast shocks, the X-ray emission from an SNR interior will also induce dissociation of pre-shock gas."
+ These all iuiply that tle molecular gas associated with Herbig-Haro objects probaby subject to weaker pliXodissociatiou. which results in less H ancl a lower efficiency of para-to-ortli) conversion.," These all imply that the molecular gas associated with Herbig-Haro objects probably subject to weaker photodissociation, which results in less H and a lower efficiency of para-to-ortho conversion."
+ We note. i0wever. tha for measurements of near-infrared Ho ro-vibraional transiious toward Herbig-HaὉ objects. the OPR values obtained are in most cases cousistent with ΠΠ Davis Lioure 1997).," We note, however, that for measurements of near-infrared $_2$ ro-vibrational transitions toward Herbig-Haro objects, the OPR values obtained are in most cases consistent with 3 (Smith, Davis Lioure 1997)."
+ With energy levels lying above 6000 Ix. these vibrationaly-excited liles orleinate mostly [ron the hottes part of the gas. which. probably has a higher atomic fraction.," With energy levels lying above 6000 K, these vibrationally-excited lines originate mostly from the hottest part of the gas, which probably has a higher atomic fraction."
+ These two [actors — high temperature and high atomic H fraction — add up to a fast. elficieut. ortlio-para COLLVerslon.," These two factors – high temperature and high atomic H fraction – add up to a fast, efficient ortho-para conversion."
+ The maps of best-fit parameters i Figures 11 — 19 were derived from the IRS H» fluxes ouly., The maps of best-fit parameters in Figures 14 – 19 were derived from the IRS $_2$ fluxes only.
+ So they may not represeut the real average value at each pixel. but we expect that they carry uxelul information about spatial variatious in the physical coucditious in these regions.," So they may not represent the real average value at each pixel, but we expect that they carry useful information about spatial variations in the physical conditions in these regions."
+ Maps of all six sources exhibit a spatial variation in n(H») larger than a factor 2. aud a variation in 5 larger," Maps of all six sources exhibit a spatial variation in $n$ $_2$ ) larger than a factor 5, and a variation in $b$ larger"
+Fig.l. (,Fig.1. (
+a) (Top-left) Iuage of UCC 1395 in the No band.,a) (Top-left) Image of UGC 1395 in the K' band.
+ Contour levels are 20.6. 20.1. 18.9. 15.6.I8. 1. 17.5. bee17.(- 17.1. 16.3 and 15.Linag/aresec?.," Contour levels are 20.6, 20.4, 18.9, 18.6, 18.4, 17.8, 17.6, 17.1, 16.3 and 15.4 $^2$."
+ The horizoutal scale corresponds to 10 arcsec: edebj(2ud-left) Sharp-«ivided I mage: (ο) CGxrd-ett) Ellise Lsubtracted RC lage: (d) (Botton-lIcft) G-I) color image (not available for UGC 1395): (0) (Top-rigit). PA (open circles) and € (solid triangles) ellipse major axis Gn aresec) m EK., The horizontal scale bar corresponds to 10 arcsec; (b)(2nd-left) Sharp-divided K' image; (c) (3rd-left) Ellipse model subtracted K' image; (d) (Bottom-left) (J-K') color image (not available for UGC 1395); (e) (Top-right) PA (open circles) and $\epsilon$ (solid triangles) ellipse major axis (in arcsec) in K'.
+ The corresponding ctrvoes for J and UST iHALES. when available. are plotted as short-dashecd lines iux thick full lines respectively: (f) (2udaight) Surface brightuess profiles in I (open circles} aud J (solid triangles). when available.," The corresponding curves for J and HST images, when available, are plotted as short-dashed lines and thick full lines respectively; (f) (2nd-right) Surface brightness profiles in K' (open circles) and J (solid triangles), when available."
+ Dulge aud disk componeuts together wih the bulge|disk profile are given as solid lines: (e) (3xd-rneht) Residuals from the bulge|disk decomposition. in percentage = (Fu Fyin)/Pons: sviubols as iu (f): (li) (DBottomaisht) (J-I) color eradieut iieasured in circular res. caleulated. from the cieulu aperture curves of erowtl (not available for UCC 1395).," Bulge and disk components together with the bulge+disk profile are given as solid lines; (g) (3rd-right) Residuals from the bulge+disk decomposition, in percentage = $(F_{obs}-F_{fit})/F_{obs}$; symbols as in (f); (h) (Bottom-right) (J-K') color gradient measured in circular rings, calculated from the circular aperture curves of growth (not available for UGC 1395)."
+ As in all following figures. north is to the top aud cast to the left.," As in all following figures, north is to the top and east to the left."
+ N-axis in Fies. (, X-axis in Figs. (
+ο) to (h) refers to the semiauajor axis (radius) iu arcsecouds.,e) to (h) refers to the semi-major axis (radius) in arcseconds.
+ Fie., Fig.
+ 2. (, 2. (
+a) (Top-left) Tage of IC ts tin the IC band.,a) (Top-left) Image of IC 184 in the K' band.
+ Contour levels are 19.3. 15.1. 17.5. 17.1. 16.5. 16.5. 16.3. 15.5 and 15.0 2ον (," Contour levels are 19.3, 18.1, 17.5, 17.1, 16.8, 16.5, 16.3, 15.8 and 15.0 $^2$. ("
+b). (c). (d) (not available). (0). (£). (8) and (a) (uot available) as in Fig.,"b), (c), (d) (not available), (e), (f), (g) and (h) (not available) as in Fig."
+ 1, 1.
+ Fig., Fig.
+ 3. (, 3. (
+a) CIop-left) huage of IC 1816 in the K baud.,a) (Top-left) Image of IC 1816 in the K' band.
+ Contour levels are 19.5. 19.0. 15.7. 18.3. 18.0. 17.7. 17.3. 17.0. 16.7. 16.5. 16.0. 15.6 aud 11.5 'yesec?. (," Contour levels are 19.5, 19.0, 18.7, 18.3, 18.0, 17.7, 17.3, 17.0, 16.7, 16.5, 16.0, 15.6 and 14.5 $^2$. ("
+b). (c). (d). (0). (£). (2) and (1) as in Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+1, 1.
+ Fig., Fig.
+ L (, 4. (
+a) (Top-deft) huage of UC in the Ik baud.,a) (Top-left) Image of UGC 3223 in the K' band.
+" Coutour evels are 19.G. 18.6. 18.2.u Smle).17"" JL. τι, 16.6. 16.2 and 15.5 /aresc"," Contour levels are 19.6, 18.6, 18.2, 17.9, 17.4, 17.1, 16.6, 16.2 and 15.5 $^2$. ("
+e?. n | (not available). (0). (£). (8) and (hi)(iot available) as iu Fie.,"b), (c), (d) (not available), (e), (f), (g) and (h) (not available) as in Fig."
+ 1., 1.
+ Fig., Fig.
+ 5. (, 5. (
+a) CTop-left) huage of NGC 2639 in the IC baud.,a) (Top-left) Image of NGC 2639 in the K' band.
+" Coutour levels are 18.9. 18.0. 17.3. 16.1. 16.0. 15.6. 15.2. Lis and 112 aresec?t, ("," Contour levels are 18.9, 18.0, 17.3, 16.4, 16.0, 15.6, 15.2, 14.8 and 14.2 $^2$. ("
+by. (0). (6d). (0). (GE). (9) aud (h) as in Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fig., Fig.
+ 6. (, 6. (
+a) (Top-lett) Tage of Ic 2 in the I baud.,a) (Top-left) Image of IC 2510 in the K' band.
+" Contour levels ire 18.7. 18.3. 17.6. 17.2. ""n16.7. 16.3. 15.5 and 11.5 hurcese"," Contour levels are 18.7, 18.3, 17.6, 17.2, 16.7, 16.3, 15.5 and 14.5 $^2$. ("
+"c?, (b). (e). ""me). ( (ο) aud (li) as in Fie.","b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ 7. (, 7. (
+a) CTop-left) huage of NGC 3281 in the IC band.,a) (Top-left) Image of NGC 3281 in the K' band.
+ Coutour levels are 18.5. 18.0. 17.5. 17.0. 16.8. 16.6. 16.1. 15.9. 15.1. 11.7 aud 11.0 yesec?. (," Contour levels are 18.5, 18.0, 17.5, 17.0, 16.8, 16.6, 16.4, 15.9, 15.4, 14.7 and 14.0 $^2$. ("
+by. (0). (6d). (0). (£). Ce) and (a) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ &. (, 8. (
+a) CTop-left) huage of NCC 3660 in the IC baud.,a) (Top-left) Image of NGC 3660 in the K' band.
+ Contour levels are 18.6. 18.2. 17.1. 16.8. 16.1 and 15.9 arcsoect2 (," Contour levels are 18.6, 18.2, 17.4, 16.8, 16.4 and 15.9 $^2$. ("
+by. (ο). GL. (ο). {fy. (g) and (hj) as iu Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ 9. (, 9. (
+a) CTop-left) huage of NCC 1253 in the IC band.,a) (Top-left) Image of NGC 4253 in the K' band.
+" Coutour levels aye 18.7. 17.7. 17.3. 16.7. 16.1. 16.0. 15.2. 11.2 aud 12.3 2'arcsect, ("," Contour levels are 18.7, 17.7, 17.3, 16.7, 16.4, 16.0, 15.2, 14.2 and 12.3 $^2$. ("
+b). (ο). Cd). Ce). (C£). (ο) aud (h) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1., 1.
+ Fie., Fig.
+ 10. (, 10. (
+a) (Top-left) huage of NCC 1507 in the IN? baud.,a) (Top-left) Image of NGC 4507 in the K' band.
+" Coutour levels are 18.9. 18.2. 17.7. 17.5.17.2. 16.8. 16.1. 15.6 and 11.5 nag/arcesec?, ("," Contour levels are 18.9, 18.2, 17.7, 17.5, 17.2, 16.8, 16.1, 15.6 and 14.5 $^2$. ("
+b). Cc). Cd). (ο). (£). (ο) and (h) as » Fis.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+1m, 1.
+ Fie., Fig.
+ nm1l. (, 11. (
+a) (left) Tage of NGC 1785 in the KC baud.,a) (Top-left) Image of NGC 4785 in the K' band.
+" C levels are Ts.l. 17.9. 17.1. 16.8. 16.1. 16.1. 15.9. 15.1 and 11.1 nag/arcesec?, ("," Contour levels are 18.4, 17.9, 17.4, 16.8, 16.4, 16.1, 15.9, 15.4 and 14.4 $^2$. ("
+b). Cc). Cd). (ο). (£). (ο) and (11) as iu Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ 12. (, 12. (
+a) (Toleft) Tage of NGC 5317 in the IC baud.,a) (Top-left) Image of NGC 5347 in the K' band.
+ Contour levels are 19.0. Ls.lL. 17.7. 17.3. 16.5. 16.0 and ↽∖1LS /aresecs.> (," Contour levels are 19.0, 18.4, 17.7, 17.3, 16.5, 16.0 and 14.8 $^2$. ("
+b). (c). (d). fe). CE}. (ο) and (hi) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ 13. (, 13. (
+a) (Top-left) Image of NGC n28 inthe IC baud.,a) (Top-left) Image of NGC 5728 inthe K' band.
+" Contour levels are 18.3. 7.7. 17.0. id).. 16.0. 15.3. ↽Los.∖↽ 11.5 aud ,11.0 mag/arc"," Contour levels are 18.3, 17.7, 17.0, 16.5, 16.0, 15.3, 14.8, 14.5 and 14.0 $^2$. ("
+sec?D(by. text i(e). (£). (5) and (11) as iu Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ LL. (, 14. (
+à) (To]left) huaee of ESO 139-12 in the I baud.,a) (Top-left) Image of ESO 139-12 in the K' band.
+" Coutour levels are 18.9. 18.14. 18.0. Έτ lT.3. 16.9. 16.5. 16.0 aud 11.5 nag/aresec?, ("," Contour levels are 18.9, 18.4, 18.0, 17.7, 17.3, 16.9, 16.5, 16.0 and 14.5 $^2$. ("
+b). Cc). Cd). (ο). (£). (ο) and (11) as iu Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fig., Fig.
+ 15. (, 15. (
+a) (Tolett) huage of NGC o in the id baud.,a) (Top-left) Image of NGC 6814 in the K' band.
+ Contour levels are 18.1. 17.7. 17.1. 17.2. 17.0. 16(c) 16.5. 16.2. 16.0. 15.G. 15.2 aud 11.5 yese," Contour levels are 18.1, 17.7, 17.4, 17.2, 17.0, 16.7, 16.5, 16.2, 16.0, 15.6, 15.2 and 14.5 $^2$. ("
+"c?, (b). (d). (ο). €). (9) zx (h) as iu Fie.","b), (c), (d), (e), (f), (g) and (h) as in Fig."
+""" 1", 1.
+ Fie., Fig.
+ 16. (, 16. (
+a) (Tolett) huage NGC 6860 iu ‘he I baud.,a) (Top-left) Image of NGC 6860 in the K' band.
+" Contour level sare 19.0. 18.5. 17.7.(""17.0. "" 16.5.ud 16.0. 15.1. 11.2 ec. O"," Contour levels are 19.0, 18.3, 17.7, 17.0, 16.5, 16.0, 15.1, 14.2 $^2$. ("
+b}. (e). td). (f). a (h) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+1, 1.
+ Fie., Fig.
+ 17. (, 17. (
+a) (Top-left) Image of NGC 6890 in he Iv’ baud.,a) (Top-left) Image of NGC 6890 in the K' band.
+ Contour level sare 19.1. Is.1. 17-:3. 16.7. 16.1. 16.1. 15.7. 15.1. L8 and L0 /aresce?. (," Contour levels are 19.1, 18.1, 17.3, 16.7, 16.4, 16.1, 15.7, 15.4, 14.8 and 14.0 $^2$. ("
+h). (ο). GL. (o). (f). (2) and (li) as in Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ l8. (, 18. (
+a)(Top-left) me of NGC .. in he K baud.,a) (Top-left) Image of NGC 151 in the K' band.
+" Contour level sare 18.3. 17.7.(e 17.0. 16""n. 16.0. 15.6. 15.1 and 116 Parcse"," Contour levels are 18.3, 17.7, 17.0, 16.5, 16.0, 15.6, 15.1 and 14.6 $^2$. ("
+"ct, (bi. d(cl). { (f). (2) and (li) as in ""","b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fig., Fig.
+ 19. , 19. (
+m uiTop-left) Tage of IC. 15112 the I. baud.,a) (Top-left) Image of IC 454 in the K' band.
+ Contour are 18.23. 17.6. 17.2. 16.9. 16.5. 16.2. 16.0 aud 118 ALCSCCTP (," Contour levels are 18.3, 17.6, 17.2, 16.9, 16.5, 16.2, 16.0 and 14.8 $^2$. ("
+by. 60). GU. (0). C). (8) aud (Bh) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fig., Fig.
+ 20. (, 20. (
+a) (Top-left) Image of NGC 22712 in u I baud.,a) (Top-left) Image of NGC 2712 in the K' band.
+" Coutour levels are 18.5. 18.1. 17.5. if.17.1. ""m16.6. 16.0 and 11858 hwesec?. ("," Contour levels are 18.5, 18.1, 17.5, 17.1, 16.6, 16.0 and 14.8 $^2$. ("
+b). (e). (6d). fe).{ (e) (h) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+ Fie., Fig.
+ 21. (, 21. (
+a) (Top-left) huage of NGC 2811 in the IN? baud.,a) (Top-left) Image of NGC 2811 in the K' band.
+ Contour levels are 18.0. 17.3. 16.8. 16.1. 16.1. 15.9. 15.6. 15.3. 1L7 aud 110 hucsec?. (," Contour levels are 18.0, 17.3, 16.8, 16.4, 16.1, 15.9, 15.6, 15.3, 14.7 and 14.0 $^2$. ("
+ha). (ο). (dj. (0). (£). (2) and (li) as in Fie.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ J, 1.
+ Fig., Fig.
+ 22. (, 22. (
+a) (Top-lett) me of NGC 3571 iu the IK baud.,a) (Top-left) Image of NGC 3571 in the K' band.
+ Contour levels are 18.6. 17.7. 17.2. (b).16.6. 16.3. 16.0. 15.7. 15.1. 1LS aud 110 hurese," Contour levels are 18.6, 17.7, 17.2, 16.6, 16.3, 16.0, 15.7, 15.4, 14.8 and 14.0 $^2$. ("
+"c?, (c). (d). (e). (£). (2) and (h) as in Fig.","b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1., 1.
+ Fig., Fig.
+ 23. (, 23. (
+a) (Top-left) huage of NGC 3835 in the IC baud.,a) (Top-left) Image of NGC 3835 in the K' band.
+ Coutour levels are 19.1. 18.6. 15.1. p.SN. 17.3. 17.0. 16.G. 16.2. 16.0. 15.2 and 18.5 inag/aresec?. (," Contour levels are 19.4, 18.6, 18.1, 17.8, 17.3, 17.0, 16.6, 16.2, 16.0, 15.2 and 18.8 $^2$ . ("
+b). (e).(d). (0). €). (9) and (hi) as iu Fig.,"b), (c), (d), (e), (f), (g) and (h) as in Fig."
+ 1, 1.
+The study of A-ray pusags has been an area ol active research for almost 302 vears.,The study of X-ray pulsars has been an area of active research for almost 30 years.
+ La spite of tiis there remain some significant short[àIs in the understanding of these objects., In spite of this there remain some significant shortfalls in the understanding of these objects.
+ An example is the persistent pulsar GN 1)4., An example is the persistent pulsar GX 1+4.
+ At the time of its discovery (Lewin.RickerandAcClintock it was one of the brighest objects in the X-ray sky., At the time of its discovery \cite{lew71} it was one of the brightest objects in the X-ray sky.
+ The companion to the neutron star is an M6 giant (Davidsen.Alalinaand3owver 1977)., The companion to the neutron star is an M6 giant \cite{dav77}.
+ GN LId s he onlv known X-rav pulsar in a symbiotic system., GX 1+4 is the only known X-ray pulsar in a symbiotic system.
+ Measurements of the average spin-up rate curing the 1970s gave he largest value recorded for any. pulsar (or in fact any astronomical object) ab zm2 per cent per vear., Measurements of the average spin-up rate during the 1970s gave the largest value recorded for any pulsar (or in fact any astronomical object) at $\approx 2$ per cent per year.
+ Inexplicably. the average spin-up trend. reversed. around. 1983. switching to spin-down at approximately the same rate.," Inexplicably, the average spin-up trend reversed around 1983, switching to spin-down at approximately the same rate."
+ Since that reversal a number of changes in the sign of the torque (as inferred. (rom the rate ofcange of the pulsar spin period) have been observed (Chakravarityetal.1997).., Since that reversal a number of changes in the sign of the torque (as inferred from the rate of change of the pulsar spin period) have been observed \cite{chak97b}.
+ Several estimates (Beurle οἱ al., Several estimates (Beurle et al.
+ 1OS4. Dotani et al.," 1984, Dotani et al."
+ 1989. Creenhill et al.," 1989, Greenhill et al."
+" 1993. Cui. 1997) indicate a neutron star surface magnetic field strength. of E3 Lott, The X-ray [lux from the source is extremely variable on inic-scaCS8 O0 seconds to decades."," 1993, Cui 1997) indicate a neutron star surface magnetic field strength of $2-3 \times 10^{13}$ G. The X-ray flux from the source is extremely variable on time-scales of seconds to decades."
+ Two principal Dux states rave been οἱserved. a thigh’ state which persisted during he spin-up »viod of t1ο LOTOs. and a ‘low state since.," Two principal flux states have been observed, a `high' state which persisted during the spin-up period of the 1970s, and a `low' state since."
+ Althoug ithe mean [Hux las been increasing steacilv curing he current ‘low state i has not vet returned to the level of the 1970s., Although the mean flux has been increasing steadily during the current `low' state it has not yet returned to the level of the 1970s.
+ Superimpx«med on these long-term variations are smooth ¢ianges in the flux on time-scales of order hours ο davs., Superimposed on these long-term variations are smooth changes in the flux on time-scales of order hours to days.
+ On he shortes time-scales the periodic variation due to the neutron stars rotation period at around 2 min is observed., On the shortest time-scales the periodic variation due to the neutron star's rotation period at around 2 min is observed.
+Figure 1l shows the polarization map of \lvk 3 with the E275W filter.,Figure \ref{fig_mrk3_PA_L} shows the polarization map of Mrk 3 with the F275W filter.
+ The polarizations are calculated in LO pixel (0.757) bins. and the regions with statistical S/N in P larger than 5 are shown.," The polarizations are calculated in 40 pixel $0.''57$ ) bins, and the regions with statistical S/N in $P$ larger than 5 are shown."
+ The position angele errors indicated in the fleure are the total sumi in quadrature of all the error sources described in the previous section., The position angle errors indicated in the figure are the total sum in quadrature of all the error sources described in the previous section.
+ The position anele (PA) distribution is centrosviunietrie. supporting the idea that the polarization is caused bv the scatteriug of the radiation from a compact unclear source.," The position angle (PA) distribution is centrosymmetric, supporting the idea that the polarization is caused by the scattering of the radiation from a compact nuclear source."
+ The lengths of the lines at cach position 1u the figure indicate the degrees of polarization., The lengths of the lines at each position in the figure indicate the degrees of polarization.
+ We note that in the outermost regions where uncertainty from the backeround subtraction dominates. the polarization degree is highly uncertain accordingh.," We note that in the outermost regions where uncertainty from the background subtraction dominates, the polarization degree is highly uncertain accordingly."
+ However. at least the PA directions are consistent with the ones in the inner regions. suggesting that we have detected the polarization even out to 3” frou the center.," However, at least the PA directions are consistent with the ones in the inner regions, suggesting that we have detected the polarization even out to $\sim
+3''$ from the center."
+ We also show the polavization map on a sinaller scale in Figures 2. and 3.., We also show the polarization map on a smaller scale in Figures \ref{fig_mrk3_P_S275} and \ref{fig_mrk3_P_S342}.
+ The poluizations are calculated in 10 pixel bins (0.711). and onlv. the regions with statistical S/N in P larger than 5 are shown.," The polarizations are calculated in 10 pixel bins $0.''14$ ), and only the regions with statistical S/N in $P$ larger than 5 are shown."
+ The leneths of the lines are proportional to the polarization degrees also in these figures., The lengths of the lines are proportional to the polarization degrees also in these figures.
+" The ceutral ~1"" radius reeion is very ον. aud the PA pattern is centrosyummetric down to this central region. where Pis 1ο20% and 515% level in the F275W aud EF312W filter. respectively (though note the dilution of P by the narrow line contamination: sec refsoc-obs))."," The central $\sim 1''$ radius region is very clumpy, and the PA pattern is centrosymmetric down to this central region, where $P$ is $10-20$ and $5-15$ level in the F275W and F342W filter, respectively (though note the dilution of $P$ by the narrow line contamination; see \\ref{sec-obs}) )."
+ From the observed. ceutrosvummetrie pattern. we can robustly determine the position of tlic hidden uueleus as a svuumetric center of this pattern.," From the observed centrosymmetric pattern, we can robustly determine the position of the hidden nucleus as a symmetric center of this pattern."
+ Specifically. assuminug a ceutrosviinuetric model of the PA distribution. a least square fit is iurpleaneuted using the PA measurements with even errors at cach point of the image.," Specifically, assuming a centrosymmetric model of the PA distribution, a least square fit is implemented using the PA measurements with given errors at each point of the image."
+ The ouly paralcters are the position of the svuunetric ceuter., The only parameters are the position of the symmetric center.
+ The most probable location of the nucleus js determined as the point of the iuiuinmuuv with: an error circleH defined* by a certain- confidence level.," The most probable location of the nucleus is determined as the point of the minimum, with an error circle defined by a certain confidence level."
+ The method is described in detail in Wishimoto (1999)., The method is described in detail in Kishimoto (1999).
+ We have implemented this fit to the data shown in Figure 2. aud 3..," We have implemented this fit to the data shown in Figure \ref{fig_mrk3_P_S275}
+ and \ref{fig_mrk3_P_S342}."
+ The reduced wwas found to be 66.53/56=Lis aud 18.3/15=1.07 for the F275W and E312W data. respectively.," The reduced was found to be $66.3/56=1.18$ and $48.3/45=1.07$ for the F275W and F342W data, respectively."
+ The mininaun ppoint is shown as a plus sign in Figure 1. together with the error circle of confidence level. for cach of the filters.," The minimum point is shown as a plus sign in Figure \ref{fig_mrk3_nucpos} together with the error circle of confidence level, for each of the filters."
+ The two results are sliehtlv different. though. they are not Incousistent with each otler.," The two results are slightly different, though they are not inconsistent with each other."
+" The central ~2"" of the UV image of Mi 3 cousists of several resolved. clouds.", The central $\sim 2''$ of the UV image of Mrk 3 consists of several resolved clouds.
+ There is uo uresolved. overwhelmuinely bright poiut source iu ——2he image. ax expected from the unified. model.," There is no unresolved, overwhelmingly bright point source in the image, as expected from the unified model."
+ Towever. the error eircle for the unclear location includes (or mareinally includes) a rather bright cloud.," However, the error circle for the nuclear location includes (or marginally includes) a rather bright cloud."
+ The nucleus is thought to be hidden from. direct view in this wavelength region. so we would not expect the radiation of this cloud to be the direct unclear light.," The nucleus is thought to be hidden from direct view in this wavelength region, so we would not expect the radiation of this cloud to be the direct nuclear light."
+ The FOC data rather prefer that the nucleus resides at the western side (SW or NW) of the cloud. but the data do not strongly constrain which side the nucleus is on. aud do not even exclude the possibility that the nucleus is within the cloud.," The FOC data rather prefer that the nucleus resides at the western side (SW or NW) of the cloud, but the data do not strongly constrain which side the nucleus is on, and do not even exclude the possibility that the nucleus is within the cloud."
+" The size of this cloud is about ""06 (— 20pc) in EWIIM in both of the F275W and E312W images. where the diffraction limit of he HIST at this waveleneth is about 0.703."," The size of this cloud is about $0.''06$ $\sim 20$ pc) in FWHM in both of the F275W and F342W images, where the diffraction limit of the HST at this wavelength is about $0.''03$."
+ The total fux color of this cloud is very red : he synthetic photometiy with a circular aperture of 0.715 diameter gives £AQSOOAT)/F40100A)) = 1.55+£0.02. which corresponds toa=5.020.2 where FoxvU.," The total flux color of this cloud is very red : the synthetic photometry with a circular aperture of $0.''15$ diameter gives $F_{\lambda}$ $F_{\lambda}$ ) = $0.55\pm0.02$, which corresponds to $\alpha = -5.0 \pm 0.2$ where $F_{\nu} \propto
+\nu^{\alpha}$."
+ The measured fux is Fy (2800A)) = (3.6£01).101T OCTO 2 41A land FAQOIOUA)) = (65202)«10I erg 2s + (hut note the emission line contamination: see previous section)., The measured flux is $F_{\lambda}$ ) = $(3.6\pm0.1) \times 10^{-17}$ erg $^{-2}$ $^{-1}$ $^{-1}$ and $F_{\lambda}$ ) = $(6.5\pm0.2) \times 10^{-17}$ erg $^{-2}$ $^{-1}$ $^{-1}$ (but note the emission line contamination; see previous section).
+" These are after the correctiou for the Galactic reddening. using the reddeniug curve of Cardelli, Clayton. Mathis (1989) with Egc—0488 (NED: Schlegel et al."," These are after the correction for the Galactic reddening, using the reddening curve of Cardelli, Clayton, Mathis (1989) with $E_{B-V} = 0.188$ (NED; Schlegel et al."
+ 1998)., 1998).
+ The calculated polarization for this cloud is not much larger than the uncertainty of the FOC polarization incasirementwith small apertures., The calculated polarization for this cloud is not much larger than the uncertainty of the FOC polarization measurementwith small apertures.
+ The svuthetic circular aperture polariuctry with 0.715 diameter for this blob. after the iuterstellar polarization correction (see," The synthetic circular aperture polarimetry with $0.''15$ diameter for this blob, after the interstellar polarization correction (see"
+It is. thus. uccessary to populate tlhe DAL halos of an N-body simulation with modeledgalarics.. and colmpare this virtual sample against the ALFALFA measurement.,"It is, thus, necessary to populate the DM halos of an N-body simulation with modeled, and compare this virtual sample against the ALFALFA measurement."
+ Modeling of the atomic lLydrogen content of the svuthetic galaxies is particularly desirable. because it ereatly facilitates the comparison between theoretical and observed distributions.," Modeling of the atomic hydrogen content of the synthetic galaxies is particularly desirable, because it greatly facilitates the comparison between theoretical and observed distributions."
+ have sinmated the ΠΠ]Πιο profiles for he galaxies in the semi-analvtie catalog. created by oost-processime the ΠΠ N-body simulation ).," have simulated the HI-line profiles for the galaxies in the semi-analytic catalog, created by post-processing the Millennium N-body simulation ."
+. Figure 9 displavs the WF (evan solid line) resulting from(7? xojecetius their modeled edge-on lnewidths. assumniug random galaxy inclinations.," Figure \ref{fig:obreschkow} displays the WF (cyan solid line) resulting from projecting their modeled edge-on linewidths, assuming random galaxy inclinations."
+ The O09 WF is iu fairly eood agreement with the ALFALFA measurement. but als to display an exponcutial cutoff at lieh widths aud herefore predicts too many high-width galaxies.," The O09 WF is in fairly good agreement with the ALFALFA measurement, but fails to display an exponential cutoff at high widths and therefore predicts too many high-width galaxies."
+ This issue has been also poiuted out in Zwl0. who areue that he disagrecment is caused bv the fact that the O09 catalog overestimates the UT masses of massive carly-ype galaxies.," This issue has been also pointed out in Zw10, who argue that the disagreement is caused by the fact that the O09 catalog overestimates the HI masses of massive early-type galaxies."
+ They found that restricting themselves to svuthetic galaxies classified as late-types (based ou their nulee-to-total stellar mass ratios in the DeLucia catalog) xoduced a iich better fit to their data., They found that restricting themselves to synthetic galaxies classified as late-types (based on their bulge-to-total stellar mass ratios in the DeLucia catalog) produced a much better fit to their data.
+" However. Figure 9 sneeests that applving the ""morphological cut of Zwl0 results in too few galaxies at intermediate widths (200 iwc600 1)."," However, Figure \ref{fig:obreschkow} suggests that applying the “morphological” cut of Zw10 results in too few galaxies at intermediate widths $200$ $< w < 600$ )."
+ The reds solid line in Figure 9 is the WF correspouding to an indirect observational estimate of the velocity distribution of spiral galaxies by?., The red solid line in Figure \ref{fig:obreschkow} is the WF corresponding to an indirect observational estimate of the velocity distribution of spiral galaxies by.
+. Their determination of the spiral ealaxy VE. was producedby combining the Southern Sky Redshift Survey DB-baud LF for spirals in conjunction with the Tullv-Fisher parameters iu the Br-baud., Their determination of the spiral galaxy VF was producedby combining the Southern Sky Redshift Survey $B$ -band LF for spirals in conjunction with the Tully-Fisher parameters in the $B_T$ -band.
+ This iudirect method. based on galaxy scaling relations. is reliable oulv for relatively massive spirals (ρω2770 13) and suffers from muuerous sources of uucertaintv (6.8. scatter in the TF relation. uucertainties related to the correction of the LF for extinction. bandpass conversion uncertaiuties. ete).," This indirect method, based on galaxy scaling relations, is reliable only for relatively massive spirals $v_{rot} > 70$ ) and suffers from numerous sources of uncertainty (e.g. scatter in the TF relation, uncertainties related to the correction of the LF for extinction, bandpass conversion uncertainties, etc.)."
+ CDAL predictions start diverging from the observational results at low widths. and so the behavior of the theoretical WE for e<200 iis of ereat importance.," CDM predictions start diverging from the observational results at low widths, and so the behavior of the theoretical WF for $w < 200$ is of great importance."
+ Unfortunately. the very interesting work of O09 is ouly reliable for w>100 ddue to the limitations iu the mass resolution of the \Gllennimiu simulation.," Unfortunately, the very interesting work of O09 is only reliable for $w \gtrsim 100$ due to the limitations in the mass resolution of the Millennium simulation."
+ We curplov instead two recent lieh-resolutioun CDM. simulations. that lack however modeling of the III component of their virtual galaxy salples.," We employ instead two recent high-resolution CDM simulations, that lack however modeling of the HI component of their virtual galaxy samples."
+ Figure 10 compares the ALFALFA measurement with the WF of the galaxy population corresponding to the Bolshoi sinulatiou(?).. as modeled byTC10.," Figure \ref{fig:trujillo} compares the ALFALFA measurement with the WF of the galaxy population corresponding to the Bolshoi simulation, as modeled by."
+. Each Bolshoi halo was assigued. realistic stellar and cold gas masses. based ou empirical relatious.," Each Bolshoi halo was assigned realistic stellar and cold gas masses, based on empirical relations."
+ Subsequenth-. two models were considered. oue where the eravitational potential of the barvous ijs simply superimposed on the DAL poteutial (solid. exeen line) aud oue where the DAL halo acliabatically coutracts diu response to the presence of the barvous (dash-dotted ereecn line).," Subsequently, two models were considered, one where the gravitational potential of the baryons is simply superimposed on the DM potential (solid green line) and one where the DM halo adiabatically contracts in response to the presence of the baryons (dash-dotted green line)."
+" Note that TC10 define 0,5; as the value of the sinulated rotation curve at a radius of 10 kpe.", Note that TG10 define $v_{rot}$ as the value of the simulated rotation curve at a radius of 10 kpc.
+ The authors argue that their iiodeliug scheme aud use of Clogpe provide a good approximation of the measured velocity for galaxies with both flat aud rising rotation curves., The authors argue that their modeling scheme and use of $v_{10kpc}$ provide a good approximation of the measured velocity for galaxies with both flat and rising rotation curves.
+ Also plotted in Figure 160 is the WF of simulated ealaxies based ou the constrained N-body simmlation (blue solid line)., Also plotted in Figure \ref{fig:trujillo} is the WF of simulated galaxies based on the constrained N-body simulation (blue solid line).
+ Zaü9 perform a iodoest volume (612.+ Mpe ou a side) but very high-resolution (ousm21i 19) constrained simulation. designed to reproduce the large-scale structure of the local universe.," Za09 perform a modest volume $64 \: h^{-1}$ Mpc on a side) but very high-resolution $v_{lim} = 24$ ) constrained simulation, designed to reproduce the large-scale structure of the local universe."
+ Virtual galaxies are modeled according to the analytical results of7.. assuming a disk-to-virial mass ratio of faseτοMuadM0.03 independent of halo size.," Virtual galaxies are modeled according to the analytical results of, assuming a disk-to-virial mass ratio of $f_{disk} \equiv M_{disk} / M_{vir} =0.03$ independent of halo size."
+ Lastly. the maximum amplitude of the rotation curve (εμας) for cach galaxy is calculated. by combining the disk aud DAL halo coutributions.," Lastly, the maximum amplitude of the rotation curve $v_{rot,max}$ ) for each galaxy is calculated, by combining the disk and DM halo contributions."
+" Since neither model cousiders the distribution of the velocity field tracer (e. HD) in simulated galaxies. we convert rotational velocities iuto III velocity widths by assuniue the relationship Calaxies are assuned to be randomly orieuted with respect to the lme-ofsighlt (cos/ is uniformly distributed in the |0.1] interval). while (5g is a small ""effective"" ern used to reproduce the broadening cffect of urbuleuce aud non-circular motions on II linewidths."," Since neither model considers the distribution of the velocity field tracer (i.e. HI) in simulated galaxies, we convert rotational velocities into HI velocity widths by assuming the relationship Galaxies are assumed to be randomly oriented with respect to the line-of-sight $\cos i$ is uniformly distributed in the $[0,1]$ interval), while $w_{eff}$ is a small “effective” term used to reproduce the broadening effect of turbulence and non-circular motions on HI linewidths."
+ The use of equ., The use of eqn.
+" Lo is justified if the TT disk is extended enough to sample the value of ο adopted by he model uuder consideration (οιο, πρι for TOLO aud Ceoranas tor 240),"," \ref{eqn:vtow} is justified if the HI disk is extended enough to sample the value of $v_{rot}$ adopted by the model under consideration (e.g. $v_{10kpc}$ for TG10 and $v_{rot,max}$ for Za09)."
+" We adopt the value (v,=5 for the broadeningterui?. which is added linearly for ealaxies with Όρο50 and in quadrature for ower velocity ealaxies."," We adopt the value $w_{eff} = 5$ for the broadening, which is added linearly for galaxies with $v_{rot} > 50$ and in quadrature for lower velocity galaxies."
+ Figure 10 puts in evidence the marked departure of the theoretical distributious from the ALFALFA neastrement ata<200i. which becomes wore dramatic with decreasing width.," Figure \ref{fig:trujillo} puts in evidence the marked departure of the theoretical distributions from the ALFALFA measurement at $w < 200$, which becomes more dramatic with decreasing width."
+" According to the TOTO WE. the difference is approximately a factor of ~| at w=100ον, exhibiting au increasing trend."," According to the TG10 WF, the difference is approximately a factor of $\sim4$ at $w = 100$, exhibiting an increasing trend."
+ The Ζα0WE. nuplies a difference of a factor of ~8 at the lowest width where the simulation is complete (ae50 iJ) and displays a much steeper low-width slope than the ALFALFA iecasurement.," The Za09, implies a difference of a factor of $\sim 8$ at the lowest width where the simulation is complete $w \approx 50$ ), and displays a much steeper low-width slope than the ALFALFA measurement."
+" Au extrapolation of the Za09 WE to the ALFALFA width limit (€=20 13), would result in a discrepancy of a factor of ~100."," An extrapolation of the Za09 WF to the ALFALFA width limit $w = 20$ ), would result in a discrepancy of a factor of $\sim 100$."
+ The ALFALFA measurement of the WF confirms the results of the ΠΑΟ survey (?).. which obtained its WE at lower sensitivity aud velocity resolution.," The ALFALFA measurement of the WF confirms the results of the HIPASS survey , which obtained its WF at lower sensitivity and velocity resolution."
+ This fact excludes the possibility that the CDAL overabundance woblem is an artifact of the limited performance characteristics of past blind 21 cur surveys., This fact excludes the possibility that the CDM overabundance problem is an artifact of the limited performance characteristics of past blind 21 cm surveys.
+ The reason, The reason
+"is the situation in between jaand kmax4, as the largest mode is 2pc, which corresponds to a k—2 mode in a 4pc domain.","is the situation in between and 4, as the largest mode is $2\pc$, which corresponds to a $k=2$ mode in a $4\pc$ domain."
+ As the density distribution and therefore the density contrast is self-similar in all three cases we expect that the same regions of the initial conditions will form the dominant structures., As the density distribution and therefore the density contrast is self-similar in all three cases we expect that the same regions of the initial conditions will form the dominant structures.
+" Only the size of the encompassed region and therefore the size and mass of the pillars should change, if the process is indeed scale free."," Only the size of the encompassed region and therefore the size and mass of the pillars should change, if the process is indeed scale free."
+ In Fig., In Fig.
+" 3 all three simulations (panel 1, panel 11, panel 12) show a clear sign of the largest k—1 mode."," \ref{FIG_compare} all three simulations (panel 1, panel 11, panel 12) show a clear sign of the largest $k=1$ mode."
+ The size of the structures formed is linearly dependent on the initial box-length or size of the largest k-mode., The size of the structures formed is linearly dependent on the initial box-length or size of the largest k-mode.
+ The values of Table 2. confirm the importance of the largest mode., The values of Table \ref{TAB_compare} confirm the importance of the largest mode.
+ In the assembled structures the estimated diameters are roughly a factor of two different and the masses vary by a factor of four., In the assembled structures the estimated diameters are roughly a factor of two different and the masses vary by a factor of four.
+ This is as expected since the regions initially encompassed by the radiation should differ by a factor of two in the y- and the z-direction., This is as expected since the regions initially encompassed by the radiation should differ by a factor of two in the y- and the z-direction.
+" Taking these results on the turbulent scale into account, we conclude that the mass and size of the pillars is directly dependent on the input scale of the turbulence, e.g the size of the driving process or the size of the pre-existing molecular cloud."," Taking these results on the turbulent scale into account, we conclude that the mass and size of the pillars is directly dependent on the input scale of the turbulence, e.g the size of the driving process or the size of the pre-existing molecular cloud."
+" On average, the most prominent structures in our simulations with an intermediate flux are dpi&πο where Zturb is the largest turbulent inputmode®.."," On average, the most prominent structures in our simulations with an intermediate flux are $d_\mathrm{pil}\approx\frac{1}{40}x_\mathrm{turb}$, where $x_\mathrm{turb}$ is the largest turbulent input."
+ In several simulations triggered dense regions form cores and are driven into gravitational collapse., In several simulations triggered dense regions form cores and are driven into gravitational collapse.
+ Since star formation is not the main goal of this study we do not replace them by sink particles., Since star formation is not the main goal of this study we do not replace them by sink particles.
+ Instead we remove the particles forming a core from the simulation to avoid a considerable slowdown of the calculation., Instead we remove the particles forming a core from the simulation to avoid a considerable slowdown of the calculation.
+" Following G09a we define a core as all gas with a density above Perit=10""mycm in the region around the density peak and remove the? particles representing this core.", Following G09a we define a core as all gas with a density above $\rho_\mathrm{crit}=10^7\dens$ in the region around the density peak and remove the particles representing this core.
+" 'The core formation is not a numerical effect, since the resolution limit as given by ? is Pnum=3x105m,cm3 in the lowest resolution case."," The core formation is not a numerical effect, since the resolution limit as given by \cite{1997MNRAS.288.1060B} is $\rho_\mathrm{num}=3\times
+10^8\dens$ in the lowest resolution case."
+" We give the simulation,mass',, formation time, average speed away from the source and positions in Table 3.."," We give the simulation, formation time, average speed away from the source and positions in Table \ref{star_form}."
+ If we assume the cores to be decoupled from the rest of the cold gas then their position at the end of the simulation can be estimated by these position and velocities., If we assume the cores to be decoupled from the rest of the cold gas then their position at the end of the simulation can be estimated by these position and velocities.
+" All of them are still close to the prominent structures, some are traveling further inside the structures, some are lagging behind and would by now be slightly outside of the pillar, closer to the source."," All of them are still close to the prominent structures, some are traveling further inside the structures, some are lagging behind and would by now be slightly outside of the pillar, closer to the source."
+" As it can be expected, inbox,, where the compressed structures are most massive, star formation is most frequent and happens earliest."," As it can be expected, in, where the compressed structures are most massive, star formation is most frequent and happens earliest."
+ There is an age spread present as well - the earlier a core forms the closer to the source it will be., There is an age spread present as well - the earlier a core forms the closer to the source it will be.
+ The other simulations where cores form during the first 500kyr after the ignition of the O star are and(G09b)., The other simulations where cores form during the first $500\kyr$ after the ignition of the O star are and.
+. The higher flux leads to a higher compression and thus an earlier formation of a core influx., The higher flux leads to a higher compression and thus an earlier formation of a core in.
+. due to the higher photo-evaporation rate this core is also less massive and moves faster compared to the core in(G09b).., due to the higher photo-evaporation rate this core is also less massive and moves faster compared to the core in.
+" Altogether, triggered star formation is very likely in this scenario."," Altogether, triggered star formation is very likely in this scenario."
+" The cores form at the center of the structures, but since their velocities differ from the velocity of their parental structure they can be decoupled"," The cores form at the center of the structures, but since their velocities differ from the velocity of their parental structure they can be decoupled"
+The complexity of the spectral range around the 14.32 iis shown in Fig. 3..,The complexity of the spectral range around the 14.32 is shown in Fig. \ref{fig:average-FOV-NeV}.
+ An unknown feature at 13.95 sseems to be present just next to the first bending mode of aat 14.02µπι., An unknown feature at $\sim$ 13.95 seems to be present just next to the first bending mode of at $14.02$.
+" Depending on the temperature of the surrounding molecular gas, ccan be vibrationally excited by absorbing the infrared photons at 14.0µπι."," Depending on the temperature of the surrounding molecular gas, can be vibrationally excited by absorbing the infrared photons at $14.0$."
+. This will produce a subsequent cascade process that can enhance rrotational lines in the (sub-)millimeter range., This will produce a subsequent cascade process that can enhance rotational lines in the (sub-)millimeter range.
+" This corresponds to the IR-pumping scenario proposed to explain the bright J=1—0 (and higher) transition observed in some ULIRGs and Galactic star-forming regions (e.g.,?????).."," This corresponds to the IR-pumping scenario proposed to explain the bright $J=1\rightarrow 0$ (and higher) transition observed in some ULIRGs and Galactic star-forming regions \citep[e.g.,][]{aalto95, garcia06, guelin07, aalto07a, pb10}."
+" Although, due to the order mismatch (mentioned in Sec. ??))"," Although, due to the order mismatch (mentioned in Sec. \ref{sec:results}) )"
+" observed between ~13.90 aand ~14µπι,, also seen by ?,, we are unable to conclude whether we detect or not the 14.02 ffeature, either in emission or in absorption."," observed between $\sim13.90$ and $\sim14$, also seen by \citet{pereira10}, we are unable to conclude whether we detect or not the $14.02$ feature, either in emission or in absorption."
+" A larger PAH feature is observed around ~14.22µπι,, and on its rightmost (and less steep) slope the fine-structure emission lines of 14.32 aand 14.37 aare found."," A larger PAH feature is observed around $\sim$ 14.22, and on its rightmost (and less steep) slope the fine-structure emission lines of 14.32 and 14.37 are found."
+" Chlorine, which has an ionization potential of 12.97eV, is known to play an important role in characterizing the neutral gas components in the ISM."," Chlorine, which has an ionization potential of $12.97~\rm eV$, is known to play an important role in characterizing the neutral gas components in the ISM."
+" When Hp is abundant (optically thick in the FUV), it reacts (exothermically by 0.17 eV) with tto form HCI*, which leads to the formation of aand citepjura74, jura78.."," When $_2$ is abundant (optically thick in the FUV), it reacts (exothermically by $0.17~\rm eV$ ) with to form $^+$, which leads to the formation of and \\citep{jura74, jura78}."
+" This means, chlorine is predominantly ionized in rregions while it is predominantly neutral when cold H5 components are present (??).."," This means, chlorine is predominantly ionized in regions while it is predominantly neutral when cold $_2$ components are present \citep{sonnentrucker02, sonnentrucker03}."
+" The 14.37 ffine-structure emission line was clearly detected (730 in the whole region mapped with the IRS/SH module, and presents a similar (although more spreadout) distribution than that of the 12.81 Iline."," The 14.37 fine-structure emission line was clearly detected $>$ $\sigma$ ) in the whole region mapped with the IRS/SH module, and presents a similar (although more spreadout) distribution than that of the 12.81 line."
+" Because its ionization potential (97.1eV) is too large to allow production by main-sequence stars, the fine-structure line 14.32 iis commonly used to probe the narrow line region of AGNs"," Because its ionization potential $97.1~\rm eV$ ) is too large to allow production by main-sequence stars, the fine-structure line 14.32 is commonly used to probe the narrow line region of AGNs \citep[e.g.,][and references therein]{moorwood96b, genzel98, sturm02, armus07, alonsoh09, baum10, willett10}."
+ used as a diagnostic tool to unambiguously identify AGN in galaxies that have not been identified as such using optical spectroscopy (?).., It has recently been used as a diagnostic tool to unambiguously identify AGN in galaxies that have not been identified as such using optical spectroscopy \citep{goulding09}.
+" The emission fluxes were estimated by fitting a Gaussian profile to the fine-structure lines and integrating the flux above a local continuum, which was estimated from a cubic spline interpolation of selected anchor points (pivots)."," The emission fluxes were estimated by fitting a Gaussian profile to the fine-structure lines and integrating the flux above a local continuum, which was estimated from a cubic spline interpolation of selected anchor points (pivots)."
+" Since we we want to determine the actual spatial distribution of the emission lines in the nuclear region of NGC 4945, we set a strict 3c level detection for all the fine-structure emission lines, and at least two data points (or a full width of 4560 stat 14.32 um)) in the profile to consider a feature as a real emission line in the spectrum."," Since we we want to determine the actual spatial distribution of the emission lines in the nuclear region of NGC 4945, we set a strict $\sigma$ level detection for all the fine-structure emission lines, and at least two data points (or a full width of $\sim$ 560 at 14.32 ) in the profile to consider a feature as a real emission line in the spectrum."
+ For a robust estimate of the local continuum we adopted an average flux of five points around each of the selected pivots used in the cubic spline interpolation., For a robust estimate of the local continuum we adopted an average flux of five points around each of the selected pivots used in the cubic spline interpolation.
+ Figure 4 shows the Gaussian fit of the 14.32 aand 14.37 ffine-structure emission lines at three different positions of the IRS/SH map., Figure \ref{fig:Gaussian-fit-NeV} shows the Gaussian fit of the 14.32 and 14.37 fine-structure emission lines at three different positions of the IRS/SH map.
+ The uncertainty of each spectral data point is indicated by the error bars., The uncertainty of each spectral data point is indicated by the error bars.
+ The pivots used for the cubic spline interpolation of the local continuum (baseline) are shown with filled circles., The pivots used for the cubic spline interpolation of the local continuum (baseline) are shown with filled circles.
+" These pivots are at the wavelengths 14.201um,, 14.247 um,, 14.281 um,, 14.428um,, 14.531um,, 14.667µπι,, and 14.780um."," These pivots are at the wavelengths 14.201, 14.247 14.281 , 14.428, 14.531, 14.667, and 14.780."
+. The of Fig., The of Fig.
+" 4 shows the spectrum in a region of the map with relatively high uncertainties (~1.4% on average) in the spectral data and with high RMS (as computed from the baseline, and indicated in the legend of the plots) that leads to less than a 3c detection forV],, but a clear detection (2307) for the Iline."," \ref{fig:Gaussian-fit-NeV} shows the spectrum in a region of the map with relatively high uncertainties $\sim$ on average) in the spectral data and with high RMS (as computed from the baseline, and indicated in the legend of the plots) that leads to less than a $\sigma$ detection for, but a clear detection $>$ $\sigma$ ) for the line."
+" The spectrum of the was extracted from the central 2x2 pixel aperture of the map, which has a S/N ratio 9 in the line and >100 in the continuum, and is among the highest S/N in the whole IRS/SH map."," The spectrum of the was extracted from the central $\times$ 2 pixel aperture of the map, which has a S/N ratio $\sim$ 9 in the line and $>$ 100 in the continuum, and is among the highest S/N in the whole IRS/SH map."
+" This spectrum shows very low uncertainties (~0.06% on average) in the data, and clear detections of both lines."," This spectrum shows very low uncertainties $\sim$ on average) in the data, and clear detections of both lines."
+" Although with low uncertainties in the data (~0.13%)),the of Fig."," Although with low uncertainties in the data $\sim$ ),the of Fig."
+" 4 shows less than a 3c detection for ((with a S/N 1.8), but a strong detection of 14.37 um.."," \ref{fig:Gaussian-fit-NeV} shows less than a $\sigma$ detection for (with a $\sim$ 1.8), but a strong detection of 14.37 ."
+while [rom the definition of specific available euthalpy (Equation 12)) aud the thermodynamic equation (Equation 1)) we cau write: Here we have expanded out c&=DP/DIOP/Ol4+0-VP+w(OP/O0z) aud used the hydrostatie assumption: e(OP/0z)=—wpg. where w=Dz/Dl is the vertical velocity.,"while from the definition of specific available enthalpy (Equation \ref{eq:adef}) ) and the thermodynamic equation (Equation \ref{eq:dsdt}) ) we can write: Here we have expanded out $\omega \equiv DP/Dt = \partial P/\partial t + \vec{v} \cdot \nabla P + w (\partial P/\partial z)$ and used the hydrostatic assumption: $w (\partial P/\partial z) = - w \rho g$, where $w \equiv Dz/Dt$ is the vertical velocity."
+ For the kinetic energy we take the dot product of the horizontal velocity aud the staudard horizontal momentum equation. Di'/DI!=—(1/p)VP—fx&+ Fg. where f—(20sinó)k is the Coriolis parameter (Q is the planetary rotation rate and © is latitude) in the vertical direction (hr). lo get: where we have defined D as the irreversible frictional dissipatiou of kinetic euergy (=—U-Frac):," For the kinetic energy we take the dot product of the horizontal velocity and the standard horizontal momentum equation, $D\vec{v}/Dt = -(1/\rho) \vec{\nabla} P - \vec{f} \times \vec{v} + \vec{F}_{\mathrm{fric}}$ , where $\vec{f}=(2 \Omega \sin \phi) \vec{k}$ is the Coriolis parameter $\Omega$ is the planetary rotation rate and $\phi$ is latitude) in the vertical direction $\vec{k}$ ), to get: where we have defined $\mathcal{D}$ as the irreversible frictional dissipation of kinetic energy $= - \vec{v} \cdot \vec{F}_{\mathrm{fric}}$ )."
+ The derivative of the gravitational potential energy is. by celinition: We explicitly split the beatingo into its frictional f (q7) and non-fictional (mainly radiative. radGrad) components and combine. EquationsE⋅ 15-18 2into a local energy," The derivative of the gravitational potential energy is, by definition: We explicitly split the heating into its frictional $q_f$ ) and non-frictional (mainly radiative, $q_{\mathrm{rad}}$ )components and combine Equations \ref{eq:dudt}- \ref{eq:eg} into a local energy"
+Type Ta Superuovae (SNe Ia) play an muüportaut role m astrophysics as cosmiological distance indicators.,Type Ia Supernovae (SNe Ia) play an important role in astrophysics as cosmological distance indicators.
+ Additionally. they provide iron peak elements having direct consequences for the chemical evolution of galaxies (Riessetal.1995:Matteucci&Cregeio1986:deDon-der&Vanbeveren 2003)..," Additionally, they provide iron peak elements, having direct consequences for the chemical evolution of galaxies \citep{RPK95,MG86,DV03}. ."
+ Currently z2000 SNe Ia have been obscrvatiouallyconfirmed*.. some as distant as lo1.55 (Strolgeretal.2001).," Currently $\gtrsim 2000$ SNe Ia have been observationally, some as distant as $z \sim 1.55$ \citep{Str04}."
+.. Exipiricallv-derived relatiouships between light curve properties aud iutriusic huninosity (1.6. Αρης (Phillips1993). auc streteli-factor s. (Perhuutteretal. 1997))) have made it possible to ‘standardize’ absolute magnitudes of SNe Ia light curves over a wide variety of Lost galaxy Their use as standard candles” on cosmological scales has led to the realization that the expansion rate of the universe is accelerating. and has enabled accurate estimations of O4 and Qa; (e.c.Schmidtetal.1998:Ricssetal.1998:Perluutter 1999).," Empirically-derived relationships between light curve properties and intrinsic luminosity (i.e., $\Delta m_{15}$ \citep{Phi93} and stretch-factor $s$, \citep{Per97}) ) have made it possible to `standardize' absolute magnitudes of SNe Ia light curves over a wide variety of host galaxy Their use as `standard candles' on cosmological scales has led to the realization that the expansion rate of the universe is accelerating, and has enabled accurate estimations of $\Omega_{\Lambda}$ and $\Omega_{M}$ \citep[e.g.,][]{Sch98,Rie98,Per99}."
+. However. using SNe Ta in order to set the distance scale for the deteriuuation of cosmological quantities requires the (still unfouuded) assumption that the physical properties of their progenitors are unchaugins with redslüft.," However, using SNe Ia in order to set the distance scale for the determination of cosmological quantities requires the (still unfounded) assumption that the physical properties of their progenitors are unchanging with redshift."
+ Despite the contiuued use of SNe Ta as standard candle distance imdicators. them origin remains uncertain.," Despite the continued use of SNe Ia as standard candle distance indicators, their origin remains uncertain."
+ It is generally accepted that SNe Ia arise from the total disruption of a Chaudrasekliar mass (~1.4 M.) carbou-oxveen white dwart (WD) as a result of thermonuclear explosion., It is generally accepted that SNe Ia arise from the total disruption of a Chandrasekhar mass $\sim 1.4$ $_{\odot}$ ) carbon-oxygen white dwarf (WD) as a result of thermonuclear explosion.
+ This hypothesis is supported by the fact that the amount of enerev observed in the explosions (~101 ere: Thiclemannetal. (2001))) is equal to the amount which would be produced iu the conversion of carbon and oxvecn into ion (seec.g..Livio2000.forareview)..," This hypothesis is supported by the fact that the amount of energy observed in the explosions $\sim 10^{51}$ erg; \citet{Thi04}) ) is equal to the amount which would be produced in the conversion of carbon and oxygen into iron \citep[see e.g.,][for a review]{Liv00}."
+ It is natural to prestune that the exploding white dwarf must accrete matter from a close stellar courpanion unti reaching the critical Chandrasekhar mass. though the nature of the companion. the rate and cficiency at which mass is accumulated onto the white dwarf. au which array of couditious are necessary in order for the white dw to ignite explosively. are not well uuderstooc (Nomotoctal.1997).," It is natural to presume that the exploding white dwarf must accrete matter from a close stellar companion until reaching the critical Chandrasekhar mass, though the nature of the companion, the rate and efficiency at which mass is accumulated onto the white dwarf, and which array of conditions are necessary in order for the white dwarf to ignite explosively, are not well understood \citep{NIK97}."
+. Iu order to constrain the nature of SN Ia progenitors. the SN Ia rate has been studiedas a functiou of parent ealaxy stellar mass. star formation rate. colour. morphology. and radio power by several groups 2003)..," In order to constrain the nature of SN Ia progenitors, the SN Ia rate has been studiedas a function of parent galaxy stellar mass, star formation rate, colour, morphology, and radio power by several groups \citep[e.g.,][]{Man05,Sul06,CM06,DP03}. ."
+Consequently. the only contribution to στον we have to be concerned with is σῃο.,"Consequently, the only contribution to $\sigma^{2}_{\rm TOT}$ we have to be concerned with is $\sigma^{2}_{\rm BG}$."
+ Once this is removed. σὲ will be obtained from 1%).," Once this is removed, $\sigma^{2}_{\rm GC}$ will be obtained from $P_0$."
+ The following law can be assumed for the luminosity finetion of background. galaxies (BGLF): where .lis a constant lor normalization. Nec: is the total number of galaxies per pixel and 5 is the slope of the magnitude distribution.," The following law can be assumed for the luminosity function of background galaxies (BGLF): where $A$ is a constant for normalization, $N_{\rm BG}$ is the total number of galaxies per pixel and $\gamma$ is the slope of the magnitude distribution."
+" Tvson(1988) found 5zz0.39 in the H filter,", \citet{T88} found $\gamma \approx 0.39$ in the $R$ filter.
+ Using eq. 4..," Using eq. \ref{sigma},"
+ we find Once (he parameters A. Nec and y are estimated [rom an analvsis of the image. as we will see later. and inj from the photometric calibration. ej; can be caleulated.," we find Once the parameters $A$, $N_{\rm BG}$ and $m_{\rm c}$ are estimated from an analysis of the image, as we will see later, and $m_1^*$ from the photometric calibration, $\sigma_{\rm BG}^2$ can be calculated."
+ As we have seen. because stars. background galaxies. and GCs are convolved with the PSF. while the read-out ancl photon shot noise are not. if a power spectrum of (he image is computed. the non-convolved variance can be separated [rom the remaining terms.," As we have seen, because stars, background galaxies, and GCs are convolved with the PSF, while the read-out and photon shot noise are not, if a power spectrum of the image is computed, the non-convolved variance can be separated from the remaining terms."
+" The power spectrum of the image. (). has the form where E(hk) is the power spectrum of the PSF convolved with (he window function (which is unitv in all the regions. we want (o study and zero in the rest of the image).ü 77,> is the PSF-convolved variance (2) = o2ptoRetopetopet..). and P?1 is the non-PSF-convolved variance--- (P4 —= oi,g2 dou).Lg)"," The power spectrum of the image, $P(k)$, has the form where $E(k)$ is the power spectrum of the PSF convolved with the window function (which is unity in all the regions we want to study and zero in the rest of the image), $P_0$ is the PSF-convolved variance $P_0$ = $\sigma^{2}_{\rm SP}$ $\sigma^{2}_{\rm GC}$ $\sigma^{2}_{\rm 
+BG}$ $\sigma^{2}_{\rm FS}$ +...), and $P_1$ is the non-PSF-convolved variance $P_1$ = $\sigma^{2}_{\rm ro}$ $\sigma^{2}_{\rm ph}$ )."
+ Belore performing the power spectrum of the image. a preliminary image treatment has io be done.," Before performing the power spectrum of the image, a preliminary image treatment has to be done."
+is similar with the toroidal Geld being the dominant field of LID 171488 (Marsden.etal.2006:Jellers&Donati 2008).. AB Dor and LQ Ίνα (Donatietal.2003a).,"is similar with the toroidal field being the dominant field of HD 171488 \citep{Marsden06,Jeffers08}, AB Dor and LQ Hya \citep{Donati03a}."
+. Like LID 171488. this field is quite complex and much more than that operating the Sun.," Like HD 171488, this field is quite complex and much more than that operating the Sun."
+ Consicering the toroidal component. an octopole field dominates with approximately 37 c of the magnetic energy. whereas the dipole field contains only 18 c and a quadrople field. containing approximately 164 of the magneticIn] energy.," Considering the toroidal component, an octopole field dominates with approximately 37 $\pm$ of the magnetic energy whereas the dipole field contains only 18 $\pm$ and a quadrople field containing approximately 16 $\pm$ of the magnetic energy."
+D This is in contrast to the poloidal field were the respective components are roughly equal between a dipole (10 c 2). quadopole (11 + )) and octopole field (10 + 150).," This is in contrast to the poloidal field were the respective components are roughly equal between a dipole (10 $\pm$ ), quadopole (11 $\pm$ ) and octopole field (10 $\pm$ )."
+ This would. indicate that a much more complex dynamo is operating on LID 106506 and is similar to that on other voung. sun-like stars.," This would indicate that a much more complex dynamo is operating on HD 106506 and is similar to that on other young, sun-like stars."
+ In addition. the toroidal field was predominantly. axisymetric (05 cx 2%)) whereas the poloidal field. tended to be more non-axisvmetric (65 c QE," In addition, the toroidal field was predominantly axisymetric (95 $\pm$ ) whereas the poloidal field tended to be more non-axisymetric (65 $\pm$ )."
+ This is consistent with the observations of Donati&Landstreet(2009) and Marsdenetal.(2011a) where stars with significant toroidal fields often have non-axisvmetic poloidal fieles., This is consistent with the observations of \citet{Donati09} and \citet{Marsden10a} where stars with significant toroidal fields often have non-axisymetic poloidal fields.
+ Phe error. estimates on these measureswere obtained by varving a range of μαellar parameters such as inclination. radial velocity. Quy. 0. period. andesinz.," The error estimates on these measureswere obtained by varying a range of stellar parameters such as inclination, radial velocity, $\Omega$, $\delta\Omega$, period and."
+ Hence these may be refered to as a aviation bar. which is analogous to an error bar except al the measurement indicates the variation found by this methoc.," Hence these may be refered to as a $\it variation$ bar, which is analogous to an error bar except that the measurement indicates the variation found by this method."
+" When measuring the cillerential rotation using the magnetic features. the equatorial rotation rate £2, was 4.51 + 0.01Lo with a photospheric shear 6Q of 0.24 + 0.03 4."," When measuring the differential rotation using the magnetic features, the equatorial rotation rate $\Omega$ was 4.51 $\pm$ 0.01, with a photospheric shear $\delta\Omega$ of 0.24 $\pm$ 0.03 ."
+ Phed comparison with those parameters found using the Stokes / information is shown in Figure 12.., The comparison with those parameters found using the Stokes $\it{I}$ information is shown in Figure \ref{error_ellipses}.
+ However. this cillerence in the dQ values are within the error bars of the (wo measurements.," However, this difference in the $\delta\Omega$ values are within the error bars of the two measurements."
+ When compared with the cdillerences observed. on other stars of similar spectral tvpe such as LID 141948 (Marsdenetal.2011b).. these features may be anchored at similar depths in the convective zone. unlike HD 141943 which may be anchored at different depths of the convective zone.," When compared with the differences observed on other stars of similar spectral type such as HD 141943 \citep{Marsden10b}, these features may be anchored at similar depths in the convective zone, unlike HD 141943 which may be anchored at different depths of the convective zone."
+ However. this conclusion must be treated =vith caution until further Stokes V. data is obtained for this sta.," However, this conclusion must be treated with caution until further Stokes $\it{V}$ data is obtained for this star."
+ Figure 9 shows a bright spot in this ονπαο profile at 70.15 phase., Figure \ref{chromo1} shows a bright spot in this dynamic profile at $\sim$ 0.15 phase.
+ This is where the chromosphere is much denser as a result of an increase in temperature in that local region., This is where the chromosphere is much denser as a result of an increase in temperature in that local region.
+ There appears to be a reduced emission at 20.75 phase as well., There appears to be a reduced emission at $\sim$ 0.75 phase as well.
+ Phese variations in Ho. emission lie within the stellar velocity profile ancl hence arise [rom regions close to the stellar surface., These variations in $\alpha$ emission lie within the stellar velocity profile and hence arise from regions close to the stellar surface.
+ We thus interpret. these Lla variations in terms of prominences more akin to solar prominences than the more extended features knows to exist at the co-rotation radius of other active solar-type stars. such as LO Lup (Donatietal.2000).," We thus interpret these $\alpha$ variations in terms of prominences more akin to solar prominences than the more extended features knows to exist at the co-rotation radius of other active solar-type stars, such as LQ Lup \citep{Donati00}."
+. In this paper. reconstructed brightness and magnetic images of the voung solar-tvpe star HD. 106506 have been presented.," In this paper, reconstructed brightness and magnetic images of the young solar-type star HD 106506 have been presented."
+ The brightness image shows low-latituce to mid-latitude features coupled. with a predominate polar spot., The brightness image shows low-latitude to mid-latitude features coupled with a predominate polar spot.
+ Ες is very similiar to many of rapidly rotating solar-tvpe stars., This is very similiar to many of rapidly rotating solar-type stars.
+ The magnetic images reveal regions of azimuthal field near the stellar surface that suggest that the dynamo mechanism may be occuring entirely in the convective zone and close to the surface of the star itself., The magnetic images reveal regions of azimuthal field near the stellar surface that suggest that the dynamo mechanism may be occuring entirely in the convective zone and close to the surface of the star itself.
+ The photospheric shear for HD 106506 supports the findings of Barnesetal. (2005).. that effective temperature is the dominant factor inthe level of surface dillerential rotation.," The photospheric shear for HD 106506 supports the findings of \citet{Barnes05}, , that effective temperature is the dominant factor inthe level of surface differential rotation."
+ However. the reason why such a relation exists remains unclear.," However, the reason why such a relation exists remains unclear."
+However. we must be uore cautious as the possibility exists 1lal (4je. column density. is increasiug as the material behind the shock recombines.,"However, we must be more cautious as the possibility exists that the column density is increasing as the material behind the shock recombines."
+ A conclusion that the column deusity is increasing is at tliis point prenattire. resting as it would on two data poius (11ο axcdNALM spectra).," A conclusion that the column density is increasing is at this point premature, resting as it would on two data points (the and spectra)."
+ The spectra 1usulliciently constraiued the colum1 deusity to xovide measWements of tl rend over a louger tine base., The spectra insufficiently constrained the column density to provide measurements of the trend over a longer time base.
+ A future observation with ot wokd. with simil: “statistics per spectral Channel. establish the nature of aiy treud tliat nay exist.," A future observation with or would, with similar statistics per spectral channel, establish the nature of any trend that may exist."
+" e power law indices ( .s) of the density profile of the ejeca CPejecta X "") auc cireutustellar 1natter distribution (pi x ) are important Jarameters of ]uteractlon ijodels."," The power law indices (n, s) of the density profile of the ejecta ${\rho}_{\rm ejecta}$ $\propto$ $^{\rm -n}$ ) and circumstellar matter distribution ${\rho}_{\rm cs}$ $\propto$ $^{\rm -s}$ ) are important parameters of interaction models."
+ They may be calculated (rom the light cu‘ve and [rom the ratio ol the tempe‘atures of the forward ancl reverse shocks (e.g.. Fraussonοἱal. 1996)).," They may be calculated from the light curve and from the ratio of the temperatures of the forward and reverse shocks (e.g., \citealt{Fran96}) )."
+ We interpret the tempe‘atures of the two colhj»ouents as the temperatures of the forward aud reverse shock emissioi., We interpret the temperatures of the two components as the temperatures of the forward and reverse shock emission.
+" Fraussouetal.(1996). slow that the temperature ratio depends upou the cireuumistellar aud su;»ernova ejecta density. distriutious as T,/Ty=(3—s)7/(ω—32. where s = circuumstellar density «istribution exponent. usuay 2 for a constant aud homogeneous stellar win. and au SN ejecta deisity distribution exponent."," \cite{Fran96}
+ show that the temperature ratio depends upon the circumstellar and supernova ejecta density distributions as ${T_{\rm r}}/{T_{\rm f}} =
+(3 - s)^2 / ( n - 3 )^2,$ where s = circumstellar density distribution exponent, usually 2 for a constant and homogeneous stellar wind, and n = SN ejecta density distribution exponent."
+ Te ratio of the temperatures of tje (wo model fits to theChandra spectun is 0.61/3.2 = 0.19: solviue0 for u vields an index of 5.2540.25., The ratio of the temperatures of the two model fits to the spectrum is 0.61/3.2 = 0.19; solving for n yields an index of $\pm$ 0.25.
+ Repeating the same arithinetic or the spectru1 yields an index of 5.0740.10., Repeating the same arithmetic for the spectrum yields an index of $\pm$ 0.10.
+ An index of ~5 is lower than usually aopted for the ejecta clistritiou: for example. derive results coveriug the iudex range of 7-12.," An index of $\sim$ 5 is lower than usually adopted for the ejecta distribution; for example, \cite{CF94} derive results covering the index range of 7-12."
+ We argue. however. that n ~5 is reasonable.," We argue, however, that n $\sim$ 5 is reasonable."
+ Recently. Zimmermann&Aschenbach(2003) described ai observation of SN1L993J in M81 aud obtained a best-fit spectrum using the cual variable Meka moclel.," Recently, \cite{Zim03} described an observation of SN1993J in M81 and obtained a best-fit spectrum using the dual variable Mekal model."
+ They inferred n = 8.9 [rom the fitted temperatures after fitting the liel| curve to obtain s ~ 1.65., They inferred n = 8.9 from the fitted temperatures after fitting the light curve to obtain s $\sim$ 1.65.
+ However. Franssot&Bjórusson(1995) showed that for 9N1993J the cireumistellar matter density was lols ~ 1.65 as they first estimated from the radio light curve (Fraussouetal.1996).. bul was s = 2 because their initial analysis failed to include syuchrotron sell-absorptiou.," However, \cite{Fran98} showed that for SN1993J the circumstellar matter density was not s $\sim$ 1.65 as they first estimated from the radio light curve \citep{Fran96}, but was s = 2 because their initial analysis failed to include synchrotron self-absorption."
+ Their correction üunplies tliat circumstellar matter distributions really are deposited with s = 2., Their correction implies that circumstellar matter distributions really are deposited with s = 2.
+ We therefore recalculate tje index. value for 8N1993J. From their best-fit inodel. Zimmermann&Aschench{20023) obtained temperatures of [20.0[. aud. [EE keV for the spectnuu.," We therefore recalculate the index value for SN1993J. From their best-fit model, \cite{Zim03} obtained temperatures of $\pm$ 0.04 and $\pm$ 4 keV for the spectrum."
+ Ádopting s = 2. we obtain a value of u of οτι 1.5Ls whe' the large error rauge. driveu by he large eOF Ol the hieli-temiperature component. just overlaps the s ~1.65 value.," Adopting s = 2, we obtain a value of n of $\sim$ $^{+1.5}_{-1.8}$, where the large error range, driven by the large error on the high-temperature component, just overlaps the s $\sim$ 1.65 value."
+ The ejecta «istributiou iudex of $9N1993J. is slightly largere than the value or SNI97S8lx. but the error range reaches neary to the ΡΑΤΣΑ value of n =».," The ejecta distribution index of SN1993J is slightly larger than the value for SN1978K, but the error range reaches nearly to the SN1978K value of n = 5."
+ Having established that au index of ~5 is not substantially cdiferent [rom the the iudex of SN1993J. we briefly examine whether the ciffereuce is in [act seusible.," Having established that an index of $\sim$ 5 is not substantially different from the the index of SN1993J, we briefly examine whether the difference is in fact sensible."
+ From 1).. lower values for the iudex lead to smaller density contrasts beween theforward and reverse shocks. aud acliabatie shocks at early times.," From \cite{CF94}, lower values for the index lead to smaller density contrasts between theforward and reverse shocks, and adiabatic shocks at early times."
+ SN1978lx. was essentially coustanut for at least 3 years, SN1978K was essentially constant for at least 3 years
+possible for the Coma cluster alone. with objects down to an absolute Ix-band magnitude of -19.23.,"possible for the Coma cluster alone, with objects down to an absolute K-band magnitude of -19.23."
+ The colour magnitude relation of the Coma cluster showed a trend for. blucr Ix-16]. colours towards fainter magnitudes., The colour magnitude relation of the Coma cluster showed a trend for bluer K-[16] colours towards fainter magnitudes.
+ The addition of the fainter Virgo galaxies shows that this trend continues to fainter magnitudes., The addition of the fainter Virgo galaxies shows that this trend continues to fainter magnitudes.
+ There appears to be a slight olfset between the Coma and Virgo data., There appears to be a slight offset between the Coma and Virgo data.
+ Although this could. be caused by an error in the relative distance of the two clusters (taken to be 97 and 17 Alpe respectively) it may also be due to the fact that the colours were determined. within larger apertures for the Coma cluster., Although this could be caused by an error in the relative distance of the two clusters (taken to be 97 and 17 Mpc respectively) it may also be due to the fact that the colours were determined within larger apertures for the Coma cluster.
+ The sense of the ollset. (Virgo objects appearing with slightly higher than expected. Ix- is consistent with the colour gradients described in section 16])., The sense of the offset (Virgo objects appearing with slightly higher than expected K-[16]) is consistent with the colour gradients described in section \ref{sec:gradients}.
+ We also note that the 4 bright Virgo galaxies lic on the low boundary of values of Ix- found in the Coma cluster., We also note that the 4 bright Virgo galaxies lie on the low boundary of values of K-[16] found in the Coma cluster.
+ There is no obvious data analysis16] elect that could have caused such a bias., There is no obvious data analysis effect that could have caused such a bias.
+ oth age and metallicity can cllect the strength. of the integrated 16jm emission from the dusty envelopes of the AGB population., Both age and metallicity can effect the strength of the integrated $16\;\rm \mu m$ emission from the dusty envelopes of the AGB population.
+ Younger stellar populations. for example. should have a larger contribution from cdusty AGB stellar envelopes.," Younger stellar populations, for example, should have a larger contribution from dusty AGB stellar envelopes."
+ However. in order to separate the elfects of age and metallicity we will compare our results with those of studies at optical wavelengths.," However, in order to separate the effects of age and metallicity we will compare our results with those of studies at optical wavelengths."
+ As discussed in section 1.. the different responses of observables in the optical and infrared. to changes in metallicity and age. can be used to resolve the degeneracy between these two parameters.," As discussed in section \ref{sec:intro}, the different responses of observables in the optical and infrared, to changes in metallicity and age, can be used to resolve the degeneracy between these two parameters."
+ As discussed in section 3.1. it seems likely that the racial gradients in the Ix-16]. colour are. due to a population eracdient of AGB stars. relative to the general. stellar population., As discussed in section \ref{sec:point} it seems likely that the radial gradients in the K-[16] colour are due to a population gradient of AGB stars relative to the general stellar population.
+ The recdcder Ix-16]. colours toward. the centre of the galaxies. could. be due to vounger ages. higher metallicities or both.," The redder K-[16] colours toward the centre of the galaxies could be due to younger ages, higher metallicities or both."
+ Clemens ct al. (, Clemens et al. (
+2009b) have used optical line strength indices of a sample of 14.000 EPCs from the Sloan Digital Sky Survey (SDSS) to determine age anc metallicity as a function of mass. ‘vironment and. also galactic raclius.,"2009b) have used optical line strength indices of a sample of 14,000 ETGs from the Sloan Digital Sky Survey (SDSS) to determine age and metallicity as a function of mass, environment and also galactic radius."
+ ον find negative racial metallicity gracdicnts for all masses., They find negative radial metallicity gradients for all masses.
+ Although they also find positive radial age gradients for massive galaxies this trend. is less apparent at intermediate masses and may even invert for low mass galaxies., Although they also find positive radial age gradients for massive galaxies this trend is less apparent at intermediate masses and may even invert for low mass galaxies.
+ Phe 4 bright ealaxies in our Virgo cluster sample have central velocity. dispersions in the range 170-250kms+.," The 4 bright galaxies in our Virgo cluster sample have central velocity dispersions in the range $250\;\rm km\,s^{-1}$."
+ In this case little or no age gradient is expected. at least on a statistical basis.," In this case little or no age gradient is expected, at least on a statistical basis."
+ Although it must be borne in mind that the results based on the SDSS are statistical in nature. anc not necessarily applicable to any individual object. these results would favour metallicity ellects over age ellects as the cause of the mid-infrared colour gradients.," Although it must be borne in mind that the results based on the SDSS are statistical in nature, and not necessarily applicable to any individual object, these results would favour metallicity effects over age effects as the cause of the mid-infrared colour gradients."
+ Negative. metallicity gradients and negligible age eracicnts have. in fact. been found by other studies of EZECis both in the field (Annibali et al. 2007) and cluster CRawle et al.," Negative metallicity gradients and negligible age gradients have, in fact, been found by other studies of ETGs both in the field (Annibali et al, 2007) and cluster (Rawle et al."
+ 2008. Alehlert ct al.," 2008, Mehlert et al."
+ 2003. Peletier et al.," 2003, Peletier et al."
+ 1999) using optical line strength indices., 1999) using optical line strength indices.
+ Since the work of Bower. Lucey Ellis (1992) the colour-magnitude relation (its slope and narrowness) has been open," Since the work of Bower, Lucey Ellis (1992) the colour-magnitude relation (its slope and narrowness) has been open"
+us evaluate structures of such accretion flows with simple analytical discussions.,us evaluate structures of such accretion flows with simple analytical discussions.
+" The convective motion can prevent the advective energy transport when íconv«taav, which would be fulfilled when z>zc where z. should satisfy 'Therefore, the equation of energy conservation should be modified by multiplying the factor z./H to the left-hand side of Eq. (5))"," The convective motion can prevent the advective energy transport when $t_{\rm conv}<t_{\rm adv}$, which would be fulfilled when $z>z_c$ where $z_c$ should satisfy Therefore, the equation of energy conservation should be modified by multiplying the factor $z_{c}/H$ to the left-hand side of Eq. \ref{energyeq}) )"
+ to be Other effects should be much milder than this suppression as will be discussed later., to be Other effects should be much milder than this suppression as will be discussed later.
+" Because of this suppression of the advective cooling, the significant amount of internal energy of the gas stay at the same radial position in the accretion disk."," Because of this suppression of the advective cooling, the significant amount of internal energy of the gas stay at the same radial position in the accretion disk."
+" As a result the disk temperature would be much higher than in the case when the convection is neglected, and then the disk structure is expected to be significantly changed."," As a result the disk temperature would be much higher than in the case when the convection is neglected, and then the disk structure is expected to be significantly changed."
+" In order to see this, we plot the sequence of the thermal equilibrium solutions for convective NDAFs on the (X, and (X,T) plane for various values of X, fixing the radial M)position r."," In order to see this, we plot the sequence of the thermal equilibrium solutions for convective NDAFs on the $(\Sigma, \dot{M})$ and $(\Sigma, T)$ plane for various values of $\Sigma$, fixing the radial position $r$."
+" By finding the sets of physical parameters which satisfy the fundamental equations for the disk structure, (1)), (2)), (3)) and (5)), we can derive the sequence of the thermal equilibrium solutions."," By finding the sets of physical parameters which satisfy the fundamental equations for the disk structure, \ref{massconv}) ), \ref{angmom}) ), \ref{hydrostat}) ) and \ref{energyeq}) ), we can derive the sequence of the thermal equilibrium solutions."
+ In Fig., In Fig.
+" 1 we show them on the (X,M) plane for r=4rg, varying the viscosity parameter a."," 1 we show them on the $(\Sigma, \dot{M})$ plane for $r=4r_S$, varying the viscosity parameter $\alpha$."
+" Here the mass accretion rate is normalized by the critical mass accretion rate, where Lgaa is the Eddington luminosity, with proton mass m, and the Thomson cross section or~6.6x10??cm?."," Here the mass accretion rate is normalized by the critical mass accretion rate, where $L_{\rm Edd}$ is the Eddington luminosity, with proton mass $m_p$ and the Thomson cross section $\sigma_T \simeq 6.6\times 10^{-25}{\rm cm}^2$."
+" We can see that for lower values of a the thermal equilibrium curves have a part with negative gradient, i.e. which means that this accretion flow is dynamically unstable (see, e.g. Lightman Eardley 1974; Kato et al."," We can see that for lower values of $\alpha$ the thermal equilibrium curves have a part with negative gradient, i.e. which means that this accretion flow is dynamically unstable (see, e.g. Lightman Eardley 1974; Kato et al."
+ 2008)., 2008).
+" In such accretion disks, we can expect the sporadical mass accretion onto the central black hole, which may occur in dwarf novae and some microquasars with a radiation pressure-dominated accretion disk."," In such accretion disks, we can expect the sporadical mass accretion onto the central black hole, which may occur in dwarf novae and some microquasars with a radiation pressure-dominated accretion disk."
+" This can be the source of violent time variability in the GRB prompt emissions with the mass accretion rate of a few times 10Mos""!, which corresponds to 1050—10°! erg s!."," This can be the source of violent time variability in the GRB prompt emissions with the mass accretion rate of a few times $10^{-4}M_{\odot}{\rm s}^{-1}$, which corresponds to $10^{50}-10^{51}$ erg $^{-1}$."
+ We also show in Fig., We also show in Fig.
+" 2 the thermal equilibrium curves on the (X,T) plane with the same parameter sets as in Fig."," 2 the thermal equilibrium curves on the $(\Sigma, T)$ plane with the same parameter sets as in Fig."
+ 1., 1.
+" Here we can see that in the whole parameter space the thermal stability criterion, is satisfied, and so this accretion flow is thermally stable."," Here we can see that in the whole parameter space the thermal stability criterion, is satisfied, and so this accretion flow is thermally stable."
+ Here we discuss the physical origin of the dynamical instability of NDAFs with vertical convective motions presented in the last section., Here we discuss the physical origin of the dynamical instability of NDAFs with vertical convective motions presented in the last section.
+" To see this let us express the physical properties of the accretion disk in terms of (=, M) or (X, T)."," To see this let us express the physical properties of the accretion disk in terms of $\Sigma$, $\dot{M}$ ) or $\Sigma$, $T$ )."
+" As for the former set of variables, it is obvious from Eq. ("," As for the former set of variables, it is obvious from Eq. ("
+6) that the viscous heating rate QT is proportional to M.,6) that the viscous heating rate $Q^+$ is proportional to $\dot{M}$.
+ We can also express Q* with the latter set of variables as follows., We can also express $Q^+$ with the latter set of variables as follows.
+" In the accretion flows with a mass accretion rate much larger than Mexit, the pressure and entropy are dominated by radiation, i.e. praa>Peas,Pa and Srad>>Sgas."," In the accretion flows with a mass accretion rate much larger than $\dot{M}_{\rm crit}$ , the pressure and entropy are dominated by radiation, i.e. $p_{\rm rad} \gg p_{\rm gas}, p_{\rm d}$ and $s_{\rm rad} \gg s_{\rm gas}$."
+" In this case, the scale height H would be proportional to 743;7!Q-?."," In this case, the scale height $H$ would be proportional to $T^4 \Sigma ^{-1} \Omega^{-2}$."
+" Therefore, the viscous heating rate Q*zz2apHQ should be proportional to o7?X;-1Q."," Therefore, the viscous heating rate $Q^+\approx 2\alpha p H \Omega$ should be proportional to $\alpha T^8 \Sigma ^{-1} \Omega^{-1}$."
+" On the other hand, as for the neutrino cooling processes, the emissivity by the electron/positron capture process in electron-nondegeneracy limit can be approximated as while the emissivity by the electron-positron annihilation is in the nondegeneracy limit."," On the other hand, as for the neutrino cooling processes, the emissivity by the electron/positron capture process in electron-nondegeneracy limit can be approximated as while the emissivity by the electron-positron annihilation is in the nondegeneracy limit."
+ In the unstable branch found in Fig., In the unstable branch found in Fig.
+ 1 (and thecorresponding part in Fig., 1 (and thecorresponding part in Fig.
+" 2),"," 2),"
+can be determined through the nozzle function without integrating the equation of motion. as I discuss in detail in relsecyocalsteady.,"can be determined through the nozzle function without integrating the equation of motion, as I discuss in detail in \\ref{sec_localsteady}."
+ Phisprovidesausefulnumericallest [formodelsaimedalsimulatingline—drivendiskwina, This provides a useful numerical test for models aimed at simulating line-driven disk winds.
+" This work sets a detailed framework with which we will analvze more realistic models (han the ""simple"" models of Paper 1. In a following paper (Paper III). we shall apply the framework discussed here to the case of an exact Shakura-Sunvaev disk fIux distribution and show that il the accretion disk is capable of sustaining the corresponding wind mass flow. the wind driven off a stancarel disk is steady."," This work sets a detailed framework with which we will analyze more realistic models than the “simple” models of Paper I. In a following paper (Paper III), we shall apply the framework discussed here to the case of an exact Shakura-Sunyaev disk flux distribution and show that if the accretion disk is capable of sustaining the corresponding wind mass flow, the wind driven off a standard disk is steady."
+ This is important in that. in turn. it shows that the likely scenario For the Formation of absorption (roughs in CVs is a line-driven disk wind.," This is important in that, in turn, it shows that the likely scenario for the formation of absorption troughs in CVs is a line-driven disk wind."
+ It also shows that a lne-driven accretion disk wind continues to be a promising scenario to explain the broad absorption lines in QSOs., It also shows that a line-driven accretion disk wind continues to be a promising scenario to explain the broad absorption lines in QSOs.
+ In the following paper we shall also compare the results with the cataclysmic variable disk wind models of Perevra(1997) ancl &Blondin (2000).," In the following paper we shall also compare the results with the cataclysmic variable disk wind models of \citet{per97a}
+ and \citet{per00}."
+. ] wish to thank David A. Turnshek. Stanley P. Owocki. Ikenneth. G. Gavlev. Norman W. Mura. aid D. John Hillier for many useful discussions.," I wish to thank David A. Turnshek, Stanley P. Owocki, Kenneth G. Gayley, Norman W. Murray, and D. John Hillier for many useful discussions."
+ This work is supported by the National Science Foundation under Grant AST-0071193. and by the National Aeronautics and Space Administration under Grant ATPO3-0104-0144.," This work is supported by the National Science Foundation under Grant AST-0071193, and by the National Aeronautics and Space Administration under Grant ATP03-0104-0144."
+in which the secular torque is linearly dependent on the (me derivative of tidal frequency.,in which the secular torque is linearly dependent on the time derivative of tidal frequency.
+ As we noted in 83.. Efroimskys torque is linear only in vanishinely narrow intervals around resonances (Figs.," As we noted in \ref{int.sec}, Efroimsky's torque is linear only in vanishingly narrow intervals around resonances (Figs."
+ laa and b)., \ref{tide.fig}a a and b).
+ The widely used assumption that the secular Gdal torque is linear evervwhere across the range of relative spin rates 4/1 leads to well-known problems and inconsistencies., The widely used assumption that the secular tidal torque is linear everywhere across the range of relative spin rates $\dot\theta/n$ leads to well-known problems and inconsistencies.
+ For example. a slowing down Moon in the absence of quadrupole momentum is not allowed to descend into the 111 resonance (svnchronous rotation) for anv nonzero eccen(rieitv (Murray.&Dermott2000;WilliamsEfvoimsky2012).. because the linear torque. monotonouslv increasing with growing eccentricity. changes sign Irom negative {ο positive at a certain equilibrium eccentricity [or any fixed 6/5>1.," For example, a slowing down Moon in the absence of quadrupole momentum is not allowed to descend into the 1:1 resonance (synchronous rotation) for any nonzero eccentricity \citep{mur,wil}, because the linear torque, monotonously increasing with growing eccentricity, changes sign from negative to positive at a certain equilibrium eccentricity for any fixed $\dot\theta/n > 1$."
+ Therefore. a perfectly spherical Moon should have stalled in its de-spinning at a hieher than svnchronous rate. namely. DE=460LOLS.," Therefore, a perfectly spherical Moon should have stalled in its de-spinning at a higher than synchronous rate, namely, $\dot\theta_{\rm equ}/n = 1+6e^2=1.018$."
+ This is not the case for Efvoimsky’s tidal torque. which sharply decreases al supersvncehronous rotation rates. counteracting the effects of eccentric motion.," This is not the case for Efroimsky's tidal torque, which sharply decreases at supersynchronous rotation rates, counteracting the effects of eccentric motion."
+ The dependence of equilibrium eccentricity. whence (he linear torque disappears. for a range of rotation rates is depicted in Fig.," The dependence of equilibrium eccentricity, whence the linear torque disappears, for a range of rotation rates is depicted in Fig."
+ 9. with the monotonously rising dotted line., \ref{equi.fig} with the monotonously rising dotted line.
+ It implies that synchronous rotation. so commonly observed among the satellites of the Solar Svstem. is not attaimable for any nonzero eccentricity.," It implies that synchronous rotation, so commonly observed among the satellites of the Solar System, is not attainable for any nonzero eccentricity."
+ If Mercury traversed the 3:2 resonance and continued to spin down toward svnchronization. it would be stuck at an equilibrium rate ol 1.249 with the curent. value of eecentricitv.," If Mercury traversed the 3:2 resonance and continued to spin down toward synchronization, it would be stuck at an equilibrium rate of $1.24\,n$ with the current value of eccentricity."
+ The character of equilibrium. torques is profoundly different with Elroimskys. model. depicted with (he jagged dotted curve in Fig. 3..," The character of equilibrium torques is profoundly different with Efroimsky's model, depicted with the jagged dotted curve in Fig. \ref{equi.fig}."
+ The curve was obtained by finding the rools of Eq., The curve was obtained by finding the roots of Eq.
+ 2. in e for a grid of Ü/n., \ref{sec.eq} in $e$ for a grid of $\dot\theta/n$.
+ The first jump from eq=ϐ at A/n=1 to eqq20.29 just slightly above the svuchronous rate implies (hat Mercury is allowed to be captured in the 1:1 resonance if the eccentricity is not too large (e< 0.29). evolving either from slower or faster rotation rates.," The first jump from $e_{\rm equ}=0$ at $\dot\theta/n = 1$ to $e_{\rm equ}\approx 0.29$ just slightly above the synchronous rate implies that Mercury is allowed to be captured in the 1:1 resonance if the eccentricity is not too large $e<0.29$ ), evolving either from slower or faster rotation rates."
+ However. for [aster rotation rates between 15 and 1.55». Mercury can stall in its de-spinning unless €<0.2.," However, for faster rotation rates between $1\,n$ and $1.5\,n$ , Mercury can stall in its de-spinning unless $e<0.2$."
+ Likewise. entrapment in the 3:2 resonance is inevitable for 0.2<e«0.41 if Mercury reaches this rate of rotation from either direction.," Likewise, entrapment in the 3:2 resonance is inevitable for $0.2<e<0.41$ if Mercury reaches this rate of rotation from either direction."
+ This is consistent with the results of numerical simulations im 83.1. that the probability. of 3:2 capture with the current value of eccentricity is 1., This is consistent with the results of numerical simulations in \ref{prob.sec} that the probability of 3:2 capture with the current value of eccentricity is 1.
+ What happens if Mercury. reaches A/n=L5 with an eccentricity below 0.2?, What happens if Mercury reaches $\dot\theta/n = 1.5$ with an eccentricity below 0.2?
+ Capture is still possible. but. the probability declines with decreasing eccentricity. cf. 85..," Capture is still possible, but the probability declines with decreasing eccentricity, cf. \ref{dis.sec}."
+ Entrapment of the planet in the 2:1 resonance is inevitable for 0.34«e0.49. etc.," Entrapment of the planet in the 2:1 resonance is inevitable for $0.34<e<0.49$, etc."
+ The teeth of the equilibrium torque curve act as very efficient resonance (raps for a planet like Mercury. whose eccentricityv. varies in afairly wide range over billions of vears of dvnamical evolution.," The teeth of the equilibrium torque curve act as very efficient resonance traps for a planet like Mercury, whose eccentricity varies in afairly wide range over billions of years of dynamical evolution."
+but à wavelength slightly shifted.,but a wavelength slightly shifted.
+ We were unable to detect A4077.714 at all (<0.6mA)). nor did (1996) report seeing this line.," We were unable to detect $\lambda$ 4077.714 at all $\leq$ ), nor did \citet{ryab-etal96} report seeing this line."
+" The solution to this anomaly is our identification of a blend of wwith A4215.60 (Llansen&PerssonLOST).. which has an estimated strength of iin Wer from the abundance in Table 4. and an astrophysical oof — 1.06. based on the Lick observations of the ultra-sharp-lined. Sr-deficient. ""mild"" Lle\In star Xal (Cowley1980:Smith&Dworetsky 1903)."," The solution to this anomaly is our identification of a blend of with $\lambda$ 4215.60 \citep{hp87}, which has an estimated strength of in Her from the abundance in Table \ref{abunds} and an astrophysical of $-1.06$ , based on the Lick observations of the ultra-sharp-lined, Sr-deficient, “mild” HgMn star Aql \citep{cow80,sd93}."
+. The remaining strength. of this feature in 1ος is probably. due to weak lines of andCriL., The remaining strength of this feature in Her is probably due to weak lines of and.
+ Strontium is at least 1.0 dex below solar abundance in both 46XXal and Ler., Strontium is at least 1.0 dex below solar abundance in both Aql and Her.
+ The authors are grateful forthe assistance of 2006 MSc, The authors are grateful forthe assistance of 2006 MSc
+computed using point sources common to all images.,computed using point sources common to all images.
+ The nine images are mosaiced together by applying offsets in intensity to the registered images to produce the most consistent sky value possible in the overlap regions., The nine images are mosaiced together by applying offsets in intensity to the registered images to produce the most consistent sky value possible in the overlap regions.
+" A new flatfield and mosaic is constructed for each set of nine jittered frames, then all the mosaics registered and co-added to form a master mosaic."," A new flatfield and mosaic is constructed for each set of nine jittered frames, then all the mosaics registered and co-added to form a master mosaic."
+" Occasionally, significant variations in the level and/or structure of the background sky on a temporal or spatial scale smaller than that sampled by the array within the ~5 mminute period of the nine jittered frames resulted in noticeable residual structure in the ensuing mosaic, forcing us to exclude that mosaic from the master mosaic."," Occasionally, significant variations in the level and/or structure of the background sky on a temporal or spatial scale smaller than that sampled by the array within the $\sim$ minute period of the nine jittered frames resulted in noticeable residual structure in the ensuing mosaic, forcing us to exclude that mosaic from the master mosaic."
+" The total on-source exposure time, after discarding such data, is shown in Table "," The total on-source exposure time, after discarding such data, is shown in Table \ref{basicdata}."
+"For B1.CSJ observations, the six sky frames are first offset in intensity to a common modal value, then a flatfield formed from the median value at each pixel."," For CSJ observations, the six sky frames are first offset in intensity to a common modal value, then a flatfield formed from the median value at each pixel."
+" All six sky frames, and five object frames are flatfielded, then the modal pixel values of the two sky frames bracketing each object frame are averaged and subtracted from that object frame."," All six sky frames, and five object frames are flatfielded, then the modal pixel values of the two sky frames bracketing each object frame are averaged and subtracted from that object frame."
+ Image registration and mosaicing is then performed on each set of five sky-subtracted object frames just as for the JSF observations., Image registration and mosaicing is then performed on each set of five sky-subtracted object frames just as for the JSF observations.
+" These mosaics are registered and co-added to form a master mosaic, with the exception of any showing residual sky structure as described above, yielding the on-source exposure times shown in Table J."," These mosaics are registered and co-added to form a master mosaic, with the exception of any showing residual sky structure as described above, yielding the on-source exposure times shown in Table \ref{basicdata}."
+" Of the 68 galaxies observed, 11 remained undetected or could not be usefully analysed."," Of the 68 galaxies observed, 11 remained undetected or could not be usefully analysed."
+" This was either because our H-band surface brightness limit of [jim7 25mmagaarcsec? or Lo aarcsec? at a distance of 1MMpc Mg,c=3.35 mmag, was not low enough or the (adoptinggalaxy light was heavily ?))contaminated by Galactic foreground stars (see Table "," This was either because our $H$ -band surface brightness limit of $\mu_{lim}\approx 25$ $^{-2}$ or $\,L_{\odot}$ $^{-2}$ at a distance of Mpc (adopting $M_{H,\odot}=3.35$ mag, \citealt{colina96}) ) was not low enough or the galaxy light was heavily contaminated by Galactic foreground stars (see Table \ref{badgals}) )."
+In both situations the data was not processed [).further., In both situations the data was not processed further.
+" For instance, the companion galaxies NGC2784 DWI and KK98-73 were observed parallel to NGC2784."," For instance, the companion galaxies NGC2784 DW1 and KK98-73 were observed parallel to NGC2784."
+" While NGC2784 and KK98-73 can be seen, NGC2784 DWI1, located between NGC2784 and KK98-73, is barely visible (see Figure second row, middle panel where KK98-73 is visible to pl.the bottom left of the image)."," While NGC2784 and KK98-73 can be seen, NGC2784 DW1, located between NGC2784 and KK98-73, is barely visible (see Figure \ref{fig:galimages4}, second row, middle panel where KK98-73 is visible to the bottom left of the image)."
+ The images of the other 57 galaxies are shown in Figures [] to [6., The images of the other 57 galaxies are shown in Figures \ref{fig:galimages1} to \ref{fig:galimages5}.
+" On each image, instrumental magnitudes for 50-100 field stars were measured employing standard IRAF PSF fitting routines."," On each image, instrumental magnitudes for 50–100 field stars were measured employing standard IRAF PSF fitting routines."
+ Cross-correlating the stellar positions with the 2MASS Point Source Catalog provided H-band magnitudes and allowed the photometric calibration of each field (see Figure [7))., Cross-correlating the stellar positions with the 2MASS Point Source Catalog provided $H$ -band magnitudes and allowed the photometric calibration of each field (see Figure \ref{fig:zeropoint}) ).
+ The stars which deviate from the 45 degree line were usually either extremely red or blue where the transformation between 2MASS and IRIS2 H-bands (?) breaks down., The stars which deviate from the 45 degree line were usually either extremely red or blue where the transformation between 2MASS and IRIS2 $H$ -bands \citep{ryder07} breaks down.
+ The lo uncertainty in the zero point was calculated to be between 0.01 and mmag depending on the number of stars used for the calibration., The $\sigma$ uncertainty in the zero point was calculated to be between 0.01 and mag depending on the number of stars used for the calibration.
+" 'To ensure accurate galaxy surface photometry down to the faintest possible isophotes, the images were cleaned of foreground stars using procedures written within the IRAF package."," To ensure accurate galaxy surface photometry down to the faintest possible isophotes, the images were cleaned of foreground stars using procedures written within the IRAF package."
+" Thereby, stars in the field around a galaxy were carefully replaced with nearby patches of plain sky."," Thereby, stars in the field around a galaxy were carefully replaced with nearby patches of plain sky."
+" If superposed on the galaxy, the galaxy light under the star was restored by replacing the contaminated area with its mirror image with respect to the galaxy center."," If superposed on the galaxy, the galaxy light under the star was restored by replacing the contaminated area with its mirror image with respect to the galaxy center."
+ The galaxy center was defined as the center of the luminosity-weighted light distribution., The galaxy center was defined as the center of the luminosity-weighted light distribution.
+" The star removal process was monitored visually to identify small-scale structures and asymmetries, and to ensure accurate removal of foreground stars whilst not removing sources associated with the galaxy itself."," The star removal process was monitored visually to identify small-scale structures and asymmetries, and to ensure accurate removal of foreground stars whilst not removing sources associated with the galaxy itself."
+ T'he effectiveness of the cleaning procedure is illustrated in Figure |8} where pre- and post- cleaning images are shown for the two galaxies ESO468-G020 and IC1959., The effectiveness of the cleaning procedure is illustrated in Figure \ref{fig:cleaning} where pre- and post- cleaning images are shown for the two galaxies ESO468-G020 and IC1959.
+ The fields of IC5152 (Figure [4)) and UGCA438 (Figure |6)) both have a bright foreground star which obscures a large portion of the galaxy., The fields of IC5152 (Figure \ref{fig:galimages3}) ) and UGCA438 (Figure \ref{fig:galimages5}) ) both have a bright foreground star which obscures a large portion of the galaxy.
+" To obtain rough photometric parameters for these objects, the contribution of the bright stars to the total flux needed to be removed."," To obtain rough photometric parameters for these objects, the contribution of the bright stars to the total flux needed to be removed."
+" Because such a large portion of both the galaxies is obscured, it was necessary to remove the affected quarter of the image and replace it with the opposite quarter, which was rotated by 180 degrees (i.e., in the case of IC5152, the top-right quarter was replaced with the rotated bottom-left quarter of the image)."," Because such a large portion of both the galaxies is obscured, it was necessary to remove the affected quarter of the image and replace it with the opposite quarter, which was rotated by 180 degrees (i.e., in the case of IC5152, the top-right quarter was replaced with the rotated bottom-left quarter of the image)."
+" This method failed in the case of the dIrr galaxy AM0521-343, where an even brighter star (CD-34 2225, my=7.1) near the faint dwarf galaxy prevented proper cleaning (see Figure [2))."," This method failed in the case of the dIrr galaxy AM0521-343, where an even brighter star (CD-34 2225, $m_H=7.1$ ) near the faint dwarf galaxy prevented proper cleaning (see Figure \ref{fig:galimages1}) )."
+ Simulated circular aperture photometry of the star-subtracted H-band images produced a growth curve as a function of the geometric mean radius Vab (where a and b are the galaxy’s major and minor axes)., Simulated circular aperture photometry of the star-subtracted $H$ -band images produced a growth curve as a function of the geometric mean radius $\sqrt{ab}$ (where $a$ and $b$ are the galaxy's major and minor axes).
+" The asymptotic intensity corresponds to the total apparent magnitude, my, that can be recovered down to the background noise level of the image."," The asymptotic intensity corresponds to the total apparent magnitude, $m_H$, that can be recovered down to the background noise level of the image."
+ The largest source of uncertainty is the sky level., The largest source of uncertainty is the sky level.
+ By systematically varying the sky brightness we determined which growth curve converges best to a plateau as far as possible from the center of the galaxy., By systematically varying the sky brightness we determined which growth curve converges best to a plateau as far as possible from the center of the galaxy.
+" We measure the half-light geometric mean radius, at half the asymptotic intensity and calculated the reff,mean surface brightness within that radius: (µΗ)εβε."," We measure the half-light geometric mean radius, $r_{eff}$, at half the asymptotic intensity and calculated the mean surface brightness within that radius: $\langle\mu_H\rangle_{eff}$."
+" The overall uncertainty for the total magnitude, my, is between 0.05 and 0.30mmag; for the mean effective surface brightness, (uH)eff, less than 0.2mmag arcsec~?, and for the half-light radius, ref, is of the order of five percent."," The overall uncertainty for the total magnitude, $m_H$, is between 0.05 and mag; for the mean effective surface brightness, $\langle\mu_H\rangle_{eff}$, less than mag $^{-2}$, and for the half-light radius, $r_{eff}$, is of the order of five percent."
+" The image of AM0521-343 (Figure contains a bright foreground star and thus, the sky DJ),brightness plus the contribution from the stellar halo was estimated at the galaxy’s position and simulated aperture photometry performed out to the radius of asymptotic intensity (defined on side of the galaxy opposite to the contaminating "," The image of AM0521-343 (Figure \ref{fig:galimages1}) ), contains a bright foreground star and thus, the sky brightness plus the contribution from the stellar halo was estimated at the galaxy's position and simulated aperture photometry performed out to the radius of asymptotic intensity (defined on side of the galaxy opposite to the contaminating star)."
+The radial surface brightness profile of a star).sample galaxy was determined by differentiating the growth curves with respect to radius., The radial surface brightness profile of a sample galaxy was determined by differentiating the growth curves with respect to radius.
+" Depending on the total integration time of the image, the profiles could be reconstructed down to a surface brightness limit between 24.5mmag arcsec-?«[lim 26mmag arcsec~?."," Depending on the total integration time of the image, the profiles could be reconstructed down to a surface brightness limit between $24.5$ mag ${}^{-2}< \mu_{lim}<26$ mag ${}^{-2}$."
+" They are shown, with a linear radius scale in Figures [}, and [I1]."," They are shown, with a linear radius scale in Figures \ref{figsbprof1}, \ref{figsbprof2} and \ref{figsbprof3}."
+ The error bars are calculated as the rms scatter of the intensity along each isophote., The error bars are calculated as the rms scatter of the intensity along each isophote.
+" The position angle and ellipticity for each isophote, with the galaxy center fixed, were calculated with IRAF's ELLIPSE package as a function of radius."," The position angle and ellipticity for each isophote, with the galaxy center fixed, were calculated with IRAF's ELLIPSE package as a function of radius."
+" The ellipticity and position angle were, in general, settled in the outer regions of the galaxies, however, in some cases they varied significantly"," The ellipticity and position angle were, in general, settled in the outer regions of the galaxies, however, in some cases they varied significantly"
+In order (ο [ind separate temperature parameterizaüons for and). determine the ICF for obtaining N/O from.. and carefully estimate uncertainties in N/O. we computed a grid of photoionization models. representative of low metallicity svstenms.,"In order to find separate temperature parameterizations for and, determine the ICF for obtaining N/O from, and carefully estimate uncertainties in N/O, we computed a grid of photoionization models, representative of low metallicity systems."
+ In the following paragraphs. we introducePIIOENIX.. the stellar atinosphere code used to generate the spectra that served as input to our photoionization models. we explain (he choice of parameter space covered by our stellar and nebular simulations. we describe how we proceeded to obtain the temperature parameterizations and ICF from our models. aud we derive (he electron temperatures. electron densities. aud ionic abundances of our sample objects.," In the following paragraphs, we introduce, the stellar atmosphere code used to generate the spectra that served as input to our photoionization models, we explain the choice of parameter space covered by our stellar and nebular simulations, we describe how we proceeded to obtain the temperature parameterizations and ICF from our models, and we derive the electron temperatures, electron densities, and ionic abundances of our sample objects."
+ is a general stellar atanosphere code which operates in one-dimensional spherically svnunelric or plane-parallel &eometry., is a general stellar atmosphere code which operates in one-dimensional spherically symmetric or plane-parallel geometry.
+ Ho can handle static atmospheres (e.g. stars. Jupiter-like planets. brown clwarls) as well as moving alimospheres (supernovae. novae. massive hot stars. ete).," It can handle static atmospheres (e.g. stars, Jupiter-like planets, brown dwarfs) as well as moving atmospheres (supernovae, novae, massive hot stars, etc.)."
+ The code simultaneously solves the equations of radiative (ransler (in the proper geometry). structure. and level population in Non-Local Thermocdsninonmic Equilibrium CNLTE).," The code simultaneously solves the equations of radiative transfer (in the proper geometry), structure, and level population in Non-Local Thermodynamic Equilibrium (NLTE)."
+ We used PHOENIX’ss wind module package (Aufdenbergetal.2002) to generate svithetie spectra in the range [rom 100 to 8x 109A. for six windy O-cdwarl stars with solar/5. solar/20. or solar/50 metal abundances. except lor N and C. which we scaled in order to be consistent wilh twpical proportions found in low metallicity systems. i.e. log(N/O) = -1.46 and log(C/O) = -0.7 (Henryοἱal.2000).," We used s wind module package \citep{aufdenberg02} to generate synthetic spectra in the range from 100 to $\times$ $^6$, for six windy O-dwarf stars with solar/5, solar/20, or solar/50 metal abundances, except for N and C, which we scaled in order to be consistent with typical proportions found in low metallicity systems, i.e., log(N/O) = -1.46 and log(C/O) = -0.7 \citep{henry00}."
+. The set of solar abuncdances was modified in order to match that of the code we used for our photoionizationmodels*., The set of solar abundances was modified in order to match that of the code we used for our photoionization.
+. Other input parameters to the models were based on those emploved by Smithetal.(2002).. who generated. svnthetic spectra for the same tvpe of stars. wilh solar/5 and solar/20 proportions for all metals (including N and C).," Other input parameters to the models were based on those employed by \cite{smith02}, who generated synthetic spectra for the same type of stars, with solar/5 and solar/20 proportions for all metals (including N and C)."
+" We adopted a o-law for the wind velocity. Le. where r is (he distance [rom (he center of the star. £2, is (he radius of (he ""photosphere? (considered to be the base of the wind). and vr. is the terminal velocity. assumed. to be reached al LOO x H,."," We adopted a $\beta$ -law for the wind velocity, i.e., where r is the distance from the center of the star, $R_*$ is the radius of the “photosphere” (considered to be the base of the wind), and $v_{\infty}$ is the terminal velocity, assumed to be reached at 100 $\times$ $R_*$."
+ The 9-index was assumed to be 0.3 (Lamers&Cassinelli1999)., The $\beta$ -index was assumed to be 0.8 \citep{lamers99}.
+. The density structure in the dvnamic region (ος)2 0) was caleulated from the continuity equation:, The density structure in the dynamic region $v(r)>0$ ) was calculated from the continuity equation:
+with the same dimensionless (see Table 1)).,with the same dimensionless (see Table \ref{tablefreq}) ).
+ It can be seen that for PMS and MS models the main driving mechanism leading to the excitation of the selected modes occurs at the bottom of the CE., It can be seen that for PMS and MS models the main driving mechanism leading to the excitation of the selected modes occurs at the bottom of the CE.
+ Both models have the same CE depth (see Sect. 3).," Both models have the same CE depth (see Sect. \ref{internalstructure}) ),"
+ hence their non-adiabatie functions present the same behaviour 1n the external layers., hence their non-adiabatic functions present the same behaviour in the external layers.
+ The only differences between PMS and MS eigenfunctions are located in the inner layers of the models and are due to their different central structures., The only differences between PMS and MS eigenfunctions are located in the inner layers of the models and are due to their different central structures.
+ Nevertheless. in both cases. the amplitude of οoP in the central layers is small. and the differences between both stellar structures have no impact on the excitation of the modes.," Nevertheless, in both cases, the amplitude of $\frac{\delta P}{P}$ in the central layers is small, and the differences between both stellar structures have no impact on the excitation of the modes."
+ Fig., Fig.
+ 12. illustrates the variation in the term /(v) (Eq. (11))), \ref{dampingsamemass} illustrates the variation in the term $I(x)$ (Eq. \ref{dimlessradiativedamping}) ))
+ for the 1.8M.. PMS and MS models., for the $1.8 M_\odot$ PMS and MS models.
+ In this case. Cx) Cor what is the same. the radiative damping) behaves the same in both models. excepted in their very internal layers.," In this case, $I(x)$ (or what is the same, the radiative damping) behaves the same in both models, excepted in their very internal layers."
+ However. the important quantity to compare is the integral of /(x) on the radiative zone. /.e. the area under each curve.," However, the important quantity to compare is the integral of $I(x)$ on the radiative zone, $i.e.$ the area under each curve."
+ The close values of 7 1n both models explain the similarity of the period domain of unstable modes., The close values of $\eta$ in both models explain the similarity of the period domain of unstable modes.
+ We focus now on the non-adiabatic behaviour of the 2.1M.. PMS and the 1.9M. MS models. which are located close to the red border of the y Dor IS. where MS and PMS g- period ranges differ (see Fig. 9)).," We focus now on the non-adiabatic behaviour of the $2.1 M_\odot$ PMS and the $1.9 M_\odot$ MS models, which are located close to the red border of the $\gamma$ Dor IS, where MS and PMS $g$ -modes period ranges differ (see Fig. \ref{periodteff}) )."
+ Fig., Fig.
+ 13. presents the y Dor period ranges of these two models., \ref{perioddiffmass} presents the $\gamma$ Dor period ranges of these two models.
+ While the number of unstable modes is the same for both models. the period domain corresponding to these modes ts smaller for the MS model than for the PMS one.," While the number of unstable modes is the same for both models, the period domain corresponding to these modes is smaller for the MS model than for the PMS one."
+ The different N profiles lead to different periods and period spacing values (see Eq. (2)), The different $N$ profiles lead to different periods and period spacing values (see Eq. \ref{asymptoticperiod}) )
+ and Eq. (3)))., and Eq. \ref{meanperiodspacing}) )).
+ As a consequence. the mode density is higher in the MS model than in the PMS one. and the instability range is narrower.," As a consequence, the mode density is higher in the MS model than in the PMS one, and the instability range is narrower."
+ As for the 1.844 models (Fig. ΕΤ).," As for the $1.8 M_\odot$ models (Fig. \ref{worksamemass}) ),"
+ we compare in Fig., we compare in Fig.
+ 14 the work integrals and eigenfunctions for two modes having the same dimensionless frequency., \ref{workdiffmass} the work integrals and eigenfunctions for two modes having the same dimensionless frequency.
+ This mode is unstable in, This mode is unstable in
+RS Ophiuchi is a symbiotic recurrent nova (Kenyon 1986)). that has undergone a number of recorded eruptions. with an inter-outburst period of 10-20 years.,"RS Ophiuchi is a symbiotic recurrent nova (Kenyon \cite{kenyon}) ), that has undergone a number of recorded eruptions, with an inter-outburst period of 10-20 years."
+ The binary system is composed by a K4-MO red giant (RG) star in a 456 day orbit with a white dwarf (WD) of mass close to the Chandrasekhar limit (Dobrzycka Kenyon 1994:: Shore et al. 1996::, The binary system is composed by a K4-M0 red giant (RG) star in a 456 day orbit with a white dwarf (WD) of mass close to the Chandrasekhar limit (Dobrzycka Kenyon \cite{danuta}; Shore et al. \cite{shore};
+ Fekel et al. 20000 , Fekel et al. \cite{fekel}; ;
+Brandi et al. 2009))., Brandi et al. \cite{brandi}) ).
+ For this reason. RS Oph has been proposed as the prototype for one of the possible channels to Type la Supernova explosions (see for example Hachisu Kato 2000.2001:: Hernanz José 2008)).," For this reason, RS Oph has been proposed as the prototype for one of the possible channels to Type Ia Supernova explosions (see for example Hachisu Kato \cite{hachisu00,hachisu01}; Hernanz José \cite{hernanz}) )."
+ The WD is supposed to acerete part of the material lost by the RG via stellar wind. until it reaches the conditions for thermonuclear burning at its surface.," The WD is supposed to accrete part of the material lost by the RG via stellar wind, until it reaches the conditions for thermonuclear burning at its surface."
+ When this happens. the WD ejects 1077 to 107 M. at velocities of «5.000 km s! (Hachisu Kato 2001: Yaron et al. 2005)).," When this happens, the WD ejects $^{-8}$ to $^{-6}$ $_\odot$ at velocities of $\sim$ 5,000 km $^{-1}$ (Hachisu Kato \cite{hachisu01}; Yaron et al. \cite{yaron}) ),"
+ producing à shock wave that propagates into the pre-existing RG wind and gives rise to strong X-ray emission. as observed in the 1985 eruption (Mason et al. 1987)).," producing a shock wave that propagates into the pre-existing RG wind and gives rise to strong X-ray emission, as observed in the 1985 eruption (Mason et al. \cite{mason87}) )."
+ The last outburst of RS Oph took place on 2006 February 12.83 UT (Narumi et al. 2006)).," The last outburst of RS Oph took place on 2006 February 12.83 UT (Narumi et al. \cite{narumi}) ),"
+ and prompt radio observations have shown that the nova ejection was bipolar (O’Brien et al. 2006))., and prompt radio observations have shown that the nova ejection was bipolar (O'Brien et al. \cite{obrien}) ).
+ Besides confirming that the blast wave clearly deviated from spherical symmetry. the Ray data (Bode et al. 2006:," Besides confirming that the blast wave clearly deviated from spherical symmetry, the X-Ray data (Bode et al. \cite{bode06};"
+ Sokoloski et al. 2006)), Sokoloski et al. \cite{sokoloski}) )
+ have indicated a low mass for the ejected material which. in turn. favors a high mass for the WD (~1.4 M..). giving further support to the proposed connection between RS Oph and Type Ia progenitors (Sokoloski et al. 20060).," have indicated a low mass for the ejected material which, in turn, favors a high mass for the WD $\sim$ 1.4 $_\odot$ ), giving further support to the proposed connection between RS Oph and Type Ia progenitors (Sokoloski et al. \cite{sokoloski}) )."
+ One important question that remains unanswered. is whether the hot component of RS Oph is a O-Ne-Mg WD rather than a C-O WD. for there are no direct observations that could help to distinguish between the two.," One important question that remains unanswered, is whether the hot component of RS Oph is a O-Ne-Mg WD rather than a C-O WD, for there are no direct observations that could help to distinguish between the two."
+ In this case. the WD would not experience a thermonuclear explosion. leading to a SN Ia. but rather undergo an electron-capture core collapse. producing a neutron-star remnant (Nomoto Kondo 1991)). when its mass gets close to the Chandrasekhar mass.," In this case, the WD would not experience a thermonuclear explosion, leading to a SN Ia, but rather undergo an electron-capture core collapse, producing a neutron-star remnant (Nomoto Kondo \cite{nomoto}) ), when its mass gets close to the Chandrasekhar mass."
+ In addition to the clear detection of circumstellar matter (CSM) very close to the system (O'Brien. et al. 2006:, In addition to the clear detection of circumstellar matter (CSM) very close to the system (O'Brien et al. \cite{obrien};
+ Sokoloski et al. 2006)).," Sokoloski et al. \cite{sokoloski}) ),"
+ Iijima (2008.2009)) has reported the presence of a narrow D absorption component. which ts blue-shifted by about 33 km s! with respect to the systemic velocity of RS Oph.," Iijima \cite{iijima,iijima09}) ) has reported the presence of a narrow D absorption component, which is blue-shifted by about 33 km $^{-1}$ with respect to the systemic velocity of RS Oph."
+ He proposed that this feature might arise in the wind of the RG. hence arguing for its circumstellar nature.," He proposed that this feature might arise in the wind of the RG, hence arguing for its circumstellar nature."
+ Stimulated by this finding and by the detection of D time-variant features in the Type la SN 20068 (Patat et al. 2007).," Stimulated by this finding and by the detection of D time-variant features in the Type Ia SN 2006X (Patat et al. \cite{patat}) ),"
+ SN 20071e (Simon et al. 20091).," SN 2007le (Simon et al. \cite{simon}) ),"
+ 1999e] (Blondin et al. 2008)).," 1999cl (Blondin et al. \cite{blondin}) ),"
+ and SN 2006dd (Stritzinger et al. 2010) , and SN 2006dd (Stritzinger et al. \cite{max}) )
+we performed a thorough analysis using pre/post-outburst multi-epoch. high-resolution spectroscopy of RS Oph. whose results we are reporting here.," we performed a thorough analysis using pre/post-outburst multi-epoch, high-resolution spectroscopy of RS Oph, whose results we are reporting here."
+ RS Oph was observed in the pre-outburst. outburst and post-outburst phases using the Fiber-fed Extended Range Optical Spectrograph (FEROS: Kaufer et al. 1999))," RS Oph was observed in the pre-outburst, outburst and post-outburst phases using the Fiber-fed Extended Range Optical Spectrograph (FEROS; Kaufer et al. \cite{kaufer}) )"
+ mounted at the ESO-MPG 2.2m telescope., mounted at the ESO-MPG 2.2m telescope.
+ The pre-outburst data set was presented by Zamanov et al. (2005)).," The pre-outburst data set was presented by Zamanov et al. \cite{zamanov}) ),"
+ while the outburst data are going to be thoroughly discussed in a forthcoming paper. to which we refer the reader for a full report on the high- spectroscopic follow-up of RS Oph (Mason et al.," while the outburst data are going to be thoroughly discussed in a forthcoming paper, to which we refer the reader for a full report on the high-resolution spectroscopic follow-up of RS Oph (Mason et al."
+ in preparation)., in preparation).
+ In order to cover the quiescence phase following the outburst. an additional set of data was obtained in February 2008. about two years after the 2006 eruption.," In order to cover the quiescence phase following the outburst, an additional set of data was obtained in February 2008, about two years after the 2006 eruption."
+ The instrument delivers a resolving. power A4- 48.000. covering the spectral range 3700-9200A.," The instrument delivers a resolving power $\lambda/\Delta \lambda
+\sim$ 48,000, covering the spectral range 3700–9200."
+. The data were processed by the FEROS pipeline. which includes the correction for Earth’s motion.," The data were processed by the FEROS pipeline, which includes the correction for Earth's motion."
+ The first set of data was obtained on four epochs between 672 and 530 days before the February 2006 outburst. whilst the second covers 5 epochs between 1.5 and 46.6 days after the eruption.," The first set of data was obtained on four epochs between 672 and 530 days before the February 2006 outburst, whilst the second covers 5 epochs between 1.5 and 46.6 days after the eruption."
+ Finally. RS Oph was observed on three consecutive nights about 740 days after the eruption. when the systemwas well back into the quiescence phase. and the mass accretion had resumed (Worters et al. 2007)).," Finally, RS Oph was observed on three consecutive nights about 740 days after the eruption, when the systemwas well back into the quiescence phase, and the mass accretion had resumed (Worters et al. \cite{worters}) )."
+ The spectra obtained on these three nights were stacked in order, The spectra obtained on these three nights were stacked in order
+2005);; we find that this effect sets in at even lower optical depths than previously anticipated.,; we find that this effect sets in at even lower optical depths than previously anticipated.
+" Moreover, as we mentioned above, the debris disks we see around other stars are much more massive than the KB, making collisions even more important; virtually all known extrasolar debris disks show “collision-dominated” behavior."," Moreover, as we mentioned above, the debris disks we see around other stars are much more massive than the KB, making collisions even more important; virtually all known extrasolar debris disks show “collision-dominated” behavior."
+" A recent paper by Vitenseetal.(2010) illustrates the importance of these grain-grain collisions on the Kuiper belt dust distribution, thought it does not model the dynamical effects of the planets."," A recent paper by \citet{vite10} illustrates the importance of these grain-grain collisions on the Kuiper belt dust distribution, thought it does not model the dynamical effects of the planets."
+" In this paper, we take a step toward a better understanding of the analogy between the KB and extrasolar debris disks."," In this paper, we take a step toward a better understanding of the analogy between the KB and extrasolar debris disks."
+ We use our new “collisional groming” algorithm (Stark&2009) to explore the effects of grain-grain collisions and planetary perturbations together on the distribution of Kuiper Belt dust.," We use our new “collisional groming"" algorithm \citep{star09} to explore the effects of grain-grain collisions and planetary perturbations together on the distribution of Kuiper Belt dust."
+" We break the KB source population into three populations: hot, cold and plutinos."," We break the KB source population into three populations: hot, cold and plutinos."
+" We model the effects of grain-grain collisions in today’s KB, and we model how the KB dust morphology would change if the amount of dust were increased from a face-on optical depth of ~1077 to ~1079, ~107°, and ~1074."," We model the effects of grain-grain collisions in today's KB, and we model how the KB dust morphology would change if the amount of dust were increased from a face-on optical depth of $\sim 10^{-7}$ to $\sim 10^{-6}$, $\sim 10^{-5}$, and $\sim 10^{-4}$."
+ Here is a brief summary of the collision grooming algorithm., Here is a brief summary of the collision grooming algorithm.
+ Stark&Kuchner(2009) described the algorithm in depth and various numerical tests it has passed., \citet{star09} described the algorithm in depth and various numerical tests it has passed.
+" First the orbits of a set of dust grains are numerically integrated using an n-body integrator, and the positions and velocities of the particles are recorded periodically in a histogram."," First the orbits of a set of dust grains are numerically integrated using an n-body integrator, and the positions and velocities of the particles are recorded periodically in a histogram."
+ We call this histogram the “seed” model.," We call this histogram the “seed"" model."
+" Then the trajectories of each particle are re-interpreted as steady state streams of particles, with weights that define the number of particles in each stream at any given point along the trajectory."," Then the trajectories of each particle are re-interpreted as steady state streams of particles, with weights that define the number of particles in each stream at any given point along the trajectory."
+ The weights are then iteratively manipulated so that they describe a self-consistent cloud of interacting particles., The weights are then iteratively manipulated so that they describe a self-consistent cloud of interacting particles.
+" The result is a 3-D grid that contains the number density of the cloud, a self-consistent solution to both the dynamical equations that govern the particle trajectories and the number flux equation that accounts for the creation and destruction of particles in every histogram bin."," The result is a 3-D grid that contains the number density of the cloud, a self-consistent solution to both the dynamical equations that govern the particle trajectories and the number flux equation that accounts for the creation and destruction of particles in every histogram bin."
+ There have been several recent papers on kinetic treatments of collisions in debris disks (e.g.Krivovetal.2006;Wyatt2007) including the Kuiper Belt (Krivovetal.2005)..," There have been several recent papers on kinetic treatments of collisions in debris disks \citep[e.g.][]{kriv06, wyat07} including the Kuiper Belt \citep{kriv05}."
+" Some of these models involve more detailed collision physics than our simulations, e.g., time evolution and fragmentation."," Some of these models involve more detailed collision physics than our simulations, e.g., time evolution and fragmentation."
+" But our simulations have the unique capability to explore the interaction between the dynamical effects of planets, such as resonances, secular forcing, etc.,"," But our simulations have the unique capability to explore the interaction between the dynamical effects of planets, such as resonances, secular forcing, etc.,"
+ and grain-grain collisions., and grain-grain collisions.
+" These pheneomena are crucial for understanding the distribution of dust in debris disks, even in the presence of collisions, as we will show."," These pheneomena are crucial for understanding the distribution of dust in debris disks, even in the presence of collisions, as we will show."
+" To create the seed model for our study, we integrated the equations of motion for a"," To create the seed model for our study, we integrated the equations of motion for a"
+the photosphere. which is taken as an infinite plane.,"the photosphere, which is taken as an infinite plane."
+ The topology of the ends of the flux rope. which in reality. curve down to meet the photosphere in a full 3D configuration. is left out of the 2D description.," The topology of the ends of the flux rope, which in reality curve down to meet the photosphere in a full 3D configuration, is left out of the 2D description."
+ This description captures features of fIux ropes the main part of whose lengths are suspended in the atmosphere above the photosphere., This description captures features of flux ropes the main part of whose lengths are suspended in the atmosphere above the photosphere.
+ The second reason is that in 3D the MITIS Force-balance equation becomes hiehly nonlinear. and a follow-up stability analvsis becomes even more challenging as the effect of line-(vinge al the photospheric end points must be taken into account as well.," The second reason is that in 3D the MHS force-balance equation becomes highly nonlinear, and a follow-up stability analysis becomes even more challenging as the effect of line-tying at the photospheric end points must be taken into account as well."
+ Furthermore. the 3D MIIS force-balance equation changes (vpe Irom elliptic to mixed (wpe. causing the character of the boundary. value problem to change.," Furthermore, the 3D MHS force-balance equation changes type from elliptic to mixed type, causing the character of the boundary value problem to change."
+ The 2D MIIS (Grad-Shafranov) equation has a linear. elliptic differential operator so that its characteristic curves are all imaginary.," The 2D MHS (Grad-Shafranov) equation has a linear, elliptic differential operator so that its characteristic curves are all imaginary."
+ Thus. like fields© described by. e.g.. Laplaces equation V70=0 or the linear wave equation V206+A26=0. solutions i ofthe Grad-Shalranov equation are continuous functions uniquely determined by specifvine their value. normal derivative or linear combination of the two on the boundary.," Thus, like fields$\phi$ described by, e.g., Laplace's equation $\nabla^2\phi =0$ or the linear wave equation $\nabla^2\phi +k^2\phi =0$, solutions $\psi$ of the Grad-Shafranov equation are continuous functions uniquely determined by specifying their value, normal derivative or linear combination of the two on the boundary."
+ The situation with the 3D MIIS problem is very different (Parker 1979. 1994): the inclusion of (hree-climensional variations causes the force-balance equation to change (wpe from a simple elliptic equation with purely imaginary characteristics to an equation of mixed type with both imaginary and real characteristics.," The situation with the 3D MHS problem is very different (Parker 1979, 1994): the inclusion of three-dimensional variations causes the force-balance equation to change type from a simple elliptic equation with purely imaginary characteristics to an equation of mixed type with both imaginary and real characteristics."
+ Across the real characteristic eurves field derivatives may not be defined., Across the real characteristic curves field derivatives may not be defined.
+ Since these real characteristics coincide with field trajectories. tangential discontinuities are permitted there provided that (he total pressure [DB]?/8z+p is balanced across (he cdiscontinuity.," Since these real characteristics coincide with field trajectories, tangential discontinuities are permitted there provided that the total pressure $|{\bf B}|^2/8\pi +p$ is balanced across the discontinuity."
+ Ii general. equilibrium solvers break down for mixed-tvpe problems and it can been questioned whether the equilibrium and perturbation can be meaningfully distinguished in these mixed-tvpe regimes. as has been «questioned lor hyperbolic regimes in mixed-tvpe problems (Goedbloed 2003. 2004).," In general, equilibrium solvers break down for mixed-type problems and it can been questioned whether the equilibrium and perturbation can be meaningfully distinguished in these mixed-type regimes, as has been questioned for hyperbolic regimes in mixed-type problems (Goedbloed 2003, 2004)."
+ The stability properties of these new equilibria. which may be relevant to solar eruptions. can be determined in (he form of a [ull resistive MIID spectrum by solving the linearized MID equations using the companion hyperbolic stability solver PHOENIX (Bloklanel et al.," The stability properties of these new equilibria, which may be relevant to solar eruptions, can be determined in the form of a full resistive MHD spectrum by solving the linearized MHD equations using the companion hyperbolic stability solver PHOENIX (Blokland et al."
+ 2007a)., 2007a).
+ MIID waves and instabililies control the dynamics of plasma and occur as the nabural response to global excitation., MHD waves and instabilities control the dynamics of plasma and occur as the natural response to global excitation.
+ Measurement of the spectrum of MIID waves gives direct information on the internal state of the plasma., Measurement of the spectrum of MHD waves gives direct information on the internal state of the plasma.
+ This is called MIID spectroscopy in analogy with «quantum mechanical spectroscopy. which also involves eigenvalue problems of linear operators.," This is called MHD spectroscopy in analogy with quantum mechanical spectroscopy, which also involves eigenvalue problems of linear operators."
+ MIID spectroscopy entails a separate study of the nonlinear static equilibrium configuration on the one hand and the various linear wave structures that can occur on the other., MHD spectroscopy entails a separate study of the nonlinear static equilibrium configuration on the one hand and the various linear wave structures that can occur on the other.
+ The first study is the subject of this paper and the spectroscopy will be treatecl im a sequel., The first study is the subject of this paper and the spectroscopy will be treated in a sequel.
+ The paper is organized as follows., The paper is organized as follows.
+ We discuss the morphology of prominences and (heir associated magnetic fields in Section 2.., We discuss the morphology of prominences and their associated magnetic fields in Section \ref{topology}. .
+ The basic magnetohvedrostatie equilibrium problem, The basic magnetohydrostatic equilibrium problem
+Eisenstein&IIu(1999).. where A is the amplitude normalized to σς=Atryy8h'\Ipe). given by observations.,"\citet{eisen}, where $A$ is the amplitude normalized to $\sigma_8=\Delta(r_{\mathrm{M}}=8h^{-1}\mathrm{Mpc})$, given by observations."
+" The cross section for lensing is e(M.z)=πιοr;U(M—My). where ÜCr) is a step function. and Mj, is Che minimum mass of halos above which lenses can produce images with separations greater than .A9."," The cross section for lensing is $\sigma(M,z)=\pi
+y_\mathrm{cr}^2r_0^2\vartheta(M-M_{\mathrm{min}})$, where $\vartheta(x)$ is a step function, and $M_{\mathrm{min}}$ is the minimum mass of halos above which lenses can produce images with separations greater than $\Delta\theta$."
+" The minimum mass A44, can be derived directly [rom the relationship between the mass of a lens halo and the corresponding image separation. as lollows."," The minimum mass $M_{\mathrm{min}}$ can be derived directly from the relationship between the mass of a lens halo and the corresponding image separation, as follows."
+ It is obvious from Figure 1 that the separation between the outer (wo images is almost independent of the source position y., It is obvious from Figure 1 that the separation between the outer two images is almost independent of the source position $y$ .
+ Thus. similar to the analvsis for NEW profiles (Li&Ostriker2002).. the image separation Ar(y) produced by a NTIS halo [or a source al y can be well approximated by Ar(Q)=2:4. where wy is the Einstein radius determined from thelens equation with y=0.," Thus, similar to the analysis for NFW profiles \citep{li02}, the image separation $\Delta x(y)$ produced by a NTIS halo for a source at $y$ can be well approximated by $\Delta x(0)=2x_0$, where $x_0$ is the Einstein radius determined from thelens equation with $y=0$ ."
+ The angular image separation is Then. from this equation aud ecquation(2)) we have where οἱ=2997.9h !Mpe is the present. IDibble radius.," The angular image separation is Then, from this equation and \ref{r0}) ) we have where $c/H_0=2997.9h^{-1}$ Mpc is the present Hubble radius."
+ The magnification bias DCM.z) is caleulated imunericallvy with eiven by Oguriοἱal.(2002).. where ο) is the total magnification of the two outer images for a source al jy: (his can also be computed numerically.," The magnification bias $B(M,z)$ is calculated numerically with given by \citet{oguri02}, where $\mu(y)$ is the total magnification of the two outer images for a source at $y$; this can also be computed numerically."
+ The numerical results of Eq. (5)), The numerical results of Eq. \ref{prob1}) )
+ for NTIS lens halos are plotted in ligure 2 line))., for NTIS lens halos are plotted in figure 2 ).
+ For comparison. the survey results of JVAS/CLASS and the predicted probability for lensing by SIS + NEW and NEW profiles are also shown.," For comparison, the survey results of JVAS/CLASS and the predicted probability for lensing by SIS + NFW and NFW profiles are also shown."
+ A subset οἱ 8958 sources from the combined JVAS/CLASS survey form a well-defined statistical sample containing 13 multiply imaged sources (lens svstems) suitable for analysis of the lens statistics (Myersetal.2003:etal.2003:Patnaik1992:Ixing 1999).," A subset of 8958 sources from the combined JVAS/CLASS survey form a well-defined statistical sample containing 13 multiply imaged sources (lens systems) suitable for analysis of the lens statistics \citep{myers,browne03,patnaik92,king99}."
+. The observed lensing probabilities can be easily calculated (Chen2003b.20048). by P(29AP)=ING9)/8958. where 8) is the number of lenses with separation greater than A9 in 13 lenses.," The observed lensing probabilities can be easily calculated \citep{chenb,chenc} by $P_{\mathrm{obs}}(>\Delta\theta)=N(>\Delta\theta)/8958$, where $N(>\Delta\theta)$ is the number of lenses with separation greater than $\Delta\theta$ in 13 lenses."
+" The observational probability £2,,,(2».2A) is plotted as a histogram in Figure 2.", The observational probability $P_{\mathrm{obs}}(>\Delta\theta)$ is plotted as a histogram in Figure 2.
+ In the two-population model 515 + NEW. the ealaxysize ancl the cluster-size lens halos are approximated by SISand NEW? profiles. respectively (Sarbu.Rusin.&Ala2001:LiOstriker2002:Chen 2004)..," In the two-population model SIS + NFW, the galaxy-size and the cluster-size lens halos are approximated by SISand NFW profiles, respectively \citep{sarbu01,li02,
+chena,chenb,chenc,chend,zhang2004}. ."
+ In the one-populationmodel NEW. lens halos with different sizes are," In the one-populationmodel NFW, lens halos with different sizes are"
+"Data reduction from archival data to Level 1 stage was carried out using the Interactive Processing Environment (HIPE, Ott 2010)) using, however, custom reduction scripts that considerably departed from standard processing for PACS (Poglitschetal.,2010) and, to a lesser extent, for SPIRE (Griffinetal.,2010).","Data reduction from archival data to Level 1 stage was carried out using the Interactive Processing Environment (HIPE, \citealt{ott10}) ) using, however, custom reduction scripts that considerably departed from standard processing for PACS \citep{poglitsch10} and, to a lesser extent, for SPIRE \citep{griffin10}."
+. Level 1 Time Ordered Data (TODs) were exported from HIPE into FITS files., Level 1 Time Ordered Data (TODs) were exported from HIPE into FITS files.
+ Further processing including the map generation was carried out using dedicated IDL and FORTRAN codes., Further processing including the map generation was carried out using dedicated IDL and FORTRAN codes.
+ Saturation conditions were reached for all detectors only in SPIRE aand iimages in correspondence with the 3 brightest peaks in the 1=30° ffield., Saturation conditions were reached for all detectors only in SPIRE and images in correspondence with the 3 brightest peaks in the $l$ field.
+" The prescribed flux correction factorsfor PACS (Poglitschetal.,2010) and SPIRE (Swinyardetal.,2010) were applied to the maps since their photometric calibration was carried out using the default calibration tree in HIPE.", The prescribed flux correction factorsfor PACS \citep{poglitsch10} and SPIRE \citep{swinyard10} were applied to the maps since their photometric calibration was carried out using the default calibration tree in HIPE.
+" A detailed description of the entire data processing chain, including the presentation of the maps obtained in the five bands for the two observed fields, can be found in Traficanteetal. (2010)."," A detailed description of the entire data processing chain, including the presentation of the maps obtained in the five bands for the two observed fields, can be found in \citet{trafi10}."
+. In the present letter we present in Figs., In the present letter we present in Figs.
+" 1 and 2 the three-color images obtained using the 70, 160, and ddata (/=30° aand /=59°,, respectively)."," \ref{l30_color} and \ref{l59_color} the three-color images obtained using the 70, 160, and data $l$ and $l$, respectively)."
+ These amazing maps convey the immediate impression of extended filamentary structures dominating the emission on all spatial scales., These amazing maps convey the immediate impression of extended filamentary structures dominating the emission on all spatial scales.
+" Measurements of the standard deviation of the signal at all wavelengths in the lowest brightness regions of the [=59° ffield yield average values a factor two higher than the sensitivity predictions for point source sensitivity from the HSpot time estimator for all bands except at wwhere the predicted limit is effectively reached, confirming that the noise in our maps is dominated by the cirrus confusion at all wavelengths."," Measurements of the standard deviation of the signal at all wavelengths in the lowest brightness regions of the $l$ field yield average values a factor two higher than the sensitivity predictions for point source sensitivity from the HSpot time estimator for all bands except at where the predicted limit is effectively reached, confirming that the noise in our maps is dominated by the cirrus confusion at all wavelengths."
+ A more detailed quantitative analysis is presented by Martinetal.(2010)., A more detailed quantitative analysis is presented by \citet{martin10}.
+". Herschel’’s ability to observe such large areas with unprecedented wavelength coverage (and extraordinary signal dynamical range allows us to image simultaneously progenitors clouds for massive protoclusters to entire clusters of Young Stellar Objects (YSOs) in acitve star forming regions, while also measuring the effect of their strong stellar winds and powerful outflows on the surrounding medium."," 's ability to observe such large areas with unprecedented wavelength coverage and extraordinary signal dynamical range allows us to image simultaneously progenitors clouds for massive protoclusters to entire clusters of Young Stellar Objects (YSOs) in acitve star forming regions, while also measuring the effect of their strong stellar winds and powerful outflows on the surrounding medium."
+ Infrared Dark Clouds (IRDCs) have received considerable attention in recent years (e.g. Rathborneetal.2006 and Peretto&Fuller 2009)) as potential sites for precursors of cluster forming sites., Infrared Dark Clouds (IRDCs) have received considerable attention in recent years (e.g. \citealt{rat06} and \citealt{peretto09}) ) as potential sites for precursors of cluster forming sites.
+" Found in silhouette against the bright mid-IR background, they shine in emission withH"," Found in silhouette against the bright mid-IR background, they shine in emission with."
+"erschel.. Perettoetal.(2010) shows how temperature effectively decreases from ambient values (20-30K) down to T=8-15K toward the densest ( 10*3cm7*) peaks of these objects, resolving further temperature substructures that can be proxies for subsequent fragmentation."," \citet{peretto10} shows how temperature effectively decreases from ambient values (20-30K) down to T=8-15K toward the densest $\sim 10^{23}$ $^{-2}$ ) peaks of these objects, resolving further temperature substructures that can be proxies for subsequent fragmentation."
+" At the other end of the massive star formation timeline, we find W43, visible in the left portion of Fig. 1,,"," At the other end of the massive star formation timeline, we find W43, visible in the left portion of Fig. \ref{l30_color},"
+ as an outstanding case of Galactic mini-starburst., as an outstanding case of Galactic mini-starburst.
+ Detailed SED construction and luminosity estimates allow us to assess the very early evolutionary stage of the most luminous and massive YSOs in the region., Detailed SED construction and luminosity estimates allow us to assess the very early evolutionary stage of the most luminous and massive YSOs in the region.
+" It is remarkable how the same images show a prominent ridge extending southward which encompasses a ""Opc-wide large cavity excavated by the W43 cluster and which possibly triggers further star formation (Ballyetal., 2010)..", It is remarkable how the same images show a prominent ridge extending southward which encompasses a 70pc-wide large cavity excavated by the W43 cluster and which possibly triggers further star formation \citep{bally10}. .
+ Triggered star formation in less extreme environments can also, Triggered star formation in less extreme environments can also
+a good approxinatiou for oxveen which is produced onu short timescales (few uullion vears) by superuovae (SNe) II as opposed to won which is produced ou long iuescales (up to a IIubble time) by SNe Ia (Matteucci and Cregeio 1986).,a good approximation for oxygen which is produced on short timescales (few million years) by supernovae (SNe) II as opposed to iron which is produced on long timescales (up to a Hubble time) by SNe Ia (Matteucci and Greggio 1986).
+ Di order to compare theory aud data we adopt a simple relation between O and Fe reproducing he observed behaviour of these elemients in the solar icielibourliood. the sue as eiven by Pagel (1989): πμ. The continuous liue in Fig.," In order to compare theory and data we adopt a simple relation between O and Fe reproducing the observed behaviour of these elements in the solar neighbourhood, the same as given by Pagel (1989): ] The continuous line in Fig."
+ l is the best function approximating data points according to the theory discussed in the appeucdix. ), 1 is the best function approximating data points according to the theory discussed in the appendix. )
+"= ve | Woe with 060.15my.1: o,0.690: Dypa: ο 222.610 We have used the sophisticated method in the appendix to approximate the data because of the strong dependence of our results on the chosen function f() and in particular ou the value at ®=0. £(0).","= w_1 + w_2 with 0.454; _1=0.690; _2=1.06; w1=4.12; w2=2.61; We have used the sophisticated method in the appendix to approximate the data because of the strong dependence of our results on the chosen function $f(\Phi)$ and in particular on the value at $\Phi=0$, $f(0)$."
+ Iudeed. the point —(0 is outside of the range of the available data and the metallicity distribution extrapolated to 9=0 depends strongly ou the adopted method.," Indeed, the point $\Phi=0$ is outside of the range of the available data and the metallicity distribution extrapolated to $\Phi=0$ depends strongly on the adopted method."
+ Since the aim of this paper is to recover the Listory of the IME from the data in Fie., Since the aim of this paper is to recover the history of the IMF from the data in Fig.
+" l. we need a method that does not iutroduce anyprior assumption on the specific functional form of the relation (AN,ΑΦ),"," 1, we need a method that does not introduce any assumption on the specific functional form of the relation $ln(\Delta N_*/ \Delta \Phi)$."
+ Iu other words our approach nuust © iudependenut of any physical hypothesis on the star formation law at the basis of the observed data. aud this is just the essence of the method discussed iu the appendix.," In other words our approach must be independent of any physical hypothesis on the star formation law at the basis of the observed data, and this is just the essence of the method discussed in the appendix."
+ The basic asstunptions of our models are the same of the Suuple Model (Tinsley 1980) with the only exception of adopting a time-dependent IME., The basic assumptions of our models are the same of the Simple Model (Tinsley 1980) with the only exception of adopting a time-dependent IMF.
+ Therefore we consider a model with the following properties: Underthe previous asstuuptions the oxygen abundance P. (defined by eq.(1)) m the interstellar rediuni ds geverned by fracpgZyads where Zy=ns~0.51:107. g is the eas nass. pads=pds is the mass of oxyeeu produced and alinost uuediatelv returned iuto the interstellar medium when ds of interstellar material goes into stars.," Therefore we consider a model with the following properties: Underthe previous assumptions the oxygen abundance $\Phi$ (defined by eq.(1)) in the interstellar medium is governed by ds where $Z_{0}=H\frac{O_{\odot}}{H_{\odot}}\simeq 9.54 \cdot 10^{-3}$, $g$ is the gas mass, $p \alpha ds=p' ds$ is the mass of oxygen produced and almost immediately returned into the interstellar medium when $ds$ of interstellar material goes into stars."
+ À fraction fjσὲ the mass gone into stars is not returned to the iuterstellar wnedimim. but remains im loue-lved stars or stellar remmauts.," A fraction $\alpha$ of the mass gone into stars is not returned to the interstellar medium, but remains in long-lived stars or stellar remnants."
+ It has also been asstuued that OZ)<< I., It has also been assumed that $\Phi Z_{0}<<1$ .
+ We lave: GunAY Ci aud di," We have: (m,t) dm and dm"
+ We lave: GunAY Ci aud din," We have: (m,t) dm and dm"
+observations. we analyzed a 3-vear sample composed with the GFS modeling data ancl satellite observation data from ISCCP.,"observations, we analyzed a 3-year sample composed with the GFS modeling data and satellite observation data from ISCCP."
+ The results are summarized as follow: In all. we observed a good overall consistency between the GES model forecast and ISCCP observation throughout the time period of interest.," The results are summarized as follow: In all, we observed a good overall consistency between the GFS model forecast and ISCCP observation throughout the time period of interest."
+ For total cloud cover. our result suggested a salisfactorv performance of (he GES model for (he need of observation scheduling up to a week ahead.," For total cloud cover, our result suggested a satisfactory performance of the GFS model for the need of observation scheduling up to a week ahead."
+ However. lor laver and convective cloud. which can be considerably important for observatories located in certain environment (for example. high cloud for high altitude observatories). (he success rate of the GFS model is relatively less satisfving than that of (otal cloud forecast.," However, for layer and convective cloud, which can be considerably important for observatories located in certain environment (for example, high cloud for high altitude observatories), the success rate of the GFS model is relatively less satisfying than that of total cloud forecast."
+of high entropy material injected in the cluster core by jets of relativistic particles.,of high entropy material injected in the cluster core by jets of relativistic particles.
+ Detailed models of bubbles (2???) and jets (777) have been proposed in the context of cluster cores heating. usually based on spherically or planar symmetric. idealized cluster configurations.," Detailed models of bubbles \citep{Churazov:2001p2079, Ruszkowski:2004p2156, Bruggen:2005p2162} and jets \citep{Reynolds:2001p2143, Omma:2004p2154, Cattaneo:2007p420} have been proposed in the context of cluster cores heating, usually based on spherically or planar symmetric, idealized cluster configurations."
+ Dased on simple energetic arguments. it is possible to relate the growth of SMDBIIS at the centre of massive galaxy spheroids to the energy required to unbined the overcooling gas (?)..," Based on simple energetic arguments, it is possible to relate the growth of SMBHs at the centre of massive galaxy spheroids to the energy required to unbind the overcooling gas \citep{Silk:1998p941}."
+ Phe idea of star formation being regulated by AGN feedback at the high mass end of the galaxy mass function has been applied first quite successfully to hvdrodyvnamical simulations of galaxy merger (2). and then to semi-analytical moclels of galaxy formation (?????)..," The idea of star formation being regulated by AGN feedback at the high mass end of the galaxy mass function has been applied first quite successfully to hydrodynamical simulations of galaxy merger \citep{DiMatteo:2005p1020} and then to semi-analytical models of galaxy formation \citep{Croton:2006p5348, Bower:2006p1007, Cattaneo:2006p990, DeLucia:2007p1656, Cattaneo:2009p946}."
+ AGN feedback models have been included only recently in cosmological simulations of galaxy groups ancl clusters (111.," AGN feedback models have been included only recently in cosmological simulations of galaxy groups and clusters \citep{Sijacki:2007p1032, Puchwein:2008p767,McCarthy:2009p747}."
+ Although the detailed physical modeling of SALBLIs erowth usually diller (?2?7).. these simulations. exclusively xwed on the GADGET code (?).. have obtained very encouraging results regarding the global properties of the simulated clusters (????).. 1n.," Although the detailed physical modeling of SMBHs growth usually differ \citep{Sijacki:2007p1032, Booth:2009p501}, these simulations, exclusively based on the GADGET code \citep{Springel:2005p1064}, have obtained very encouraging results regarding the global properties of the simulated clusters \citep{Puchwein:2008p767, McCarthy:2009p747, Puchwein:2010p763, Fabjan:2010p787}."
+ this paper. we report on he first AMI high-resolution simulation of a Virgo-like cluster. following both ΟΛΠ] and star formation. with the associated. feedback.," In this paper, we report on the first AMR high-resolution simulation of a Virgo-like cluster, following both SMBH and star formation, with the associated feedback."
+ A companion paper is exploring the ooperties of a jet.based. GN. feedback. model. (2)...," A companion paper is exploring the properties of a jet–based, AGN feedback model \citep{Dubois:2010p5834}."
+ We use the AMIU code. RAATSES. which dillers significantly rom the other code used so far to model AGN feedback in a fully cosmological simulation. namely the GADGET| code.," We use the AMR code RAMSES, which differs significantly from the other code used so far to model AGN feedback in a fully cosmological simulation, namely the GADGET code."
+ Although AALR sehemes suller from larger advection errors than SPILL. which is a strictly Calilean invariant methocl. there are some celinitive advantages of using AAI in this context: although the total energy (kinetic | thermal | potential) is conserved. only at the percent. level. it is strictly conserving for the Εις energy (kinetic | thermal). which is very important in presenceof strong shocks. and it captures hydrodynamical instabilities more realistically (?2?77).. which is very important in presence of convective motions of buovant gas.," Although AMR schemes suffer from larger advection errors than SPH, which is a strictly Galilean invariant method, there are some definitive advantages of using AMR in this context: although the total energy (kinetic + thermal + potential) is conserved only at the percent level, it is strictly conserving for the fluid energy (kinetic + thermal), which is very important in presenceof strong shocks, and it captures hydrodynamical instabilities more realistically \citep{Agertz:2007p415, Wadsley:2008p1093, Mitchell:2009p4163}, which is very important in presence of convective motions of buoyant gas."
+ Lt also captures gas stripping of infalling satellites due to hydrodynamical instabilities more realistically than standard implementations of SPIL (7).., It also captures gas stripping of infalling satellites due to hydrodynamical instabilities more realistically than standard implementations of SPH \citep{Agertz:2007p415}.
+ We have adapted the SMDII growth model of to the sink particle method for AMI. presented by 2.., We have adapted the SMBH growth model of \cite{Booth:2009p501} to the sink particle method for AMR presented by \cite{Krumholz:2004p1079}.
+ With 7.respect to the previous work of ? ancl follow-up papers. we have mace significant improvements over the OWL simulations suites in term of mass and spatial resolution. but. onlv for one single zoom-in simulation of a Virgolike cluster.," With respect to the previous work of \cite{Booth:2009p501} and follow-up papers, we have made significant improvements over the OWL simulations suites in term of mass and spatial resolution, but only for one single zoom-in simulation of a Virgo–like cluster."
+ Recenth. ? and ? have also used the GADGET code. but a different AGN feedback. model. to perform zoom-in simulations of groups ancl clusters of galaxies. with however a mass resolution and a gravitational softening length. not as good as the one we used in our high resolution case.," Recently, \cite{Puchwein:2008p767} and \cite{Puchwein:2010p763} have also used the GADGET code, but a different AGN feedback model, to perform zoom-in simulations of groups and clusters of galaxies, with however a mass resolution and a gravitational softening length not as good as the one we used in our high resolution case."
+ Note also that the minimum smoothing length in these SPILL simulations is usually much smaller than the gravitational softening length., Note also that the minimum smoothing length in these SPH simulations is usually much smaller than the gravitational softening length.
+ In this paper. we have specifically chosen à Virgo-like cluster. in order to compare our results with the very detailed observations that are available for this well-known astronomical object.," In this paper, we have specifically chosen a Virgo-like cluster, in order to compare our results with the very detailed observations that are available for this well-known astronomical object."
+ We will focus here on the mass distribution of the three main components. namely dark meter. gas and stars.," We will focus here on the mass distribution of the three main components, namely dark matter, gas and stars."
+ The paper is organized as follows: the first. section is dedicated. to. numerical methods ancl physical ingredient (cooling. star formation and AGN feedback). while our second section presents our results. comparing our three models.," The paper is organized as follows: the first section is dedicated to numerical methods and physical ingredient (cooling, star formation and AGN feedback), while our second section presents our results, comparing our three models."
+ The final section is left for cliscussion., The final section is left for discussion.
+ In this section. we describe the numerical techniques and the initial conditions we used to model our Virgolike cluster.," In this section, we describe the numerical techniques and the initial conditions we used to model our Virgo–like cluster."
+ As it is now customary for cosmological simulations. we used a zoom-in (or volume renormalisation) technique to focus our computational resources on a specific region of a LOO Alpesh periodic box.," As it is now customary for cosmological simulations, we used a zoom-in (or volume renormalisation) technique to focus our computational resources on a specific region of a 100 $Mpc/h$ periodic box."
+" We adopt a standard ACDAL cosmology with Q,,Ξ0.3. Q4—0.7. Oy=0.045 and ff=τὸ km/s/Mpce."," We adopt a standard $\Lambda$ CDM cosmology with $\Omega_m=0.3$, $\Omega_\Lambda=0.7$, $\Omega_b=0.045$ and $H_0=70$ km/s/Mpc."
+ We have adopted the ? transfer function and the package (27) in its parallel implementation (7?) to eenerate our initial conditions., We have adopted the \cite{Eisenstein:1998p1104} transfer function and the package \citep{Bertschinger:2001p1123} in its parallel implementation \citep{Prunet:2008p388} to generate our initial conditions.
+ From a first low resolution dark matter only run. we built a catalog of candidate halos whose virial masses lie in the range 1075 to 2«1023 AL. fh.," From a first low resolution dark matter only run, we built a catalog of candidate halos whose virial masses lie in the range $10^{14}$ to $2 \times 10^{14}$ $_\odot$ /h."
+" From this catalog. we have extracted our final halo based on its mass assembly historv:itsfinal mass (AM,210+ AL.) is already. in place around z=1l. so it can be considered as à well relaxed cluster by z0."," From this catalog, we have extracted our final halo based on its mass assembly history:itsfinal mass $M_{vir} \simeq 10^{14}$ $_\odot$ ) is already in place around $z=1$, so it can be considered as a well relaxed cluster by $z=0$."
+" The final halo mass has been measured to be Asus,=1.04107 M. or Alsou.=TSO.107 AL. where indice e refers to the critieal density."," The final halo mass has been measured to be $M_{200c}=1.04 \times 10^{14}$ $_\odot$ or $M_{500c}=7.80 \times 10^{13}$ $_\odot$, where indice $c$ refers to the critical density."
+ We have performed two simulations. one at low resolution. for which the initial grid elfective size was 10247. and one at high resolution. with ellective grid. size of 20487.," We have performed two simulations, one at low resolution, for which the initial grid effective size was $1024^3$, and one at high resolution, with effective grid size of $2048^3$."
+" From this initial grid. we have extracted high resolution particles only in the Lagrangian volume of the halo. and we have used larger ancl larger mass particles to sample the cosmological volume. sothat the total number of particles in the box was 5.210"" (resp."," From this initial grid, we have extracted high resolution particles only in the Lagrangian volume of the halo, and we have used larger and larger mass particles to sample the cosmological volume, sothat the total number of particles in the box was $5.2 \times 10^6$ (resp."
+" 22 10"") for only 2.6«10° (resp.", $22\times 10^6$ ) for only $2.6\times 10^6$ (resp.
+" 19. 10"") in the central region. and 10"" (resp."," $19\times 10^6$ ) in the central region, and $10^6$ (resp."
+" SN«10"") inside the final virial radius for the low resolution (resp.", $8\times 10^6$ ) inside the final virial radius for the low resolution (resp.
+ high resolution) simulation., high resolution) simulation.
+ Ehe dark matter particle mass is therefore 6.5«107 (resp., The dark matter particle mass is therefore $6.5\times 10^7$ (resp.
+ 8.2. 107) AL. /h and the barvons resolution clement mass is 1.2104 (resp., $8.2 \times 10^6$ ) $_\odot$ /h and the baryons resolution element mass is $1.2\times 10^7$ (resp.
+ L4. 107) M. fh., $1.4 \times 10^6$ ) $_\odot$ /h.
+ The AAIR. grid was initially refined to the same level of refinement than the particle eric (1024? ancl 20487). and 7 more levels of refinements were considered.," The AMR grid was initially refined to the same level of refinement than the particle grid $1024^3$ and $2048^3$ ), and 7 more levels of refinements were considered."
+ We imposed the spatial resolution to remain roughly constant in physical units. so that the minimum grid cell staved close to Avinin=Lí2'* with fig —17 (resp.," We imposed the spatial resolution to remain roughly constant in physical units, so that the minimum grid cell stayed close to $\Delta x_{\rm min} = L/2^{\ell_{\rm max}}$ with $\ell_{\rm max}=$ 17 (resp."
+ IS} at z= 0., 18) at $z=0$ .
+" We therefore reached a spatial resolution of zMrgyi,ya1 kpe for the low resolution simulation and Avwinc500 pe for the high resolution one.", We therefore reached a spatial resolution of $\Delta x_{\rm min} \simeq 1$ kpc for the low resolution simulation and $\Delta x_{\rm min} \simeq 500$ pc for the high resolution one.
+ “Phe grid was dynamically refined up to this resolution using a quasi-Lagrangian strategy: when the dark matter or barvons mass in a cell reaches S times the initial mass resolution. it is split into S children cells.," The grid was dynamically refined up to this resolution using a quasi-Lagrangian strategy: when the dark matter or baryons mass in a cell reaches 8 times the initial mass resolution, it is split into 8 children cells."
+ We reached 20 with 14510° (resp., We reached $z=0$ with $14 \times 10^6$ (resp.
+ 68«. 107) cells for the low (resp., $68 \times 10^6$ ) cells for the low (resp.
+ high) resolution run. including split cells.," high) resolution run, including split cells."
+ Gas dynamics is modeled. using a second-order unsplit Godunov scheme (27?) based on the LILLC Riemann solver (?)..," Gas dynamics is modeled using a second-order unsplit Godunov scheme \citep{Teyssier:2002p451, Teyssier:2006p413, Fromang:2006p400} based on the HLLC Riemann solver \citep{Toro:1994p1151}. ."
+ We assume a perfect eas equation of state (EOS) with 5 5/5., We assume a perfect gas equation of state (EOS) with $\gamma=5/3$ .
+ Gas metallicity is advected as à passive scalar. and is sell-consistently accounted forin the cooling function.," Gas metallicity is advected as a passive scalar, and is self-consistently accounted forin the cooling function."
+ We also considered the standard homogeneous UV background of ?.. but we modified the starting redshift. extrapolating the average intensity [rom μαι=6 to γω= 12. in order to," We also considered the standard homogeneous UV background of \cite{Haardt:1996p1167}, , but we modified the starting redshift, extrapolating the average intensity from $z_{\rm reion}=6$ to $z_{\rm reion}=12$ , in order to"
+through the open field lines. therelore their estimate should be an upper limit.,"through the open field lines, therefore their estimate should be an upper limit."
+" Although the exact value of f,« depends on (he model details. in general it should be roughly 1—10."," Although the exact value of $f_{e^{\pm}}$ depends on the model details, in general it should be roughly $\sim 1-10$."
+ The fraction of spin-down power carried away by €7 pairs is given bv Therefore the efficiency of spin-down power Q carried away by pairs is roughly between 0.1 to LI., The fraction of spin-down power carried away by $e^{\pm}$ pairs is given by Therefore the efficiency of spin-down power $\zeta_{e^{\pm}}$ carried away by pairs is roughly between 0.1 to 1.
+ As suggested by Cheng Taam (2003). outergap may not exist for some MSPs with complicated surface magnetic lield in globular clusters. if this is (rue then G+~OL for those MSPs without outergap.," As suggested by Cheng Taam (2003), outergap may not exist for some MSPs with complicated surface magnetic field in globular clusters, if this is true then $\zeta_{e^{\pm}} \sim 0.1$ for those MSPs without outergap."
+ We would like to emphasize again that even outer eap does not exist large number of pairs can still be produced by the polar gap., We would like to emphasize again that even outer gap does not exist large number of pairs can still be produced by the polar gap.
+ In Ruderman Sutherland model. the pair multiplicity is (vpicallv ~10°.," In Ruderman Sutherland model, the pair multiplicity is typically $\sim 10^2$."
+ Therefore in eeneral number of pairs is still much higher than number of protons., Therefore in general number of pairs is still much higher than number of protons.
+ It is interesting to ask how much pulsar wind energy. will lost in (he shock region., It is interesting to ask how much pulsar wind energy will lost in the shock region.
+ In case of the Crab nebula. most of spin-down of pulsar are radiated within the nebula region.," In case of the Crab nebula, most of spin-down of pulsar are radiated within the nebula region."
+ llowever the unpulsed X-ray Iuminositv of pulsars. which is assumed to be emitted [rom the shock region. is only a small fraction of the spin-down power.," However the unpulsed X-ray luminosity of pulsars, which is assumed to be emitted from the shock region, is only a small fraction of the spin-down power."
+ Furthermore Γή. Cheng Taam (2009) have tried to identify the diffuse N-ravs of globular clusters resulting from pulsar wind shock regions. (hey conclude that there is no evidence that the diffuse N-ravs can result. from pulsar wind shock regions.," Furthermore Hui, Cheng Taam (2009) have tried to identify the diffuse X-rays of globular clusters resulting from pulsar wind shock regions, they conclude that there is no evidence that the diffuse X-rays can result from pulsar wind shock regions."
+ They argue that the formation of shock region in globular cluster may be very. difficult because the characteristic shock radius is much larger (han (he characteristic separation of stars due to the very low number density in the elobular clusters., They argue that the formation of shock region in globular cluster may be very difficult because the characteristic shock radius is much larger than the characteristic separation of stars due to the very low number density in the globular clusters.
+ In (his paper we shall assume (hat pairs can be accelerated by absorbing the low Irequenucy electromagnetic wave energy produced by Che dipole radiation of pulsars to relativistic speed., In this paper we shall assume that pairs can be accelerated by absorbing the low frequency electromagnetic wave energy produced by the dipole radiation of pulsars to relativistic speed.
+ If indeed the shock does not exist or very. weak. pairs emitted and accelerated by pulsars can be treated as mono-energetic particles with a Lorentz [actor given bv Eq. (," If indeed the shock does not exist or very weak, pairs emitted and accelerated by pulsars can be treated as mono-energetic particles with a Lorentz factor given by Eq. ("
+3).,3).
+ Since the particle energy loss in the shock is negligible. (he pulsar spin-down," Since the particle energy loss in the shock is negligible, the pulsar spin-down"
+(2009b).. we limit the strength of the Lorentz force in low- regious (2<0.05) to lin-—it the Alfvénn speed to a maximum of 31 kin .,", we limit the strength of the Lorentz force in $\beta$ regions $\beta < 0.05$ ) to limit the Alfvénn speed to a maximum of $31$ km $^{-1}$."
+ This artificial limiting of the Lorentz force maii[uny acts μι the central uilal regions of sunspots in mic-to-hieh photosphere. where local Alfvénn speeds reach up to 103 kins aud would impose prolibitively simall time-steps ou the explicit scheme of the cocle.," This artificial limiting of the Lorentz force mainly acts within the central umbral regions of sunspots in mid-to-high photosphere, where local Alfvénn speeds reach up to $10^3$ km $^{-1}$ and would impose prohibitively small time-steps on the explicit scheme of the code."
+ Recently. MacTageartMD&Hood(2009) mocleled the rise of buoyant toroidal magnetic fux tubes.," Recently, \citet*{MacTaggart:ToroidalFluxTubes} modeled the rise of buoyant toroidal magnetic flux tubes."
+ Their simulatio1 setup is different (rom ours in a nunmiber of ways., Their simulation setup is different from ours in a number of ways.
+ First of all. they used a background convection zone and aticspheric model to mimic the stratification iu tle Sun rut do not iucltde convective flows uor racdiative transfer.," First of all, they used a plane-parallel background convection zone and atmospheric model to mimic the stratification in the Sun but do not include convective flows nor radiative transfer."
+ Secoucdly. the toroidal flux tubes iu their simulations were embedded as buoyant. statiolary structures as an initial condition.," Secondly, the toroidal flux tubes in their simulations were embedded as buoyant, stationary structures as an initial condition."
+ An interesting 'esult [rom their calculations is that the choicee of a toroidal flux tube (as opposed to an originally i0rizontal tube [acilitated the emereence of tje axis of the tube into the photosphere., An interesting result from their calculations is that the choice of a toroidal flux tube (as opposed to an originally horizontal tube) facilitated the emergence of the axis of the tube into the photosphere.
+ attributed this to the ability (iu the toroidal tube case) of plasma to drain down the egs of of the tuye to enhance buoyaut at the ube apex., \citet*{MacTaggart:ToroidalFluxTubes} attributed this to the ability (in the toroidal tube case) of plasma to drain down the legs of of the tube to enhance buoyant at the tube apex.
+ Iu section ??.. we showed that convective lows acting ou he rising magnetic plasma have a siniilar effect.," In section \ref{sec:mass_discharge}, we showed that convective flows acting on the rising magnetic plasma have a similar effect."
+ Before proceeding to discuss the physics underlying the flux emergence aud. AR formation process. we desTibe how the emereine[n] [Iux region evolves in time.," Before proceeding to discuss the physics underlying the flux emergence and AR formation process, we describe how the emerging flux region evolves in time."
+ Magnetic Παν begins to emerge iuto the photosphere beginuiug at /=2 hrs and emergence progresses for several hours., Magnetic flux begins to emerge into the photosphere beginning at $t = 2$ hrs and emergence progresses for several hours.
+ A time sequence of the erey intensity aud maps of tlve vertical field (sampled at 775;= 0.1) are shown iu Fie. L.., A time sequence of the grey intensity and maps of the vertical field (sampled at $\tau_{R}=0.1$ ) are shown in Fig. \ref{fig1}.
+ In the intensity map at /—3 hr. elougated granules begiu to appear.," In the intensity map at $t=3$ hr, elongated granules begin to appear."
+ This stretching of the convective cells is due to the horizontal expansion of the risiug maguetic plasma., This stretching of the convective cells is due to the horizontal expansion of the rising magnetic plasma.
+ A more detailed discussion of the physics of horizoual expausiou is given in section ??.., A more detailed discussion of the physics of horizontal expansion is given in section \ref{sec:subsurface_rise}.
+ Due to the subsurface i0rizontal expansion and undulatioun of the fekl lines by convective flows. the emerging flux 'elon covers an extend surface area within which small-scale bipoles emerge with a systematic orientation (seealsoCheungetal.2008)..," Due to the subsurface horizontal expansion and undulation of the field lines by convective flows, the emerging flux region covers an extend surface area within which small-scale bipoles emerge with a systematic orientation \citep[see also][]{Cheung:SolarSurfaceEmergingFluxRegions}."
+ By /=6.7 lu. small pores begin to appear in inultiple ocations in the intensity map (c.fSteinetal.2010)...," By $t=6.7$ hr, small pores begin to appear in multiple locations in the intensity map \citep[c.f.][]{Stein:SolarFluxEmergenceSimulations}."
+ A corresponding synthetic magnuetoeranm (B. sampled at ro— 0.1) is shown in Fig. 2.., A corresponding synthetic `magnetogram' $B_z$ sampled at $\tau=0.1$ ) is shown in Fig. \ref{fig2}.
+ Two hours later. two spots have appeared wear he horizontal positious coincideut with where the subsurface footpoiuts of the initial serultorus are rooted.," Two hours later, two spots have appeared near the horizontal positions coincident with where the subsurface footpoints of the initial semitorus are rooted."
+ As time progresses. the intensity maps show the spots growing in size.," As time progresses, the intensity maps show the spots growing in size."
+ A synthetic uagnuetograum at a later time shows two large. vertical coucentratious of opposite polarity feld at the locations of the spot (see Fig. 3)).," A synthetic magnetogram at a later time shows two large, vertical concentrations of opposite polarity field at the locations of the spot (see Fig. \ref{fig3}) )."
+" Throughout the simulation. the two spots never attain ully developed pentunbrae which is cousistent with observatious (Zwaau19785.LOST) since the two aree magnetic"" concentrations. have (over/ the course of+ the simulation). at most 1—5-x107!> Mx."," Throughout the simulation, the two spots never attain fully developed penumbrae which is consistent with observations \citep{Zwaan:AppearanceOfMagneticFlux,Zwaan:ElementsAndPatterns} since the two large magnetic concentrations have (over the course of the simulation) at most $4-5\times10^{21}$ Mx."
+ vevertheless. transient. peuumbral filameuts. umbral dots aud even liglit bridges are found in the," Nevertheless, transient penumbral filaments, umbral dots and even light bridges are found in the"
+is in accord with other studies of star formation rates in the LAIC aud the SAIC.,is in accord with other studies of star formation rates in the LMC and the SMC.
+ Alternately. or in addition. the low metallicity of the SAIC aay inhibit wind collimation somehow.," Alternately, or in addition, the low metallicity of the SMC may inhibit wind collimation somehow."
+ We also preseut the measurement of the optical line iuteusities., We also present the measurement of the optical line intensities.
+ We find that the surface brielitucss declines with radii iu most cussion lines. adding value to the possible calibration of the surface briglituess- jiebular radius correlation seeu in the LMC.," We find that the surface brightness declines with radii in most emission lines, adding value to the possible calibration of the surface brightness- nebular radius correlation seen in the LMC."
+ Tonized masses aud electron deusities were calculated where he doublet iuteusity was available., Ionized masses and electron densities were calculated where the doublet intensity was available.
+ The resulting ionized masses do not secu to have a relation to the uorphologv. as already. fouud in the LAC PN.," The resulting ionized masses do not seem to have a relation to the morphology, as already found in the LMC PNs."
+ The ratio of Magellanic Cloud. PNs has been studied in detail., The ratio of Magellanic Cloud PNs has been studied in detail.
+ The factor of two difference tween the LMC aud SAIC PNs were modeled with Cloudy., The factor of two difference between the LMC and SMC PNs were modeled with Cloudy.
+ We fud that the cooling models strongly depend on the oxveeu content. aud that the brightest parts of the Iuninosity function shifts ποποσα SAIC and LAIC PNs.," We find that the cooling models strongly depend on the oxygen content, and that the brightest parts of the luminosity function shifts between SMC and LMC PNs."
+ We also find a relation between the brightest PNs and their norpholoey. that is differeut for cach Alaecllanic Cloud population.," We also find a relation between the brightest PNs and their morphology, that is different for each Magellanic Cloud population."
+ We thank Ceoree Jacoby and Orsola de Marco for discussing the faint target selection. Orsola de Marco for her help iu completing the Phase IT of theST program. and Massimo Stiavelli for his sueecstious about Cloudy models.," We thank George Jacoby and Orsola de Marco for discussing the faint target selection, Orsola de Marco for her help in completing the Phase II of the program, and Massimo Stiavelli for his suggestions about Cloudy models."
+ This work was supported by NASA through erant CO-08663.01-A from Space Telescope Science Iustitute. which is operated by the Association of Universities for Research in Astronomy. Iucorporated. under NASA coutract NAS -26555.," This work was supported by NASA through grant GO-08663.01-A from Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS -26555."
+mobile. iucousistent with euergv conservation.,", inconsistent with energy conservation."
+" This arguirent. although strongly areuing for the previous conclusion. gives uo insight into the local plivsies that prevents the described process {rom occuring. but instead provides au additional motivation fo""tlie analysis given above."," This argument, although strongly arguing for the previous conclusion, gives no insight into the local physics that prevents the described process from occuring, but instead provides an additional motivation for the analysis given above."
+ Let us now cousiler the application of the result to matter inside neutron stars., Let us now consider the application of the result to matter inside neutron stars.
+ Luunecliately after the initial collapse. these stars are hot aud opaque to jeutriuos (Sato(1972):Arnett(1977))). whose non-uniform distribution violates tle codition of full chemical equilibrium.," Immediately after the initial collapse, these stars are hot and opaque to neutrinos \cite{Sato,Arnett}) ), whose non-uniform distribution violates the condition of full chemical equilibrium."
+ Both this ellect aud a siguilicant eutropy gradient act towares cestabiliziug the star (Epstein(1979))). leadiug to vigorous convection (veil.Janka.&Nlueler(1996):Pousetal. (1999))).," Both this effect and a significant entropy gradient act towards destabilizing the star \cite{Epstein}) ), leading to vigorous convection \cite{Keil,Pons}) )."
+ However. ueutron stars older than a few secouds have lost essentially all tir initial entropy. aud uo longer have aly neutrinos trapped tuside (e.g.. Ponsetal.," However, neutron stars older than a few seconds have lost essentially all their initial entropy and no longer have any neutrinos trapped inside (e.g., \cite{Pons}) )."
+ (1999))). ; ime. they have settled into au esseutiallvy zero-eutropy. [ull chemica equilibrium state.," At this time, they have settled into an essentially zero-entropy, full chemical equilibrium state."
+ T in the previous section üunplies that. regardless of the large uncertainties in tlie equi ense matter (2000))). these stars are sably stratified.," The result derived in the previous section implies that, regardless of the large uncertainties in the equation of state for dense matter \cite{EOS}) ), these stars are stably stratified."
+ Maret xinciple possible. but uulikely to hold over exteuced regions.," Marginal stability is in principle possible, but unlikely to hold over extended regions."
+ Furthermore. the s ropy is distributed so as {ο give a nearly uniform temperature (Nomoto&Tsurt also has a stabilizing effect.," Furthermore, the small remaining entropy is distributed so as to give a nearly uniform temperature \cite{Nomoto}) ), which also has a stabilizing effect."
+ The implications of stable statilication for clilferent aspects of neuron star astropliysics have been mentioned in the Introduction., The implications of stable stratification for different aspects of neutron star astrophysics have been mentioned in the Introduction.
+ The writing of this Letter was motivated by a talk given by C. Thompson within the program on at the Institute for Theoretical Physics in Santa Barbara., The writing of this Letter was motivated by a talk given by C. Thompson within the program on at the Institute for Theoretical Physics in Santa Barbara.
+ Parts of it are based on discussions with P. Goldreich which took place several years before., Parts of it are based on discussions with P. Goldreich which took place several years before.
+ Thanks are also due to J. Vélliz for preparing the figure., Thanks are also due to J. Vélliz for preparing the figure.
+ This work was supported by FONDECYT (Chile) under grant no., This work was supported by FONDECYT (Chile) under grant no.
+ 8070009. Complementarias). and by the National Science Foundation (USA) uncer erant no.," 8970009 ), and by the National Science Foundation (USA) under grant no."
+ PHY9LO719L., PHY94-07194.
+objects. which move with linear speeds.,"objects, which move with linear speeds."
+ During the evolution many density maxima merge and lose their identity., During the evolution many density maxima merge and lose their identity.
+ Some clusters obtain higher velocities than the values. expected according to the linear scaling. other clusters move at lower velocities.," Some clusters obtain higher velocities than the values expected according to the linear scaling, other clusters move at lower velocities."
+ However. these effects. compensate each other.," However, these effects compensate each other."
+ Our results show that the rms peculiar velocity of density maxima at z=0 is well approximated by the linear rms velocity expected on a given scale., Our results show that the rms peculiar velocity of density maxima at $z=0$ is well approximated by the linear rms velocity expected on a given scale.
+ The rms peculiar velocity of density maxima is similar to the rms peculiar velocity for all field. points at the same scale., The rms peculiar velocity of density maxima is similar to the rms peculiar velocity for all field points at the same scale.
+ In other words. the rms peculiar. velocity does not depend much on the location of spheres used to determine the velocities.," In other words, the rms peculiar velocity does not depend much on the location of spheres used to determine the velocities."
+ In comparison with the density field. the velocity Ποιά is more heavily weighted bv modes with low values of the wavenumber A. Lec. large scales which are in the linear regine.," In comparison with the density field, the velocity field is more heavily weighted by modes with low values of the wavenumber $k$, i.e. large scales which are in the linear regime."
+ In this paper we have examined. peculiar velocities of clusters. predicted. in the ACDAL model., In this paper we have examined peculiar velocities of clusters predicted in the $\Lambda$ CDM model.
+ We analyzed. the clusters in the Virgo simulation for the ACDAL model with Og—0.3. h—0.7 and ex=0.9.," We analyzed the clusters in the Virgo simulation for the $\Lambda$ CDM model with $\Omega_0=0.3$, $h=0.7$ and $\sigma_8=0.9$."
+ To identify clusters in the simulation we used two methods: (1) the standard. (FOR) method. and (2) the method. where the clusters are defined as the maxima of the smoothed density 1ο] ΜΙΑΝ).," To identify clusters in the simulation we used two methods: (1) the standard (FOF) method and (2) the method, where the clusters are defined as the maxima of the smoothed density field (DMAX)."
+ We used a top-hat window with smoothing radii Ry=1.5%+ Alpe and Ry=10h+ Alpe., We used a top-hat window with smoothing radii $R_s=1.5h^{-1}$ Mpc and $R_s=1.0h^{-1}$ Mpc.
+ The peculiar velocity of DALAN clusters was defined to be the mean »eculiar velocity within a sphere of radius £.., The peculiar velocity of DMAX clusters was defined to be the mean peculiar velocity within a sphere of radius $R_s$.
+ We studied the rms. peculiar. velocity of clusters. for cilferent cluster masses., We studied the rms peculiar velocity of clusters for different cluster masses.
+ We found that the relation between he rms peculiar velocity of a cluster and the mass of the cluster depends on the method that is used to determine the cluster masses and. velocities., We found that the relation between the rms peculiar velocity of a cluster and the mass of the cluster depends on the method that is used to determine the cluster masses and velocities.
+ “Phe rms peculiar velocity of FOR clusters decreases as the mass of the clusters increases., The rms peculiar velocity of FOF clusters decreases as the mass of the clusters increases.
+ The rms peculiar velocity of DALAX clusters is. almost independent of the cluster mass and is well approximated by the linear rms velocity of peaks ofthe density field smoothed at the radius /?=Ry., The rms peculiar velocity of DMAX clusters is almost independent of the cluster mass and is well approximated by the linear rms velocity of peaks of the density field smoothed at the radius $R=R_s$.
+" We also studied the rms. peculiar velocities of peaks. σε. for the linear smoothing radius R=KR,(AL)."," We also studied the rms peculiar velocities of peaks, $\sigma_p$, for the linear smoothing radius $R=R_L(M)$."
+ In this approximation. the rms velocities are significantly lower than the rms velocities of clusters found in the simulation.," In this approximation, the rms velocities are significantly lower than the rms velocities of clusters found in the simulation."
+ We investigated the. distribution. functions of the cluster peculiar. velocities for different cluster masses., We investigated the distribution functions of the cluster peculiar velocities for different cluster masses.
+ The peculiar velocity distributions found for the low-mass FOL clusters have stronger tails than predicted by a Ciaussian distribution., The peculiar velocity distributions found for the low-mass FOF clusters have stronger tails than predicted by a Gaussian distribution.
+ But the velocity cistribution of massive FOL clusters is similar to the Gaussian distribution., But the velocity distribution of massive FOF clusters is similar to the Gaussian distribution.
+ The velocity distribution of DALXX clusters in cülferent mass intervals is similar., The velocity distribution of DMAX clusters in different mass intervals is similar.
+ Phe one-dimensional velocity cistribution of 2.=15h* Alpe clusters is well approximated by a Gaussian., The one-dimensional velocity distribution of $R_s=1.5h^{-1}$ Mpc clusters is well approximated by a Gaussian.
+" La the velocity distribution of R,=10h+ Mpe clusters. we found small deviations from the Gaussian cistribution."," In the velocity distribution of $R_s=1.0h^{-1}$ Mpc clusters, we found small deviations from the Gaussian distribution."
+ We also investigated the evolution of cluster velocities on the scale Re=L5h+ Ape., We also investigated the evolution of cluster velocities on the scale $R_s=1.5h^{-1}$ Mpc.
+ We found a correlation between the peculiar velocities of massive clusters at z=0 and the peculiar velocities of peaks at the initial centres of the clusters at 2=10., We found a correlation between the peculiar velocities of massive clusters at $z=0$ and the peculiar velocities of peaks at the initial centres of the clusters at $z=10$.
+ Density maxima at the scale Ro=15h| Alpe are not isolated objects that move with inear speeds., Density maxima at the scale $R_s=1.5h^{-1}$ Mpc are not isolated objects that move with linear speeds.
+ In this paper we introduced the DALAXN. method., In this paper we introduced the DMAX method.
+ This method. operates with the smoothecl density ancl velocity ields on a given scale Ry., This method operates with the smoothed density and velocity fields on a given scale $R_s$.
+ We expect that the properties of DALAXN clusters (e.g. their mass distribution. correlation function) can be determined analytically from the properties of initial density ancl velocity. fields. at least. for the quasi-inear scales.," We expect that the properties of DMAX clusters (e.g. their mass distribution, correlation function) can be determined analytically from the properties of initial density and velocity fields, at least, for the quasi-linear scales."
+ The density ane velocity distribution function or the maxima in the Gaussian density [old was derived by Bardeen et al. (, The density and velocity distribution function for the maxima in the Gaussian density field was derived by Bardeen et al. (
+1986).,1986).
+ Further study is needed to consider different properties of the DALAX clusters., Further study is needed to consider different properties of the DMAX clusters.
+" We found that the rms velocities of DALAX clusters on a given scale Ry are well described. by the linear theory at AR=£2, (equation 4).", We found that the rms velocities of DMAX clusters on a given scale $R_s$ are well described by the linear theory at $R=R_s$ (equation 4).
+ Suhhonenko Cramann (2003) studied the rms peculiar velocities of clusters in the TCDAL model., Suhhonenko Gramann (2003) studied the rms peculiar velocities of clusters in the $\tau$ CDM model.
+ They found that also in this model. the rms peculiar velocities of DALAXN clusters are close to the linear theory expectations.," They found that also in this model, the rms peculiar velocities of DMAX clusters are close to the linear theory expectations."
+ \We can use the lincar approximation to estimate rms peculiar velocities of clusters in dillerent cosmological models., We can use the linear approximation to estimate rms peculiar velocities of clusters in different cosmological models.
+"2.2109NL.vr+. and optical depths 75,7 between 4.85 and 19.6 for their sample of OH/IR stars.","$2.2\,10^{-6}\,{\rm M_{\odot}yr^{-1}}$, and optical depths $\tau_{9.7}$ between 4.85 and 19.6 for their sample of OH/IR stars."
+ In addition to model A with ει=107. refsed-I4 also displays the SED for a model with the same parameters but with an outer boundary. which is ten times smaller. με=10° (dotted line)," In addition to model A with $r_{out}/r_{1} = 10^{4}$, \\ref{sed-14} also displays the SED for a model with the same parameters but with an outer boundary, which is ten times smaller, $r_{out}/r_{1} = 10^{3}$ (dotted line)."
+ The spectra are virtually indistinguishable up to A ὄθμιι.," The spectra are virtually indistinguishable up to $\lambda \ga 60
+{\rm\mu m}$."
+ The additional cold dust due to the larger outer boundary of model A increases the far infrared fluxes (see insert of refsed-14)). but does not affect the SED at shorter wavelengths JJusttanont Tielens 1992)).," The additional cold dust due to the larger outer boundary of model A increases the far infrared fluxes (see insert of \\ref{sed-14}) ), but does not affect the SED at shorter wavelengths Justtanont Tielens \cite{JusTi92}) )."
+ Since the values of the derived properties (fi. Ma.) are determined essentially by the shape of the SED below GO µια. they do not change for Fontflyxm107.," Since the values of the derived properties $f_{\rm b}$, $\dot{M}_{\rm d}$ , ...) are determined essentially by the shape of the SED below $60\,{\rm\mu
+m}$ , they do not change for $r_{out}/r_{1} \ga 10^{3}$."
+" Figure 8 shows a comparison of the 2.11ju visibility calculated for model A with the azimuthally averaged 2.11ju visibility 1, derived from our speckle observations."," Figure \ref{vis-14} shows a comparison of the $2.11\,{\rm\mu m}$ visibility calculated for model A with the azimuthally averaged $2.11\,{\rm\mu m}$ visibility $V_{\rm obs}$ derived from our speckle observations."
+ Although model A yields a good fit to the observed SED. it fails to reproduce the observed visibility.," Although model A yields a good fit to the observed SED, it fails to reproduce the observed visibility."
+ The slope of the model visibilityis much shallower than the slope of the observed one. and the model visibility levels off at y220arcsec+ with 1544=0.6.," The slope of the model visibilityis much shallower than the slope of the observed one, and the model visibility levels off at $q \ga 20\,{\rm arcsec^{-1}}$ with $V_{2.11}=0.6$."
+" From the observed visibility we obtained an upper limit for the contribution of the star to the monochromatic flux at 2.11jon of = 104. whereas model A yields 62%, indicating that the model optical depth of my,=3 at 2.11yan is too low."," From the observed visibility we obtained an upper limit for the contribution of the star to the monochromatic flux at $2.11\,{\rm\mu m}$ of $\approx 40\,\%$ , whereas model A yields $62\,\%$, indicating that the model optical depth of $\tau_{2.11}=3$ at $2.11\,{\rm\mu m}$ is too low."
+" Furthermore the steeper slope of Την, suggests that a levelling-off should occur at a smaller spatial frequency compared to the model visibility. which means that the observed intensity distribution would be rather more extended than the intensity profile of model A. To summarize. a dust shell model. which yields a visibility in agreement with the observations. requires a larger optical depth at 2.11nn and a more extended intensity distribution compared to model A. but it has to produce the same spectral energy distribution."," Furthermore the steeper slope of $V_{\rm obs}$ suggests that a levelling–off should occur at a smaller spatial frequency compared to the model visibility, which means that the observed intensity distribution would be rather more extended than the intensity profile of model A. To summarize, a dust shell model, which yields a visibility in agreement with the observations, requires a larger optical depth at $2.11\,{\rm\mu m}$ and a more extended intensity distribution compared to model A, but it has to produce the same spectral energy distribution."
+ It turns out that these requirements can be partially fulfilled with suitably modified parameter values., It turns out that these requirements can be partially fulfilled with suitably modified parameter values.
+ In order to find a model which matches both the observed SED and the visibility. as well as exploring the sensitivity of the SED and the 2.11jan visibility of the models on the input parameters. we have calculated various sequences of models.," In order to find a model which matches both the observed SED and the visibility, as well as exploring the sensitivity of the SED and the $2.11\,{\rm\mu m}$ visibility of the models on the input parameters, we have calculated various sequences of models."
+ In each sequence one parameter was varied within a certain range. while the other parameters were kept fixed at the values of model A. except for the optical depth at the reference wavelength 7.55.," In each sequence one parameter was varied within a certain range, while the other parameters were kept fixed at the values of model A, except for the optical depth at the reference wavelength $\tau_{0.55}$."
+ This quantity was adjusted for each model in order to obtain a match of the SED with the observation. especially with the observed strength of the 9.7ju silicate feature.," This quantity was adjusted for each model in order to obtain a match of the SED with the observation, especially with the observed strength of the $9.7\,{\rm\mu m}$ silicate feature."
+ [t turns out that models. for which the SED fit the feature equally well have (almost) equal dust mass loss rates. wwith our procedure we force the models in a sequence to have equal mass loss rates instead of equal 7.55.," It turns out that models, for which the SED fit the feature equally well have (almost) equal dust mass loss rates, with our procedure we force the models in a sequence to have equal mass loss rates instead of equal $\tau_{0.55}$."
+ The effects of different dust temperatures at the inner boundary on the calculated SED and the 2.11jn visibility are displayed in refsed-Ow92-T1..," The effects of different dust temperatures at the inner boundary on the calculated SED and the $2.11\,{\rm\mu m}$ visibility are displayed in \\ref{sed-Ow92-T1}."
+" The dust temperature is varied betwee Tj=6001] and Tj,=1200K. and the derived properties of the corresponding models are given in Table 3.."," The dust temperature is varied between $T_{1}=600\,{\rm K}$ and $T_{1}=1200\,{\rm K}$, and the derived properties of the corresponding models are given in Table \ref{tab-Ow92-T1}."
+ For the values of Tj presented here. the SED is changed at wavelengths below A~9ju.," For the values of $T_{1}$ presented here, the SED is changed at wavelengths below $\lambda \approx 9\,{\rm\mu m}$."
+" Compared to model A with Tj=NOOTN a smaller value of Tj results in a higher flux below A=23,nu and a deficiency of flux at wavelengths betweet 540u and 9ju."," Compared to model A with $T_{1}=800\,{\rm K}$ a smaller value of $T_{1}$ results in a higher flux below $\lambda=3\,{\rm\mu m}$ and a deficiency of flux at wavelengths between $3\,{\rm\mu m}$ and $9\,{\rm\mu m}$."
+ Values of Tj>SOOTS produce a deficiency of flux below A=2n and an excess of flux between 3jim and 9nu.," Values of $T_{1}>800\,{\rm K}$ produce a deficiency of flux below $\lambda=3\,{\rm\mu m}$ and an excess of flux between $3\,{\rm\mu m}$ and $9\,{\rm\mu m}$."
+ As a consequence. the bolometric fluxes are different for these models.," As a consequence, the bolometric fluxes are different for these models."
+" The changes of the SED for a variation of Ti with fixec other parameters. including 7,55. have been studied already by Ivezié Elitzur (1997))."," The changes of the SED for a variation of $T_{1}$ with fixed other parameters, including $\tau_{0.55}$, have been studied already by Ivezić Elitzur \cite{IE97}) )."
+" For high 75,55>100 decreasing T: lowers the flux at shorter wavelengths. increases the strength of the silicate feature and raises the flux at larger wavelengths. similiar to thechanges producedby increasing 75,55."," For high $\tau_{0.55} > 100$ decreasing $T_{1}$ lowers the flux at shorter wavelengths, increases the strength of the silicate feature and raises the flux at larger wavelengths, similiar to thechanges producedby increasing $\tau_{0.55}$ ."
+ Since we have adjusted 7.5; foreach value of Ti to reproduce the observed strength of the 9.7jn silicate feature. the model SED are only changed at shorter wavelengths in à modified way.," Since we have adjusted $\tau_{0.55}$ foreach value of $T_{1}$ to reproduce the observed strength of the $9.7\,{\rm\mu m}$ silicate feature, the model SED are only changed at shorter wavelengths in a modified way."
+ This behaviour ean be understood from the competition of the, This behaviour can be understood from the competition of the
+reflection spectrum.,reflection spectrum.
+ À combination of strongly blurred reflection and a centrally concentrated dise emissivity is consistent with gravitational light bending., A combination of strongly blurred reflection and a centrally concentrated disc emissivity is consistent with gravitational light bending.
+ This predicts that if the X-ray source is close to the plane of the disc. continuum photons will be bent towards the accretion disc and black hole instead of escaping to infinity.," This predicts that if the X-ray source is close to the plane of the disc, continuum photons will be bent towards the accretion disc and black hole instead of escaping to infinity."
+ As a result. the intrinsic continuum is suppressed and the anisotropic illumination of the dise can produce highly smeared reflection.," As a result, the intrinsic continuum is suppressed and the anisotropic illumination of the disc can produce highly smeared reflection."
+ The strength of the reflection can also vary much more from source-to-source than in a simple Newtonian geometry. because it depends on the degree of light bending and relativistic focusing.," The strength of the reflection can also vary much more from source-to-source than in a simple Newtonian geometry, because it depends on the degree of light bending and relativistic focusing."
+ The line of sight to the dises with high inclinations could be obscured by the molecular torus. but a misalignment of the disc with the torus due to cchaotic accretion (King can result in Seyfert [s with high inclinations.," The line of sight to the discs with high inclinations could be obscured by the molecular torus, but a misalignment of the disc with the torus due to chaotic accretion \citep{ki06} can result in Seyfert 1s with high inclinations."
+ The heavily blurred iron lines are associated with strong Compton humps in the 7 narrow-line observations that are fit with Models SRC or SRI., The heavily blurred iron lines are associated with strong Compton humps in the 7 narrow-line observations that are well-fit with Models SRC or SRI.
+ Indeed in 4/7 of the observations. the reflection hump is stronger than the continuum above 10 keV. While the bandpass of iis limited to 12 keV. iis suited to observing the high energy X-ray spectra above 10 keV. The only two narrow-line objects that have been observed with aare NGC 7469 (Patricketal.2011) and NGC 2110 (Reevesal. 2006).," Indeed in 4/7 of the observations, the reflection hump is stronger than the continuum above 10 keV. While the bandpass of is limited to 12 keV, is suited to observing the high energy X-ray spectra above 10 keV. The only two narrow-line objects that have been observed with are NGC 7469 \citep{pa11} and NGC 2110 \citep{re06}."
+. The hard X-ray fluxes in both are consistent with the predictions of Model SRI using the parameters listed in Table 3.., The hard X-ray fluxes in both are consistent with the predictions of Model SRI using the parameters listed in Table \ref{tab:par}.
+ Part of the difticulty in identifying the blurred dise reflection in some of the narrow-line observations is their short observation length tthe low signal-to-noise ratio., Part of the difficulty in identifying the blurred disc reflection in some of the narrow-line observations is their short observation length the low signal-to-noise ratio.
+ It is possible that longer observations of the same objects may reveal the dise reflection., It is possible that longer observations of the same objects may reveal the disc reflection.
+ To test this. we took the Model SRI best-fit spectra for the 7 observations that are well-fit with this model and created fake spectra with exposure times of [00 Ks.," To test this, we took the Model SRI best-fit spectra for the 7 observations that are well-fit with this model and created fake spectra with exposure times of 100 ks."
+ Fitting the fake spectra with Model ION gave a good fit to Mrk 6. NGC SS06¢1) and NGC 7469(1) at confidence level.," Fitting the fake spectra with Model ION gave a good fit to Mrk 6, NGC 5506(1) and NGC 7469(1) at confidence level."
+ Therefore. in these three observations. blurred dise reflection can be detected without the requirement of a spinning black hole and emissivity slope greater than g=3.," Therefore, in these three observations, blurred disc reflection can be detected without the requirement of a spinning black hole and emissivity slope greater than $q = 3$."
+ However. in the + remaining observations. Model ION does not give a good fit to spectra with exposure times as large as 500 ks.," However, in the 4 remaining observations, Model ION does not give a good fit to spectra with exposure times as large as 500 ks."
+ This suggests that the iron line will not be resolved in some cases even if the S/N is extremely good unless extreme blurring effects are included in models., This suggests that the iron line will not be resolved in some cases even if the S/N is extremely good unless extreme blurring effects are included in models.
+ Incorporating our results with those of NO7. a relativistically or non-relativistically broadened reflection component is now detected in 34/37 observations of the total sample.," Incorporating our results with those of N07, a relativistically or non-relativistically broadened reflection component is now detected in 34/37 observations of the total sample."
+ The three narrow observations that do not show evidence for a broad iron line component with the ionised and/or blurred iron line models are: Mrk 590. IC 4329A(1) and NGC SS48¢1).," The three narrow observations that do not show evidence for a broad iron line component with the ionised and/or blurred iron line models are: Mrk 590, IC 4329A(1) and NGC 5548(1)."
+ A high ionisation parameter of£=4437 erg em s.! with Model SRI could explain why no dise reflection is detected in NGC S548¢1)., A high ionisation parameter of $\xi = 4437$ erg cm $^{-1}$ with Model SRI could explain why no disc reflection is detected in NGC 5548(1).
+ Mrk 590 and IC 4329A¢1) are similar to Mrk 6 and NGC 5506(l) in that they are best-tit with a small reflection strength and an emissivity index close to gq=3., Mrk 590 and IC 4329A(1) are similar to Mrk 6 and NGC 5506(1) in that they are best-fit with a small reflection strength and an emissivity index close to $q=3$.
+ However. while Mrk 590 has a moderate disc inclination. our view of the dise in IC 4329A(1) is edge-on in Model SRI.," However, while Mrk 590 has a moderate disc inclination, our view of the disc in IC 4329A(1) is edge-on in Model SRI."
+ In this case. if it is true that the dise inclination is high but the relativistic effects are not strong. the reflection and iron line would be expected to be faint and extremely difficult to identify in Sspectra.," In this case, if it is true that the disc inclination is high but the relativistic effects are not strong, the reflection and iron line would be expected to be faint and extremely difficult to identify in spectra."
+ Overall. our results provide a solution to the long-standing woblem of why some sources do not exhibit broad iron lines.," Overall, our results provide a solution to the long-standing problem of why some sources do not exhibit broad iron lines."
+ Such objects have posed a difficulty for models of AGN as a relativistically broadened iron line and reflection component should be a natural outcome of matter accreting onto a supermassive black ywole via an accretion dise that is illuminated by hard X-rays., Such objects have posed a difficulty for models of AGN as a relativistically broadened iron line and reflection component should be a natural outcome of matter accreting onto a supermassive black hole via an accretion disc that is illuminated by hard X-rays.
+" We iive tested if the reflection apparently appears to be ""missing"" because the accretion dise is ionised and/or the reflection is heavily smeared."," We have tested if the reflection apparently appears to be “missing"" because the accretion disc is ionised and/or the reflection is heavily smeared."
+ The latter effect can be implemented by both a steep emissivity profile and a spinning black hole., The latter effect can be implemented by both a steep emissivity profile and a spinning black hole.
+ The combination of hese two factors can account for the missing blurred reflection in δ/ΓΕ sources although all 11 observations can allow for this type of reflection., The combination of these two factors can account for the missing blurred reflection in 8/11 sources although all 11 observations can allow for this type of reflection.
+" We tind that a steep emissivity protile is essential in smearing the reflection to such an extent that it is difficult to separate from the continuum,", We find that a steep emissivity profile is essential in smearing the reflection to such an extent that it is difficult to separate from the continuum.
+ Furthermore. the stronger Doppler effects associated with a high dise inclination around 60 enhances the blending of the line into the continuum.," Furthermore, the stronger Doppler effects associated with a high disc inclination around $60\de$ enhances the blending of the line into the continuum."
+ However. we are still able to identify this spectral component as the observations generally have reflection strengths that are larger than expected in a geometry where the dise is illuminated from above by a point source.," However, we are still able to identify this spectral component as the observations generally have reflection strengths that are larger than expected in a geometry where the disc is illuminated from above by a point source."
+ The strong and heavily smeared reflection is consistent with the expectation of the gravitational light bending models., The strong and heavily smeared reflection is consistent with the expectation of the gravitational light bending models.
+ This paper is based on observations taken with the XMM-Newton Satellite. an ESA science mission with contributions by ESA Member States and USA.," This paper is based on observations taken with the XMM-Newton Satellite, an ESA science mission with contributions by ESA Member States and USA."
+ KN thanks the Royal society for support and SB acknowledges funding from a Science and Technology Facilities Council (STFC) studentship., KN thanks the Royal society for support and SB acknowledges funding from a Science and Technology Facilities Council (STFC) studentship.
+formation region without imposing unphysical boundary conditions (Tenorio-Tagle&Munoz-Tunón1997:Strickland&Stevens 2000).,"formation region without imposing unphysical boundary conditions \citep{tenorio-tagle97,ss2000}."
+. Lligher spatial resolution X-ray or optical narrow-band imaging of NGC 6810. would be advantageous in. showing whether the extra-planar filament and/or soft. X-ray emission physically joins up with the disk at ον6.5 kpe. or whether this is an optical illusion caused by the projection of a hourglass-shaped nuclear outflow onto the background disk of the galaxy.," Higher spatial resolution X-ray or optical narrow-band imaging of NGC 6810 would be advantageous in showing whether the extra-planar filament and/or soft X-ray emission physically joins up with the disk at $r\sim 6.5$ kpc, or whether this is an optical illusion caused by the projection of a hourglass-shaped nuclear outflow onto the background disk of the galaxy."
+ or CO mapping is required to confirm the presence of a bar. and. would also better constrain the relative gcometry of the hotter ionized. gas.," or CO mapping is required to confirm the presence of a bar, and would also better constrain the relative geometry of the hotter ionized gas."
+ An observation of NGC 6810 with reveals the oesence of extended soft N-rav emission within and beyond he optical cisk of the galaxy., An observation of NGC 6810 with reveals the presence of extended soft X-ray emission within and beyond the optical disk of the galaxy.
+ The emission seen above and xdow the plane of the galaxy might be emission. seen in oojection from large galactic radius (r~30 kpe). but is more plausibly extra-planar emission extending to heights of z~7 kpe from the plane of the galaxy.," The emission seen above and below the plane of the galaxy might be emission seen in projection from large galactic radius $r \sim 30$ kpc), but is more plausibly extra-planar emission extending to heights of $z \sim 7$ kpc from the plane of the galaxy."
+ This. along with eviously known filamentation and peculiar minor-axis ionized gas kinematics. strongly suggest that NGC 6810 is lost toa galactic-scale superwind.," This, along with previously known filamentation and peculiar minor-axis ionized gas kinematics, strongly suggest that NGC 6810 is host to a galactic-scale superwind."
+ The soft X-ray emission seen in projection within the optical disk of the galaxy is closely. associated: with orming regions visible in erounc-basec ancl Optical Monitor. near-UV imaging., The soft X-ray emission seen in projection within the optical disk of the galaxy is closely associated with star-forming regions visible in ground-based and Optical Monitor near-UV imaging.
+ The actively star forming regions. and apparently also the base of the ivpothesized outflow. extend out to a radius of ~6.5 kpc rom the nucleus. and are associated with either a stellar bar or à partially obscurecl star-forming ring.," The actively star forming regions, and apparently also the base of the hypothesized outflow, extend out to a radius of $\sim 6.5$ kpc from the nucleus, and are associated with either a stellar bar or a partially obscured star-forming ring."
+ The apparently extra-planar soft. N-ray. emission is spatially associated with emission. in particular with a prominent. 5-kpc-Iong filament on the north-west of the clisk.," The apparently extra-planar soft X-ray emission is spatially associated with emission, in particular with a prominent 5-kpc-long filament on the north-west of the disk."
+ The soft. diffuse X-ray. emission is presumably thermal xemsstrablung and line emission from a collisionally-ionized asma. based on combined EPIC PN ancl MOS spectral ing. and has an intrinsic luminosity of Ly~(1.5 22).10Meres in the £=0.3 S.0 keV energy. band.," The soft diffuse X-ray emission is presumably thermal bremsstrahlung and line emission from a collisionally-ionized plasma, based on combined EPIC PN and MOS spectral fitting, and has an intrinsic luminosity of $L_{\rm X} \sim (1.5$ $2.2)\times 10^{40}\ergps$ in the $E=0.3$ – 8.0 keV energy band."
+ The characteristic temperature &T0.6 keV. o-element-o-iron abundance ratio Zí/Zye22.../Zpe.. close soft spatial association ancl X-ray to far-HH& llux ratio ἐννα~10.7 are all consistent with local galaxies iosting supernova-driven winds.," The characteristic temperature $kT \sim 0.6$ keV, $\alpha$ -element-to-iron abundance ratio $Z_{\alpha}/Z_{\rm Fe} \sim 2
+Z_{\alpha, \odot}/Z_{\rm Fe, \odot}$, close soft spatial association and X-ray to far-IR flux ratio $f_{\rm X}/f_{\rm FIR} \sim 10^{-3.6}$ are all consistent with local galaxies hosting supernova-driven winds."
+" The apparently large base radius of the outflow: (~ kpe) makes NGC 0510 one of the few ""disk-wicde"" superwinds currently known. as most local superwinds appear to arise in S1 kpe-seale nuclear starburst regions."," The apparently large base radius of the outflow $\sim 6.5$ kpc) makes NGC 6810 one of the few “disk-wide” superwinds currently known, as most local superwinds appear to arise in $\la 1$ kpc-scale nuclear starburst regions."
+ lard X-ray emission from NGC 6810 is weak., Hard X-ray emission from NGC 6810 is weak.
+" Phe total E-2 Oke luminosity is Lx~(442).107""Cres+ ancl the spectrum best fit with a mocderatelv-absorbed E~1.9 power law.", The total $E=2$ – 10 keV luminosity is $L_{\rm X} \sim (4\pm{2}) \times 10^{39} \ergps$ and the spectrum best fit with a moderately-absorbed $\Gamma \sim 1.9$ power law.
+ Based on the stellar mass and star formation rate ol NGC 6810 the observed hard X-ray emission can be fully accounted for with expected level of N-rav binary emission [rom from the old and voung stellar population., Based on the stellar mass and star formation rate of NGC 6810 the observed hard X-ray emission can be fully accounted for with expected level of X-ray binary emission from from the old and young stellar population.
+ The X-ray observations provide no evidence of any AGN activity Consistent with the previous classification of this galaxy as a Seyfert 2., The X-ray observations provide no evidence of any AGN activity consistent with the previous classification of this galaxy as a Seyfert 2.
+ We find that the X-ray. optical. I1 and radio properties of NGC 0510 are all consistent with a starburst galaxy.," We find that the X-ray, optical, IR and radio properties of NGC 6810 are all consistent with a starburst galaxy."
+ Vhe old classification of this galaxy as a Sevfert 2 galaxy is probably incorrect. ancl can be explained based on the mis-identification of an unrelated HZAO-1 X-pav source with NGC 6810 and the broacer-than-normal optical emission lines caused by a supernova-driven outflow.," The old classification of this galaxy as a Seyfert 2 galaxy is probably incorrect, and can be explained based on the mis-identification of an unrelated HEAO-1 X-ray source with NGC 6810 and the broader-than-normal optical emission lines caused by a supernova-driven outflow."
+ We conclude that NGC 6810. is an early-type spiral of roughly equivalent mass to the Milkv. Was. hosting a superwind with an intriguinglv large base radius.," We conclude that NGC 6810 is an early-type spiral of roughly equivalent mass to the Milky Way, hosting a superwind with an intriguingly large base radius."
+ lt is a pleasure to thank the referee. Michael. Dahlem. for the constructive comments that improved the content of this paper.," It is a pleasure to thank the referee, Michael Dahlem, for the constructive comments that improved the content of this paper."
+ “Phis work was funded by grant. NNGOJCGISICG from by NASA., This work was funded by grant NNG04G181G from by NASA.
+ This research has mace use of the NASA/IPAC Extragalactic Database (NIED) which is operated by the Jet Propulsion Laboratory. California. Institute of Technology. under contract with the National Aeronautics and. Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+We present here an estimate of the angular momentum of cores initially supported by turbulent motions. generalizing the results of MILO5 to arbitrary ni width-size relations.,"We present here an estimate of the angular momentum of cores initially supported by turbulent motions, generalizing the results of ML05 to arbitrary line width-size relations."
+" The analytical treatment of this problem rests on several idealizations about turbulent velocities: (1) alea.that they are isotropic and homogeneous: and (2) that the Cartesian velocity components are neither correlated with each nor (to anv appreciable degree)"" with density Huüctuations."," The analytical treatment of this problem rests on several idealizations about turbulent velocities: (1) that they are isotropic and homogeneous; and (2) that the Cartesian velocity components are neither correlated with each other, nor (to any appreciable degree) with density fluctuations."
+ For MNitv we also assume (3) that the core density. profile is spherically svnunetric and can captured in a single function p(r).ration., For simplicity we also assume (3) that the core density profile is spherically symmetric and can be captured in a single function $\rho(r)$.
+ When evaluating our formulae we assume (4) that the velocity dillerence between two points scales as à power of their s and (5). for the purpose of computing Iluctuations. that the velocity components are Gaussian random: fields.," When evaluating our formulae we assume (4) that the velocity difference between two points scales as a power of their separation, and (5), for the purpose of computing fluctuations, that the velocity components are Gaussian random fields."
+ One might object that condition (1) is inconsistent with the turbulent. support of an inhomogeneous density profile., One might object that condition (1) is inconsistent with the turbulent support of an inhomogeneous density profile.
+" Consider.]. however. that ifIp the core profile⋅ is. a power law p(r)xrn"". then the turbulent line. width. must scale as ai)xri where A velocity field with this scaling is consistent with our conditions (1). (2). and (4) if where angle brackets represent an ensemble average and & is a normalization constant related to the virial parameter a."," Consider, however, that if the core profile is a power law $\rho(r)\propto r^{-k_\rho}$, then the turbulent line width must scale as $\sigma(r)\propto r^\beta$ where A velocity field with this scaling is consistent with our conditions (1), (2), and (4) if where angle brackets represent an ensemble average and $k$ is a normalization constant related to the virial parameter $\alpha$ ."
+" MEOS considered the case 3=1/2 appropriate for giant molecular clouds (Αρ= 1): we generalize their formulae to other values of ?.including the fiducial value 3=1/4 corresponding to dispersion&,=3/2."," ML05 considered the case $\beta=1/2$ appropriate for giant molecular clouds $k_\rho=1$ ); we generalize their formulae to other values of $\beta$, including the fiducial value $\beta=1/4$ corresponding to $k_\rho=3/2$."
+ Our goalis to compute the expectation value and of the core specific angular momentum j. normalized to f£. where σ is the one-dimensional line width of the core.," Our goal is to compute the expectation value and dispersion of the core specific angular momentum $j$, normalized to $R_c \sigma$ where $\sigma$ is the one-dimensional line width of the core."
+ We find that. under our assumptions. and that where AMod={pad?r ancl all integrals are restricted to the region of interest (tvally the core interior).," We find that, under our assumptions, and that where $M=\int \rho d^3 \rvec$ and all integrals are restricted to the region of interest (typically the core interior)."
+" These equations. which woe will below. agree with the formulaein —9MEO»s Appendix but allow picο)6LGu)P to be an arbitrary function of Ir,νο "," These equations, which we will prove below, agree with the formulae in ML05's Appendix but allow $\langle[v_x(\rvec_1)-v_x(\rvec_2)]^2\rangle$ to be an arbitrary function of $|\rvec_1-\rvec_2|$ ."
+"It is important to note that the two are identical up to the [actor 2r,sr» in the integrand."," It is important to note that the two formulae are identical up to the factor $-2
+\rvec_1\cdot\rvec_2$ in the integrand."
+" The negative sign in equation (44)) ensures that £73 is positive. since Gr.(ry)ορ1) takes higher values when |rre] is large. hence when r,Pe is negative."," The negative sign in equation \ref{<j^2>}) ) ensures that $\langle j^2\rangle$ is positive, since $\langle[v_x(\rvec_1)-v_x(\rvec_2)]^2\rangle$ takes higher values when $|\rvec_1-\rvec_2|$ is large, hence when $\rvec_1\cdot\rvec_2$ is negative."
+ To evaluate equations (44)) and (45)) we make use of spherical symmetry and. impose equation (43)) for the velocity correlations., To evaluate equations \ref{<j^2>}) ) and \ref{<sigma^2>}) ) we make use of spherical symmetry and impose equation \ref{dv-dr}) ) for the velocity correlations.
+ Defining (|2vj:re(ο) as the cosine of the angle between ry and re. and q =2ryre/Od|r5).as in MEOS.," Defining $\mu = \rvec_1\cdot\rvec_2/(r_1r_2)$ as the cosine of the angle between $\rvec_1$ and $\rvec_2$ , and where $q = 2 r_1 r_2/(r_1^2+r_2^2)$, as in ML05."
+ Tablex 7 and figure 7 provide values of £732 Ria347)for severaldensity. proliles. both for the internal spectrum given equation (42)). and for ο=1/2. which may better represent the background spectrum.," Table \ref{jOnRSigma} and figure \ref{jForCores} provide values of $\langle j^2\rangle^{1/2}/(R\langle \sigma^2\rangle^{1/2})$ forseveral density profiles, both for the internal spectrum given by equation \ref{beta}) ), and for $\beta=1/2$, which may better represent the background spectrum."
+ Ligure 7. plots the ratio (jL7f(Ria737)as given by equations (46)) and. (47)) for aturbulent core profile. a critical Jonnor-Lbert sphere. and a thinshell.," Figure \ref{jForCores} plots the ratio $\langle j^2\rangle^{1/2}
+/(R\langle \sigma^2\rangle^{1/2})$as given by equations \ref{<j^2>-angular}) ) and \ref{<sigma^2>-angular}) ) for aturbulent core profile, a critical Bonnor-Ebert sphere, and a thinshell."
+ The results are very elose to power laws in «2:, The results are very close to power laws in $\beta$ :
+"Alogg=log(GM/R;)-log(GM/R2,vμιrotReg)9704 and ATay53000K (assuming Tor)~e7(9). 1994).","$\Delta\log g=
+\log\zav{GM/R_\text{p}^2}-\log\zav{GM/R_\text{eq}^2-v_\text{rot}^2/R_\text{eq}}
+\approx-0.4$ and $\Delta T_\text{eff} \approx3000\,\text{K}$ (assuming $T_\text{eff}(\vartheta)\sim g^{1/4}(\vartheta)$, )."
+ The variations of local surface gravity and effective temperature cornected with a fast rotation by itself cannot raise any light variability (assuming fixed axis of rotation) because of their axial symmetry., The variations of local surface gravity and effective temperature connected with a fast rotation by itself cannot raise any light variability (assuming fixed axis of rotation) because of their axial symmetry.
+ However. even axisymmetric surface varlations may modify the light variations caused by inhomogenecmus elemental surface distribution via several effects.," However, even axisymmetric surface variations may modify the light variations caused by inhomogeneous elemental surface distribution via several effects."
+ First. surface layers with different abundances may respond differently to the temperature and surface gravity variations. modifying the flux distribution.," First, surface layers with different abundances may respond differently to the temperature and surface gravity variations, modifying the flux distribution."
+ Second. the area of the surface element is modified through the oblateness of the surface of a rotating star.," Second, the area of the surface element is modified through the oblateness of the surface of a rotating star."
+ Finally. the direction of the beam pointing to the observer relative to the local outward normal and its cross-section are different on spherical and oblate surface.," Finally, the direction of the beam pointing to the observer relative to the local outward normal and its cross-section are different on spherical and oblate surface."
+ To test the influence of these effects. we calculated an additional grid of model atmospheres corresponding to the stellar equator (τη=11500K. logg.y= 3.71) and to the pole (Τρ=14300K. logge;= 4.16) with different abundances of silicon and iron (after Table 2)).," To test the influence of these effects, we calculated an additional grid of model atmospheres corresponding to the stellar equator $T_\text{eff,eq}=11\,500\,\text{K}$, $\log g_\text{eq}=3.71$ ) and to the pole $T_\text{eff,p}=14\,300\,\text{K}$, $\log g_\text{p}=4.16$ ) with different abundances of silicon and iron (after Table \ref{esit}) )."
+ The polar and equatorial effective temperature and gravity were calculated assuming that the pparameters Za. e derived by represent some kind of mean over the stellar surface and that these parameters correspond to some particular surface region.," The polar and equatorial effective temperature and gravity were calculated assuming that the parameters $T_\text{eff}$, $g$ derived by represent some kind of mean over the stellar surface and that these parameters correspond to some particular surface region."
+ In this case Tay)=[e/g+Tr.," In this case $T_\text{eff}(\vartheta)=
+ \hzav{g(\vartheta)/g}^{1/4}T_\text{eff}$."
+" For simplicity we assumed à fixed abundance of helium and chromium here (2j,=-]. &c,= 5.9) because these elements are not the maur sources of the light variability."," For simplicity we assumed a fixed abundance of helium and chromium here $\varepsilon_\text{He}=-1$, $\varepsilon_\text{Cr}=5.9$ ) because these elements are not the main sources of the light variability."
+ The resultinglight curve was derived using Eq. 5..," The resultinglight curve was derived using Eq. \ref{velik},"
+ now interpolating the intensities in Eq., now interpolating the intensities in Eq.
+ 4+ also between the models corresponding to the pole and to the equator., \ref{barint} also between the models corresponding to the pole and to the equator.
+ The inclusion of the gravity darkening only leads to the decrease of predicted amplitude of light variations by up to 0.0] mag.," The inclusion of the gravity darkening only leads to the decrease of predicted amplitude of light variations by up to $0.01\,$ mag."
+" This 1s because for both silicon and iron the values of Am, Eq.", This is because for both silicon and iron the values of $\Delta m_\lambda$ Eq.
+ 5. decrease with decreasing temperature.," \ref{velik}
+ decrease with decreasing temperature."
+ The inclusion of different surface areas caused by the oblateness of the star leads to another small modification of the light curve. which is mostly negligible (the difference is up to one millimagnitude).," The inclusion of different surface areas caused by the oblateness of the star leads to another small modification of the light curve, which is mostly negligible (the difference is up to one millimagnitude)."
+ We did not consider the effect of the difference of beams pointing to the observer (1n an oblate star compared to the spherical one). but given the small influence of other effects. we expect that this is also negligible.," We did not consider the effect of the difference of beams pointing to the observer (in an oblate star compared to the spherical one), but given the small influence of other effects, we expect that this is also negligible."
+ We conclude that the effect of the gravity. darkening modifies the light curves. but it does not seem to be the main reason for the discrepancy between predicted and observed light curves.," We conclude that the effect of the gravity darkening modifies the light curves, but it does not seem to be the main reason for the discrepancy between predicted and observed light curves."
+ But the effect of gravity darkening on the light curves introduces a new possibility to test the theory of gravity darkening in the future., But the effect of gravity darkening on the light curves introduces a new possibility to test the theory of gravity darkening in the future.
+ To this end. the influence of this effect on the abundances derived from the Doppler imaging has also to be accounted for.," To this end, the influence of this effect on the abundances derived from the Doppler imaging has also to be accounted for."
+ iis known to host a large-scale magnetic field. which could. in principle. contribute to the light variability via the rotational modulation of the magnetic field intensity and thus opacity in Zeeman broadened spectral features.," is known to host a large-scale magnetic field, which could, in principle, contribute to the light variability via the rotational modulation of the magnetic field intensity and thus opacity in Zeeman broadened spectral features."
+" Assuming a dipolar field geometry based on phase-resolved longitudinal magnetic field measurements from determined a polar magnetic field B,=3 kG and the angle between the stellar rotation axis and the line-of-sight f=74°.", Assuming a dipolar field geometry based on phase-resolved longitudinal magnetic field measurements from determined a polar magnetic field $B_{\rm p}=3$ kG and the angle between the stellar rotation axis and the line-of-sight $\beta=74\degr$.
+ Regarding magnetic CP stars. this is a fairly moderate magnetic field and. as demonstrated by based on detailed model atmosphere calculations with canomalous Zeeman effect and polarized radiative transfer. such eoc field does not produce significant effects on parameters observed in different photometric systems.," Regarding magnetic CP stars, this is a fairly moderate magnetic field and, as demonstrated by based on detailed model atmosphere calculations with anomalous Zeeman effect and polarized radiative transfer, such a field does not produce significant effects on parameters observed in different photometric systems."
+ In addition. with the inclination angle ;/=30° the surface average magnetic field modulus varies only by A(|B])=300 G during the rotational cycle. which ts far too small to induce any detectable changes i light-curves.," In addition, with the inclination angle $i=30\degr$ the surface average magnetic field modulus varies only by $\Delta\,\langle\,|\boldsymbol{B}|\,\rangle\approx300$ G during the rotational cycle, which is far too small to induce any detectable changes in light-curves."
+ On the other hand. as shown by(2008).. the combined impact of the magnetic field and the realistic chemistry could be more important than the effect of using only individual abundances.," On the other hand, as shown by, the combined impact of the magnetic field and the realistic chemistry could be more important than the effect of using only individual abundances."
+ Thus. the effect of including the magnetic field in the computations of photometric parameters would be more stronger in regions of spots with enhanced abundances compared to the case of a homogeneously stratified atmosphere.," Thus, the effect of including the magnetic field in the computations of photometric parameters would be more stronger in regions of spots with enhanced abundances compared to the case of a homogeneously stratified atmosphere."
+ Because we are mostly iterested in the estimate of the maximum possible amplitude that the magnetic. field could introduce in the light curves relative to the inhomogeneously distributed abundances. we computed several magnetic model atmospheres using the code2004).," Because we are mostly interested in the estimate of the maximum possible amplitude that the magnetic field could introduce in the light curves relative to the inhomogeneously distributed abundances, we computed several magnetic model atmospheres using the code."
+ Anomalous Zeeman splitting and polarised radiative transfer were included as described 1(2006)., Anomalous Zeeman splitting and polarised radiative transfer were included as described in.
+.. We found that the highest impact of the magnetic field for a spot with enhanced Fe. Si. and Cr. located at the pole with |B}=3 KG and equator with |B]=1.5 kG amounts to Aux0.005 mag.," We found that the highest impact of the magnetic field for a spot with enhanced Fe, Si, and Cr, located at the pole with $|\boldsymbol{B}|=3$ kG and equator with $|\boldsymbol{B}|=1.5$ kG amounts to $\Delta u\approx0.005$ mag."
+ This difference subsequently decreases for other Strómmgren filters., This difference subsequently decreases for other Strömmgren filters.
+ Accounting for the surface averaged magnetic field modulus results in AC|B|)=2.2 kG at the phase when the magnetic pole is visible and Ας=1.9 kG at the phase of magnetic equator respectively.," Accounting for the surface averaged magnetic field modulus results in $\Delta\,\langle\,|\boldsymbol{B}|\,\rangle=2.2$ kG at the phase when the magnetic pole is visible and $\Delta\,\langle\,|\boldsymbol{B}|\,\rangle=1.9$ kG at the phase of magnetic equator respectively."
+ The effect 1s ther reduced to Aux0.001 mag., The effect is then reduced to $\Delta u\approx0.001$ mag.
+ All this ts of an order of magnitude less than needed to explain the deviations between observed and predicted non-magnetic amplitudes of the «-parameter shown in Fig. 5.., All this is of an order of magnitude less than needed to explain the deviations between observed and predicted non-magnetic amplitudes of the $u$ -parameter shown in Fig. \ref{cuvir_hvvel}.
+ We therefore conclude that the magnetic field has little or negligible effect on the light-curve appearance anc can be ignored in the present study., We therefore conclude that the magnetic field has little or negligible effect on the light-curve appearance and can be ignored in the present study.
+ In addition to the central convection zone. a hot star may have subsurface tron and helium convective zones 2010)..," In addition to the central convection zone, a hot star may have subsurface iron and helium convective zones ."
+ Our models, Our models
+eeV for the profile and LW=11322eeV for the broad Gaussian emploved in Model A. The results obtained here are consistent with a recent analvsis of Ark 120. by Narcini et al. (,eV for the profile and $EW=113^{+46}_{-52}$ eV for the broad Gaussian employed in Model A. The results obtained here are consistent with a recent analysis of Ark 120 by Nardini et al. (
+2010. submitted) who initially fit the features in the KK region. with a prolile.,"2010, submitted) who initially fit the features in the K region with a profile."
+ In the Narcini et al. (, In the Nardini et al. (
+"2010) analysis. both the inclination of the accretion disk and cniissivity are frozen at typical values of ;=40 and q=3.0 respectively. finding an inner radius of emission of ry,=13Iro.","2010) analysis, both the inclination of the accretion disk and emissivity are frozen at typical values of $i=40^{\circ}$ and $q=3.0$ respectively, finding an inner radius of emission of $r_{\rm in}=13^{+19}_{-7}\,r_{\rm g}$."
+ These values are consistent with those obtained here with \lodel €. finding MP=3k fog=3.0PP and rin=25D ," These values are consistent with those obtained here with Model C, finding $i^{\circ}=34^{+7}_{-5}$, $q=3.0^{+1.6}_{-1.0}$ and $r_{\rm in}=25^{+19}_{-7}\,r_{\rm g}$."
+Also detected is the narrow ionized emission line. again consistent with Narcdini ct al. (," Also detected is the narrow ionized emission line, again consistent with Nardini et al. ("
+2010).,2010).
+ As in Model A. there are no narrow ionized emission lines found. in SWIFT J2127.415654. corresponding to and.," As in Model A, there are no narrow ionized emission lines found in SWIFT J2127.4+5654 corresponding to and."
+FeXXNVLE. This implies that. the excesses modelled. as narrow components in. Model D. may. instead be due to relativistically broacened FelxIxo— emission., This implies that the excesses modelled as narrow components in Model B may instead be due to relativistically broadened ${\alpha}$ emission.
+ When modelled. as a single. profile. all significant excesses in the Fe Ix. band. are removed with the profile centroid rest [rame energv at GAkkeV and the blue-wing peaking at  kkeV in the rest frame., When modelled as a single profile all significant excesses in the Fe K band are removed with the profile centroid rest frame energy at keV and the blue-wing peaking at $\sim$ keV in the rest frame.
+ Consequently the observed excess at kkeV. observed in. Model A and more significantly in Model D is more convincinglv described as a blue-wing., Consequently the observed excess at keV observed in Model A and more significantly in Model B is more convincingly described as a blue-wing.
+ Additionally. the observed. intermediate feature in MCC-02-14-0090 is modelled well by the blue-wing of the profile which occurs at 6.58 kkeV. prompting the removal of the previously emploved narrow Ciaussian.," Additionally, the observed intermediate feature in MCG-02-14-009 is modelled well by the blue-wing of the profile which occurs at $\sim6.58$ keV, prompting the removal of the previously employed narrow Gaussian."
+ Replacing the mocel with a (Brenneman Revnolds 2006) line profile further allows a more physically motivated fit to the data., Replacing the model with a (Brenneman Reynolds 2006) line profile further allows a more physically motivated fit to the data.
+ The model allows the spin parameter to be varied between ο«e0.998 with the ain of determining the extent to which we can rule in or out a non-rotating or maximally spinning black hole., The model allows the spin parameter to be varied between $0<a<0.998$ with the aim of determining the extent to which we can rule in or out a non-rotating or maximally spinning black hole.
+ Similarly o Model €. the outer radius of the disc is fixed at rry. while assuming a cise of uniform emissivity.," Similarly to Model C, the outer radius of the disc is fixed at $_{\rm ms}$, while assuming a disc of uniform emissivity."
+ Phroughout the its using Model D the inner radius of emission is assumed to extend down to the innermost stable circular orbit (rjco)., Throughout the fits using Model D the inner radius of emission is assumed to extend down to the innermost stable circular orbit $_{\rm ISCO}$ ).
+ The line energy is confined to kkeV in the rest rame. being frozen at keV. i£ it reaches its lower limit.," The line energy is confined to keV in the rest frame, being frozen at keV if it reaches its lower limit."
+ Alodel D provides good fits to all six objects. producing he best fitting physically. motivated: model for Mrk 345 and NGC 7469 prior to considering a blurred. reflection component from the inner regions of the aceretion disc (i.e. Alodels I2 E). see Table 7 and Figure 5..," Model D provides good fits to all six objects, producing the best fitting physically motivated model for Mrk 335 and NGC 7469 prior to considering a blurred reflection component from the inner regions of the accretion disc (i.e. Models E F), see Table \ref{tab:ModelD} and Figure \ref{fig:ModelD_eeuf}."
+ The presence of ionized. emission lines is entirely consistent with Mocel C. adding weight to the probability that these lines. are present in the spectra.," The presence of ionized emission lines is entirely consistent with Model C, adding weight to the probability that these lines are present in the spectra."
+ Further to this. the inclinations of the acerction disces to the observer are all comparable to those obtained. in. Model €. as are the equivalent widths of the relativistic line profiles.," Further to this, the inclinations of the accretion discs to the observer are all comparable to those obtained in Model C, as are the equivalent widths of the relativistic line profiles."
+ Given that the spin of the central black hole is a free parameter ancl the assumptions within this model. the emissivity. indicios are expected. to vary. from those obtained. with the profile since the measured emissivity is degenerate to some extent with the inner radius of emission.," Given that the spin of the central black hole is a free parameter and the assumptions within this model, the emissivity indicies are expected to vary from those obtained with the profile since the measured emissivity is degenerate to some extent with the inner radius of emission."
+ Therefore as the spin. parameter varies so does rico. consequently alfecting the measured q.," Therefore as the spin parameter varies so does $_{\rm ISCO}$, consequently affecting the measured $q$."
+" The results obtained. from the spin parameter e for these six objects suggest that a maximally spinning central black hole can be ruled out at the confidence level. however for some objects such as Ark 120 and. ALCCG-02-14-009 the spin is unconstrained ancl only an upper limit can be measured (@«0.94 and a<OLSS. suggesting that emission. does not occur within 2.02r, and 2.45ήν respectively. consistent with the Nardini et al."," The results obtained from the spin parameter $a$ for these six objects suggest that a maximally spinning central black hole can be ruled out at the confidence level, however for some objects such as Ark 120 and MCG-02-14-009 the spin is unconstrained and only an upper limit can be measured $a<0.94$ and $a<0.88$, suggesting that emission does not occur within $r_{\rm g}$ and $r_{\rm g}$ respectively, consistent with the Nardini et al."
+ 2010 analysis of Ark 120)., 2010 analysis of Ark 120).
+ llere for ανα 9 the best fitting. model. gives a measured spin value of @=O44OO), Here for Fairall 9 the best fitting model gives a measured spin value of $a=0.44^{+0.04}_{-0.11}$.
+ A previous spin constraint for ανα 9 by Schmoll et. al. (, A previous spin constraint for Fairall 9 by Schmoll et al. (
+2009). found a—0.65Q2 using a blurred. rellector model(&ERIRCONV.. Brenneman Ievnolds (2006). convolved with REFLIONX)),"2009) found $a=0.65^{+0.05}_{-0.05}$ using a blurred reflector model, Brenneman Reynolds (2006), convolved with )."
+ The measurement here also gives an intermediate value. rowever it is only consistent with their findings when they ignore the spectra below kkeV. to ensure that the soft excess is not the component driving the main part of the it.," The measurement here also gives an intermediate value, however it is only consistent with their findings when they ignore the spectra below keV to ensure that the soft excess is not the component driving the main part of the fit."
+ Schmoll οἱ al. (, Schmoll et al. (
+2009) quote a worse spin Constraint when hese conditions↔ are upheld @=0.5t(1.,2009) quote a worse spin constraint when these conditions are upheld $a=0.5^{+0.1}_{-0.3}$.
+ Within error bars at the level these results are consistent with the value of the spin parameter found. here., Within error bars at the level these results are consistent with the value of the spin parameter found here.
+ However. here we find he emissivity index is constrained 0o q—2.7.(4τ whereas Schmoll ct al.," However, here we find the emissivity index is constrained to $q=2.7^{+0.7}_{-0.4}$ whereas Schmoll et al."
+ only constrain this to qo>4.9., only constrain this to $q>4.9$.
+ Alrk 335 has been noted previously as an object with a broad. relativistic FelIxIx line (Gondoin et al., Mrk 335 has been noted previously as an object with a broad relativistic K line (Gondoin et al.
+ 2002)., 2002).
+ Lore we measure an intermediate spin e=0.70.(35. ruling out maximally spinning and non-rotating black holes only at the confidence level.," Here we measure an intermediate spin $a=0.70^{+0.12}_{-0.01}$, ruling out maximally spinning and non-rotating black holes only at the confidence level."
+ The emissivity index is measured ab a moderate value of g=2.652 in agreement. with previous studies of this object. for example by Longinotti ot al. (," The emissivity index is measured at a moderate value of $q=2.6^{+0.5}_{-0.3}$ in agreement with previous studies of this object, for example by Longinotti et al. ("
+2007).,2007).
+ In. agreement with Model €. the component also features a relatively broad equivalent width EW=1463) with an improvement of Ay?40 (for four aclelitional free parameters) over purely distant emission.," In agreement with Model C, the component also features a relatively broad equivalent width $EW=146^{+39}_{-39}$ with an improvement of $\triangle\chi^{2}\sim40$ (for four additional free parameters) over purely distant emission."
+ In NCC 7469 we measure a broad Fe line featuring a low emissivity index of q21.7ο. an inclination of /223D and EW=91 2eeV. providing good agreement with the previous measurements made within Model €. We obtain a black hole spin of «—0.69My0.00Gyo. ruling. out a Schwarzschild1. black hole with conlidence and. a maximally rotating black hole at the confidence level.," In NGC 7469 we measure a broad Fe line featuring a low emissivity index of $q=1.7^{+0.4}_{-0.6}$, an inclination of $i^{\circ}=23^{+15}_{-7}$ and $EW=91^{+9}_{-8}$ eV, providing good agreement with the previous measurements made within Model C. We obtain a black hole spin of $a=0.69^{+0.09}_{-0.09}$, ruling out a Schwarzschild black hole with confidence and a maximally rotating black hole at the confidence level."
+" SWIRT J2127.415654. also. shows a broad and statistically significant component. M=17san coy, improving the fit by Ayo~37 (similarly to Model. C) for four additional [ree parameters."," SWIFT J2127.4+5654 also shows a broad and statistically significant component, $EW=178^{+82}_{-69}$ eV, improving the fit by $\triangle\chi^{2}\sim37$ (similarly to Model C) for four additional free parameters."
+ The emissivity and the inclination of the accretion disc are typical of the objects in this sample., The emissivity and the inclination of the accretion disc are typical of the objects in this sample.
+ Ehe spin parameter measured here is consistent with that measured. bv. Miniutti et al. (, The spin parameter measured here is consistent with that measured by Miniutti et al. (
+2009). 0.70- compared to 0.6(03 found using a blurred reflection model.,"2009), $0.70^{+0.10}_{-0.14}$ compared to $0.6^{+0.2}_{-0.2}$ found using a blurred reflection model."
+ Lere we reject a maximally rotating central black hole at ereater than confidence and a non-rotating black hole with confidence., Here we reject a maximally rotating central black hole at greater than confidence and a non-rotating black hole with confidence.
+ The spin constraints obtained with Mocel D for Merk 335. NGC 7469 and SWIET 2127.4|5654 can be seen in Figure 7..," The spin constraints obtained with Model D for Mrk 335, NGC 7469 and SWIFT J2127.4+5654 can be seen in Figure \ref{fig:spin_contours}."
+ In contrast to the previous models. Model 1 does. not include a contribution from the Comptonization of soft photons (i.e. the model).," In contrast to the previous models, Model E does not include a contribution from the Comptonization of soft photons (i.e. the model)."
+ Alternatively. the excess seen at soft. energies is modelled: using a blurred. rellector in addition to rellection consistent with distant. emission.," Alternatively, the excess seen at soft energies is modelled using a blurred reflector in addition to reflection consistent with distant emission."
+ The ionized reflection. model is convolved with, The ionized reflection model is convolved with
+direct ionization by primary protons turus out to be signilicaut. the other coutributious being at least oue order of magnitude sinaller.,"direct ionization by primary protons turns out to be significant, the other contributions being at least one order of magnitude smaller."
+" The ionization rate of Πο by primary protous reduces to with £N,/dE, as the CR spectrum at a given kinetic energy Ej.", The ionization rate of $_2$ by primary protons reduces to with $dN_p/dE_p$ as the CR spectrum at a given kinetic energy $E_p$.
+ The upper integration limit is set to Ey= 1 GeV. as the rapidly decreasing cross section makes contributions at higher euergies ueelicrible.," The upper integration limit is set to $E_{\max} =$ 1 GeV, as the rapidly decreasing cross section makes contributions at higher energies negligible."
+" There are two tain [actors allecting the ionizatiou rate: We chose the following approach to account for these uncertainties: we first assume that the spect""unm continues as a power-law clown to MeV energies.", There are two main factors affecting the ionization rate: We chose the following approach to account for these uncertainties: we first assume that the spectrum continues as a power-law down to MeV energies.
+ We theu introduce an artificial and sharp cutoll in the spectrum ata iuinunuim kinetie energy Ey. below which the CR flux is set lo ze'0.," We then introduce an artificial and sharp cutoff in the spectrum ata minimum kinetic energy $E_{\min}$, below which the CR flux is set to zero."
+" We then vary the minimum energy between Egi,=1 MeV ancl 100 MeV. As the CR ioulz:Mon Cross section decreases with energy. increasing the thuresholcl will remove much of the poteltial σα. haviug an even more extreme effect thau flattening the CR spectrum."," We then vary the minimum energy between $E_{\min}=1$ MeV and $100$ MeV. As the CR ionization cross section decreases with energy, increasing the threshold will remove much of the potential signal, having an even more extreme effect than flattening the CR spectrum."
+ luicreasine the tweshold at the same time accounts for consicering dillerent stopping ranges., Increasing the threshold at the same time accounts for considering different stopping ranges.
+ At this stage. due to the lack of knowledge of the lower part of the spectrum. we will calculate the possible rauge of ioiization intensity rather thai one specific value.," At this stage, due to the lack of knowledge of the lower part of the spectrum, we will calculate the possible range of ionization intensity rather than one specific value."
+ Figure 1. shows the ionization rate. uormalized to Galactic average. versus the minimum euerev.," Figure \ref{fig:all} shows the ionization rate, normalized to Galactic average, versus the minimum energy."
+ For comparison. the stopping rauge is shown as a fuuction of Ey.," For comparison, the stopping range is shown as a function of $E_{\min}$ ."
+ For the most optimistic, For the most optimistic
+niass-loss process.,mass-loss process.
+ Iu particular. VY ολίαν one of the RSG with the highest observed mass-loss rate. exhibits an asviuuetrie nebula. with ares probably resulting from multiple ejectious.," In particular, VY CMa, one of the RSG with the highest observed mass-loss rate, exhibits an asymmetric nebula, with arcs probably resulting from multiple ejections."
+ Smith et al. (, Smith et al. (
+2001) argue that these arcs were most likely produced by ejection events localised ou the stellar surface. sugecsting a link with elanut convective cells.,"2001) argue that these arcs were most likely produced by ejection events localised on the stellar surface, suggesting a link with giant convective cells."
+ Strong circtustellar inhomogencitics such as plumes or clumps have also been observed in οἱ Ori shell (Plez Lambert 2002)., Strong circumstellar inhomogeneities such as plumes or clumps have also been observed in $\alpha$ Ori shell (Plez Lambert 2002).
+ We thus examined whether the observed velocity fields preseut any relation with miass-loss rates., We thus examined whether the observed velocity fields present any relation with mass-loss rates.
+ Unfortunately. because of the weak emissiou of eircumstellaz molecular eas (paper I) and the difficulty of observing atomic gas. oulydist wass-loss rates are available for our sample.," Unfortunately, because of the weak emission of circumstellar molecular gas (paper I) and the difficulty of observing atomic gas, only mass-loss rates are available for our sample."
+ The circiunstellar dust-to-gas ratio is not known aud Lay vary within RSC., The circumstellar dust-to-gas ratio is not known and may vary within RSG.
+ We show in Fie., We show in Fig.
+ 13 the relation between the dust mass-loss rate aud ὁο» as well as between he ratio Λερ. ie. the rate of flow. aud Oat3/Come," \ref{massloss} the relation between the dust mass-loss rate and $\delta v_{atm \, 3}$, as well as between the ratio $\dot{M}_{dust}/\rho$, i.e. the rate of flow, and $\delta v_{atm \, 3}/v_{esc}$."
+ A clear tendency toward higher rates witli ieher values of οων appears. again sugeestingOO that convection may play a key role in the mass-loss process operating in RSC.," A clear tendency toward higher rates with higher values of $\delta v_{atm \, 3}$ appears, again suggesting that convection may play a key role in the mass-loss process operating in RSG."
+ The typical photospheric escape velocity for a RSG @vith a mass of 15 M. ane a radius of 500 Rj is 100 km 1., The typical photospheric escape velocity for a RSG (with a mass of 15 $_\odot$ and a radius of 500 $_\odot$ ) is $\sim$ 100 km $^{-1}$.
+ Because of the short mean free path of atomis in the atinosphiere of RSC compared to the atmospheric extent. a simple evaporation. such as the Darwin's mechanisi operating in planetary atimospheres. cannot explain the observed tass-loss rates alouc.," Because of the short mean free path of atoms in the atmosphere of RSG compared to the atmospheric extent, a simple evaporation, such as the Darwin's mechanism operating in planetary atmospheres, cannot explain the observed mass-loss rates alone."
+ We recall here that. as oulv dust niass-loss rates are available for most of the RSC in our siuuple. no definite conclusion cau be drawn as long as total mass-loss rates are not evaluated.," We recall here that, as only dust mass-loss rates are available for most of the RSG in our sample, no definite conclusion can be drawn as long as total mass-loss rates are not evaluated."
+ DIudeed. the dust condensation effiiiencw. aud thus dust-to-gas ratio. may also be iuflueuced. by couvective velocities aud by related shock waves (through local density enliauceimoenuts).," Indeed, the dust condensation efficiency, and thus dust-to-gas ratio, may also be influenced by convective velocities and by related shock waves (through local density enhancements)."
+ Iu order to understand the mupact convection has ou the atimospheric structure. we iuav try to follow a very suuple aremment. based on the propertv of ivdrostatie atmospheres.," In order to understand the impact convection has on the atmospheric structure, we may try to follow a very simple argument, based on the property of hydrostatic atmospheres."
+ Additional details can e found in Custafsson Plez (1992). who discuss the impact of urbuleu pressure on density inversions aud mass loss iu ivpergiaut stars.," Additional details can be found in Gustafsson Plez (1992), who discuss the impact of turbulent pressure on density inversions and mass loss in hypergiant stars."
+ We have shown above that the velocities we measured most probably reflect the convective velocity auplitude., We have shown above that the velocities we measured most probably reflect the convective velocity amplitude.
+" If we suppose these couvective motions are urbuleut. they eive rise to a turbulent pressure of the order of (.5pc2,4."," If we suppose these convective motions are turbulent, they give rise to a turbulent pressure of the order of $0.5 \rho v_{turb}^2$."
+" We asstuned further that ci,4 (nu -u/s) is a constant throughout the atinosphiere (This may rot be true. but still the eradicut of turbulent pressure will be dominate by the deusitv eradicut) aud that the imosphere is isothermal."," We assumed further that $v_{turb}$ (in km/s) is a constant throughout the atmosphere (This may not be true, but still the gradient of turbulent pressure will be dominated by the density gradient) and that the atmosphere is isothermal."
+ Iuclicing this pressure teria in he hydrostatic equilibrium equation leads then to where g is the surface eravitv of the star. aud µ the mean molecular mass.," Including this pressure term in the hydrostatic equilibrium equation leads then to where $g$ is the surface gravity of the star, and $\mu$ the mean molecular mass."
+ For temperatures of 2000 to 3000 Ix. θε dS Casily half the value of g. or lower. for values of ρω~~Das and one teuth or less of g for νι715knis f|.," For temperatures of 2000 to 3000 K, $g_{eff}$ is easily half the value of $g$, or lower, for values of $v_{turb}\approx 5{\rm kms^{-1}}$, and one tenth or less of $g$ for $v_{turb}\approx 15{\rm kms^{-1}}$ ."
+ We may thus expect that the effective eravity of RSGs is much ower in their atimosphere than the value computed from their masses and radi., We may thus expect that the effective gravity of RSGs is much lower in their atmosphere than the value computed from their masses and radii.
+ Custafssou Plez (1992) also computed the ratio of radiative to eravitational acceleration iu hwdrostatie NARC'S models (Ciustafssoun et al., Gustafsson Plez (1992) also computed the ratio of radiative to gravitational acceleration in hydrostatic MARCS models (Gustafsson et al.
+ 1975: Plez et al., 1975; Plez et al.
+ 1992)., 1992).
+ It. turus out that the radiative acceleration ou molecular lines iu the outer atinospheric lavers is of the order of g/100 for static models at Tig=3500Is., It turns out that the radiative acceleration on molecular lines in the outer atmospheric layers is of the order of $g/100$ for static models at $T_{\rm eff}=3500~{\rm K}$.
+ Iu a dyaiunie atmosphere it nav become higher. due to he combined effect of line desaturation by velocity fields aud of cooling of the gas by adiabatic expausion.," In a dynamic atmosphere it may become higher, due to the combined effect of line desaturation by velocity fields and of cooling of the gas by adiabatic expansion."
+ It is possible to estimate the maxinuuu rate of kinetic cherey flux injected by turbulent motions as For the stars iu our sample. using our doeteriiünation of <OCgt3> for Cup. We find typical nuubers between LO! and 107 W. We find that the power of convection as nieasured by our indicators is about equal to the huninosity of the stars (Fie. 11)).," It is possible to estimate the maximum rate of kinetic energy flux injected by turbulent motions as For the stars in our sample, using our determination of $<\delta v_{atm,3}>$ for $v_{turb}$, we find typical numbers between $10^{30.5}$ and $10^{32}$ W. We find that the power of convection as measured by our indicators is about equal to the luminosity of the stars (Fig. \ref{kinetic}) )."
+ The wind power is, The wind power is
+the main-sequence 0.05 mag redder than the TO.,the main-sequence 0.05 mag redder than the TO.
+" The sensitivitv of this new absolute age iudicator (Z4V9"") and its reliability have been investigated iu great detail in BCP.", The sensitivity of this new absolute age indicator $\triangle V^{0.05}$ ) and its reliability have been investigated in great detail in BCP.
+ Iu total. a sample of 11 clusters covering [Fe/H]=2.2...0.5 was defined. whose absolute ages agreed within £1 Gyr.," In total, a sample of 11 clusters covering ${\rm [Fe/H] } = -2.2 \cdots -0.5$ was defined, whose absolute ages agreed within $\pm$ 1 Gyr."
+ While the absolute age of this eroup depended on the set of iochrones used (D'Autoua et 11997: VaudeuBere 1996. private comunumnicatiou o BCP: Stranicro Chicfh 1996. tto. BCP). nembership within the eroup was basically invariant.," While the absolute age of this group depended on the set of isochrones used (D'Antona et 1997; VandenBerg 1996, private communication to BCP; Straniero Chieffi 1996, to BCP), membership within the group was basically invariant."
+ From this. the empirical relation of ACBWyRGB Fe /U|(attheconsnonageo fthegroupo felusters)isobtained( Fig.," From this, the empirical relation of $\triangle (B-V)_{\rm TO}^{\rm RGB}$ $\rm [Fe/H]$ (at the common age of the group of clusters) is obtained (Fig."
+ ΜΜ BCP). whichinturneanbecomparcdtothe predictedone.," 11 of BCP), which in turn can be compared to the predicted one."
+ Le formation," Here, BCP noted that none of the adopted isochrone sets agreed with the empirical relation."
+scmploycdandth asthe neousistencycouldarise fromthethcorctical inodeIst neο, It was also demonstrated that the discrepancy is independent of the colour-transformations employed and thus the inconsistency could arise from the theoretical models themselves.
+ Tn the second step 6 clusters of definitely lower age hau that of the calibrating group have Όσοι. used to determine the change in A(Byee with age at Ooeiveu |Fe/II|.," In the second step 6 clusters of definitely lower age than that of the calibrating group have been used to determine the change in $\triangle
+(B-V)_{\rm TO}^{\rm RGB}$ with age at given $\rm[Fe/H]$."
+" The agesOo of the vouneC» clusters relative o the calibrating ones were obtained by means of the AV!"" inethod.", The ages of the young clusters relative to the calibrating ones were obtained by means of the $\triangle V^{0.05}$ method.
+" Again. theoretical models have to be used or this step: ouly theoretical predictious for the stellar απο]ους, but not for colours. are used."," Again, theoretical models have to be used for this step; only theoretical predictions for the stellar brightness, but not for colours, are used."
+ This finally vielded an empirical expression relating age differences with two observational quautities. namely A(BpRGP aud iuetallicity (Eq.," This finally yielded an empirical expression relating age differences with two observational quantities, namely $\triangle (B-V)_{\rm TO}^{\rm RGB}$ and metallicity (Eq."
+ 3 aud Fie., 3 and Fig.
+ 16 of BCP)., 16 of BCP).
+ This relationship is then applied for deriving relativo ages for GC whose ZAIID level caunot be properly defined (blue or scarcely populated IIB)., This relationship is then applied for deriving relative ages for GC whose ZAHB level cannot be properly defined (blue or scarcely populated HB).
+" For all other clusters in the BCP sample the average of the AV!"" aud the ALBΤΟΝ relative ages is used."," For all other clusters in the BCP sample the average of the $\triangle V^{0.05}$ and the $\triangle (B-V)_{\rm
+TO}^{\rm RGB}$ relative ages is used."
+ Cousidering that the BCP relationship depends itself on theoretical isochrones which were demonstrated to be inconsistent iu at least one respect. it is worthwhile to question the reliability of the BCP-velation. which cau be valid ouly if the theoretical isochrones do not suffer from svsteniatfic errors in the definition of the coeval clusters.," Considering that the BCP relationship depends itself on theoretical isochrones which were demonstrated to be inconsistent in at least one respect, it is worthwhile to question the reliability of the BCP-relation, which can be valid only if the theoretical isochrones do not suffer from systematic errors in the definition of the coeval clusters."
+ Second. since SW97 and SW98 claiu the internal cousisteney of their results. both iu terms of absolute and relative ages. one could ask whether this result depeuds on the smaller sample of clusters selected by SW97 and SWOs.," Second, since SW97 and SW98 claim the internal consistency of their results, both in terms of absolute and relative ages, one could ask whether this result depends on the smaller sample of clusters selected by SW97 and SW98."
+ In the following we willtherefore re-derive the DCP.— relationship using the SW9s isochrones where theoretical Input is necessary., In the following we will therefore re-derive the BCP relationship using the SW98 isochrones where theoretical input is necessary.
+ ⊺↕∐∖↥⋅↸∖↕≯↸∖↥⋅↸∖∐↸⊳↸∖≼∶≼⊲≺↧⋜↧↑⋜↧↴∖↴↸∖↑↕↴∖↴↑∐⋜↧↑∪↕≯↕≧≼⊲↕⋟↖↖↽↕↑∐↑∐↸∖↕↴∖↴⋯⊳↕∐⋅∪↕∐∖↴∖↴↸⊳∪↕↴∖↴↕≼∐∖↥⋅↸∖≼↧∙↽∕∏∐∖⋯⋜↧↕↕⊔⋅↸∖⋜↧↴∖↴∪∐↕≯∪↥⋅∐↕↴∖↴⋯⊳↸⊳↿∐⋅↥⋅↸∖∐⊳↸∖ ⋜∥∐∐↑↕∪∐∪↕∟∖⊽≼∶≼⊲⊓↕⊤↕⋖⋜↧⋯↸∖↑⋜↧↕≓↥⋅↕↸⊳∐∐⋜↧↕∪≼∶≼⊲↕∐↸⊳↕⋯∐∖≼⊔∐↕↴∖↴↑∐↸∖↕∪↖↖↽↸∖↥⋅⋜↧⋝↴∖↴∪↕∏↑↸∖⋜↧∶↴∙⊾↸∖↕≯∪∏∐≼∐≯∪↥⋅↑∐↸∖↕∐≼∐↖⇁↕≼↧∏⋜↧↕≼∶≼⊲↖↖↽↕↑∐ SW97) aud the exclusion of the three disk GC I7 Tuc. AL TL and NOC 6352.," The reference GC data set is that of BCP with the addition of NGC 6171 (a metal-rich halo GC included in SW97) and the exclusion of the three disk GC 47 Tuc, M 71 and NGC 6352."
+ As extensively discussed in Alonso etal. (1997). ACBV)SCDPTO and AVESTO Vicld significantlyiz differeut ages for these clusters. if the typical CC helinn couteut (0.21) is assumed.," As extensively discussed in Alonso et (1997), $\triangle (B-V)_{\rm TO}^{\rm RGB}$ and $\triangle V^{\rm HB}_{\rm
+TO}$ yield significantly different ages for these clusters, if the typical GC helium content $0.24$ ) is assumed."
+ The problem cau be resolved with a nearly solar value (0.27: SW9s) for the disk clusters. but this would add an unwanted additional deeree of freedom) for the senmi-enipirical relationship investigated here.," The problem can be resolved with a nearly solar value $0.27$; SW98) for the disk clusters, but this would add an unwanted additional degree of freedom for the semi-empirical relationship investigated here."
+ Moreover. given the strong depeudence of the borizontal-brzuich Iuninositv on the He abuudauce. the iuclusiou of these clusters would affect the preseut coluparison with the models because it would affect the selection of the sample of coeval clusters.," Moreover, given the strong dependence of the horizontal-branch luminosity on the He abundance, the inclusion of these clusters would affect the present comparison with the models because it would affect the selection of the sample of coeval clusters."
+ The reader should be well aware of this cdiffercuce between BCP and the present paper. because at the metal-rich eud of the relationship AUBWR)Pes|Fe/T BCP fond the largest discrepancies between observations aud models.," The reader should be well aware of this difference between BCP and the present paper, because at the metal-rich end of the relationship $\triangle (B-V)_{\rm TO}^{\rm RGB} vs {\rm [Fe/H]}$ BCP found the largest discrepancies between observations and models."
+" Adopting this data set we have selected a siauple of coeva custers, on the base of the observed AV? values ac their errors; aud the absolute ages as deteriuiued from the SW98 isochvoucs."," 	 Adopting this data set we have selected a sample of coeval clusters, on the base of the observed $\triangle V^{0.05}$ values and their errors, and the absolute ages as determined from the SW98 isochrones."
+ The largest possible sample of cocva clusters is different frou BCP aud it is shown in Fig., The largest possible sample of coeval clusters is different from BCP and it is shown in Fig.
+ (sec also data in Tab. 1)):, \ref{f:1} (see also data in Tab. \ref{t:1}) );
+ individual ages rauge from —LO to 12 Car and overlap within the errors. providing au average age of 10.9 η," individual ages range from 10 to 12 Gyr and overlap within the errors, providing an average age of 10.9 Gyr."
+ The sanuple spaus the metallicity Τάσο -2.0<[PeΠΠ <-1.0. typical of the bulk of halo GC. :xd. coutrarily to BCP. it does not include any of the most ictal poor clusters with «-2.0.," The sample spans the metallicity range $<[Fe/H]<$ -1.0, typical of the bulk of halo GC, and, contrarily to BCP, it does not include any of the most metal poor clusters with $<$ -2.0."
+ This is in agreeiueut with the results by SW97 aud SW9s who fond that the ost metal poor GC are on average older than the other GC., This is in agreement with the results by SW97 and SW98 who found that the most metal poor GC are on average older than the other GC.
+ The result that the coeval clusters are not the same as in BCP is at odds with them conclusion that the ⋅ ⊔↸∖∐∐⋝↸∖↥⋅↴∖↴↕∐↻∪↑∐↸∖↸⊳∪↸∖↖↽⋜↧↕∶↴∙⊾↥⋅⋯∏≻≼↧∪↸∖↴∖↴∐∪↑≼∐∖↻↸∖∐≼↧∪∐↕∐∖, The result that the coeval clusters are not the same as in BCP is at odds with their conclusion that the membership to the coeval group does not depend on the isochrones considered.
+ he SW9s isochrones., The main reason for this occurrence is the lower absolute age found for the individual GC with the SW98 isochrones.
+ As previously described. the averageo age of the coeval GC is 10.9 Cove. while BCP considered au averageo aee of 15 Gyr.," As previously described, the average age of the coeval GC is 10.9 Gyr, while BCP considered an average age of 15 Gyr."
+ As it is evident from Fies., As it is evident from Figs.
+↖↰ 7 aud 8 in BCP. the Izshape of the theoretical relation between AW0.05 and |Fo/TI] depends ou the absolute age of the isochrones.," 7 and 8 in BCP, the shape of the theoretical relation between $\triangle V^{0.05}$ and [Fe/H] depends on the absolute age of the isochrones."
+ For exiniple. at ages around 12 Car. the derivative 82V9.9?(age) at low inetallicitics is quite different from the case of LE or 16 Cir. independently of the set of colour transformations aud bolometric corrections used.," For example, at ages around 12 Gyr, the derivative $\delta \triangle
+V^{0.05}/\delta({\rm age})$ at low metallicities is quite different from the case of 14 or 16 Gyr, independently of the set of colour transformations and bolometric corrections used."
+AU between 0 and 20 AU.,AU between 0 and 20 AU.
+ The gravitational potential of the spherical eloud is obtained by the (tridiagonal matrix inversion technique., The gravitational potential of the spherical cloud is obtained by the tridiagonal matrix inversion technique.
+ The boundary conditions are chosen at both 0 and 20 AU to absorb radial velocity perturbations. (hough we shall see that in the new test cases. the velocities are assumed to be zero. so the hycdrodvunamics boundary conditions are irrelevant for testing the radiative transfer routines.," The boundary conditions are chosen at both 0 and 20 AU to absorb radial velocity perturbations, though we shall see that in the new test cases, the velocities are assumed to be zero, so the hydrodynamics boundary conditions are irrelevant for testing the radiative transfer routines."
+" For (he axisvnunetric tests CE models). the equations are solved on a spherical coordinate ο] with NV,=101 and Ny=22 lor 7/2>00."," For the axisymmetric tests (T models), the equations are solved on a spherical coordinate grid with $N_r = 101$ and $N_{\theta} = 22$ for $\pi/2 \ge \theta > 0$."
+ Sxmmeltry through the midplane is assumed lor π>8 zx/2., Symmetry through the midplane is assumed for $\pi \ge \theta > \pi/2$ .
+ The radial grid (r) is uniformly spaced with Ar=0.16 AU between 4 and 20 AU. the same configuration as is typically assumed in the Boss disk instability models.," The radial grid $r$ ) is uniformly spaced with $\Delta r = 0.16$ AU between 4 and 20 AU, the same configuration as is typically assumed in the Boss disk instability models."
+" The Poisson equation for the disk's exavitational potential is solved by a spherical harmonic expansion (Y5,) including terms up to Aj,=16.", The Poisson equation for the disk's gravitational potential is solved by a spherical harmonic expansion $Y_{lm}$ ) including terms up to $N_{lm} = 16$ .
+ Because ol the axisvanmetrie assumption. terms with m+40 are dropped. reducing the solution to an expansion in Legendre polvnonmüals.," Because of the axisymmetric assumption, terms with $m \ne 0$ are dropped, reducing the solution to an expansion in Legendre polynomials."
+ The boundary conditions are chosen at both 4 and 20 AU to absorb radial velocity perturbations. though as noted above. this is irrelevant [or the radiative transler.," The boundary conditions are chosen at both 4 and 20 AU to absorb radial velocity perturbations, though as noted above, this is irrelevant for the radiative transfer."
+ Dolev οἱ al. (, Boley et al. (
+2006. 2007b) developed test problems based on the approach taken. by ]Iubeny (1990) in deriving simplified analvtical models for the vertical styucture of accretion disks.,"2006, 2007b) developed test problems based on the approach taken by Hubeny (1990) in deriving simplified analytical models for the vertical structure of accretion disks."
+ Hubeny (1990) studied. one dimensional. plane-parallel atinospheres in hwdrostatic equilibrium (i.e.. at rest). wilh grev opacities and Eddington approximation radiative Gransfer. while neglecting convection ancl external radiation.," Hubeny (1990) studied one dimensional, plane-parallel atmospheres in hydrostatic equilibrium (i.e., at rest), with grey opacities and Eddington approximation radiative transfer, while neglecting convection and external radiation."
+ Boley et al. (, Boley et al. (
+2007b) made similar assumptions for testing their flux-limited diffusion approximation munerical models. while also assuming constant verlical gravitv. constant opacity. and constant densitv.,"2007b) made similar assumptions for testing their flux-limited diffusion approximation numerical models, while also assuming constant vertical gravity, constant opacity, and constant density."
+ We will make similar approximations here., We will make similar approximations here.
+ All of the author's disk instability models since Boss (2001) have emploved. radiative transfer in the diffusion approximation. through the solution of the equation determining the evolution of the specific internal energy. 7 where p is the total gas and dust mass density. / is (me. v is the velocity of the gas and cust (considered to be a single fluid). p is the eas pressure. & is the Rosselancl mean opacitv. σ isthe Stefan-Doltzmann constant. and7 is the gas and dust temperature.," All of the author's disk instability models since Boss (2001) have employed radiative transfer in the diffusion approximation, through the solution of the equation determining the evolution of the specific internal energy $E$ where $\rho$ is the total gas and dust mass density, $t$ is time, ${\bf v}$ is the velocity of the gas and dust (considered to be a single fluid), $p$ is the gas pressure, $\kappa$ is the Rosseland mean opacity, $\sigma$ isthe Stefan-Boltzmann constant, and$T$ is the gas and dust temperature."
+ The, The
+in contact. the area subtended by the (then main sequence star was on the order of 0.001 percent of the surface area of the supernova precursor.,"in contact, the area subtended by the then main sequence star was on the order of 0.001 percent of the surface area of the supernova precursor."
+ Even if the supernova ejecta exceeded several solar masses of oxvgen. the amount accreted by the main sequence star would have resulted in undetectable surface abundance changes alter the main sequence surface lias been subduetecl in the giant stage.," Even if the supernova ejecta exceeded several solar masses of oxygen, the amount accreted by the main sequence star would have resulted in undetectable surface abundance changes after the main sequence surface has been subducted in the giant stage."
+ Defininge characteristics of high-massex and low-mass X-ray binaries are exgiven by VerbuntHeuvel(1995)., Defining characteristics of high-mass and low-mass X-ray binaries are given by \citet{vh95}.
+. In low-mass X-ray binaries. the companion to the compact object has amass στ 1.0 M...," In low-mass X-ray binaries, the companion to the compact object has a mass $\lesssim$ 1.0 $_{\sun}$."
+ These svstems belong to a very. old stellar population and do not show runaway characteristics., These systems belong to a very old stellar population and do not show runaway characteristics.
+ Dased solely on our upper limit to the mass of the red eiant. V2116 Oph/GX 1-24 would marginally meet the criteria for a low-mass svstem.," Based solely on our upper limit to the mass of the red giant, V2116 Oph/GX 1+4 would marginally meet the criteria for a low-mass system."
+ However. the abundances of the red giant are nearly solar and hence the svstem does not belong to a very old population.," However, the abundances of the red giant are nearly solar and hence the system does not belong to a very old population."
+ The center-o[-mass velocity of the svstem is quite large. 2177 km +.," The center-of-mass velocity of the system is quite large, $-$ 177 km $^{-1}$."
+ While ihe V2116 Oph svstem fits Cie mass criteria for low-mass svstenis. (e near-solar abundances make V2116 Oph a deviant member.," While the V2116 Oph system fits the mass criteria for low-mass systems, the near-solar abundances make V2116 Oph a deviant member."
+ Furthermore. the large velocity of V2116 Oph suggests that it is a runaway svsten.," Furthermore, the large velocity of V2116 Oph suggests that it is a runaway system."
+ On 2003 February 16 the 1.564 jan spectrum of a star separated [rom V2116 Oph by 19 areseconds. Star 6 in Chakrabarty&Roche(1997) Table 2. was observed wilh a spectral resolving power of 50000 with Phoenix at Gemini South.," On 2003 February 16 the 1.564 $\mu$ m spectrum of a star separated from V2116 Oph by 19 arcseconds, Star 6 in \citet{cr97}
+ Table 2, was observed with a spectral resolving power of 50000 with Phoenix at Gemini South."
+ This star. which has a similar visual magnitude to V2116 Oph. proved to be an M star with a velocity of +118 km s...," This star, which has a similar visual magnitude to V2116 Oph, proved to be an M star with a velocity of $+$ 118 km $^{-1}$."
+ The observation of one star of unknown distance does not demonstrate (he kinematics of the V2116 Oph field. but the 295 km ! velocity difference between the (wo svstenis suggests that a significant velocity was imparted to the V2116 Oph system by the supernova explosion. eiving it a runaway velocity.," The observation of one star of unknown distance does not demonstrate the kinematics of the V2116 Oph field, but the 295 km $^{-1}$ velocity difference between the two systems suggests that a significant velocity was imparted to the V2116 Oph system by the supernova explosion, giving it a runaway velocity."
+ Considerable discussion exists in the literature on wavs to produce X-ray binaries., Considerable discussion exists in the literature on ways to produce X-ray binaries.
+ As reviewed by Canaletal.(1990).. Ibenetal. (1995).. Verbunt&vandenIHeuvel (1995).. and others. while there is consensus that high mass svstems are the result of massive star evolution. four schemes have been advanced for the formation of low-mass X-ray binary svslenms.," As reviewed by \citet{cil90}, \citet{ity95}, \citet{vh95}, and others, while there is consensus that high mass systems are the result of massive star evolution, four schemes have been advanced for the formation of low-mass X-ray binary systems."
+ These four are: (1) capture of a neutron star by another star in a dense stellar svstem (e.g. globular cluster). (2) the direct evolution of one star of a primordial binary into a neutron star. (3) triple star evolution of a svstem including a massive star binary. and (4) a low mass binary where one member first becomes a white dwarf then undergoes mass accretion resulting in collapse to a neutron star.," These four are: (1) capture of a neutron star by another star in a dense stellar system (e.g. globular cluster), (2) the direct evolution of one star of a primordial binary into a neutron star, (3) triple star evolution of a system including a massive star binary, and (4) a low mass binary where one member first becomes a white dwarf then undergoes mass accretion resulting in collapse to a neutron star."
+" For a field star. tidal capture by a neutron star is unlikely,"," For a field star, tidal capture by a neutron star is unlikely."
+ The neutron star must approach the star to be captured by 2 (0 3 (mes (he target star radius. although the odds may be increased by involving a binary svstem (Canaletal.1990).," The neutron star must approach the star to be captured by 2 to 3 times the target star radius, although the odds may be increased by involving a binary system \citep{cil90}."
+. Since the neutron star and (he M giant have nearly equal masses (he capture conditions are very limited., Since the neutron star and the M giant have nearly equal masses the capture conditions are very limited.
+ The tidal, The tidal
+Our decomposition of the MHD turbulence confirmed the results in CLO3 in terms of spectra. namely. that the Alfvénnic and slow modes are anisotropic and consistent with the predictions of the GL95 model of incompressible turbulence. while the fast modes are mostly isotropic and form an acoustic turbulence cascade 2002).,"Our decomposition of the MHD turbulence confirmed the results in CL03 in terms of spectra, namely, that the Alfvénnic and slow modes are anisotropic and consistent with the predictions of the GL95 model of incompressible turbulence, while the fast modes are mostly isotropic and form an acoustic turbulence cascade ."
+ As these results were used in studies of cosmic ray scattering and acceleration as well as charged dust acceleration this is an encouraging development., As these results were used in studies of cosmic ray scattering and acceleration as well as charged dust acceleration this is an encouraging development.
+accelerratior in external galaxies.,ration in external galaxies.
+The EW of the [OJ] A6300 line in the out-of-eclipse spectrum is about 0.17A. which is «quite large lor a WITS.,"The EW of the [OI] $\lambda$ 6300 line in the out-of-eclipse spectrum is about 0.17, which is quite large for a WTTS."
+ Onlv one of the ten WTTSs in Table 3 of has a detectable [OI] AG300 line (with an EW = 0.5 A)., Only one of the ten WTTSs in Table 3 of \citet{Hartigan95} has a detectable [OI] $\lambda$ 6300 line (with an EW = 0.5 ).
+ The others have upper limits of about 0.06A., The others have upper limits of about 0.06.
+ Most. CTTSs in that table. on the other hand. have EWs of 0.5 - 3A.," Most CTTSs in that table, on the other hand, have EWs of 0.5 - 3."
+ The profiles obtained during mic-eclipse and eeress suggest (hat we have (wo enission peaks. ab about -20 and +18 km/s with emission wines extending -60 to +50 km/s. respectively.," The profiles obtained during mid-eclipse and egress suggest that we have two emission peaks, at about -20 and +18 km/s with emission wings extending -60 to +50 km/s, respectively."
+ These profiles are quite different from the [OI] A6300 profile seen in most strong-emission CTTSs (with strong veiling). which often have a hieh-velocity component al -100 to -150 kin/s (resulting trom the jet) aud a low-velocily component at about -20 km/s (Ilartiganal.1997).," These profiles are quite different from the [OI] $\lambda$ 6300 profile seen in most strong-emission CTTSs (with strong veiling), which often have a high-velocity component at -100 to -150 km/s (resulting from the jet) and a low-velocity component at about -20 km/s \citep{Hartigan95, Hirth97}."
+. In addition. the prolile is quite different Irom that of a CTTS with small veiling aud UR excess (see (the bottom profile in Figure 11 of (1995))).," In addition, the profile is quite different from that of a CTTS with small veiling and IR excess (see the bottom profile in Figure 11 of \citet{Hartigan95}) )."
+ These latter profiles are usually single-peaked and unshilted in velocity., These latter profiles are usually single-peaked and unshifted in velocity.
+ Although the profile shape for KIT 15D is not fully clear due to imperfect subtraction of the emission of [OI] at 6300 in the night sky. it is most likely double-peaked.," Although the profile shape for KH 15D is not fully clear due to imperfect subtraction of the emission of [OI] at 6300 in the night sky, it is most likely double-peaked."
+ Such a profile can be most easily explained bv a bipolar jet moving nearly perpendicular to the line of sight. quite consistent will our expectation that the disk associated with KIL 15D is viewed nearly edee on. resulting in a verv small radial velocity separation between (he (wo jets.," Such a profile can be most easily explained by a bipolar jet moving nearly perpendicular to the line of sight, quite consistent with our expectation that the disk associated with KH 15D is viewed nearly edge on, resulting in a very small radial velocity separation between the two jets."
+ The prolile catalog of Hartiganetal.(1995) contains several examples of bipolar jet sources (RW Aur. AS 353A. DD Tan). where the two jet components are clearly sepawated due to favorable inclination angles.," The profile catalog of \citet{Hartigan95} contains several examples of bipolar jet sources (RW Aur, AS 353A, DD Tau), where the two jet components are clearly separated due to favorable inclination angles."
+ It is important to keep in mind that we do not vet know the geometry of the 1Η 15D svstem., It is important to keep in mind that we do not yet know the geometry of the KH 15D system.
+ In particular. i( is not certain whether the occultation proceeds along a line perpendicular to or parallel to the rotation axis of the INT star (or neither) or to what extent the rotation axis. the presumed stellar magnetic axis and (he orbital plane are aligned.," In particular, it is not certain whether the occultation proceeds along a line perpendicular to or parallel to the rotation axis of the K7 star (or neither) or to what extent the rotation axis, the presumed stellar magnetic axis and the orbital plane are aligned."
+ It is not even known whether. relative to the svstem's center of mass. it is the INT star or the oceulling matter which is primarily in motion.," It is not even known whether, relative to the system's center of mass, it is the K7 star or the occulting matter which is primarily in motion."
+ Detailed modeling of the system must obviously awalt clarification of these basic issues., Detailed modeling of the system must obviously await clarification of these basic issues.
+ ILowever. a preliminary (qualitative interpretation of the emission line variations is possible and provided in this section.," However, a preliminary qualitative interpretation of the emission line variations is possible and provided in this section."
+" We believe (hat the observed spectral variations of KIT 15D can be understood qualitatively in ternis of a weakly accreting. ""scaled down” classical TD Tauri star (see e.g.. Muzerolle (2001))) whose photosphere and magnetosphere are periodically occulted by a relatively sharp knife-edge (as described in Herbstοἱal...(2002).. their Fie."," We believe that the observed spectral variations of KH 15D can be understood qualitatively in terms of a weakly accreting, “scaled down"" classical T Tauri star (see e.g., \citet{Muz01}) ) whose photosphere and magnetosphere are periodically occulted by a relatively sharp knife-edge (as described in \citet{Herbst02}, their Fig."
+ 5)., 5).
+ We expect such a star {ο have a bipolar jet region. which we assume. by analogy with the CTTS. is revealed by the," We expect such a star to have a bipolar jet region, which we assume, by analogy with the CTTS, is revealed by the"
+The CoRoT satellite was successfully launched in 2006.,The CoRoT satellite was successfully launched in 2006.
+ Onboard CoRoT. there is à small 27em telescope feeding two science channels to study astroseismology and transits. respectively (?)..," Onboard CoRoT, there is a small 27cm telescope feeding two science channels to study astroseismology and transits, respectively \citep{2000JApA...21..319B}."
+" The CoRoT has a field of view (FOV) of - 2.7"" x 3.05"".", The CoRoT has a field of view (FOV) of $\sim$ $2.7^{o}$ $x$ $3.05^{o}$.
+" In its first observed field (IRa01 - a = 646""53 ó = -00""12007). CoRoT observed continuously for 60 days. producing uninterrupted light curves for the first time."," In its first observed field (IRa01 - $\alpha$ = $6^{h}46^{m}53^{s}$ $\delta$ = $-00^{o}12^{'}00{""}$ ), CoRoT observed continuously for 60 days, producing uninterrupted light curves for the first time."
+ The data for the IRaO1 field have been public since December 2008 and the astronomical community has access to these data., The data for the IRa01 field have been public since December 2008 and the astronomical community has access to these data.
+" Unfortunately. the CoRoT light curves are affected by a variety of instrumental problems that severely hamper the ""Sata interpretation."," Unfortunately, the CoRoT light curves are affected by a variety of instrumental problems that severely hamper the data interpretation."
+ To overcome these difficulties. we have developed the CoRot detrend algorithm (CDA)).," To overcome these difficulties, we have developed the CoRot detrend algorithm )."
+ In this paper. the algorithm is presented and we demonstrate its functionality on some typical CoRoT data sets.," In this paper, the algorithm is presented and we demonstrate its functionality on some typical CoRoT data sets."
+ The CoRoT data files contain multi-color light curves that were produced by inserting a low-resolution dispersing prism into the telescope beam., The CoRoT data files contain multi-color light curves that were produced by inserting a low-resolution dispersing prism into the telescope beam.
+ This set-up is desinged to provide simultaneous light curves in the red (R). green (G). and blue (B) bands. although. these bands do not correspond to true photometric filters and. the bands may indeed differ from star to star.," This set-up is desinged to provide simultaneous light curves in the red (R), green (G), and blue (B) bands, although, these bands do not correspond to true photometric filters and, the bands may indeed differ from star to star."
+ In this paper. we study the multi-color data. but also consider the total (white) flux obtained by summing up the individual light curves through W=R+G+B.," In this paper, we study the multi-color data, but also consider the total (white) flux obtained by summing up the individual light curves through $W=R+G+B$."
+ Figure 1. are typical CoRoT light curves from [RaQ]., Figure \ref{fig1} are typical CoRoT light curves from IRa01.
+ The first panel of Fig., The first panel of Fig.
+ | shows a typical exponential jump very similar to a flare star., \ref{fig1} shows a typical exponential jump very similar to a flare star.
+ A trend is also evident., A trend is also evident.
+ In the second light curve. there appears to be a box-shape jump. and in the third and fourth light curves one can discern features similar to those i1 the first and second light curves. except that the jumps are downwards.," In the second light curve, there appears to be a box-shape jump, and in the third and fourth light curves one can discern features similar to those in the first and second light curves, except that the jumps are downwards."
+ We note that the downward jump in the third light curve is very similar to a transit event. thus making the detection of true transits difficult.," We note that the downward jump in the third light curve is very similar to a transit event, thus making the detection of true transits difficult."
+ Combinations of all the above features appear to produce a rather typical CoRoT light curve., Combinations of all the above features appear to produce a rather typical CoRoT light curve.
+ There are. two basic instrumental problems with all CoRoT light curves.," There are, two basic instrumental problems with all CoRoT light curves."
+ First. there ts a long-term trend. forcing a secular decrease in the light curve intensity over the full observing period of 60 days.," First, there is a long-term trend, forcing a secular decrease in the light curve intensity over the full observing period of 60 days."
+ The strengths of the trends in different sources may be different: the physical cause of these trends is not well understood., The strengths of the trends in different sources may be different; the physical cause of these trends is not well understood.
+ The second and evei more serious problem is the instrumental jumps in the light curves., The second and even more serious problem is the instrumental jumps in the light curves.
+" The term ""jump"" refers to a sudden variation in intensity without any obvious reason.", The term “jump” refers to a sudden variation in intensity without any obvious reason.
+ Many of these jumps resemble stellar flares. although. the vast majority are clearly instrumental nature.," Many of these jumps resemble stellar flares, although, the vast majority are clearly instrumental nature."
+ Their physical explanation could be cosmic radiation and the time evolution of bright pixels (2).., Their physical explanation could be cosmic radiation and the time evolution of bright pixels \citep{2008MNRAS.384.1337P}.
+ These jumps are a random phenomenon and affect each filter differently., These jumps are a random phenomenon and affect each filter differently.
+ An inspection of hundreds of CoRoT light curves similar to those presented in Fig., An inspection of hundreds of CoRoT light curves similar to those presented in Fig.
+ | allows us to classify the observed shapes of jumps into five groups: A statistical analysis of [RaQ] field (visual inspection) shows that only a small minority (Table 1)) of all jumps are powerful enough to simultaneously appear in each colour., \ref{fig1} allows us to classify the observed shapes of jumps into five groups: A statistical analysis of IRa01 field (visual inspection) shows that only a small minority (Table \ref{tab1}) ) of all jumps are powerful enough to simultaneously appear in each colour.
+ Most of the light curves are affected not only by one single jump. but by many jumps occurring in the different filters at different times.," Most of the light curves are affected not only by one single jump, but by many jumps occurring in the different filters at different times."
+ In Table Τ.. we show the results of a statistical study of the appearance and the shapes of Jumps using data for IRaOI.," In Table \ref{tab1}, we show the results of a statistical study of the appearance and the shapes of jumps using data for IRa01."
+ The three first columns of Table 1. show the number of light curves affected by jumps in the respective filter filter and the fourth column indicates the total size of the effect., The three first columns of Table \ref{tab1} show the number of light curves affected by jumps in the respective filter filter and the fourth column indicates the total size of the effect.
+ It is quite difficult to describe all the features perturbing a CoRoT light curve with a given function. since there are many," It is quite difficult to describe all the features perturbing a CoRoT light curve with a given function, since there are many"
+Ilexion. formalism.,flexion formalism.
+ For à more complete discussion. of the theory ancl description of the various techniques used. to measure Hexion. the reader is referred to Goldberg Bacon (2005). Bacon et al. (," For a more complete discussion of the theory and description of the various techniques used to measure flexion, the reader is referred to Goldberg Bacon (2005), Bacon et al. ("
+2006). Goldberg Leonarcl (2007). Leonard οἱ al. (,"2006), Goldberg Leonard (2007), Leonard et al. ("
+2008) and Okura οἱ al. (,2008) and Okura et al. (
+2007. 2008).,"2007, 2008)."
+ Cravitational flexion arises when the lens field varies significantIy over the scale of the lensed image., Gravitational flexion arises when the lens field varies significantly over the scale of the lensed image.
+ In this case. the lens equation taken to be linear in weak lensing stuclics must be extended to second order: where 8 is the coordinate in the source plane ancl @ is the coordinate in the lens plane.," In this case, the lens equation – taken to be linear in weak lensing studies – must be extended to second order: where $\bb$ is the coordinate in the source plane and $\bt$ is the coordinate in the lens plane."
+ These coordinates are related by 4-6αίθ). where a(@) is the dellection angle induced by the lens potential.," These coordinates are related by $\bb=\bt-\ba(\bt)$, where $\ba(\bt)$ is the deflection angle induced by the lens potential."
+ is the magnification matrix. which is related to the convergence. &. and the complex eravitational shear ?=l-22ie. and D;ji=0v; is related o the first Ilexion. ν=FleOs. and second. [exion. G=|Gere—s. only.," is the magnification matrix, which is related to the convergence, $\kappa$, and the complex gravitational shear $\gamma=|\gamma|e^{2i\phi}$, and $D_{ijk}=\partial_kA_{ij}$ is related to the first flexion, $\f=|\f|e^{i\phi}=\partial\kappa$, and second flexion, $\g=|\g|e^{3i\phi}=\partial\gamma$, only."
+ The first Hexion signal gives rise to a skewness in the rightness distribution of the lensed galaxy image. and. is a direct. probe of the local gradient. of the convergence.," The first flexion signal gives rise to a skewness in the brightness distribution of the lensed galaxy image, and is a direct probe of the local gradient of the convergence."
+ As such. it olfers an ideal probe of cluster substructure. as it ends to be more sensitive to small-scale structures than to he large-scale cluster potential outside of the critical region of the cluster.," As such, it offers an ideal probe of cluster substructure, as it tends to be more sensitive to small-scale structures than to the large-scale cluster potential outside of the critical region of the cluster."
+ The second Uexion signal gives rise to a bending or arciness in the lensed galaxy image., The second flexion signal gives rise to a bending or arciness in the lensed galaxy image.
+ Although we expect he second Uexion signal to be larger for most mass profiles han the first flexion signal (see. e.g.. Goldberg Bacon. 2005: Lasky Fluke. 2009). Goldberg Leonard. (2007) and Leonard et al. (," Although we expect the second flexion signal to be larger for most mass profiles than the first flexion signal (see, e.g., Goldberg Bacon, 2005; Lasky Fluke, 2009), Goldberg Leonard (2007) and Leonard et al. ("
+2008) have found second Uexion to »f slenificanthy more diüllicult to measure.,2008) have found second flexion to be significantly more difficult to measure.
+ We therefore restrict. our discussion for the remainder of this paper to measurements of first Hexion., We therefore restrict our discussion for the remainder of this paper to measurements of first flexion.
+ The convergence is related to he first Dlexion through: where , The convergence is related to the first flexion through: where and $X=x_1+ix_2$.
+To date. two groups have presented: gravitational Uexion studies of a cluster of galaxies with encouraging results.," To date, two groups have presented gravitational flexion studies of a cluster of galaxies with encouraging results."
+ Leonard οἱ al. (, Leonard et al. (
+2008) presented the first measurements of Ilexion in a cluster of galaxies. and described a parametric modelling technique to reconstruct the distribution of mass within the ealaxy cluster Abell 1689.,"2008) presented the first measurements of flexion in a cluster of galaxies, and described a parametric modelling technique to reconstruct the distribution of mass within the galaxy cluster Abell 1689."
+ Whilst this method showed good results in the periphery of the cluster. allowing mapping of the substructure content of the cluster. it was not able to constrain the central mass profile at all as a result of various masking techniques emploved by the authors.," Whilst this method showed good results in the periphery of the cluster, allowing mapping of the substructure content of the cluster, it was not able to constrain the central mass profile at all as a result of various masking techniques employed by the authors."
+ Okura ct al. (, Okura et al. (
+2008) demonstrated a nonparanietric Hexion reconstruction technique on the same cluster. with good results. beine obtained using only 5. background sources per square arcminute and a positive detection in the central region. of the cluster.,"2008) demonstrated a nonparametric flexion reconstruction technique on the same cluster, with good results being obtained using only 5 background sources per square arcminute and a positive detection in the central region of the cluster."
+ Their method involves bin-averaging the [lexion signal ancl reconstructing the convergence according to Equation 2.., Their method involves bin-averaging the flexion signal and reconstructing the convergence according to Equation \ref{eq:kapx}.
+ This is performed most simply and cllicienth in. Fourier. space. where the relationship is expressed as (Bacon et al..," This is performed most simply and efficiently in Fourier space, where the relationship is expressed as (Bacon et al.,"
+ 2006): where .X represents the Fourier. transform of Y., 2006): where $\tilde{X}$ represents the Fourier transform of $X$.
+ In taking the Fourier transform of & one recovers #8o. where Ho ds à constant of integration usually set by requiring the convergence to &o to zero at large cistances from the cluster centre.," In taking the Fourier transform of $\tilde{\kappa}$ one recovers $\kappa-\kappa_0$, where $\kappa_0$ is a constant of integration usually set by requiring the convergence to go to zero at large distances from the cluster centre."
+ Such reconstructions may be problematic. however. as the resolution one can obtain in the output. convergence map is stronely dependent on the background source density ancl distribution.," Such reconstructions may be problematic, however, as the resolution one can obtain in the output convergence map is strongly dependent on the background source density and distribution."
+ Moreover. in the presence of significant substructure. bin-averaging over large pixels may. have the ellect. of washing out the llexion. signal.," Moreover, in the presence of significant substructure, bin-averaging over large pixels may have the effect of washing out the flexion signal."
+ In. addition. any masking schemes that result in a deficiency of sources in the central region of the cluster will result in the reconstructed convergence in this region being underestimated.," In addition, any masking schemes that result in a deficiency of sources in the central region of the cluster will result in the reconstructed convergence in this region being underestimated."
+ His for this reason that a parametric technique was chosen. by Leonard. ct al. (, It is for this reason that a parametric technique was chosen by Leonard et al. (
+2008) as a more appropriate method for reconstructing the mass distribution in Abell 1689. from Hexion measurements.,2008) as a more appropriate method for reconstructing the mass distribution in Abell 1689 from flexion measurements.
+ Aperture measures have long been used. in weak lensing. and the clleet of aperture filtering is to provide a robust measure of the mass distribution within lens systems whilst. simultaneously. providing a straightforward: method [or determining the noise properties of the mass maps generated.," Aperture measures have long been used in weak lensing, and the effect of aperture filtering is to provide a robust measure of the mass distribution within lens systems whilst simultaneously providing a straightforward method for determining the noise properties of the mass maps generated."
+ In this way. aperture filtering ollers an alternative nonparametric reconstruction method. which allows one to detect and map out mass concentrations within lens svstems without requiring knowledge or assumptions regarding the shape of the mass profile of the lens.," In this way, aperture filtering offers an alternative nonparametric reconstruction method, which allows one to detect and map out mass concentrations within lens systems without requiring knowledge or assumptions regarding the shape of the mass profile of the lens."
+ The aperture mass statistic is defined as (Schneider 1996: LIXW09): Using Equation 2. above. we can re-define this in ternis of the measured [Dexion as This can be made independent ofthe constant quantity Ho dM we require that the mass filter. function. wr) is compensated. i6. Alaking the transformation y=x/ Xy. ancl writing," The aperture mass statistic is defined as (Schneider 1996; LKW09): Using Equation \ref{eq:kapx} above, we can re-define this in terms of the measured flexion as This can be made independent ofthe constant quantity $\kappa_0$ if we require that the mass filter function $w(x)$ is compensated, i.e. Making the transformation $\by=\bx^\prime-\bx_0$ , and writing"
+"i They are exceedingly luminous. comparable with liminosity of a whole galaxy: they can. (hus. be detected aud observed with high S/N ratio even al cosmological distances. H) """,") They are exceedingly luminous, comparable with luminosity of a whole galaxy; they can, thus, be detected and observed with high S/N ratio even at cosmological distances. ) “"
+"Normal"" SNe Ia have small variations among their peak absolute magnitudes (around ος",Normal” SNe Ia have small variations among their peak absolute magnitudes (around 0.3).
+ The accelerated expansion of our Universe was discovered as a result of (wo projects: the IHigh-z SN Search (Schmidtetal.1993:Riess1998) and the Supernova Cosmology Project (Perlmutteretal.1999.," The accelerated expansion of our Universe was discovered as a result of two projects: the High-z SN Search \citep{sch98,rie98} and the Supernova Cosmology Project \citep{per99}."
+a).. In fact. a new tvpe of matter was preclicted many vears ago by A. Einstein. who included a A-term into his considerations (Einstein1917).," In fact, a new type of matter was predicted many years ago by A. Einstein, who included a $\Lambda$ -term into his considerations \citep{ein17}."
+. At the beginning of the past century. the A-term was just a new fundamental eonstant. and only much later it was really considered as a [ormidable challenge by both observational and theoretical cosmologists (Weinberg1939:2000:Rubakoy 2000).," At the beginning of the past century, the $\Lambda$ -term was just a new fundamental constant, and only much later it was really considered as a formidable challenge by both observational and theoretical cosmologists \citep{wein89,zel90,zel92,cpt92,carr00,star00,rub00}."
+". Moreover. during the last 20 vears cosmologists understood that this constant can be replaced with a scalar field. which induces the repulsive gravitational force dvnamically,"," Moreover, during the last 20 years cosmologists understood that this constant can be replaced with a scalar field, which induces the repulsive gravitational force dynamically."
+ Accordingly. several models were proposed (RatraancPeebles1933:Bineltruy2000:RubanoandSeuclellaro 2001).. in order to explain the observed. present acceleration of our Universe.," Accordingly, several models were proposed \citep{rat88,pee88,wet95,wet98,cop98,fer98,lid99,ste99,bra99,sah00,bin00,rub01}, in order to explain the observed present acceleration of our Universe."
+ Two of these models continued to be developed after the discovery of acceleration. aud were also roughly elaborated and adapted. for present-day data 2001).," Two of these models continued to be developed after the discovery of acceleration, and were also roughly elaborated and adapted for present-day data \citep{rub01}."
+. Here. we again use these models to such an end. but in a much more refined way: the eoal is now to fit the observed data of apparent magnitude and redshilt of the supernovae la. and test (he models themselves.," Here, we again use these models to such an end, but in a much more refined way: the goal is now to fit the observed data of apparent magnitude and redshift of the supernovae Ia, and test the models themselves."
+ As said above. in this paper we discuss two models for quintessence. both based on a scalar field with a special type of potential.," As said above, in this paper we discuss two models for quintessence, both based on a scalar field with a special type of potential."
+ The field is minimally coupled with pressureless matter. and the total density parameter O of the Universe is lixed to be 1.," The field is minimally coupled with pressureless matter, and the total density parameter $\Omega$ of the Universe is fixed to be 1."
+ A detailed discussion of the consequences of assuming such models in cosmology is given in 2001).. so that we limit ourselves here only to à short summary of the results we need for our purpose.," A detailed discussion of the consequences of assuming such models in cosmology is given in \citep{rub01}, so that we limit ourselves here only to a short summary of the results we need for our purpose."
+ The main attractive leature of these models is that they allow a general exact solution, The main attractive feature of these models is that they allow a general exact solution
+production of good mock catalogs.,production of good mock catalogs.
+ Since the CiC method only uses information on scales where the one-halo term dominates. our method is less sensitive to knowing the two-halo correlation function on  Mpce// scales.," Since the CiC method only uses information on scales where the one-halo term dominates, our method is less sensitive to knowing the two-halo correlation function on $\sim$ $h$ scales."
+ We have demonstrated that the deficit of 6vo-halo pairs al ~1 Mpc/h is problematic for HOD fitting methods using 2 and 3-pt statistics and FoF catalogs because the WOD must accommodate the missing (wo halo pairs at this separation. and thus predict groups containing many more LRGs (han are Finally. the low number density of the LRGs makes them particularly. well-suited for ihis The low munber density of SDSS LRGs allows us to partially separate LRG pairs occupying (he same dark matter halo Lom pairs occupying distinct. dark matter halos.," We have demonstrated that the deficit of two-halo pairs at $\sim 1$ $h$ is problematic for HOD fitting methods using 2 and 3-pt statistics and FoF catalogs because the HOD must accommodate the missing two halo pairs at this separation, and thus predict groups containing many more LRGs than are Finally, the low number density of the LRGs makes them particularly well-suited for this The low number density of SDSS LRGs allows us to partially separate LRG pairs occupying the same dark matter halo from pairs occupying distinct dark matter halos."
+ Candidate one-halo pairs are identified using simple cuts in the transverse and LOS separations., Candidate one-halo pairs are identified using simple cuts in the transverse and LOS separations.
+ We eroup these pairs using the FoF algorithm (o compute the CiC eroup multiplicity function., We group these pairs using the FoF algorithm to compute the CiC group multiplicity function.
+" We measure (he CiC group multiplicity Iunction lor the subsample of SDSS LRGs saldslving —23.2<AL,—21.2 and 0.16<20.36. carefully accounting for the effects of fiber collisions and survey boundaries. holes. and In order to derive ΠΟ constraints [rom our measurement. we calibrated (he relation between the CiC and true one-halo group multiplicity functions using mock LRG catalogs."," We measure the CiC group multiplicity function for the subsample of SDSS LRGs satisfying $-23.2 < M_{g} < -21.2$ and $ 0.16 < z < 0.36$, carefully accounting for the effects of fiber collisions and survey boundaries, holes, and In order to derive HOD constraints from our measurement, we calibrated the relation between the CiC and true one-halo group multiplicity functions using mock LRG catalogs."
+ The variance about the mean relation is comparable to the Poisson sampling variance of the CiC multiplicity function and must be properly accounted for in the maximum likelihood parameter The CiC eroup multiplicity Duiction places strong constraints on the satellite LRG 110D. Ny(AL).," The variance about the mean relation is comparable to the Poisson sampling variance of the CiC multiplicity function and must be properly accounted for in the maximum likelihood parameter The CiC group multiplicity function places strong constraints on the satellite LRG HOD, $N_{sat}(M)$."
+" When we fix σι=0.7 and nig;=10! (Mpc/h)?. the maximum likelihood HOD parameters and marginalized one-dimensional and confidence intervals are Mou=5.0EC32)xLOMAL M,=4.95OSC3)XLOMA anda=1.035(C31).with Moin=8.05xLOMAL. at the maximum likelihood point."," When we fix $\sigma_{log M} = 0.7$ and $\bar{n}_{LRG} = 10^{-4}$ $h)^{-3}$, the maximum likelihood HOD parameters and marginalized one-dimensional and confidence intervals are $M_{cut} = 5.0^{+1.5}_{-1.3} (^{+2.9}_{-2.6}) \times 10^{13} M_{\sun}$, $M_1 = 4.95^{+0.37}_{-0.26} (^{+0.79}_{-0.53}) \times 10^{14} M_{\sun}$, and $\alpha = 1.035^{+0.10}_{-0.17} (^{+0.24}_{-0.31})$,with $M_{min} = 8.05 \times 10^{13} M_{\sun}$ at the maximum likelihood point."
+ We tightly constrain the satellite ⋅⋅⋅ 9Bey↽⋅ ⋅ ∐⋅≀↧↴≺∢⊔∪∐∩≻∙∕⊽∖⊓∕∶∩⋅∩≺≻⋮↽⊰↻⋮⋮⋮⋮⋮⋮≜∣∣⋖⋮⋮⋮⋮⋮⋮∖⋝⋝⊺↥, We tightly constrain the satellite fraction to $f_{sat} = 0.0636^{+0.0019}_{-0.0020} (^{+0.0038}_{-0.0039})$.
+∐↲↕≻↕⋅∪∙↿≼↲≺∢∩↲≼⇂≺∢∪↕⋅↕⋅≼↲↥≀↕↴⊔∪∐↓∏∐≺∢⊔∪∐∣∣⋅↙⋣⋖↜∣⋮↙↾∣⋝∪↓∐↓∪≺∢↳↽ ≺∢≀↧↴↥≀↧↴↥∪≸≟⊳∖⇁≺⇂≼↲∏∖↽≼↲≼⇂∐⋅∪∐↓≀↧↴∐⊳↔⊲↻↥⋯↴↥∪≺∢≀↧↴↥≀↧↴↥∪≸≟↥⊳∖⊽≀↧↴∐≼↲⇀↸≺∢≼↲∐≼↲∐↥≀↧↴≸↽↔↴↕⋅≼↲≼↲∐∐↲∐↥∖∖⊽↥⊔↥⊔∐↲∐∐↲≀↧↪∖⊽⋯⋅≼↲∐∐↲↕∐⊳∖⇁ ∪↓⇝∙↗≀↧↴∐≼⇂∊∙↗∖∖⊽↥∐↲∐⊔∐↲↥≀↧↴↕⋅≸↽↔↴≼↲⋝∖⊽≺∢≀↧↴↥≼↲≺∢↥∏⋝∖⊽∩↲↕⋅↕∐≸≟↕⋝∖⊽∏⊳∖, The projected correlation function $w_p(r_p)$ of mock catalogs derived from an SO halo catalog is an excellent agreement with the measurements of \citet{masjedi/etal:2006} and \citet{zehavi/etal:2005a} when the large scale clustering is used to fix $\sigma_{log M}$ .
+⇁≼↲≺⇂, Fig.
+∩≻∐⇀↸⊔∣∩∙↙∕⋏⊔⋅∏≸≟⋅≤∍⊳∖⊽∐∪∖∖↽⊳∖⇁⊔⋯↥⊟⊞, \ref{fig:fofcompare} shows that FoF
+"investigated whether a non-constant tensile strength value, defined by a power-law with slope of s, will affect the steady state particle size distribution.","investigated whether a non-constant tensile strength value, defined by a power-law with slope of $s$, will affect the steady state particle size distribution."
+" They found a linear relationship between the two power-law slopes, yielding a steady state solution producing more small particles than the traditional Dohnanyi(1969) solution, when considering realistic tensile strength slopes."," They found a linear relationship between the two power-law slopes, yielding a steady state solution producing more small particles than the traditional \cite{dohnanyi69} solution, when considering realistic tensile strength slopes."
+" An analytic solution similar to the one by &Greenberg(2003) was presented in Wyattetal. who also analyzed their equations numerically, revealing(2011),, substructures superimposed on the power-law size distributions."," An analytic solution similar to the one by \cite{obrien03} was presented in \cite{wyatt11}, who also analyzed their equations numerically, revealing substructures superimposed on the power-law size distributions."
+" Belyaev&Rafikov(2011) expanded on the O'Brien&Greenberg(2003) model by allowing a variable largest fragment mass to target mass ratio (X/M in our model), following the Fujiwaraetal.(1977) experimental results."," \cite{belyaev11} expanded on the \cite{obrien03} model by allowing a variable largest fragment mass to target mass ratio $X/M$ in our model), following the \cite{fujiwara77} experimental results."
+" Belyaev&Rafikov found particle distribution that could not be described(2011) by a asingle power-law, rather by a slowly varying, mass dependent power-law function."," \cite{belyaev11} found a particle distribution that could not be described by a single power-law, rather by a slowly varying, mass dependent power-law function."
+ We summarize the characteristic properties of these analytic models in Table 2.., We summarize the characteristic properties of these analytic models in Table \ref{tab:companal}.
+" A common issue with all of these analytic solutions is that they assume a steady state system, meaning that the mass flux through mass m is zero [OF(m)/8m= 0]."," A common issue with all of these analytic solutions is that they assume a steady state system, meaning that the mass flux through mass $m$ is zero $\partial F(m) / \partial m = 0$ ]."
+" In reality, without mass input at the high mass end the systems lose mass and decay, meaning the mass flux is never zero in a system, but rather is"," In reality, without mass input at the high mass end the systems lose mass and decay, meaning the mass flux is never zero in a system, but rather is"
+"ie.. (he blazar direction) has been chosen as an input (o our cascade simulation to evaluate the cascade emission in the nuclear infrared radiation field of Cen A observed αἱ the given angle of 8,4,=70°.","i.e., the blazar direction) has been chosen as an input to our cascade simulation to evaluate the cascade emission in the nuclear infrared radiation field of Cen A observed at the given angle of $\theta_{\rm obs} = 70^o$."
+ For the cascade simulation. we assumed a blackbocly temperature of 2300 Ix. resulting in a peak [requency in the Ix-band.," For the cascade simulation, we assumed a blackbody temperature of $2300$ K resulting in a peak frequency in the K-band."
+ The external radiation field is parameterized (hroughΠω=1.5x107 erg * and Ray=3xLOM em., The external radiation field is parameterized through$u_{\rm ext} = 1.5 \times 10^{-3}$ erg $^{-3}$ and $R_{\rm ext} = 3 \times 10^{16}$ cm.
+ These parameters combine to an IR luminosity of Lup;=4drR2cusa5x108 erg 1+. in agreement with mid-IR. flux observed for Cen A (Marconietal.2000).," These parameters combine to an IR luminosity of $L_{\rm BLR} = 4 \pi R_{\rm ext}^2 \, c \, u_{\rm ext}
+= 5 \times 10^{41}$ erg $^{-1}$, in agreement with mid-IR flux observed for Cen A \citep{marconi00}."
+. The magnetic field is 2=I mG. oriented at an angle of 05=4.," The magnetic field is $B = 1$ mG, oriented at an angle of $\theta_B = 4^o$."
+ The cascade spectrum shown in Figure 7 pertains to the angular bin Q.28<pe0.38 (corresponding to 67°<@ τοῦ. appropriate for the known orientation ol Cen A and consistent with our broadband SED fit parameters.," The cascade spectrum shown in Figure \ref{fit} pertains to the angular bin $0.28 < \mu < 0.38$ (corresponding to $67^o \lesssim
+\theta \lesssim 73^o$ ), appropriate for the known orientation of Cen A and consistent with our broadband SED fit parameters."
+ The cascade spectrum is shown bv (he maroon curve in Figure 7.. while the total observed spectrum is (he solid red curve.," The cascade spectrum is shown by the maroon curve in Figure \ref{fit}, while the total observed spectrum is the solid red curve."
+ The figure illustrates that the cascade contribution in the range substantially improves the fit. while still allowing physically reasonable parameters for the broadband SED fil.," The figure illustrates that the cascade contribution in the range substantially improves the fit, while still allowing physically reasonable parameters for the broadband SED fit."
+ We investigated (he signatures of Compton-supportecl pair cascades initiated by the interaction of nuclear VIIE 5-ravs with the thermal infrared radiation field of a cireiumnuclear dust torus in AGNs., We investigated the signatures of Compton-supported pair cascades initiated by the interaction of nuclear VHE $\gamma$ -rays with the thermal infrared radiation field of a circumnuclear dust torus in AGNs.
+ We follow the spatial development of the cascade in [ull 3-dimensional eeomeltry and study the dependence of the radiative output on various parameters pertaining to the infrared radiation field and the magnetic field in (he cascade region., We follow the spatial development of the cascade in full 3-dimensional geometry and study the dependence of the radiative output on various parameters pertaining to the infrared radiation field and the magnetic field in the cascade region.
+ We confirm the results of our previous study. of cascades in monoenergetic radiation fields that small (2 µία) perpendicular (to the primary VIIE 4-rav beam) magnetic field components lead to efficient isotropization of the cascade emission oul to ILE 5-ray energies., We confirm the results of our previous study of cascades in monoenergetic radiation fields that small $\gtrsim \mu$ G) perpendicular (to the primary VHE $\gamma$ -ray beam) magnetic field components lead to efficient isotropization of the cascade emission out to HE $\gamma$ -ray energies.
+ The cascade intensity as well as the location of a high-energv turnover due to inefficient isolropizalion also depend sensitivelv on the energy density. and temperature of the solt blackbody radiation field. as long as the cascade is not saturated in the sense that not all VIIE 5-ravs are absorbed.," The cascade intensity as well as the location of a high-energy turnover due to inefficient isotropization also depend sensitively on the energy density and temperature of the soft blackbody radiation field, as long as the cascade is not saturated in the sense that not all VHE $\gamma$ -rays are absorbed."
+ The shape of the low-frequency tail of the cascade emission is a result of the interplay between radiative cooling and escape., The shape of the low-frequency tail of the cascade emission is a result of the interplay between radiative cooling and escape.
+" For environments characterized bv efficient radiative cooling. the canonical pF,xpU? spectrum results."," For environments characterized by efficient radiative cooling, the canonical $\nu F_{\nu} \propto \nu^{1/2}$ spectrum results."
+ IL radiative cooling is inefficient compared Lo escape. the low-Irequenevy. cascade spectra are harder than pl.," If radiative cooling is inefficient compared to escape, the low-frequency cascade spectra are harder than $\nu^{1/2}$ ."
+ We provide a model fit to the broadband SED of the dust-rich. 5-rav loud radio galaxy," We provide a model fit to the broadband SED of the dust-rich, $\gamma$ -ray loud radio galaxy"
+The MDI instrument aboard SOMO spacecraft provides the full-disk magnetogram with a pixel size of 2.,The MDI instrument aboard SOHO spacecraft provides the full-disk magnetogram with a pixel size of 2”.
+ Ir order to obtain a low noise level. only those 5-niin average magnetogranmis are selected in the study.," In order to obtain a low noise level, only those 5-min average magnetograms are selected in the study."
+ We extract one observed full-disk magnetogram per day. and thusly we totallv select 3764 magnetograms [rom 1996 September to 2010 February. which include the complete 23rd solar evele.," We extract one observed full-disk magnetogram per day, and thusly we totally select 3764 magnetograms from 1996 September to 2010 February, which include the complete 23rd solar cycle."
+" In order to further reduce (he noise level. we apply à boxcar smoothing bhuncetion to each magnetogram by a width of x6""."," In order to further reduce the noise level, we apply a boxcar smoothing function to each magnetogram by a width of $\times$ 6”."
+ There are (wo groups of authors who first pointed out the under-estimation of magnetic [lux bv earlier MDI magnetogram calibration (Berger et al., There are two groups of authors who first pointed out the under-estimation of magnetic flux by earlier MDI full-disk magnetogram calibration (Berger et al.
+ 2003: Wang et al., 2003; Wang et al.
+ 2003)., 2003).
+ All the magnetograms used in this study are that retrieved alter. recalibration of December 2008., All the magnetograms used in this study are that retrieved after recalibration of December 2008.
+ For a better understanding about the evelie behavior of solar minima of Cveles 22 ancl 23. we extend the MIDI data base by adding Ixitt Peak full-disk magnetograms from August 1996 back to January 1994.," For a better understanding about the cyclic behavior of solar minima of Cycles 22 and 23, we extend the MDI data base by adding Kitt Peak full-disk magnetograms from August 1996 back to January 1994."
+ The data merging is made based on a least-square fitting of the mean flux density of Ixitt Peak magnetograms to that of MDI magnetograms lor the common interval ol 1996., The data merging is made based on a least-square fitting of the mean flux density of Kitt Peak magnetograms to that of MDI magnetograms for the common interval of 1996.
+ We estimate the noise level of these smoothed 5-min average magnetograms according io the method described by Lagenaar (2001) and Hagenaar et al. (, We estimate the noise level of these smoothed 5-min average magnetograms according to the method described by Hagenaar (2001) and Hagenaar et al. (
+2003).,2003).
+ Based on these magnetograms. we analvze their histograms of magnetic flix density.," Based on these magnetograms, we analyze their histograms of magnetic flux density."
+" The core of the distribution function is fitted by a Gaussian [unction. Fr)=Ej;,0rp(—27/207). where the width σ of the Gaussian—. function.B. about 6"" >Mx/enr'. isB defined. as the noiseB level."," The core of the distribution function is fitted by a Gaussian function $
+F(x)=F_{max}exp(-x^2/2\sigma^2)$ , where the width $\sigma$ of the Gaussian function, about 6 $^{2}$, is defined as the noise level."
+ We assume that the observed line of sight magnetic flux densitv is a projection of the intrinsic flix density normal to the solar surface. so (he magnetic flux density for each pixel is corrected(see Hagenaar 2001 and Hagenaar et al.," We assume that the observed line of sight magnetic flux density is a projection of the intrinsic flux density normal to the solar surface, so the magnetic flux density for each pixel is corrected(see Hagenaar 2001 and Hagenaar et al."
+" 2003) as D,=Du(o)/costa).", 2003) as $B_{cal}=B_{obs}(\alpha)/\cos(\alpha)$.
+ The angle o of each pixel is defined by sin(a)=ya?+y?/IH Where x and Y are the pixel position relerring to the disk center. al which x and v is equal to 0. and the Rois the solar disk radius.," The angle $\alpha$ of each pixel is defined by $\sin(\alpha)=\sqrt{x^{2}+y^{2}}/R$ Where x and Y are the pixel position referring to the disk center, at which x and y is equal to 0, and the R is the solar disk radius."
+ After the correlation. the magnetogram shows the magnetic [αν density normal to solar surface.," After the correlation, the magnetogram shows the magnetic flux density normal to solar surface."
+ Alter the angle is greater GO degrees. (here are less and less magnetic signals due to (he lower magnetic sensitivity aud spatial resolution of MDI/SOIIO magnetogranms. ancl the magnetic noise level would increase according to the magnetic correction 1/cos(o).," After the angle is greater 60 degrees, there are less and less magnetic signals due to the lower magnetic sensitivity and spatial resolution of MDI/SOHO magnetograms, and the magnetic noise level would increase according to the magnetic correction $\alpha$ )."
+ Therefore. we only analyze these pixels with angle a less (han 60 degree. ie.. the region included by the black circle in Che left panels of Fig.," Therefore, we only analyze these pixels with angle $\alpha$ less than 60 degree, i.e., the region included by the black circle in the left panels of Fig."
+ 1., 1.
+ The flux densitv of the pixel with 60 < a < 90° is sel (o zero., The flux density of the pixel with $^{o}$ $\leq$ $\alpha$ $\leq$ $^{o}$ is set to zero.
+ For each smoothed aud. corrected. fulldisk magnetogram of MDI/SOLIO. we apply a magnetic. [luxdensity. ofB 15 Alx/em-L7 as a threshold to celinen (he active. regions. and (heir," For each smoothed and corrected full-disk magnetogram of MDI/SOHO, we apply a magnetic fluxdensity of 15 $^{2}$ as a threshold to define the active regions and their"
+ For each smoothed aud. corrected. fulldisk magnetogram of MDI/SOLIO. we apply a magnetic. [luxdensity. ofB 15 Alx/em-L7 as a threshold to celinen (he active. regions. and (heir.," For each smoothed and corrected full-disk magnetogram of MDI/SOHO, we apply a magnetic fluxdensity of 15 $^{2}$ as a threshold to define the active regions and their"
+Young solar-tvpe E. €. and [x stars provide proxies for the Sun's early evolution and insight into its intense pre-main sequence ancl zero-age main sequence magnetic activity.,"Young solar-type F, G, and K stars provide proxies for the Sun's early evolution and insight into its intense pre-main sequence and zero-age main sequence magnetic activity."
+ In this regard. observations of starspots and surface magnetic fieles are important. as they provide the empirical basis for understanding the stars magnetic dvnamo and its internal structure.," In this regard, observations of starspots and surface magnetic fields are important, as they provide the empirical basis for understanding the star's magnetic dynamo and its internal structure."
+ Dillerential rotation is one of the primary drivers of the dynamo process that produces magnetic fields within the Sun., Differential rotation is one of the primary drivers of the dynamo process that produces magnetic fields within the Sun.
+ In the solar convection zone. differential rotation is constant down to the base of this zone in the region known as the tachocline.," In the solar convection zone, differential rotation is constant down to the base of this zone in the region known as the tachocline."
+ Strong shear forces are formed at the interface between the solicd-body rotation of the radiative zone and the cdilferentially rotating convective laver., Strong shear forces are formed at the interface between the solid-body rotation of the radiative zone and the differentially rotating convective layer.
+ This interaction converts the poloidal magnetic field to a toroidal field., This interaction converts the poloidal magnetic field to a toroidal field.
+" This clleet is known as the ""O-ellect""", This effect is known as the $\Omega$ -effect”.
+ Llowever. on young solar-tvpe stars. the operation of the magnetic dynamo process may be occurring throughout the convective zone itself.," However, on young solar-type stars, the operation of the magnetic dynamo process may be occurring throughout the convective zone itself."
+ “Phis is known as a clistributect dynamo (Brancenburectal.1989:Mosset1995).," This is known as a distributed dynamo \citep{Brandenburg89,Moss95}."
+. Jarnesetal.(2005). have investigated the link between dillerential rotation and elfective temperature. and. hence convective turnover time.," \citet{Barnes05} have investigated the link between differential rotation and effective temperature, and hence convective turnover time."
+ They were able to fit an empirical power law to recent observations of a small sample of stars., They were able to fit an empirical power law to recent observations of a small sample of stars.
+ As the temperature increases. so does the rotational shear.," As the temperature increases, so does the rotational shear."
+ A small number of late E-/early. G-type stars have been investigated. using Doppler imaging (DI) to determine their, A small number of late F-/early G-type stars have been investigated using Doppler imaging (DI) to determine their
+The surveys of Ixasperetal.(2007) and Billeretal.(2007).. lave set constraints on the distributions of extrasolar planets similar to those we present he‘ell. while Nielsenetal.(2008) aud especially Lalreuiéreetal.(2007) have set stronger coustraints.,"The surveys of \citet{kasper} and \citet{biller1}, have set constraints on the distributions of extrasolar planets similar to those we present herein, while \citet{nielsen} and especially \citet{GDPS} have set stronger constraints."
+ Moο recent aualyses by NielseuClose(2009) aud Chauvinetal.(2010) also provide coustraints ci! the planetary distribution.," More recent analyses by \citet{nielsenclose} and \citet{chauvin}
+ also provide constraints on the planetary distribution."
+ For example. Nielsen&Close(2009) provide a confidence tha he Cumingetal.(2008) αι»ributiou can be excluded for a truncation radius of 28 AU.," For example, \citet{nielsenclose} provide a confidence that the \citet{cumming} distribution can be excluded for a truncation radius of 28 AU."
+ However. if different models are used this number jumps to 83 AU.," However, if different models are used this number jumps to 83 AU."
+ Chauvin et al., Chauvin et al.
+ indicate a similar limi rour analyzing their results usiig Baralle et al., indicate a similar limit from analyzing their results using Baraffe et al.
+ 2003 moclels., 2003 models.
+ For the staudard parameters they indicate a uaximnunm periuitted truucation radius of approximately 35 AU., For the standard parameters they indicate a maximum permitted truncation radius of approximately 35 AU.
+ In this context. the results presented here provide looser COLstraints ou the planet distribution. but provide an iudependeut. €jeck on tlie model-clepeucent SVStematic errors wxüch may exist with shorter wavelength data. due to incorrect model brightuess estμιαίος or age cleermiuation.," In this context, the results presented here provide looser constraints on the planet distribution, but provide an independent check on the model-dependent systematic errors which may exist with shorter wavelength data, due to incorrect model brightness estimates or age determination."
+ Theoretical speCtra o£ self-=uinous extrasolar planets are very poorly coustrained observationally., Theoretical spectra of self-luminous extrasolar planets are very poorly constrained observationally.
+ The recent detectiois of possible planets around HR 8799 (Maroisetal.2008).. Fomalbaut 2008).. and 3 Pic (Lagrangeetal.2009) are either single-band (2 Pic) or only beginuiug to be evaluated at inutiple wavelengths (HR 8799. Foinalhaut).," The recent detections of possible planets around HR 8799 \citep{hr8799}, Fomalhaut \citep{fomalhaut}, and $\beta$ Pic \citep{betapic} are either single-band $\beta$ Pic) or only beginning to be evaluated at multiple wavelengths (HR 8799, Fomalhaut)."
+ The caudidate planets orbiting HR S799 and ο Pic are hotter than we would expect to find orbiting wicddle-aged stars such as those iu our survey. wlile HST photometry of Fomalhaut b suggests nuucrol its |xightuess is starlight reflected Grom a circtunplanetary dust disk.," The candidate planets orbiting HR 8799 and $\beta$ Pic are hotter than we would expect to find orbiting middle-aged stars such as those in our survey, while HST photometry of Fomalhaut b suggests much of its brightness is starlight reflected from a circumplanetary dust disk."
+ Our survey. aud otler exojxlanet strveys. ust therefore be iuterpreted using models of planetary spectra that are not νο weI-testecl agalust observations.," Our survey, and other exoplanet surveys, must therefore be interpreted using models of planetary spectra that are not yet well-tested against observations."
+ Such models predict brightuesses in the A band. and particular!y il lar'ow spectral wiudows wilhin the A band. that are enormously in excess of black body |uxes.," Such models predict brightnesses in the $H$ band, and particularly in narrow spectral windows within the $H$ band, that are enormously in excess of black body fluxes."
+ Tle coustraints set by Masceladrietal.(2005):Biller(2007):LafreniereNjelsen(2008): (2009):: and €‘hauvinetal.(2010) depeud ou the accuracy oft1056 precictious of remarkable xightuess iu the A baud.," The constraints set by \citet{masciadri,biller1,GDPS,nielsen,nielsenclose}; and \citet{chauvin}
+ depend on the accuracy of these predictions of remarkable brightness in the $H$ band."
+ The £/ aud AZ bands that we have used are nearer the blackbody peaks ol low-temperature sell-Iuminous planets. aud might be expected to be more reliable.," The $L'$ and $M$ bands that we have used are nearer the blackbody peaks of low-temperature self-luminous planets, and might be expected to be more reliable."
+ However. Leggettetal.(2007) aud Reid&Cruz(2002) suggest that the AZ baud brightness least of hotter extrasolar planets will be less thau predicted by Burrowsetal.(2003) due to 1X»ve-equilibrium concentrations of CO [rom convective mixing.," However, \citet{L07} and \citet{reid} suggest that the $M$ band brightness at least of hotter extrasolar planets will be less than predicted by \citet{bur} due to above-equilibrium concentrations of CO from convective mixing."
+ Hubeuy&Burrows(2007) present leyv inocdels indi‘ating the effect is present [for planets with ranging [roi 600 to 15001. The luaxinum AL band flux supression is about1056... and flux supression disappears completely. for below 001. Based on Burrowsetal.(2003)... this value corresponds to planets of about 3.5. 6.5. 12. and 15 at ages of 100 Myr. 300 Myr. 1: Cyr. aud 2 Cyr. respectively.," \citet{NCE} present new models indicating the effect is present for planets with ranging from 600 to 1800K. The maximum $M$ band flux supression is about, and flux supression disappears completely for below 500K. Based on \citet{bur}, this value corresponds to planets of about 3.5, 6.5, 12, and 15 at ages of 100 Myr, 300 Myr, 1 Gyr, and 2 Gyr, respectively."
+ In may cases our AL band observations were sensitive to plauets at lower masses than these values. aud therefore lower than 000. implying that the CO [lux supression would have uo effect ou our iuass limits.," In many cases our $M$ band observations were sensitive to planets at lower masses than these values, and therefore lower than 500K, implying that the CO flux supression would have no effect on our mass limits."
+ In other cases our AL band seusitivity did not extend so low., In other cases our $M$ band sensitivity did not extend so low.
+ However. given that AL band. observatious formed a relatively πια. part of our survey. aud. CO supression would allect," However, given that $M$ band observations formed a relatively small part of our survey, and CO supression would affect"
+adjusting the direction of the angular momentum vectors of both the hole spin ancl the clise.,adjusting the direction of the angular momentum vectors of both the hole spin and the disc.
+ The LE elect is. very important in AGN. as the accretion episodes all initially have random directions with respect to the hole spin.," The LT effect is very important in AGN, as the accretion episodes all initially have random directions with respect to the hole spin."
+ As it has a large lever arm. it operates more quickly than the aceretion torque. which cannot ercathy change the hole spin until the hole has significantly increased its mass.," As it has a large lever arm, it operates more quickly than the accretion torque, which cannot greatly change the hole spin until the hole has significantly increased its mass."
+ Accorclinegly. the ET. elfect determines whether the disc and hole angular momenta are co or counteraligned long before the accretion torque has had a real effect.," Accordingly, the LT effect determines whether the disc and hole angular momenta are co– or counter–aligned long before the accretion torque has had a real effect."
+ Early studies of the LT. ellect. (c.g. Scheuer Feiler. 1996) concluded that it always produced coalignment.," Early studies of the LT effect (e.g. Scheuer Feiler, 1996) concluded that it always produced co–alignment."
+ Accorcingly studies of SMDII growth. (c.g. Volonteri ct al.," Accordingly studies of SMBH growth (e.g. Volonteri et al.,"
+ 2005: Volonteri Rees. 2005) argued that the main ellect of the accretion torque was always to produce spinup.," 2005; Volonteri Rees, 2005) argued that the main effect of the accretion torque was always to produce spinup."
+ Llowever Ixing et al. (, However King et al. (
+"2005: hereafter KLOP) showed on very general grounds that an initially retrograde accretion disc with total angular momentum J, would end up stably counteraligned with a hole of angular momentunm Jy provided that JJ;«24, and the initial angle 8€ between the angular momenta satisfied (Vhroughout this paper Jy.J, and the Kerr parameter α denote the absolute values of these quantities.)","2005; hereafter KLOP) showed on very general grounds that an initially retrograde accretion disc with total angular momentum $J_d$ would end up stably counter–aligned with a hole of angular momentum $J_h$ provided that $J_d < 2J_h$ and the initial angle $\theta$ between the angular momenta satisfied (Throughout this paper $J_d, J_h$ and the Kerr parameter $a$ denote the absolute values of these quantities.)"
+ Thus Scheuer Feiler’s (1996) conclusion that coalignment always. occurs depended. on their implicit assumption that the outer disc was fixed. Le. JaXJn.," Thus Scheuer Feiler's (1996) conclusion that co--alignment always occurs depended on their implicit assumption that the outer disc was fixed, i.e. $J_d \gg J_h$."
+ v contrast. WLOD’s result implies that counteralignment occurs in a fraction of accretion episodes.," By contrast, KLOP's result implies that counter–alignment occurs in a fraction of accretion episodes."
+" In their analysis. ΙΓΟΙ considered a disc of fixed total angular momentum J, misaligned with the hole."," In their analysis, KLOP considered a disc of fixed total angular momentum $J_d$ misaligned with the hole."
+ In reality any misaligned. cise must be warped at some radius i. in general rather smallerthan its outer edge.," In reality any misaligned disc must be warped at some radius $R_{\rm warp}$, in general rather smallerthan its outer edge."
+ Only the part of the disc around. Zea exerts an LP torque on the hole. so the correct interpretation of JJ; in such cases is the total angular momentum passing through £44 during the alignment process.," Only the part of the disc around $\rw$ exerts an LT torque on the hole, so the correct interpretation of $J_d$ in such cases is the total angular momentum passing through $\rw$ during the alignment process."
+ This interpretation agrees with numerical simulations of the alignment process in such cases (Locdato Pringle. 2006).," This interpretation agrees with numerical simulations of the alignment process in such cases (Lodato Pringle, 2006)."
+" Of course if the whole disc is warped JJ, is simply the total angular momentum of the disc as before.", Of course if the whole disc is warped $J_d$ is simply the total angular momentum of the disc as before.
+" Equation (2)) shows that a random sequence of episodes with Jr«—J, must have fc1/2. so that co and counteralignment are almost equally. common."," Equation \ref{hemi}) ) shows that a random sequence of episodes with $J_d \ll J_h$ must have $f \simeq 1/2$, so that co– and counter–alignment are almost equally common."
+ Since the ISCO for retrograde accretion carries higher specific angular momentum. than for prograde accretion if the herr e parameter is significant (in the ratio 11:3 for maximal spins) we see that spindown is more effective than spinup. ultimately producing a slowlyspinning hole (Llughes Dlandford. 2003 reach a similar conclusion for the effect of repeated coalescences of spinning holes. for the same reason: see also Colbert Wilson. 1995).," Since the ISCO for retrograde accretion carries higher specific angular momentum than for prograde accretion if the Kerr $a$ parameter is significant (in the ratio 11:3 for maximal spins) we see that spindown is more effective than spinup, ultimately producing a slowly–spinning hole (Hughes Blandford, 2003 reach a similar conclusion for the effect of repeated coalescences of spinning holes, for the same reason; see also Colbert Wilson, 1995)."
+ Ixing Pringle (2006) used this argument to show how suitably random accretion episodes could allow the growth of very large SAIBLE (~5.10 AL.) at high redshift +~6 from even stellar.mass seed black holes.," King Pringle (2006) used this argument to show how suitably random accretion episodes could allow the growth of very large SMBH $\sim
+5\times 10^9\, \msun$ ) at high redshift $z \sim 6$ from even stellar–mass seed black holes."
+ In this paper we show that the fact that self.&ravity cuts olf the aceretion disces in AGN feeding episodes implies Jy24 in almost all cases., In this paper we show that the fact that self–gravity cuts off the accretion discs in AGN feeding episodes implies $J_d < 2J_h$ in almost all cases.
+ A sequence of repeated episodes of this type with random orientations therefore spins the hole down to low spin parameters. and we find a mean value (o90.2.0.3.," A sequence of repeated episodes of this type with random orientations therefore spins the hole down to low spin parameters, and we find a mean value $\bar a \sim 0.2-0.3$."
+ We first require a simple description of the properties of the accretion disc during a feeding episode of the type described above., We first require a simple description of the properties of the accretion disc during a feeding episode of the type described above.
+ One can regard an evolving clisc as passing through a sequence of steady states. so we follow INKDOT in adopting the disc properties derived by CollinSoullrin Dumont (1990) in the context of AGN.," One can regard an evolving disc as passing through a sequence of steady states, so we follow KP07 in adopting the disc properties derived by Collin–Souffrin Dumont (1990) in the context of AGN."
+ These are essentially the same as those derived by Shakura Sunvaev (1973) lor steady clises in close binaries., These are essentially the same as those derived by Shakura Sunyaev (1973) for steady discs in close binaries.
+ The disc surface density is 1ος αι is the Shakura Sunvaey (1973). viscosity parameter. ¢ is the accretion cllicieney. with the luminosity L and accetion rate AZ related. by Further is the Edelineton luminosity. AM is the black hole mass. AZ. in units of 107 M... E is the radius aad is the Schwarzschild ractius ofthe central black hole.," The disc surface density is Here $\alpha_1$ is the Shakura Sunyaev (1973) viscosity parameter, $\epsilon$ is the accretion efficiency, with the luminosity $L$ and accetion rate $\dot{M}$ related by Further is the Eddington luminosity, $M_8$ is the black hole mass, $M$, in units of $10^8$ $_\odot$, $R$ is the radius and is the Schwarzschild radius of the central black hole."
+ The mass of the disc M(«HR) inside racius 2 is then ‘The disc semithickness // obeys ‘The disc becomes selfgravitating when its mass reaches (c.g. Pringle. 1981) which occurs at radii f?2 Roe. where 1n line with our discussion. in the Introduction. we identify Af as the mass AAsince dn an individual accretion episode.," The mass of the disc $M(<R)$ inside radius $R$ is then The disc semi–thickness $H$ obeys The disc becomes self–gravitating when its mass reaches (e.g. Pringle, 1981) which occurs at radii $R \ge R_{\rm sg}$ , where In line with our discussion in the Introduction we identify $M_{\rm
+ sg}$ as the mass $\dme$ in an individual accretion episode."
+ The evolution timescale of the disc. aud," The evolution timescale of the disc, and"
+According to Hale's polarity law we expect quadrants 2 and 4 to be unpopulated. while quadrants 1 and 3 should be populated in a mutually exclusive way.,"According to Hale's polarity law we expect quadrants 2 and 4 to be unpopulated, while quadrants 1 and 3 should be populated in a mutually exclusive way."
+ Thus. for a given solar cvcle. only one of the N or $ hemispheres should be populated in quadrant no.," Thus, for a given solar cycle, only one of the N or S hemispheres should be populated in quadrant no."
+ 1. while only the opposite hemisphere should be populated in quadrant no.," 1, while only the opposite hemisphere should be populated in quadrant no."
+ 3., 3.
+ As the next cevcle begins. (he hemispheres that are populated get exchange.," As the next cycle begins, the hemispheres that are populated get exchanged."
+ The pattern shown in Fie., The pattern shown in Fig.
+ 6 illustrates these properties of Hale's polarity law in a very explicit wav. while indicating that it is a law that is not strictly obeved but is fairly olten violated.," \ref{fig:butflynum} illustrates these properties of Hale's polarity law in a very explicit way, while indicating that it is a law that is not strictly obeyed but is fairly often violated."
+ The familiar active-region belts that migrate from higher to lower latitudes are Clearly seen. but only in the N hemisphere in euadrant no.," The familiar active-region belts that migrate from higher to lower latitudes are clearly seen, but only in the N hemisphere in quadrant no."
+ 1 (before 2008). while the 5 hemisphere belt occurs in equadrant no.," 1 (before 2008), while the S hemisphere belt occurs in quadrant no."
+ 3., 3.
+ After 2008 this pattern reverses as the new cvcle starts., After 2008 this pattern reverses as the new cycle starts.
+ In comparison. quadrants 2 and 4 are nearly empty. as expected from ILale's law.," In comparison, quadrants 2 and 4 are nearly empty, as expected from Hale's law."
+ The grev-scale cuts in Fig., The grey-scale cuts in Fig.
+ 6. have been chosen such that white represents zero. black elobal maximum value.," \ref{fig:butflynum} have been chosen such that white represents zero, black global maximum value."
+" While Hale's law is obeved in the great majority of cases. we nolice that quadrants 2 and 4 are not entirely empty. and Chat the ""wrong hemispheres in quadrants 1 and 3 are also weakly populated."," While Hale's law is obeyed in the great majority of cases, we notice that quadrants 2 and 4 are not entirely empty, and that the “wrong” hemispheres in quadrants 1 and 3 are also weakly populated."
+ As these apparent violations of Hale's law appear to be rather randomly distributed. (he question arises (ο what extent (hev represent noise fluctuations and errors in our automatic region idenüifications. or are real. physical violations of Hale's law.," As these apparent violations of Hale's law appear to be rather randomly distributed, the question arises to what extent they represent noise fluctuations and errors in our automatic region identifications, or are real, physical violations of Hale's law."
+ We address this question in the next section., We address this question in the next section.
+ Let us first note that the frequency of violations of Hale's polarity law increases rather steeplv as we eo to regions of smaller size., Let us first note that the frequency of violations of Hale's polarity law increases rather steeply as we go to regions of smaller size.
+ This behavior is illustrated in Fig. 7..," This behavior is illustrated in Fig. \ref{fig:wrongquad},"
+" where we have plotted the fraction of bipolar regions assigned to the ""wrong quadrants 2 and 4."," where we have plotted the fraction of bipolar regions assigned to the “wrong” quadrants 2 and 4,"
+The modeling of surface inhomogeneities in terms of spots. or caps and belts. can perhaps lead one to believe that the origin of these surface features must be directly related to the presence of a magnetic field al (he stellar surface.,"The modeling of surface inhomogeneities in terms of spots, or caps and belts, can perhaps lead one to believe that the origin of these surface features must be directly related to the presence of a magnetic field at the stellar surface."
+ In that picture. GD 323 would be a analog of Feige 7. à DAB white dwarf with a 35 MG cipolar field (Achilleosetal.1992).," In that picture, GD 323 would be a low-field analog of Feige 7, a DAB white dwarf with a 35 MG dipolar field \citep{achilleosetal92}."
+. In GD 323. the null circular polarization measurement of Angelοἱal.(1051) vields a 36 upper limit of 3 AIG on the longitudinal field component at the stellar surface.," In GD 323, the null circular polarization measurement of \citet{abl81} yields a $3\sigma$ upper limit of 3 MG on the longitudinal field component at the stellar surface."
+ However. the idea (hat surface inhomogeneities are related to other processes. such as (he convective dilution of a (hin overlving hydrogen envelope. should also be entertained.," However, the idea that surface inhomogeneities are related to other processes, such as the convective dilution of a thin overlying hydrogen envelope, should also be entertained."
+ Indeed. it has been pointed out a while back (Liebertetal.1987). Chat GD 323 is slightly cooler than the red edge of the DD gap. that peculiar region between 45.000 Ix. and 30.000 IX (hat appears to be devoid of helium-atmosphere degenerates.," Indeed, it has been pointed out a while back \citep{lfw87} that GD 323 is slightly cooler than the red edge of the DB gap, that peculiar region between 45,000 K and 30,000 K that appears to be devoid of helium-atmosphere degenerates."
+ Models of the spectral evolution of white cdwarfs (e.g..Fontaine&Wesemael1987). attempt (o account for (his gap bv. calling on the dilution of a thin hydrogen laver at the surface of a 30.000 IX. DA star with the more massive helium envelope. dilution that would lead to à DA—DD transition near that ellective temperature.," Models of the spectral evolution of white dwarfs \citep
+[e.g.,][]{fw87} attempt to account for this gap by calling on the dilution of a thin hydrogen layer at the surface of a 30,000 K DA star with the more massive helium envelope, dilution that would lead to a $\to$ DB transition near that effective temperature."
+ If GD 323 were understood in those terms. an idea that we favor. ils surface might be characterized by the presence of horizontal abundance gradients.," If GD 323 were understood in those terms, an idea that we favor, its surface might be characterized by the presence of horizontal abundance gradients."
+ Unfortunately. 1 is not possible to be more specific since a complete picture of the mixing process still eludes us al this stage.," Unfortunately, it is not possible to be more specific since a complete picture of the mixing process still eludes us at this stage."
+ The geometries considered in the exploratory investigations of Beauchamp(1993) and Ixoesteretal.(1994) appear reasonable starting points for a renewed ancl more thorough study of that. possibility., The geometries considered in the exploratory investigations of \citet{beauchampetal93} and \citet{kls94} appear reasonable starting points for a renewed and more thorough study of that possibility.
+ We thank the director and stalf of the Steward Observatory for the use of their facilities. and A. Beauchamp for his essential contribution to the caleulation of models with surface inhomogeneities.," We thank the director and staff of the Steward Observatory for the use of their facilities, and A. Beauchamp for his essential contribution to the calculation of models with surface inhomogeneities."
+" This work was supported in part bx the NSERC Canada aid by the Fund FORNT (Québbec),", This work was supported in part by the NSERC Canada and by the Fund FQRNT (Québbec).
+RS becomes strong at rnsatrongv0rv. (see Eq. (2.3))).,"RS becomes strong at $r _{RS, {\rm strong}} \sim \sigma r _{N}$ (see Eq. \ref{RS1}) ))."
+ If the outside medium is stellar wiud. strong RS forms immediately if πο... (a3 where Όρμος is the velocity of the progenitors wind iu thousands kilometers per ποσο.," If the outside medium is stellar wind, strong RS forms immediately if _w > = 220 ( where $ v_{wind, 8}$ is the velocity of the progenitors wind in thousands kilometers per second."
+" For stnaller ~,.. no RS forms ever (for weak shocks. oue should put a—1)."," For smaller $\gamma_w$, no RS forms ever (for weak shocks, one should put $\sigma \rightarrow 1$ )."
+ The reverse shock could stall (in tlie observer [frane) at (assumiug oX1) ≓↓∕∕∶≀⋡⊥∿↓∣⊽∖≓ Since μοι ds typically larger than ἐς. RS does not stall duriug the relativistic expansion phase: thus. the thin shell approximation is geuerally applicable.," The reverse shock could stall (in the observer frame) at (assuming $\sigma \gg 1$ ) = 3 Since $r_{RS,stall}$ is typically larger than $l_S$, RS does not stall during the relativistic expansion phase; thus, the thin shell approximation is generally applicable."
+" Iu the unmaguetized case. au inportant quantity is the radius when the RS crosses the ejectaEx) bLqQUPXUIIq0gUqqy, Iu the magnetized case. ra is still a good approximation for the RS crossing radius. but witli two cavities."," In the unmagnetized case, an important quantity is the radius when the RS crosses the ejecta = l_S = In the magnetized case, $r_\Delta$ is still a good approximation for the RS crossing radius, but with two cavities."
+ First. a delayed ouset of the RS. see (3.2)). delays the RS crossing moment.," First, a delayed onset of the RS, see \ref{rN}) ), delays the RS crossing moment."
+" Since ry/ry=£77, this delay is not important for £<1 (the thick shell case. generally applicable to the magnetized ejecta)."," Since $r_N /r_\Delta = \xi^{3/2}$, this delay is not important for $\xi < 1$ (the thick shell case, generally applicable to the magnetized ejecta)."
+ Also. [or subsonic outflows (see below). r isthe distance wherethe back ol the outflowcatcheswith the CD. see 83.3..," Also, for subsonic outflows (see below), $r_\Delta$ isthe distance wherethe back of the outflowcatcheswith the CD, see \ref{subsonic}. ."
+The galaxies of the sample do not show variability on a night-to-night basis. within the limits of the photometry.,"The galaxies of the sample do not show variability on a night-to-night basis, within the limits of the photometry."
+ The only ones which may not have varie at all wae 22110 and. 1103-C35., The only ones which may not have varied at all are 2110 and 103-G35.
+ Most. show smooth »«haviour on a timescale of vears. but some show shorter-erm fluctuations.," Most show smooth behaviour on a timescale of years, but some show shorter-term fluctuations."
+ The most conspicuous of the latter is 33783., The most conspicuous of the latter is 3783.
+ An active galaxy is thus more likely to be revealed by LR variability than by broad. emission. lines. which may be hidden from view by extinction within the jost galaxy.," An active galaxy is thus more likely to be revealed by IR variability than by broad emission lines, which may be hidden from view by extinction within the host galaxy."
+ As already mentioned. two very interestingὃν cases of aree amplitude short-term events are seen in 22992. which seems to have experienced: a substantial outburst (Glass. 1997a) and 77460. whose nuclear source virtually shut olf for a short. period. (Class. 1998).," As already mentioned, two very interesting cases of large amplitude short-term events are seen in 2992, which seems to have experienced a substantial outburst (Glass, 1997a) and 7469, whose nuclear source virtually shut off for a short period (Glass, 1998)."
+ Some galaxies show major changes over the long term., Some galaxies show major changes over the long term.
+ For example. the archetype S2 galaxy 11068 doubled in L [lux over about 20 vears.," For example, the archetype S2 galaxy 1068 doubled in $L$ flux over about 20 years."
+ 44748 has become fainter over 2000d and ESO3323-G7T7 over 5000d., 4748 has become fainter over 2000d and 323-G77 over 5000d.
+ δις-2-58-22 ias declined. by over a magnitude in. 5000d. though not monotonically.," MCG-2-58-22 has declined by over a magnitude in 5000d, though not monotonically."
+ Otherwise. major variations. of more than one mag. aave been seen in Fairall 9 and 5526a.," Otherwise, major variations, of more than one mag, have been seen in Fairall 9 and 526a."
+ ανα 9 shows a delay between its UV output and its £ Dux. which can να detected even between the / and. £ bands.," Fairall 9 shows a delay between its UV output and its $L$ flux, which can be detected even between the $J$ and $L$ bands."
+ This has oen satisfactorily. explained in terms of dust-reprocessing w Clavel. Wamsteker Class (1989). using à model put orward by Barvainis (1987).," This has been satisfactorily explained in terms of dust-reprocessing by Clavel, Wamsteker Glass (1989), using a model put forward by Barvainis (1987)."
+ Phe delay continues to exist in data obtained since the completion of the earlier work., The delay continues to exist in data obtained since the completion of the earlier work.
+ In many cases. the variability is most apparent at the onger wavelengths. especially £L.," In many cases, the variability is most apparent at the longer wavelengths, especially $L$."
+ This can have two causes. swamping at short wavelengths by the underlving galaxy in the 12 aresee diameter aperture (as mentioned). and circumnuclear reddening. which is almost certainly present in 11065.," This can have two causes, swamping at short wavelengths by the underlying galaxy in the 12 arcsec diameter aperture (as mentioned) and circumnuclear reddening, which is almost certainly present in 1068."
+ Some constraint on the size of the emitting regions can be obtained from the timescale on which substantial variations can occur., Some constraint on the size of the emitting regions can be obtained from the timescale on which substantial variations can occur.
+ Loa dust re-emission model is accepted. 1e material may be distributed at dillerent. distances from. 1e nucleus of the active galaxy and at arbitrary angles.," If a dust re-emission model is accepted, the material may be distributed at different distances from the nucleus of the active galaxy and at arbitrary angles."
+ Thus. the infrared. response to an instantaneous change in re ultraviolet. e.g... Bea(f)x 9600] will usually be smeared ut.," Thus, the infrared response to an instantaneous change in the ultraviolet [e.g., $F_{\nu \rm{uv}}(t) \propto \delta(t)$ ] will usually be smeared out."
+ Only in the case of a ring of dust. perpendicular to 1 line of sight can a very sharp response be expected (Barvainis. 1992).," Only in the case of a ring of dust perpendicular to the line of sight can a very sharp response be expected (Barvainis, 1992)."
+ A number of the galaxies in the present paper have been monitored at shorter wavelengths for various lengths of time and some of this work is included in Fig 2., A number of the galaxies in the present paper have been monitored at shorter wavelengths for various lengths of time and some of this work is included in Fig 2.
+ Phe variations in U tend to have higher amplitude and. possibly shorter time scales than those at infrared. wavelengths., The variations in $U$ tend to have higher amplitude and possibly shorter time scales than those at infrared wavelengths.
+ Delavs between the short-wavelength light. eurve(s). thought to be indicative of the UV. Dux from the central engine. and the £L [lux originating in the dust that it is heating. have been looked for using a simple cross-correlation programme based on the interpolation method of Gaskell Peterson (1987).," Delays between the short-wavelength light curve(s), thought to be indicative of the UV flux from the central engine, and the $L$ flux originating in the dust that it is heating, have been looked for using a simple cross-correlation programme based on the interpolation method of Gaskell Peterson (1987)."
+ LO should be emphasized that there are many opinions as to the best method for finding delays in cases where the data are ireeularly spaced., It should be emphasized that there are many opinions as to the best method for finding delays in cases where the data are irregularly spaced.
+ A more detailed analvsis will be carried out in a future paper., A more detailed analysis will be carried out in a future paper.
+ The Gaskell Peterson method was applied to those ealaxies with significant variability and many data points., The Gaskell Peterson method was applied to those galaxies with significant variability and many data points.
+ Ideally. the cata should have been sampled at. regular intervals. shorter than the expected. delay times. for this procedure to be reliable.," Ideally, the data should have been sampled at regular intervals, shorter than the expected delay times, for this procedure to be reliable."
+ Phe programme first. interpolates the longest wavelength. photometry to give values for each day as well as extrapolating its starting and stopping values as constants for 1000. cays before. and. after the series., The programme first interpolates the longest wavelength photometry to give values for each day as well as extrapolating its starting and stopping values as constants for 1000 days before and after the series.
+ Each light curve is shifted. by a constant to have average value O and divided so as to have rms deviation 1., Each light curve is shifted by a constant to have average value 0 and divided so as to have rms deviation 1.
+ The Cposs-correlation is performed for the dates of the shorter-wavelength observations. using à rangeν of 1000 days in cach direction.," The cross-correlation is performed for the dates of the shorter-wavelength observations, using a range of 1000 days in each direction."
+ The correlogram is cliviclb by its maximunr value so that the maximum: of the peak is 1.0., The correlogram is divided by its maximum value so that the maximum of the peak is 1.0.
+ Some of the causes that mav distort the correlogram. are (a) insullicient sampling (b) long gaps in the data. particularly if these coincide with enhanced activity (deviation from [lat behaviour) (c) closely-clusterecl measurements. which cause over-weightine of certain sections of the non-interpolated sequences (d) an amplitude of variation that is not. much ereater than the noise. such as occurs at the J wavelength," Some of the causes that may distort the correlogram are (a) insufficient sampling (b) long gaps in the data, particularly if these coincide with enhanced activity (deviation from flat behaviour) (c) closely-clustered measurements, which cause over-weighting of certain sections of the non-interpolated sequences (d) an amplitude of variation that is not much greater than the noise, such as occurs at the $J$ wavelength"
+cold classical Kuiper belt objects.,cold classical Kuiper belt objects.
+ While all of tlose possibilities are supported 1o» the current laboraory data. ouly experimental verification will allow supyort or refutation of this hypothesis for explaining the colors of objects in the Ixuiper belt.," While all of these possibilities are supported by the current laboratory data, only experimental verification will allow support or refutation of this hypothesis for explaining the colors of objects in the Kuiper belt."
+ This research has been. supported bv. grant NNNOQABLOC. from the NASA Planetary Astronomy program., This research has been supported by grant NNX09AB49G from the NASA Planetary Astronomy program.
+ We thank Tal Levisou. Rosario Brunetto. aud Ikoustautiu Datveiu for euliehteniug conversations.," We thank Hal Levison, Rosario Brunetto, and Konstantin Batygin for enlightening conversations."
+"renioved. leaving only Poisson noise (74s;=fexp). or because the PSF structure was already uncorrelated between consecutive images in the first place (Tied,Xlexp=Tdeorrleap).","removed, leaving only Poisson noise $\tau_{\rm{dcorr}} = t_{\rm{exp}}$ ), or because the PSF structure was already uncorrelated between consecutive images in the first place $\tau_{\rm{speck}} \le t_{\rm{exp}} \Rightarrow \tau_{\rm{dcorr}} = t_{\rm{exp}}$ )."
+ In this case. the speckle attenuation increases with the square root of the total number of image differences.," In this case, the speckle attenuation increases with the square root of the total number of image differences."
+" The decorrelation (üimescale is simply In the second regime. since ο 15 longer than /,4 but shorter than 7. the residuals of consecutive differences are partially correlated and do not average-oul as efficiently."," The decorrelation timescale is simply In the second regime, since $\tau_{\rm{speck}}$ is longer than $t_{\rm{exp}}$ but shorter than $\tau$, the residuals of consecutive differences are partially correlated and do not average-out as efficiently."
+ Assuming that only the noise having a spatial scale ~1 PSF EWIIM limits point source detection. the time needed to decorrelate (his noise is the shortest time between Τι. the time for a one EWIIM FOV rotation. and 74.," Assuming that only the noise having a spatial scale $\sim 1$ PSF FWHM limits point source detection, the time needed to decorrelate this noise is the shortest time between $\tau_{\rm{FWHM}}$, the time for a one FWHM FOV rotation, and $\tau_{\rm{speck}}$."
+ The decorrelation timescale is thus We note that if 7 is several minutes and if individual ADI differences are dominated bv random short-lived speckles. then (he noise attenuation can [follow Eq.," The decorrelation timescale is thus We note that if $\tau$ is several minutes and if individual ADI differences are dominated by random short-lived speckles, then the noise attenuation can follow Eq."
+ - [for several images bul ultimately converges to Eq., \ref{eq1} for several images but ultimately converges to Eq.
+ Hr5) when sullicient averaging of short-lived speckles has occurred., \ref{eq2} when sufficient averaging of short-lived speckles has occurred.
+" It is important to known that reducing /;4, to increase the number of images for a given inge sequence of length / would not result in better detection limits.", It is important to known that reducing $t_{\rm{exp}}$ to increase the number of images for a given image sequence of length $t$ would not result in better detection limits.
+ For the first regime. shorter /;; would mean a smaller companion 5/N per image: the final 5/N would thus be (he same (neglecting losses due to overheads).," For the first regime, shorter $t_{\rm{exp}}$ would mean a smaller companion S/N per image; the final S/N would thus be the same (neglecting losses due to overheads)."
+ For the second regime. a shorter fe. would veild quasi-static speckles correlated over more consecutive images. also resulting in the same performance.," For the second regime, a shorter $t_{\rm{exp}}$ would yeild quasi-static speckles correlated over more consecutive images, also resulting in the same performance."
+ We emphasize that ADI guarantees a gain in detection with increasing observing time for both regimes., We emphasize that ADI guarantees a gain in detection with increasing observing time for both regimes.
+ In the worst case. the speckle attenuation efficiency is limited bv field rotation.," In the worst case, the speckle attenuation efficiency is limited by field rotation."
+ This is an enormous advantage over Classical observation in which the quasi-static aberrations prevent significant gain after a relatively short observing time., This is an enormous advantage over classical observation in which the quasi-static aberrations prevent significant gain after a relatively short observing time.
+ An equation similar to Eq., An equation similar to Eq.
+ 5. also applies to classical observations. but in this case (here is no Twy since the field is not rotating.," \ref{eq2} also applies to classical observations, but in this case there is no $\tau_{\rm{FWHM}}$ since the field is not rotating."
+" Speckle attenuation is then governed by z,,44. Which can be verv laree. reducing drastically the efficiency of the observations."," Speckle attenuation is then governed by $\tau_{\rm{speck}}$, which can be very large, reducing drastically the efficiency of the observations."
+ Section G.4 presents a comparison of ADI and classical observations., Section \ref{adivscla} presents a comparison of ADI and classical observations.
+where E(z) describes the evolution of the Hubble parameter as a function of ther,where $E(z)$ describes the evolution of the Hubble parameter as a function of the.
+"edshift! In the definition of the comoving distance.. there is the implicit hypothesis that peculiar velocities give a negligible contribution to the measured redshift in the fforest (Rauchetal.2005,PaperI)..", In the definition of the comoving distance there is the implicit hypothesis that peculiar velocities give a negligible contribution to the measured redshift in the forest \citep[][Paper I]{rauch05}.
+" The Cross CF of the transmitted flux between two lines of sight at angular separation A is defined as where, Ar=MRro—2ryirj;2cosA0) is the spatial separation between pixel 1 at Τι along one aand pixel 2 at rj. along the pairedsight."," The Cross CF of the transmitted flux between two lines of sight at angular separation $\Delta\theta$ is defined as where, $\Delta r = \sqrt{(r_{\pa,1}^2 + r_{\pa,2}^2 - 2\,
+r_{\pa,1}\,r_{\pa,2}\,\cos\Delta\theta)}$ is the spatial separation between pixel 1 at $r_{\pa,1}$ along one and pixel 2 at $r_{\pa,2}$ along the paired."
+. We also computed the Cross CF for the sample of mock spectra., We also computed the Cross CF for the sample of mock spectra.
+ The simulated spectra are characterized by the redshift of the output box and a velocity extent., The simulated spectra are characterized by the redshift of the output box and a velocity extent.
+" In order to assign a redshift value to every pixel, we gave the central pixel of every spectrum the redshift of the corresponding output box, then we numbered the pixels one by one, transforming the velocity size of the pixel into a redshift size."," In order to assign a redshift value to every pixel, we gave the central pixel of every spectrum the redshift of the corresponding output box, then we numbered the pixels one by one, transforming the velocity size of the pixel into a redshift size."
+" Once the redshifts were determined pixel by pixel, we followed the same procedure adopted for the observed spectra for the 50 simulated samples and computed the average Cross CF and its 1 o standard deviation."," Once the redshifts were determined pixel by pixel, we followed the same procedure adopted for the observed spectra for the 50 simulated samples and computed the average Cross CF and its 1 $\sigma$ standard deviation."
+ The result of our computation is shown in Fig., The result of our computation is shown in Fig.
+ 3 compared with the observed Cross CF of the pairs., \ref{fig:cross_CF} compared with the observed Cross CF of the pairs.
+ The two functions show very good agreement (at the 1 c level) and significant clustering signal up to comoving separations of ~a4h! comoving Mpc., The two functions show very good agreement (at the 1 $\sigma$ level) and a significant clustering signal up to comoving separations of $\sim 4\ h^{-1}$ comoving Mpc.
+ We also compared the Cross CF with the previously computed Auto CF., We also compared the Cross CF with the previously computed Auto CF.
+ As can be seen in Fig., As can be seen in Fig.
+" 4 the two data series show a consistency at the 1 σ level, supporting the hypothesis that the peculiar velocity along the iis negligible."," \ref{fig:auto-cross} the two data series show a consistency at the 1 $\sigma$ level, supporting the hypothesis that the peculiar velocity along the is negligible."
+" A measure of the transverse clustering properties of the IGM which is less affected by peculiar velocities is the set of the flux Cross correlation coefficients (CCC), where, every pixel along one iis correlated with the one face-to-face in redshift space along the paired aand the result is averaged over all the pixels in the common redshift interval."," A measure of the transverse clustering properties of the IGM which is less affected by peculiar velocities is the set of the flux Cross correlation coefficients (CCC), where, every pixel along one is correlated with the one face-to-face in redshift space along the paired and the result is averaged over all the pixels in the common redshift interval."
+" Every pair of QSOs at angular separation Δθ gives one value of xl(A0), and a sample with several pairs at different separations, as in our sample, gives an estimate of the correlation function."," Every pair of QSOs at angular separation $\Delta\theta$ gives one value of $\chi^f_{\times}(\Delta\theta)$, and a sample with several pairs at different separations, as in our sample, gives an estimate of the correlation function."
+" At a given redshift, the angular separation A0 corresponds to a velocity separation Av,=cF(z) A0, where c denotes the speed of light, and F(z) is a dimensionless function of redshift that includes all the dependence on the global cosmological metric."," At a given redshift, the angular separation $\Delta\theta$ corresponds to a velocity separation $\Delta v_{\perp}=c\,F(z)\,\Delta\theta$ , where $c$ denotes the speed of light, and $F(z)$ is a dimensionless function of redshift that includes all the dependence on the global cosmological metric."
+ In the cosmological model that we have adopted: where z is the mean redshift of the overlapping rregion from a pair., In the cosmological model that we have adopted: where $z$ is the mean redshift of the overlapping region from a pair.
+ We computed the range of velocity separations covered by each of our pairs of spectra then we grouped the pairs in velocity bins of variable width and computed the average CCC for every group., We computed the range of velocity separations covered by each of our pairs of spectra then we grouped the pairs in velocity bins of variable width and computed the average CCC for every group.
+ Given the small number of pairs in each group (a maximum of three QSO pairs) the uncertainties associated with these determinations can be only computed by applying a simple error propagation., Given the small number of pairs in each group (a maximum of three QSO pairs) the uncertainties associated with these determinations can be only computed by applying a simple error propagation.
+" However, these uncertainties would account only for the pixel statistics and the noise"," However, these uncertainties would account only for the pixel statistics and the noise"
+"Comparisous between VIRGO/SPM data. SIM alcl our model are shown in ο, ","Comparisons between the VIRGO/SPM data, SIM and our model are shown in Fig. \ref{fig:VIR_comp}."
+"We find that a ΠΙΟΣ: of «lavs stand out in all filers,", We find that a number of days stand out in all filters.
+" As already fouud for the TSI. the strongest discrepaicv regarding SIM data occurs aronud June 25. just after the Jne suuspot massage,"," As already found for the TSI, the strongest discrepancy regarding SIM data occurs around June 25, just after the June sunspot passage."
+ SIM apyxweutlv fails to pick up the facular brieleniis ax the actIve region is near the liub: this is particularly salieut iu the blue and ereen filters., SIM apparently fails to pick up the facular brightening as the active region is near the limb; this is particularly salient in the blue and green filters.
+ Note that the red filters nuryor sole of he problems seen in Sec. Ll. παλιοί’," Note that the red filters mirror some of the problems seen in Sec. \ref{sec:tsi_comps},"
+ au chhancement i nid May just before the first spot eroup appears aand a rise at the eud of .hue before f1C data gap., namely an enhancement in mid May just before the first spot group appears and a rise at the end of June before the data gap.
+ T1050 nro evel nore pronounced im the original red filter (¢ata not shown here) aud we thus conclule that they are largelyo duc o an incomplete removal of t lustreital tempecratire chauges that affect the longer waveleugh regions particularVv strorely., These are even more pronounced in the original red filter (data not shown here) and we thus conclude that they are largely due to an incomplete removal of the instrumental temperature changes that affect the longer wavelength regions particularly strongly.
+ The largest difference beween the model aud bo SIN [ane VIRGO data is the ach larger facular xiehteniug before aud after the Ju vosunspot passageo., The largest difference between the model and both SIM and VIRGO data is the much larger facular brightening before and after the July sunspot passage.
+"""0 This is verv noticeable 1n all three fiters. as indeed aSO or the TSI (see Sec. 1.) )."," This is very noticeable in all three filters, as indeed also for the TSI (see Sec. \ref{sec:tsi_comps}) )."
+ By contrast. we find that the sunuspot «arkening measured by VIRGO and SIM agrees very well with the iodeIs. iu particuar iu the ereen axd ue filters: in the red heuid there is a slight tendency for he darkening to be overestimated.," By contrast, we find that the sunspot darkening measured by VIRGO and SIM agrees very well with the models, in particular in the green and blue filters; in the red band there is a slight tendency for the darkening to be overestimated."
+ Au exception to the eood fit is the very small spot at the eid of Alay where t1ο node nuderestimates the sunspot darkening in all filters., An exception to the good fit is the very small spot at the end of May where the model underestimates the sunspot darkening in all filters.
+ Oie wav to judee the tightness of the correlation )otwoee1 two data sets is to consider listograms of their SEfknal differences as shown in Fie. 5.., One way to judge the tightness of the correlation between two data sets is to consider histograms of their fractional differences as shown in Fig. \ref{fig:histos}.
+" Because of the res]20ise problems of SIM/VISI for wavelength iu excess of:ixmut 820 112. we have used the short red filter (reda) o obain the fluxes for the SIA-to-mmodel comALISO,"," Because of the response problems of SIM/VIS1 for wavelength in excess of about 820 nm, we have used the short red filter $_S$ ) to obtain the fluxes for the SIM-to-model comparisons."
+ The original ceutral waveleneth was used for the mode-to-VIRGO comparison aud the two ciffereut filters are used when comparing VIRGO to SI., The original central wavelength was used for the model-to-VIRGO comparison and the two different filters are used when comparing VIRGO to SIM.
+" While he imber of data poiuts considered here is relaively stall. we fid hat uost of the histograms resenible Gaussians, some oftweld with noticeable skew."," While the number of data points considered here is relatively small, we find that most of the histograms resemble Gaussians, some of them with noticeable skew."
+ The hugest deviatioIs are seen for he blue filter for the SIM vs SATIRE comyarison. Where he distribution is very broad anc may be double-peaked.," The largest deviations are seen for the blue filter for the SIM vs SATIRE comparison, where the distribution is very broad and may be double-peaked."
+ Iu most of the cases. however. we can use the width of he standard deviation of a Gaussian fit O he histograms Qtcharacterise the scatter betwee ithe different data sets.," In most of the cases, however, we can use the width of the standard deviation of a Gaussian fit to the histograms to characterise the scatter between the different data sets."
+ These are listed in Tab. 3.., These are listed in Tab. \ref{tab:virgo_correls}.
+ Tn order to quantiv the fits further. we Hist f correlation coefficients in. Tab. 3..," In order to quantify the fits further, we list the correlation coefficients in Tab. \ref{tab:virgo_correls}."
+ To calculate t correlations. we binned he SORCTE/SIM and model data Olo the same grid as the daily VIRGO/SPM results.," To calculate the correlations, we binned the SORCE/SIM and model data onto the same grid as the daily VIRGO/SPM results."
+ Fig., Fig.
+o 6 shows the correlaion plots of the VIRGO/SPA aud the SIM filter observatious with respect to the mioe calculations., \ref{fig:VIR_corr} shows the correlation plots of the VIRGO/SPM and the SIM filter observations with respect to the model calculations.
+ The red aud black lines show the hest-fi lineS assunniue that all data sets suffer from equal relative erexs.," The red and black lines show the best-fit lines, assuming that all data sets suffer from equal relative errors."
+ These errors were estimated to be of the order of 10! ppia in the ereen aAC red filter. aud of the order of Ls! pp iu the blue filter (see also Tab. 33).," These errors were estimated to be of the order of 100 ppm in the green and red filter, and of the order of 180 ppm in the blue filter (see also Tab. \ref{tab:virgo_correls}) )."
+ The das1ος blue line indicates a slope of uuitv as would be expected for a perfect natch., The dashed blue line indicates a slope of unity as would be expected for a perfect match.
+ When considering all flcrs together. the correlation cocficients and histogram widths sugeest tha the best aereclent is ound betwee‘ll he VIRGO data and SATIRE caleulations. while comparisous fare least we1 for SIM versus SATIRE.," When considering all filters together, the correlation coefficients and histogram widths suggest that the best agreement is found between the VIRGO data and SATIRE calculations, while comparisons fare least well for SIM versus SATIRE."
+ The correlation cocfficieuts aud plots indicate that the agreenmienu between the model aud the observations is best for the erecn filter. but sightly less eood for both the red aud bue filters.," The correlation coefficients and plots indicate that the agreement between the model and the observations is best for the green filter, but slightly less good for both the red and blue filters."
+ Iu the folowing. we discuss the fits for the individual filters iu more detail.," In the following, we discuss the fits for the individual filters in more detail."
+ The correlation cocfiicicuts iu fιο red filer show a relatively large variation. raneine froni a tielt correlalon with cocficicut 0.96 G7=0.90) beween VIRGO data ancl model caleuations. down to a cocfiicicut of O.77 (7220.59) for SIAre to model comparisois.," The correlation coefficients in the red filter show a relatively large variation, ranging from a tight correlation with coefficient 0.96 $r^2$ =0.90) between VIRGO data and model calculations, down to a coefficient of 0.77 $r^2$ =0.59) for SIM/red to model comparisons."
+ This latter value can larecly he οςdlained by the eniperaturednuduced seusitivitv probleus {κx STAL/VISL above about 850 lii., This latter value can largely be explained by the temperature-induced sensitivity problems for SIM/VIS1 above about 850 nm.
+ A comparison between the model aud. SIA data for the slightly shorter reda filter where this is less of am issue eives a significantly higher correlation cocticieut of 0.92 D (r-081)., A comparison between the model and SIM data for the slightly shorter $_S$ filter where this is less of an issue gives a significantly higher correlation coefficient of 0.92 $r^2$ =0.84).
+* Despite.B t1C filtera shiftκ. aud 1C associated change im responsivitv to spot and facular passages. there is also aui ducrease in the correlation coefficieut (from 0.89 o 0.96) when comparing VIRGO red with reds.," Despite the filter shift and the associated change in responsivity to spot and facular passages, there is also an increase in the correlation coefficient (from 0.89 to 0.96) when comparing VIRGO red with $_S$."
+ Tn terms of the correlation coeficicut. there is no significant difference between using the nioel calculations or the red or reds filter to compare with a ei3 observed ine series.," In terms of the correlation coefficient, there is no significant difference between using the model calculations for the red or $_S$ filter to compare with a given observed time series."
+ This is expected. as the basic features of the nodel dier only very sligtly betweei the two waveleneth xus.," This is expected, as the basic features of the model differ only very slightly between the two wavelength bands."
+ TI16 züuplitude of the variabiliv. however. mate1ο uuch beter when the uxxlel calculalons are carried out or the fer appropriate or the comxurlson data.," The amplitude of the variability, however, matches much better when the model calculations are carried out for the filter appropriate for the comparison data."
+ The sopes derived for the red correlations show tlat SATIRE overestimates je amplitude of the variahiitv w abou (SIM) ux (VIRGO)., The slopes derived for the red correlations show that SATIRE overestimates the amplitude of the variability by about (SIM) and (VIRGO).
+ For the SIM analysis. we used the altcnative short red filter as this is essentially unaffected w the temperature effects.," For the SIM analysis, we used the alternative `short' red filter as this is essentially unaffected by the temperature effects."
+" Note hat the xory suuall exaclieit of 088 derived for the VIRGO vs 1nodel conrparisolns ds in part due to a sinele outlviug nodel 1xiut on June 2] shown as the red triangle at (0.99935, 0.99975) in the bottom plot of Fie."," Note that the very small gradient of 0.88 derived for the VIRGO vs model comparisons is in part due to a single outlying model point on June 21, shown as the red triangle at (0.99935, 0.99975) in the bottom plot of Fig."
+ 6 (sce also Fie. D)., \ref{fig:VIR_corr} (see also Fig. \ref{fig:VIR_comp}) ).
+ Tf this outher is excluded. the eradieut increases o 0.91 which is in better agreeimeut with the results for SIM.," If this outlier is excluded, the gradient increases to 0.91 which is in better agreement with the results for SIM."
+ A further reason for the low exadieut could be that he red filter baud includes the recdimost line of the Ca IR triplet that might not be modeled well by SATIRE., A further reason for the low gradient could be that the red filter band includes the redmost line of the Ca IR triplet that might not be modelled well by SATIRE.
+ The correlations between all the data sets are excelleut in the ereen filter., The correlations between all the data sets are excellent in the green filter.
+ In fact. they are better han those cleenuined for the TSI comparisons. τοπια]Mv because they are limited to the similar narrow προςral region and therefore sample a welbdefiuied region of the solar atnosphere.," In fact, they are better than those determined for the TSI comparisons, presumably because they are limited to the similar narrow spectral region and therefore sample a well-definied region of the solar atmosphere."
+ We also find that the eracicuts of the best-fit lues are near unity. indicating that the ampliude of the variabilitv agrees between all three datasets.," We also find that the gradients of the best-fit lines are near unity, indicating that the amplitude of the variability agrees between all three datasets."
+ In the blue filcr. the agreement between the model calcuations ad," In the blue filter, the agreement between the model calculations and"
+in ELH03) double-peaked line emitters in 503 and EIIO3.,in EH03) double-peaked line emitters in S03 and EH03.
+" Under the prescription οἱ elliptical disk models. the ellipticities of both objects are (vpical for. double-peaked line emitters,"," Under the prescription of elliptical disk models, the ellipticities of both objects are typical for double-peaked line emitters."
+" The optical luminosities [vL,(5100A)=1.9xI0!'erg for SDSS J09422-0000. and L.5x10Perg for SDSS J1417--6141]. and the Balmer decrements are well within the range for line emitters in the SDSS sample."," The optical luminosities $\nu L_\nu(5100{\mathrm \AA})=1.9\times
+10^{44}$ for SDSS J0942+0900, and $\times 10^{43}$ for SDSS J1417+6141], and the Balmer decrements are well within the range for double-peaked line emitters in the SDSS sample."
+ By comparing the UV line profiles of Arp 102D with those of the Balmer lines. Halpern et al. (," By comparing the UV line profiles of Arp 102B with those of the Balmer lines, Halpern et al. ("
+1996) found that the contribution of a disk component to CHI]. CIV. and Lyo emission is neelieible. suggesting that high-1onization lines are produced in a dillerent region: aud that an additional broad. non-disk component is required to fit the Ha and IL? line profile in Arp 102D. But it is hard (o distinguish an additional broad. component Irom the asvinmetrical disk model because of the closeness of the two peaks and their blending with narrow lines.,"1996) found that the contribution of a disk component to CIII], CIV, and $\alpha$ emission is negligible, suggesting that high-ionization lines are produced in a different region; and that an additional broad, non-disk component is required to fit the $\alpha$ and $\beta$ line profile in Arp 102B. But it is hard to distinguish an additional broad component from the asymmetrical disk model because of the closeness of the two peaks and their blending with narrow lines."
+ In SDSS JO942+0900 and SDSS JI4I7+6141. the separation between the two peaks of the lla line is much larger. which makes it easier (to identify anv additional component.," In SDSS J0942+0900 and SDSS J1417+6141, the separation between the two peaks of the $\alpha$ line is much larger, which makes it easier to identify any additional component."
+ In both objects. (he normal broad component exists apparently albeit weak.," In both objects, the normal broad component exists apparently albeit weak."
+ With FWIIM being around 3300—6000 kins 4. this component cannot be the broad wing of the narrow lines. and should come [rom a different. line-emitting region (e.g.. the ionized wind).," With FWHM being around $-$ 6000 km $^{-1}$, this component cannot be the broad wing of the narrow lines, and should come from a different line-emitting region (e.g., the ionized wind)."
+ The EW of {his component is only  (rest. frame)., The EW of this component is only $\simeq$ (rest frame).
+ since only (wo are known. extreme double-peaked line emitters are certainly very rare.," Since only two are known, extreme double-peaked line emitters are certainly very rare."
+ llowever. we note that such objects might be not as rare as thev appear since Che detection of such objects is technically difficult.," However, we note that such objects might be not as rare as they appear since the detection of such objects is technically difficult."
+ First. such a line can only be detected in spectra with relative high S/N ratio because of its low height.," First, such a line can only be detected in spectra with relative high S/N ratio because of its low height."
+ Since the IL? line is a factor of three weaker. ilis even more difficult to find them if Ha is shifted out of the spectral coverage.," Since the $\beta$ line is a factor of three weaker, it is even more difficult to find them if $\alpha$ is shifted out of the spectral coverage."
+ Second. the SDSS pipeline is not designed to detect a line as broad as this.," Second, the SDSS pipeline is not designed to detect a line as broad as this."
+ In fact. the very broad lines in both SDSS J0942--0900 and SDSS JI4172-6142 were taken as part of the continuum by the SDSS pipeline. and only the narrow cores of the lines were measured.," In fact, the very broad lines in both SDSS J0942+0900 and SDSS J1417+6142 were taken as part of the continuum by the SDSS pipeline, and only the narrow cores of the lines were measured."
+ Selections based on the line parameters given bv the SDSS pipeline will miss such objects., Selections based on the line parameters given by the SDSS pipeline will miss such objects.
+ A further problem is that when the contribution of a stellar continuum is significant. anv weak broad peaks mav be mixed with stellar feature.," A further problem is that when the contribution of a stellar continuum is significant, any weak broad peaks may be mixed with stellar feature."
+ Proper subtraction of stellar light with the presence of such broad lines is challenging., Proper subtraction of stellar light with the presence of such broad lines is challenging.
+ Combining all these [actors together. {here may be more such objects still to be sought out in the SDSS spectroscopic sample.," Combining all these factors together, there may be more such objects still to be sought out in the SDSS spectroscopic sample."
+ llistoricallv. double-peaked lines were also interpreted as emission from a binary black," Historically, double-peaked lines were also interpreted as emission from a binary black"
+ One of the major goals of cosmology has been to measure accurately the cosmological deceleration parameter go.,] 	 One of the major goals of cosmology has been to measure accurately the cosmological deceleration parameter $q_0$.
+ Determining go from the count magnitude relation for galaxies is notoriously difficult due to evolutionary effects (e.g.. Sandage 1988).," Determining $q_0$ from the count magnitude relation for galaxies is notoriously difficult due to evolutionary effects (e.g., Sandage 1988)."
+ However. recent studies of galaxy counts from redshift surveys provide reasonable constraints on do.," However, recent studies of galaxy counts from redshift surveys provide reasonable constraints on $q_0$."
+ Schuecker et al. (, Schuecker et al. (
+1998) find a best fit of go = 0.1 with an upper limit of go< 0.75 at confidence assuming A=0.,1998) find a best fit of $q_0$ = 0.1 with an upper limit of $q_0 <$ 0.75 at confidence assuming $\Lambda = 0$.
+ Studies of high-redshift Type Ia supernovae (Perlmutter al., Studies of high-redshift Type Ia supernovae (Perlmutter et al.
+ 1997. Garnavich et al.," 1997, Garnavich et al."
+ 1998. Riess et al.," 1998, Riess et al."
+ 1998) suggest that etgy may be negative due to a non-zero cosmological constant., 1998) suggest that $q_0$ may be negative due to a non-zero cosmological constant.
+ Finally. a study of gravitational lensing statistics suggests that qo7—2 (Park Gott 1998).," Finally, a study of gravitational lensing statistics suggests that $q_0 > -2$ (Park Gott 1998)."
+ These constraints do not. however. allow us to rule out any of the most popular cosmological models (e.g.. closed universes with qo>0.5 or OcA=1 universes with qo~ —0.5).," These constraints do not, however, allow us to rule out any of the most popular cosmological models (e.g., closed universes with $q_0 > 0.5$ or $\Omega + \Lambda = 1$ universes with $q_0 \sim
+-0.5$ )."
+ It has been proposed that a constant Ly—T relation would allow cosmological parameters to be measured (Henry Tucker 1979)., It has been proposed that a constant $L_X-T$ relation would allow cosmological parameters to be measured (Henry Tucker 1979).
+ Another interesting method for improving or reinforcing these constraints is to analyze the gas mass fractions of rich galaxy clusters., Another interesting method for improving or reinforcing these constraints is to analyze the gas mass fractions of rich galaxy clusters.
+ Since this gas mass fraction ts likely to be constant and its calculated value depends on cosmological parameters. comparisons of gas mass fractions of nearby and distant clusters can constrain go (Sasaki 1996. Pen 1997).," Since this gas mass fraction is likely to be constant and its calculated value depends on cosmological parameters, comparisons of gas mass fractions of nearby and distant clusters can constrain $q_0$ (Sasaki 1996, Pen 1997)."
+ In this paper. we test the possibility of using the constancy of the gas mass fraction to constrain go.," In this paper, we test the possibility of using the constancy of the gas mass fraction to constrain $q_0$."
+ We use the available data with suitable analyses to evaluate plausible uncertainties to determine 1f the method is sufficiently sensitive to allow accurate measurements of go using future and observations of large cluster samples., We use the available data with suitable analyses to evaluate plausible uncertainties to determine if the method is sufficiently sensitive to allow accurate measurements of $q_0$ using future and observations of large cluster samples.
+ The approach we follow to constrain go requires that the mean cluster gas mass fraction be a universal constant. independent of redshift.," The approach we follow to constrain $q_0$ requires that the mean cluster gas mass fraction be a universal constant, independent of redshift."
+ Rich clusters have masses 10'°M... which suggests that the matter they contain ought to be a fair sample of the distribution of baryons and dark matter in the universe.," Rich clusters have masses $\sim 10^{15}
+M_\odot$, which suggests that the matter they contain ought to be a fair sample of the distribution of baryons and dark matter in the universe."
+ Detatled analyses for hot. massive clusters (White et al.," Detailed analyses for hot, massive clusters (White et al."
+ 1993. Evrard. Metzler. Navarro 1996; hereafter EMN) show that there is no known mechanism to significantly concentrate baryons (both galaxies and X-ray gas) in the center of the cluster or to expel them from the potential.," 1993, Evrard, Metzler, Navarro 1996; hereafter EMN) show that there is no known mechanism to significantly concentrate baryons (both galaxies and X-ray gas) in the center of the cluster or to expel them from the potential."
+ This is supported by numerical simulations (e.g.. White et al.," This is supported by numerical simulations (e.g., White et al."
+ 1993. Lubin et al.," 1993, Lubin et al."
+" 1996). which show that the X-ray gas mass fraction f,2M,/M,,, averaged within ~| Mpe should be constant. independent of redshift or evolution (Danos Pen 1998. Frenk et al."," 1996), which show that the X-ray gas mass fraction $f_g = M_{gas}/M_{tot}$ averaged within $\sim 1$ Mpc should be constant, independent of redshift or evolution (Danos Pen 1998, Frenk et al."
+ 1998)., 1998).
+ In a flat universe. there is a degeneracy between redshift and σε. or equivalently. epoch and mass.," In a flat universe, there is a degeneracy between redshift and $\sigma_8$, or equivalently, epoch and mass."
+ Evrard (1997) has shown that the gas mass fraction is independent of total cluster mass to an accuracy of1., Evrard (1997) has shown that the gas mass fraction is independent of total cluster mass to an accuracy of.
+5%.. This ts equivalent to an independence of the gas mass fraction in redshift to this accuracy., This is equivalent to an independence of the gas mass fraction in redshift to this accuracy.
+ Alternatively. in an extremely open universe. clusters formation ceases early in the evolution of the universe.," Alternatively, in an extremely open universe, clusters formation ceases early in the evolution of the universe."
+ Afterwards. there is little evolution imposed on clusters by their environments.," Afterwards, there is little evolution imposed on clusters by their environments."
+" Most simulations (Thomas Couchmann 1992, Katz. Hernquist. Weinberg 1992, Cen Ostriker 1993. Kang et al."," Most simulations (Thomas Couchmann 1992, Katz, Hernquist, Weinberg 1992, Cen Ostriker 1993, Kang et al."
+ 1994. Lubin et al.," 1994, Lubin et al."
+ 1996. Pen 1998) show the gas to be slightly more extended than the dark matter within ~| Mpe due to shock heating and other transfers of energy.," 1996, Pen 1998) show the gas to be slightly more extended than the dark matter within $\sim 1$ Mpc due to shock heating and other transfers of energy."
+ This raises a possible concern. namely that the gas fraction may depend on radius.," This raises a possible concern, namely that the gas fraction may depend on radius."
+ At the virial radius. inflow and outflow of mass will affect hot gas and dark matter identically. so that the gas mass fraction should," At the virial radius, inflow and outflow of mass will affect hot gas and dark matter identically, so that the gas mass fraction should"
+Equation then reduces to: Jo satisfv this requirement. it is necessary to have many sources in the solid angle AQ.,Equation then reduces to: To satisfy this requirement it is necessary to have many sources in the solid angle $\Delta\Omega$.
+ However. this may be not sullicicnt because not all sources contribute the same flux at Earth.," However, this may be not sufficient because not all sources contribute the same flux at Earth."
+ The situation is thus complicated by the dependence on distance 0., The situation is thus complicated by the dependence on distance $D$.
+ Introducing the fraction [GD) of the total Lux produced by the sources closer than D. and the number of these sources in the sample (1). equation can be rewritten in the following way. defining the quantity Y: The number of sources 0(D) is readily. caleulated [rom equation(2): The fractional llux f(2)=BUD)fF(Duns). where represents the total Dux from sources closer than 1).," Introducing the fraction $f(D)$ of the total flux produced by the sources closer than $D$, and the number of these sources in the sample $n(D)$, equation can be rewritten in the following way, defining the quantity $\Upsilon$: The number of sources $n(D)$ is readily calculated from equation: The fractional flux $f(D) \equiv F(D) / F (D_{\rm max})$, where represents the total flux from sources closer than $D$."
+ Note the appearance of the weights wCD) assigned. by the sliding-box method in the last equation., Note the appearance of the weights $w(D)$ assigned by the sliding-box method in the last equation.
+ Neglecting deviations of #(2) from 1. which is reasonable on cosmological scales. we use the above equations to find that where Ny(D)=N(D.AL(2)) is the number of galaxies in the voltune-limited sample at distance {9 (cf.," Neglecting deviations of $\nu(D)$ from 1, which is reasonable on cosmological scales, we use the above equations to find that where $N_V(D) = N(D,M_*(D))$ is the number of galaxies in the volume-limited sample at distance $D$ (cf."
+ eqs. (2)))., eqs. \ref{eq:defNDM}) )).
+ We stress that equation is valid for [uxes modeled. using the sliding-box technique., We stress that equation is valid for fluxes modeled using the sliding-box technique.
+ For a single volume-limitecl subsample that is valid. up to Dy. a similar computation vields: Comparing eqs and(13).. we see that the number of sources contributing to a given Lux fraction at distance D ids increased. by a factor NyGD)DILUZNvUDQSDA.," For a single volume-limited subsample that is valid up to $D_{\rm max}$, a similar computation yields: Comparing eqs and, we see that the number of sources contributing to a given flux fraction at distance $D$ is increased by a factor $N_V (D) D_{\rm max}^3 / N_V (D_{\rm max}) D^3$."
+ This factor is unity at D=Daas (where the sliding-box method ollers no improvement). but may. become very large at small clistances.," This factor is unity at $D=D_{\rm max}$ (where the sliding-box method offers no improvement), but may become very large at small distances."
+ Inserting equation into brings us to the final expression for our criterion of small (uctuations: The three factors in equation respectively encode the dependence of Y on the angular scale. on the statistics of the Hux-limited sample. and on the attenuation of the moclel Dux.," Inserting equation into brings us to the final expression for our criterion of small fluctuations: The three factors in equation respectively encode the dependence of $\Upsilon$ on the angular scale, on the statistics of the flux-limited sample, and on the attenuation of the model flux."
+ The last factor reduces to unity in the case of no attenuation., The last factor reduces to unity in the case of no attenuation.
+ The equation has to hold for all values of D: when it is violated (1) lluctuations in Επιχ may occur in regions of angular size zMQ due to the contribution of a single source., The equation has to hold for all values of $D$; when it is violated $\mathcal{O}(1)$ fluctuations in flux may occur in regions of angular size $\Delta\Omega$ due to the contribution of a single source.
+ We now discuss some of the quantities entering equation(14)., We now discuss some of the quantities entering equation.
+. Phe number of galaxies IN(D) becomes small at both very. large. ancl very small distances. potentially leading to large. Uuetuations.," The number of galaxies $N_V(D)$ becomes small at both very large and very small distances, potentially leading to large fluctuations."
+ Lor large distances this can be prevented by considering only sources up to a maximum clistance Dyas and assuming an isotropic [lux from sources bevond that distance., For large distances this can be prevented by considering only sources up to a maximum distance $D_{\rm max}$ and assuming an isotropic flux from sources beyond that distance.
+ Alternatively. particle horizons may provide a natural maximum clistance (sec below).," Alternatively, particle horizons may provide a natural maximum distance (see below)."
+ In the case of small distances. the number of nearby sources is small while their contribution may he important due to. their proximity.," In the case of small distances, the number of nearby sources is small while their contribution may be important due to their proximity."
+ Unlike theDuctuations at large distances which are due to our poor knowledge of the galaxy clistribution at those scales. the fluctuations at small distances. are physical and may represent the actual [ux variations due to close sources.," Unlike the fluctuations at large distances which are due to our poor knowledge of the galaxy distribution at those scales, the fluctuations at small distances are physical and may represent the actual flux variations due to close sources."
+ Their complete treatment may require a case-by-case study of the most nearby objects., Their complete treatment may require a case-by-case study of the most nearby objects.
+ The last factor in equation encodes the effect of Hux attenuation. which can play an important role in modeling the Hux of UllECTIs and of verv-high. energy eamma ravs.," The last factor in equation encodes the effect of flux attenuation, which can play an important role in modeling the flux of UHECRs and of very-high energy gamma rays."
+ Focusing on the case of ULIECHTIUs. we show in figure 5the flux attenuation [actor J as a function of distance for ULUECH. protons.," Focusing on the case of UHECRs, we show in figure \ref{fig:ffunc} the flux attenuation factor $J$ as a function of distance for UHECR protons."
+ Phe attenuation factor is obtained using à numerical cosmic-ray propagation code described in ??..," The attenuation factor is obtained using a numerical cosmic-ray propagation code described in \citet{Koers:2008hv,Koers:2008ba}."
+ For comparison the attenuation factor due to redshift only is also shown in the figure., For comparison the attenuation factor due to redshift only is also shown in the figure.
+ In figure 6 we show the quantity which accounts for [ux attenuation in equation(14)., In figure \ref{fig:attfactor} we show the quantity which accounts for flux attenuation in equation.
+ Note that. as indicated in figure 5.. the horizon [or ULLECTTU protons above 00 IeV. is around 200 Alpe.," Note that, as indicated in figure \ref{fig:ffunc}, the horizon for UHECR protons above 60 EeV is around 200 Mpc."
+ Since sources bevond this distance do not contribute to the observed Iux. the requirement Yc1 should be satisfied automatically.," Since sources beyond this distance do not contribute to the observed flux, the requirement $\Upsilon \gg 1$ should be satisfied automatically."
+ We observe from figure 6. that zl indeed blows up around 200 Mpc. which guarantees that Yc1 for any value o£ AQ or IN.," We observe from figure \ref{fig:attfactor}
+ that $A$ indeed blows up around 200 Mpc, which guarantees that $\Upsilon \gg 1$ for any value of $\Delta \Omega$ or $N_V$ ."
+The cexiremely red objects! (EROs) are a populationof very red galaxies (e.g. with /? -5or£.-dvc 4) which appear züntward of A~17.5.,The `extremely red objects' (EROs) are a population of very red galaxies (e.g. with $R-K>5$ or $I-K>4$ ) which appear faintward of $K\sim 17.5$.
+ EROs are of special1 interest in the study. of galaxy. evolution. in that their colours ancl other properties suggest that they are the high-redshift (z=1 2) counterparts and. progenitors of local elliptical and SO ealaxies. and are amongst the oldest galaxies present at hese recdshifts.," EROs are of special interest in the study of galaxy evolution, in that their colours and other properties suggest that they are the high-redshift $z\simeq 1$ --2) counterparts and progenitors of local elliptical and S0 galaxies, and are amongst the oldest galaxies present at these redshifts."
+ Recent spectroscopy of some of the brightest. EROs (Cimatti et al., Recent spectroscopy of some of the brightest EROs (Cimatti et al.
+ 2002a: Saracco et al., 2002a; Saracco et al.
+ 2003) has revealed a mixture of old. passive galaxies (hereafter. plZltOs) and vounger star-forming or post-starburst galaxies with strong clustv-recddening (hereafter dstEBtOs).," 2003) has revealed a mixture of old, passive galaxies (hereafter pEROs) and younger star-forming or post-starburst galaxies with strong dusty-reddening (hereafter dsfEROs)."
+ Some EROs also host AGN activity., Some EROs also host AGN activity.
+ Recent surveys lind that ~ 1030 per cent of EROs on the Chandra Deep Field North (CDEN) are detected in. X-rays and. while in some of these the nav be produced. by starbursts. the majority have X- properties indicative of obscured AGN CAlexander et al.," Recent surveys find that $\sim
+10$ –30 per cent of EROs on the Chandra Deep Field North (CDFN) are detected in X-rays and, while in some of these the X-rays may be produced by starbursts, the majority have X-ray properties indicative of obscured AGN (Alexander et al."
+ 2002: Vienali et al., 2002; Vignali et al.
+ 2003)., 2003).
+" In our previous paper (Roche et al,", In our previous paper (Roche et al.
+ 2002. hereafter also Paper D). we combined our deep Ix-band survey of the “ELALS:N2° Biekd with the ULL (reaching fy 21) and Inerid (AV 20) cameras with a deep Zi-band image. to select à sample of 158 EROs as RoA75.0 galaxies.," 2002, hereafter also Paper I), we combined our deep K-band survey of the `ELAIS:N2' field with the UFTI (reaching $K\simeq 21$ ) and Ingrid $K\simeq 20$ ) cameras with a deep $R$ -band image, to select a sample of 158 EROs as $R-K>5.0$ galaxies."
+ The, The
+The relation between bolometric luminosity {μου and jet power is plotted in Fig. 3..,The relation between bolometric luminosity $L_{\rm bol}$ and jet power $Q_{\rm jet}$ is plotted in Fig. \ref{fig4}.
+" The jet power is lower than the Qj,bolometric luminosity for all sources in this sample. which would still hold even if the bolometric luminosity is overestimated by less than an order of magnitude for spectrum sources."," The jet power is lower than the bolometric luminosity for all sources in this sample, which would still hold even if the bolometric luminosity is overestimated by less than an order of magnitude for flat-spectrum sources."
+" In order to test relation (22)) between and total X- luminosity Lx from the disk coronas QWpredicted by the scenario of the jets being magnetically accelerated by the fields created 1n the disk coronas. we roughly convert X-ray lummosity in 0.1—2.4 keV κο5gv detected by ROSAT (Vogesetal.1999) to the total X-ray luminosity assuming the X-ray continuum emission to be in the range from 100 keV with a mean energy spectral index ay=1 (fxx E's),"," In order to test relation \ref{qjetbol}) ) between $Q_{\rm
+jet}^{\rm max}$ and total X-ray luminosity $L_{\rm X}$ from the disk coronas predicted by the scenario of the jets being magnetically accelerated by the fields created in the disk coronas, we roughly convert X-ray luminosity in $0.1-2.4$ keV $L_{\rm X,0.1-2.4keV}$ detected by ROSAT \citep{v99} to the total X-ray luminosity assuming the X-ray continuum emission to be in the range from $0.01-100$ keV with a mean energy spectral index $\alpha_{\rm X}=1$ $f_{\rm X}\propto E^{-\alpha_{\rm X}}$ )."
+ This is consistent with the general features of the theoretical calculations on the X-ray spectra from the disk coronas (e.g.. Nakamura Osaki 1993).," This is consistent with the general features of the theoretical calculations on the X-ray spectra from the disk coronas (e.g., Nakamura Osaki 1993)."
+ In this case. the total X-ray luminosity Ly~2.9Lxo2v.," In this case, the total X-ray luminosity $L_{\rm X}\simeq 2.9 L_{\rm X,0.1-2.4keV}$."
+ In Fig. 4..," In Fig. \ref{fig5},"
+ the relation between A-corrected core radio luminosity Les and X-ray luminosity Ly.," the relation between $K$ -corrected core radio luminosity $L_{\rm
+c,5G}$ and X-ray luminosity $L_{\rm X}$ ."
+ The linear regression considering errors in both coordinates gives (Pressetal.1992) + 7.51727for 40. flat-spectrum sources logLx(the solid line in 0.866logFig. 4)).," The linear regression considering errors in both coordinates gives \citep{p92} +7.517 for 40 flat-spectrum sources (the solid line in Fig. \ref{fig5}) ),"
+" and * 29.272028) 18 steep-spectrum sourceslog, (the dotted 20.371log,line in Fig. 4)).", and +29.272 for 18 steep-spectrum sources (the dotted line in Fig. \ref{fig5}) ).
+ The Spearmannfor correlation analyses (Pressetal.1992) show that the correlation for flat-spectrum sources 1s at the significant level of 99.98 per cent. while it becomes 98.5 per cent for steep-spectrum sources.," The Spearmann correlation analyses \citep{p92} show that the correlation for flat-spectrum sources is at the significant level of 99.98 per cent, while it becomes 98.5 per cent for steep-spectrum sources."
+ The different slopes of the correlations between Lx and Lysc for steep-spectrum and flat-spectrum sources are obvious. which may imply that theobserved X-ray emission from these two kinds of AGNs may have different origins.," The different slopes of the correlations between $L_{\rm X}$ and $L_{\rm c,5G}$ for steep-spectrum and flat-spectrum sources are obvious, which may imply that theobserved X-ray emission from these two kinds of AGNs may have different origins."
+Nyquist limit of (he original images. but rather by the spacing of the dithers of the combined images.,"Nyquist limit of the original images, but rather by the spacing of the dithers of the combined images."
+ À gentle taper in Fourier space will best suppress this ringing., A gentle taper in Fourier space will best suppress this ringing.
+ However. to pass all ol the information allowed by the optics. while limiting noise correlation. one prefers a sharp cuttoff near the Nvquist Irequencey. of hανA/D.," However, to pass all of the information allowed by the optics, while limiting noise correlation, one prefers a sharp cuttoff near the Nyquist frequency of $| \vec{k} | \approx \lambda/D$."
+ For the examples shown here a circular mask (in Fourier space) that [alls from. to transmission over a width of 0.1A/D is used., For the examples shown here a circular mask (in Fourier space) that falls from to transmission over a width of $0.1 \lambda/D$ is used.
+ The sampling of the Drizzled image was nearly twice the Nvquist sampling of the band-limited image., The sampling of the Drizzled image was nearly twice the Nyquist sampling of the band-limited image.
+ This was found to give good suppression of ringing in the simulated WFEPC? and ACS images. while limiting correlated noise and minimizing the loss of true intormation in the image.," This was found to give good suppression of ringing in the simulated WFPC2 and ACS images, while limiting correlated noise and minimizing the loss of true information in the image."
+ The center of the taper function. was sliehtlv to the blue of the Nyquist Irequeney., The center of the taper function was slightly to the blue of the Nyquist frequency.
+ This essentially eliminates correlated noise (when a final image is created wilh Nyquist sampling). at (he expense of allowing a slightly greater amount of noise through the filter.," This essentially eliminates correlated noise (when a final image is created with Nyquist sampling), at the expense of allowing a slightly greater amount of noise through the filter."
+ No claim is made that (his particular choice of cutoff. apodization. or sampling is optimal. but it was found effective.," No claim is made that this particular choice of cutoff, apodization, or sampling is optimal, but it was found effective."
+ The algorithm described here is part of a class of techniques that vary. in (heir use of specific weightngs or interpolation schemes to iteratively arrive at a solution of (he irregular sampling problem., The algorithm described here is part of a class of techniques that vary in their use of specific weightings or interpolation schemes to iteratively arrive at a solution of the irregular sampling problem.
+ These all tend to converge geometrically to the correct solution with the speed of conversion being dependent on the value of 2/1. where / is (he maximal distance between samples and vy is (e Nveuist Irequeney of the data.," These all tend to converge geometrically to the correct solution with the speed of conversion being dependent on the value of $2 l \nu_0$, where $l$ is the maximal distance between samples and $\nu_0$ is the Nyquist frequency of the data."
+ Where 2/154 1. convergence is rapid: however. where 2/14;~1. convergence is slow or [ails (see Werther 1999. Eq 50 for a specific case. and Werther 2004 and Feichtinger.," Where $2 l \nu_0 \ll 1$ , convergence is rapid; however, where $2 l \nu_0 \sim 1$, convergence is slow or fails (see Werther 1999, Eq 50 for a specific case, and Werther 2004 and Feichtinger."
+ Grócchening 1994 for broader discussions)., Gröcchening 1994 for broader discussions).
+ In the case being discussed here. however. (here is no unique solution because of the presence of noise.," In the case being discussed here, however, there is no unique solution because of the presence of noise."
+ Further. because of geometric distortion the value of /. and thus 2/14. is likely to vary across the image. and thus so too will the rapiditv of local convergence.," Further, because of geometric distortion the value of $l$, and thus $2 l \nu_0$, is likely to vary across the image, and thus so too will the rapidity of local convergence."
+ In practice. it is [airlv simple for the user to test whether a particular observing plan with a given instrument is likely (o produce acceptable images.," In practice, it is fairly simple for the user to test whether a particular observing plan with a given instrument is likely to produce acceptable images."
+ To do (his the user should creale a well sampled “output” image., To do this the user should create a well sampled “output” image.
+ In the test cases shown here. PSFs placed al regular intervals with a small random cdither were used.," In the test cases shown here, PSFs placed at regular intervals with a small random dither were used."
+ The user can vary the amplitude of the PSF as well. to see how the technique works in dillerent noise regimes.," The user can vary the amplitude of the PSF as well, to see how the technique works in different noise regimes."
+ This simulated image is (hen mapped to the input image scale using Drizzle package routine.blot.," This simulated image is then mapped to the input image scale using Drizzle package routine,."
+ will. with the correct distortion files. map the image back as it would have appeared on the distorted lield of the detector.," will, with the correct distortion files, map the image back as it would have appeared on the distorted field of the detector."
+ The user can create as many input images as he or she desires. with a dither pattern emulating the one expected to be obtained from the observations.," The user can create as many input images as he or she desires, with a dither pattern emulating the one expected to be obtained from the observations."
+ Noise can then be added (o each image to match the level expected in the observations., Noise can then be added to each image to match the level expected in the observations.
+Finally one,Finally one
+combine them with virial Mp (from Magellan/INACS and zCOSMOS/VLT/VIMOS) for 10 AGN iu the redshift iuterval 1.06<21.92. 3.25.5 Chivas after the Big Bane.,"combine them with virial $M_\mathrm{BH}$ (from Magellan/IMACS and zCOSMOS/VLT/VIMOS) for 10 AGN in the redshift interval $1.06<z<1.92$, 3.2–5.5 Gyrs after the Big Bang."
+" Throughout we use AB zero-poiuts unless otherwise noted and a cosmology of Hy=T0lans!Mpe 1,03; 0.3. and O4=0.7."," Throughout we use AB zero-points unless otherwise noted and a cosmology of $H_0=70\,\mathrm{km\,s^{-1} Mpc^{-1}}$ , $\Omega_M=0.3$ , and $\Omega_\Lambda=0.7$."
+ Iu order to control selection effects on AL.. Mp or any special relation between these two. we require a transparent sample definition.," In order to control selection effects on $M_\mathrm{*}$, $M_\mathrm{BH}$ or any special relation between these two, we require a transparent sample definition."
+ Our selection of tvpe-1 AGN is based ou N-ray detection in the NAIAECOSAIOS survey (IHasiugerctal.2007:Cappellutiet2009) and subsequent identification of their optical counterparts (Brusaetal.2007).," Our selection of type-1 AGN is based on X-ray detection in the XMM-COSMOS survey \citep{hasi07,capp09} and subsequent identification of their optical counterparts \citep{brus07}."
+. Classification as tvpe-1 ACN for this study uses both spectroscopic ideutification of broad cussion lines in the Alaeclan/TXIACS (Trampetal.2009a) and zCOSMOS (Lillyetal.2007). surveys as well as photometric classification using the loug spectral baseline SED covered in COSMOS (Capaletal.2007:in prep...," Classification as type-1 AGN for this study uses both spectroscopic identification of broad emission lines in the Magellan/IMACS \citep{trum09a} and zCOSMOS \citep{lill07} surveys as well as photometric classification using the long spectral baseline SED covered in COSMOS \citep[Capak et al., in
+ prep.]{capa07,salv09,ilbe09}."
+ As such. only objects with hieh-coufidence classification cuter our sample.," As such, only objects with high-confidence classification enter our sample."
+ For ~550 προ AGN selected this wav we require the following data. resulting iu a raucous subsauple: (1) coverage by ACS ESLIW (Scovilleetal.2007:Noekeuocretal. 2007).. and (2) NICALOSparallels’? as well as (3) spectra for virial Mpjpestinates from the ΑΙΟΠ broad cluission line.," For $\sim$ 550 type-1 AGN selected this way we require the following data, resulting in a random subsample: (1) coverage by ACS F814W \citep{scov07b,koek07}, and (2) NICMOS, as well as (3) spectra for virial $M_\mathrm{BH}$ -estimates from the MgII broad emission line."
+ The laniting factors are the relative NICS coverage of of the ACS area (20. ACN) and the status of ongoing spectroscopic survevs., The limiting factors are the relative NIC3 coverage of of the ACS area $\sim$ 20 AGN) and the status of ongoing spectroscopic surveys.
+ Black hole masses have been caleulated by Truupetal.(2009)D) οι IMACS spectra and by Merlonietal.(2009) using ZCOSMOS/VIMOS., Black hole masses have been calculated by \citet{trum09b} from IMACS spectra and by \citet{merl09} using zCOSMOS/VIMOS.
+ For this letter we have BIT masscs available for LO ACN. spanning the redshift range 1<2 <2.," For this letter we have BH masses available for 10 AGN, spanning the redshift range $1<z<2$ ."
+" Two AGN were observed by both instiunceuts the A/py-estimates are cousisteut within 0.2 dex in both cases,", Two AGN were observed by both instruments – the $M_\mathrm{BH}$ -estimates are consistent within 0.2 dex in both cases.
+ Sample information is listed in Table 1.. inclaiding DII masses and galaxy parameters derived below.," Sample information is listed in Table \ref{tab:all}, including BH masses and galaxy parameters derived below."
+ We obtain information on the host galaxies usine broad-band photometry from the high-resolution UST ACS/WFC images in the Faliw (=f) filter with 07003(pixel saaupling (IXoekeinocretal.2007).. aud the NICALIOS/NIC3 parallels in the FIGOW (277) filter at 1101/pixel. both iuteerated for oue orbit.," We obtain information on the host galaxies using broad-band photometry from the high-resolution HST ACS/WFC images in the F814W $I$ ) filter with 03/pixel sampling \citep{koek07}, and the NICMOS/NIC3 parallels in the F160W $H$ ) filter at 101/pixel, both integrated for one orbit."
+ Iu the NICMOS Z-buud the host salaxies of all AGN are clearly resolved. visible already to the unaided eve.," In the NICMOS $H$ -band the host galaxies of all AGN are clearly resolved, visible already to the unaided eye."
+ We heuce extract the host ealaxy fux from the composite galaxy|ACN NICMOSimage by modelling the two-dimensional light distribution of each object using GALFIT (Pengetal.2002:Pene2009) iu Version 2.0 (C. PPene. ccouun.).," We hence extract the host galaxy flux from the composite galaxy+AGN NICMOSimage by modelling the two-dimensional light distribution of each object using GALFIT \citep{peng02,peng09} in Version 3.0 (C. Peng, comm.)."
+ We restrict our models to a point-source plus a single elliptical Sérrsie (1968)) profile., We restrict our models to a point-source plus a single elliptical Sérrsic \citeyear{sers68}) ) profile.
+ Previous siauulatious show that at our resolution and depth it is uureliable to use many and/or complex galaxy components (Smuuous&ταν2008:Sáuchezetal.," Previous simulations show that at our resolution and depth it is unreliable to use many and/or complex galaxy components \citep{simm08,sanc04a}."
+ 2001).. We carry out several passes of GALFIT. first with free Séórsie p)». aud subsequently fixed i=1. 2 or L. depending ou the best free-o-ft.," We carry out several passes of GALFIT, first with free Sérrsic $n$, and subsequently fixed $n = 1$, 2 or 4, depending on the best $n$ -fit."
+ We require CALFIT to converge on a sensible solution. indicating that au actual host galaxy is beiug described and not uuncertamties in the poiut spread function (PSF).," We require GALFIT to converge on a sensible solution, indicating that an actual host galaxy is being described and not uncertainties in the point spread function (PSF)."
+ This 1118 we require à CALFIT uucleus-to-hiost contrast of «3.25 imag. a lidt-light radius of k&72 pixcl. and 0.5«0«8.," This means we require a GALFIT nucleus-to-host contrast of $<3.25$ mag, a half-light radius of $r>2$ pixel, and $0.5<n<8$."
+ NID 2261 and 5230 both show unrealistically conrpact host ealaxv aodel scale-leusths. iudicatiug an unsuccessful CALFIT model.," XID 2261 and 5230 both show unrealistically compact host galaxy model scale-lengths, indicating an unsuccessful GALFIT model."
+ For these two cases we suuply subtract the best fitting sinele poiut-source inodel — without a Sérrsic coniponenut from the original nuage. resulting im an only sliehtlv oversubtracted/underestimated lost) galaxy.," For these two cases we simply subtract the best fitting single point-source model – without a Sérrsic component – from the original image, resulting in an only slightly oversubtracted/underestimated host galaxy."
+ We use aperture photometry on the host ealaxy in these cases. and the GALFEIT galaxy model magnitude for the other eight.," We use aperture photometry on the host galaxy in these cases, and the GALFIT galaxy model magnitude for the other eight."
+ Iu total we resolvee 10 host galaxies in the sample uo object drops frou the sample due to high uucleus-to-host contrast or other reasons., In total we resolve 10 host galaxies in the sample – no object drops from the sample due to high nucleus-to-host contrast or other reasons.
+ Extracted host galaxy nuages are shown in Figures 1| 2.., Extracted host galaxy images are shown in Figures \ref{fig:allobj}+ \ref{fig:allobj2}.
+ Iu the ACS 7-baud data the contrast between ACN and host galaxy is less favorable than in the Παπά. as expected. from the near-UW SED.," In the ACS $I$ -band data the contrast between AGN and host galaxy is less favorable than in the $H$ -band, as expected from the near-UV SED."
+ Therefore we take a two-step approach., Therefore we take a two-step approach.
+" First. we carry out ""peak subtraction” removal of the nuclear component for the f-haud. by scaling a PSF to the central 1 pixel aperture flux."," First, we carry out “peak subtraction” removal of the nuclear component for the $I$ -band, by scaling a PSF to the central 4 pixel aperture flux."
+ From statistics on the expected random residuals after subtraction of this scaled PSF for several 1000 stars (Jahnkeetal.200Πο]. we require a residual flux of 2554 for a host ealaxy to have a ligh probability of being resolved.," From statistics on the expected random residuals after subtraction of this scaled PSF for several 1000 stars \citep[][]{jahn04b}, we require a residual flux of $>$ for a host galaxy to have a high probability of being resolved."
+ This is the case for 7/10 objects — the hosts of XIDs 59. 2261. and 5230 remain unresolved in the ACS uae.," This is the case for 7/10 objects – the hosts of XIDs 59, 2261, and 5230 remain unresolved in the ACS image."
+ The seven resolved objects are again modeled using GALFTIT. with Sérrsic » fixed from the Z7-baud fit. thus minimuizimeS/N-«depeudeut biases in our extracted colors.," The seven resolved objects are again modeled using GALFIT, with Sérrsic $n$ fixed from the $H$ -band fit, thus minimizing S/N-dependent biases in our extracted colors."
+ As for the ZI-baud we check the models for successful convergence. Which is the case for alb seven resolved objects. and again use the host model maguitude.," As for the $H$ -band we check the models for successful convergence, which is the case for all seven resolved objects, and again use the host model magnitude."
+ The resulting host ealaxy photometry is also Listed iu Table 1.., The resulting host galaxy photometry is also listed in Table \ref{tab:all}.
+ Our general approach is to coustrain the mass-to-leht ratio (ML) ofeach host galaxy by a single optical color., Our general approach is to constrain the mass-to-light ratio (M/L) ofeach host galaxy by a single optical color.
+ We have successfully. eiiploved. this method before to study. stellar populations of low-: quasar host galaxies (Jahnkeetal.200la) or stellar ages aud masses of quasar hosts at 0.5<2«3 (Sáuchezetal.2001:Jalinke 2008)..," We have successfully employed this method before to study stellar populations of $z$ quasar host galaxies \citep{jahn04a} or stellar ages and masses of quasar hosts at $0.5<z<3$ \citep{sanc04a,jahn04b,kuhl05,schr08}. ."
+ Tere we compute stellar masses from the rest-frame V- Iuuimositv iu combination with the M/L from the rest-frame (2V) color:, Here we compute stellar masses from the rest-frame $V$ -band luminosity in combination with the M/L from the rest-frame $(B-V)$ color:
+lüpt components of the magnetic field have Strongbeen multipole to an role in the racio emissionlong from t,10pt Strong multipole components of the magnetic field have long been thought to play an important role in the radio emission from pulsars.
+hought play fields importanthave been invoke: for the of pulsars.electron. Multipole pairs in the pulsar generat, Multipole fields have been invoked for the generation of electron positron pairs in the pulsar magnetosphere.
+ionFor the positronRuderman Sutherlane (1975) moclel magnetosphere. that example.the radius of curvature of the requires surface be of of the Ποιάstellarlines nearracius theto stellarsustain pair should in thelong orderperio," For example, the Ruderman Sutherland (1975) model requires that the radius of curvature of the field lines near the stellar surface should be of the order of the stellar radius to sustain pair production in long period pulsars."
+ This is much smaller than the production radius of pulsars.curvature of the field., This is much smaller than the expected radius of curvature of the dipole field.
+ Such a small expected.racius of curvature could be a signaturedipole of either an olfse (Arons. 1998) or of a field. extremely(Barnard dipole.Arons. 1982).," Such a small radius of curvature could be a signature of either an extremely offset dipole (Arons, 1998) or of a multipolar field (Barnard Arons, 1982)."
+ Further. soft. X-ray multipolarobservations of pulsars show non-uniform surface which can be attributed to the presence of a temperatures.quadrupolar field (Page Sarmiento. 1996).," Further, soft X-ray observations of pulsars show non-uniform surface temperatures which can be attributed to the presence of a quadrupolar field (Page Sarmiento, 1996)."
+ multipole structure at. and. near the polar cap is also Alagnetic to be responsible for the unique pulse wrolile of a thought (Vivekanand Radhakrishnan 1980. ντος 199 pulsarRankin Rathnasree 1995).," Magnetic multipole structure at and near the polar cap is also thought to be responsible for the unique pulse profile of a pulsar (Vivekanand Radhakrishnan 1980, Krolik 1991, Rankin Rathnasree 1995)."
+ The recent estimates that there should be several tens of sparks ≻∪↓≻⊔↓⋜∐⊲↓⊔⋏∙≟⇂↓↥∢⊾↓≻∪↓⋜⊔⋅≼⇍∥↓≻⊲↓⊳∖⋜↧≱∖∪∣⋊⋅≱∖↿⋖⋅⇀∖↓≻⇂⋜↧↕↓↥⋜↧∣⋡↓⋖⋅↕⇂⋅ fields dictate the near the surface multipole(Deshpande Rankin 1998. sparkRankin geometryDeshpande 1998. Seiradakis ," The recent estimates that there should be several tens of sparks populating the polar cap is also best explainable if multipole fields dictate the spark geometry near the surface (Deshpande Rankin 1998, Rankin Deshpande 1998, Seiradakis 1998)."
+Significant evolution in the structure of the 1998).field. the lifetime of a pulsar may herefore leave magneticobservable during, Significant evolution in the structure of the magnetic field during the lifetime of a pulsar may therefore leave observable signatures.
+ I£ the grow weaker in comparison signatures.to the multipolesthen one can oogressively.pulse profiles to with age and dipolevice versa., If the multipoles grow progressively weaker in comparison to the dipole then one can expect pulse profiles to simplify with age and vice versa.
+ The expectevolution of the simplifvfields in neutron stars in is still a magneticopen question., The evolution of the magnetic fields in neutron stars in general is still a relatively open question.
+ During the ast generaldecade. two relativelyalternative scenarios for the field evolution have major," During the last decade, two major alternative scenarios for the field evolution have emerged."
+ One of these assumes that. the ielel of the neutron emerged.star permeates the whole star at Xurth. and its evolution is dictated. by the interaction »etween. superluicl vortices (carrving angular momentum)," One of these assumes that the field of the neutron star permeates the whole star at birth, and its evolution is dictated by the interaction between superfluid vortices (carrying angular momentum) and superconducting fluxoids (carrying magnetic flux) in the stellar interior."
+ un4 ," As the star spins down, the outgoing vortices may drag and expel the field from the interior leaving it to decay in the crust (Srinivasan 1990)."
+",∖ dipole field carried."," In a related model, plate tectonic motions driven by pulsar spindown drags the magnetic poles together, reducing the magnetic moment (Ruderman The other scenario assumes that most of the field is generated in the outer crust after the birth of the neutron star (Blandford, Applegate Hernquist 1983)."
+ by., The later evolution of this field is governed entirely by the ohmic decay of currents in the crustal layers.
+ such. currents has been investigated in some detail in the recent literature (Sang LOST. Urpin 1994. Urpin GeppertChanmugam 1995. 1996. lxonar GeppertBhattacharva 1997. 1905).," The evolution of the dipole field carried by such currents has been investigated in some detail in the recent literature (Sang Chanmugam 1987, Geppert Urpin 1994, Urpin Geppert 1995, 1996, Konar Bhattacharya 1997, 1998)."
+ These studies include evolution isolated. neutron stars as well as those Ποιάaccreting from theirin binary companions., These studies include field evolution in isolated neutron stars as well as those accreting from their binary companions.
+ The results show interesting agreements with observations lending some credence to the crustal picture., The results show interesting agreements with observations lending some credence to the crustal picture.
+ In this paper. we the ohmic evolution of higher order in isolated. exploreneutron stars assuming the currents to be multipolesoriginally confined in the crustal region.," In this paper, we explore the ohmic evolution of higher order multipoles in isolated neutron stars assuming the currents to be originally confined in the crustal region."
+ Our goal is to fined whether there would be observable οσοι on the pulse of radio emission from aavisolated pulsars as a result shapeevolution., Our goal is to find whether there would be any observable effect on the pulse shape of radio emission from isolated pulsars as a result of this evolution.
+ In section 2 we discuss the details of the ofthis and in section 3 we present our results and discuss the computationimplications., In section 2 we discuss the details of the computation and in section 3 we present our results and discuss the implications.
+ The evolution of the magnetic field. due to ohmic dillusion. is governed by the equation (Jackson 1975): where co(r.4) ds the electrical. conductivity of the medium.," The evolution of the magnetic field, due to ohmic diffusion, is governed by the equation (Jackson 1975): = - ), where $\sigma(r,t)$ is the electrical conductivity of the medium."
+" Following Wendell. Van Horn. Sargent. (1987) we introduce a vector potential A=(0.0.24,,) assuming the field to be purely poloidal. such that: where SC.6.4) is the Stokes’ stream function."," Following Wendell, Van Horn Sargent (1987) we introduce a vector potential ${\bf A} = (0, 0, A_{\phi})$ assuming the field to be purely poloidal, such that: where $S(r, \theta, t)$ is the Stokes' stream function."
+ S can be separated in απο @ in the form: where Pr1(cos(@)) is. the associated. polynomial. of degree one ancl is the multipole radial Legendrefunction., $S$ can be separated in $r$ and $\theta$ in the form: where $P_{l}^{1}(cos(\theta))$ is the associated Legendre polynomial of degree one and$R_{l}$ is the multipole radial function.
+ From equation (1)) we obtain{1 :, From equation \ref{e_induction}) ) we obtain :
+" 21cm (F+G)*H=I,(F+6)*HF GH. G H / F, G, H I. (H) performance(F). ", $21$ $(F+G)*H~=~I$$(F+G)*H$$F$ $G$$H$ $F$ $G$ $H$ $I$ $F$ $G$ $H$ $I$ $H$ $F$ 
+hickuess of the tachocline remains nich smaller thaw unity (D.~10 23 ,thickness of the tachocline remains much smaller than unity $P_e \approx 10^{-2}$ ).
+These estimates sussest to use the sinall-Pécclet-ΠΙΟ approxinatiou to investigate the property of sanall scale turbulent motions in stellar radiative zones., These estimates suggest to use the small-P\'ecclet-number approximation to investigate the property of small scale turbulent motions in stellar radiative zones.
+ A first possible investigation could be the homogeneous urbuleuce in presence of a ΠΠΣΟΥ mean shear aud mean eniperature gradient., A first possible investigation could be the homogeneous turbulence in presence of a uniform mean shear and mean temperature gradient.
+ With ecoplysical applications in wind. this configuration has already been extcusively studied for large Pécclet απσος (see Schuniumn 1996. for a review).," With geophysical applications in mind, this configuration has already been extensively studied for large Pécclet numbers (see Schumann 1996, for a review)."
+ A comparison with the small-Pécclet-ummber case should be very instructive., A comparison with the small-Pécclet-number case should be very instructive.
+ Another important topic concerns the awisotropy between vertical auc horizoutal notions iu au atmosphere dominated by thermal diffusion., Another important topic concerns the anisotropy between vertical and horizontal motions in an atmosphere dominated by thermal diffusion.
+ Our linear study sugeests that this anisotropy ca- stronecr than in a non-diffusive atmosphere., Our linear study suggests that this anisotropy can be stronger than in a non-diffusive atmosphere.
+ The ratio vetween vertical aud horizontal turbulent viscosities coul ve affected and this can estimated through umuuerica simulations ofthe smallPécclet-ummber equations., The ratio between vertical and horizontal turbulent viscosities could be affected and this can estimated through numerical simulations of the small-Pécclet-number equations.
+ Before concliding. it mist be noted that the lanit of aree diffusivities has already been considered (Spicec 1962.. Thual 1992)) iu the coutext of the Ravleigh-Béennard convection.," Before concluding, it must be noted that the limit of large diffusivities has already been considered (Spiegel \cite{spiegel}, Thual \cite{thual}) ) in the context of the Rayleigh-B\'ennard convection."
+" Ποπονα, au Huportaut plysica xopertv of thermal convection is lost iu this Lit."," However, an important physical property of thermal convection is lost in this limit."
+ The thermal stratification is indeed assumed uuchauge and this is uot compatible with the general observation hat convective motions transform the initial uustable stratification iuto an adiabatic stratification., The thermal stratification is indeed assumed unchanged and this is not compatible with the general observation that convective motions transform the initial unstable stratification into an adiabatic stratification.
+ There is uo such an inconsistency for the case of stably stratified radiative zones we considered here., There is no such an inconsistency for the case of stably stratified radiative zones we considered here.
+usually assumed.,usually assumed.
+ Presently there is no sound scientific basis on which one can either confirm or reject the hypothesis. so that it remains a serious possibility. of fundamental significance for compact stars (Weberetal...1997b:Weberetal.1997a:Maclsen..1998). plus various other physical phenomena (Madsen&Llaensel.1991).," Presently there is no sound scientific basis on which one can either confirm or reject the hypothesis, so that it remains a serious possibility of fundamental significance for compact stars \cite{Weber96a,Weber96b,Madsen98a} plus various other physical phenomena \cite{Aarhus91}."
+. Below we will explore the implications of the existence of strange stars with respect to the QPO phenomena., Below we will explore the implications of the existence of strange stars with respect to the QPO phenomena.
+ This enables one to test the possible existence of strange stars and thus draw definitive conclusions about the true ground state of strongly. interacting matter., This enables one to test the possible existence of strange stars and thus draw definitive conclusions about the true ground state of strongly interacting matter.
+ As pointed out by Alcock et al. (1986).," As pointed out by Alcock et al. \shortcite{Alcock86},"
+.. à strange star can carey a solicl nuclear crust whose densitv is. strictly limited by neutron drip., a strange star can carry a solid nuclear crust whose density is strictly limited by neutron drip.
+ This possibility is caused by the displacement of electrons at the of surfacethe strange matter core. which generates an electric dipole laver there.," This possibility is caused by the displacement of electrons at the surface of the strange matter core, which generates an electric dipole layer there."
+ Lt can be sullicientIy strong to produce a gap between ordinary atomic matter (crust) and the quark-matter surface. which prevents a conversion of ordinary atomic matter into the assumed Iower-Iving ground state of strange matter.," It can be sufficiently strong to produce a gap between ordinary atomic matter (crust) and the quark-matter surface, which prevents a conversion of ordinary atomic matter into the assumed lower-lying ground state of strange matter."
+ Obviously. free neutrons. being electrically. charge neutral. cannot exist in the crust. because they do not feel the Coulomb barrier and thus would gravitate toward the strange-quark matter core. where they are converted by hypothesis into strange matter.," Obviously, free neutrons, being electrically charge neutral, cannot exist in the crust, because they do not feel the Coulomb barrier and thus would gravitate toward the strange-quark matter core, where they are converted by hypothesis into strange matter."
+ Consequently. the density at the base of the crust (inner crust density) must always be smaller than neutron clip densitv.," Consequently, the density at the base of the crust (inner crust density) must always be smaller than neutron drip density."
+ The situation is illustrated. in Fig., The situation is illustrated in Fig.
+ 2. where the EOS of a strange star with crust is shown., \ref{fig:sm:eos1} where the EOS of a strange star with crust is shown.
+ The strange-star models presented in the following are constructed for an EOS of strange matter derived by Farhi Jalle (1984)., The strange-star models presented in the following are constructed for an EOS of strange matter derived by Farhi Jaffe \shortcite{Farhi84}.
+. We shall study models with a fixed strange quark mass. m.=10:) MeV. and the two bag constants. B3=145 MeV and BE?1=158 MeV. In Table 2. the observations of kilohertz QPOs [rom 16 sources are summarised.," We shall study models with a fixed strange quark mass, $m_{\rm s}=100$ MeV, and the two bag constants, $B^{1/4}=145$ MeV and $B^{1/4}=158$ MeV. In Table \ref{tab:qpo.data} the observations of kilohertz QPOs from 16 sources are summarised."
+ Phe sources can be classified into two classes. the Atoll- and. the Z-sources. depending. on the form of their colour-colour diagram.," The sources can be classified into two classes, the Atoll- and the Z-sources, depending on the form of their colour-colour diagram."
+ Almost all sources show simultaneously two kilohertz QPOs., Almost all sources show simultaneously two kilohertz QPOs.
+ Although the frequencies of both QPOs vary by several hundred. Hertz with the acerction rate. the variation of the frequency separation is only small.," Although the frequencies of both QPOs vary by several hundred Hertz with the accretion rate, the variation of the frequency separation is only small."
+ Therefore the observations strongly favour some kind of xit-Irequeney. model., Therefore the observations strongly favour some kind of beat-frequency model.
+ La such models the upper kilohertz QPO corresponds to the Kepler rotation at the inner edge of the accretion disk., In such models the upper kilohertz QPO corresponds to the Kepler rotation at the inner edge of the accretion disk.
+ The lower QPO corresponds to the recat [requeney between the upper QPO and the spin of the neutron star., The lower QPO corresponds to the beat frequency between the upper QPO and the spin of the neutron star.
+ This interpretation is supported by the Lact hat the frequeney separation is equal to the QDO frequency (or half of it) in type E X-ray bursts observed in some of the Atoll sources (see Tab. 2))., This interpretation is supported by the fact that the frequency separation is equal to the QPO frequency (or half of it) in type I X-ray bursts observed in some of the Atoll sources (see Tab. \ref{tab:qpo.data}) ).
+ The question remains however where the upper QPO is generated., The question remains however where the upper QPO is generated.
+ Strohmaver et al., Strohmayer et al.
+ (1996) suggested that this radius is identical to the magnetospheric radius. whereas Miller et al.," \shortcite{Strohmayer96a}
+ suggested that this radius is identical to the magnetospheric radius, whereas Miller et al."
+ (1998€)— propose that dt is identical to the sonic racdus (see van der Whis 1997. for a critical discussion of these mocdoels)., \shortcite{Miller96a} propose that it is identical to the sonic radius (see van der Klis 1997 \nocite{VanDerKlis97a} for a critical discussion of these models).
+ In both cases. the orbital radius has to be larger than the radius of the innermost stable orbit. (see Sect. 2.2)).," In both cases, the orbital radius has to be larger than the radius of the innermost stable orbit (see Sect. \ref{ssec:stable}) ),"
+" or equivalentlv. the frequency vovo, of the upper ΟΡΟ has to be smaller than the Ixepler frequeney at the innermost stable orbit: where the approximate expression was used."," or equivalently, the frequency $\nu_{\rm QPO1}$ of the upper QPO has to be smaller than the Kepler frequency at the innermost stable orbit: where the approximate expression was used."
+ The exact expression depends on the LOS and on the spin period of the neutron star., The exact expression depends on the EOS and on the spin period of the neutron star.
+ This inequality sets an upper limit to the mass of the star., This inequality sets an upper limit to the mass of the star.
+ Obviously. the orbital radius has to be larger than the star's radius. which is larger than the radius of the innermost stable orbit for low star masses.," Obviously, the orbital radius has to be larger than the star's radius, which is larger than the radius of the innermost stable orbit for low star masses."
+ This constraint sets a lower limit to the star mass which again depends on the spin period and the LOS., This constraint sets a lower limit to the star mass which again depends on the spin period and the EOS.
+ Figure 3 shows the maximally allowed Kepler frequeney as function of the gravitational mass for two LOss. aand ΙΟΣ.," Figure \ref{fig:qpo3} shows the maximally allowed Kepler frequency as function of the gravitational mass for two EOSs, and RHF8."
+ Wo the radius of the innermost stable orbit is smaller than the stars radius the maximally allowed Ixepler frequeney is given by the Kepler frequency. at the stars equator., If the radius of the innermost stable orbit is smaller than the star's radius the maximally allowed Kepler frequency is given by the Kepler frequency at the star's equator.
+ Vhis kind of solution is represented. by the rising branch of vy(M). since my increases with increasing eravitational mass.," This kind of solution is represented by the rising branch of $\nu_{\rm K}(M)$, since $\nu_{\rm K}$ increases with increasing gravitational mass."
+ Lf the innermost stable orbit is located outside the star the maximally allowed. Ixepler. frequency decreases again with increasing mass., If the innermost stable orbit is located outside the star the maximally allowed Kepler frequency decreases again with increasing mass.
+ Phe two LOSs diller for MLOAL. in the composition.," The two EOSs differ for $M\gtrsim 1.0\,M_\odot$ in the composition."
+ In RIES hyperons are included. whereas aassumes pure nucleonic matter.," In RHF8 hyperons are included, whereas assumes pure nucleonic matter."
+ RUGS is therefore. softer than, RHF8 is therefore softer than
+The extracted radial velocity data was fitted simultaneously with the transit photometry available at the time of analysis.,The extracted radial velocity data was fitted simultaneously with the transit photometry available at the time of analysis.
+ Three models are adjusted to the data: a Keplerian radial velocity orbit (?).. a photometric planetary transit (?).. and à spectroscopic transit. also known as Rossiter-MeLaughlin effect (2)..," Three models are adjusted to the data: a Keplerian radial velocity orbit \citep{Hilditch:2001p773}, a photometric planetary transit \citep{Mandel:2002p768}, and a spectroscopic transit, also known as Rossiter-McLaughlin effect \citep{Gimenez:2006p31}."
+ This combined approach ts very useful for taking into account all of the possible contributions to. the uncertainties due to correlations among all relevant parameters., This combined approach is very useful for taking into account all of the possible contributions to the uncertainties due to correlations among all relevant parameters.
+ A single set of parameters describes both the photometry and the radial velocities., A single set of parameters describes both the photometry and the radial velocities.
+ We use a Markov Chain Monte Carlo (MCMC) approach to optimize the models and estimate the uncertainties of the fitted parameters., We use a Markov Chain Monte Carlo (MCMC) approach to optimize the models and estimate the uncertainties of the fitted parameters.
+ The fit of the model to the data is quantified using the y statistic., The fit of the model to the data is quantified using the $\chi^2$ statistic.
+ The code ts described in detail by ?.. has been used several times (eg: ?)) and is similar to the code deseribed in 2..," The code is described in detail by \citet{Triaud:2009p3865}, has been used several times (eg: \citet{Gillon:2009p3869}) ) and is similar to the code described in \citet{Cameron:2007p2879}."
+ We fitted up to 10 parameters. namely the depth of the primary transit D. the radial velocity (RV) semi-amplitude K. the impact parameter >. the transit width W. the period P. the epoch of mid-transit 7o. ecosc esinw. Vsin/cosB. and Vsin7sinfl.," We fitted up to 10 parameters, namely the depth of the primary transit $D$, the radial velocity (RV) semi-amplitude $K$, the impact parameter $b$, the transit width $W$ , the period $P$ , the epoch of mid-transit $T_0$, $e \cos \omega$, $e \sin \omega$, $V\sin I\,\cos \beta$, and $V\sin I\,\sin \beta$."
+" Here eis the eccentricity and w the angle between the line of sight and the periastron. Vsin/ is the sky-projected rotation velocity of the while β is the sky-projected angle between the stellar rotation axis (??) and the planet's orbitalaxis""."," Here $e$ is the eccentricity and $\omega$ the angle between the line of sight and the periastron, $V\sin I$ is the sky-projected rotation velocity of the while $\beta$ is the sky-projected angle between the stellar rotation axis \citep{Hosokawa:1953p2002, Gimenez:2006p31} and the planet's orbital."
+ These parameters have been chosen to reduce correlations between then., These parameters have been chosen to reduce correlations between then.
+ The use of uncorrelated parameters allows to explore parameter space more efficiently since the correlatio length between jumps is smaller., The use of uncorrelated parameters allows to explore parameter space more efficiently since the correlation length between jumps is smaller.
+ Eccentricity and periastro angle were paired as were Vsin/ and f., Eccentricity and periastron angle were paired as were $V\sin I$ and $\beta$.
+ This breaks à£5 correlation between them (the reader is invited to compare Figs., This breaks a correlation between them (the reader is invited to compare Figs.
+ 244 3 for à clear illustration for choosing certai jump parameters as opposed to others)., \ref{fig:WASP4dis}d d \ref{fig:betas} for a clear illustration for choosing certain jump parameters as opposed to others).
+ This way we also explore solutions around zero more easily: ecosco and esinco move in the |-I.1[ range while e could only be floating i J0.1[.," This way we also explore solutions around zero more easily: $e\cos \omega$ and $e \sin \omega$ move in the ]-1,1[ range while $e$ could only be floating in ]0,1[."
+ For exploring particular solutions such as a circular orbit. parameters can be fixed to certain values.," For exploring particular solutions such as a circular orbit, parameters can be fixed to certain values."
+ We caution that. as noted by ? that the choice of ecosc and esinco as jump variables implicitly imposes a prior that is proportional to e.," We caution that, as noted by \citet{Ford:2006p7023} that the choice of $e\cos\omega$ and $e\sin\omega$ as jump variables implicitly imposes a prior that is proportional to $e$."
+ This approach thus has a tendency to yield a higher eccentricity than would be obtained with a uniform prior. in cases when e is poorly-constrained by the data.," This approach thus has a tendency to yield a higher eccentricity than would be obtained with a uniform prior, in cases when $e$ is poorly-constrained by the data."
+ A similar argument applies to the useof Vsin/cosp and Vsin/sinf as jump variables. in cases where the impact parameter 1s low and there is a strong degeneracy between VYsin /and f in modelling the Rossiter-McLaughlin effect.," A similar argument applies to the useof $V\sin I \cos \beta$ and $V \sin I \sin \beta$ as jump variables, in cases where the impact parameter is low and there is a strong degeneracy between $V\sin I\,$ and $\beta$ in modelling the Rossiter-McLaughlin effect."
+ In such cases. however. the tendency to overestimate Vsin/ from the Rossiter-McLaughlin effect can effectively be curbed by imposing an independent. spectroscopically-determined vsin7 prior on Vsin7. as we have done here.," In such cases, however, the tendency to overestimate $V \sin I$ from the Rossiter-McLaughlin effect can effectively be curbed by imposing an independent, spectroscopically-determined $v \sin I$ prior on $V \sin I$, as we have done here."
+ In addition to the physical free floating parameters. we need to use one y velocity for each RV set and one normalisation factor for each lighteurve as adjustment parameters.," In addition to the physical free floating parameters, we need to use one $\gamma$ velocity for each RV set and one normalisation factor for each lightcurve as adjustment parameters."
+ These are found by using optimal averaging and optimal scaling., These are found by using optimal averaging and optimal scaling.
+ y velocities represent the mean radial velocity of the star in space with respect to the barycentre of the Solar System., $\gamma$ velocities represent the mean radial velocity of the star in space with respect to the barycentre of the Solar System.
+ Since our analysis had many datasets. the results for these adjustment parameters have been omitted. not adding anything to the discussion.," Since our analysis had many datasets, the results for these adjustment parameters have been omitted, not adding anything to the discussion."
+ During these initial analyses we also fitted an additional acceleration in the form of an RV drift y but on no occasion was 1t significantly different from zero., During these initial analyses we also fitted an additional acceleration in the form of an RV drift $\dot{\gamma}$ but on no occasion was it significantly different from zero.
+ We therefore assumed there was no drift for any of our objects., We therefore assumed there was no drift for any of our objects.
+ We will give upper limits for each star in the following sections., We will give upper limits for each star in the following sections.
+ The MCMC algorithm perturbs the fitting parameters at each step ¢ with a simple formula: where P; is a free parameter. G 1s aGaussian random number of unit standard deviation and zero mean (meaning a Gaussian prior on each parameter). while c is the step size for each parameter.," The MCMC algorithm perturbs the fitting parameters at each step $i$ with a simple formula: where $P_j$ is a free parameter, $G$ is aGaussian random number of unit standard deviation and zero mean (meaning a Gaussian prior on each parameter), while $\sigma$ is the step size for each parameter."
+ A factor f is used to control the chain and ensures that 256c of steps are being accepted via a Metropolis-Hastings algorithm. as recommended in ? to give an optimal exploration of parameter space.," A factor $f$ is used to control the chain and ensures that $25\,\%$ of steps are being accepted via a Metropolis-Hastings algorithm, as recommended in \citet{Tegmark:2004p793} to give an optimal exploration of parameter space."
+ The step size is adapted by doing several initial analyses., The step size is adapted by doing several initial analyses.
+ They are adjusted to produce as small a correlation length as possible., They are adjusted to produce as small a correlation length as possible.
+ Once the value is chosen. it remains fixed.," Once the value is chosen, it remains fixed."
+ Only f fluctuates., Only $f$ fluctuates.
+ A phase of 500000 accepted steps 1s used to make the chain converge., A phase of 000 accepted steps is used to make the chain converge.
+ This 1s detected when the correlation length of each parameter is small and that the average y does not improve anymore (?).., This is detected when the correlation length of each parameter is small and that the average $\chi^2$ does not improve anymore \citep{Tegmark:2004p793}.
+ Then starts the real chain. of 0000 accepted steps. from which results will be extracted.," Then starts the real chain, of 000 accepted steps, from which results will be extracted."
+ This number of steps 1s used as a compromise between computation time and exploration., This number of steps is used as a compromise between computation time and exploration.
+ Statistical tests. notably by comparing y are used to estimate the significance of the results.," Statistical tests, notably by comparing $\chi^2$ are used to estimate the significance of the results."
+ Bayesian penalties acting as prior probability. distribution can be added to y to account for any prior information that we might have on any fitted or derived parameter., Bayesian penalties acting as prior probability distribution can be added to $\chi^2$ to account for any prior information that we might have on any fitted or derived parameter.
+" Stellar mass M, can notably be inserted via a prior in the MCMC in order to propagate its error bars on the planet's mass.", Stellar mass $M_{\star}$ can notably be inserted via a prior in the MCMC in order to propagate its error bars on the planet's mass.
+ We also inserted the found by spectral analysis as priors 1n some of our fits to control how much the fit was dependent on them and the resulting value of Vsin7 and whether this influenced the fitted value of B., We also inserted the found by spectral analysis as priors in some of our fits to control how much the fit was dependent on them and the resulting value of $V\sin I$ and whether this influenced the fitted value of $\beta$.
+ Because of the random nature of an MCMC. sometimes a step with an impact parameter close to zero Is taken.," Because of the random nature of an MCMC, sometimes a step with an impact parameter close to zero is taken."
+ This can cause Vsin/ to wander to unphysical values because of the degeneracy between Vsin/ andfgat lowimpact parameters.," This can cause $V\,\sin\,I$ to wander to unphysical values because of the degeneracy between $V\,\sin\,I$ and $\beta$at lowimpact parameters."
+ This ts controlled by imposing a prior., This is controlled by imposing a prior.
+ The prior values are in Table 2.., The prior values are in Table \ref{tab:SpecParam}. .
+ We use a quadratic limb-darkening law with fixed values for the two limb darkening coefficients appropriate to the stellar effective temperature., We use a quadratic limb-darkening law with fixed values for the two limb darkening coefficients appropriate to the stellar effective temperature.
+ They were extracted for the photometry from tables published in ?.., They were extracted for the photometry from tables published in \citet{Claret:2000p364}. .
+ For the radial velocity, For the radial velocity
+The O star shock is buried so deeply in its wind acceleration region that a small difference in its position considerably alters the preshock velocity.,The O star shock is buried so deeply in its wind acceleration region that a small difference in its position considerably alters the preshock velocity.
+" Hence, as soon as the O star's wind is given some room to accelerate the acquired preshock ram pressure greatly increases, helping the WCR to separate from the O star."," Hence, as soon as the O star's wind is given some room to accelerate the acquired preshock ram pressure greatly increases, helping the WCR to separate from the O star."
+" Subsequently, the higher preshock velocities lead to higher postshock gas temperatures/pressures, which increase the stability of the WCR apex against NTSIs."," Subsequently, the higher preshock velocities lead to higher postshock gas temperatures/pressures, which increase the stability of the WCR apex against NTSIs."
+ The snapshots in Fig., The snapshots in Fig.
+ 7 show the gradual transition from collapse to recovery., \ref{fig:recovery_images} show the gradual transition from collapse to recovery.
+" Notice that the recovery is not instantaneous and, as with the collapse, NTSIs vigorously disrupt the WCR and sporadic collisions of the WCR against the O star occur (e.g. in Fig. 7))."," Notice that the recovery is not instantaneous and, as with the collapse, NTSIs vigorously disrupt the WCR and sporadic collisions of the WCR against the O star occur (e.g. in Fig. \ref{fig:recovery_images}) )."
+" However, as the amplitude of the oscillations increases, the O star's wind is allowed to accelerate to higher velocities before it shocks and, as such, the postshock gas pressure sees a commensurate increase."," However, as the amplitude of the oscillations increases, the O star's wind is allowed to accelerate to higher velocities before it shocks and, as such, the postshock gas pressure sees a commensurate increase."
+ Following ¢=1.10 the WCR maintains a clear separation from the O star.," Following $\phi
+=1.10$ the WCR maintains a clear separation from the O star."
+" On the basis of the wind-wind momentum balance alone, the separation of the WCR from the O star is inevitable (Fig.2))."," On the basis of the wind-wind momentum balance alone, the separation of the WCR from the O star is inevitable \ref{fig:chi}) )."
+" However, by initiating the winds in model B within a radial distance of ~1.15R, cells— we are essentially enforcing a lower limit to the wind ram pressure (i.e. the O star's wind velocity never reaches zero) and influencing the time at which the shocks detach from the O star."," However, by initiating the winds in model B within a radial distance of $\sim
+1.15\;{\rm R_{\ast}}$ cells we are essentially enforcing a lower limit to the wind ram pressure (i.e. the O star's wind velocity never reaches zero) and influencing the time at which the shocks detach from the O star."
+" For instance, in models of Iota Orionis, found that the orbital phase at which the WCR recovered was tied to the size of the annulus used to initiate the stellar wind with a smaller annulus leading to a more prolonged collapse."," For instance, in models of Iota Orionis, found that the orbital phase at which the WCR recovered was tied to the size of the annulus used to initiate the stellar wind with a smaller annulus leading to a more prolonged collapse."
+" In reality, the incoming wind may interfere with the initiation of the wind and prolong a collapse/disruption of the WCR."," In reality, the incoming wind may interfere with the initiation of the wind and prolong a collapse/disruption of the WCR."
+" Hence, although the exact time at which the WCR separates from the O star may be resolution dependent, its occurrence in model B is qualitatively correct."," Hence, although the exact time at which the WCR separates from the O star may be resolution dependent, its occurrence in model B is qualitatively correct."
+ We note that numerical heat conduction between hot tenuous gas and cool dense gas will occur in the simulations2010).., We note that numerical heat conduction between hot tenuous gas and cool dense gas will occur in the simulations.
+" In model B this will affect the temperature of postshock gas, and hence the rate at which it cools."," In model B this will affect the temperature of postshock gas, and hence the rate at which it cools."
+" For instance, during the collapse of the WCR the postshock WR wind will be caused to cool more rapidly and the O star wind will be artificially heated, potentially to X-ray emitting temperatures."," For instance, during the collapse of the WCR the postshock WR wind will be caused to cool more rapidly and the O star wind will be artificially heated, potentially to X-ray emitting temperatures."
+" This effect is also present in model A, but to a lesser extent due to the lower contrasts between adjacent postshock gas (i.e. there are cool clumps in the postshock gas in model B but not in model A)."," This effect is also present in model A, but to a lesser extent due to the lower contrasts between adjacent postshock gas (i.e. there are cool clumps in the postshock gas in model B but not in model A)."
+" However, the conclusions drawn from this work are not affected by the undesirable influence of numerical heat conduction."," However, the conclusions drawn from this work are not affected by the undesirable influence of numerical heat conduction."
+" The collision of the winds in WR22 generates hot (T>10"" K) plasma which emits radiation at X-ray wavelengths, providing a direct probe of the postshock winds, and an indirect probe of the preshock winds."," The collision of the winds in WR22 generates hot $T>10^{7}\;$ K) plasma which emits radiation at X-ray wavelengths, providing a direct probe of the postshock winds, and an indirect probe of the preshock winds."
+" However, a recent analysis of observations of WR22 by revealed a number of discrepancies between predictions from a wind-wind collision model and observations."," However, a recent analysis of observations of WR22 by revealed a number of discrepancies between predictions from a wind-wind collision model and observations."
+" To determine whether the contrasting flow dynamics from models A and B can shed any light on these issues we have performed X-ray radiative transfer calculations on the simulation output, details of which can be found in 2.3.."," To determine whether the contrasting flow dynamics from models A and B can shed any light on these issues we have performed X-ray radiative transfer calculations on the simulation output, details of which can be found in \ref{subsec:xray_emission}."
+" We define our line of sight geometry as follows: the inclination angle, i, is measured against the z-axis (i=0° would view the system from directly above the orbital plane), and the angle 0 is measuredagainst the x- (O star in front at apastron) such that 0 increases in the prograde direction (0=90? would align the line of sight with the negative y-axis)."," We define our line of sight geometry as follows: the inclination angle, $i$, is measured against the $z$ -axis $i=0^{\circ}$ would view the system from directly above the orbital plane), and the angle $\theta$ is measuredagainst the $x$ -axis (O star in front at apastron) such that $\theta$ increases in the prograde direction $\theta =90^{\circ}$ would align the line of sight with the negative $y$ -axis)."
+" We adopt viewing angles of i=85° and 0= 0°, in agreement with the system orientations determined by and (1999).."," We adopt viewing angles of $i=85^{\circ}$ and $\theta = 0^{\circ}$ , in agreement with the system orientations determined by and ."
+"maximum value of 1 if they are exactly the same and no filtering is applied), the maximum value in the final peak heights plot indicates which template best resembles the galaxy spectrum.","maximum value of 1 if they are exactly the same and no filtering is applied), the maximum value in the final peak heights plot indicates which template best resembles the galaxy spectrum."
+" This method takes advantage of the full information contained in the spectrum, instead of being constrained to certain features as in the case of the line-strength indices."," This method takes advantage of the full information contained in the spectrum, instead of being constrained to certain features as in the case of the line-strength indices."
+" This type of full-spectrum fitting technique works in the Fourier space and, as an additional advantage, it does not need for a previous emission correction of the data."," This type of full-spectrum fitting technique works in the Fourier space and, as an additional advantage, it does not need for a previous emission correction of the data."
+ The application of this technique to the derivation of the stellar populations was already introduced by Vazdekis(1999) and Vazdekis&Arimoto(1999) and in this work it is adopted to double-check the results obtained from the analysis of the line-strength indices., The application of this technique to the derivation of the stellar populations was already introduced by \citet{1999ApJ...513..224V} and \citet{1999ApJ...525..144V} and in this work it is adopted to double-check the results obtained from the analysis of the line-strength indices.
+" 'To compute the errors for the cross-correlation method we perform a set of 100 Monte Carlo simulations by randomly adding Gaussian noise to the data, keeping the original S/N ratio."," To compute the errors for the cross-correlation method we perform a set of 100 Monte Carlo simulations by randomly adding Gaussian noise to the data, keeping the original S/N ratio."
+ The 1-σ regions are then overplotted to the cross-correlation peak height curves., The $\sigma$ regions are then overplotted to the cross-correlation peak height curves.
+" In order to test the requirements of this method we run some simulations with the Vazdekisetal.(2010) models by artificially changing the different parameters involved: spectral resolution, spectral range, and S/N ratio, taking as a reference the values measured for our observed spectra."," In order to test the requirements of this method we run some simulations with the \citet{2010MNRAS.404.1639V} models by artificially changing the different parameters involved: spectral resolution, spectral range, and S/N ratio, taking as a reference the values measured for our observed spectra."
+" We find that the S/N of the data obtained for the outer bar is not enough to disentangle its age and metallicity, since the error bars for this case are too large (see Appendix C))."," We find that the S/N of the data obtained for the outer bar is not enough to disentangle its age and metallicity, since the error bars for this case are too large (see Appendix \ref{sec:apcc}) )."
+ This prevents us from providing a more constrained solution for the outer bar through the cross-correlation analysis., This prevents us from providing a more constrained solution for the outer bar through the cross-correlation analysis.
+ We have to note that the results obtained here are aimed to be complementary to those from the line-strength analysis., We have to note that the results obtained here are aimed to be complementary to those from the line-strength analysis.
+" Therefore, we restrict the application of the cross-correlation to the bulge and inner bar."," Therefore, we restrict the application of the cross-correlation to the bulge and inner bar."
+" These results are shown in Figures 7 and 8, and will be analysed in Sections ?? and ??,, respectively."," These results are shown in Figures \ref{fig:ccBB} and \ref{fig:ccIB}, and will be analysed in Sections \ref{sec:bulge} and \ref{sec:gradients}, respectively."
+ Figure 5 shows the analysis of the line-strength indices of the bulge of 3357., Figure \ref{fig:spBB} shows the analysis of the line-strength indices of the bulge of 357.
+" Four age indicators are plotted against several well selected metallicity indicators and the G--band, using features from both the red and blue spectral ranges available for this component."," Four age indicators are plotted against several well selected metallicity indicators and the -band, using features from both the red and blue spectral ranges available for this component."
+" Balmer lines might be affected by emission line contamination, although the amplitude of this emission decays from the top (HG. and H8) to the bottom (Hór) indices of Figure 5.."," Balmer lines might be affected by emission line contamination, although the amplitude of this emission decays from the top $\beta_o$ and $\beta$ ) to the bottom $\delta_F$ ) indices of Figure \ref{fig:spBB}. ."
+" In fact, we find no emission in the Hó feature for the bulge."," In fact, we find no emission in the $\delta$ feature for the bulge."
+" For the case of Ηβο and Hf, where the emission contamination is larger, we find some disagreements in the resulting ages with respect to those derived from Hy and Hd."," For the case of $\beta_o$ and $\beta$, where the emission contamination is larger, we find some disagreements in the resulting ages with respect to those derived from $\gamma$ and $\delta$."
+" Although the emission correction has been performed very carefully (see Section ??)), these inconsistencies are probably due to an underestimation of the emission affecting the H9 feature within the noise level."," Although the emission correction has been performed very carefully (see Section \ref{sec:emission}) ), these inconsistencies are probably due to an underestimation of the emission affecting the $\beta$ feature within the noise level."
+" As an attempt to take this effect into account, the HG and Hf errorbars are obtained by computing the maximum and minimum (but no null) 1-c emission correction."," As an attempt to take this effect into account, the $\beta$ and $\beta_o$ errorbars are obtained by computing the maximum and minimum (but no null) $\sigma$ emission correction."
+" Moreover, we find subtle differences between the ages obtained with the HG, and Hj indices."," Moreover, we find subtle differences between the ages obtained with the $\beta_o$ and $\beta$ indices."
+" These differences might be attributed to either the emission correction, that affects the two definitions in a different manner, or to the greater sensitivity of HG, to the a-enhancement in comparison to that of H8 (Cervantes&Vazdekis 2009).."," These differences might be attributed to either the emission correction, that affects the two definitions in a different manner, or to the greater sensitivity of $\beta_o$ to the $\alpha$-enhancement in comparison to that of $\beta$ \citep{2009MNRAS.392..691C}. ."
+" Indeed, these differences are more evident in the a-enhanced indicators (Mgb and CN), whereas the results from (Fe) or"," Indeed, these differences are more evident in the $\alpha$ -enhanced indicators (Mgb and $_2$ ), whereas the results from $\langle$ $\rangle$ or"
+where ος) is given by the standard. computation.,where $\phi_k(\eta)$ is given by the standard computation.
+ We just add the extra contributions due to Eq.(21)) to the sources from) which to compute the correlators as deseribed iu the previous section., We just add the extra contributions due to \ref{extra}) ) to the sources from which to compute the correlators as described in the previous section.
+ Therefore it becomes straightforward to determine the features in the CMD sky in this prescription., Therefore it becomes straightforward to determine the features in the CMB sky in this prescription.
+ This parametrization describes a gradieut-tvpoe nodification of the effective. CAIB sources., This parametrization describes a gradient-type modification of the effective CMB sources.
+ We note hat ? has determined the implicatious of a supcr-torizon perturbation. and ? considered such a spatial variation of the curvaton field at inflation.," We note that \cite{Tangen:2009mw} has determined the implications of a super-horizon perturbation, and \cite{Erickcek:2008sm} considered such a spatial variation of the curvaton field at inflation."
+ Iu our nodel. the anisotropy is formed dvuamically and )oecomies niportaut with the dominance of dark cuerey.," In our model, the anisotropy is formed dynamically and becomes important with the dominance of dark energy."
+ Explicitly we have a spontaneous modification of the effective CMD sources through the eradicut operator as ollows: e((x) = Jotx)(25) when |n]=zulpprE/paj).," Explicitly, we have a spontaneous modification of the effective CMB sources through the gradient operator as follows: ) = ] ) when $|n| = \varpi_0 (\rho_{DE}/\rho_M)$."
+ This ΠΡ to shifting the Fourier modes of the perturbations exactly as prescribed in (20)) aud (23)). = Since the anisotropic part develops as a result of he evolution of the universe. it is a property of the rauster functions aud not of the primordial spectrum of of perturbations.," This amounts to shifting the Fourier modes of the perturbations exactly as prescribed in \ref{psi}) ) and \ref{varpi}) ), = Since the anisotropic part develops as a result of the evolution of the universe, it is a property of the transfer functions and not of the primordial spectrum of of perturbations."
+ One does not expect odd Af couplings roni primordial origi. since they violate paritv.," One does not expect odd $\Delta\ell$ couplings from primordial origin, since they violate parity."
+ The xesence of the nuaginary unit is necessarv for reality of the perturbations. which cau be checked as follows.," The presence of the imaginary unit is necessary for reality of the perturbations, which can be checked as follows."
+ Since the physical perturbation is a convolution of he primordial physical perturbation aud the Fourier ransformation of the trausfer function. oue notes that he latter should also be real.," Since the physical perturbation is a convolution of the primordial physical perturbation and the Fourier transformation of the transfer function, one notes that the latter should also be real."
+ We ect that where off) is the (veal. isotropic) transfer functiou which depends ouly on the maenitude of the wavevector. and the right hand side is the anisotropic transfer function in the configuration space that retains it’s reality.," We get that where $\phi(k)$ is the (real, isotropic) transfer function which depends only on the magnitude of the wavevector, and the right hand side is the anisotropic transfer function in the configuration space that retains it's reality."
+ To recap our approach. we introduced a mismatch of the two eravitational potentials iu the Newtonian gauge.," To recap our approach, we introduced a mismatch of the two gravitational potentials in the Newtonian gauge."
+ This nismatch. quantified by ze is eiven by a eracdicut along a preferred axis ».," This mismatch, quantified by $\varpi$, is given by a gradient along a preferred axis $\hat{n}$."
+ Iu the following. we will also allow scale dependence of this effect by introduci the spectral iudex ρα.," In the following, we will also allow scale dependence of this effect by introducing the spectral index $n_{\mathrm{de}}$."
+ This paranmetnüsation can then be used to coustrain the presence of the such eradicuts iu he late universe. since they would be seen iu the CMD (practically oulv) through their impact on the time-evolution of the gravitational potentials because of the ISW effect.," This parametrisation can then be used to constrain the presence of the such gradients in the late universe, since they would be seen in the CMB (practically only) through their impact on the time-evolution of the gravitational potentials because of the ISW effect."
+ Phvsicallv. these exadieuts could be caused *oa laree-scale inhomoecucity cuterime our horizon. spontaneous formation due to ce. coherent magnetic ficlds or simply the possible imperfect nature of the dark enerev field.," Physically, these gradients could be caused by a large-scale inhomogeneity entering our horizon, spontaneous formation due to e.g. coherent magnetic fields or simply the possible imperfect nature of the dark energy field."
+ To clarify the difference of our approach to all xevious literature. let us iieution that odd modulations mav be considered to occur at three clistinet levels.," To clarify the difference of our approach to all previous literature, let us mention that odd modulations may be considered to occur at three distinct levels."
+ The temperature field itself cau be modulated. for a recent example see 7..," The temperature field itself can be modulated, for a recent example see \cite{Aluri:2011wv}."
+ This would effectively describe sole systematics in the data., This would effectively describe some systematics in the data.
+ Strangely enough. the primordial spectrum itself could contain parity violating coutrbution ?..," Strangely enough, the primordial spectrum itself could contain parity violating contribution \cite{Koivisto:2010fk}."
+ That is cousisteut only iu the coutext of noncommutative quanti field theory. aud thus provides a unique signal for such ligh energy modifications of the standard model ?..," That is consistent only in the context of noncommutative quantum field theory, and thus provides a unique signal for such high energy modifications of the standard model \cite{Groeneboom:2010fn}."
+ Finally. the cosinological structures nav evolve statistically oddly. which is the case we focus upon here.," Finally, the cosmological structures may evolve statistically oddly, which is the case we focus upon here."
+ Iu this section we describe in detail the confrontation of the model with the WALAP data., In this section we describe in detail the confrontation of the model with the WMAP data.
+ We will now discuss the method used to obtain the set of parameters which gives the best fit between our model aud the observations., We will now discuss the method used to obtain the set of parameters which gives the best fit between our model and the observations.
+ Oi basis is the evaluation of the Likelihood fuuctiou ina Lcdiaensioual parameter space., Our basis is the evaluation of the Likelihood function in a 4-dimensional parameter space.
+ After explaining the eeneral likelihood. procedure. we explain m more detail the different steps taken to calculate and maximize the likelihood.," After explaining the general likelihood procedure, we explain in more detail the different steps taken to calculate and maximize the likelihood."
+ Given a set of data [d;] our goal is to find the set of parameters which maximizes the posterior., Given a set of data $\{d_i\}$ our goal is to find the set of parameters which maximizes the posterior.
+ For case of notation let a=(6.0.mp.nq.)(hh.cy.Mac) denote the set of paraimcters to be deteriuned.," For ease of notation let $\alpha = (\theta,\phi,\varpi_0,n_{\mathrm{de}}) = (\hbn,\varpi_0,n_{\mathrm{de}})$ denote the set of parameters to be determined."
+ By Baves’ theorem we know that the posterior distribution Piald)xP(djajPla)=Z(o)P(o) where £(o) is the ikchhood and (a) is a prior.," By Bayes' theorem we know that the posterior distribution $P(\alpha|\bfd) \propto P(\bfd |\alpha)
+P(\alpha) = \mL(\alpha)P(\alpha)$ where $\mL(\alpha)$ is the likelihood and $P(\alpha)$ is a prior."
+ We take a conservative approach and assume that we know nothing prior about he anisotropic parameters. aud thus P(a|d)=La) up o a normalization constant.," We take a conservative approach and assume that we know nothing prior about the anisotropic parameters, and thus $P(\alpha|\bfd) = \mL(\alpha)$ up to a normalization constant."
+ If we manage to compute he likelihood function iu all of the parameter space then we automatically have the posterior distribution aud our job is essentially. done., If we manage to compute the likelihood function in all of the parameter space then we automatically have the posterior distribution and our job is essentially done.
+" Although our model is anisotropic we still assume that he underlying distribution is Caussian. and thus the ikelihood is ο where the data VdetCla)vector d cousists of the a,,, of the observed masked map."," Although our model is anisotropic we still assume that the underlying distribution is Gaussian, and thus the likelihood is ) where the data vector $\bfd$ consists of the $a_{\ell m}$ of the observed masked map."
+ The correlation matrix C=S|N is the stun of the CAIB signal covariance matiix S aud he noise covariance matrix IN., The correlation matrix $\bfC=\bfS+\bfN$ is the sum of the CMB signal covariance matrix $\bfS$ and the noise covariance matrix $\bfN$.
+" Our analysis is performed in harmonic space where the signal covariance 5,42""=(mye) is computed frou Eq. (12))", Our analysis is performed in harmonic space where the signal covariance $S_{\ell m;\ell' m'} = \langle a_{\ell m}a^*_{\ell' m'}\rangle$ is computed from Eq. \ref{corr}) )
+ and thus coutains ron-diagonal anisotropic contributions from dark energy., and thus contains non-diagonal anisotropic contributions from dark energy.
+ This matrix gives the dependence on the cosmological xuameters., This matrix gives the dependence on the cosmological parameters.
+" Finally the observed data vector d may be written iu jurluoldes space as —bw,,. τα where b; is the imstrumneutal beam. wy is the pixel window fiction and nj, is the ΚΙ ΤΠΕ teri."," Finally the observed data vector $\bfd$ may be written in harmonics space as = + where $_\ell$ is the instrumental beam, $_\ell$ is the pixel window function and $_{\ell m}$ is the (Gaussian) noise term."
+ Since there is no correlation between the signal aud the noise we lave — Ε, Since there is no correlation between the signal and the noise we have = +
+ Since there is no correlation between the signal aud the noise we lave — Επ, Since there is no correlation between the signal and the noise we have = +
+ Since there is no correlation between the signal aud the noise we lave — Επ., Since there is no correlation between the signal and the noise we have = +
+frequency values were obtained for more than 5 thousand acoustic mode frequencies νι.,"frequency values were obtained for more than 5 thousand acoustic mode frequencies $\nu_{l,n}$."
+" This achievement is the result of the combined efforts of several observational networks, such as BISON(Chaplinetal.1996) and GONG (Harveyetal.1996),, followed by the helioseismic experiments from the SoHO mission: GOLF(Gabrieletal. 1995),, VIRGO(Fróhlichetal. and MDI-SOI (Scherreretal.1995)."," This achievement is the result of the combined efforts of several observational networks, such as \citep{art-BISON1996} and GONG \citep{art-GONG1996Sci}, followed by the helioseismic experiments from the SoHO mission: \citep{art-GOLF1995}, \citep{art-VIRGO1995} and MDI-SOI \citep{art-MDI1995}."
+. Seismology1995) has provided the most powerful tool to probe the Sun’s interior to date., Seismology has provided the most powerful tool to probe the Sun's interior to date.
+" Accurate frequency measurements of large numbers acoustic modes have been made, including the radial and dipole global acoustic modes which penetrate more deeply into the core."," Accurate frequency measurements of large numbers acoustic modes have been made, including the radial and dipole global acoustic modes which penetrate more deeply into the core."
+" Unfortunately, the diagnostics of the Sun’s core provided by these acoustic modes are still insufficient."," Unfortunately, the diagnostics of the Sun's core provided by these acoustic modes are still insufficient."
+" However, the possible discovery of gravity modes by current ongoing experiments(Garcíaetal.2007) or future ones could be a breakthrough towards fully understanding the physics of the Sun's core and establishing the possible existence of dark matter."," However, the possible discovery of gravity modes by current ongoing \citep{art-Garcia2007Sci} or future ones could be a breakthrough towards fully understanding the physics of the Sun's core and establishing the possible existence of dark matter."
+" The impact of dark matter particles in the Sun's core (Cox and its consequences onthe acoustic spectrum 1991)were previously analysed by Lopesetal.(2002b) among others, and more recently by Cumberbatch and Taosoetal."," The impact of dark matter particles in the Sun's core \citep{art-Coxetal1990ApJ...353..698C,art-Dearbornetal1990,art-Kaplanetal1991}
+ and its consequences onthe acoustic spectrum were previously analysed by \citet{art-LopesSH2002} among others, and more recently by \citet{art-Cumberbatch2010} and \citet{art-Taoso2010}."
+" In this Letter, we (2010)..have computed a series of solar models evolving within dark matter halos for which we have estimated the impact of dark matter particles in the present-day Sun."," In this Letter, we have computed a series of solar models evolving within dark matter halos for which we have estimated the impact of dark matter particles in the present-day Sun."
+" As a reference, we use an updated solar standard model that provides a seismic diagnostic similar to other solar standard models found in the literature."," As a reference, we use an updated solar standard model that provides a seismic diagnostic similar to other solar standard models found in the literature."
+" Furthermore, we have computed the gravity mode oscillation spectrum for such dark matter scenarios."," Furthermore, we have computed the gravity mode oscillation spectrum for such dark matter scenarios."
+" 'The capture rate of dark matter particles in the Sun was computed in detail by taking into account their scattering with the different chemical species of the solar plasma (Gould1987;Gondoloetal.2004;Casanellas&Lopes 2010,b)."," The capture rate of dark matter particles in the Sun was computed in detail by taking into account their scattering with the different chemical species of the solar plasma \citep{art-Gould1987,art-GondoloEdsjoDarkSusy2004,art-CL2010procIU,art-astrophCL}."
+". The presence of dark matter inside the star can change its evolution by two mechanisms: (i) changing the transport of energy inside the star (Spergel&Press1985;GouldRaffelt1990a),, or (ii) by contributing to the energy production that sustains the gravitational collapse of the star (Ioccoetal.2008;Freeseetal.2008;Casanellas&Lopes 2009)."," The presence of dark matter inside the star can change its evolution by two mechanisms: (i) changing the transport of energy inside the star \citep{art-SpergelPress1985,art-GouldRaffelt1990a}, or (ii) by contributing to the energy production that sustains the gravitational collapse of the star \citep{art-Ioccoetal2008,art-FreeseBSG08,art-CasanellasLopes2009}."
+". In the case of the Sun, the latter mechanism is negligible (Lopesetal."," In the case of the Sun, the latter mechanism is negligible \citep{art-LopesSH2002}."
+" The efficiency of the energy transport by dark matter 2002b)..particles depends, among other factors, on the scattering cross-sections of dark matter particles on baryons."," The efficiency of the energy transport by dark matter particles depends, among other factors, on the scattering cross-sections of dark matter particles on baryons."
+" Dark matter particles might have a total scattering cross-section σι with several components: spin-independent interactions with all chemical elements in the Sun, and spin-dependent interactions with hydrogen."," Dark matter particles might have a total scattering cross-section $\sigma_t$ with several components: spin-independent interactions with all chemical elements in the Sun, and spin-dependent interactions with hydrogen."
+ The presence of such particles contributes to the local energy transport., The presence of such particles contributes to the local energy transport.
+" The efficiency of the heat transport by dark matter particles is dependent on the value σε, relative to the fiducially critical cross-section Oc=(my/M)R?~ 810?9cm?, where mj is the proton mass and M and R are the mass and radius of the Sun."," The efficiency of the heat transport by dark matter particles is dependent on the value $\sigma_t$, relative to the fiducially critical cross-section $\sigma_c = (m_p/M) R^2\sim$ $8\;10^{-36} {\rm cm^2}$, where $m_p$ is the proton mass and $M$ and $R$ are the mass and radius of the Sun."
+ The optimal heat transport depends on the dark matter particle scattering cross-section σι., The optimal heat transport depends on the dark matter particle scattering cross-section $\sigma_t$.
+ For particles with small mean free paths (diffusion regime) or with large cross-sections (σι> the conductivity falls as o-/o+.," For particles with small mean free paths (diffusion regime) or with large cross-sections $\sigma_t > \sigma_c$ ), the conductivity falls as $\sigma_c/\sigma_t$."
+" Similarly, for particles σε),with large mean free path (Knudsen regime) or lower cross-sections (a;« σε), the conductivity falls as "," Similarly, for particles with large mean free path (Knudsen regime) or lower cross-sections $\sigma_t < \sigma_c$ ), the conductivity falls as $\sigma_t /\sigma_c$ ."
+"Furthermore, a suppression factoris required in this σε/σ..latter transport regime (Gould 1990a,b).."," Furthermore, a suppression factoris required in this latter transport regime \citep{art-GouldRaffelt1990a,art-GouldRaffelt1990b}. ."
+ Both regimes are included in our stellar evolution code, Both regimes are included in our stellar evolution code
+"still concerned with finding unambiguous vortex identification schemes (e.g.,??)..","still concerned with finding unambiguous vortex identification schemes \citep[e.g.,][]{2004Haller,2007Kolar}."
+" Given its frequent use in the literature, it should be noted that a high vorticity w=sxi is not a sufficient indicator for the presence of a vortex, because w can also be high in shear flows without rotation."," Given its frequent use in the literature, it should be noted that a high vorticity $\omega = \abs{\Nabla \times \vec{v}}$ is not a sufficient indicator for the presence of a vortex, because $\omega$ can also be high in shear flows without rotation."
+" To detect vortices, we use the so-called “swirling strength”, or Ag, criterion (?),, which is based on an eigenanalysis of the velocity gradient tensor U;;:=(0;v;)."," To detect vortices, we use the so-called “swirling strength”, or $\lci$, criterion \citep{1999Zhou}, which is based on an eigenanalysis of the velocity gradient tensor $U_{ij}\coloneqq(\partial_j v_i)$."
+ Vortices are defined to be regions where U has a pair of complex conjugate eigenvalues., Vortices are defined to be regions where $U$ has a pair of complex conjugate eigenvalues.
+" A large unsigned imaginary part λα implies a region of “strong swirling”, i.e., (part of) a vortex."," A large unsigned imaginary part $\lci$ implies a region of “strong swirling”, i.e., (part of) a vortex."
+" In the case of rigid rotation, the value of τς:=2z/A gives the revolution period of the rotating flow."," In the case of rigid rotation, the value of $\taus
+\coloneqq 2\pi/\lci$ gives the revolution period of the rotating flow."
+" Unlike many other methods for vortex detection, the A, criterion is applicable also in the case of compressible hydrodynamics (?).."," Unlike many other methods for vortex detection, the $\lci$ criterion is applicable also in the case of compressible hydrodynamics \citep{2009Kolar}."
+" Depending on the signs of the real eigenvalues 4, and the real part of the complex eigenvalues A4, we discriminate four types of vortices (?):: Here, converging and diverging refer to the fluid motion along the direction of the eigenvector corresponding to 4;, which is identified as the direction of the vortex."," Depending on the signs of the real eigenvalues $\lr$ and the real part of the complex eigenvalues $\lcr$, we discriminate four types of vortices \citep{1999Haimes}: Here, converging and diverging refer to the fluid motion along the direction of the eigenvector corresponding to $\lr$, which is identified as the direction of the vortex."
+" Note that this direction is not necessarily orthogonal to the “swirling plane"" spanned by the real and imaginary parts of the complex eigenvectors.", Note that this direction is not necessarily orthogonal to the “swirling plane” spanned by the real and imaginary parts of the complex eigenvectors.
+ We distinguish vortices by their inclination angle L With respect to the vertical direction., We distinguish vortices by their inclination angle $\iota$ with respect to the vertical direction.
+" The identification of vortices via A; is Galilean invariant, because it only relies on derivatives of the velocity field."," The identification of vortices via $\lci$ is Galilean invariant, because it only relies on derivatives of the velocity field."
+" Streamlines, however, are not invariant under Galilean transformations."," Streamlines, however, are not invariant under Galilean transformations."
+" Hence the streamlines in a particular frame of reference (e.g., the one which is at rest with respect to the computational box) do not always show circular orbits where vortices are detected."," Hence the streamlines in a particular frame of reference (e.g., the one which is at rest with respect to the computational box) do not always show circular orbits where vortices are detected."
+" A strong vortex, however, is expected to produce vortical streamlines for a large range of observer velocities."," A strong vortex, however, is expected to produce vortical streamlines for a large range of observer velocities."
+" Although the definition of the swirling strength is purely local (as is the definition of vorticity), it can be shown that for a volume of rigidly rotating fluid, the swirling strength is uniform over the entire interior of the volume."," Although the definition of the swirling strength is purely local (as is the definition of vorticity), it can be shown that for a volume of rigidly rotating fluid, the swirling strength is uniform over the entire interior of the volume."
+" In all our simulations, the regions of strong swirl form a network of highly tangled filaments."," In all our simulations, the regions of strong swirl form a network of highly tangled filaments."
+ An illustrative example of this is shown in Fig. 1.., An illustrative example of this is shown in Fig. \ref{fig:Im3D}.
+" Some of the filaments protrude from the optical surface, forming variously shaped vortex tubes with different inclinations, sometimes bridging horizontally separated points in an arc."," Some of the filaments protrude from the optical surface, forming variously shaped vortex tubes with different inclinations, sometimes bridging horizontally separated points in an arc."
+ The optical surface is usually depressed at the footpoints of vertically emerging vortex tubes., The optical surface is usually depressed at the footpoints of vertically emerging vortex tubes.
+" The horizontal distribution of the swirls is not homogeneous, the granular upflows being mostly devoid"," The horizontal distribution of the swirls is not homogeneous, the granular upflows being mostly devoid"
+Kwee&vanWoerden(1956) method for all these minima.,\cite{Kwee} method for all these minima.
+ The mean linear light elements suitable for observations are which were also used for deriving the proper epochs and types of the minima times written in Table 1.., The mean linear light elements suitable for observations are which were also used for deriving the proper epochs and types of the minima times written in Table \ref{Minima}.
+ If we plot these data points in the ο—C diagram. the difference between primary and secondary minima ts clearly visible.," If we plot these data points in the $O-C$ diagram, the difference between primary and secondary minima is clearly visible."
+ Following the method of apsidal motion analysis as described in Giménez&Garcfa-Pelayo(1983).. we tried the computation with orbital inclination ¢=90° for the first time.," Following the method of apsidal motion analysis as described in \cite{1983Ap&SS..92..203G}, we tried the computation with orbital inclination $i=90^\circ$ for the first time."
+" After that we used for the second attempt the melination i=87.88"" (see Sect. 3)."," After that we used for the second attempt the inclination $i=87.88^\circ$ (see Sect. \ref{LCanal}) ),"
+ which resulted in almost the same parameters (owing to the term cot(D. which is nearly 0 for our inclination).," which resulted in almost the same parameters (owing to the term $\cot^2(i)$, which is nearly 0 for our inclination)."
+ In Fig., In Fig.
+ 1. we plot the final fit with the apsidal motion hypothesis on all used data points., \ref{Fig3} we plot the final fit with the apsidal motion hypothesis on all used data points.
+ This leads to the parameters of the motion of apsides given in Table 2.., This leads to the parameters of the motion of apsides given in Table \ref{OCparam1}.
+ All uncertainties of the parameters were calculated from the covariance matrix. of the fit and from the uncertainty of the inclination., All uncertainties of the parameters were calculated from the covariance matrix of the fit and from the uncertainty of the inclination.
+" Evidently the apsidal period ts very short. about only 23 yr. which places this system among one those with the shortest apsidal motion,"," Evidently the apsidal period is very short, about only 23 yr, which places this system among one those with the shortest apsidal motion."
+ The system V490 Cyg has much lower published times of minima observations., The system V490 Cyg has much lower published times of minima observations.
+ The first rough. times-of-minima estimates are those by Wachmann(1948) from his photometry in the 1930°s. but these have such a large scatter that they cannot be used for any reliable analysis.," The first rough times-of-minima estimates are those by \cite{Wachmann2} from his photometry in the 1930's, but these have such a large scatter that they cannot be used for any reliable analysis."
+ However. he also noticed that the secondary minimum is not symmetric with regards to the primary one.," However, he also noticed that the secondary minimum is not symmetric with regards to the primary one."
+ Nevertheless. a possible eccentricity and apsidal motion have never been studied since then.," Nevertheless, a possible eccentricity and apsidal motion have never been studied since then."
+" The more precise photoelectric and CCD observations have been measured since 1999, but there are only 12 published minima."," The more precise photoelectric and CCD observations have been measured since 1999, but there are only 12 published minima."
+ A few new CCD observations were obtained in the Ondieejov observatory with the same telescope as for V456 Oph. and two new minima times were also derived from the INTEGRAL/OMC data.," A few new CCD observations were obtained in the Ondřeejov observatory with the same telescope as for V456 Oph, and two new minima times were also derived from the INTEGRAL/OMC data."
+ The new measurements and the already published ones are presented in Table 3.., The new measurements and the already published ones are presented in Table \ref{Minima2}.
+ The suitable linear ephemerides for observations are The minima times presented in Table 3. were used for the period analysis. which we did by applying the apsidal motion hypothesis.," The suitable linear ephemerides for observations are The minima times presented in Table \ref{Minima2} were used for the period analysis, which we did by applying the apsidal motion hypothesis."
+ The only difference in analysis between V490 Cyg and V456 Oph was the assumption of an inclination /=90° for V490 Cyg because we had no light curve analysis., The only difference in analysis between V490 Cyg and V456 Oph was the assumption of an inclination $i=90^\circ$ for V490 Cyg because we had no light curve analysis.
+ The difference between primary and secondary is clearly visible. reaching up to 47 minutes. which is surprisingly high for a binary with such a short orbital period.," The difference between primary and secondary is clearly visible, reaching up to 47 minutes, which is surprisingly high for a binary with such a short orbital period."
+ The analysis led to the parameters of the apsidal motion presented in Table 2.., The analysis led to the parameters of the apsidal motion presented in Table \ref{OCparam1}.
+ Obviously the resulting value of the apsidal motion period of about 18.8 years is even shorter than for V456 Oph: we, Obviously the resulting value of the apsidal motion period of about 18.8 years is even shorter than for V456 Oph; we
+of ους respectively.,of $R_{200}$ respectively.
+ We expect the central black hole to be located. at the center of the main galaxy in the halo since this is the place of highest gas density and thus maximal black hole growth., We expect the central black hole to be located at the center of the main galaxy in the halo since this is the place of highest gas density and thus maximal black hole growth.
+" The fact that most ""central AGN’ are found. near the group. center shows that. our technique for identifving the central ACN is. reasonably adequate.", The fact that most `central AGN' are found near the group center shows that our technique for identifying the central AGN is reasonably adequate.
+ Due to the limited. volume of our. simulation we could only probe the faint end of the AGN luminosity function (?).., Due to the limited volume of our simulation we could only probe the faint end of the AGN luminosity function \citep{degrafetal10}.
+ For calculating the mean occupation we stucy AGN samples with bolometric luminosities above 1077 ergs/s. corresponding to the observational limit of X-ray selected ACN in deep surveys (e.g...227)...," For calculating the mean occupation we study AGN samples with bolometric luminosities above $10^{42}$ ergs/s, corresponding to the observational limit of X-ray selected AGN in deep surveys \citep[e.g.,][]{uedaetal03, boyleetal93, cowieetal03}."
+ We also show the distribution of AGN with luminosities above 107 eres/s and 107 ergs/s for theoretical interests to see the Luminosity dependence of HOD. parameters., We also show the distribution of AGN with luminosities above $10^{40}$ ergs/s and $10^{38}$ ergs/s for theoretical interests to see the luminosity dependence of HOD parameters.
+ In à companion paper. ? (hereafter Paper D). we use he same SPI simulation to characterize the HOD of black roles as à function of black hole mass.," In a companion paper, \citet{degrafetal11b} (hereafter Paper I), we use the same SPH simulation to characterize the HOD of black holes as a function of black hole mass."
+ Our present study is fundamentally different from that in Paper lL. We aim o study the HOD of AGN in terms of ACN: luminosity., Our present study is fundamentally different from that in Paper I. We aim to study the HOD of AGN in terms of AGN luminosity.
+ Dx definition the HOD is always studied based on physical oxoperties of the object., By definition the HOD is always studied based on physical properties of the object.
+ In the case of galaxy clustering. the ealaxy HOD has been extensively studied in the literature xwed on color. Luminosity. stellar mass (see ?.. 2)).," In the case of galaxy clustering, the galaxy HOD has been extensively studied in the literature based on color, luminosity, stellar mass (see \citealt{zehavietal05}, \citealt{zhengetal07}) )."
+ Black vole mass and luminosity (accretion rates) are clilferent »xwameters and there does not exist an obvious one-to-one correspondence between mass and accretion rates (e.g. 11)., Black hole mass and luminosity (accretion rates) are different parameters and there does not exist an obvious one-to-one correspondence between mass and accretion rates (e.g. 1).
+ Additional barvonie physies is needed to eo from black dole mass to aceretion rate (and thus luminosity)., Additional baryonic physics is needed to go from black hole mass to accretion rate (and thus luminosity).
+ There is some correlation between mass and accretion rates. but he scatter is huge and the correlation depends strongly on redshift (2?)..," There is some correlation between mass and accretion rates, but the scatter is huge and the correlation depends strongly on redshift \citep{c&d08,chatetal08}."
+ In practice. it is barely. possible to have accurate mass measurement. of a [arge sample of black holes bevond. the local universe.," In practice, it is barely possible to have accurate mass measurement of a large sample of black holes beyond the local universe."
+ Therefore. the study in. Paper Lis mostly of theoretical interest.," Therefore, the study in Paper I is mostly of theoretical interest."
+ By investigating the occupation properties of black holes at the subhalo level. Paper 1 develops an analytic mechanism to populate cark matter halos with black holes in simulation.," By investigating the occupation properties of black holes at the subhalo level, Paper I develops an analytic mechanism to populate dark matter halos with black holes in simulation."
+ On the contrary. AGN luminosity is more readily measured. observationally and there are existing and ongoing ellorts in measuring Αλομαδας clustering as a function. of. luminosity.," On the contrary, AGN luminosity is more readily measured observationally and there are existing and ongoing efforts in measuring AGN/quasar clustering as a function of luminosity."
+ Our study in this paper will provide the necessary ingredient in interpreting these observations., Our study in this paper will provide the necessary ingredient in interpreting these observations.
+ Similar to the mass function of black holes and luminosity function of AGN. the black-hole-mass based LLIOD and ACN-luminosity based HOD are complementary tools to advance our understanding of black hole evolution.," Similar to the mass function of black holes and luminosity function of AGN, the black-hole-mass based HOD and AGN-luminosity based HOD are complementary tools to advance our understanding of black hole evolution."
+ Since the goal of this paper is to. provide the LOD framework to interpret observations we adopt some technical dillerences in the definition of dark. matter. halos and the choice of our. sample from Paper L1. In. Paper 1: we differentiate between central and satellite black holes at the subhalo level., Since the goal of this paper is to provide the HOD framework to interpret observations we adopt some technical differences in the definition of dark matter halos and the choice of our sample from Paper I. In Paper I we differentiate between central and satellite black holes at the subhalo level.
+ Black holes residing in the central subgroup are called central black holes and. black holes residing in satellite subgroups are called satellite black holes., Black holes residing in the central subgroup are called central black holes and black holes residing in satellite subgroups are called satellite black holes.
+ In. the nomenclature in Paper E there can be multiple central black holes., In the nomenclature in Paper I there can be multiple central black holes.
+ We emphasize that. the model in Paper LE will be useful for populating halos (N body simulations) with black holes using semi-analvtic approaches and for studying the clistribution of black holes in cosmological simulations., We emphasize that the model in Paper I will be useful for populating halos (N body simulations) with black holes using semi-analytic approaches and for studying the distribution of black holes in cosmological simulations.
+ On the other hand. our approach in this paper is to perform a theoretical study on the relation between ACN (selected xwed on Luminosity) ancl dark matter halos and. examine he assumptions in the HOD for modeling AGN clustering data., On the other hand our approach in this paper is to perform a theoretical study on the relation between AGN (selected based on luminosity) and dark matter halos and examine the assumptions in the HOD for modeling AGN clustering data.
+ Our identification of central (1.0. most. massive black role within Mogg) and satellite (rest of the black holes within Royy) AGN is equivalent to primary ane secondary black roles in Paper I (see 22 of Paper 1) and is similar to 1e standard terminology. usec for galaxy HOD., Our identification of central (i.e. most massive black hole within $R_{200}$ ) and satellite (rest of the black holes within $R_{200}$ ) AGN is equivalent to primary and secondary black holes in Paper I (see 2 of Paper I) and is similar to the standard terminology used for galaxy HOD.
+ We note iu the friends-oí-Driends. halos were used in. Paper E for μαudving black hole occupation., We note that the friends-of-friends halos were used in Paper I for studying black hole occupation.
+ Llere we choose {ου to cline the halo boundary. which is commonly. adopted. in clinine galaxy clusters.," Here we choose $R_{200}$ to define the halo boundary, which is commonly adopted in defining galaxy clusters."
+ Since the LOD describes bias at 16 level of svstems near dynamical equilibrium. traditional 7ürial methods of mass estimation can take advantage of —us and provide a direct. measurement of the probability istribution PONV[AZ) (2)...," Since the HOD describes bias at the level of systems near dynamical equilibrium, traditional virial methods of mass estimation can take advantage of this and provide a direct measurement of the probability distribution $P(N|M)$ \citep{b&w02}."
+ We have verified that the two ilferent halo definitions show a good agreement in assigning black holes to host halos within 1054., We have verified that the two different halo definitions show a good agreement in assigning black holes to host halos within $10\%$.
+ In this section we present the distribution of accretion rates as à function of halo mass ancl black hole mass., In this section we present the distribution of accretion rates as a function of halo mass and black hole mass.
+ From this we calculate the conditional luminosity functions and. the mean occupation distribution of central and satellite AGN., From this we calculate the conditional luminosity functions and the mean occupation distribution of central and satellite AGN.
+ Finally we show the racial cistribution of the satellite AGN ancl provide the best-fit hod. parameters., Finally we show the radial distribution of the satellite AGN and provide the best-fit hod parameters.
+ Figure 1 shows the distribution of the LEclclington scaled accretion rate (y=Al Aliepp). as a function of halo mass for three recdshifts.," Figure 1 shows the distribution of the Eddington scaled accretion rate $\chi = \dot{M}/\dot{M}_{\rm {EDD}}$ ), as a function of halo mass for three redshifts."
+ Phe black points are central AGN and the red. open circles. represent satellites., The black points are central AGN and the red open circles represent satellites.
+ The mean rate is roughly independent. of halo mass. which is commonly assumed in scnii-analytic studies.," The mean rate is roughly independent of halo mass, which is commonly assumed in semi-analytic studies."
+ However. the mean value of N ata given halo mass varies with redshift.," However, the mean value of $\chi$ at a given halo mass varies with redshift."
+ AGN at higher redshifts (left. panel of Ll showing += 5.0) tend to accrete at a higher value than their low redshift counterparts (right panel of 11 showing z= 1.0)., AGN at higher redshifts (left panel of 1 showing $z=5.0$ ) tend to accrete at a higher value than their low redshift counterparts (right panel of 1 showing $z=1.0$ ).
+ Figure 2 shows the distribution of X as a function of black hole mass., Figure 2 shows the distribution of $\chi$ as a function of black hole mass.
+ It is defined. as the conditional probability distribution of x for a given black hole mass P(x[AMpg)., It is defined as the conditional probability distribution of $\chi$ for a given black hole mass $P(\chi|M_{\rm{BH}})$.
+ The solid and the dotted lines denote central ancl satellite AGN. respectively.," The solid and the dotted lines denote central and satellite AGN, respectively."
+ The mean rates are again roughly independent: of black hole mass but vary substantially with redshift., The mean rates are again roughly independent of black hole mass but vary substantially with redshift.
+ Phe central ancl satellite clistributions tend to trace each other at all redshifts., The central and satellite distributions tend to trace each other at all redshifts.
+ In 33 we present the spatial distribution of the satellite ACN., In 3 we present the spatial distribution of the satellite AGN.
+ The sizes of the circles are. proportional to the black hole mass (left panel). gas mass (middle panel). ancl stellar mass (right panel) respectively.," The sizes of the circles are proportional to the black hole mass (left panel), gas mass (middle panel), and stellar mass (right panel) respectively."
+ “Phe eas and the stellar mass are computed within a spherical region of radius 10 kpe surrounding the black hole., The gas and the stellar mass are computed within a spherical region of radius $10$ kpc surrounding the black hole.
+ Phe green triangles represent black holes with zero stellar mass within 10 kpc., The green triangles represent black holes with zero stellar mass within $10$ kpc.
+ The top. middle. and bottom panels are for z=1.0. 3.0. and 5.0 respectively.," The top, middle, and bottom panels are for $z=1.0$, $3.0$, and $5.0$ respectively."
+ Phe AGN accretion rates do not show any dependence on radius or black hole mass., The AGN accretion rates do not show any dependence on radius or black hole mass.
+ Also we do not see any correlation between black hole mass and --s location within the halo., Also we do not see any correlation between black hole mass and its location within the halo.
+ The trends of accretion rates with gas and stellar mass are more prominent., The trends of accretion rates with gas and stellar mass are more prominent.
+ The low, The low
+with an azinmthallv svinmetrie density enhancement on its midplane and a solar (vpe star al ils center.,with an azimuthally symmetric density enhancement on its midplane and a solar type star at its center.
+ The results of our numerical simulations of (he motions of solids ranging from nmicron-sized particles to meter-sized objects indicated that solids tend to migrate toward the locations of local maximum pressure., The results of our numerical simulations of the motions of solids ranging from micron-sized particles to meter-sized objects indicated that solids tend to migrate toward the locations of local maximum pressure.
+ In our model nebula. such locations correspond to places where the density of the gas is locally enhanced.," In our model nebula, such locations correspond to places where the density of the gas is locally enhanced."
+ Our results indicate that depending on the physical properties of solids. gas drag induced migration can be quite rapid.," Our results indicate that depending on the physical properties of solids, gas drag induced migration can be quite rapid."
+ Such rapid migrations ean increase (he rates of collision and coalescence of particles and enliance their growth-rates to larger bodies., Such rapid migrations can increase the rates of collision and coalescence of particles and enhance their growth-rates to larger bodies.
+ The results of the extension of this study to include mass-growth ancl collisional coagulation are currentiv under preparation lor publication., The results of the extension of this study to include mass-growth and collisional coagulation are currently under preparation for publication.
+ It is necessary to emphasize that the assumptions of an isothermal ancl turbulence-free nebula were made merely (ο focus attention on the sole effect of gas drag on the rate of migration of solids., It is necessary to emphasize that the assumptions of an isothermal and turbulence-free nebula were made merely to focus attention on the sole effect of gas drag on the rate of migration of solids.
+ Also. in this study. we focused our attention on the motion of objects without considering their mutual interactions.," Also, in this study, we focused our attention on the motion of objects without considering their mutual interactions."
+ In a realistic scenario. one has to allow lor the changes of the eas temperature on both radial anc vertical directions and also take turbulence ancl collision into consideration.," In a realistic scenario, one has to allow for the changes of the gas temperature on both radial and vertical directions and also take turbulence and collision into consideration."
+ Ilt is expected that bv including turbulence. in parücular in the vertical direction. the rate of approach of particles to the midplane of the nebula may decrease.," It is expected that by including turbulence, in particular in the vertical direction, the rate of approach of particles to the midplane of the nebula may decrease."
+ It would then be of utmost. value to investigate to what degree the combined effect of turbulence and gas drag induced rapid migration will enhance the rate of the formation of larger bodies., It would then be of utmost value to investigate to what degree the combined effect of turbulence and gas drag induced rapid migration will enhance the rate of the formation of larger bodies.
+ Such studies are currently underway., Such studies are currently underway.
+ This work is partially supported bv the NASA Origins of the Solar System Program under Grant NAG5-10547 and bv the NASA Astrobiologv Institute under. Cooperative Agreement. NCC2-LO5G., This work is partially supported by the NASA Origins of the Solar System Program under Grant NAG5-10547 and by the NASA Astrobiology Institute under Cooperative Agreement NCC2-1056.
+ , 
+ r, 
+ re, 
+ ret, 
+ ret , 
+"We aocel the pulsar beam shape as elliptical in eeucral and express it analytically as. While a. 9 aud o, can be estimated. directly or indirectly. from observations. AR2 and p, are the two paraiueters which in turn define the beam shape and ze and the available data set of 'T and M-profiles is expected to sauple most of the H/py| range (01) with reasonableuniforimüty.","We model the pulsar beam shape as elliptical in general and express it analytically as, While $\alpha$ , $\beta$ and $\phi_{\nu}$ can be estimated, directly or indirectly, from observations, $R$ and $\rho_{\nu}$ are the two parameters which in turn define the beam shape and size— and the available data set of $\bf{T}$ and $\bf{M}$ -profiles is expected to sample most of the $\mid \beta/\rho_{\nu}\mid$ range (0–1) with reasonableuniformity."
+ The implicit assuniptiou in this, The implicit assumption in this
+"observations centred on the same frequency, the one with the largest bandwidth will be more affected.","observations centred on the same frequency, the one with the largest bandwidth will be more affected."
+ Pulsars are used to determine Faraday Rotation Measures (RM) which provide information on the Galactic magnetic field in the direction of the line of sight., Pulsars are used to determine Faraday Rotation Measures $RM$ ) which provide information on the Galactic magnetic field in the direction of the line of sight.
+" RM measurements are based on the assumption that the only frequency dependence of PA arises from Faraday Rotation (taking into account potential orthogonal jumps), which, in the presence of scattering, is not true."," $RM$ measurements are based on the assumption that the only frequency dependence of PA arises from Faraday Rotation (taking into account potential orthogonal jumps), which, in the presence of scattering, is not true."
+" A figure showing significant variation in phase resolved RM measurements across the pulse, constitutes the basis of an argument on non-orthogonal modes in pulsar B2016+28 in Ramachandran et al. (2004)."," A figure showing significant variation in phase resolved $RM$ measurements across the pulse, constitutes the basis of an argument on non-orthogonal modes in pulsar B2016+28 in Ramachandran et al. \nocite{rbr+04}."
+". The observed Faraday Rotation is an effect of propagation through the interstellar medium and not the pulsar magnetosphere, therefore resolved RM variations are not expected."," The observed Faraday Rotation is an effect of propagation through the interstellar medium and not the pulsar magnetosphere, therefore resolved $RM$ variations are not expected."
+" In that work, an RM is determined for each pulse phase bin, by fitting the simple equation: where c the speed of light and v the frequency."," In that work, an $RM$ is determined for each pulse phase bin, by fitting the simple equation: where $c$ the speed of light and $\nu$ the frequency."
+ We considered interstellar scattering as an alternative explanation to non-orthogonal modes., We considered interstellar scattering as an alternative explanation to non-orthogonal modes.
+" Simulated scattered data for a band of 16 16-MHz channels at 1.4 GHz were fitted for an RM, with no Faraday Rotation included in the simulation."," Simulated scattered data for a band of 16 16–MHz channels at 1.4 GHz were fitted for an $RM$, with no Faraday Rotation included in the simulation."
+ To match the Ramachandran et al., To match the Ramachandran et al.
+" analysis, such a fit was carried out for each pulse phase bin (as shown in the"," analysis, such a fit was carried out for each pulse phase bin (as shown in the"
+on the 24 jun map is shown in Figure 7..,"on the $24\,\mu$ m map is shown in Figure \ref{fig_EVLA_MIPS_24}."
+ We extracted eighteen cuts from each map and determined an azimuthallv-averaged source profile. shown in (Figure 3)).," We extracted eighteen cuts from each map and determined an azimuthally-averaged source profile, shown in (Figure \ref{profiles}) )."
+ The radio shell has a smaller thickness ancl a sharp decreasing trend in (he outer part of the nebula., The radio shell has a smaller thickness and a sharp decreasing trend in the outer part of the nebula.
+ The dust nebula is more extended and shows a smoother distribution in the outer regions., The dust nebula is more extended and shows a smoother distribution in the outer regions.
+" The second shell is clearly evident in the 24jan profile. at a distance of about 2:200"" from the center."," The second shell is clearly evident in the $24\,\mu$ m profile, at a distance of about $ \pm 200^{\prime \prime}$ from the center."
+ At the same distance. less delined peaks are visible in the 70jan profile. smoothed out by the coarser spatial resolution.," At the same distance, less defined peaks are visible in the $70\,\mu$ m profile, smoothed out by the coarser spatial resolution."
+ In this letter we present new radio EVLA observations of the nebula surrounding the LBV G79.29+0.46., In this letter we present new radio EVLA observations of the nebula surrounding the LBV G79.29+0.46.
+ In particular. we have analvzed the 6 em map ancl. for the first time. compared the radio free-free. which (races the spatial distribution of the ionized gas. with maps recently obtained in the mid-IR.9 which traces the spatial distribution of the dust component.," In particular, we have analyzed the 6 cm map and, for the first time, compared the radio free-free, which traces the spatial distribution of the ionized gas, with maps recently obtained in the mid-IR, which traces the spatial distribution of the dust component."
+ The dust distribution observed at 8 jn appears different than the distribution at 24 and 70;0n. The warmer dust. traced by the & san emission. is more concentrated in the southwest and southeast regions.," The dust distribution observed at $8\,\mu$ m appears different than the distribution at $24$ and $70\,\mu$ m. The warmer dust, traced by the $8\,\mu$ m emission, is more concentrated in the southwest and southeast regions."
+ Moreover. when compared with the radio nebula. the 24 and 70 sau emission appears more extended. while the 8 sm emission is well-contined inside the radio nebula.," Moreover, when compared with the radio nebula, the $24$ and $70\,\mu$ m emission appears more extended, while the $8\,\mu$ m emission is well-contained inside the radio nebula."
+ This is consistent with (he existence of two dust components: a cooler component traced by the 24 and 70 san emission. and a warmer dust component. probably consisting of smaller grams. traced by the 8 jm emission.," This is consistent with the existence of two dust components: a cooler component traced by the $24$ and $70\,\mu$ m emission, and a warmer dust component, probably consisting of smaller grains, traced by the $8\,\mu$ m emission."
+ This warmer component. with morphological properties quite different from thoseobserved al longer wavelengths. could be related to the PAIL emission that has been claimed by Jimenez-Estebanetal. (2010)..," This warmer component, with morphological properties quite different from thoseobserved at longer wavelengths, could be related to the PAH emission that has been claimed by \citet{Jimenez_2010}. ."
+"expansion. would be the sui of] a large-scale structure ter ας(LLS)x!0.1 anda confusion noise terii which isa finction of the flux 5,4; down to which sources have been excise frou the maps for the uniltipole. analysis. «,[Noise)nX55.l1cH Where > Is. the slope of theintegral source οςnts in the eunergv band used (N(7ο)X597).","expansion would be the sum of a large-scale structure term $a_{lm}^{(LLS)}\propto l^{-0.4}$ and a confusion noise term which is a function of the flux $S_{cut}$ down to which sources have been excised from the maps for the multipole analysis $a_{lm}^{(Noise)}\propto S_{cut}^{\gamma-1}$, where $\gamma$ is the slope of the source counts in the energy band used $N(>S)\propto S^{-\gamma}$ )."
+ Trever e al (1998) perform this aualvsis on the WEAQO-1 A2 all sky data by removing regions around the Piccinotti et al (1982) sample aud the galactic plane., Treyer et al (1998) performed this analysis on the HEAO-1 A2 all sky data by removing regions around the Piccinotti et al (1982) sample and the galactic plane.
+ They find evidence for a growth of the spherical harmonic cocticicuts erowing at low values o fi a manner roughv consistent with the predictions., They find evidence for a growth of the spherical harmonic coefficients growing at low values of $l$ in a manner roughly consistent with the predictions.
+ The significance of the signa is difficult to assess as the harmonic coefficients are not independent due to cross-talls between differeit orders introduced by the masking., The significance of the signal is difficult to assess as the harmonic coefficients are not independent due to cross-talk between different orders introduced by the masking.
+ Assmuine a redslift dependent bias parameter for the N-rav sources parametrized as by(:)byο)τον)1| Gwhich assiunes that all galaxies form at some past epoch. Fry 1996). they estimate a rather modest jas parameter (4.0<by(0)1.6).," Assuming a redshift dependent bias parameter for the X-ray sources parametrized as $b_X(z)=b_X(0)+z[b_X(0)-1]$ (which assumes that all galaxies form at some past epoch, Fry 1996), they estimate a rather modest bias parameter $1.0<b_X(0)<1.6$ )."
+ In their dciaeranis it is also seen hat the dipole (/=1) has an unusually large amplitude compared to higher harinonics., In their diagrams it is also seen that the dipole $l=1$ ) has an unusually large amplitude compared to higher harmonics.
+ The way to go is Τιςecd to have precise measurements of the NRB intensity ou laree angular scales. mt with the possibility of excludiug sources down to the faintest possible levels.," The way to go is indeed to have precise measurements of the XRB intensity on large angular scales, but with the possibility of excluding sources down to the faintest possible levels."
+ Trever et al (1998) sueeestOO that a1 all-sky map with XRB lutensities measured wit ra precision and with souCs excised down to eni?s8l dio. 100 times £üuter than the Picciuo tict al catalogue) would be ideal for a spherical harmonic analysis.," Treyer et al (1998) suggest that an all-sky map with XRB intensities measured with a precision and with sources excised down to $3\times 10^{-13}\,
+{\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ (i.e., 100 times fainter than the Piccinotti et al catalogue) would be ideal for a spherical harmonic analysis."
+ An alternative wav that has been devised to look for structure in the N-rav s£V is to cross-correlate the unresolved NRB intensity with caalogues of galaxies., An alternative way that has been devised to look for structure in the X-ray sky is to cross-correlate the unresolved XRB intensity with catalogues of galaxies.
+" The auplitude of the cross-correlation function (CCE) between he N-rav intensity Lygp aud the galaxy surface deusity Wy. Wv,(0)=OxnpNqioIvepiiNyg)1 αἲ lag (0= 0) provides an approximate nieasurenment of the fraction of the NRB arsiis either in the catalogued galaxies or in sources clustered wih them within a scale of the beam with which X-ray observations have been obtained (Lahavy ct al 1993)."," The amplitude of the cross-correlation function (CCF) between the X-ray intensity $I_{XRB}$ and the galaxy surface density $N_g$, $W_{Xg}(\theta)=\langle I_{XRB}N_g\rangle_{\theta}/\langle
+I_{XRB}\rangle\langle N_g\rangle-1$ at zero-lag $\theta=0$ ) provides an approximate measurement of the fraction of the XRB arising either in the catalogued galaxies or in sources clustered with them within a scale of the beam with which X-ray observations have been obtained (Lahav et al 1993)."
+ Positive signals have been found for Wy. typically of the order of when the ealaxies are optically or infrared selected. aid up to >1 when active ealaxi are selected. (see table 2).," Positive signals have been found for $W_{Xg}$, typically of the order of when the galaxies are optically or infrared selected, and up to $>10$ when active galaxies are selected (see table 2)."
+ The iuterpretatio1 of this signal requires to model t clustering of N-rayv sources around the catalogued galaxies. which is iudeed modulate by the bias parameter by.," The interpretation of this signal requires to model the clustering of X-ray sources around the catalogued galaxies, which is indeed modulated by the bias parameter $b_X$."
+ Usine by=1. it is fouud that the local volume eiuissiviti optically selected. galaxies amounts to ~1P)heresMpe? (Labav et al 1993: ivaji ct al 1991: Carrera ct al 1995). most of which is contributed bv Seyfert daxies and QSOs (Barcous et al 1995).," Using $b_X=1$, it is found that the local volume emissivity of optically selected galaxies amounts to $\sim 10^{39}\, h\, {\rm erg}\, {\rm s}^{-1}\, {\rm
+Mpc}^{-3}$ (Lahav et al 1993; Miyaji et al 1994; Carrera et al 1995), most of which is contributed by Seyfert galaxies and QSOs (Barcons et al 1995)."
+ When this volune eniüssivitv is extrapoated to higher redshifts. the fraction | the NRB intensity due to the precursors of the catalogued galaxies can be predicted (Lahav et al 1993).," When this volume emissivity is extrapolated to higher redshifts, the fraction of the XRB intensity due to the precursors of the catalogued galaxies can be predicted (Lahav et al 1993)."
+ Carrera ct al (1995) find that of the hard X- backeround might be produced by optically selected galaxies without excecding the upper lits ou the autocorrelation fuuctiou of the NRD., Carrera et al (1995) find that of the hard X-ray background might be produced by optically selected galaxies without exceeding the upper limits on the autocorrelation function of the XRB.
+ This value is similar to the result of cross-correlation analyses of deep ROSAT. Nav images with deep, This value is similar to the result of cross-correlation analyses of deep $ROSAT$ X-ray images with deep
+We compiled a total of 113 Southern Hemisphere Vega-like candidate stars from the literature. based on their infrared or submillimetric excess emissions (Backman&Paresce1993;SylvesterHabingetal.2001:Laureijs2002:Sheret 2004).,"We compiled a total of 113 Southern Hemisphere Vega-like candidate stars from the literature, based on their infrared or submillimetric excess emissions \citep{backman-paresce93,sylvester96,mannings-barlow98,fajardo-acosta99,sylvester-mannings00,
+habing01,laureijs02,sheret04}."
+. This compilation also includes G dwarfs with infrared excess recently detected by Spitzer (Beichmanetal.2005.2006:Brydenetal.2006:SuTrilling 2008).," This compilation also includes G dwarfs with infrared excess recently detected by Spitzer \citep{beichman05,beichman06,bryden06,su06,trilling08}."
+. Specifically the list comprises objects with BAFGK spectral types (22. 38. 28. 17 and 8. respectively).," Specifically the list comprises objects with BAFGK spectral types (22, 38, 28, 17 and 8, respectively)."
+ All the stars are lumimosity class V (Hipparcos catalogue) and have distances between 5 and 300 pe., All the stars are luminosity class V (Hipparcos catalogue) and have distances between 5 and 300 pc.
+ Table |lists the observed objects., Table \ref{muestra} lists the observed objects.
+ Table 1. includes a sub-sample of stars that were originally selected by IRAS as candidate Vega-like stars., Table \ref{muestra} includes a sub-sample of stars that were originally selected by IRAS as candidate Vega-like stars.
+ However. when observed by Spitzer the infrared excesses were deemed to be of little significance.," However, when observed by Spitzer the infrared excesses were deemed to be of little significance."
+ These objects are: HD 10800. HD 20794. HD 38393. HD 41700. HD 68456. HD 160691. HD 169830. HD 203608. and HD 216437 (Beichmanetal.2005.2006:Brydenetal.2006:Hillenbrand2008:Trilling 2008)..," These objects are: HD 10800, HD 20794, HD 38393, HD 41700, HD 68456, HD 160691, HD 169830, HD 203608, and HD 216437 \citep{beichman05,beichman06,bryden06,hill08,trilling08}. ."
+ For example. Brydenetal.(2006) found that for HD 10800 Duces=].3 (the observed flux over the photospheric emission at 70 pam) and ANS;ens=|.2 for HD 68456.," For example, \citet{bryden06} found that for HD 10800 ${f_{\rm MIPS\_70 \mu m} \over f_*} = 1.3$ (the observed flux over the photospheric emission at 70 $\mu$ m) and ${f_{\rm MIPS\_70 \mu m} \over f_*} = 1.2$ for HD 68456."
+ This group of objects should be considered with caution., This group of objects should be considered with caution.
+ The stellar spectra were obtained at the Complejo Astronomico ΕΙ Leoncito (CASLEO). using the 2.15-m telescope equipped with a REOSC echelle and a TEK 10241024 CCD detector.," The stellar spectra were obtained at the Complejo Astronomico El Leoncito (CASLEO), using the 2.15-m telescope equipped with a REOSC echelle and a TEK $\times$ 1024 CCD detector."
+ The REOSC spectrograph uses gratings as cross dispensers., The REOSC spectrograph uses gratings as cross dispensers.
+" We used a grating with 400 lines mm!. covering the spectral range ,13500-6500. giving a resolving power of ~ 12500."," We used a grating with 400 lines $^{-1}$, covering the spectral range $\lambda\lambda$ 3500–6500, giving a resolving power of $\sim$ 12500."
+ Three individual spectra for each object were obtained in four observing runs: August 05-08 2005. August 18-22 2005. February 18-25 2006 and May 04-07 2007 and have S/N ratio of about 300.," Three individual spectra for each object were obtained in four observing runs: August 05–08 2005, August 18–22 2005, February 18–25 2006 and May 04–07 2007 and have S/N ratio of about 300."
+ The spectra were reduced using IRAF standard procedures for echelle spectra., The spectra were reduced using IRAF standard procedures for echelle spectra.
+ We applied bias and flat corrections and then normalized order by order with the task. using order Chebyshev polynomials.," We applied bias and flat corrections and then normalized order by order with the task, using 7--9 order Chebyshev polynomials."
+ We also corrected by the scattered light in the spectrograph task)., We also corrected by the scattered light in the spectrograph task).
+ We fitted the background with a linear function on both sidesof the echelle apertures. using the task αρ.," We fitted the background with a linear function on both sidesof the echelle apertures, using the task ."
+ The resolution of the reduced spectra is 0.17 A//pix., The resolution of the reduced spectra is 0.17 /pix.
+seenis (oo contrived. though SN explosion mechanism is still uncertain and (hus we cannot conclusively deny such a possibility.,"seems too contrived, though SN explosion mechanism is still uncertain and thus we cannot conclusively deny such a possibility."
+ llere we discuss the additional cases with the dillerent Gimescales. the time it takes the flow to cool from some temperature to lower temperature.," Here we discuss the additional cases with the different timescales, the time it takes the flow to cool from some temperature to lower temperature."
+ As a result. [Co/Fe]| is changed by the timescale below T= 4 x 10? I (Pruet et al.," As a result, [Co/Fe] is changed by the timescale below $T$ = 4 $\times$ $^9$ K (Pruet et al."
+ 2005)., 2005).
+ The timescale below 7 = 4 x 10 K multiplied by a factor of 1/2 and 2 gives [Co/Fe]| ~ 0.5 and 0.2. respectively. while [Co/Fe| ~ 0.4 in the case with the original timescale as mentioned in Section 5.," The timescale below $T$ = 4 $\times$ $^9$ K multiplied by a factor of 1/2 and 2 gives [Co/Fe] $\sim$ 0.5 and 0.2, respectively, while [Co/Fe] $\sim$ 0.4 in the case with the original timescale as mentioned in Section 5."
+ However. Ye range and ejected mass needed is not so changed: the upper limit of Ye becomes a little larger. Ye ~ 0.502. and ejected mass. a little smaller  0.03 A..," However, Ye range and ejected mass needed is not so changed; the upper limit of Ye becomes a little larger, Ye $\sim$ 0.502, and ejected mass, a little smaller $\sim$ 0.03 $M_\odot$."
+ There is another reason why SN models are not favored to explain large |Co.Zn/Fe| stars.," There is another reason why SN models are not favored to explain large [Co,Zn/Fe] stars."
+ As mentioned in Introduction. [Co/Fe| decreases with higher metallicity. and many stars show [Co/Fe] ~ 0 when [Fe/I] = -2.5.," As mentioned in Introduction, [Co/Fe] decreases with higher metallicity, and many stars show [Co/Fe] $\sim$ 0 when [Fe/H] $\gsim$ -2.5."
+ M this age. only core-collapse SNe are considered to contribute to interstellar matter [rom the observation of [a-elements/Fe).," At this age, only core-collapse SNe are considered to contribute to interstellar matter from the observation of $\alpha$ -elements/Fe]."
+ If the typical value of [Co/Fe] in normal SN is ~ 0.3. [Co/Fe| does not decrease to ~ 0 at [Fe/1I] e -2.5.," If the typical value of [Co/Fe] in normal SN is $\sim$ 0.3, [Co/Fe] does not decrease to $\sim$ 0 at [Fe/H] $\sim$ -2.5."
+ On the other hand. IMF average of the supernova ejecta can be [Co.Zn/Fe] ~ 0 even though ΗΝ models have large [Co.Zn/Fe| (Tominaga et al.," On the other hand, IMF average of the supernova ejecta can be [Co,Zn/Fe] $\sim$ 0 even though HN models have large [Co,Zn/Fe] (Tominaga et al."
+ 2007)., 2007).
+ Therefore. we suggest that Co and Zn in ENP stars originated from IINe. not [rom normal SNe.," Therefore, we suggest that Co and Zn in EMP stars originated from HNe, not from normal SNe."
+ As lor weak r-process elements. it is not clear (hat Sr-Zr rich weak r-process stars are alwavs Alo-Rh rich.," As for weak r-process elements, it is not clear that Sr-Zr rich weak r-process stars are always Mo-Rh rich."
+ Observational data of both Sr-Zr ancl Mo-Ili are obtained in only two weak r-process stars IID122563 and Η55009 (Lloneda et al., Observational data of both Sr-Zr and Mo-Rh are obtained in only two weak r-process stars HD122563 and HD88609 (Honda et al.
+ 2007)., 2007).
+" Si-Zr are produced in neutron-rich hot-bubble of 5/4, ~ 15. while both Si-Zr and Mo-Rh are produced in hot-bubble of ο, ~ 150. that may be the part of neutrino-driven wind. rather than hot-bubble component (Roberts et al."," Sr-Zr are produced in neutron-rich hot-bubble of $s$ $k_{\rm b}$ $\sim$ 15, while both Sr-Zr and Mo-Rh are produced in neutron-rich hot-bubble of $s$ $k_{\rm b}$ $\sim$ 150, that may be the part of neutrino-driven wind, rather than hot-bubble component (Roberts et al."
+ 2010)., 2010).
+ Of course. (his work is only parametric search. and some other possibilities nav exist in realitv.," Of course, this work is only parametric search, and some other possibilities may exist in reality."
+ Further investigation is needed to disclose the origin of the abundances in EMD stars., Further investigation is needed to disclose the origin of the abundances in EMP stars.
+ This work was partly supported by Grants-in-Aid for JSPS Fellows., This work was partly supported by Grants-in-Aid for JSPS Fellows.
+operated by the Association of Universities for Research in Astronomy. Ine.. for NASA under constract NAS 5-26555.,"operated by the Association of Universities for Research in Astronomy, Inc., for NASA under constract NAS 5-26555."
+(the Rossiter-MeLaughlin effect is also dependent on limb darkening) we use values for the V band.,(the Rossiter-McLaughlin effect is also dependent on limb darkening) we use values for the $V$ band.
+ ? showed that HARPS is centred on the V band., \citet{Triaud:2009p3865} showed that HARPS is centred on the $V$ band.
+ The coefficients were chosen for atmospheric parameters close to those presented in Table 2.., The coefficients were chosen for atmospheric parameters close to those presented in Table \ref{tab:SpecParam}.
+ For each star. we performed four analyses. each using a MCMC chain with 0000 accepted steps: This is to assess the sensitivity of the model parameters to a small but uncertain orbital eccentricity and to the value found by spectral analysis which. as demonstrated in ?.. can seriously affect the fitting of the Rossiter-MeLaughlin effect.," For each star, we performed four analyses, each using a MCMC chain with 000 accepted steps: This is to assess the sensitivity of the model parameters to a small but uncertain orbital eccentricity and to the value found by spectral analysis which, as demonstrated in \citet{Triaud:2009p3865}, can seriously affect the fitting of the Rossiter-McLaughlin effect."
+ The comparative tables holding the results of these various fits are available in the appendices to support the conclusions we reach while allowing readers to form their own opinion., The comparative tables holding the results of these various fits are available in the appendices to support the conclusions we reach while allowing readers to form their own opinion.
+ In addition. we also conducted control chains fixing the parameters controlled by the photometry 1n order to check whether this was a limiting source of errors in the determination of our most important parameter: B.," In addition, we also conducted control chains fixing the parameters controlled by the photometry in order to check whether this was a limiting source of errors in the determination of our most important parameter: $\beta$."
+ The results from these chains are in the appendices as well., The results from these chains are in the appendices as well.
+ Although different in their starting hypotheses all the chains are also useful at checking their respective. convergence., Although different in their starting hypotheses all the chains are also useful at checking their respective convergence.
+ Our final results are presented in Table 3.., Our final results are presented in Table \ref{tab:params}.
+ The best solution 1s found in the best of the four fits by comparing y and using Ockham's principle of minimising the number of parameters for similar results: for fits with similar -- we usually choose a circular solution with πο prior on Vsin/., The best solution is found in the best of the four fits by comparing $\chi^2$ and using Ockham's principle of minimising the number of parameters for similar results: for fits with similar $\chi^2_\mathrm{reduced}$ we usually choose a circular solution with no prior on $V\sin I$.
+ Results are extracted from the best fit by taking the median of the posterior probability distribution for each parameter. determined from the Markov chain.," Results are extracted from the best fit by taking the median of the posterior probability distribution for each parameter, determined from the Markov chain."
+ Errors bars are estimated from looking at the extremes of the distribution comprising the In the following section and m tables. several statistical values are used: y is the value found for all the data. while gives the value of y solely for the radial velocities.," Errors bars are estimated from looking at the extremes of the distribution comprising the In the following section and in tables, several statistical values are used: $\chi^2$ is the value found for all the data, while $\chi^2_\mathrm{RV}$ gives the value of $\chi^2$ solely for the radial velocities."
+ The reducedView γ΄ for the radial velocities. denoted by aar is used to estimate how well a model fits the data and to compare various fits and their respective significance.," The reduced $\chi^2$ for the radial velocities, denoted by $\chi^2_\mathrm{reduced}$, is used to estimate how well a model fits the data and to compare various fits and their respective significance."
+ In addition we will also use the residuals. denoted as O—C.," In addition we will also use the residuals, denoted as $O-C$."
+ These estimates are only for radial velocities., These estimates are only for radial velocities.
+ The results from photometry are not mentioned since they are not new., The results from photometry are not mentioned since they are not new.
+ They are only here to constrain the shape of the Rossiter-MeLaughlin effect., They are only here to constrain the shape of the Rossiter-McLaughlin effect.
+ When giving bounds. for eccentricity and long term radial velocity drift. we quote the confidence interval. for exclusion.," When giving bounds, for eccentricity and long term radial velocity drift, we quote the confidence interval for exclusion."
+ A sequence of 26 RV measurements was taken on WASP-2 using HARPS on 2008 October 15. with additional observations made outside transit as given 1n the journal of observations presented in the appendices.," A sequence of 26 RV measurements was taken on WASP-2 using HARPS on 2008 October 15, with additional observations made outside transit as given in the journal of observations presented in the appendices."
+ The cadence during transit was close to a point every 430s., The cadence during transit was close to a point every 430s.
+ The average photon noise error of that sequence is 5.7 ss!., The average photon noise error of that sequence is $5.7$ $^{-1}$.
+ We made additional observations with CORALIE to refine the orbital. solution obtained by ? using the SOPHIE instrument on the mm telescope at Observatoire de Haute-Provence. and to look for long-term variability of the orbit.," We made additional observations with CORALIE to refine the orbital solution obtained by \citet{Cameron:2007p2879} using the SOPHIE instrument on the m telescope at Observatoire de Haute-Provence, and to look for long-term variability of the orbit."
+ 20 measurements were taken with a mean precision of 13.9 ss! over close to 11 months between 2008 October 25 and 2009 September 23., 20 measurements were taken with a mean precision of $13.9$ $^{-1}$ over close to 11 months between 2008 October 25 and 2009 September 23.
+ All the RV data is available in the appendices along with exposure times., All the RV data is available in the appendices along with exposure times.
+ To establish the photometric ephemeris and the transit geometry. we fitted the photometric datasets of ? (3 seasons by SuperWASP in the unfiltered WASP bandpass). ? (a z band Keplercam lighteurve) and ?. (a William Herschel Telescope R band transit curve).," To establish the photometric ephemeris and the transit geometry, we fitted the photometric datasets of \citet{Cameron:2007p2879} (3 seasons by SuperWASP in the unfiltered WASP bandpass), \citet{Charbonneau:2007p1458} (a $z$ band $Keplercam$ lightcurve) and \citet{Hrudkova:2009p1459} (a William Herschel Telescope $R$ band transit curve)."
+ WASP-2b's data were fitted with up to 10 free parameters plus 8 independent adjustment parameters: three y velocities for the three RV data sets and five normalisation factors for photometry., WASP-2b's data were fitted with up to 10 free parameters plus 8 independent adjustment parameters: three $\gamma$ velocities for the three RV data sets and five normalisation factors for photometry.
+ This sums up to 58 RV measurements and 895] photometric observations., This sums up to 58 RV measurements and 8951 photometric observations.
+ nudus] COS not Improve significantly between circular and eccentric models., $\chi^2_\mathrm{reduced}$ does not improve significantly between circular and eccentric models.
+" We therefore Impose a circular solution,", We therefore impose a circular solution.
+ The presence of a prior on Vsin/ does not affect the results., The presence of a prior on $V\sin I$ does not affect the results.
+ We find Vsin/.=0.997] kkmss! in accordance with the value found in section ?2.., We find $V\sin I = 0.99 ^{+0.27}_{-0.32}$ $^{-1}$ in accordance with the value found in section \ref{sec:SpecAn}.
+ The fit delivers p=1537 7i.," The fit delivers $\beta = 153^{\circ\,+11}_{\,\,\,\,-15}$ ."
+ The overall root-mean-square (RMS) scatter of the spectroscopic residuals about the fitted model is !., The overall root-mean-square (RMS) scatter of the spectroscopic residuals about the fitted model is $^{-1}$ .
+ During the HARPS transit sequence these residuals are at ss., During the HARPS transit sequence these residuals are at $^{-1}$.
+ Fig., Fig.
+ 3 shows the resulting distribution às. VsirI vs. ., \ref{fig:betas} shows the resulting distribution as $V\sin I$ vs. $\beta$.
+ We detect Vsin/ significantly above zero with confidence interval showing that (37) of the posterior probability function has Vsin/>0.2 kkmss! whileB>77.267.," We detect $V\,\sin\,I$ significantly above zero with confidence interval showing that $3\,\sigma$ ) of the posterior probability function has $V\sin I > 0.2$ $^{-1}$ while $\beta > 77.26^{\circ}$."
+ We have computed 6 additional chains in order to test the strength of our conclusions., We have computed 6 additional chains in order to test the strength of our conclusions.
+ Table AI. shows the comparison between the various fits; we invite the reader to refer to it as only important results are given in the text., Table \ref{tab:WASP1comp} shows the comparison between the various fits; we invite the reader to refer to it as only important results are given in the text.
+ In all cases. eccentricity is not detected being below a 3c significance fromcircular which ts likely affected from the poor coverage of the phase by the HARPS points.," In all cases, eccentricity is not detected being below a $3\,\sigma$ significance fromcircular which is likely affected from the poor coverage of the phase by the HARPS points."
+ Circular solutions are therefore adopted., Circular solutions are therefore adopted.
+ We fix the eccentricity’s upper limit to e<0.070., We fix the eccentricity's upper limit to $e < 0.070$.
+ In addition no significant long term drift wasdetected in the spectroscopy: |y]«36 mmss! yyr!., In addition no significant long term drift wasdetected in the spectroscopy: $| \dot{\gamma} | < 36 $ $^{-1}$ $^{-1}$ .
+ Using the spectroscopically-determined vsin/ value of ss! and forcingβ to zero. -- changes from," Using the spectroscopically-determined $\vsini$ value of $^{-1}$ and forcing $\beta$ to zero, $\chi^2_\mathrm{reduced}$ changes from"
+12pt,= 12pt
+12pt-,= 12pt
+"Through most of the following analysis we take a fiducial model with d,=3 and z,=8.9, corresponding to 7=0.1 as already well determined by WMAPS.","Through most of the following analysis we take a fiducial model with $d_\eta=3$ and $z_{\rm r}=8.9$, corresponding to $\tau=0.1$ as already well determined by WMAP5."
+" Flat priors are assumed on z, and d,,, over ranges of 6—30 and 0-10 respectively (the figures show that values of d; larger than this result in scenarios very close to instantaneous)."," Flat priors are assumed on $z_{\rm r}$ and $d_\eta$, over ranges of 6–30 and 0–10 respectively (the figures show that values of $d_\eta$ larger than this result in scenarios very close to instantaneous)."
+" Figure 1 also motivates us to try a logarithmic prior on d,,, for which we take the range 0.3-30."," Figure \ref{f:models3} also motivates us to try a logarithmic prior on $d_\eta$, for which we take the range 0.3–30."
+" However a 7<0.3 prior is also imposed, so that the one additional parameter as compared to instant reionization does not correspond to one additional degree of freedom."," However a $\tau<0.3$ prior is also imposed, so that the one additional parameter as compared to instant reionization does not correspond to one additional degree of freedom."
+" For this reason regular likelihood ratio tests, of the kind performed in Kaplinghat et al. ("," For this reason regular likelihood ratio tests, of the kind performed in Kaplinghat et al. ("
+2003) and Holder et al. (,2003) and Holder et al. (
+"2003), will not be valid.","2003), will not be valid."
+" Figure 3 shows the induced prior on 7 resulting from linear priors on z, and d, (their own priors are not perfectly flat due to the extra imposition of the 7«0.3 prior).", Figure \ref{f:testpriors} shows the induced prior on $\tau$ resulting from linear priors on $z_{\rm r}$ and $d_\eta$ (their own priors are not perfectly flat due to the extra imposition of the $\tau < 0.3$ prior).
+" The uncertainty on 7 from Planck is of order 0.01, and over such widths the prior on 7 is roughly uniform; it is even more so around the fiducial value of 7."," The uncertainty on $\tau$ from Planck is of order 0.01, and over such widths the prior on $\tau$ is roughly uniform; it is even more so around the fiducial value of $\tau$."
+" The same applies to the case of a log prior on d,,.", The same applies to the case of a log prior on $d_\eta$.
+" Other models of reionization have been proposed (Haiman Holder 2003; Cen 2003), for instance the double reionization scenario considered in Lewis et al. ("," Other models of reionization have been proposed (Haiman Holder 2003; Cen 2003), for instance the double reionization scenario considered in Lewis et al. ("
+"2006), which would in general require a third parameter.","2006), which would in general require a third parameter."
+" From a model selection point of view ttaking into account parameter uncertainties within each model) results in this paper indicate that discriminating such a model from a smooth transition model would be beyond the scope of Planck, though perhaps within the scope of a closer to cosmic variance limited experiment."," From a model selection point of view taking into account parameter uncertainties within each model) results in this paper indicate that discriminating such a model from a smooth transition model would be beyond the scope of Planck, though perhaps within the scope of a closer to cosmic variance limited experiment."
+ Here our focus is mainly on clarifying what Planck can learn about reionization., Here our focus is mainly on clarifying what Planck can learn about reionization.
+ Model selection forecasting assesses a given experiment's ability to distinguish between different cosmological models., Model selection forecasting assesses a given experiment's ability to distinguish between different cosmological models.
+ This ability necessarily depends on the true model and on its parameter values (and of course on the overriding assumption that one of the, This ability necessarily depends on the true model and on its parameter values (and of course on the overriding assumption that one of the
+kpe from the Galactic plauc.,kpc from the Galactic plane.
+ Thanks to nearly complete flux-liuited SDSS stellar samples with both rotational and metallicity. iieasuremenuts. the expected correlation was stronglv ruled out (Ivezió et al.," Thanks to nearly complete flux-limited SDSS stellar samples with both rotational and metallicity measurements, the expected correlation was strongly ruled out (Ivezić et al."
+ 2008)., 2008).
+ On the other haud. the MW. simulation we preseuted here. while rot a detailed model of the Milky Way. produces ood qualitative aercement with the data and sheds new ight ou the origin and evolution of the observed disk structure.," On the other hand, the MW simulation we presented here, while not a detailed model of the Milky Way, produces good qualitative agreement with the data and sheds new light on the origin and evolution of the observed disk structure."
+ Of particular importance is the role of radial weration iu nüxiug stars born throughout the disk iuto he solar neighborhood., Of particular importance is the role of radial migration in mixing stars born throughout the disk into the solar neighborhood.
+ The properties of the Roskkar-IZ et al., The properties of the Roškkar et al.
+ model for the overall spatial. metallicity and kincmatic distributions of he Milkv. Way stars are in qualitative agreement with he SDSS data (-Jurié et al.," model for the overall spatial, metallicity and kinematic distributions of the Milky Way stars are in qualitative agreement with the SDSS data (Jurić et al."
+ 2008: Ivezic et al., 2008; Ivezić et al.
+ 2008: Dond et al., 2008; Bond et al.
+ 2010)., 2010).
+ While there are quantitative differences iu he spatial exadieuts of these distributions. as well as in heir detailed behavior. even this qualitative agreement is remarkable because the simulation was not fine tuned o inateh the Milkv Way," While there are quantitative differences in the spatial gradients of these distributions, as well as in their detailed behavior, even this qualitative agreement is remarkable because the simulation was not fine tuned to match the Milky Way."
+ Not only does. the model reproduce the observed chanee of slope iu the counts of disk stars as a function of distance from the Calactic aue (the original motivation for iutroducing a separate hick disk component). but it also predicts the eradieuts in metallicity aud rotational velocity.," Not only does the model reproduce the observed change of slope in the counts of disk stars as a function of distance from the Galactic plane (the original motivation for introducing a separate thick disk component), but it also predicts the gradients in metallicity and rotational velocity."
+ Furthermore. the netallicity and rotational velocity are uncorrelated iu his model for appropriately volumic-selected subsamples of stars. in agreement with observations.," Furthermore, the metallicity and rotational velocity are uncorrelated in this model for appropriately volume-selected subsamples of stars, in agreement with observations."
+ The robust qualitative agreements between the data and model predictions motivate the use of model quantities inaccessible to observations. such as the stellar age and the ISAL metallicity at the time and position of stellar birth. to interpret recent SDSS results.," The robust qualitative agreements between the data and model predictions motivate the use of model quantities inaccessible to observations, such as the stellar age and the ISM metallicity at the time and position of stellar birth, to interpret recent SDSS results."
+ Iu particular. the lack of correlation between the metallicity and rotational velocity at 1 kpe from the Galactic plane can be understood as due to complex interplay between the ISM aetallicity at the time and position of stellar birth. aud the subsequent secular evolution largely drive- by spiral arius.," In particular, the lack of correlation between the metallicity and rotational velocity at $\sim$ 1 kpc from the Galactic plane can be understood as due to complex interplay between the ISM metallicity at the time and position of stellar birth, and the subsequent secular evolution largely driven by spiral arms."
+ Nopriori assumptions about the disk’s structure are incorporated in the model vet it reproduces the main observational results which —motivated decomposition of the disk iuto two presuuablv distinct componcuts., No assumptions about the disk's structure are incorporated in the model – yet it reproduces the main observational results which motivated decomposition of the disk into two presumably distinct components.
+ The absence of niergers in this model implies that they are not required to explain the overall disk structure., The absence of mergers in this model implies that they are not required to explain the overall disk structure.
+ While Ισος. renimants are detected within the Milkv Way disk (e.g. JOS). thoi influence is appareutly well localized.," While merger remnants are detected within the Milky Way disk (e.g. J08), their influence is apparently well localized."
+ lustead. the Milkv. Way disk cau be viewed as a single conmrplex structure with an age distribution that is a strong function of position iu the Galaxy due to racial migration of stars.," Instead, the Milky Way disk can be viewed as a single complex structure with an age distribution that is a strong function of position in the Galaxy due to radial migration of stars."
+ Thus even if a primordial thick disk is preseut. having formed via accretion/exterual heating. it is likely to be substantially polluted by mierating disk stars.," Thus even if a primordial thick disk is present, having formed via accretion/external heating, it is likely to be substantially polluted by migrating disk stars."
+ The same mixing effects are likely responsible for the observed differences in à clement abundance betweenselected thin and thick disk stars., The same mixing effects are likely responsible for the observed differences in $\alpha$ element abundance between thin and thick disk stars.
+ By adopting kincmatic selection criteria used by observers. we are able to reproduce distinctive |a/Fe| and APOGEE (Alajewskietal.2010).," By adopting kinematic selection criteria used by observers, we are able to reproduce distinctive $\alpha$ /Fe] and APOGEE \citep{Majewski2010}."
+. Kev to garnering a deeper uuderstaudiug of the importance of radial migratiou iu the Milkv Way evolution is eathering both precise age determinations and detailed chemical compositions., Key to garnering a deeper understanding of the importance of radial migration in the Milky Way evolution is gathering both precise age determinations and detailed chemical compositions.
+ We are optimistic that this study will lead to further observational aud theoretical work., We are optimistic that this study will lead to further observational and theoretical work.
+ We thank our uuuerous SDSS (www.sdss.ore) collaborators for their valuable contributious aud helpfu discussions., We thank our numerous SDSS (www.sdss.org) collaborators for their valuable contributions and helpful discussions.
+ S. Lochiman aud Ivezié acknowledge support by NSF erauts AST-0707901 and AST-100878 to the University of Washington. aud by NSF eraut AST-0551161 to LSST for design aud development activity.," S. Loebman and Ivezić acknowledge support by NSF grants AST-0707901 and AST-1008784 to the University of Washington, and by NSF grant AST-0551161 to LSST for design and development activity."
+ Ivezió acknowledges support bv the Croatian Nationa Science Foundation erant O-1518-2009., Ivezić acknowledges support by the Croatian National Science Foundation grant O-1548-2009.
+ This research was supported in part by the NSF through TeraCuwic resources provided by TACC and PSC., This research was supported in part by the NSF through TeraGrid resources provided by TACC and PSC.
+ BR. Roskkar anc T. R. Quinn were supported by the NSF ITR evant at the University of Washington., R. Roškkar and T. R. Quinn were supported by the NSF ITR grant PHY-0205413 at the University of Washington.
+Shocks are ubiquitous in the Universe aud occur at various scales: from planetary shocks iu our Solar Svsteum to large scale shocks in clusters of ealaxics.,Shocks are ubiquitous in the Universe and occur at various scales: from planetary shocks in our Solar System to large scale shocks in clusters of galaxies.
+ Astrophysical shocks are rather different from shocks on Earth. as typical mncan-free-patls for particle-particle interactions in the interstellar mediuni are of the order of parsecs. which is larger than the scale of the shock trausitious of mterstellar shocks.," Astrophysical shocks are rather different from shocks on Earth, as typical mean-free-paths for particle-particle interactions in the interstellar medium are of the order of parsecs, which is larger than the scale of the shock transitions of interstellar shocks."
+ The fact that we observe sharp transitions. iuplies that the existence of the shock mist be communicated on scales shorter than this 1iean-frec-path.," The fact that we observe sharp transitions, implies that the existence of the shock must be communicated on scales shorter than this mean-free-path."
+ From conservation of mass. omentum aud cuerey over the shock frout. aud assuniug an infinite Mach nuuber and nuo cosmic-ray acceleration. i which ος denotes the shock velocity. aud > the equation of state of the plasma (>=5/3 for a now-relativistic eas).," From conservation of mass, momentum and energy over the shock front, and assuming an infinite Mach number and no cosmic-ray acceleration, in which $v_{\rm s}$ denotes the shock velocity, and $\gamma$ the equation of state of the plasma $\gamma=5/3$ for a non-relativistic gas)."
+" This equation holds for a plasma maximally out of thermal cequilibrimm. miplving an electron temperature over proton temperature ΕΤ Την) of the ratio of the electron and proton mass (eft,= 115060),"," This equation holds for a plasma maximally out of thermal equilibrium, implying an electron temperature over proton temperature $T_{\rm e}/T_{\rm p}$ ) of the ratio of the electron and proton mass $m_{\rm e}/m_{\rm p} = 1/1836$ )."
+ For a plasma in thermal equilibrium. equation reads: iu which pp ds the mean particle mass (0.6 for a fully ionized plasma with solar abundances).," For a plasma in thermal equilibrium, equation reads: in which $\mu$ is the mean particle mass (0.6 for a fully ionized plasma with solar abundances)."
+ However. if a shock is accelerating cosmic ravs. à ΟμΤαν precursor will develop. which heats the incoming imediun and lowers the effective Mach umuber of the shock. in which case equation or are no lounger valid etal.2009:Vink 2010).," However, if a shock is accelerating cosmic rays, a cosmic-ray precursor will develop, which heats the incoming medium and lowers the effective Mach number of the shock, in which case equation or are no longer valid \citep[see also,][]{Drury2009,Vink2010}."
+. Iu nost empirical studies on temperature equilibration οσοιά supernova remnant shocks. Το aud T are determined from Bahluer dominated shocks.," In most empirical studies on temperature equilibration behind supernova remnant shocks, $T_{\rm e}$ and $T_{\rm p}$ are determined from Balmer dominated shocks."
+ These shocks are characterized by bydrogen line cuiission im the Balmer series. leading to a line profile consisting of two superimposed lines.," These shocks are characterized by hydrogen line emission in the Balmer series, leading to a line profile consisting of two superimposed lines."
+" These shocks are non-radiative: their enerev losses iu the form of radiation are dywnanücallv icelieible (thishappensforshockvelocities=200ausbhee,Drame&Melee 19953)..", These shocks are non-radiative; their energy losses in the form of radiation are dynamically negligible \citep[this happens for shock velocities $>$ 200 \kms; e.g. ][]{Draine}.
+ This nuplies hat the hwdroseu lines of these shocks are caused by inupact excitation. not by recombination.," This implies that the hydrogen lines of these shocks are caused by impact excitation, not by recombination."
+ The narrow component of the spectrum is ciitted dy neutral wdrogen atoms. excited shortly after cutering the shock ront: its width reflects the temperature of the ambicut neditun.," The narrow component of the spectrum is emitted by neutral hydrogen atoms, excited shortly after entering the shock front; its width reflects the temperature of the ambient medium."
+ The broad componcut is enuütted after charec exchange between incoming neutral hydrogen and hot xotonus behind the shock trout., The broad component is emitted after charge exchange between incoming neutral hydrogen and hot protons behind the shock front.
+ The width reflects the cluperature of the protous behind the shock frout., The width reflects the temperature of the protons behind the shock front.
+ Combining the width with the ratio of the flux in the warow to the flux in the broad componeut. one cau determine the shock velocity and hence the clectron eniperature (forarecentreview.seeTeng2010).," Combining the width with the ratio of the flux in the narrow to the flux in the broad component, one can determine the shock velocity and hence the electron temperature \citep[for a recent review, see][]{Heng2010}."
+. This xocedure works for shocks which have negligible cosiic-rav acceleration., This procedure works for shocks which have negligible cosmic-ray acceleration.
+ If à shock is an efficient accelerator. he cosmiüc-rav precursor jonizes aud heats the coming uediu aud therewith iucereases the fux of the narrow component (Ravinoudctal.2011).," If a shock is an efficient accelerator, the cosmic-ray precursor ionizes and heats the incoming medium and therewith increases the flux of the narrow component \citep{Raymond2011}."
+. The electrou. temperature for most voung superuova relunauts can also be measured from the N-rav spectrum. as the thermal component of the N-rav spectrum is shaped by the electron temperature.," The electron temperature for most young supernova remnants can also be measured from the X-ray spectrum, as the thermal component of the X-ray spectrum is shaped by the electron temperature."
+ Electron temperatures nieasured from X-ray spectra are generally uot the electron temperatures rieht behind the shock frout. but rather somewhat further to the inside of the shock. where teiiperatures have already. (partially) equilibrated through Coulomb interactions.," Electron temperatures measured from X-ray spectra are generally not the electron temperatures right behind the shock front, but rather somewhat further to the inside of the shock, where temperatures have already (partially) equilibrated through Coulomb interactions."
+ The amount of equilibration is a function of nf: the ionization, The amount of equilibration is a function of $n_{\rm e}t$ : the ionization
+that range fron) wer up to slightly less than the eravitomaenctic precession frequency wey; of a particle at the inner edge of the disk. which is ~1.2«109 for the parameter values used iu the αποσα] computatious reported here.,"that range from $\omega_{\rm crit}$ up to slightly less than the gravitomagnetic precession frequency $\omega_{\rm gm,i}$ of a particle at the inner edge of the disk, which is $\sim1.2\ee{6}$ for the parameter values used in the numerical computations reported here."
+ The TFCOAL modes are essentially unaffected by any raciation-warping torque., The HFGM modes are essentially unaffected by any radiation-warping torque.
+ The damping rates and frequencies of the modes iu both families are chauged by <104 if the l/r Newtoman potential is replaced by auv of several »seudopotentials that mumic the steeper effective eravitational »otentials of general relativity. because the properties of the modes at the iuner edge of the disk are determuiued ΠΑΝ by the eravitomiagnetie torque.," The damping rates and frequencies of the modes in both families are changed by $\lta10$ if the $1/r$ Newtonian potential is replaced by any of several pseudopotentials that mimic the steeper effective gravitational potentials of general relativity, because the properties of the modes at the inner edge of the disk are determined primarily by the gravitomagnetic torque."
+ The vast majority of the modes that lie between he lowest-frequeney LFCAIR modes aud the highest-requency TFGOAL modes are νον shortavaveleneth corrugations of the disk aud are very highly overdamped., The vast majority of the modes that lie between the lowest-frequency LFGMR modes and the highest-frequency HFGM modes are very short-wavelength corrugations of the disk and are very highly overdamped.
+ Iu particular. the modes with frequencies iu the interval have corrmeation wavelengths of the radius at which the mode peaks.," In particular, the modes with frequencies in the interval have corrugation wavelengths $\lta$ of the radius at which the mode peaks."
+ When the wavelength becomes less than the vertical thickness of the disk.H our treatinent oft the diskH as infinitesimalHa: is: uo longer valid (sec. e.g... Papaloizon Pringle 1983).," When the wavelength becomes less than the vertical thickness of the disk, our treatment of the disk as infinitesimal is no longer valid (see, e.g., Papaloizou Pringle 1983)."
+ Tu this regime. modes of the type we have found may be even more strongly damped or even ionexistent.," In this regime, modes of the type we have found may be even more strongly damped or even nonexistent."
+ The halfdozeu lowest-order LFCMR nodes have csscutially the same shapes as the previously kuown low-order radiatiou warping modes (Alaloney et 11996). except that the very innermost part of the disk is forced iuto the rotation equator of the central object bv the combined action of the eravitomaguetic and internal viscous torques. as first discussed by Bardeen Petterson (1975).," The half-dozen lowest-order LFGMR modes have essentially the same shapes as the previously known low-order radiation warping modes (Maloney et 1996), except that the very innermost part of the disk is forced into the rotation equator of the central object by the combined action of the gravitomagnetic and internal viscous torques, as first discussed by Bardeen Petterson (1975)."
+ This is the reason that. unlike the radiation-wiurpiug modes. the low-order LECMB modes do not depeud ou the boundary condition at the inner edge of the cixk.," This is the reason that, unlike the radiation-warping modes, the low-order LFGMR modes do not depend on the boundary condition at the inner edge of the disk."
+ The lowest-order LEGMRB 120des. precess iu the progradedirection with precession of frequencieseither 0.5 or 1.0iu dimensionless wits. depeudingou whether the outer edge of the disk is free or pinned in the disk plane.," The lowest-order LFGMR modes precess in the prograde direction with precession frequencies of either 0.5 or 1.0 in dimensionless units, depending on whether the outer edge of the disk is free or pinned in the disk plane."
+ This is the same precession behavior displaved by the low-order radiation warping modes., This is the same precession behavior displayed by the low-order radiation warping modes.
+ For a compact object of solar iuass. these modes precess with frequeucies ~10.7 Tz.," For a compact object of solar mass, these modes precess with frequencies $\sim 10^{-5}$ Hz."
+ The lowest-order LFCGMR modes. like the low-order radiation warping modes. are erowine modes.," The lowest-order LFGMR modes, like the low-order radiation warping modes, are growing modes."
+ Typically two or three low-order LFCAIR modes have propertics intermediate between the eravitomaguetic modes and the radiation-warpiug modes., Typically two or three low-order LFGMR modes have properties intermediate between the gravitomagnetic modes and the radiation-warping modes.
+ These lybrid modes are simular iu shape to the raciation-warping modes of the same order. except that thei tilt functious eo to zero at the inner οσο of the disk because of the combined action of the eravitomiaenetie aud internal viscous torques.," These hybrid modes are similar in shape to the radiation-warping modes of the same order, except that their tilt functions go to zero at the inner edge of the disk because of the combined action of the gravitomagnetic and internal viscous torques."
+ The lybrid modes generally (but not ahwavs) precess iu the prograde direction and are either weakly damped or weakly erowiug., The hybrid modes generally (but not always) precess in the prograde direction and are either weakly damped or weakly growing.
+ For a compact object of solar mass. the precession frequencies of these inodes lie between 10.6 ITz aud ~10 Πε. which is the frequency of the lowest-order LFGAIR modes.," For a compact object of solar mass, the precession frequencies of these modes lie between $\sim10^{-6}$ Hz and $\sim
+10^{-5}\,$ Hz, which is the frequency of the lowest-order LFGMR modes."
+" With increasing mode προ, the frequeucies of the LFOAIR 10des at first increase steadily but eventually approach weir asviuptotically from below."," With increasing mode number, the frequencies of the LFGMR modes at first increase steadily but eventually approach $\omega_{\rm crit}$ asymptotically from below."
+ This happeus even if the disk is driven at its mner edee., This happens even if the disk is driven at its inner edge.
+ The LFOAIR modes that have precession frequencies =lO Te ave essentially unaffected ly even a strong racdiation-warping torque aud have damping rates zLO funes their frequencies., The LFGMR modes that have precession frequencies $\gta10^{-4}$ Hz are essentially unaffected by even a strong radiation-warping torque and have damping rates $\gta10$ times their frequencies.
+ The highest-frequency, The highest-frequency
+2«0.5. and thus has one of the brightest radio Duxes of an RQQ.,"$z<0.5$, and thus has one of the brightest radio fluxes of an RQQ."
+ The optical. X-ray and near-infrared. properties of the quasar have been studied by Ixolman et ((1993.. 1991).," The optical, X-ray and near-infrared properties of the quasar have been studied by Kolman et (1993, 1991)."
+ The long-term: variability study of Ixolman ct ((1993) shows that the optical and UV flux from the quasar declined between LOST and 1991: our Z£-band data is consistent with a decline from 1988 April 25. when Ixolman et mmoeasured ff=12.19. to our observations in 1993 May 8 (44= 12.6). although due to the errors in the photometry this result should be considered. tentative.," The long-term variability study of Kolman et (1993) shows that the optical and UV flux from the quasar declined between 1987 and 1991; our $H$ -band data is consistent with a decline from 1988 April 25, when Kolman et measured $H=12.19$, to our observations in 1993 May 8 $H=12.6$ ), although due to the errors in the photometry this result should be considered tentative."
+ A further radio image of this object 11) was taken to search for any association of the three radio sources discussed in Paper Ll. None was found. ancl heneclorth they are treated as separate objects.," A further radio image of this object 1) was taken to search for any association of the three radio sources discussed in Paper I. None was found, and henceforth they are treated as separate objects."
+ A very faint. identification has been found for SCT1826|660A 33)., A very faint identification has been found for 8C1826+660A 3).
+ In 55 we plot the #2 relation for the SC galaxies ancl quasars., In 5 we plot the $R-z$ relation for the 8C galaxies and quasars.
+ Although such diagrams as this and the more famous A/—z relation are somewhat “blunt instruments? for studying radio galaxy evolution. where a number of dillerent. cllects may be contributing to the light as a funetion of redshift. they have some useful features which can be exploited if they are used in conjunction with detailec studies of smaller samples of objects.," Although such diagrams as this and the more famous $K-z$ relation are somewhat “blunt instruments” for studying radio galaxy evolution, where a number of different effects may be contributing to the light as a function of redshift, they have some useful features which can be exploited if they are used in conjunction with detailed studies of smaller samples of objects."
+ As expected. the scatter in the &2 relation is fairly low (vith ao few notable exceptions) unti 28O06. where the breakmoresthroughtheR band," As expected, the scatter in the $R-z$ relation is fairly low (with a few notable exceptions) until $z\approx 0.6$, where the $\;$ break moves through the $R$ -band."
+ Aborelhisredshifllhereisawiderangeintheamonntofrest [react Vlighl.ashasbeennoled forthe3C sam ple(Litly& Longairl9084). bulg: ," Above this redshift there is a wide range in the amount of rest-frame UV light, as has been noted for the 3C sample (Lilly Longair 1984), but generally the galaxies are less luminous than their 3C counterparts, consistent with the findings of Dunlop Peacock (1993) for their PSR sample."
+Upto at least z= Lthere seems to be an upper envelope to the magnitudes which traces the expected locus of an old stellar population. consistent with the idea that in many objects the host galaxy is not strongly evolving out to these redshifts.," Up to at least $z=1$ there seems to be an upper envelope to the magnitudes which traces the expected locus of an old stellar population, consistent with the idea that in many objects the host galaxy is not strongly evolving out to these redshifts."
+ Phe presence of a weak alignment ellect in the 7€, The presence of a weak alignment effect in the 7C
+In recent vears many fields of astrophysics have seen a transition towards increasingly [aree experiments and surveys.,In recent years many fields of astrophysics have seen a transition towards increasingly large experiments and surveys.
+ The level of complexity and the costs are rising alongside. requiring careful planning and assessment of he expected performance of the envisaged. project at all stages.," The level of complexity and the costs are rising alongside, requiring careful planning and assessment of the expected performance of the envisaged project at all stages."
+ In forecasts of the statistical constraints on moclel xwameters by future experiments the Fisher matrix (Fisher1935:Teemarketal.1997) has proven to be indispensable (c.&.Albrechtetal.2006).," In forecasts of the statistical constraints on model parameters by future experiments the Fisher matrix \citep{fisher35,tegmark97} has proven to be indispensable \citep[e.g.][]{albrecht06}."
+ lis ubiquity can largely be attributed to. the low computational cost of a Fisher matrix calculation compared o a full mock likelihood. analysis. in particular if the data set to be analysed and the number of parameters to ος inferred are large.," Its ubiquity can largely be attributed to the low computational cost of a Fisher matrix calculation compared to a full mock likelihood analysis, in particular if the data set to be analysed and the number of parameters to be inferred are large."
+ This simplicity comes at. the price of a twofold assumption of Gaussianity in the derivation of the Fisher matrix expressions. (Tegmarkctal.1997)., This simplicity comes at the price of a twofold assumption of Gaussianity in the derivation of the Fisher matrix expressions \citep{tegmark97}.
+. First. the data are assumed. to be distributed according to a multivariate. Gaussian.," First, the data are assumed to be distributed according to a multivariate Gaussian."
+ However. this assumption is shared. with the majority of full likelihood analyses to date although precision measurements may require. more complicated. forms (e.g.Bonectal.2000:Llartlapet 2009).," However, this assumption is shared with the majority of full likelihood analyses to date although precision measurements may require more complicated forms \citep[e.g.][]{bond00,hartlap09}."
+. Second. since the Fisher matrix is defined as the expectation value of the Llessian of the log-likelihood. in parameter space. it can only fully represent a Ciaussian posterior whose logarithm has constant curvature.," Second, since the Fisher matrix is defined as the expectation value of the Hessian of the log-likelihood in parameter space, it can only fully represent a Gaussian posterior whose logarithm has constant curvature."
+ Therefore the confidence levels on model paranieters derived from a Fisher matrix are inevitably elliptical., Therefore the confidence levels on model parameters derived from a Fisher matrix are inevitably elliptical.
+ They can describe the posterior distribution close to the point of maximum likelihood and indicate linear degeneracies among parameters via the ellipticitv of confidence regions., They can describe the posterior distribution close to the point of maximum likelihood and indicate linear degeneracies among parameters via the ellipticity of confidence regions.
+ Fisher matrix analyses fail to identify the shape of the posterior away [rom its maximum. as well as to detect. non-linear dependencies of parameters.," Fisher matrix analyses fail to identify the shape of the posterior away from its maximum, as well as to detect non-linear dependencies of parameters."
+ However. non-linear model parameter degeneracies are common. and the attempt to minimise or break them can drive the design of experiments.," However, non-linear model parameter degeneracies are common, and the attempt to minimise or break them can drive the design of experiments."
+ llence it is desirable to £o. bevond the assumption of a Gaussian posterior in forecasts for the advanced. stages of upcoming precision measurements., Hence it is desirable to go beyond the assumption of a Gaussian posterior in forecasts for the advanced stages of upcoming precision measurements.
+velocity limits.,velocity limits.
+" The Fornax Cluster is isolated in redshift space from both Galactic stars and background galaxies, so with some confidence we define cluster membership by velocity limits of 500 to 2500kms""!."," The Fornax Cluster is isolated in redshift space from both Galactic stars and background galaxies, so with some confidence we define cluster membership by velocity limits of 500 to $2500 \; \mbox{km} \, \mbox{s}^{-1}$."
+" The Virgo Cluster is embedded in the Virgo Supercluster which stretches to include the Local Group, and being less dynamically relaxed than Fornax it even includes galaxies with negative recession velocities."," The Virgo Cluster is embedded in the Virgo Supercluster which stretches to include the Local Group, and being less dynamically relaxed than Fornax it even includes galaxies with negative recession velocities."
+" A few Virgo objects with ‘star-like’ velocities may therefore be cluster CSSs, but considering the 500 to kms""! velocity limits applied in the previous Virgo CSS survey (?), we decided to adopt the same recession velocity limits as Fornax."," A few Virgo objects with `star-like' velocities may therefore be cluster CSSs, but considering the 500 to $\; \mbox{km} \, \mbox{s}^{-1}$ velocity limits applied in the previous Virgo CSS survey \citep{Jones..2006}, we decided to adopt the same recession velocity limits as Fornax."
+" Having identified the cluster members, we rechecked their redshifts with cross-correlationAO software."," Having identified the cluster members, we rechecked their redshifts with cross-correlation software."
+" In general, we achieved velocity errors <30kms""."," In general, we achieved velocity errors $\le\!30 \; \mbox{km} \, \mbox{s}^{-1}$."
+ We examined SDSS or CTIO (?) CCD images of each detected cluster object to confirm their compact nature and local environment., We examined SDSS or CTIO \citep*{Karick..2007} CCD images of each detected cluster object to confirm their compact nature and local environment.
+ Two cluster objects in the Virgo intracluster field were parts of the spiral galaxy NGC 4451 which SDSS incorrectly identified as point sources., Two cluster objects in the Virgo intracluster field were parts of the spiral galaxy NGC 4451 which SDSS incorrectly identified as point sources.
+" One new cluster object located in the Fornax NGC 1399 field was catalogued by APM but not identifiable on the CTIO image, and two other CSSs were outside the CTIO image areas."," One new cluster object located in the Fornax NGC 1399 field was catalogued by APM but not identifiable on the CTIO image, and two other CSSs were outside the CTIO image areas."
+" All other objects were compact, circular and featureless, except for one object in each Virgo field (see Table 4)) which showed evidence of a possible outer stellar envelope."," All other objects were compact, circular and featureless, except for one object in each Virgo field (see Table \ref{table:aaovirgolist}) ) which showed evidence of a possible outer stellar envelope."
+ Table 3 summarises our redshift measurement results and extinction-corrected photometry — our selection criteria have successfully targeted samples of the CSS populations in both clusters., Table \ref{table:results} summarises our redshift measurement results and extinction-corrected photometry – our selection criteria have successfully targeted samples of the CSS populations in both clusters.
+ Tables 4 and 5 provide details of the redshift-confirmed cluster objects., Tables \ref{table:aaovirgolist} and \ref{table:aaofornaxlist} provide details of the redshift-confirmed cluster objects.
+" We found 18 new CSSs in the core regions of Virgo and Fornax, and 3 IGCs in the Virgo intracluster field."," We found 18 new CSSs in the core regions of Virgo and Fornax, and 3 IGCs in the Virgo intracluster field."
+" We found no CSSs in the Fornax intracluster/galaxy merger field, allowing an upper limit to be placed on their distribution."," We found no CSSs in the Fornax intracluster/galaxy merger field, allowing an upper limit to be placed on their distribution."
+" As already explained, for each cluster we used a different method of reducing the large numbers of point source targets in our search for CSSs."," As already explained, for each cluster we used a different method of reducing the large numbers of point source targets in our search for CSSs."
+ The results of our colour- selection method in the Virgo Cluster are evident in Fig., The results of our colour-colour selection method in the Virgo Cluster are evident in Fig.
+ 6 which compares the distribution of CSSs with that of foreground stars and background galaxies located by our observations., \ref{fig:Virgogri} which compares the distribution of CSSs with that of foreground stars and background galaxies located by our observations.
+ The Galactic stellar locus evident in both Virgo fields contrasts with the widely distributed background galaxy colours., The Galactic stellar locus evident in both Virgo fields contrasts with the widely distributed background galaxy colours.
+ CSSs detected in our survey are not distributed along the stellar locus., CSSs detected in our survey are not distributed along the stellar locus.
+" Based on the positions of detected CSSs, we could improve our targeting efficiency in future surveys by restricting targets to g—r 0.8."," Based on the positions of detected CSSs, we could improve our targeting efficiency in future surveys by restricting targets to $g - r \le 0.8$ ."
+" In the Fornax Cluster, Fig."," In the Fornax Cluster, Fig."
+ 7 shows the results of the colour-magnitude selection method we adopted., \ref{fig:Fornax_52} shows the results of the colour-magnitude selection method we adopted.
+" In the previously observed NGC 1399 field we extended our existing 2dF redshift data to the upper end of the GC distribution at b;x21.2 (Με,x -—10.2), whereas for the NGC 1316 field with limited prior point-source redshift data we concentrated on the brighter targets to obtain 7996 "," In the previously observed NGC 1399 field we extended our existing 2dF redshift data to the upper end of the GC distribution at $b_J \le 21.2$ $M_{b_J} \le -10.2$ ), whereas for the NGC 1316 field with limited prior point-source redshift data we concentrated on the brighter targets to obtain $79\%$ "
+Clusters of galaxies are characterized by X-ray emission from a hot intra-cluster medium (ICM. 7~2-10 keV).,"Clusters of galaxies are characterized by X-ray emission from a hot intra-cluster medium (ICM, $T \sim 2{-}10$ keV)."
+ Thermal emission is à common property of all clusters of galaxies and has been detected even in poor groups as well as in optical filaments connecting rich clusters., Thermal emission is a common property of all clusters of galaxies and has been detected even in poor groups as well as in optical filaments connecting rich clusters.
+ The ICM ts believed to be shock heated by merging during the structure formation., The ICM is believed to be shock heated by merging during the structure formation.
+ Diffuse non-thermal radio sources with steep spectra have been detected in the central regions of à large number of clusters of galaxies (see e.g. Giovanninietal.2009.. Venturietal.2008.— (and refs. therein) and review papers such as Peretti2005... and Ferrarietal. 2008)).," Diffuse non-thermal radio sources with steep spectra have been detected in the central regions of a large number of clusters of galaxies (see e.g. \cite{gio09}, \cite{ven08} (and refs therein) and review papers such as \cite{fer05}, and \cite{ferra08}) )."
+ They are giant radio sources with a spatial extent similar to that of the hot ICM. which are called radio halos.," They are giant radio sources with a spatial extent similar to that of the hot ICM, which are called radio halos."
+ In addition. diffuse radio sources classified as radio relics are detected in cluster peripheral regions (see e.g. Giovannint&Feretti2004 and vanWeerenetal. 2011)).," In addition, diffuse radio sources classified as radio relics are detected in cluster peripheral regions (see e.g. \cite{gio04} and \cite{vwe11}) )."
+ These sources are not directly associated with the activity of individual galaxies and are related to physical properties of the whole cluster., These sources are not directly associated with the activity of individual galaxies and are related to physical properties of the whole cluster.
+" There is a substantial evidence that radio halos are found in clusters with significant substructures in X-ray images and complex gas temperature distributions. both signatures of cluster mergers (e.g. Feretti1999,, Govonietal. 2004.. Giovanninietal. 2009.. Cassanoetal. 2010))."," There is a substantial evidence that radio halos are found in clusters with significant substructures in X-ray images and complex gas temperature distributions, both signatures of cluster mergers (e.g. \cite{fer99}, \cite{gov04}, \cite{gio09}, \cite{cas10}) )."
+ Classical radio halos show a centrally located and regular radio morphology., Classical radio halos show a centrally located and regular radio morphology.
+ However. in. several cases the radio halos are irregular and elongated as in A209. A401. and A2034.," However, in several cases the radio halos are irregular and elongated as in A209, A401, and A2034."
+ A few radio halos such as A85] and A2218 are highly asymmetric. being located mostly to one side with respect to the cluster center.," A few radio halos such as A851 and A2218 are highly asymmetric, being located mostly to one side with respect to the cluster center."
+ Giovanninietal.2009. interpreted these structures as possibly reflecting the geometry of the merger., \cite{gio09} interpreted these structures as possibly reflecting the geometry of the merger.
+ Feretti20060 discussed that the radio power of halos at 1.4 GHz increases with the X-ray luminosity of the parent cluster. implying a direct connection between the radio and X-ray plasmas.," \cite{fer00} discussed that the radio power of halos at 1.4 GHz increases with the X-ray luminosity of the parent cluster, implying a direct connection between the radio and X-ray plasmas."
+ Since the cluster X-ray luminosity and mass are correlated. it follows that radio halo power correlates with cluster mass (Feretti2000: Govonietal. 2001)).," Since the cluster X-ray luminosity and mass are correlated, it follows that radio halo power correlates with cluster mass \cite{fer00}; \cite{gov01}) )."
+" The strong correlation between the radio halo total power at 1.4 GHz and the X-ray cluster luminosity was recently discussed using higher quality statisticsby Brunettietal.2007 and Giovanninietal.2009... who show that the correlation is also present for low power radio halos (Pj. < 107 W/Hz and L, ~ 107 erg/s)."," The strong correlation between the radio halo total power at 1.4 GHz and the X-ray cluster luminosity was recently discussed using higher quality statisticsby \cite{bru07} and \cite{gio09}, who show that the correlation is also present for low power radio halos $_{1.4}$ $<$ $^{24}$ W/Hz and $_x$ $\sim$ $^{44}$ erg/s)."
+ Giovanninietal.2009 were the first to discuss whether a radio halo was present in AI213. an under-Iuminous X-ray cluster te. well outside the above correlation.," \cite{gio09} were the first to discuss whether a radio halo was present in A1213, an under-luminous X-ray cluster i.e. well outside the above correlation."
+ A similar case was found in the cluster 0217-70 by Brownetal.2011.., A similar case was found in the cluster 0217+70 by \cite{bro11}.
+ We do not consider radio relics because present models connect radio halos to turbulence in the ICM. while relie sources are correlated to the presence of shock waves.," We do not consider radio relics because present models connect radio halos to turbulence in the ICM, while relic sources are correlated to the presence of shock waves."
+ Brownetal.2011 discussed whether the low X-ray luminosity in 0217470 could be due to an underestimate of the X-ray luminosity because of the hydrogen column density. which might be patchy at these low Galactic latitudes.," \cite{bro11} discussed whether the low X-ray luminosity in 0217+70 could be due to an underestimate of the X-ray luminosity because of the hydrogen column density, which might be patchy at these low Galactic latitudes."
+ Therefore. further X-ray observations are required to measure the absorption directly.," Therefore, further X-ray observations are required to measure the absorption directly."
+ We note that Giovanninietal.2009. discussed the peculiarity of the diffuse emission in Al213., We note that \cite{gio09} discussed the peculiarity of the diffuse emission in A1213.
+ We present new VLA images from archive data of diffuse radio emission. found by chance by inspecting NVSS images. in A523. one more radio halo in à low X-ray luminosity cluster.," We present new VLA images from archive data of diffuse radio emission, found by chance by inspecting NVSS images, in A523, one more radio halo in a low X-ray luminosity cluster."
+ We present in Sect., We present in Sect.
+ 2 the radio data and the X-ray and optical information., 2 the radio data and the X-ray and optical information.
+ Discussion and conclusions are given in Sect., Discussion and conclusions are given in Sect.
+ 3., 3.
+" The intrinsic parameters quoted in this paper are computed for a ACDM cosmology with Ho 2 71 km ! Mpe!. Q,, = 0.27. and O4 = 0.73."," The intrinsic parameters quoted in this paper are computed for a $\Lambda$ CDM cosmology with $H_0$ = 71 km $^{-1}$ $^{-1}$, $\Omega_m$ = 0.27, and $\Omega_{\Lambda}$ = 0.73."
+ At z =0.10. the luminosity distance is 455 Mpc. and the angular conversion factor is 1.82 kpe/aresec.," At z =0.10, the luminosity distance is 455 Mpc, and the angular conversion factor is 1.82 kpc/arcsec."
+ A523 was observed at 1.4 GHz in line mode in 2005. with the VLA in the C (5 hrs) and D (2 hrs) configuration (project ABIISO).," A523 was observed at 1.4 GHz in line mode in 2005, with the VLA in the C (5 hrs) and D (2 hrs) configuration (project AB1180)."
+ Calibration and imaging were performed with the NRAO Astronomical Image Processing System (AIPS)., Calibration and imaging were performed with the NRAO Astronomical Image Processing System (AIPS).
+ After the editing of bad points. self-calibration was applied to both datasets. to remove residual phase and gain variations.," After the editing of bad points, self-calibration was applied to both datasets, to remove residual phase and gain variations."
+ In Fig. 1. ," In Fig. \ref{fig:1}, ,"
+we show the final image obtained with C configuration data and natural weights (contours). superimposed on the X-ray image (colors) obtained with the Rosat satellite.," we show the final image obtained with C configuration data and natural weights (contours), superimposed on the X-ray image (colors) obtained with the Rosat satellite."
+ The diffuse emission ts also visible at this, The diffuse emission is also visible at this
+the other TTS in the cloud have been previously detected nu X-ravs.,the other TTS in the cloud have been previously detected in X-rays.
+ We present the satellite observations iu Sect., We present the satellite observations in Sect.
+ 2 that show the distance to MDMI2 is not as well constrained as previously thought aud our high resolutiou observations of two TTS in the cloud that indicate the TTS are probably at the sane distance as the cloud., \ref{hipsec} that show the distance to MBM12 is not as well constrained as previously thought and our high resolution observations of two TTS in the cloud that indicate the TTS are probably at the same distance as the cloud.
+ Iu Sect., In Sect.
+ 3. we describe theROSAT All-Skyv Survey (RASS) andROSAT pointed observations investigated., \ref{rassobs} we describe the All-Sky Survey (RASS) and pointed observations investigated.
+ In Sect., In Sect.
+ 1 we prescut the results of our low-resolution spectroscopic observations of the T Tauri star candidates., \ref{cafos} we present the results of our low-resolution spectroscopic observations of the T Tauri star candidates.
+ In Sect., In Sect.
+ we cliscuss the N-rav variability of the TTS aud im Sect., \ref{xrayvar} we discuss the X-ray variability of the TTS and in Sect.
+ 6 we derive an N-ray luuinosity function for the TTS iu AIDMI2.," \ref{xlf}
+ we derive an X-ray luminosity function for the TTS in MBM12."
+ In Sect., In Sect.
+ 7 we present our conclusious., \ref{conclusions} we present our conclusions.
+ The distauce to MDBMI2 was first estimated by Duerr Craine (L982a) using Wolff diagrams., The distance to MBM12 was first estimated by Duerr Craine (1982a) using Wolff diagrams.
+ They found evidence for two clouds along this hue of sieht: oue cloud at a distance of 200.300 pc with a vpical visual extinction of ~23 mae and another cloud at a distance of 500SOU pe with a typical visual extinction of  1.5 mae., They found evidence for two clouds along this line of sight: one cloud at a distance of 200–300 pc with a typical visual extinction of $\sim$ 2–3 mag and another cloud at a distance of 500–800 pc with a typical visual extinction of $\sim$ 1.5 mag.
+ However. more recent observations show that this is probably one coniplex of molecular gas at a 11ο sinaller distance.," However, more recent observations show that this is probably one complex of molecular gas at a much smaller distance."
+ A μονο accurate estimate ο ‘the distance was reported bv Hobbs et al. (, A more accurate estimate of the distance was reported by Hobbs et al. (
+1986) aud Hobbs et al. (,1986) and Hobbs et al. (
+19858).,1988).
+ The inethod used was ο identify 10 bright stars in the direction of the cloud for which a parallax could be deteriuued aud look for iuterstcllar Nal D lines in the spectra of these stars., The method used was to identify 10 bright stars in the direction of the cloud for which a parallax could be determined and look for interstellar NaI D lines in the spectra of these stars.
+ They found hat the star IIDIslol (distance ~ 660 pe) showed uo interstellar absorption features and is therefore presumev àin front of the cloud and the star IIDI8519/20 (distauce ~ 770 pe) did show interstellar absorption features aud is therefore behiud the cloud., They found that the star HD18404 (distance $\sim$ 60 pc) showed no interstellar absorption features and is therefore presumably in front of the cloud and the star HD18519/20 (distance $\sim$ 70 pc) did show interstellar absorption features and is therefore behind the cloud.
+ Since observations of the other 8 stars are consistent with these two stars. most investigators lave assed a distance of~ 65+5 pe for the cloud.," Since observations of the other 8 stars are consistent with these two stars, most investigators have assumed a distance of $\sim$ $65\pm5$ pc for the cloud."
+ Simce the satellite measured theΙΕ: parallax of LOS of these stars df ds uo longer necessary to assiune a spectral type or intrinsic luumositv (as is necessary for a spectroscopic parallax) to measure their distance., Since the satellite measured the parallax of most of these stars it is no longer necessary to assume a spectral type or intrinsic luminosity (as is necessary for a spectroscopic parallax) to measure their distance.
+ According to Lpparcos. the distance to IIDISlol is ~ 3241 pe aud the distance to IID18519/20 is — 90412 pe.," According to , the distance to HD18404 is $\sim$ $\pm$ 1 pc and the distance to HD18519/20 is $\sim$ $\pm$ 12 pc."
+ The stars used by Hobbs et al. (, The stars used by Hobbs et al. (
+1986) aud. Hobbs et al. (,1986) and Hobbs et al. (
+1988) to establish the distance to MDMI2 are listed in Table 2 alone with their apparent magnitude. spectral type. distance based on spectroscopic parallax. distance based on the parallax. aud whether the spectrum presented in Iobbs et al. (,"1988) to establish the distance to MBM12 are listed in Table 2 along with their apparent magnitude, spectral type, distance based on spectroscopic parallax, distance based on the parallax, and whether the spectrum presented in Hobbs et al. ("
+1986) aud Hobbs et al. (,1986) and Hobbs et al. (
+1988) showed iuterstellax absorption lines.,1988) showed interstellar absorption lines.
+ Although theος results indicate the distance to NDMMI2 is not as well coustrained as thought. it is consistent with the previous result (he. ~ 65435 pe).," Although the results indicate the distance to MBM12 is not as well constrained as thought, it is consistent with the previous result (i.e., $\sim$ $\pm$ 35 pc)."
+ However. since the revised distauce estimate has ~ uncertainty it cannot be used to derive accurate parameters for the TTS iu the cloud.," However, since the revised distance estimate has $\sim$ uncertainty it cannot be used to derive accurate parameters for the TTS in the cloud."
+ Zunuunieruuuis Cneerechts (1990). poiut out that. although the distance estimate based on the two bright stars from Hobbsetal.(1986). and Hobbsetal.(1988). is valid for the northern section of the cloud. radio observations of the molecular eas show there aro at least Ll velocity coniponents that nav all be at cüffereut clistances.," Zimmermann Ungerechts (1990) point out that, although the distance estimate based on the two bright stars from \cite{hob86} and \cite{hob88} is valid for the northern section of the cloud, radio observations of the molecular gas show there are at least 4 velocity components that may all be at different distances."
+ Towever. the polarization map of MDMI2 produced bx Bhatt Jain (1992) shows that he polarization of the stars in the upper region of the cloud (where the Πο]pa et al.," However, the polarization map of MBM12 produced by Bhatt Jain (1992) shows that the polarization of the stars in the upper region of the cloud (where the Hobbs et al."
+ 19856 aud IHobbs et al., 1986 and Hobbs et al.
+ 1988 stars axe located) is the sale at iat of the lower region othe cloud., 1988 stars are located) is the same at that of the lower region of the cloud.
+ Thus. they areuc that the eutire complex is at the same cistance.," Thus, they argue that the entire complex is at the same distance."
+ Ziuunnermiuun Unuescerechts (1990) ραo find that the ratio of the CO mass to the viria mass λίαν = 0.03 for 1e whole cloud indicating the eutire cloud is not eravitatiouallv bound., Zimmermann Ungerechts (1990) also find that the ratio of the CO mass to the virial mass $_{\rm vir}$ = 0.03 for the whole cloud indicating the entire cloud is not gravitationally bound.
+ However. they poit out that some of the cores in the larger clumps may be eravitationally bound.," However, they point out that some of the cores in the larger clumps may be gravitationally bound."
+ If the cloud is found to be at a somewhat ercater distance. these results will support the livpothesis that a few of the large chumps are evavitationally bound.," If the cloud is found to be at a somewhat greater distance, these results will support the hypothesis that a few of the large clumps are gravitationally bound."
+ Although the TTS and the cloud are projected along the same line of sight. they may be at different locations aud thus not associated with each other.," Although the TTS and the cloud are projected along the same line of sight, they may be at different locations and thus not associated with each other."
+ However. we can check whether the radial velocities of the stars are simular with the eas to find out if it is likely that the T. Tami stars are assoclatec with the cloud.," However, we can check whether the radial velocities of the stars are similar with the gas to find out if it is likely that the T Tauri stars are associated with the cloud."
+ The distribution of cloud velocities are quite large with at least. 1 coniponeuts with velocities | 5.0 kan 1 (D. 3 kau 1 (ID. l.l kn average(IID. and Pouud5.0 kau (IV) (ο”.m Cueerechts 1990: et al.," The distribution of cloud velocities are quite large with at least 4 components with average velocities $-$ 5.0 km $^{-1}$ (I), $-$ 2.3 km $^{-1}$ (II), 1.4 km $^{-1}$ (III), and 5.0 km $^{-1}$ (IV) (Zimmermann Ungerechts 1990; Pound et al."
+ 1990: Moriarty-Schievenu et al.1997)., 1990; Moriarty-Schieven et al.1997).
+ We obtained Ligh resolution spectra of two of the T Tauri stars in MDMI2 with FOCES at the Calar, We obtained high resolution spectra of two of the T Tauri stars in MBM12 with FOCES at the Calar
+value of the central asymptotic behaviour of the halo density profile (c.g. Fukushige&Makino1997:Mooreοἱal.1998:llavashietal. 2004)).,"value of the central asymptotic behaviour of the halo density profile (e.g. \citealt{FM97,Moore98,Gea00,JS00,P03,Hayaea04}) )."
+ More recently. the Einasto(1965) profile was shown to give even better fits down to smaller paci (Navarroetal.2004.2010:Merrittct2005.2006).," More recently, the \citet{E65} profile was shown to give even better fits down to smaller radii \citep{Navea04,Navea10,M05,M06}."
+. The origin of this densitv profile remains to. be understood., The origin of this density profile remains to be understood.
+ Certainly. it can only arise from the way haloes aggregate their mass. which. in hierarchical cosmologies. is through continuous mergers.," Certainly, it can only arise from the way haloes aggregate their mass, which, in hierarchical cosmologies, is through continuous mergers."
+ But the dynamical elfect of mergers depends. on the relative mass of the captured and capturing haloes. AAZ/AL.," But the dynamical effect of mergers depends on the relative mass of the captured and capturing haloes, $\Delta M/M$."
+ For this reason. it is usually distinguished between minor and. major mergers (AALSAL &reater or less than about 0.3: Salvaclor-Soléctal. 2007)).," For this reason, it is usually distinguished between minor and major mergers $\Delta M/M$ greater or less than about 0.3; \citealt{smgh07}) )."
+ Alajor mergers have a dramatic effect each on the structure of the object., Major mergers have a dramatic effect each on the structure of the object.
+ While minor mergers contribute jointly. together with the capture of dilluse matter (if any). to the so-called accretion that vields only a small secular etlect on the accreting object.," While minor mergers contribute jointly, together with the capture of diffuse matter (if any), to the so-called accretion that yields only a small secular effect on the accreting object."
+ Some authors (Sver&White1998:Raigetal.1998:Subramanian2000:Dekel2003) studied the possibility that the non-power-law density profile for simulated haloes is the result of repeated major mergers.," Some authors \citep{SW98,RGS98,Suea00,Dea03} studied the possibility that the non-power-law density profile for simulated haloes is the result of repeated major mergers."
+ Others (Avila-Reeseetal.1998:LIuss1999:DelPopolo concentrated on the ellects of pure accretion (PA). that proceeds according to the simple SI model above from peaks (secondary maxima) in the primordial random. Gaussian density field (Doroshkevich1970:Bardeenetal.," Others \citep{ARea98,Hea99,DPea00,As04,McM06} concentrated on the effects of pure accretion (PA), that proceeds according to the simple SI model above from peaks (secondary maxima) in the primordial random Gaussian density field \citep{D70,BBKS}."
+1986).. Phe density. profile found in this latter scenario appears to be similar. indeed. to that of simulated. haloes.," The density profile found in this latter scenario appears to be similar, indeed, to that of simulated haloes."
+ Furthermore. recent cosmological simulations have confirmed that major mergers plav no central role in setting the structure and kinematies of dark matter haloes (Wang&White2009).," Furthermore, recent cosmological simulations have confirmed that major mergers play no central role in setting the structure and kinematics of dark matter haloes \citep{WW09}."
+. Llowever. the reason why major mergers would not alter the density. profile set by P is not understood.," However, the reason why major mergers would not alter the density profile set by PA is not understood."
+ In the present. paper. we develop an accurate model of halo density profile that clarifies all these issues.," In the present paper, we develop an accurate model of halo density profile that clarifies all these issues."
+ For simplicity. we consider pure dark matter svstems. that is we neglect the ellects of barvons on halo structure.," For simplicity, we consider pure dark matter systems, that is we neglect the effects of baryons on halo structure."
+ This will allow us to cürectlv compare the theoretical predictions of the model to the results of N-body simulations., This will allow us to directly compare the theoretical predictions of the model to the results of $N$ -body simulations.
+ This model is built within the SL framework and. uses the peak formalism. which assumes there is a one-to-one mapping between haloes and the density peaks.," This model is built within the SI framework and uses the peak formalism, which assumes there is a one-to-one mapping between haloes and the density peaks."
+ However. peaks are not assumed to be spherically svnimetric and endowed with their typicalfiltered density. profile calculated by Bardeenetal. (1986... hereafter the BDBINS. profile) as usual. but we take into account that they are triaxial (Doroshkevich1970:Bardeenetal.1986). and consider their accurate spherically averageduniconmvoteed density profile.," However, peaks are not assumed to be spherically symmetric and endowed with their typical density profile calculated by \citeauthor{BBKS} \citeyear{BBKS}, hereafter the BBKS profile) as usual, but we take into account that they are triaxial \citep{D70,BBKS} and consider their accurate spherically averaged density profile."
+ In addition. we account for the initial peculiar velocities believed to alTect the central halo density. profile (Lliotelis2002:Ascasibaretal.2004:Mikheevact 2007).," In addition, we account for the initial peculiar velocities believed to affect the central halo density profile \citep{H02,As04,MDL07}."
+. Finally. instead. of making use of the usual adiabatic invariant. poorly justified. near turnaround and eeeatIy complicating the analvtie derivation of the final clensitw profile. we take advantage of the fact that accreting haloes develop outwardly. keeping their instantaneous inner structure unaltered.," Finally, instead of making use of the usual adiabatic invariant, poorly justified near turnaround and greatly complicating the analytic derivation of the final density profile, we take advantage of the fact that accreting haloes develop outwardly, keeping their instantaneous inner structure unaltered."
+ Phe valicity of such an assumption. hereafter simply. referred. to as inside-out &rowth. is accurately checked by verifving that the trends observed in simulated haloes evolving by smooth accretion are reproduced by an analytic model where the haloes are forced to grow inside-out.," The validity of such an assumption, hereafter simply referred to as inside-out growth, is accurately checked by verifying that the trends observed in simulated haloes evolving by smooth accretion are reproduced by an analytic model where the haloes are forced to grow inside-out."
+ And what about the fact that major mergers are ignored in SI?, And what about the fact that major mergers are ignored in SI?
+ This important caveat allects not only the modelling through SI of the halo density profile. as mentioned. but also of halo statistics in the so-called peak formalism (Peacock&Solé1995.1996.. hereafter. MSSa and. MSSb. respectively: Alanriqueetal.1998)) which relies on the Ansatz that here is à one-to-one correspondence between haloes and »aks as in sphericalcollapse.," This important caveat affects not only the modelling through SI of the halo density profile, as mentioned, but also of halo statistics in the so-called peak formalism \citealt{PH90,bm,MSS95,MSS96}, hereafter MSSa and MSSb, respectively; \citealt{metal98}) ) which relies on the Ansatz that there is a one-to-one correspondence between haloes and peaks as in spherical."
+.. Although no such one correspondence is actually found in simulations (Ixatzetal.1993).. this is because peaks show nested configurations (AISSa) which are not corrected for.," Although no such one-to-one correspondence is actually found in simulations \citep{Kea93}, this is because peaks show nested configurations (MSSa) which are not corrected for."
+ I£ one concentrates in simply checking whether or not halo seeds. hereafter. also called. protohaloes. coincide with peaks. then the answer is »ositive in all but à few percent of cases compatible with the requeney of non-Lully virialised haloes (Porcianietal2002).," If one concentrates in simply checking whether or not halo seeds, hereafter also called protohaloes, coincide with peaks, then the answer is positive in all but a few percent of cases compatible with the frequency of non-fully virialised haloes \citep{Pea}."
+. Llowever. there is still the problem that. in principle. major mergers should blur the correspondence between peaks and woes arising from SL.," However, there is still the problem that, in principle, major mergers should blur the correspondence between peaks and haloes arising from SI."
+ In the present paper. we argue that he reason why halo statistics and the spherically average density profile of virialised haloes seem to be unalfected » major mergers is that virialisation is a real relaxation orocess.," In the present paper, we argue that the reason why halo statistics and the spherically average density profile of virialised haloes seem to be unaffected by major mergers is that virialisation is a real relaxation process."
+ We stress that the mocdoel deals withfull virialised woes., We stress that the model deals with virialised haloes.
+ This is important because. after à major merger. the density. profile of à halo spends some time to adopt a stable density profile.," This is important because, after a major merger, the density profile of a halo spends some time to adopt a stable density profile."
+ The outline of the paper is as follows., The outline of the paper is as follows.
+" In Section οον, we present some e&eneral relations holding for all svstemis regardless of their svmmetry."," In Section \ref{general}, we present some general relations holding for all systems regardless of their symmetry."
+ Taking into account these relations. we develop in Section ??. the model for triaxial haloes grown by PA.," Taking into account these relations, we develop in Section \ref{model} the model for triaxial haloes grown by PA."
+ In Section. ??.. we calculate. the properties of typical protohaloes ancl use them in Section 2? to derive the typical spherically averaged density profile for CDAL haloes.," In Section \ref{peaks}, we calculate the properties of typical protohaloes and use them in Section \ref{haloes} to derive the typical spherically averaged density profile for CDM haloes."
+ In Section ??.. we discuss the cllects of major mergers.," In Section \ref{discussion}, we discuss the effects of major mergers."
+ The main results are summarised in Section ?2.., The main results are summarised in Section \ref{summ}.
+ Throughout the paper. we adopt the ACDAL Winap7 (Ixomatsuetal.2011) concordance model.," Throughout the paper, we adopt the $\Lambda$ CDM Wmap7 \citep{Km11} concordance model."
+" In spite of this, we use Newtonian dvnamies with null cosmological constant as its effects. are irrelevant at. the scale of Virialised haloes."," In spite of this, we use Newtonian dynamics with null cosmological constant as its effects are irrelevant at the scale of virialised haloes."
+ A package with the numerical codes used. in the present paper is publicly available fromwww., A package with the numerical codes used in the present paper is publicly available from.
+am.ub.es/~cosmo/haloeskpeaks.tgz. In the present section. we derive some general relations or spherically averaged quantities that hold. for any arbitrary system regardless of its svmunetry and that will ater be used to build the mocel.," In the present section, we derive some general relations for spherically averaged quantities that hold for any arbitrary system regardless of its symmetry and that will later be used to build the model."
+ Consider a sell-gravitating svstem with arbitrary mass distribution. aggregation history and dynamical state.," Consider a self-gravitating system with arbitrary mass distribution, aggregation history and dynamical state."
+ The ocal density ancl gravitational potential can be split in the, The local density and gravitational potential can be split in the
+The following svstem of equations (cf.,The following system of equations (cf.
+ equations (1)) (2))) describes the propagation of the r mode and n=0 intermediate mode. coupled by the warp. and needs to be solved for the growth rate to be determined: where are the coupling terms. arising from non-lincaritics in the basic equations.," equations \ref{motion1}) \ref{energy1}) )) describes the propagation of the r mode and $n=0$ intermediate mode, coupled by the warp, and needs to be solved for the growth rate to be determined: where are the coupling terms, arising from non-linearities in the basic equations."
+ Since we are interested in studying the axisvmmetric r mode. the azimuthal mode number for the intermediate mocle is 1.," Since we are interested in studying the axisymmetric r mode, the azimuthal mode number for the intermediate mode is $1$."
+" Also. since the simplest possible r mode has one node in the vertical direction and the intermediate mode has n=0. we use the following separation of variables: This separation of variables results in having the coupling terms (AS)) and. (A9)) proportional to Le, only. while the coupling terms CXIO)) CXI1)) eive rise to terms proportional to both Heo,=1 and He»=SH?1."," Also, since the simplest possible r mode has one node in the vertical direction and the intermediate mode has $n=0$, we use the following separation of variables: This separation of variables results in having the coupling terms \ref{coup1}) ) and \ref{coup2}) ) proportional to $\textrm{He}_1$ only, while the coupling terms \ref{coup3}) \ref{coup4}) ) give rise to terms proportional to both $\textrm{He}_0=1$ and $\textrm{He}_2=z^2/H^2-1$."
+ Since we are interested in the terms that influence the mode with η=0. we need to project these forcing terms on to Lleg.," Since we are interested in the terms that influence the mode with $n=0$, we need to project these forcing terms on to $\textrm{He}_0$."
+ Also. the separation of variables results in coupling terms of the form which means that the interaction of the m=0 r mode with the m=1 warp results in two new moces. one with m=1 and one with m=1: Since we are only interested in the action of this coupling on the intermediate mode with nm=1. because the one with m=1 coes not have any. Lindblad resonances (location where the wave should be excited) in the disc. these forcing terms," Also, the separation of variables results in coupling terms of the form which means that the interaction of the $m=0$ r mode with the $m=1$ warp results in two new modes, one with $m=1$ and one with $m=-1$: Since we are only interested in the action of this coupling on the intermediate mode with $m=1$, because the one with $m=-1$ does not have any Lindblad resonances (location where the wave should be excited) in the disc, these forcing terms"
+"C,,,í, uncertainties are in light gray, the ones induced by observational uncertaintiesin Rey and Y,,,; are in medium light gray and the theoretical error envelope is in medium dark gray.","$c_{sound}$ uncertainties are in light gray, the ones induced by observational uncertaintiesin $R_{CZ} $ and $Y_{surf}$ are in medium light gray and the theoretical error envelope is in medium dark gray."
+ The total error envelope is represented in dark gray in the bottom panel., The total error envelope is represented in dark gray in the bottom panel.
+ Our detailed procedure is described below., Our detailed procedure is described below.
+" In order to estimate the effect of the different physics ingredients onto the c,,,,; profile. we changed one at a time and adjusted the composition in the same manner as in DPO6 to recover the seismic Rey and Y,,,,."," In order to estimate the effect of the different physics ingredients onto the $c_{sound}$ profile, we changed one at a time and adjusted the composition in the same manner as in DP06 to recover the seismic $R_{CZ} $ and $Y_{surf}$."
+" The differences in c,,,,; between these calibrated models are then measures of the impact of theoretical uncertainties on the c,,,, profile.", The differences in $c_{sound}$ between these calibrated models are then measures of the impact of theoretical uncertainties on the $c_{sound}$ profile.
+ We adopt the same error estimates and central values as described in Table I of DPO6 with the following exceptions., We adopt the same error estimates and central values as described in Table I of DP06 with the following exceptions.
+ The errors and central values of the nuclear reaction cross sections have been taken from the sources cited above and for the equation of state the errors are derived from the differences in the prediction between the EOS2001 and EOS2006 tables., The errors and central values of the nuclear reaction cross sections have been taken from the sources cited above and for the equation of state the errors are derived from the differences in the prediction between the EOS2001 and EOS2006 tables.
+ We used the differences in the calibrated model sound speed profiles to infer the envelope of uncertainties., We used the differences in the calibrated model sound speed profiles to infer the envelope of uncertainties.
+" From the top panel of Figure 1, we see the uncertainties in the diffusion coefficients for gravitational settling (dot-dashed line) are the most significant error source, followed by the uncertainties in the opacities (solid line)."," From the top panel of Figure 1, we see the uncertainties in the diffusion coefficients for gravitational settling (dot-dashed line) are the most significant error source, followed by the uncertainties in the opacities (solid line)."
+" The observational uncertainties in the solar c,,,, are deduced from the range of inverted solar structures in Basuetal.(2000) (see Figure |, bottom panel. light gray envelope)."," The observational uncertainties in the solar $c_{sound}$ are deduced from the range of inverted solar structures in \citet{Basu2000} (see Figure 1, bottom panel, light gray envelope)."
+" Errors in Rey and Y,,, also impact the c,,,; profile.", Errors in $R_{CZ} $ and $Y_{surf}$ also impact the $c_{sound}$ profile.
+ The fractional depth of the convection zone is Rez=0.7133£0.0005R. (Basu&Antia2004)., The fractional depth of the convection zone is $R_{CZ}=0.7133 \pm 0.0005\ R_{\sun}$ \citep{BA2004}.
+. The theoretical location of the convection zone depth can change by up to 0.001 with different opacity interpolation methods 2004); we therefore treat the latter as the uncertainty in Rey., The theoretical location of the convection zone depth can change by up to 0.001 with different opacity interpolation methods \citep{BSP04}; we therefore treat the latter as the uncertainty in $R_{CZ}$.
+" The helium abundance in the surface convection zone (Y,,,,=0.2483+ 0.0046) is measured from the strength of the depression in the adiabatic gradient in the helium ionization zone (Basu1998;Richardetal.Di discussion)..."," The helium abundance in the surface convection zone $Y_{surf}=0.2483 \pm 0.0046$ ) is measured from the strength of the depression in the adiabatic gradient in the helium ionization zone \citep[][see DP06 for discussion]{Basu1998,Richard1998,DiMauro2002}. ."
+" The observational errors in Rey and ες permit a range of derived compositions which will then generates different predicted c,,,, profiles.", The observational errors in $R_{CZ}$ and $Y_{surf}$ permit a range of derived compositions which will then generates different predicted $c_{sound}$ profiles.
+ The, The
+" Figure] is a plot of the mean values of the alpha-particle temperature anisotropy, Τια/Τια. the proton temperature anisotropy, Τιρ/Τι». and the mean relative wave power versus Vap/Va.","Figure \ref{fig.4} is a plot of the mean values of the alpha-particle temperature anisotropy, $T_{\perp \alpha}/T_{\parallel \alpha}$, the proton temperature anisotropy, $T_{\perp p}/T_{\parallel p}$, and the mean relative wave power versus $V_{\alpha p}/V_\mathrm{A}$."
+" The mean values are obtained by averaging the data within bins of a width Δ(γαρ/γα)=0.1, and the vertical error bars indicate the related uncertainties of the mean values."," The mean values are obtained by averaging the data within bins of a width $\Delta (V_{\alpha
+p}/V_\mathrm{A}) = 0.1$, and the vertical error bars indicate the related uncertainties of the mean values."
+" As mentioned above, it turns out that the proton temperature ratio is strictly correlated with the normalized wave power, which indicates perpendicular heating of the protons whenever the power of the transverse fluctuations is enhanced."," As mentioned above, it turns out that the proton temperature ratio is strictly correlated with the normalized wave power, which indicates perpendicular heating of the protons whenever the power of the transverse fluctuations is enhanced."
+ An interesting behaviour of the temperature ratio of the alpha particlescan be inferred from Figure H]., An interesting behaviour of the temperature ratio of the alpha particlescan be inferred from Figure \ref{fig.4}.
+ This ratio increases with increasing normalized wave power as long as the normalized ion differential speed stays below about 0.5., This ratio increases with increasing normalized wave power as long as the normalized ion differential speed stays below about 0.5.
+" Beyond this value of Vy,/Va, the alpha-particle temperature ratio becomes roughly constant, until V,,/V4 exceeds a value of about 0.7, but then it decreases towards a value below unity when V,/V4 reaches one."," Beyond this value of $V_{\alpha p}/V_\mathrm{A}$, the alpha-particle temperature ratio becomes roughly constant, until $V_{\alpha
+p}/V_\mathrm{A}$ exceeds a value of about 0.7, but then it decreases towards a value below unity when $V_{\alpha p}/V_\mathrm{A}$ reaches one."
+" In the slow solar wind region, where the collisions are expected to be relatively high, the proton temperature anisotropy is ranging between 0.8 and 1, however, the alpha temperature anisotropy is not in the same range but higher, and varies between 0.9 and 1.2."," In the slow solar wind region, where the collisions are expected to be relatively high, the proton temperature anisotropy is ranging between 0.8 and 1, however, the alpha temperature anisotropy is not in the same range but higher, and varies between 0.9 and 1.2."
+" Most likely strong collisionality at low γαρ/γα tends to isotropize the alphas, whereas the weakening of the resonance at higher alpha/proton speeds leads to a decreasing of the alpha anisotropy, as in [Garyetal.(2005)."," Most likely strong collisionality at low $V_{\alpha p}/V_\mathrm{A}$ tends to isotropize the alphas, whereas the weakening of the resonance at higher alpha/proton speeds leads to a decreasing of the alpha anisotropy, as in \citet{gary2005}."
+ The monotonic increase of the proton temperature anisotropy with the normalized differential ion speed appears to be consistent with the findings of and with the ACE observations of (2005)))., The monotonic increase of the proton temperature anisotropy with the normalized differential ion speed appears to be consistent with the findings of \citet{kasper2008} and with the ACE observations of \citet{gary2005}) ).
+" However, there is clear difference between the trend of the alpha temperature anisotropy in Figure 4] and the results found by [Kasperetal.| (2008).."," However, there is clear difference between the trend of the alpha temperature anisotropy in Figure \ref{fig.4} and the results found by \citet{kasper2008}. ."
+" In their paper, the alpha temperature"," In their paper, the alpha temperature"
+redshifts z~6.,redshifts $\rm z\sim 6$.
+ If the Fell/Mell flux ratio is tracing the Fe/a abundance ratio. this result would imply that the tron abundance is nearly constant or even slightly increasing from redshift zero to 6.4. e. when the age of the universe was about 800 Myr.," If the FeII/MgII flux ratio is tracing the $\alpha$ abundance ratio, this result would imply that the iron abundance is nearly constant or even slightly increasing from redshift zero to 6.4, i.e. when the age of the universe was about 800 Myr."
+ While this finding is difficult to reconcile with the scenarios expecting a delay of 1 Gyr or more for the iron enrichment. models of rapid and strong star formation at high redshift offer a plausible explanation.," While this finding is difficult to reconcile with the scenarios expecting a delay of 1 Gyr or more for the iron enrichment, models of rapid and strong star formation at high redshift offer a plausible explanation."
+ Indeed. Matteucci&Recchi(2001) have shown that models of early formation of ellipticals. with very efficient star formation and lasting ~0.4 Gyr. imply a peak of SNla production (hence of iron enrichment) already at 0.3 Gyr after the onset of star formation.," Indeed, \citet{matrec01} have shown that models of early formation of ellipticals, with very efficient star formation and lasting $\sim$ 0.4 Gyr, imply a peak of SNIa production (hence of iron enrichment) already at 0.3 Gyr after the onset of star formation."
+ They have also shown that in the case of an instantaneous burst the SNIa production peaks at only 50 Myr after the burst., They have also shown that in the case of an instantaneous burst the SNIa production peaks at only 50 Myr after the burst.
+ As a consequence. our results can be explained with two possible scenarios: It is interesting to note that the first scenario (early formation of ellipticals) implies a redshift for the first major episode of star formation (z279) which is consistent with the redshift of the re-ionization (at z20c10) inferred from the recent Wilkinson Microwave Anisotropy Probe (WMAP) results (Bennettetal.2003).," As a consequence, our results can be explained with two possible scenarios: It is interesting to note that the first scenario (early formation of ellipticals) implies a redshift for the first major episode of star formation $\geq$ 9) which is consistent with the redshift of the re-ionization (at $\approx 20\pm 10$ ) inferred from the recent Wilkinson Microwave Anisotropy Probe (WMAP) results \citep{benet03}."
+. While the second scenario would be more consistent with the secondary re-ionization (at z~6) proposed by and Wyithe&Loeb(2003)., While the second scenario would be more consistent with the secondary re-ionization (at $\sim$ 6) proposed by \citet[][]{cen03} and \citet[][]{wylo03}.
+. We also find a population of quasars at z~4.3-5.1 with low iron emission (Table 2))., We also find a population of quasars at $\sim$ 4.3–5.1 with low iron emission (Table \ref{tab1}) ).
+ These objects are responsible for the apparently lower Fell bump in the average spectrum of the quasars in this redshift bin (Fig.2))., These objects are responsible for the apparently lower FeII bump in the average spectrum of the quasars in this redshift bin \ref{avers}) ).
+ Other quasars with such low Fe emission in this redshift range were also found by etal.(2003) and by Iwamuroetal.(2002) (see also Fioreetal.2003)).," Other quasars with such low Fe emission in this redshift range were also found by \citet{dieet03}
+ and by \citet{iwaet02} (see also \citealt{fiore03}) )."
+ Statistically we cannot make any strong statement on the fraction of quasars with low Fell., Statistically we cannot make any strong statement on the fraction of quasars with low FeII.
+ However. they could represent a population of quasars hosted in stellar systems which evolved more slowly than the quasars with higher Fe emission.," However, they could represent a population of quasars hosted in stellar systems which evolved more slowly than the quasars with higher Fe emission."
+ An alternative explanation for the FellI/Mgll trends discussed above is that the physical conditions of the Broad Line Region in high redshift quasars are significantly different than in quasars at lower redshift., An alternative explanation for the FeII/MgII trends discussed above is that the physical conditions of the Broad Line Region in high redshift quasars are significantly different than in quasars at lower redshift.
+ In this case the observed Fell/Mgll ratio at high redshift would reflect the combined effects of abundance variationsand variations of other physical parameters such as density. ionization. parameter. and microturbulence.," In this case the observed FeII/MgII ratio at high redshift would reflect the combined effects of abundance variations variations of other physical parameters such as density, ionization parameter, and microturbulence."
+ Disentangling these effects requires our observations to be complemented with the detection of the optical iron bump (Verneretal.2003).. which is shifted beyond the K-band at z-5.," Disentangling these effects requires our observations to be complemented with the detection of the optical iron bump \citep{veret03}, which is shifted beyond the K-band at $>$ 5."
+ Observations sensitive enough to detect the optical iron bump at these redshifts will probably have to await the James Webb Space Telescope (JWST)., Observations sensitive enough to detect the optical iron bump at these redshifts will probably have to await the James Webb Space Telescope (JWST).
+ We thank the referee M. Dietrich for helpful comments., We thank the referee M. Dietrich for helpful comments.
+ Part of this work was developed during the Guillermo Haro Workshop in 2003., Part of this work was developed during the Guillermo Haro Workshop in 2003.
+ RM acknowledges partial financial support from the Italian Space Ageney (ASI)., RM acknowledges partial financial support from the Italian Space Agency (ASI).
+ YJ and RM acknowledge support from CONACYT grant J32178-E., YJ and RM acknowledge support from CONACyT grant J32178-E.
+dark frames were used to linearise all [rames. the sky vackerouncd was subtractecl [rom the object. frames. and jixel-to-pixel variations were corrected.,"dark frames were used to linearise all frames, the sky background was subtracted from the object frames, and pixel-to-pixel variations were corrected."
+ Lhe chip distortion was corrected. in order to align the dispersion and. spatial directions along rows and columns of the chip: the sky vackerouncd was then subtracted) ancl spectra. extracted., The chip distortion was corrected in order to align the dispersion and spatial directions along rows and columns of the chip; the sky background was then subtracted and spectra extracted.
+ A low-order polynomial was fit to the argon lines. and hese calibrations applied to. the. object. spectra.," A low-order polynomial was fit to the argon lines, and these calibrations applied to the object spectra."
+ Flux calibration was achieved. by dividing the observed. spectra »w the spectrum. of a nearby. mid. C-tvpe star and then multiplving by a model for the absolute fux. distribution of the calibrator., Flux calibration was achieved by dividing the observed spectra by the spectrum of a nearby mid G-type star and then multiplying by a model for the absolute flux distribution of the calibrator.
+ Residual terrestrial atmospheric features were then corrected using an eariy-tvpe star., Residual terrestrial atmospheric features were then corrected using an early-type star.
+ The optical spectrum is shown in Fig. l.., The optical spectrum is shown in Fig. \ref{fig:spectrum}.
+ The spectrum clearly shows three narrow emission lines lla À65623. A6G6TS. ancl AT065 — as well as a broad component to the Ha Hine which is bluc-shiftec with respect to the narrow component.," The spectrum clearly shows three narrow emission lines — $\alpha$ $\lambda$ 6563, $\lambda$ 6678, and $\lambda$ 7065 — as well as a broad component to the $\alpha$ line which is blue-shifted with respect to the narrow component."
+ We fit gaussian models to these lines. using the package inIRAP.," We fit gaussian models to these lines, using the package in."
+ After normalising the spectrum. by a low-order volvnomial fit to the continuum. we fit a single. gaussian o each of the helium lines and two gaussians to the Ia line.," After normalising the spectrum by a low-order polynomial fit to the continuum, we fit a single gaussian to each of the helium lines and two gaussians to the $\alpha$ line."
+ “Phe velocities ancl widths of the three narrow lines are consistent with being the same: we therefore constrained hem to be equal., The velocities and widths of the three narrow lines are consistent with being the same; we therefore constrained them to be equal.
+ Ehe details of the fit are shown in Table 1.. The narrow lines show a velocity of ~|378kms.+ aand a width of ~9.5A. or 400kms+. while the broad component o the Ha Hine has a very dillerent velocity310kms +. or 688kinsI wwith respect to the narrow lines) and width (46A.. or 2200kms. +).," The details of the fit are shown in Table \ref{tab:optspec}.. The narrow lines show a velocity of $\sim +378\kms$ and a width of $\sim 9.5\;$, or $400\kms$, while the broad component to the $\alpha$ line has a very different velocity$-310\kms$ , or $-688\kms$ with respect to the narrow lines) and width $46\;$, or $2200\kms$ )."
+ No stellar features can be seen in the spectrum., No stellar features can be seen in the spectrum.
+ In xwtieular. no absorption features can be discerned. which might sugegeest the nature of the binary companion.," In particular, no absorption features can be discerned which might sugggest the nature of the binary companion."
+ Phere is no significant. variability apparent between the five spectra., There is no significant variability apparent between the five spectra.
+ Mignani. Caraveo Dignami obtained a spectrum of Cir N-1. using the LST near. periastron (phase 0.18) in June 1995.," Mignani, Caraveo Bignami obtained a spectrum of Cir X-1 using the HST near periastron (phase 0.18) in June 1995."
+ The line profile of the Lla Hine in their spectrum was very cillerent to our. 1997 spectrum., The line profile of the $\alpha$ line in their spectrum was very different to our 1997 spectrum.
+ They found the broad component of the gaussian to be centered at approximately its rest wavelength. while the narrow component was observed at a velocity of*.," They found the broad component of the gaussian to be centered at approximately its rest wavelength, while the narrow component was observed at a velocity of."
+. They interpreted the narrow component as arising in an accretion disc. with the 400kmsn vvelocity. οη the rotation. velocity. ο re disc material: the absence of the blue component they --ntorpretecd as a phase-cdependent shadowing elfect.," They interpreted the narrow component as arising in an accretion disc, with the $\sim 400\kms$ velocity representing the rotation velocity of the disc material; the absence of the blue component they interpreted as a phase-dependent shadowing effect."
+. Under us model. the velocity. of the narrow component shouk gaift with orbital phase. while the broad component remains ixecd.," Under this model, the velocity of the narrow component should shift with orbital phase, while the broad component remains fixed."
+ Our spectrum was taken at apastron. ancl does no garow the prediced shift of the narrow component.," Our spectrum was taken at apastron, and does not show the predicted shift of the narrow component."
+ In fact. je velocities we observe for the narrow component of the llo Hine (and 1e helium lines. not observed. in the LIST spectrum) are identical with the HIST spectrum. while thebroad component has shifted in velocity.," In fact, the velocities we observe for the narrow component of the $\alpha$ line (and the helium lines, not observed in the HST spectrum) are identical with the HST spectrum, while the component has shifted in velocity."
+ This would seen to indicate that the model proposed by Mignani et al., This would seem to indicate that the model proposed by Mignani et al.
+ for the svslem is wrong., for the system is wrong.
+ ‘To investigate this. we used data from the AAT archive. which contains nearly all observations taken with the AAT since its inception.," To investigate this, we used data from the AAT archive, which contains nearly all observations taken with the AAT since its inception."
+ Cir N-1 has been observed using the AAT several times by cüllerent observers. with most observations occurring about 20 vears ago.," Cir X-1 has been observed using the AAT several times by different observers, with most observations occurring about 20 years ago."
+ We obtained spectra at a total of five epochs which showed detectable emission: these are listed in Table 2. and shown in Fig. 2.., We obtained spectra at a total of five epochs which showed detectable emission: these are listed in Table \ref{tab:archival-obs} and shown in Fig. \ref{fig:archivalfit}.
+ Phe May. 1976 data were published by Whelan17: the rest of the data are unpublished., The May 1976 data were published by Whelan; the rest of the data are unpublished.
+ We obtained the original data from the AAT archive and reduced. then using standard. techniques ancl the ata reduction package., We obtained the original data from the AAT archive and reduced them using standard techniques and the data reduction package.
+ Note that for all these observations. lightight [rom stars[ 2 and 133 was presumablyunabl in the slit.," Note that for all these observations, light from stars 2 and 3 was presumably in the slit."
+" BBecausel of this. no attempt has been made to Dux. calibrate the μαvectra: instead. cach spectrum was binned to a spectral sampling of two data points per resolution clement. and )0 resulting spectrum, was normalised by a low-order polynomial fit to the spectrum."," Because of this, no attempt has been made to flux calibrate the spectra; instead, each spectrum was binned to a spectral sampling of two data points per resolution element, and the resulting spectrum was normalised by a low-order polynomial fit to the spectrum."
+ We performed. the same double gaussian fit to these μα»ectra as described in Table 1.. and the results are shown in ‘Table 2..," We performed the same double gaussian fit to these spectra as described in Table \ref{tab:optspec}, and the results are shown in Table \ref{tab:archival-obs}."
+ Several points are immediately. clear: despite re wide range of orbital phases. the broad. component is never observed. rechvard of the narrowcomponent'.," Several points are immediately clear: despite the wide range of orbital phases, the broad component is never observed redward of the narrow."
+". The (weighted) mean velocity of the narrow component is [350kms ο,that of the broad: component 100kis"," The (weighted) mean velocity of the narrow component is $+350\kms$ ,that of the broad component $-100\kms$ ."
+ The scatter of the velocities of the broad component is much, The scatter of the velocities of the broad component is much
+theories. such as those of Fabrycky Winn (2009) and Triaud et al. (,"theories, such as those of Fabrycky Winn (2009) and Triaud et al. ("
+2010). should consider only hot stars. because cool stars may have been affected by subsequent tidal evolution.,"2010), should consider only hot stars, because cool stars may have been affected by subsequent tidal evolution."
+ Finally. we interpret the results. as did Triaud et al. (," Finally, we interpret the results, as did Triaud et al. ("
+2010). as a blow against the theory of disk migration. which would yield low obliquities as a general rule.,"2010), as a blow against the theory of disk migration, which would yield low obliquities as a general rule."
+ Disk migration probably does play a role in sculpting exoplanetary orbits. and convergent migration of multiple planets may occasionally produce tilted orbits (Yu Tremame 2001).," Disk migration probably does play a role in sculpting exoplanetary orbits, and convergent migration of multiple planets may occasionally produce tilted orbits (Yu Tremaine 2001)."
+ But if obliquity truly depends on the present-day convective zone of the host star. then hot Jupiters likely arrived after the pre-main sequence convective phase ceased. tens of Myr after disk dispersal.," But if obliquity truly depends on the present-day convective zone of the host star, then hot Jupiters likely arrived after the pre-main sequence convective phase ceased, tens of Myr after disk dispersal."
+ Few-body gravitational dynamies (scattering or Kozai cycles) followed by tidal dissipation in the planet is compatible with this timescale. and it naturally produce misalignments. so this mechanism might account for most or all hot Jupiters.," Few-body gravitational dynamics (scattering or Kozai cycles) followed by tidal dissipation in the planet is compatible with this timescale, and it naturally produce misalignments, so this mechanism might account for most or all hot Jupiters."
+ We thank Jeremy Goodman. Matija Cuk. Bill Cochran. Sally Dodson-Robinson. Andrew Howard. Geoff Marcy. Tim Morton. Fred Rasio. Dimitar Sasselov. and Richard Wade for helpful discussions.," We thank Jeremy Goodman, Matija Cuk, Bill Cochran, Sally Dodson-Robinson, Andrew Howard, Geoff Marcy, Tim Morton, Fred Rasio, Dimitar Sasselov, and Richard Wade for helpful discussions."
+ geratefully acknowledges support from the NASA Origins program through award NNX09AD36G and the MIT Class of 1942. as well as the Tinsley Scholars program and the hospitality of the Astronomy Department at the University of Texas. Austin.," gratefully acknowledges support from the NASA Origins program through award NNX09AD36G and the MIT Class of 1942, as well as the Tinsley Scholars program and the hospitality of the Astronomy Department at the University of Texas, Austin."
+ DF acknowledges support from the Michelson Fellowship. supported by the National Aeronautics and Space Administration and administered by the Michelson Science Center.," DF acknowledges support from the Michelson Fellowship, supported by the National Aeronautics and Space Administration and administered by the Michelson Science Center."
+ SA acknowledges support from à NWO Rubicon, SA acknowledges support from a NWO Rubicon
+reference source J1755-2232. which was then subjected. to iterative imaging anc self-calibration.,"reference source J1755-2232, which was then subjected to iterative imaging and self-calibration."
+ The final image was used as a model for bxuxclpass calibration before transferring bandpass. amplitude and phase solutions to NPE J1752-223.," The final image was used as a model for bandpass calibration before transferring bandpass, amplitude and phase solutions to XTE J1752-223."
+ The strong scatering along the line of sight towards the source led us to use only the shortest baselines (up to ATALA(/5CLiz). wrere £ is the observing frequenev) in making the images. apering the weights of the visibility. cata with a Gaussian 'unction of full width at half maximum to further downweight the long baselines.," The strong scattering along the line of sight towards the source led us to use only the shortest baselines (up to $\lambda(\nu/5{\rm GHz})$, where $\nu$ is the observing frequency) in making the images, tapering the weights of the visibility data with a Gaussian function of full width at half maximum $\lambda$ to further downweight the long baselines."
+ We performed astrometry on à D5-second. απ exposure., We performed astrometry on a 5-second $i^{\prime}$ -band exposure.
+ We compared the positions of the stars against entries from the third U.S. Naval Observatory CCD Astrograph Catalog οςAC:Zachariasetal.9010)., We compared the positions of the stars against entries from the third U.S. Naval Observatory CCD Astrograph Catalog \citep[UCAC3;][]{Zac10}.
+.. An astrometric solution was computed by fitting for the reference. point. position. the scale and the position angle. considering all the sources that are not saturated. and appear stellar ancl unblencded.," An astrometric solution was computed by fitting for the reference point position, the scale and the position angle, considering all the sources that are not saturated and appear stellar and unblended."
+ We obtain a solution with rroot-mean-square (rms) residuals (rom 26 stars (UCAC [it model magnitudes 1416.5) well distributed on the ~4/. 4 iimage., We obtain a solution with root-mean-square (rms) residuals from 26 stars (UCAC fit model magnitudes 14–16.5) well distributed on the $\sim$ $\times$ image.
+" Phe positional accuracy of UCACS is estimated to be ~0.01"" oon stars of such magnitude (Zachariasetal.2010).", The positional accuracy of UCAC3 is estimated to be $\sim0.01$ on stars of such magnitude \citep{Zac10}.
+. In addition. the systematic uncertainty in tving the UCACS stars to the International Celestial Reference Svstem (ICT) is mmas (Zachariasetal.2010).," In addition, the systematic uncertainty in tying the UCAC3 stars to the International Celestial Reference System (ICRS) is mas \citep{Zac10}."
+. For the accuracy on our stellar positions we adopt the linear sum of the residuals of the astrometry and the accuracy of the catalogue. (as the latter is probably a systematic error): the resulting positional accuracy at La is oon both right ascension anc declination., For the accuracy on our stellar positions we adopt the linear sum of the residuals of the astrometry and the accuracy of the catalogue (as the latter is probably a systematic error): the resulting positional accuracy at $\sigma$ is on both right ascension and declination.
+ Thus our best optical position is This is mmas to the south east of the position of component A (corrected. for the error in the assumed calibrator position) as identified by Yangetal.(2010). in their VLBA observations of 2010 February 11., Thus our best optical position is This is mas to the south east of the position of component A (corrected for the error in the assumed calibrator position) as identified by \citet{Yan10} in their VLBA observations of 2010 February 11.
+ In the VLBA observations of 2010 April 25. a point-like radio source was detected with the VLBA. at a level of 0.46+0.07 at a position Consistent with the optical position of NPE J1752-223 (Fig. 2)).," In the VLBA observations of 2010 April 25, a point-like radio source was detected with the VLBA, at a level of $0.46\pm0.07$, at a position consistent with the optical position of XTE J1752-223 (Fig. \ref{fig:image}) )."
+ After correcting for the error in the assumed. calibrator position. the position of the radio source. relative to the VC'S-3 calibrator J1755-2232 (with an assumecl position of 22000) 175552(:22845.. 22°332'1 LO6616). was where the quoted. uncertainties are only the statistical errors from the Gaussian fitting algorithm (JMETE within AIvS).," After correcting for the error in the assumed calibrator position, the position of the radio source, relative to the VCS-3 calibrator J1755-2232 (with an assumed position of 2000) $^{\rm h}$ $^{\rm m}$ 2845, $-22$ 616), was where the quoted uncertainties are only the statistical errors from the Gaussian fitting algorithm (JMFIT within )."
+ The svstematic error clue to the cumulative effect. of 10 error in the assumed calibrator position and the hrow between calibrator and target5 (PracelCharlot&Lestrade2006). is estimated to be mmas., The systematic error due to the cumulative effect of the error in the assumed calibrator position and the throw between calibrator and target \citep*{Pra06} is estimated to be mas.
+ This VLBA position dillers from the optically-derivedl core position by mamas. along a position angle EI of N. With the derived uncertainty of mmas in both co-ordinates of the optical position. the radio source ds well within the optical error circle.," This VLBA position differs from the optically-derived core position by mas, along a position angle E of N. With the derived uncertainty of mas in both co-ordinates of the optical position, the radio source is well within the optical error circle."
+ Given the positional agreement. and since NTE 1752-223 was in the hard spectral state at the time of the radio observations. we conclude that this radio source is indeed a compact jet from he core of the svstem.," Given the positional agreement, and since XTE J1752-223 was in the hard spectral state at the time of the radio observations, we conclude that this radio source is indeed a compact jet from the core of the system."
+ We note that the ~ 6-month. time olfset. between the observations from which the optical and radio positions were determined. will likely lead to a small »ositional shift between epochs owing to the parallax and oper motion signatures of the source. of order. mumas if he source is at à distance of a few kpc and. participates in he Galactic rotation.," We note that the $\sim6$ -month time offset between the observations from which the optical and radio positions were determined will likely lead to a small positional shift between epochs owing to the parallax and proper motion signatures of the source, of order mas if the source is at a distance of a few kpc and participates in the Galactic rotation."
+ Furthermore. since we see the radio emission [rom optical depth Lat each frequency in à compact jet. the true core could be slightly olfset along the jet axis rom the position of the radio source.," Furthermore, since we see the radio emission from optical depth 1 at each frequency in a compact jet, the true core could be slightly offset along the jet axis from the position of the radio source."
+ On 2010 April 29. we detected a radio source at. this same position. but it had faded to 0.202E0.08 aat οι.," On 2010 April 29, we detected a radio source at this same position, but it had faded to $0.29\pm0.08$ at GHz."
+ On 2010 June 17. the source was only marginallv detected. at a level of 0.2540.08 aat GLIZ.," On 2010 June 17, the source was only marginally detected, at a level of $0.25\pm0.08$ at GHz."
+ Phe persistent. variable radio emission at this locationCi also supports our conclusion that this radio source does indeed correspond to a compact jet from the core of the system.," The persistent, variable radio emission at this location also supports our conclusion that this radio source does indeed correspond to a compact jet from the core of the system."
+ In none of these three hard state observations was, In none of these three hard state observations was
+ideal combination.,ideal combination.
+" We define as the mean standard deviation (MSD) the mean value of these five packages for each filter where k; denotes 5 different random data points of the light curve and Mean, is the mean value of the flux of each package.", We define as the mean standard deviation $(MSD)$ the mean value of these five packages for each filter where $k_j$ denotes 5 different random data points of the light curve and $Mean_{min}$ is the mean value of the flux of each package.
+" In general, each filter has a different MSD value, which is compared with the standard deviation of each filter TSD defined to be Finally, the relative standard deviation of each filter is computed and defined to be At the end of this process, we have three normalized light curves NFr;, NFg,;, and NFpj, and three values for the relative standard deviation RSDs, RSDG, and RSDg for each filter light curve, respectively."," In general, each filter has a different $MSD$ value, which is compared with the standard deviation of each filter $TSD$ defined to be Finally, the relative standard deviation of each filter is computed and defined to be At the end of this process, we have three normalized light curves $NF_{R,i}$, $NF_{G,i}$, and $NF_{B,i}$, and three values for the relative standard deviation $RSD_{R}$, $RSD_{G}$, and $RSD_{B}$ for each filter light curve, respectively."
+ The algorithm compares these three numbers and refers to the light curve with the minimum as the base and the light curve with the maximum as the target., The algorithm compares these three numbers and refers to the light curve with the minimum as the base and the light curve with the maximum as the target.
+" To help illustrate the procedure, we continue with an example."," To help illustrate the procedure, we continue with an example."
+ We assume that the base is the blue light curve (NFp) and the target is the red (VFr;) light curve.," We assume that the base is the blue light curve $(NF_{B,i})$ and the target is the red $(NF_{R,i})$ light curve."
+" Using both the base and the target, calculates a new mean light curve (AFj); in our example, computes and then refers to AF; as the light curve with the maximum (in this example, it defines AF; to be NFg;)."," Using both the base and the target, calculates a new mean light curve $(AF_{i})$; in our example, computes and then refers to $AF_{i}$ as the light curve with the maximum (in this example, it defines $AF_{i}$ to be $NF_{R,i}$ )."
+" According to assumptions 2 and 3, in the AF; light curve any possible"," According to assumptions 2 and 3, in the $AF_{i}$ light curve any possible"
+of203608.. with a reasonable amount of mixine. the early convective core can survive until the present age.,"of, with a reasonable amount of mixing, the early convective core can survive until the present age."
+ The agreement between observations aud models is then ereatly improved., The agreement between observations and models is then greatly improved.
+ Iu Sect. 3..," In Sect. \ref{sect_discussion},"
+ we explain why the convective core of ccan survive. even when the burning of ?Ile is no louger capable of sustaining it.," we explain why the convective core of can survive, even when the burning of $^3$ He is no longer capable of sustaining it."
+ For their modcling of203608... NOS. adopted. the ollowius stellar paramicters: Tig=6070+100A (Jouselletal.2005.. delPelosoetal. 2005)). WISL. Crom the Ilippare5 parallax 7;=107.98c imas. vanLeeuwen 2007)) and au observed surface uetallicity of |Z/X]g=450.1 dex. foin. detailed analvsis (soe MOS).," For their modeling of, M08 adopted the following stellar parameters: $T\ind{eff}=6070\pm100\,K$ \cite{2005A&A...440..321J}, \cite{2005A&A...441.1149D}) ), $L=1.39\pm0.13\,L_\odot$ (from the Hipparcos parallax $\pi=107.98\pm0.19$ mas, \cite{2007ASSL..250.....V}) ) and an observed surface metallicity of $[Z/X]\ind{f}=-0.5\pm0.1$ dex, from detailed analysis (see M08)."
+ They searched for an optimal model * adjusting the age of the star. its mass. the initial abundance of helium Y. ae the initial metallicity X]i o ft the following observational constraints: Tig. L. Z/X |g. Av. 14. aud yo. where Av represents a 1iean value of the Luge spacing. and avy. drys are the small spacings defined as Their best model fits the observational constraints on the star quite well. except for ὅτι. for which siguificaut disaerecmients remain (see Fig. 1)).," They searched for an optimal model by adjusting the age of the star, its mass, the initial abundance of helium $Y\ind{i}$ , and the initial metallicity $\ind{i}$ to fit the following observational constraints: $T\ind{eff}$, $L$, $[Z/X]\ind{f}$ , $\overline{\Delta\nu}$, $\zeroun$, and $\zerodeux$, where $\overline{\Delta\nu}$ represents a mean value of the large spacing, and $\zeroun$ , $\zerodeux$ are the small spacings defined as Their best model fits the observational constraints on the star quite well, except for $\zeroun$, for which significant disagreements remain (see Fig. \ref{fit_d01}) )."
+" As shown in Appendix AppendixAs. the simall spacing Avy, is very seusitive to the ceutral part of the stellar interior."," As shown in Appendix \ref{app_d01}, the small spacing $\zeroun$ is very sensitive to the central part of the stellar interior."
+ The present disagreement sugeests that the way we model these inuer regions should be recousidered., The present disagreement suggests that the way we model these inner regions should be reconsidered.
+ As already mentioned. the models computed im MOS lave no convective core.," As already mentioned, the models computed in M08 have no convective core."
+ We tried to obtain models of wwith a convective core here by generating extra musing atthe edge of the core. as sugeested by Boxbureh (1985)...," We tried to obtain models of with a convective core here by generating extra mixing atthe edge of the core, as suggested by \cite{1985SoPh..100...21R}. ."
+ We performed a new modeling of includius mixing ονομα the )ouudarv of the convective core by allowing do.20 im our models, We performed a new modeling of including mixing beyond the boundary of the convective core by allowing $d\ind{ov}>0$ in our models.
+ We computed stellar models using CESAM?KEK. (Morel 1997)). and we derived the mode frequencies from these models with the Liege oscillation code (LOSC. Scuflaireetal. 2008)).," We computed stellar models using CESAM2k \cite{1997A&AS..124..597M}) ), and we derived the mode frequencies from these models with the Liege oscillation code (LOSC, \cite{2008Ap&SS.316..149S}) )."
+ We used the same plysics as in MUS for our models; apart from the treatineu of convection.," We used the same physics as in M08 for our models, apart from the treatment of convection."
+ We preferred to use the more realistic ομαδα of Cauuto&Mazzitelli(1991).., We preferred to use the more realistic formalism of \cite{1991ApJ...370..295C}.
+ For calibration. we computed a solar model with this treatment of convection and found OC0.01.," For calibration, we computed a solar model with this treatment of convection and found $\alpha\ind{CM}=0.94$."
+ As in MUS. we adopted this value for our modelue.," As in M08, we adopted this value for our modeling."
+ Overshooting is described. as explained before. as au extension of the adiabatic and mixed region associated to the convective core. over a distance dc defined as where re is the radius of the convective core. {1 the pressure scale height. aud ogy the overshooting parameter.," Overshooting is described, as explained before, as an extension of the adiabatic and mixed region associated to the convective core, over a distance $d\ind{ov}$ defined as where $r\ind{c}$ is the radius of the convective core, $H_P$ the pressure scale height, and $\aov$ the overshooting parameter."
+ The models computed with overshooting showed that the star can still have a couvective core at its present age. provided the paraincter agy is lage enough.," The models computed with overshooting showed that the star can still have a convective core at its present age, provided the parameter $\aov$ is large enough."
+ We looked for an optimal model fitting the elobal paraiucters of203608: Tig. log(L/L..). aud [Z/X]g (eiven in Table 1)). as well as its seismic paraueters.," We looked for an optimal model fitting the global parameters of: $T\ind{eff}$, $\log(L/L_\odot)$, and $[Z/X]_{\hbox{\rm\small f}}$ (given in Table \ref{fit}) ), as well as its seismic parameters."
+ For the seismic constraints. we adopted the mode frequencies obtained in the analysis of MOS.," For the seismic constraints, we adopted the mode frequencies obtained in the analysis of M08."
+" We used a mean value of thelarge separation Av and the coefficieuts of a least-squares linear reeression of the small spacings dvy, and Avy.", We used a mean value of thelarge separation $\overline{\Delta\nu}$ and the coefficients of a least-squares linear regression of the small spacings $\zeroun$ and $\zerodeux$.
+ The frequency rangeof the identified modes is indeed s3inall chough to linearize them iu the form: with my the frequency of the imiaxiunuau of the signal (Myc2.7% illz), The frequency rangeof the identified modes is indeed small enough to linearize them in the form: with $\nu_0$ the frequency of the maximum of the signal $\nu_0\simeq 2.75$ mHz).
+ The observationalvalues of these cocficicnts are given inTable 1. , The observationalvalues of these coefficients are given inTable\ref{fit}. .
+We thus obtain a set of NO=8 observational constraiuts referred to as pU. i= LN.," We thus obtain a set of $N=8$ observational constraints referred to as $p_i^{\hbox{\rm\tiny obs}}$ , $i=1,N$ ."
+" The free pumuneters are the stellar lnass, age. initial helium: abundance. metallicity. audthe overshooting coefficient ayy (sec Table 1))."," The free parameters are the stellar mass, age, initial helium abundance, metallicity, andthe overshooting coefficient $\aov$ (see Table \ref{fit}) )."
+he secondary pair production mechanisni iu ordinary oulsars ?7..,the secondary pair production mechanism in ordinary pulsars \cite{Lyubarsky2008}.
+ For estimates of the emcereius radio flux. we adopt a oower-Lwwe radio spectrn with index a=2: F(f)xTu⋅ ," For estimates of the emerging radio flux, we adopt a power-law radio spectrum with index $\alpha =-2$: $F(f)\propto f^{-2}$ ."
+"our estimates we assume⋅ ZZ 10 τσMEIUsingoreso,1LOFA", In our estimates we assume $\dot{E}=10^{50\div 52}~\mathrm{erg~s^{-1}}$.
+RM Then: the expected radio: flux density: at the fiducial⋅− ower cut-off frequency LOO MITz will be F, Then the expected radio flux density at the fiducial lower cut-off frequency 100 MHz will be F )^2.
+esdgrE 30) wher Lay=£/10⋅↽↣ergs+, where $\dot{E}_{50}=\dot{E}/10^{50} \mathrm{erg s^{-1}}$.
+ The: duration. of. the radio: pulse will: be shorter than the time. between the mereine. aud the final collapse to BIL. i.c. a few teus ofmst.," The duration of the radio pulse will be shorter than the time between the merging and the final collapse to BH, i.e. a few tens of."
+ Tt is not excluded that the cliginatic nis radio trausicut repor “by ο actually originated from the binary NS coalescence eveut., It is not excluded that the enigmatic ms radio transient reported by \cite{Lorimer2007} actually originated from the binary NS coalescence event.
+ Can such a bright short radio transient eveut be detected by LOFAR?, Can such a bright short radio transient event be detected by LOFAR?
+ Bearing in ind the assuned power-law spectrum we are primarily iuterested iu tlie 120 MIIz band of the Wigh-Band Anteuna (IBA) of LOFAR., Bearing in mind the assumed power-law spectrum we are primarily interested in the 120 MHz band of the High-Band Antenna (HBA) of LOFAR.
+ Low-frequency radio observations are hinuperedl with effects of propagation in the interstellar and iutergalactie medi. especially due to strong dispersioul aud time scattering.," Low-frequency radio observations are hampered with effects of propagation in the interstellar and intergalactic medium, especially due to strong dispersion and time scattering."
+ The first effect cau be treated by niens of de-dispersion. at the same time providing conviueiue argument about extragalactic origin of the transient: moreover. only time delay due to the interstellar/iiterealactic dispersion malses the obscrvationg of the precursor GRB-alert possible.," The first effect can be treated by means of de-dispersion, at the same time providing convincing argument about extragalactic origin of the transient; moreover, only time delay due to the interstellar/intergalactic dispersion makes the observation of the precursor GRB-alert possible."
+" On the contrary, the interstellar/interealactic scattering las purely devastating effect outhe observations. irreversibly sinearing short pulses."," On the contrary, the interstellar/intergalactic scattering has purely devastating effect on the observations, irreversibly smearing short pulses."
+ This broadening very strongly depends on the frequency «. . » ≺⋔↴∖↴⊏∖↥⋅↖⇁⋮↕⊓↭∐↴∖↴≼∙↗≻↙−∖↙∖ Ft.," This broadening very strongly depends on the frequency of observations \citep{Bhatetal2004}
+ $\tau_{\mathrm sc} \propto f^{-4}$ ."
+ So a pulse with intrinsic width ~10ws and flux density (11)) will be broadened to me~102(D/1Cpe}? 4l and the resulting fux density will be reduced: F(120 MITZI) ο. parameters2. from. ?.. we. .obtain for fo⇁⊐ o>. ↑∐∖⋯∐∩⊾↿∐⋅⋪↕⊓∪∐↖↖⇁↕↑↕↽∶≻⋯↥⋅↸∖↖∣↥⋅↸∖⋯∪↑↸∖↴∖↴↑⋪↧⊓∪∐↴∖↴↑∐∖: thon13 - . 1 ‘ following sensitivity:m UEud. where Af is the bandwidth.," So a pulse with intrinsic width $\sim10~\mathrm{ms}$ and flux density \ref{flux_density_1}) ) will be broadened to $\tau_{\mathrm{sc}}\sim10^2(D/1~\mathrm{Gpc})^2$ s and the resulting flux density will be reduced: (120 ) )^4 Using LOFAR parameters from \cite{Nijboer2009}, we obtain for the configuration with 13 core + 7 remote stations the following sensitivity: = 40, where $\Delta f$ is the bandwidth."
+ The maximal distance D from which radiovate precursorSCR cau be detected bv LOFAR Ὁejves us thehe rate‘ of °‘{Bs with the ‘assumed tota‘total euerevU release:]‘ According to the DATSE catalog. the SGRD rate is about (0 per vear (?).. around 30 per cent of them are found at redshifts smaller than +=0.2 (D=1 Cpe) (302 We have used these figures to obtain the expected rate of LOFAR detections shown in Fig. 2..," The maximal distance $D$ from which radio precursor can be detected by LOFAR gives us the rate of SGRBs with the assumed total energy release: According to the BATSE catalog, the SGRB rate is about 170 per year \citep{Meegan1997}, around 30 per cent of them are found at redshifts smaller than $z=0.2$ $D=1$ Gpc) \citep{Guetta2006}; we have used these figures to obtain the expected rate of LOFAR detections shown in Fig. \ref{gamma_rate}."
+ A binary NS coalescence leads to the formation of a differentially rotating imassive object which can CVCually collapse DII possibly surrounded bv a magnetized torus., A binary NS coalescence leads to the formation of a differentially rotating massive object which can eventually collapse to BH possibly surrounded by a magnetized torus.
+ Numericalto simulations show (7) that the magnetic field strouely amplified at the pre-collapse stage np to LOis107 G. This configuration is favorable for the formation of a short gamuna-rayv burst., Numerical simulations show \citep{Duez2006} that the magnetic field is strongly amplified at the pre-collapse stage up to $10^{15}-10^{16}$ G. This configuration is favorable for the formation of a short gamma-ray burst.
+ At the ∖↴⋠⊾∖stage preceding⋅∖⊳∖⋅⊾ the ∖⊳⊀↴∖∢collapse aud GRD.∢≻≱⊀ a relativistic plasma outflow can be produced by the differentially rotating strongly maguctized configurationl," At the stage preceding the collapse and GRB, a relativistic plasma outflow can be produced by the differentially rotating strongly magnetized configuration."
+i Some fraction of the total power of this outflow can be couverted to electromagnetic radio enission., Some fraction of the total power of this outflow can be converted to electromagnetic radio emission.
+ The main uncertainty in the estimation of the radio uuinositv before the collapse of this configuration is the unknown efBcieney 5 of the conversion of the rotational cucrey losses from the differcutially rotating me-collapse object iuto radio cussion. which we assuned to depend onu the total energv loss rate as yxCE) in analogy with millisecoud radio pulsars.," The main uncertainty in the estimation of the radio luminosity before the collapse of this configuration is the unknown efficiency $\eta$ of the conversion of the rotational energy losses from the differentially rotating pre-collapse object into radio emission, which we assumed to depend on the total energy loss rate as $\eta\propto(\dot{E})^{\gamma}$ in analogy with millisecond radio pulsars."
+ We wave shownthat iu the most optimistic case the planned LOFAR scusitivity at120 MITz allows the potential detection at the signal-to-noise‘ ratio level 210 of short, We have shownthat in the most optimistic case the planned LOFAR sensitivity at120 MHz allows the potential detection at the signal-to-noise ratio level $>10$ of short
+cavity which presents a clear conical morphology with well defined edees aud apex ling ~3” from tlie I peak.,cavity which presents a clear conical morphology with well defined edges and apex lying $\simeq 3\arcsec$ from the K peak.
+ Due to the above mentioned reddening eradieut. the apex of the cone ects closer to the nucleus with increasing wavelength (compareτσ with R baud aud Pao /II» tages see below).," Due to the above mentioned reddening gradient, the apex of the cone gets closer to the nucleus with increasing wavelength (compare with R band and $\alpha$ $_2$ images – see below)."
+ The coutinmuu-subtracted Paschen o image (Fig., The continuum-subtracted Paschen $\alpha$ image (Fig.
+" 2aa) shows the presence of several strong cluission line knots along the ealactic plane. verv likely resulting from a cireunmuclear ring of star formation secu alinost edge-on. """," \ref{fig:line}a a) shows the presence of several strong emission line knots along the galactic plane, very likely resulting from a circumnuclear ring of star formation seen almost edge-on. """
+huot D of Moorwood et al. (19962)),"Knot B"" of Moorwood et al. \cite{moorwood96a}) )"
+ 18 clearly observed South East of the uncleus while νο C. North-West of he nucleus. is barelydetected.," is clearly observed South East of the nucleus while ""Knot C"", North-West of the nucleus, is barely detected."
+ Both knots are also marked ou the figure., Both knots are also marked on the figure.
+ Dust extinction strouely aects the muorphology making very diicult to trace the ring and ocate its center: a likely consequence is the appareut wisaligminent between the galaxy nucleus aud the ring ceuter., Dust extinction strongly affects the morphology making very difficult to trace the ring and locate its center; a likely consequence is the apparent misalignment between the galaxy nucleus and the ring center.
+ The observed ring of star formation ds similar o what has been found iu other starburst galaxies (cf., The observed ring of star formation is similar to what has been found in other starburst galaxies (cf.
+ AMoorwood 19965)j., Moorwood \cite{moorwood96b}) ).
+ The starburst ring could result from wo alternative scenarios: either the starburst origiuates at the imcleus. aud then propagates outward forming a ving in the galactic disk: or the rine corresponds to the xositiou of the imuecr Lindblad resonance where the gas deusitv is naturally increased by flow from both sides (see he review in Moorwood 199605)).," The starburst ring could result from two alternative scenarios: either the starburst originates at the nucleus, and then propagates outward forming a ring in the galactic disk; or the ring corresponds to the position of the inner Lindblad resonance where the gas density is naturally increased by flow from both sides (see the review in Moorwood \cite{moorwood96b}) )."
+ Panel b shows the contimunrsubtracted IT? inage which traces the edges of the wind-blown cavity., Panel b shows the continuum-subtracted $_2$ image which traces the edges of the wind-blown cavity.
+ As expected. the morphology is completely different from that of wwhich traces mainly starburst activity.," As expected, the morphology is completely different from that of which traces mainly starburst activity."
+ Note the strong II» cimission close to the nucleus at the apex of the cavity with au elougated. arc-like morphology.," Note the strong $_2$ emission close to the nucleus at the apex of the cavity with an elongated, arc-like morphology."
+" The ffiux in a 6""« aperture centered ou the I& baud peak ix aand corresponds tfo --70% of the total iutegrated enissiou in the NICMOS field of view.", The flux in a $6\arcsec\times 6\arcsec$ aperture centered on the K band peak is and corresponds to $\sim 70\%$ of the total integrated emission in the NICMOS field of view.
+ This is in good agreement with the 1.29-40.05 found by Isoormnect Ixrael (1996)) iu an equally sized aperture aud the inteerated ffrom the map by Moorwood Oliva (199 133)., This is in good agreement with the $\pm 0.05$ found by Koornneef Israel \cite{koorn96}) ) in an equally sized aperture and the integrated from the map by Moorwood Oliva \cite{moorwood94a}) ).
+ We roiiuk that contaimiuatiou of ecnissiou bv Hs unlikely since the line was detected neither by Ixooruneef (1993)) nor bv Moorwood Oliva (199 laj) aud froii their spectra we can sot au upper lait of tto the ratio., We remark that contamination of emission by is unlikely since the line was detected neither by Koornneef \cite{koorn}) ) nor by Moorwood Oliva \cite{moorwood94a}) ) and from their spectra we can set an upper limit of to the ratio.
+ Pane c in Figure 2 shows hat the equivalent width of iis up to in the star foriune regious. bu much lower in the winel-blown cavity.," Panel c in Figure \ref{fig:line} shows that the equivalent width of is up to in the star forming regions, but much lower in the wind-blown cavity."
+ Siuce near-IR coutimaiun cussion witlin the cone is not significantly hieher than in the surounudiug medi. the low equivalent width within the cone is duc to weaker eenission. the likely cousequence of low σας deusity.," Since near-IR continuum emission within the cone is not significantly higher than in the surrounding medium, the low equivalent width within the cone is due to weaker emission, the likely consequence of low gas density."
+ Panel din Figure 2 is the Pao /II» ratio image which also traces the wind-blown cavity., Panel d in Figure \ref{fig:line} is the $\alpha$ $_2$ ratio image which also traces the wind-blown cavity.
+ Note that the cone raced by Pao aud IL is offset with respect to the light cone observed ina BR: this is a result of the reddening eradicut iu the direction perpendicular to the ealactic lane., Note that the cone traced by $\alpha$ and $_2$ is offset with respect to the light cone observed in R: this is a result of the reddening gradient in the direction perpendicular to the galactic plane.
+ Fig., Fig.
+ 2ff is a true color ROB represeutatiou of ino and confimun images (Red-E222ML(ασοαι--».," \ref{fig:line}f f is a true color RGB representation of line and continuum images (Red=F222M, )."
+. DlIue-DPaa)). L/ aud N baud erouud based images are show iu Fie., $^\prime$ and N band ground based images are shown in Fig.
+ a aud b with contours overlaved ou the NICMOS K baud nage.," \ref{fig:midir}
+ a and b with contours overlayed on the NICMOS K band image."
+ No obvious poit source is detected at the location of the nucleus aud the extended emissiou is smooth and regular. elougated as the galactic disk.," No obvious point source is detected at the location of the nucleus and the extended emission is smooth and regular, elongated as the galactic disk."
+ The colmission shown in coutours m Fig., The emission shown in contours in Fig.
+ 2bb deviates from the Mage ina ΠΡΟ of interesting wavs., \ref{fig:line}b b deviates from the image in a number of interesting ways.
+ First. the northern edge of the cavity outlined most clearly in eendssion is also detected. although more fainth. iuu].. prestumably excited by the shocks resulting from the superwind.," First, the northern edge of the cavity outlined most clearly in emission is also detected, although more faintly, in, presumably excited by the shocks resulting from the superwind."
+ Otherwise. the eenission displavs two prouüneut peaks in the starburst region traced byPaa.. oue peak close to the uucleus and one offset at a position angle of about 250 degrees (counterclockwise frou North).," Otherwise, the emission displays two prominent peaks in the starburst region traced by, one peak close to the nucleus and one offset at a position angle of about 250 degrees (counterclockwise from North)."
+ In both of these reeious the ivafio is nimuch higher than in the rest of the starburst region., In both of these regions the ratio is much higher than in the rest of the starburst region.
+ The eenidssion likely originates iu radiativo Supernova remnants (SNBs)., The emission likely originates in radiative supernova remnants (SNRs).
+ Iu the dense nuclear region of NGCL915 the radiative phase of the SNRs will be short. and hence the cenussion will be immuch more strongly affected bv the stochastic nature of supernova explosions iu the starburst ring thanPao.," In the dense nuclear region of NGC4945 the radiative phase of the SNRs will be short, and hence the emission will be much more strongly affected by the stochastic nature of supernova explosions in the starburst ring than."
+.. The regious of high irafios thus simply trace recent supernova activity., The regions of high ratios thus simply trace recent supernova activity.
+ A lower limit aud a reasonable estimate of reddening cau be obtained from the IT[I color image in the case of, A lower limit and a reasonable estimate of reddening can be obtained from the H–K color image in the case of
+"Parity and angular momentum conservation exclude the S-wave A.\V-intermecliate state from the reaction pp—[ppl.s"".",Parity and angular momentum conservation exclude the S-wave $\Delta N$ -intermediate state from the reaction $pp\to \{pp\}_s\pi^0$.
+ In similar way. the MI (transition. dominating in the d—pn reaction al several hundred. MeV. via excitation of the A-isobar. is forbidden in the reaction pp—{pp}.o.," In similar way, the M1 transition, dominating in the $\gamma d\to pn$ reaction at several hundred MeV via excitation of the $\Delta$ -isobar, is forbidden in the reaction $pp\to \{pp\}_s\gamma$."
+ This suppression is similar to that in the reaction pd—{pp}.n d.6|]. as compared to the pd—dp process.," This suppression is similar to that in the reaction $pd\to \{pp\}_sn$ \cite{uz2002,jhuz2003} as compared to the $pd\to dp$ process."
+ Therefore. one could expect that some features of the short-range cdvnamics. which. perhaps. are not visible in the reactions with deuteron. pp—dx and 4d—pn. max reveal themselves in the corresponding reactions with the diproton.," Therefore, one could expect that some features of the short-range dynamics, which, perhaps, are not visible in the reactions with deuteron, $pp\to d\pi^+$ and $\gamma d \to pn$, may reveal themselves in the corresponding reactions with the diproton."
+" The OPE calculations. in agreement with the data show. however. that the A- is still significant in the reactions pp+(pp).s"" and pp—{pp}."," The OPE calculations, in agreement with the data show, however, that the $\Delta$ -contribuition is still significant in the reactions $pp\to \{pp\}_s\pi^0$ and $pp\to \{pp\}_s\gamma$."
+ HC would mean (hat short-range effect is rather minor itself in (hese reactions in the considered region., It would mean that short-range effect is rather minor itself in these reactions in the considered region.
+Inspection of the EEPIC images reveals an excess of X-ray emission north of18333-0832..PSR E the position of the middle-aged radio Pulsar BBI830-08 (spin-down age 7.~150 kyr: Taylor.Manchester&Lyne 1993).,"Inspection of the EPIC images reveals an excess of X-ray emission north of, encompassing the position of the middle-aged radio pulsar B1830–08 (spin-down age $\tau_c \sim 150$ kyr; \citealt*{tml93}) )."
+ BBIS830-08 is located outside. and possibly associated to. the asymmetric shell-type supernova remnant W41 (G23.3-0.3) despite an angular separation of tto the centre of the remnant (Kassim1992:Gaensler&John-ston 1995).," B1830–08 is located outside, and possibly associated to, the asymmetric shell-type supernova remnant W41 (G23.3–0.3) despite an angular separation of to the centre of the remnant \citep{kassim92, gaensler95}."
+. Its distance estimated through the dispersion measure (DM=411 pe em?) and the NE2001 model of the Galactic electron density eaten by Cordes&Lazio(2002). is 4.6 kpe., Its distance estimated through the dispersion measure $\rm DM=411$ pc $^{Ð3}$ ) and the NE2001 model of the Galactic electron density distribution by \citet{cordes02} is 4.6 kpc.
+ BBI830-08HESS might be also linked to the unidentitied source e34—087 lying inside W41 (Aharonianetal.e2006:owlAlbertetal.kr 2006).," B1830-08 might be also linked to the unidentified TeV source J1834–087 lying inside W41 \citep{aab06short,albert06short}."
+. In this case. the extended TeV source be a relie wind nebula (PWN) generated by electrons ejected by BBIS30-08 at its birth place and up-scattering the surrounding low-energy photons (e.g..Bartko&Bednarek 2008).," In this case, the extended TeV source would be a relic pulsar wind nebula (PWN) generated by electrons ejected by B1830–08 at its birth place and up-scattering the surrounding low-energy photons \citep[e.g.,][]{bartko08}."
+ However. different explanations have been suggested: the TeV emission of 155-057 could result from the aw’ decay of shock-accelerated hadrons interacting with a giant molecular cloud along the line of sight (Tianetal.2007). or from the inverse Compton radiation of electrons in a different PWN recently discovered near the centre of W41 (Mukherjeeetal. 2009: Misanovic.Kargaltsev&Pavlov 20113).," However, different explanations have been suggested: the TeV emission of J1834–087 could result from the $\pi^0$ decay of shock-accelerated hadrons interacting with a giant molecular cloud along the line of sight \citep{tian07} or from the inverse Compton radiation of electrons in a different PWN recently discovered near the centre of W41 \citealt{mukherjee09}; \citealt*{misanovic11}) )."
+ The detection of a large-scale X-ray nebula trailing the pulsar in the direction of the supernova remnant and of the TeV source would support the association between BBIS30-08 and 111834—087., The detection of a large-scale X-ray nebula trailing the pulsar in the direction of the supernova remnant and of the TeV source would support the association between B1830–08 and J1834–087.
+ We refined the imaging analysis by considering only EPIC-MOS data from the observations 0605852001. and 0605852101. whereas PSRBBI830-08 was located in the unread regions of the MOS detectors operating in Large Window mode during observation 0605851901.," We refined the imaging analysis by considering only EPIC-MOS data from the observations 0605852001 and 0605852101, whereas B1830–08 was located in the unread regions of the MOS detectors operating in Large Window mode during observation 0605851901."
+ The EPIC-pn data of the three observations were not used either. because of a higher level of residual soft proton contamination indicated by the ratio of the surface brightness inside the field of view to the one outside the field of view (DeLuca&Molendi2004).," The EPIC-pn data of the three observations were not used either, because of a higher level of residual soft proton contamination indicated by the ratio of the surface brightness inside the field of view to the one outside the field of view \citep{deluca04}."
+. We selected only 1- and 2-pixel events with the default flag mask., We selected only 1- and 2-pixel events with the default flag mask.
+ We built background maps using the Filter Wheel Closed (FWC) event lists provided by the bbackground working As the protile of the FWC maps changed since the launch of the mission. we filtered only the events collected after January 2007 to produce a more accurate template of the particle-induced background.," We built background maps using the Filter Wheel Closed (FWC) event lists provided by the background working As the profile of the FWC maps changed since the launch of the mission, we filtered only the events collected after January 2007 to produce a more accurate template of the particle-induced background."
+ The background maps were re-scaled in narrow energy bands according to the counts in the non-vignetted MOS corners., The background maps were re-scaled in narrow energy bands according to the counts in the non-vignetted MOS corners.
+ Finally. the images were mosaiced. background-subtracted. exposure-corrected. and adaptively smoothed to achieve a MANN ratio of 6.," Finally, the images were mosaiced, background-subtracted, exposure-corrected, and adaptively smoothed to achieve a signal-to-noise ratio of 6."
+ Diffuse emission around BBI830-08 is apparent in the 2-10 keV band mosaic. extending —2' nnorth-east and south of the pulsar (Fig. 6.," Diffuse emission around B1830–08 is apparent in the 2–10 keV band mosaic, extending $\sim$ north-east and south of the pulsar (Fig. \ref{B1830mosa},"
+ left panel)., left panel).
+ A closer inspection shows a compact structure at the pulsar position embedded in a halo of lower surface brightness (Fig. 6..," A closer inspection shows a compact structure at the pulsar position embedded in a halo of lower surface brightness (Fig. \ref{B1830mosa},"
+ right panel)., right panel).
+ The compact structure is 1' eelongated towards south-east. whereas the halo is mostly extending in the perpendicular direction.," The compact structure is $\sim$ elongated towards south-east, whereas the halo is mostly extending in the perpendicular direction."
+ The bulk of the halo emission lies closer to PSRBBIS30-08. while the the south-western tip is apparently connected with SGR 1833-0832: this is most likely due to the smoothing procedure.," The bulk of the halo emission lies closer to B1830–08, while the the south-western tip is apparently connected with SGR 1833–0832; this is most likely due to the smoothing procedure."
+ Both the halo and the compact structures are hardly visible in the soft band (0.5—2 keV). suggesting an absorbed non-thermal emission.," Both the halo and the compact structures are hardly visible in the soft band (0.5–2 keV), suggesting an absorbed non-thermal emission."
+ This is confirmed by the spectral analvsis., This is confirmed by the spectral analysis.
+ We produced the spectra also including the three EPIC-pn data sets., We produced the spectra also including the three EPIC-pn data sets.
+ The source counts were extracted from the mutually exclusive elliptical regions shown in Fig., The source counts were extracted from the mutually exclusive elliptical regions shown in Fig.
+ 6. (right panel)., \ref{B1830mosa} (right panel).
+ The background counts were estimated from apertures placed on the same CCD of each camera. and at the same distance to the readout node as the source for the EPIC-pn.," The background counts were estimated from apertures placed on the same CCD of each camera, and at the same distance to the readout node as the source for the EPIC-pn."
+ For each EPIC camera. we merged the individual spectra from the different observations into a single spectrum and combined the response matrices.," For each EPIC camera, we merged the individual spectra from the different observations into a single spectrum and combined the response matrices."
+ The compact and halo regions contained 560c30 and 2800=160 background-subtracted counts. respectively. corresponding to an average surface brightness of 0.2 and 0.085 counts 7.," The compact and halo regions contained $560 \pm 30$ and $2800 \pm 160$ background-subtracted counts, respectively, corresponding to an average surface brightness of 0.2 and 0.085 counts $^{-2}$."
+ The spectra Were rebinned with a minimum of 20 counts per spectral channel for the compact emission and 50 for the halo., The spectra were rebinned with a minimum of 20 counts per spectral channel for the compact emission and 50 for the halo.
+ The EPIC pn and MOS spectra of the compact structure (Fig. 7)), The EPIC pn and MOS spectra of the compact structure (Fig. \ref{B1830spec}) )
+ are well fit by an absorbed power-law model. yielding photon index D —L9 55. unabsorbed 2-10 KeV flux Piqug=(1.6a)101. and hydrogen-equivalent column density =4+2)107? ? (reduced v=L0 for 46 dof).," are well fit by an absorbed power-law model, yielding photon index $\Gamma = 1.9^{+0.7}_{-0.6}$ , unabsorbed 2–10 keV flux $F_{\mathrm{(2-10\,keV)}} = (1.6^{+0.4}_{-0.2})\times 10^{-13}$, and hydrogen-equivalent column density $ = (4 \pm 2) \times 10^{22}$ $^{-2}$ (reduced $\chi^{2} = 1.0$ for 46 dof)."
+ The halo spectra are also well fit by an absorbed power-law model with Po—IJ- unabsorbed 2-10 keV flux Pio.quis=(5.9HE10“ergemband =(2914)1077 em.> (reduced 47=1.1 for 153 dof).," The halo spectra are also well fit by an absorbed power-law model with $\Gamma = 1.7^{+0.5}_{-0.4}$ , unabsorbed 2–10 keV flux $F_{\mathrm{(2-10\,keV)}} = (5.9^{+0.7}_{-0.6}) \times 10^{-13}$, and $ = (2.9^{+1.3}_{-0.9})\times 10^{22}$ $^{-2}$ (reduced $\chi^{2} = 1.1$ for 153 dof)."
+ By fitting simultaneously the compact and halo spectra with the same column density. we found =(3.4yo)L1)107 7. slightly higher than the neutral hydrogen column density from radio observations M22107 ? (Dickey&Lockman1990).. which might indicate additional absorbing material along the line of sight.," By fitting simultaneously the compact and halo spectra with the same column density, we found $ = (3.4^{+1.1}_{-0.9})\times 10^{22}$ $^{-2}$, slightly higher than the neutral hydrogen column density from radio observations $\sim$$2 \times 10^{22}$ $^{-2}$ \citep{dickey90}, which might indicate additional absorbing material along the line of sight."
+ The hard diffuse emission around PSRBBIS30-08 is suggestive of a PWN., The hard diffuse emission around B1830–08 is suggestive of a PWN.
+ M a spin-down power of b=5S107 synchrotron4+. PSRBBIS3--—08 is energetic enough to produce a nebula in X-rays for a broad range of distances.," With a spin-down power of $\dot{E}= 5.8\times 10^{35}$ , B1830–08 is energetic enough to produce a synchrotron nebula detectable in X-rays for a broad range of distances."
+ For an assumed distance of 4.6 kpe. the measured fluxes imply a luminosity in the 2-10 keV energy range of —410 ffor the compact structure and of ~1.510 ffor the halo.," For an assumed distance of 4.6 kpc, the measured fluxes imply a luminosity in the 2–10 keV energy range of $\sim$$4 \times 10^{32}$ for the compact structure and of $\sim$$1.5 \times 10^{33}$ for the halo."
+ This corresponds to a conversion efficiency Lx//: of To10 land310/7. within the observed range of men systems (Possentieetal.2002:Kargaltsev&Pav," This corresponds to a conversion efficiency $L_{\mathrm{X}}/\dot{E}$ of $7 \times 10^{-4}$ and $3 \times 10^{-3}$, within the observed range of pulsar/PWN systems \citep{possenti02, kargaltsev08}."
+lov2008 The inner compaet structure. looks -- and its brightness peak is coincident with the pulsar position.," The inner compact structure looks one-sided, and its brightness peak is coincident with the pulsar position."
+ Such a morphology is reminiscent of a bow-shock PWN. in which the wind of a supersonically moving pulsar is contined by the oressure of the surrounding medium (seeGaensler&Slane2006ora review).," Such a morphology is reminiscent of a bow-shock PWN, in which the wind of a supersonically moving pulsar is confined by the ram-pressure of the surrounding medium \citep[see][for a review]{gaensler06}."
+ This phase occurs once the pulsar escapes ifs remnant. several ten thousand years after the supernova event. which is compatible with the age of BB1830-08.," This phase occurs once the pulsar escapes its remnant, several ten thousand years after the supernova event, which is compatible with the age of B1830–08."
+" Furthermore. he proper motion of PSRBBIS30-08 (3345 mas yr+. corresponding to a linear velocity of 730d, km 1: Hobbset 20043) lies in the high-velocity end of the 4;pulsar population and it is most likely supersonic."," Furthermore, the proper motion of B1830–08 $33\pm5$ mas $^{-1}$, corresponding to a linear velocity of $730 d_{4.6}$ km $^{-1}$; \citealt{hobbs04}) ) lies in the high-velocity end of the pulsar population and it is most likely supersonic."
+ The compact emission seems to be aligned with the pulsar propermotion (Fig. 6..," The compact emission seems to be aligned with the pulsar propermotion (Fig. \ref{B1830mosa},"
+ right). the brightness profile declining gradually from the pulsar position to south-east and sharply to north-west.," right), the brightness profile declining gradually from the pulsar position to south-east and sharply to north-west."
+ The location opposite to the directionof the pulsar motion. the hard extended emission. and the column density suggest that the halo is also linked to ΒΒΙ520-405.," The location opposite to the directionof the pulsar motion, the hard extended emission, and the column density suggest that the halo is also linked to B1830–08."
+ However. its morphology is," However, its morphology is"
+he inner slope parameter is not seusitive to the fit. as for he 1-D fitting. we confirmed that the resultant model xwanmeters for all three compoucuts do not change within statistical errors if 0.64<1.4.,"the inner slope parameter is not sensitive to the fit, as for the 1-D fitting, we confirmed that the resultant $\beta$ -model parameters for all three components do not change within statistical errors if $0.6
+\le \beta_{1} \le 1.5$."
+ We estimated the Iuninosities for the three compoucuts x dividing the total luminosity obtained from the spectral Bt according to the ratio of photou counts between he three., We estimated the luminosities for the three components by dividing the total luminosity obtained from the spectral fit according to the ratio of photon counts between the three.
+ Here we assunued isothermality for the three components., Here we assumed isothermality for the three components.
+ The results are also shown in Table 3.., The results are also shown in Table \ref{tab:2dfit}.
+" The uuninositv of the ""eluster compoucut is L5<1015 ere/s and dominates the total luminosity.", The luminosity of the “cluster” component is $4.5\times10^{44}$ erg/s and dominates the total luminosity.
+ We consider that tle cluission of G2 can be attributed to an elliptical galaxy ecause our clerived bhpuuiuositv is within the scatter of other ellipticals (Fabbiauoetal.1992:Matsushita2001) and there is a good positional coincidence with the member elliptical galaxy. £3362.," We consider that the emission of G2 can be attributed to an elliptical galaxy because our derived luminosity is within the scatter of other ellipticals \citep{FABBIANO_ETAL.1992, MATSUSHITA.2001} and there is a good positional coincidence with the member elliptical galaxy, 362."
+ On the other haud the luninosity of the Gl component is higher for a typical elliptical galaxy and the core radius is also comparable to that of other cluster with a all core radius (Ota&Mitzuda.2002)., On the other hand the luminosity of the G1 component is higher for a typical elliptical galaxy and the core radius is also comparable to that of other cluster with a small core radius \citep{OTA_ETAL.2002}.
+. We will discuss the properties of the Cl component iu more detail later., We will discuss the properties of the G1 component in more detail later.
+ Furthernore. we tested the significance of the ellipticitv of the cluster damage by substituting the third. component iu equation 1. for the elliptical ποσο].," Furthermore, we tested the significance of the ellipticity of the cluster image by substituting the third component in equation \ref{eq:threebeta} for the elliptical $\beta$ -model."
+ We found the ellipticity is not significant: the upper limit was 0.2.," We found the ellipticity is not significant; the upper limit was $\epsilon <
+0.2$ ."
+ Ou the other haud. Souciiletal.(2000). aud Dóliriueeretal.(2000) reported that the IRI image is sliehtlv clongated in a northeast-soutliwvest direction and can be fitted with a ellipse with an ellipticity of ~0.2.0.3.," On the other hand, \cite{Soucail_etal_2000} and \cite{BOHRINGER_ETAL.2000} reported that the HRI image is slightly elongated in a northeast-southwest direction and can be fitted with a ellipse with an ellipticity of $\sim
+0.2-0.3$."
+ Soucailetal.(2000) also mentioned that there is a significant twist of the position angle at about 100 kpe from the center., \cite{Soucail_etal_2000} also mentioned that there is a significant twist of the position angle at about 100 kpc from the center.
+ We consider that the significant shift of the center position of cluster cussion by SO kpe aud the preseuce of the secoud peak at 100 ος off of the strongest peak in theChandra nuage can account for the moderate ellipticitv and the twist derived by the elliptical model analysis with the WRI., We consider that the significant shift of the center position of cluster emission by 80 kpc and the presence of the second peak at 100 kpc off of the strongest peak in the image can account for the moderate ellipticity and the twist derived by the elliptical model analysis with the HRI.
+ Iu this section we will derive the projected cluster παν» profiles based ou the results of the X-ray analysis and estimate the cluster mass cuclosed within the arc radius., In this section we will derive the projected cluster mass profiles based on the results of the X-ray analysis and estimate the cluster mass enclosed within the arc radius.
+ We will directly compare these with the results from the lensing mass reconstruction., We will directly compare these with the results from the lensing mass reconstruction.
+ From the spectral and spatial analysis withChandra. we found that the cluster eas is consistent with beime isothermal and can be described with the spherical .i-nocdel after removing the local enissiou of Gl and G2.," From the spectral and spatial analysis with, we found that the cluster gas is consistent with being isothermal and can be described with the spherical $\beta$ -model after removing the local emission of G1 and G2."
+ Iu what follows. we regard the “cluster” component. /=3 in he three model fit. as diffuse enissiou comiue from the ICAL confined iu the cluster dark matter poteutial.," In what follows, we regard the “cluster” component, $i=3$ in the three $\beta$ -model fit, as diffuse emission coming from the ICM confined in the cluster dark matter potential."
+ Since he Gd core radius is much siualler than the arce radius (roaSE rue). While the cluster core radius 757raj. the Cl component is thought to be a minor contributor to he N-rav niass estimation.," Since the G1 core radius is much smaller than the arc radius $r_{c,1} \ll r_{\rm arc}$ ), while the cluster core radius $r_{c,3}\sim r_{\rm arc}$, the G1 component is thought to be a minor contributor to the X-ray mass estimation."
+ The effect of including the CU nass will be discussed later iu this section., The effect of including the G1 mass will be discussed later in this section.
+ Asunius the eas ds isothermal aud spherically distributed aud iu hvdrostatie equilibrimm. the density of uatter at a radius kr is estimated from the .anocel as The projected X-ray unass density profile is giveu by Iu Figure 7 we show the mass density profile calculated for model," Assuming the gas is isothermal and spherically distributed and in hydrostatic equilibrium, the density of matter at a radius $r$ is estimated from the $\beta$ -model as The projected X-ray mass density profile is given by In Figure \ref{fig:densityprof} we show the mass density profile calculated for $\beta$ -model."
+ The evlindzical cluster mass within a certain radius A is calculated by integrating the projected density profile. which can be written as follows (Otaetal.1998).," The cylindrical cluster mass within a certain radius $R$ is calculated by integrating the projected density profile, which can be written as follows \citep{Ota_etal_1998}."
+". Then from the results of 2D fitting for the cluster component we estimated the N-rayv mass within the arc radius of Raye=35""220T kpe to be Because the statistical errors of 2 and r. are coupled with oue another (Figure 6)) though AL is independently determined from the 4nodel parameters. we determined the error associated with the N-ray mass by evaluating it in the statistically allowed domain of the three dimensional paraueter space."," Then from the results of 2D fitting for the cluster component we estimated the X-ray mass within the arc radius of $R_{\rm 
+arc}=35\arcsec=220~h_{50}^{-1}$ kpc to be Because the statistical errors of $\beta$ and $r_{c}$ are coupled with one another (Figure \ref{fig:contour}) ) though $kT$ is independently determined from the $\beta$ -model parameters, we determined the error associated with the X-ray mass by evaluating it in the statistically allowed domain of the three dimensional parameter space."
+ Ou the other hand. adopting the are redshift of 1.675. the strong lensing mass is estimated to be youn equation (2) m Tysonetal.(1998).," On the other hand, adopting the arc redshift of $z_{\rm arc}=1.675$ , the strong lensing mass is estimated to be from equation (2) in \cite{Tyson_etal_1998}."
+. Broadlimrstetal.(2000) derived the stroug leusimg mass as Aat . which is aller wv about than equation 6..," \cite{Broadhurst_etal_2000} derived the strong lensing mass as $M_{\rm
+lens}(<200~h_{50}^{-1} {\rm kpc})=(2.22\pm0.06)\times10^{14}~
+h_{50}^{-1}{\rm M_{\sun}}$ , which is smaller by about than equation \ref{eq:lensmass_tyson}."
+" However. comparing Mx, ο Mis. a mass discrepancy of a factor of 3 is evident."," However, comparing $M_{\rm X,\beta}$ to $M_{\rm lens}$, a mass discrepancy of a factor of 3 is evident."
+" As shown in Figure 7.. since the surface mass density of the cluster compoucut obtained by the isothermal .- nodel is less than the critical surface mass deusitv. N44=GA/AxG)/CDI/DGD4,)=2.3<100M/pe? everywhere. he existence of the eravitationally multiple images im this cluster means that the mass distribution of tle cluster is ar frou the isothermal ιο aud/or the existence of an extra nass conrponent along the line of sight toward he cluster ceutral region is required."," As shown in Figure \ref{fig:densityprof}, , since the surface mass density of the cluster component obtained by the isothermal $\beta$ -model is less than the critical surface mass density, $\Sigma_{\rm crit}=(c^2/4\pi G)/(D_s/D_d D_{ds})=2.1\times10^3{\rm
+M_{\sun}/pc^2}$ everywhere, the existence of the gravitationally multiple images in this cluster means that the mass distribution of the cluster is far from the isothermal $\beta$ -model and/or the existence of an extra mass component along the line of sight toward the cluster central region is required."
+" The X-rav surface brightuess profile of the NEW. potential (Navarroetal.1996). is simular to that of the - model aud cau be couverted from the >-model paramcters through the relations of r,=022 and B=155 (Makinoetal.1998).", The X-ray surface brightness profile of the NFW potential \citep{Navarro_etal_1996} is similar to that of the $\beta$ -model and can be converted from the $\beta$ -model parameters through the relations of $r_s=r_c/0.22$ and $B=15\beta$ \citep{MAKINO_ETAL.1998}.
+". We thus derived the NEW density profile based on the results of ποσο] fitting. where pskTB/leGuir2,"," We thus derived the NFW density profile based on the results of $\beta$ -model fitting, where $\rho_s = kT B/4\pi G \bar{m} r_s^2$."
+ We then calenlated the projected mass density profile in the same ΠΙΑΤΟ as equation 3.., We then calculated the projected mass density profile in the same manner as equation \ref{eq:massdensity}.
+ The mass density profile is πιο steeper at the innermost region than that of the model and the central surface mass density can be higher than the critical surface mass density. however. those two models are not dcistinguishable at the are radius (Figure 7)).We estimated the N-raw mass to be This shows that the NEW model does not improve the enclosed mass within the arc radius.," The mass density profile is much steeper at the innermost region than that of the $\beta$ -model and the central surface mass density can be higher than the critical surface mass density, however, those two models are not distinguishable at the arc radius (Figure \ref{fig:densityprof}) ).We estimated the X-ray mass to be This shows that the NFW model does not improve the enclosed mass within the arc radius."
+ Therefore. adopting the NEW inodel can not be the main solution for the mass discrepancy problem of this cluster.," Therefore, adopting the NFW model can not be the main solution for the mass discrepancy problem of this cluster."
+equally spaced in semi-major axis. but rather define (heir spacing such (hat all wires have equal mass: (he wires must (herelore be more closely packed in regions of higher surlace densitv.,"equally spaced in semi-major axis, but rather define their spacing such that all wires have equal mass; the wires must therefore be more closely packed in regions of higher surface density."
+ This decomposition bv mass rather (han by semi-major axis is critical to our proof of stability., This de-composition by mass rather than by semi-major axis is critical to our proof of stability.
+ Denote the total ring mass by M. so (hat each wire has mass A/.N.," Denote the total ring mass by $M$, so that each wire has mass $M/N$."
+ Let wire j have semi-major axis e;. longitude of ascending node O;. and inclination J;.," Let wire $j$ have semi-major axis $a_j$, longitude of ascending node $\Omega_j$, and inclination $I_j$ ."
+" Deline mean elements bv à2σαΝΔ. 08320,/N. and 1=ΣΤΗΝ: let 7 denote the Keplerian mean molion at semi-major axis d."," Define mean elements by $\bar{a} \equiv \sum a_j/N$, $\bar{\Omega} \equiv \sum
+\Omega_j /N$, and $\bar{I} \equiv \sum I_j/N$; let $\bar{n}$ denote the Keplerian mean motion at semi-major axis $\bar{a}$."
+" As before. we assume that variations in orbital elements across (he ring are small: for any wire j. |a;—à|«&à. 4;—I|<I. and οn—QO)«1, "," As before, we assume that variations in orbital elements across the ring are small: for any wire $j$, $|a_j - \bar{a}| \ll \bar{a}$, $|I_j - \bar{I}| \ll \bar{I}$, and $|\Omega_j -
+\bar{\Omega}| \ll 1$ ."
+We adapt ((26a)) and (26b)) to write down the time rates of change of O; and J; due to the gravitational attraction of wire 4: Interparticle collisions are assimilated as follows., We adapt \ref{c_avgo}) ) and \ref{c_avgi}) ) to write down the time rates of change of $\Omega_j$ and $I_j$ due to the gravitational attraction of wire $k$: Interparticle collisions are assimilated as follows.
+ According to equation (2)). the collisional repulsion between (wo neiehboring wires is inversely proportional to the distance between then. and the resulting acceleration is directed oppositelv to that of inter-wire gravity.," According to equation \ref{collacc}) ), the collisional repulsion between two neighboring wires is inversely proportional to the distance between them, and the resulting acceleration is directed oppositely to that of inter-wire gravity."
+ This prescription implies Chat collisional lorees act effectively as anti-gravitational forces. having the same dependence on inter-wire separation as selb-eravitv. but with an opposite sign.," This prescription implies that collisional forces act effectively as anti-gravitational forces, having the same dependence on inter-wire separation as self-gravity but with an opposite sign."
+ Thus. we can account Lor collisions bv prependinga factor gj; to the right-hand sides of ((28a)) and (28b)).," Thus, we can account for collisions by prependinga factor $g_{jk}$ to the right-hand sides of \ref{c_ndo1}) ) and \ref{c_ndi1}) )."
+" The factor gj, represents the effects of collisions: the particular value of gj, depends on the local velocity dispersion: it is always <L. and it is negative where collisional repulsion is stronger than gravitational attraction."," The factor $g_{jk}$ represents the effects of collisions; the particular value of $g_{jk}$ depends on the local velocity dispersion; it is always $< 1$, and it is negative where collisional repulsion is stronger than gravitational attraction."
+ Finally. consider (he central planets «quadrupole field. which causesno secular change in {£;} but induces precessionof the nodes {Q;}.," Finally, consider the central planet's quadrupole field, which causesno secular change in $\{I_j\}$ but induces precessionof the nodes $\{\Omega_j\}$ ."
+ The effect is described to leading order by, The effect is described to leading order by
+Navarro et al (1996) have argued that CDM haloes have a universal density profile.,Navarro et al (1996) have argued that CDM haloes have a universal density profile.
+" In this section, we argue that the same holds true for voids, in qualitatively agreement with Van Kampen Van de Weygaert (1993)."," In this section, we argue that the same holds true for voids, in qualitatively agreement with Van Kampen Van de Weygaert (1993)."
+" Using our samples from the GIF, GIF2, and VLS simulations, we have computed the mass profiles of voids, using the actual particles in the voids instead of the smoothed density grid (we found the difference was important)."," Using our samples from the GIF, GIF2, and VLS simulations, we have computed the mass profiles of voids, using the actual particles in the voids instead of the smoothed density grid (we found the difference was important)."
+" Because of the different lower thresholds of the samples, we only compute void profiles for voids that have effective radii of 5h! MMpc or more."," Because of the different lower thresholds of the samples, we only compute void profiles for voids that have effective radii of $5\,h^{-1}$ Mpc or more."
+ Figure 11 shows the averaged enclosed density in z—0 voids as a function of radius., Figure \ref{CumDensityProfiles} shows the averaged enclosed density in $z=0$ voids as a function of radius.
+" For each void, we re-scaled the length scales by dividing by the effective radius, and we re-scaled densities by dividing by the enclosed density at the effective radius."," For each void, we re–scaled the length scales by dividing by the effective radius, and we re–scaled densities by dividing by the enclosed density at the effective radius."
+" We truncated the profiles at small radii, where numerical resolution effects begin to dominate (these will be discussed in more detail below)."," We truncated the profiles at small radii, where numerical resolution effects begin to dominate (these will be discussed in more detail below)."
+" For almost the entire range, the average density profiles of voids in the three simulation sets agree very well."," For almost the entire range, the average density profiles of voids in the three simulation sets agree very well."
+ We also computed density profiles for the Q—1 7CDM GIF simulation., We also computed density profiles for the $\Omega=1$ $\tau$ CDM GIF simulation.
+ These agree with the profiles of the ACDM simulations., These agree with the profiles of the $\Lambda$ CDM simulations.
+ This finding indicates that the form of void density profiles is indeed universal., This finding indicates that the form of void density profiles is indeed universal.
+" 'The different mass resolution of the three simulations affects the profiles in a systematic way: the higher the resolution, the lower the profile."," The different mass resolution of the three simulations affects the profiles in a systematic way: the higher the resolution, the lower the profile."
+ T'his effect is strongest at small radii., This effect is strongest at small radii.
+ The mass resolution also affects the centers of the voids., The mass resolution also affects the centers of the voids.
+" For example, in the VLS simulation, a single particle in a sphere of radius 1! MMpc corresponds to an overdensity of —0.8."," For example, in the VLS simulation, a single particle in a sphere of radius $1\,h^{-1}$ Mpc corresponds to an overdensity of $-0.8$."
+" Therefore, we cannot resolve the density profiles in the innermost regions."," Therefore, we cannot resolve the density profiles in the innermost regions."
+ If we plot the profiles all the way to the centers we find that the profiles all rise — individual particles contribute too much mass., If we plot the profiles all the way to the centers we find that the profiles all rise – individual particles contribute too much mass.
+" Therefore, we truncate the profiles in the void centers."," Therefore, we truncate the profiles in the void centers."
+ We cross-checked the effect of mass resolution by down-sampling the GIF2 simulation and producing void profiles., We cross–checked the effect of mass resolution by down–sampling the GIF2 simulation and producing void profiles.
+ The down-sampled simulation shows the trend visible in Figure 11.., The down–sampled simulation shows the trend visible in Figure \ref{CumDensityProfiles}.
+" Figure 11 is very encouraging: except for the effect of mass resolution, there are no systematic differences in the void samples."," Figure \ref{CumDensityProfiles} is very encouraging: except for the effect of mass resolution, there are no systematic differences in the void samples."
+" Figures 12,, shows the scaled enclosed density profiles of voids in the VLS simulation."," Figures \ref{CumDensityProfilesVLS}, shows the scaled enclosed density profiles of voids in the VLS simulation."
+ The curves lie fairly nicely on top of each other., The curves lie fairly nicely on top of each other.
+ The cumulative profile shown in Figure is quite well described by The fits only start to deviate somewhat beyond r/regg=1. (, The cumulative profile shown in Figure \ref{CumDensityProfilesVLS} is quite well described by The fits only start to deviate somewhat beyond $r/r_{eff} = 1$. (
+"Although the density run around halos is usually studied using the differential profile, void centers","Although the density run around halos is usually studied using the differential profile, void centers"
+radio jet.,radio jet.
+ Lt is also the smallest known Einstein ring with a diameter of only 335 milliaresee (mas). equal to. the separation between the two compact components.," It is also the smallest known Einstein ring with a diameter of only 335 milliarcsec (mas), equal to the separation between the two compact components."
+ Because of its small size. producing reliable maps of the brightness clistribution in the ring that have high resolution and hieh sensitivity is dillicult.," Because of its small size, producing reliable maps of the brightness distribution in the ring that have high resolution and high sensitivity is difficult."
+ The previous best map was made a a frequeney of 5 Gllz with the AHERLIN array (Patnaikeal. 1993).. but due to the incomplete aperture coverage of he observations no attempt was mace to interpret details of he Einstein ring in a lensing context.," The previous best map was made at a frequency of 5 GHz with the MERLIN array \cite{patnaik93}, but due to the incomplete aperture coverage of the observations no attempt was made to interpret details of the Einstein ring in a lensing context."
+ This paper presents a new map of 0915|357 mace from data obtained: with roth ALERLIN (in multi-frequency mode) and the VLA a requencies around 5 Gllz., This paper presents a new map of B0218+357 made from data obtained with both MERLIN (in multi-frequency mode) and the VLA at frequencies around 5 GHz.
+ The excellent aperture coverage when these data are combined has allowed us to produce he best map of this svstem at this frequency to date., The excellent aperture coverage when these data are combined has allowed us to produce the best map of this system at this frequency to date.
+ Work is currently underway to exploit the new map for. moclel constraints (Wucknitz et al..," Work is currently underway to exploit the new map for model constraints (Wucknitz et al.,"
+ in preparation), in preparation).
+ Observations were taken in June 1995 with the MEIRLIN array in multi-frequency synthesis (MES). mode. (see Conway. Cornwell Wilkinson (1990). or Sault Conway (1999) [or a full. discussion. of this. technique).," Observations were taken in June 1995 with the MERLIN array in multi-frequency synthesis (MFS) mode (see Conway, Cornwell Wilkinson \shortcite{conway90} or Sault Conway \shortcite{sault99} for a full discussion of this technique)."
+ The advantage of this technique is that for a given. baseline separation. changing the observing frequency. changes the spatial frequeney of the measured. visibility ancl hence the position ofthe visibility in the (i.0) plane.," The advantage of this technique is that for a given baseline separation, changing the observing frequency changes the spatial frequency of the measured visibility and hence the position of the visibility in the $(u,v)$ plane."
+ Vhis allows the (vor) plane to be filled in more fully than would be possible for observations taken at a single frequency with the same array and covering the same period. of time.," This allows the $(u,v)$ plane to be filled in more fully than would be possible for observations taken at a single frequency with the same array and covering the same period of time."
+" This in turn allows the sky brightness distribution to be reconstruct more accurately,", This in turn allows the sky brightness distribution to be reconstructed more accurately.
+ A complication arises from. the spectra index of the emission which. if uncorrected for. will produce imagine artelacts and significantly. degrade the dynamic range (the ratio of peak brightness to oll-source error) of a nap.," A complication arises from the spectral index of the emission which, if uncorrected for, will produce imaging artefacts and significantly degrade the dynamic range (the ratio of peak brightness to off-source error) of a map."
+ For these observations three frequencies around 5 Cllz separated by 320 MlIA were used., For these observations three frequencies around 5 GHz separated by 320 MHz were used.
+ These were 4546. 4866 ane 5186 MllIEZ. à frequencey range of ET per cent.," These were 4546, 4866 and 5186 MHz, a frequency range of $\pm$ 7 per cent."
+ Phe bandwidth of the individual frequencies was 16 MllIz., The bandwidth of the individual frequencies was 16 MHz.
+ DB0218|357 was observed alternately with the nearby JVAS calibrator source 30233[359 (Patnaiketal.1992) for phase-referencing purposes. a complete evele of these two sources through all requencics taking about half an hour before being repeated.," B0218+357 was observed alternately with the nearby JVAS calibrator source B0233+359 \cite{patnaik92} for phase-referencing purposes, a complete cycle of these two sources through all frequencies taking about half an hour before being repeated."
+ Observations were taken irregularly over the period. 1720 June totalling approximately 36 hours of data., Observations were taken irregularly over the period 17–20 June totalling approximately 36 hours of data.
+ Full Stokes »vameters were Observed so that. polarization maps could »e mace., Full Stokes parameters were observed so that polarization maps could be made.
+ Initial data reduction followed standard. procedures for ALERLIN data., Initial data reduction followed standard procedures for MERLIN data.
+ After the Hageing of corrupted. data. the lux clensity scale was set relative to ὃς 286.," After the flagging of corrupted data, the flux density scale was set relative to 3C 286."
+ Subsequent calibration was accomplished using the NRAO Astronomical Image Processing Software. (Avs)., Subsequent calibration was accomplished using the NRAO Astronomical Image Processing Software ).
+ Phe phase-calibrator was found to be substantially resolved and was therefore iteratively mapped ancl self-calibrated: several times to produce an accurate model of the source structure., The phase-calibrator was found to be substantially resolved and was therefore iteratively mapped and self-calibrated several times to produce an accurate model of the source structure.
+ This was then used to derive antenna-basecd complex gain solutions which were then applied to the 1209215|357 data., This was then used to derive antenna-based complex gain solutions which were then applied to the B0218+357 data.
+ 30233|359 was also used to derive instrumental polarization corrections and. 3€ 286 was used to calibrate the polarization position angle., B0233+359 was also used to derive instrumental polarization corrections and 3C 286 was used to calibrate the polarization position angle.
+ The data from cach frequency were next. individually mapped. ancl self-calibrated using the package (Shepherd1997).. solving for both phase and amplitucle antenna corrections.," The data from each frequency were next individually mapped and self-calibrated using the package \cite{shepherd97}, solving for both phase and amplitude antenna corrections."
+ A map mace from the 4546 ΜΗ data is shown in Fig. l., A map made from the 4546 MHz data is shown in Fig. \ref{mer1}.
+" Before the three frequencies. could x combined into a single data set. i was necessary to correct the data for the steep spectral index of the Einstein ring (a=0.6. where S,x £7)."," Before the three frequencies could be combined into a single data set, it was necessary to correct the data for the steep spectral index of the Einstein ring $\alpha = -0.6$, where $S_{\nu} \propto \nu^{\alpha}$ )."
+ Simply. sealing each pequencys data was not possible due to the compact cores wing a significantly Latter spectral index than the ring (a= 0.2)., Simply scaling each frequency's data was not possible due to the compact cores having a significantly flatter spectral index than the ring $\alpha = -0.2$ ).
+ This problem was solved. by subtracting the compact cores [rom the (roe) data.," This problem was solved by subtracting the compact cores from the $(u,v)$ data."
+ This was achieved. by (in AIPS) making a map of the calibrated cata output ron and subtracting the Fourier transform of he CLEAN components corresponding to the positions of components A and D. The cata were then remapped and the process repeated until all the core emission. had xen removed., This was achieved by (in ) making a map of the calibrated data output from and subtracting the Fourier transform of the CLEAN components corresponding to the positions of components A and B. The data were then remapped and the process repeated until all the core emission had been removed.
+" The total Dux density that remained in each subtracted map was then used to scale the (η,60) data sets o a consistent. amplitude scale."," The total flux density that remained in each subtracted map was then used to scale the $(u,v)$ data sets to a consistent amplitude scale."
+ This correction was as expected. the higher frequencies having to be scaled upwards in amplitude relative to the lower.," This correction was as expected, the higher frequencies having to be scaled upwards in amplitude relative to the lower."
+ Finally. the CLEAN components removed from the central frequency were added," Finally, the CLEAN components removed from the central frequency were added"
+At the time of a planet-lens event. the planet and lensecl star are nearly coincident.,"At the time of a planet-lens event, the planet and lensed star are nearly coincident."
+ After a time (hat can be as short as several vears. the planet may have moved Jar enough from the event location. that. it may be possible to detect it.," After a time that can be as short as several years, the planet may have moved far enough from the event location, that, it may be possible to detect it."
+ If the majority of free-floating planets have masses comparable to the Solar Systems terrestrial planets. it is unlikely that thev will be detected in the foreseeable Duture.," If the majority of free-floating planets have masses comparable to the Solar System's terrestrial planets, it is unlikely that they will be detected in the foreseeable future."
+ If they are more massive. however. direct detection. will be possible lor some of them.," If they are more massive, however, direct detection will be possible for some of them."
+ The feasibilitv of direct detection has been established by recent detections of wide-orbit planets., The feasibility of direct detection has been established by recent detections of wide-orbit planets.
+ For example. three planets. with masses between 5 and 13 times (he mass of Jupiter have been discovered orbiting the star IIR. 8799. which is 39.4 pe away (Alarois οἱ 2008).," For example, three planets with masses between $5$ and $13$ times the mass of Jupiter have been discovered orbiting the star HR 8799, which is 39.4 pc away (Marois et 2008)."
+ These planets have projected distances from the star ranging from 24 to GS AU., These planets have projected distances from the star ranging from $24$ to $68$ AU.
+ At these distances (the radiation they receive [rom the star contributes little to their luminosity. which is instead. dominated by cooling.," At these distances the radiation they receive from the star contributes little to their luminosity, which is instead dominated by cooling."
+ In addition. a planet with mass estimated to be no greater (han a few times that ol Jupiter has been discovered in a wide orbit (119 AU) around a star. Formalhaut. which is 7.7 pe away (INalas et 2008).," In addition, a planet with mass estimated to be no greater than a few times that of Jupiter has been discovered in a wide orbit (119 AU) around a star, Formalhaut, which is $7.7$ pc away (Kalas et 2008)."
+ Planets with radii comparable {ο or lareer than that of Jupiter. and temperatures of 1000 Ix. can be imaged if their distances [rom us are nol too great.," Planets with radii comparable to or larger than that of Jupiter, and temperatures of $\sim 1000$ K, can be imaged if their distances from us are not too great."
+ Imaging planets. even in these wide orbits. requires deep observations which can only be undertaken if there is a good chance of successfully [finding a planet or else placing meaningful limits.," Imaging planets, even in these wide orbits, requires deep observations which can only be undertaken if there is a good chance of successfully finding a planet or else placing meaningful limits."
+ For planets orbiting stars. such as those around IIR. 8799 or Formalhaut. (he star itself provides a well-defined region within which to conduct a search.," For planets orbiting stars, such as those around HR 8799 or Formalhaut, the star itself provides a well-defined region within which to conduct a search."
+ For free-floating planets. the occurrence of a lensing event provides a location about which to search.," For free-floating planets, the occurrence of a lensing event provides a location about which to search."
+ If 7j and 9j; are both measured. then w is known. hence (he angular separation ol the lens from the lensed source can be computed for all later times.," If $\tau_E$ and $\theta_E$ are both measured, then $\omega$ is known, hence the angular separation of the lens from the lensed source can be computed for all later times."
+ Observers can then plan observations to occur several vears alter (hie event. in order (ο image (he lens.," Observers can then plan observations to occur several years after the event, in order to image the lens."
+ This procedure will image a subset of nearby free-floating planets., This procedure will image a subset of nearby free-floating planets.
+" If enough photons are collected to estimate the planets temperature. then (he range of possible radii will provide a range of possible luminosities. allowing the distance D, (ad (he uncertaintv in ils value) to be estimated."," If enough photons are collected to estimate the planet's temperature, then the range of possible radii will provide a range of possible luminosities, allowing the distance $D_L$ (and the uncertainty in its value) to be estimated."
+ In addition. a sequence of images may be able to establish the geometric parallax.," In addition, a sequence of images may be able to establish the geometric parallax."
+ Armed with an estimate of Dj. the lens mass can be determined in those cases in which 85 has been measured through finite-source-size To distinguish between models in which the lens is a planet bound (to a star or else is a [ree-[loating planet. we search for evidence of a star that could be associated with the lens.," Armed with an estimate of $D_L,$ the lens mass can be determined in those cases in which $\theta_E$ has been measured through finite-source-size To distinguish between models in which the lens is a planet bound to a star or else is a free-floating planet, we search for evidence of a star that could be associated with the lens."
+ Even il a star is not detected. however. it is possible that the event was caused by a high-velocitv stellax-miass object rather than by a Iree-Hoating planet.," Even if a star is not detected, however, it is possible that the event was caused by a high-velocity stellar-mass object rather than by a free-floating planet."
+" Consider a short event (sav Tp=1 day). and a measured small Einstein angle (sav 8j=1.56x 10.7"")."," Consider a short event (say $\tau_E=1$ day), and a measured small Einstein angle (say ${\theta_E=1.56\times 10^{-5}}^{\prime\prime}$ )."
+ If (he lens is not a planet. how fast must the lens be moving and how far away must it be?," If the lens is not a planet, how fast must the lens be moving and how far away must it be?"
+ For, For
+that the extrapolation does not go lower than the primordial cooling rate.,that the extrapolation does not go lower than the primordial cooling rate.
+ Phe primordial abundance pattern is. used where ο/η] ] (Wheelerctal.1989:Bessellοἱ1991). and the solar abundance pattern is used where οΗ] x0 CXnders&Grevesse1980).," The primordial abundance pattern is used where [Fe/H] $\le -1$ \citep{Wheeler.etal:89,Bessell.etal:91}, and the solar abundance pattern is used where [Fe/H] $\ge 0$ \citep{Anders.Grevesse:89}."
+. For 1« Fel] «0. the abundance pattern is computed. by interpolating between the primordial and. solar abundance patterns.," For $-1 <$ [Fe/H] $ < 0$, the abundance pattern is computed by interpolating between the primordial and solar abundance patterns."
+"Note that. οΗ] - log(nposnp - login,""np.) where n is the number clensitv.","Note that, [Fe/H] = $\log (n_{\textup{Fe}}/n_{\textup{H}})$ - $\log (n_{\textup{Fe}_{\sun}}/n_{\textup{H}_{\sun}})$, where $n$ is the number density."
+ Figure 1. shows the relationships between the metal parameters ο). X. Y. and 12|log(O/11)) and the total metallicity Z.," Figure \ref{fig:XYZ_table} shows the relationships between the metal parameters ([Fe/H], $X$, $Y$, and $12 + \log(O/H)$ ) and the total metallicity $Z$."
+ Because the abundance pattern is fixed for à given Z. Fe/L] is a monotonically increasing function of Z.," Because the abundance pattern is fixed for a given $Z$, [Fe/H] is a monotonically increasing function of $Z$."
+ The main dilference between the solar and the primordial abundance pattern is the oxygen or a-particle enhancement., The main difference between the solar and the primordial abundance pattern is the oxygen or $\alpha$ -particle enhancement.
+ In the early phase of metal enrichment. Tvpe LL SNe are the main source of metals and they tend to generate more a-particles than SN la. Therefore. it is expected. that a low metallicity gas to have more a-particles than the high metallicity one.," In the early phase of metal enrichment, Type II SNe are the main source of metals and they tend to generate more $\alpha$ -particles than SN Ia. Therefore, it is expected that a low metallicity gas to have more $\alpha$ -particles than the high metallicity one."
+ As mentioned above. we use two cilferent chemical abundance pattern (at low- and high-end. of the Z values) to model this evolution in chemical composition.," As mentioned above, we use two different chemical abundance pattern (at low- and high-end of the $Z$ values) to model this evolution in chemical composition."
+ We compute the oxvgen abundance based on the Table 4 of 8D93., We compute the oxygen abundance based on the Table 4 of SD93.
+ In our code. we incorporate the metal cooling elect by adding the additional cooling contribution [rom metals on top of the primordial cooling rate. which is computed in a similar fashion as described in Ixatzetal.(1996).," In our code, we incorporate the metal cooling effect by adding the additional cooling contribution from metals on top of the primordial cooling rate, which is computed in a similar fashion as described in \citet{Katz.etal:96}."
+. Metal enrichment changes the mean molecular weight peo which is needed to compute the temperature and internal enereyv of eas.," Metal enrichment changes the mean molecular weight $\mu$, which is needed to compute the temperature and internal energy of gas."
+ In order to compute fr we need to know the electron number density. a... which varies with metallicity.," In order to compute $\mu$, we need to know the electron number density $n_e$, which varies with metallicity."
+ We obtain the values of 5. from SD93 in the same wav as we obtained the metal cooling rates., We obtain the values of $n_e$ from SD93 in the same way as we obtained the metal cooling rates.
+ Our electron number density estimate ignores the electrons. from. photoionised metals bv the UV background radiation. but we expect that the number of electrons [rom ionised metals will be small.," Our electron number density estimate ignores the electrons from photoionised metals by the UV background radiation, but we expect that the number of electrons from ionised metals will be small."
+" Varving the mean molecular weight so also inlluences the SE threshold censity py, in the multiphase LSAT model of Springel&Llernquist(2003a) eas well as the cold and hot gas fractions of multiphase gas particles."," Varying the mean molecular weight $\mu$ also influences the SF threshold density $\rhoth$ in the multiphase ISM model of \citet{Springel.Hernquist:03}, as well as the cold and hot gas fractions of multiphase gas particles."
+ In this paper. we consider two dillerent treatments of pi.," In this paper, we consider two different treatments of $\rhoth$."
+" In the ""constant pi nmiocel. we take the value of pij as in the original formulation by Springel&Llernquist(2003a) and do not modulate pin bv the varving f."," In the `constant $\rhoth$ ' model, we take the value of $\rhoth$ as in the original formulation by \citet{Springel.Hernquist:03} and do not modulate $\rhoth$ by the varying $\mu$."
+" In the ""varying pi model. we take into account of the modulation of pi, by the change in i."," In the `varying $\rhoth$ ' model, we take into account of the modulation of $\rhoth$ by the change in $\mu$."
+ We use three series of simulations with varving box size and resolution., We use three series of simulations with varying box size and resolution.
+ The elfects of metal enrichment. and metal cooling are tested by comparing the runs listed. in ‘Table 1.., The effects of metal enrichment and metal cooling are tested by comparing the runs listed in Table \ref{tbl:Simulation}.
+ The ellects of galactic wind are also tested by turning on and olf the wind., The effects of galactic wind are also tested by turning on and off the wind.
+ The adopted: cosmological parameters of all simulations are (ην04.Om.η).=(0.3.0.7 .0.04.0.85.0.7).. which mosth concur with the results of the Wilkinson Alicrowave Anisotropy Probe (WALAPL) (Spergeletal. 2003).," The adopted cosmological parameters of all simulations are $(\Omega_{m}\, , \Omega_{\Lambda}\, , \Omega_{b}\, , \sigma_{8}\, , h) \: = \: (0.3\, ,0.7\, ,0.04\, ,0.85\, ,0.7)$ , which mostly concur with the results of the Wilkinson Microwave Anisotropy Probe (WMAP1) \citep{WMAP1}. ."
+. Our naming conventions for the simulations are as follows., Our naming conventions for the simulations are as follows.
+ “Phe first part of the run name denotes the particle number and box size used for the simulation series., The first part of the run name denotes the particle number and box size used for the simulation series.
+" The run with the extension of ""me (for ""metal cooling’) adopts the ‘constant pin model for metal cooling. the run with the extension ""ms (for ""metal. νανήσο} adopts the ‘varying pi model. the run with the extension ""nw has no galactic wind. and the run without any extension uses the cooling rates for the primordial chemical composition (i.c. only LE and Le)."," The run with the extension of `mc' (for `metal cooling') adopts the `constant $\rhoth$ ' model for metal cooling, the run with the extension `mv' (for `metal, varying') adopts the `varying $\rhoth$ ' model, the run with the extension `nw' has no galactic wind, and the run without any extension uses the cooling rates for the primordial chemical composition (i.e., only H and He)."
+" The ""N216L10 series is the fiducial simulation set in our study.", The `N216L10' series is the fiducial simulation set in our study.
+ We implement one simulation without galactic wind CN216LIO0nw run)., We implement one simulation without galactic wind (`N216L10nw' run).
+ Phe ;N216L10' run has the same physical models as the O4. run used in Springel&Hernquist(2003b) and Nagamineetal.(2004) with no metal cooling., The `N216L10' run has the same physical models as the `Q4' run used in \citet{Springel.Hernquist:03:S} and \citet{Nagamine.etal:04} with no metal cooling.
+" To examine the resolution effect. we implement lower resolution simulations. the ""NIJA4LIO. series."," To examine the resolution effect, we implement lower resolution simulations, the `N144L10' series."
+ Owing to the comparably small box size and missing long wavelength perturbations. we evolve the NIJ4LIO0 ancl N216LIO0 series only down to z=2.75.," Owing to the comparably small box size and missing long wavelength perturbations, we evolve the N144L10 and N216L10 series only down to $z=2.75$."
+ To evolve the simulation toward lower redshifts. we need a simulation with a larger box size.," To evolve the simulation toward lower redshifts, we need a simulation with a larger box size."
+" Therefore we implement the ""IN2SSLS4. series. which has a larger box size than other series. but à lower resolution."," Therefore we implement the `N288L34' series, which has a larger box size than other series, but a lower resolution."
+ This series is evolved. down to z=1., This series is evolved down to $z=1$.
+ The runs within the same series use the identical initial condition. which makes the comparison more robust and enables halo-by-halo comparison if needed.," The runs within the same series use the identical initial condition, which makes the comparison more robust and enables halo-by-halo comparison if needed."
+ 1n cosmological simulations. dark matter. gas. and stars are represented by particles. and these particles are Monte Carlo representations of matter distribution.," In cosmological simulations, dark matter, gas, and stars are represented by particles, and these particles are Monte Carlo representations of matter distribution."
+ In. order to study ealaxy formation using a cosmological simulation. we need a working definition for simulated: galaxies.," In order to study galaxy formation using a cosmological simulation, we need a working definition for simulated galaxies."
+ In this paper the simulated. galaxies are defined as isolated groups of star and gas particles. identified by a simplified: variant of the algorithm developed by Springelοἱal.(2001).," In this paper the simulated galaxies are defined as isolated groups of star and gas particles, identified by a simplified variant of the algorithm developed by \citet{Springel.etal:01}."
+ In more detail. the code first computes a smoothed barvonic density field to identify candidate galaxies with high density peaks.," In more detail, the code first computes a smoothed baryonic density field to identify candidate galaxies with high density peaks."
+ The full extent of these groups are found by acdcding gas and star particles to the groups in the order of declining density., The full extent of these groups are found by adding gas and star particles to the groups in the order of declining density.
+ Wall Nau! nearest neighbour particles have lower densities. this group of particles is considered. as a new eroup.," If all $\Nmin$ nearest neighbour particles have lower densities, this group of particles is considered as a new group."
+ If there is a denser neighbour. the particle is attached o the group to which its nearest denser neighbour already xdongs to.," If there is a denser neighbour, the particle is attached to the group to which its nearest denser neighbour already belongs to."
+ Lf two nearest neighbours belong to cilferent groups and one of them has less than μμ particles. these he two groups are merged.," If two nearest neighbours belong to different groups and one of them has less than $\Nmin$ particles, these the two groups are merged."
+ LE two nearest neighbours belong o dillerent. groups and both of them has more than Ius xwiicles. the particles are. attached to the larger eroup. caving the other group intact.," If two nearest neighbours belong to different groups and both of them has more than $\Nmin$ particles, the particles are attached to the larger group, leaving the other group intact."
+" In addition. the gas particles in groups should be denser than 0.01pi,. In this paper we are not concerned. with groups without star. particles."," In addition, the gas particles in groups should be denser than $0.01 \rhoth$ In this paper we are not concerned with groups without star particles."
+ The raction of such groups are very small. therefore they do not allectour conclusions.," The fraction of such groups are very small, therefore they do not affectour conclusions."
+ Figure 2. shows the elfect. of metal cooling ancl galactic wind on the cosmic SE history for the N216L10 series., Figure \ref{fig:sfr_N216L10} shows the effect of metal cooling and galactic wind on the cosmic SF history for the N216L10 series.
+ The, The
+the true orbital period to be 17.11 hours.,the true orbital period to be 17.11 hours.
+ Using the most up-to-date ephemeris for AC211I (Ioannou et al., Using the most up-to-date ephemeris for AC211 (Ioannou et al.
+ 2001). we have re-examüned the original spectroscopy obtained by DAC and TN in three sets of observing runs: August. 1986 (published in Navlor ct al.," 2001), we have re-examined the original spectroscopy obtained by PAC and TN in three sets of observing runs: August 1986 (published in Naylor et al."
+ 1988). August LOST (published in Navlor Charles 1989). and August 1988 (unpublished).," 1988), August 1987 (published in Naylor Charles 1989), and August 1988 (unpublished)."
+ A summary of the observations is given in Table 1., A summary of the observations is given in Table 1.
+ The data were obtained with the RGO Intermediate Dispersion Speetrograph and LPCS (1986 1987) or CCD (1988) on the Isaac Newton Telescope at the Roque de los Muchachos Observatory on La Palma., The data were obtained with the RGO Intermediate Dispersion Spectrograph and IPCS (1986 1987) or CCD (1988) on the Isaac Newton Telescope at the Roque de los Muchachos Observatory on La Palma.
+ Each dataset consists of three or four nights of 600-s exposures., Each dataset consists of three or four nights of 600-s exposures.
+ The mean spectrum for each vear is plotted in Fig., The mean spectrum for each year is plotted in \ref{fig:3specs}.
+ A detailed description of how the spectra were extracted 3..from the LPCS data can be found in Navlor et al. (, A detailed description of how the spectra were extracted from the IPCS data can be found in Naylor et al. (
+1988).,1988).
+ ΑΟΣΤ]. located. in the core of MI5. is. situated in. an exceptionally croweled field. which makes observing it a considerable technical challenge. and requires superb secing.," AC211, located in the core of M15, is situated in an exceptionally crowded field, which makes observing it a considerable technical challenge, and requires superb seeing."
+ CGround-based. spectra of AC21I1 are heavily contaminated with light from the cluster core., Ground-based spectra of AC211 are heavily contaminated with light from the cluster core.
+ Variations in secing between and during exposures leads to changes in the relative contribution of the eluster background in AC21II's spectra. which makes it. impossible to determine AC21I1's continuum Hux and hence obtain reliable equivalent widths. [or its absorption and emission lines.," Variations in seeing between and during exposures leads to changes in the relative contribution of the cluster background in AC211's spectra, which makes it impossible to determine AC211's continuum flux and hence obtain reliable equivalent widths for its absorption and emission lines."
+ The cluster. contamination is clearly visible in. the Balmer lines., The cluster contamination is clearly visible in the Balmer lines.
+ In Fie.4 we have for clarity merged the 11.1. 1I5 ancl L9 lines. to reduce visual errors in order to highlight orbital-phase-dependent changes in the Balmer line profiles.," In \ref{fig:stack} we have for clarity merged the ${\beta}$, ${\gamma}$ and ${\delta}$ lines, to reduce visual errors in order to highlight orbital-phase-dependent changes in the Balmer line profiles."
+ As noted. by Llovaisky (1989). the Balmer line profiles are complex in structure.," As noted by Ilovaisky (1989), the Balmer line profiles are complex in structure."
+ The strongest. feature in the Dine profiles is the relatively narrow (ENIM ~3 A)) absorption component (at e=O km s7 Leasin 4)) which. corresponds, The strongest feature in the line profiles is the relatively narrow (FWHM 3 ) absorption component (at $v=0$ km $^{-1}$ in \ref{fig:stack}) ) which corresponds
+This work is the fourth part iu a series describing siauultaneous iuultifrequency observations of radio pulsus.,This work is the fourth part in a series describing simultaneous multi-frequency observations of radio pulsars.
+ The project had its beeimniug in a collaboration known as the Europeau Pulsar Network (EPN) but has erown bevoud its initial participating partuers (Lorimer et al. 1998)), The project had its beginning in a collaboration known as the European Pulsar Network (EPN) but has grown beyond its initial participating partners (Lorimer et al. \cite{ljs98}) )
+ aud now also utilizes telescopes outside Europe such as the Cüaut Moetre Badio Telescope CGMBT). India. which provided low frequency coverage for the study preseuted here.," and now also utilizes telescopes outside Europe such as the Giant Metre Radio Telescope (GMRT), India, which provided low frequency coverage for the study presented here."
+ The project's observations are aimed at providing information and constraiuts ou the vet to be determined cussion mechanism of radio pulsars., The project's observations are aimed at providing information and constraints on the yet to be determined emission mechanism of radio pulsars.
+ While the properties of average profiles are deteruined bv geometrical factors to a large extent. the study of suele pulses promises to be the best wav to investigate enüssion properties produced by the radiatiou process.," While the properties of average profiles are determined by geometrical factors to a large extent, the study of single pulses promises to be the best way to investigate emission properties produced by the radiation process."
+ The comparison of data obtained at several frequeucies at the same time offers the possibility of distinguishing between propagation effects aud of properties inherited from the original radiation process., The comparison of data obtained at several frequencies at the same time offers the possibility of distinguishing between propagation effects and of properties inherited from the original radiation process.
+ This has 30011 demonstrated diu the first three paper of this series bv dEarastergiou et al. (, This has been demonstrated in the first three paper of this series by Karastergiou et al. (
+2001. 2002. 2003).. where we Ciceutrated on the polarization properties of individual pulses.,"2001, 2002, \nocite{khk+01,kkj+02, kjk03}, where we concentrated on the polarization properties of individual pulses."
+ In this paper. we shift our focus towards the fiux density spectra of single pulses aud their componeuts.," In this paper, we shift our focus towards the flux density spectra of single pulses and their components."
+ Detailed kuowlelee about the spectral behaviour of individual cussion entities can help to put coustraiuts ou theoretical models aud the energy. distribution of the cluitting particles., Detailed knowledge about the spectral behaviour of individual emission entities can help to put constraints on theoretical models and the energy distribution of the emitting particles.
+ Generally. ouly the spectral behaviour of average pulse profiles is well known. but obviously all iuformation about ανασα. processes in the pulsar magnetosphere is lost iu the averaging process.," Generally, only the spectral behaviour of average pulse profiles is well known, but obviously all information about dynamical processes in the pulsar magnetosphere is lost in the averaging process."
+ Ince. it is desirable to study the radio spectrum of the single pulses themselves.," Hence, it is desirable to study the radio spectrum of the single pulses themselves."
+ Suuultaucous ΙΙΠοιος. observations are difficult to realize. iu particular if total power information is to be obtained with a single radio telescope.," Simultaneous multi-frequency observations are difficult to realize, in particular if total power information is to be obtained with a single radio telescope."
+ This enipliasizes the need for coordinated iuulti-statiou experiments such as reported were., This emphasizes the need for coordinated multi-station experiments such as reported here.
+ In the past. rather few such studies have appeared in the literature (see Isarastereiou ot al.," In the past, rather few such studies have appeared in the literature (see Karastergiou et al."
+ 2001)., 2001).
+ Most of the studies on sinultauecous flux ceusity mcasurements were doue in the carly davs of pulsar research., Most of the studies on simultaneous flux density measurements were done in the early days of pulsar research.
+" These were stuuimarized bx Dacker Fisher 1971). who provided the larecst frequency COVCTALC SO far frou, 250 to ssh MIIz. although they were linites to sinele-polarization information for most of their frequencies."," These were summarized by Backer Fisher \nocite{bf74}, who provided the largest frequency coverage so far from 250 to 8085 MHz, although they were limited to single-polarization information for most of their frequencies."
+ However. iu coutrast to Robinson et al. (1968).," However, in contrast to Robinson et al. \nocite{rcg+68},"
+. who covered a wide range from 85 to 1110 MITz or PSR D1919|21. Backer Fisher did not present the spectra of single oulses but averaged over a certain time iun order to combat ιο effects of interstellar sciutillation (ISS).," who covered a wide range from 85 to 1410 MHz for PSR B1919+21, Backer Fisher \nocite{bf74} did not present the spectra of single pulses but averaged over a certain time in order to combat the effects of interstellar scintillation (ISS)."
+ Robinson et 1., Robinson et al.
+ did uot correct for ISS. while we know today that in je strong scattering regiae (i.e. at observing frequencies low a few CGIIz for low dispersion measure pulsars). liffractive scintillation results iu stroug imodulatious of 16 observed signal.," did not correct for ISS, while we know today that in the strong scattering regime (i.e. at observing frequencies below a few GHz for low dispersion measure pulsars), diffractive scintillation results in strong modulations of the observed signal."
+ ISS is therefore au iiportaut problem ο overcome When trviug to derive the iutriusic radio spectrum. and heuce we pay particular attention to this aspect before deriving and discussing the flux. deusity spectra of individual radio pulses.," ISS is therefore an important problem to overcome when trying to derive the intrinsic radio spectrum, and hence we pay particular attention to this aspect before deriving and discussing the flux density spectra of individual radio pulses."
+ The plui for the rest of the paper is as follows: after brieflv sununuiarizuig the observatious of two pulsus. PSRs D0329|51 aud D1133]16. we inspect the fux densities measured and correct for the effects of ISS.," The plan for the rest of the paper is as follows: after briefly summarizing the observations of two pulsars, PSRs B0329+54 and B1133+16, we inspect the flux densities measured and correct for the effects of ISS."
+ We then derive he individual flux deusity spectra and discuss the results iu the last section., We then derive the individual flux density spectra and discuss the results in the last section.
+ The data presented here were obtained ou 2000 January l1. usine the Effelsbere 100-20 telescope at 1500 MIIZ. the Lovell 76-11 telescope at Jodrell Bank at 1112 \UIz and the GAIRT at 626 MIIZ.," The data presented here were obtained on 2000 January 4, using the Effelsberg 100-m telescope at 4850 MHz, the Lovell 76-m telescope at Jodrell Bank at 1412 MHz and the GMRT at 626 MHz."
+ The GAIRT also observed PSR D0320|51 at 235 MIIz and PSR D1133|16 at, The GMRT also observed PSR B0329+54 at 238 MHz and PSR B1133+16 at
+In Fiewwe 2 we show the SED computed from the same Gimeslice used in Figure 1. (thick line) along with the average SED (thin line).,In Figure \ref{fig:fidspect} we show the SED computed from the same timeslice used in Figure \ref{fig:physvar} (thick line) along with the average SED (thin line).
+ The figure also shows a selection of radio. NIR and N-ray. observational data points (relerences given in the figure caption).," The figure also shows a selection of radio, NIR and X-ray observational data points (references given in the figure caption)."
+ The SED has a peak al ~690GIIz due to thermal svnchrotron enission., The SED has a peak at $\sim 690 \GHz$ due to thermal synchrotron emission.
+ Below ~100GIIz it fails to fit the data because that emission is produced outside the simulation volume., Below $\sim 100\GHz$ it fails to fit the data because that emission is produced outside the simulation volume.
+ A second peak in the [ar UV is due to the first Compton scattering order. and at >LOMLz (=40keV) the photons are produced by (vo or more scatterings.," A second peak in the far UV is due to the first Compton scattering order, and at $\gtrsim 10^{19} {\rm Hz}$ $\gtrsim 40\keV$ ) the photons are produced by two or more scatterings."
+ Where do the photons originate in each band?, Where do the photons originate in each band?
+ Figure 3. maps the points of origin for photons in the svnchrotron peak (230 to 690GIIz). in the NIR (2—10jm) and in the X-ray (2-8 keV).," Figure \ref{fig:emmap} maps the points of origin for photons in the synchrotron peak $230$ to $690
+\GHz$ ), in the NIR $2-10\mum$ ) and in the X-ray (2-8 keV)."
+ Most of the submillimeter emission originates near the midplane αἱ ολ.," Most of the submillimeter emission originates near the midplane at $4 \Rg < r < 6
+\Rg$ ."
+" NIR photons are produced in hot regions with high magnetic field intensity and high gas temperature. aid (hese are concentrated close to the innermost stable circular orbit. (ISCO: rjseo&26Mc Dor a,= 0.94) ie. they come [rom between r«66M/e."," NIR photons are produced in hot regions with high magnetic field intensity and high gas temperature, and these are concentrated close to the innermost stable circular orbit (ISCO; $r_{ISCO} \approx 2 \Rg$ for $a_* = 0.94$ ) i.e. they come from between $2 \Rg < r < 6 \Rg$ ."
+ All photons responsible for the formation of the first Compton bump are up- between 2GM/e?<r«8GM/cC bul the 2-8 keV emission is produced mainly by scalterings in the hottest parts of the disk al 2GAL/e7<r«36M/cC.," All photons responsible for the formation of the first Compton bump are up-scattered between $2 \Rg < r < 8 \Rg$ but the 2-8 keV emission is produced mainly by scatterings in the hottest parts of the disk at $2
+\Rg < r < 3 \Rg$."
+ Emission can be seen around (he borders of the funnel in each panel. but (his is at a low level and is associated with unreliable temperatures assigned by ito the [unnel region.," Emission can be seen around the borders of the funnel in each panel, but this is at a low level and is associated with unreliable temperatures assigned by to the funnel region."
+" The photons that form the submillimeter peak (690 GlIz) originate at (r)zzάλλος. where πι)24x10em7 (D)zz28 Gauss. and (0,)z17."," The photons that form the submillimeter peak $690 \GHz$ ) originate at $\<r\> \approx 4 \Rg$, where $\<n_e\> \approx 4 \times 10^6 \cm^{-3}$ $\<B\> \approx 28$ Gauss, and $\<\Theta_e\> \approx 17$."
+" The emissivitv vj,.(¢=x/2) calculated. [rom these averaged values peaks αἱ νο8x10! Hz. which is quite close to the actual peak in pL, (where pL,=Ax D7rF,)."," The emissivity $\nu
+j_\nu(\theta=\pi/2)$ calculated from these averaged values peaks at $\nu
+\sim 8 \times 10^{11}$ Hz, which is quite close to the actual peak in $\nu L_\nu$ (where $\nu L_\nu = 4 \pi D^2 \nu F_\nu$ )."
+ The averaged values also vield 0.25; the submillimeter peak plasma is dominated bv gas pressure rather (han magnetic pressure., The averaged values also yield $B^2/(8 \pi P) \simeq 0.25$ ; the submillimeter peak plasma is dominated by gas pressure rather than magnetic pressure.
+" The optically thin svnchrotron. photons in the energy range between the svnchrotron peak and first-order Compton scattering bump (6.9x10<<1.510! Iz) tend to arise in current sheets at 26M/c6<r<6GM/cC. where 1.5€n,/(10em*)<4. 30—35 Gauss. and 9,280 or higher."," The optically thin synchrotron photons in the energy range between the synchrotron peak and first-order Compton scattering bump $6.9 \times
+10^{11} < \nu < 1.5 \times 10^{14}$ Hz) tend to arise in current sheets at $2 \Rg < r < 6 \Rg$, where $1.5 < n_e/(10^6 \cm^{-3}) < 4$, $\<B\> =
+30-35$ Gauss, and $\Theta_e \approx 80$ or higher."
+" The corresponding vj, peaks in the mid-IR. at 12jm."," The corresponding $\nu j_\nu$ peaks in the mid-IR, at $12 {\rm \mum}$."
+ These current sheets have higher entropy than the surrounding plasma. consistent. with the idea that they are heated bv (numerical) reconnection.," These current sheets have higher entropy than the surrounding plasma, consistent with the idea that they are heated by (numerical) reconnection."
+ The prominence of (he current sheets is likely an artifact of axisvinmetry., The prominence of the current sheets is likely an artifact of axisymmetry.
+ Similar current sheets are seen in axisvinmetrie shearing, Similar current sheets are seen in axisymmetric shearing
+"As pointed out by ?,, the value of dm/dt is also an indicator of the speed class.","As pointed out by \cite{1989clno.conf....1W}, the value of dm/dt is also an indicator of the speed class."
+ We thus use the dm/dt values of our sample to derive the frequency of different speed classes of nova in M31., We thus use the dm/dt values of our sample to derive the frequency of different speed classes of nova in M31.
+" A comparison of the dm/dt distribution in our sample to the novae presented by ? is shown in Fig. 15,"," A comparison of the dm/dt distribution in our sample to the novae presented by \cite{2004MNRAS.353..571D} is shown in Fig. \ref{fig.n_fq},"
+ the two M31 samples agree rather well within the statistical errors (shown are only the J/n number count errors)., the two M31 samples agree rather well within the statistical errors (shown are only the $\sqrt{n}$ number count errors).
+ The Milky Way data set, The Milky Way data set
+alpha clement eraclicuts (7).. sizes of spiral disks (?).. and he correlation of post starburst ealaxy positions with the location of X-ray substructures (2?) and infall regions (?)..,"alpha element gradients \citep{smi09}, sizes of spiral disks \citep{agu04}, and the correlation of post starburst galaxy positions with the location of X-ray substructures \citep{pog04} and infall regions \citep{mah11}."
+ A population of recently quenched elliptical galaxies would also explain the observations of ? who studied the color-magnitude diagram of Coma at Mid-IR wavelengths fiucding that of ellipticals were truly passive. but showed evidence for very recent activity.," A population of recently quenched elliptical galaxies would also explain the observations of \citet{cle09} who studied the color-magnitude diagram of Coma at Mid-IR wavelengths finding that of ellipticals were truly passive, but showed evidence for very recent activity."
+ This picture is consistent with results from cosmological simulations that show the build-up of massive clusters through: accretion of infalliug ealaxies and galaxy eroups., This picture is consistent with results from cosmological simulations that show the build-up of massive clusters through accretion of infalling galaxies and galaxy groups.
+o The uunucerical simulations of? show higher star formation rates (SFRs) in the accreted galaxies. in agreement with recent observational results.," The numerical simulations of \citet{bek99} show higher star formation rates (SFRs) in the accreted galaxies, in agreement with recent observational results."
+ ? found that most of the obscured star formation in the complex Abell 9=1Abcll 902 comes from galaxies between aud at tle edges ofthe two clusters., \citet{gal09} found that most of the obscured star formation in the complex Abell 901-Abell 902 comes from galaxies between and at the edges of the two clusters.
+ Our work on Abell 176: (?7) revealed a ~12 Alpe fbuueut o| Mid-IR enüttine galaxies connecting this cluster to the uearby Abel L770.," Our work on Abell 1763 \citep{edw10,fad08} revealed a $\sim$ 12 Mpc filament of Mid-IR emitting galaxies connecting this cluster to the nearby Abell 1770."
+ For that case. the fraction of star fornüng flhluueut ealaxics was found to be double that in the cluster regions.," For that case, the fraction of star forming filament galaxies was found to be double that in the cluster regions."
+ Work bv ?.. 7? and ? has shown that the dwiurf star forming galaxy population iu Coma is quenched im the high deusitv core.," Work by \citet{jen07}, \citet{bai06} and \citet{mah10} has shown that the dwarf star forming galaxy population in Coma is quenched in the high density core."
+ However. the dwarf galaxies do not necessarily have very inteuse starbursts. therefore. a study of the star formation rates is miportant.," However, the dwarf galaxies do not necessarily have very intense starbursts, therefore, a study of the star formation rates is important."
+ This requires an accurate method for determining ealaxy stellar mass determinations., This requires an accurate method for determining galaxy stellar mass determinations.
+ Thus. of particular inuportauce to this study are our new Near-IR observations as they bracket the stellar dass bump. allowing for accurate stellar mass measurements.," Thus, of particular importance to this study are our new Near-IR observations as they bracket the stellar mass bump, allowing for accurate stellar mass measurements."
+ Tn this paper. we combine these original Near-IR photomoetrv with public Mid-Iufired photometry from the IRAC and MIPS instruments aboard theSpitzer Space Telescope. in addition to optical photometry from the Sloan Digital Sky Survey (SDSS).," In this paper, we combine these original Near-IR photometry with public Mid-Infrared photometry from the IRAC and MIPS instruments aboard the Space Telescope, in addition to optical photometry from the Sloan Digital Sky Survey (SDSS)."
+ We eather optical spectra from the SDSS as well as frou our own campaign., We gather optical spectra from the SDSS as well as from our own campaign.
+ With these data we build) spectral enerev distributions (SEDs) of Coma cluster members., With these data we build spectral energy distributions (SEDs) of Coma cluster members.
+ Star formation rates are calculated frou the total iufrared Iuuinositv derived frou best-fit template galaxy. SEDs., Star formation rates are calculated from the total infrared luminosity derived from best-fit template galaxy SEDs.
+" Iu the next section. we present the observations and reduction methods for the deep Near-IR J. IL. and Ik, data which was collected from the Palomar WIRC instirmment. as well as theSpetcer MOPS 21. 70. aud 1607420 and TRAC 3.6. 15. he. and μι data which we have re-reduced."," In the next section, we present the observations and reduction methods for the deep Near-IR J, H, and $_{s}$ data which was collected from the Palomar WIRC instrument, as well as the MIPS 24, 70, and $\mu$ m and IRAC 3.6, 4.5, 5.8, and $\mu$ m data which we have re-reduced."
+ We also include the optical photometry from the SDSS Data Release 7. the available spectra from that program. as well as 338 spectra of {μαι selected. galaxies we gathered from the ITYDRA iustimuent ou the WIYN telescope.," We also include the optical photometry from the SDSS Data Release 7, the available spectra from that program, as well as 338 spectra of $\mu$ m selected galaxies we gathered from the HYDRA instrument on the WIYN telescope."
+ Tn Section 3. we prescut the catalog of Yau selected. galaxics in Coma. alone with the associated optical and infrared fluxes iud. redshifts.," In Section 3, we present the catalog of $\mu$ m selected galaxies in Coma, along with the associated optical and infrared fluxes and redshifts."
+ We release our fully reduced images and spectra to the public via the NASA/IPAC Tnuftared Science Archive (IRSA) and the NASA Extragalactic Database. (NED)., We release our fully reduced images and spectra to the public via the NASA/IPAC Infrared Science Archive (IRSA) and the NASA Extragalactic Database (NED).
+ These data allow us to determine the amount of ACN contamination. measure stellar mass. and derive the total infrared luminosities required to calculate specific star formation rates (SFRs}. which we present Section .," These data allow us to determine the amount of AGN contamination, measure stellar mass, and derive the total infrared luminosities required to calculate specific star formation rates (SFRs), which we present Section 4."
+ We compare these obscured SFRs to unobscured rates measured from nebular enussion lines in Section 5. where we also compare other proxies for the obscured SER aud discuss the specific SFR as a function of location.," We compare these obscured SFRs to unobscured rates measured from nebular emission lines in Section 5, where we also compare other proxies for the obscured SFR and discuss the specific SFR as a function of location."
+ We conclude. in Section 6. that the starburst galaxies are blue cavarfs aud that there is a paucity of these in the deuse cluster core.," We conclude, in Section 6, that the starburst galaxies are blue dwarfs and that there is a paucity of these in the dense cluster core."
+ This is iu line with the hypothesis that starburst galaxies have been recently accreted onto the cluster aud will be quenched as they make their wav to the COLC., This is in line with the hypothesis that starburst galaxies have been recently accreted onto the cluster and will be quenched as they make their way to the core.
+ Tu this section. we describe the observations aud data reduction.," In this section, we describe the observations and data reduction."
+ Table 1 lists our observational dataset and Figure | shows the coverage of the infrared data and the location of spectroscopic members., Table \ref{obs} lists our observational dataset and Figure \ref{spitzer_cov} shows the coverage of the infrared data and the location of spectroscopic members.
+" The photometric data was acquired from recent WIRC Near-IR anc archival Maid ane Fax-IRSpitser nagging (PIs: Fazio. Ricke. ανα), optical catalogs from the SDSS DR7. aud. radio source cataloes from the Faint Images ofthe Radio Sky at Tweutv-Centineters (FIRST)."," The photometric data was acquired from recent WIRC Near-IR and archival Mid and Far-IR imaging (PIs: Fazio, Rieke, Kitayama), optical catalogs from the SDSS DR7, and radio source catalogs from the Faint Images of the Radio Sky at Twenty-Centimeters (FIRST)."
+ Because SED fitting is more accurate when the MIPS Yau datapoiut exists. we include only galaxies which are detected at an to at least the 30 level.," Because SED fitting is more accurate when the MIPS $\mu$ m datapoint exists, we include only galaxies which are detected at $\mu$ m to at least the $\sigma$ level."
+ The AIIPS data field of view covers the full core of the cluster. (0.2 - «rogo. (where rogo is the radius," The MIPS data field of view covers the full core of the cluster, (0.2 - $\times$ $_{200}$ (where $_{200}$ is the radius"
+unseen planets.,unseen planets.
+ Modeling indicates that Neptune to Jupiter mass planets at several tens of AUs are typically needed to reproduce the observations., Modeling indicates that Neptune to Jupiter mass planets at several tens of AUs are typically needed to reproduce the observations.
+ Although the current picture is still far from complete. all of these results offer a very interesting first glimpse of the full extent of orbital separations and masses of giant planets and low-mass brown dwarf companions around stars.," Although the current picture is still far from complete, all of these results offer a very interesting first glimpse of the full extent of orbital separations and masses of giant planets and low-mass brown dwarf companions around stars."
+ The planetary mass candidate companion to the young solar mass starJ160929.1-210524.. which is a member of the 5 Myr-old Upper Scorpius association (Preibischetal.2002) is the subject of this paper.," The planetary mass candidate companion to the young solar mass star, which is a member of the 5 Myr-old Upper Scorpius association \citep{preibisch02} is the subject of this paper."
+" The candidate companion. lying only 2.2""(-.330 AU projected) from the primary star. appears to be a young planetary mass (0.008 ..)) object in the Upper Sco association. based on near-infrared imaging (HK) and spectroscopy (HK)."," The candidate companion, lying only $\sim$ 330 AU projected) from the primary star, appears to be a young planetary mass $0.008$ ) object in the Upper Sco association, based on near-infrared imaging $JHK$ ) and spectroscopy $HK$ )."
+" In particular. the triangular shape of its H-band spectrum clearly indicates low surface gravity, and the overall spectrum is in excellent agreement with the spectra of known low-mass brown dwarfs in the same association as well as with model spectra with Z;j~1800 K and logg 4. consistent with the inferred age and mass."," In particular, the triangular shape of its $H$ -band spectrum clearly indicates low surface gravity, and the overall spectrum is in excellent agreement with the spectra of known low-mass brown dwarfs in the same association as well as with model spectra with $T_{\rm eff}\sim1800$ K and $\log{g}\sim4$ , consistent with the inferred age and mass."
+ While at the time of discovery. we had no proper motion measurements and henee could not confirm it was gravitationally bound. the proximity and youth of the candidate companion offered very compelling evidence that the two form a bound pair.," While at the time of discovery, we had no proper motion measurements and hence could not confirm it was gravitationally bound, the proximity and youth of the candidate companion offered very compelling evidence that the two form a bound pair."
+ In this paper. we present follow-up astrometric observations that confirm the common proper motion of the candidate companion with the primary star.," In this paper, we present follow-up astrometric observations that confirm the common proper motion of the candidate companion with the primary star."
+ In addition. to improve the characterization of this system. we present new photometric observations at 3.05 jim and 3.78 jim (L). new spectroscopic observations in the J-band. and new high-contrast angular differential imaging (ADI. Maroisetal. 2006)) observations.," In addition, to improve the characterization of this system, we present new photometric observations at 3.05 $\mu$ m and 3.78 $\mu$ m $L^\prime$ ), new spectroscopic observations in the $J$ -band, and new high-contrast angular differential imaging (ADI, \citealp{marois06}) ) observations."
+ All observations used for this were obtained with the Gemini North telescope and the ALTAIR adaptive optics system (Herriotetal.2000)., All observations used for this were obtained with the Gemini North telescope and the ALTAIR adaptive optics system \citep{herriot00}.
+. The target star itself was used for wave front sensing and the field lens of ALTAIR was used to reduce off-axis Strehl degradation due to anisoplanatism., The target star itself was used for wave front sensing and the field lens of ALTAIR was used to reduce off-axis Strehl degradation due to anisoplanatism.
+ All imaging observations were obtained with the NIRI camera (Hodappetal.2003) in f/32 mode., All imaging observations were obtained with the NIRI camera \citep{hodapp03} in $f/32$ mode.
+" Used with the AO system and the field lens. this yields a pixel scale of 21.4 mas and a field of view of 21.9""«21.9” for the full detector array."," Used with the AO system and the field lens, this yields a pixel scale of 21.4 mas and a field of view of $21.9\arcsec \times 21.9\arcsec$ for the full detector array."
+ The initial ΗΝ imaging observations of wwere obtained in 2008 April and June and are detailed in Paper 1. For astrometric follow-up of the candidate companion. the star was observed again on 2009 April 6 and 2009 July 1.," The initial $JHK$ imaging observations of were obtained in 2008 April and June and are detailed in Paper I. For astrometric follow-up of the candidate companion, the star was observed again on 2009 April 6 and 2009 July 1."
+" These observations were made in the K' filter and used a 5-point dither pattern corresponding to the center and four comers of a square of side10"".", These observations were made in the $K^\prime$ filter and used a 5-point dither pattern corresponding to the center and four corners of a square of side.
+.. At each dither position. we obtained one co-addition of thirty 0.3 s integrations in fast. high read-noise mode. and one 10 s integration in slow. low read-noise mode.," At each dither position, we obtained one co-addition of thirty 0.3 s integrations in fast, high read-noise mode, and one 10 s integration in slow, low read-noise mode."
+ At each position this provides an unsaturated image of the target star and a much deeper image of the field that can be spatially registered and scaled in flux without ambiguity., At each position this provides an unsaturated image of the target star and a much deeper image of the field that can be spatially registered and scaled in flux without ambiguity.
+ These imaging data were reduced using custom routines., These imaging data were reduced using custom routines.
+ A sky frame was constructed by taking the median of the images at all dither positions after masking out the regions dominated by the target’s signal., A sky frame was constructed by taking the median of the images at all dither positions after masking out the regions dominated by the target's signal.
+ After subtraction of this sky frame. the images were divided by à normalized flat-field.," After subtraction of this sky frame, the images were divided by a normalized flat-field."
+ Then isolated bad pixels were replaced by the interpolated value of a third-order polynomial surface fit to the good pixels in à 7« pixels box while clustered bad pixels were simply masked out., Then isolated bad pixels were replaced by the interpolated value of a third-order polynomial surface fit to the good pixels in a $7\times7$ pixels box while clustered bad pixels were simply masked out.
+ Next the images were distortion corrected using the distortion solution provided by the Gemini Finally. the long-exposure Images were properly scaled in intensity and merged with their corresponding short-exposure images.," Next the images were distortion corrected using the distortion solution provided by the Gemini Finally, the long-exposure images were properly scaled in intensity and merged with their corresponding short-exposure images."
+ Follow-up observations of the system were also obtained at 3.05 jm (H:O ice filter. bandpass) and 3.78 jim (£L filter) on 2009 May 2 and 2009 June 8. respectively. to constrain the spectral energy distribution of the candidate companion further.," Follow-up observations of the system were also obtained at 3.05 $\mu$ m $_2$ O ice filter, bandpass) and 3.78 $\mu$ m $L^\prime$ filter) on 2009 May 2 and 2009 June 8, respectively, to constrain the spectral energy distribution of the candidate companion further."
+" For 3.05 jm. 20 co-additions of ten 5 s integrations were obtained over 20 dither positions covering a square of side8""."," For 3.05 $\mu$ m, 20 co-additions of ten 5 s integrations were obtained over 20 dither positions covering a square of side."
+". To enable a shorter integration time and avoid saturation for 3.78 jm. only the central 768« pixels of the array was read and 30 co-additions of three hundred 0.2 s yointegrations were obtained over 30 dither positions covering a quare of side6""."," To enable a shorter integration time and avoid saturation for 3.78 $\mu$ m, only the central $768\times768$ pixels of the array was read and 30 co-additions of three hundred 0.2 s integrations were obtained over 30 dither positions covering a square of side."
+ For the long wavelength imaging. the data reduction ivolved first removing the sky background by subtracting. for each image. the mean of the two images acquired closest in time for which the star is not located in the vicinity of the candidate companion in the image considered.," For the long wavelength imaging, the data reduction involved first removing the sky background by subtracting, for each image, the mean of the two images acquired closest in time for which the star is not located in the vicinity of the candidate companion in the image considered."
+ The images were then divided by a normalized flat field image created by taking the median of all images after masking out the star in each one and normalizing them to a median value of one., The images were then divided by a normalized flat field image created by taking the median of all images after masking out the star in each one and normalizing them to a median value of one.
+ A ~121« pixels median-filtered image was subsequently subtracted from each image to remove any low spatial frequency residual background signal. which was occasionally present: to avoid any potential bias in doing this. the median filter was computed after masking out à 9« pixel box over the star and companion.," A $\sim 121 \times 121$ pixels median-filtered image was subsequently subtracted from each image to remove any low spatial frequency residual background signal, which was occasionally present; to avoid any potential bias in doing this, the median filter was computed after masking out a $9\times 9$ pixel box over the star and companion."
+ All images were finally aligned and combined using a 3o clipped average., All images were finally aligned and combined using a $\sigma$ clipped average.
+ An angular differential imaging (ADI. Maroisetal.2006)) sequence of observations was obtained on 2009 April 6 to search for additional. closer and/or fainter companions in the system.," An angular differential imaging (ADI, \citealt{marois06}) ) sequence of observations was obtained on 2009 April 6 to search for additional, closer and/or fainter companions in the system."
+" These observations were done in the A"" filter and only the central 512« pixel region of the detector was read.", These observations were done in the $K^\prime$ filter and only the central $512\times512$ pixel region of the detector was read.
+ The main ADI sequence consisted of 320 exposures of 10 s in low read noise mode: after every 20 of these Images. two short. unsaturated co-additions of ten 0.3 s exposures in medium read noise mode were obtained to monitor the star position. the Strehl ratio. and the sky transparency.," The main ADI sequence consisted of 320 exposures of 10 s in low read noise mode; after every 20 of these images, two short, unsaturated co-additions of ten 0.3 s exposures in medium read noise mode were obtained to monitor the star position, the Strehl ratio, and the sky transparency."
+ The total duration of the sequence is 1.9 hours and the total field of view rotation 1s397., The total duration of the sequence is 1.9 hours and the total field of view rotation is.
+. Both before and after the main ADI sequence. a 5-position dither pattern of short unsaturated observations (10. 0.35) was obtained to provide a sky frame for the reduction of the short exposures as well as additional image of the point-spread-function (PSF) core.," Both before and after the main ADI sequence, a 5-position dither pattern of short unsaturated observations $10\times0.3 s$ ) was obtained to provide a sky frame for the reduction of the short exposures as well as additional image of the point-spread-function (PSF) core."
+ The short. unsaturated exposures for the ADI observations were reduced as above for the K-band imaging.," The short, unsaturated exposures for the ADI observations were reduced as above for the $K$ -band imaging."
+ The long. saturated exposures were dark subtracted. divided by a normalized flat field image. corrected for distortion using the solution provided by the Gemini staff. and padded with zeros to increase their size to 750 pixels and thus avoid loosing parts of the field of view during subsequent shifts and rotations.," The long, saturated exposures were dark subtracted, divided by a normalized flat field image, corrected for distortion using the solution provided by the Gemini staff, and padded with zeros to increase their size to $750\times 750$ pixels and thus avoid loosing parts of the field of view during subsequent shifts and rotations."
+ Bad pixels were corrected or masked as explained, Bad pixels were corrected or masked as explained
+ab /=0 proved sullicient in our simulations).,at $t=0$ proved sufficient in our simulations).
+" Hi the initial 5,,5,«100. no extra time delay of the ERC emission is possible."," If the initial $\gamma_{{{}} min}\ll100$, no extra time delay of the ERC emission is possible."
+ In addition. if the ERC emission exhibits an extra delay compared to the svnchrotron variability al lower frequencies. the spectral index of the ERC emission must be positive for a sienilicant fraction of the duration of the flare.," In addition, if the ERC emission exhibits an extra delay compared to the synchrotron variability at lower frequencies, the spectral index of the ERC emission must be positive for a significant fraction of the duration of the flare."
+ This is because (his extra delay is only possible for the ERC light curves at [requencies below (he initial turn-over Irequency of the SED., This is because this extra delay is only possible for the ERC light curves at frequencies below the initial turn-over frequency of the SED.
+ On the other haud. the spectral index of the SSC emission will be negative even at the beginning of the SSC! flare because the turn-over frequency of the SSC spectrum is determined by the svnchrotron self-absorption frequency. which is typically much lower (han the characteristic frequency of dust emission. ορ.," On the other hand, the spectral index of the SSC emission will be negative even at the beginning of the SSC flare because the turn-over frequency of the SSC spectrum is determined by the synchrotron self-absorption frequency, which is typically much lower than the characteristic frequency of dust emission, $\nu_{{{}} tor}$."
+ Similar behavior of the ERC flares has been found by Sikoraetal.(2001)., Similar behavior of the ERC flares has been found by \citet{sik01}.
+. These authors explained (he delavs that they found by appealing to the gradient in the magnetic field along (he jet. which leads (o a decrease in (he critical frequency of svuchrotron emission.," These authors explained the delays that they found by appealing to the gradient in the magnetic field along the jet, which leads to a decrease in the critical frequency of synchrotron emission."
+ In their model the flare is produced bv a plasma in a eeometrically thin shell energized by forward/reverse shocks produced lollowing the collision between portions of the jet propagating with different relativistic speeds., In their model the flare is produced by a plasma in a geometrically thin shell energized by forward/reverse shocks produced following the collision between portions of the jet propagating with different relativistic speeds.
+ In contrast with our model. (hese authors assume that the injection of relativistic electrons is uniform throughout the shell.," In contrast with our model, these authors assume that the injection of relativistic electrons is uniform throughout the shell."
+ They aseribe the time delavs between ERC and svnchrotron flares to gradients in magnetic field., They ascribe the time delays between ERC and synchrotron flares to gradients in magnetic field.
+ Their modeling is thus sultable to more prolonged flares than the ones considered in this paper., Their modeling is thus suitable to more prolonged flares than the ones considered in this paper.
+ The results reported in Paper I are based on (he assumption that (he energy. losses ol electrons are dominated by svinchrotron emission and that ERC emission is negligible compared to the SSC flux levels at the [requencies of interest., The results reported in Paper I are based on the assumption that the energy losses of electrons are dominated by synchrotron emission and that ERC emission is negligible compared to the SSC flux levels at the frequencies of interest.
+ These assumptions are valid when the properties of the sources of external emission as well as the location of the emitting plasma in (he jet are as in Fig. 10.., These assumptions are valid when the properties of the sources of external emission as well as the location of the emitting plasma in the jet are as in Fig. \ref{mesp3.0.less}.
+" Decreasing the size of the torus and/or placing the emitting plasma farther down the jet substantially decreases the Doppler boosting; even if I, remains the same.", Decreasing the size of the torus and/or placing the emitting plasma farther down the jet substantially decreases the Doppler boosting even if $\Gamma'_p$ remains the same.
+ For the parameters used in Fig., For the parameters used in Fig.
+ LO the ERC losses of electrons are only a fraction of the svnehrotron losses.," \ref{mesp3.0.less}
+ the ERC losses of electrons are only a fraction of the synchrotron losses."
+ Also. the SSC {lux exceeds the ERC flux by a factor of ad least a few at all [requencies and by several orders of magnitude near 10/5Hz (in the plasma rest [rame). al which the SSC flave has been shown in Paper I to have an extra delay due to the (travel time of the seed photons.," Also, the SSC flux exceeds the ERC flux by a factor of at least a few at all frequencies and by several orders of magnitude near $10^{16}\,\mbox{Hz}$ (in the plasma rest frame), at which the SSC flare has been shown in Paper I to have an extra delay due to the travel time of the seed photons."
+" The results of simulations in the plasma rest [rame can be translormed to the observer's frame according to the following expressions: FLUobs!o5oO(DoD*ΤΕ(μι). 2). and obs=(142)lops/Oone. Where the asterisk denotes quantities in the observer's frame. Oops Is the Doppler factor determined by I, and the viewing angle. and z and D* are the redshift and distance to the blazar (24 is the reference distance used in the simulations)."," The results of simulations in the plasma rest frame can be transformed to the observer's frame according to the following expressions: $F^{*}_{\nu_{*}}(t^{*}_{{{}} obs})
+=
+\delta_{{{}} obs}^3(D_0/D^{*})^2
+F^E_{{{}} \nu}(t_{{{}} obs})$, $\nu_{*}=\delta_{{{}} obs}\nu/(1+z)$ , and $t^{*}_{{{}} obs}=(1+z)t_{{{}} obs}/\delta_{{{}}
+obs}$, where the asterisk denotes quantities in the observer's frame, $\delta_{{{}} obs}$ is the Doppler factor determined by $\Gamma'_p$ and the viewing angle, and $z$ and $D^{*}$ are the redshift and distance to the blazar $D_0$ is the reference distance used in the simulations)."
+ This implies a particular dependence of (he observed quantities on the bulk Lorentz factor of the emitting plasma., This implies a particular dependence of the observed quantities on the bulk Lorentz factor of the emitting plasma.
+ Iowever. in (he case where (he energy losses of electrons are dominated by," However, in the case where the energy losses of electrons are dominated by"
+gaars need not inevitably be blue objects ?]]. many authors argeted their search on. earlv-tvpe. stars.,"stars need not inevitably be blue objects \cite{Sterken_etal2008}] ], many authors targeted their search on early-type stars."
+ Some authors ooked for 4443-like objects ?]]. ?]. ?]].," Some authors looked for 443-like objects \cite{Neeseetal91}] ], \cite{Calzetti1995}] ], \cite{FabShol95}] ]."
+ The spectrum of 4433 resembles those of late WR stars. ?]]., The spectrum of 433 resembles those of late WR stars \cite{Fab04}] ].
+ The team of Massey πα 2] made à new step oward the study of the massive stellar population of the ADJ333 galaxy., The team of Massey al \cite{Massey2006}] ] made a new step toward the study of the massive stellar population of the 33 galaxy.
+ The above authors used CCD images. of he galaxy to produce a catalog of 146622 stars down a limiting magnitude of 23 containing broad-band iotometric measurements with an accuracy of24., The above authors used CCD images of the galaxy to produce a catalog of 146622 stars down to a limiting magnitude of 23 containing broad-band photometric measurements with an accuracy of.
+ In heir next paper τῇ. the team reported a list of emission stars in Local Group galaxies including 333.," In their next paper \cite{Massey2007_Ha}] ], the team reported a list of emission stars in Local Group galaxies including 33."
+ The main »urpose of this study was to find new LBV star candidates., The main purpose of this study was to find new LBV star candidates.
+ ‘To this end. the above authors adopted the following criteria: a constraint on the line lux: line Πας had to be greater than the SEL] line Hux: the OLLI] line flux hac to exceed the continuum Εαν. and a constraint was imposed on the star's color (sce. ?]] for details).," To this end, the above authors adopted the following criteria: a constraint on the line flux; line flux had to be greater than the [SII] line flux; the [OIII] line flux had to exceed the continuum flux, and a constraint was imposed on the star's color (see \cite{Massey2007_Ha}] ] for details)."
+ Massey et al. ?]], Massey et al. \cite{Massey2007_Ha}] ]
+ report 37 LDV stars ancl LBV candidates in MXj., report 37 LBV stars and LBV candidates in 33.
+ monie of the candidates reported by Massey et al... 2] , Some of the candidates reported by Massey et al. \cite{Massey2007_Ha}] ]
+have been identified before. ?]]. ?]]. την ?]]. ?]].," have been identified before \cite{Corall1996}] ], \cite{Shol_etal97}] ], \cite{FabShol95}] ], \cite{Shol_Fab2000}] ], \cite{Fabrika2000}] ]."
+ We performed. independent aperture photometry of bright objects with V.«1875 in all filters (UDVIU and Lla)) on the CCD frames of Massey et al. 2]., We performed independent aperture photometry of bright objects with $V <$ in all filters (UBVRI and ) on the CCD frames of Massey et al. \cite{Massey2006}] ].
+ We used the same methods to select LBY candidates as Fabrika et al. ?]., We used the same methods to select LBV candidates as Fabrika et al. \cite{Fabrika_etal1997}] ].
+ I5mission-line objects are easily identifiable on the 7Làa--line Hux vs. V-band Uux” diagram. because they lic above the linear relation outlined by objects without the emission (see below for details).," Emission-line objects are easily identifiable on the -line flux vs. V-band flux” diagram, because they lie above the linear relation outlined by objects without the emission (see below for details)."
+ As a result of this work we composed a list of blue LBY- star candidates in 333., As a result of this work we composed a list of blue LBV-like star candidates in 33.
+ In order not to overlook highly redcdened. objects we composed an additional list of red emission objects selected with less stringent color criteria., In order not to overlook highly reddened objects we composed an additional list of red emission objects selected with less stringent color criteria.
+ We downloaded from the NOAÀO science archive //www.archive.noao.edu/nsa/) all primarily— reduced 333 frames taken by. Massey et al. ]], We downloaded from the NOAO science archive ) all primarily reduced 33 frames taken by Massey et al. \cite{Massey2006}] ]
+ with the coordinate grid applied., with the coordinate grid applied.
+ Observations were made in October 2000 and September 2001 with the the μι telescopes of the KPNO and C°PLO telescopes using UDVIU filters and narrow-band (50 AY)lla., Observations were made in October 2000 and September 2001 with the the 4-m telescopes of the KPNO and CTIO telescopes using UBVRI filters and narrow-band (50 ).
+ The UBVRI anc Llo--filter images were taken during photometricnights when the secing was 66 to ος. and OYSS to. 1700. respectively.," The UBVRI and -filter images were taken during photometricnights when the seeing was 6 to 8, and 8 to 0, respectively."
+ The detector criplolved was a mosaic of eight. 4096 CCDs., The detector emplolyed was a mosaic of eight $\times$ 4096 CCDs.
+ Each image has a field of view of 36’ancl a scale of 07226/pix (after primary reduction and distortion correction)., Each image has a field of view of $\times$ and a scale of 26/pix (after primary reduction and distortion correction).
+ Phe \1333 galaxy was subdivided: into three zones with each zone observed [ive times with a small olfset to fill the gaps between the CCDs in the mosaic., The 33 galaxy was subdivided into three zones with each zone observed five times with a small offset to fill the gaps between the CCDs in the mosaic.
+ In this paper we use 15 mosaic images in each filter., In this paper we use 15 mosaic images in each filter.
+ more detailed: description. of the data and. primary. reduction steps can be found in 7]]., A more detailed description of the data and primary reduction steps can be found in \cite{Massey2006}] ].
+ Our aim was to select stars with emission the LBV candidates., Our aim was to select stars with emission — the LBV candidates.
+ The images contain both point and extended objects aud therefore in the process of photometry we must check that measurements in dillerent filters refer to the same object., The images contain both point and extended objects and therefore in the process of photometry we must check that measurements in different filters refer to the same object.
+ This condition should be satisfied even in crowded fields and in complex regions. something that is dillicult to achieve in automatic photometry of several hundred thousand objects ?]].," This condition should be satisfied even in crowded fields and in complex regions, something that is difficult to achieve in automatic photometry of several hundred thousand objects \cite{Massey2007_Ha}] ]."
+ Aperture photometry has an important advantage over PSE photometry in that it makes it possible to determine the Hus from an object without making any assumptions about its structure ancl the form. ancl parameters of the point-spreacl function (PSE) in the Earth atmosphere., Aperture photometry has an important advantage over PSF photometry in that it makes it possible to determine the flux from an object without making any assumptions about its structure and the form and parameters of the point-spread function (PSF) in the Earth atmosphere.
+ The use of the averaged PSE may introduce an error in the estimated flux due to the variation of the form of the PSE across the field., The use of the averaged PSF may introduce an error in the estimated flux due to the variation of the form of the PSF across the field.
+ A problem with aperture photometry in crowded fields is that the aperture radius should be chosen individually for cach object so as to keep all other objects outside the aperture., A problem with aperture photometry in crowded fields is that the aperture radius should be chosen individually for each object so as to keep all other objects outside the aperture.
+ “Phe determination and. application of aperture correction is vet another problem to be faced when performing aperture photometry., The determination and application of aperture correction is yet another problem to be faced when performing aperture photometry.
+ However. if we need not determine the total Dux. but only the (us cillerence in different filters. the same aperture size can be used. in all ilters.," However, if we need not determine the total flux, but only the flux difference in different filters, the same aperture size can be used in all filters."
+ Lt is also important that seeing is approximately the same in allfilters., It is also important that seeing is approximately the same in all filters.
+ Phe observations of Massey et al. 2]], The observations of Massey et al. \cite{Massey2006}] ]
+ meet his condition., meet this condition.
+ We used standard. tools. for. photometric measurcments and all performed auxiliary operations in the xwch mode using programs written in Python language., We used standard tools for photometric measurements and all performed auxiliary operations in the batch mode using programs written in Python language.
+ To choose the optimum aperture size. we performed he photometry of each object varving the aperture radius t from 0755 to 6700 with a step of 07225.," To choose the optimum aperture size, we performed the photometry of each object varying the aperture radius R from 5 to 0 with a step of 25."
+" We estimated he background level using the median averaging over a ring with the inner and outer radii equal to It and AR. respectively, where2""."," We estimated the background level using the median averaging over a ring with the inner and outer radii equal to R and $\Delta$ R, respectively, where."
+ In this approach the [lux rom a single object first smoothly increases with increasing radius It and then remains approximately constant (the Xdateau in ret Iux000119aa)., In this approach the flux from a single object first smoothly increases with increasing radius R and then remains approximately constant (the plateau in \\ref{Flux000119}a a).
+ Phe vertical coordinate of the point where he plateau begins corresponds to the full lux from the star in the filter considered., The vertical coordinate of the point where the plateau begins corresponds to the full flux from the star in the filter considered.
+ After reaching the plateau the Dux comes. practically ceases to vary with aperture radius. whereas the cependenee is very strong when the radius approaches the plateau.," After reaching the plateau the flux becomes practically ceases to vary with aperture radius, whereas the dependence is very strong when the radius approaches the plateau."
+ We consider the optimum radius of the object to be equal to Ro.o. the abscissa of the point corresponding to of the total flux.," We consider the optimum radius of the object to be equal to $_{0.9}$, the abscissa of the point corresponding to of the total flux."
+ We interpolate the curve by drawing a cubic spline through the measured data »oints., We interpolate the curve by drawing a cubic spline through the measured data points.
+ In the presence of a nearby neighboring object. (see Fig., In the presence of a nearby neighboring object (see Fig.
+ Ibb) the ux continues to increase even after reaching he plateau., \ref{Flux000119}b b) the flux continues to increase even after reaching the plateau.
+ We chose the [ux and optimum aperture radius o be same for all filters at the supposed. location of the plateau so as to keep the Dux from the neighboring objects outside the measurement. aperture., We chose the flux and optimum aperture radius to be same for all filters at the supposed location of the plateau so as to keep the flux from the neighboring objects outside the measurement aperture.
+ In the case of a two-dimensional Gaussian. of he total Εαν corresponds to Ito= 1.95m. where σ is he parameter of the Gaussian.," In the case of a two-dimensional Gaussian of the total flux corresponds to $_{0.9}=1.95\sigma$ , where $\sigma$ is the parameter of the Gaussian."
+ The. full width at half maximum (ENLIM) ofa two-dimensional Gaussian (which. w definition. corresponds to the size of a point source) is," The full width at half maximum (FWHM) of a two-dimensional Gaussian (which, by definition, corresponds to the size of a point source) is"
+Solar models have beeji very successful in describing the detailec structure of he solar interior that eierged frou helioseisnuic stucicSs (Clhiisteuseu-Dalseaudοἳ2001 ).,"Solar models have been very successful in describing the detailed structure of the solar interior that emerged from helioseismic studies \citep{modelS,bp01,turck:2001}."
+". Further STLCCOSS came with the first direct ΠιοΟλοit of the tota ""D neutrino flux bv the SNO CollaborationΚΙ (Almacetal2002). which proved to be in exceleut agreenmeaw with standard solar model (SSAL) preoc.‘tions (Balealletal.2001)."," Further success came with the first direct measurement of the total $^8$ B neutrino flux by the SNO Collaboration \citep{sno:2002}, which proved to be in excellent agreement with standard solar model (SSM) predictions \citep{bp01}."
+. Solar models are sensitive o the cetais of their constitutive plysics., Solar models are sensitive to the details of their constitutive physics.
+" IIelioseisiic xedictioIs, for example. depend. critically on racdiative opacities :id nmücroscopic diffusion. and ire also sensitive to unccYtauties dmn input parameters. such as abuncdaxOR deternnimed from absorption line analvses of t1ο Suus atuosphere."," Helioseismic predictions, for example, depend critically on radiative opacities and microscopic diffusion, and are also sensitive to uncertainties in input parameters, such as abundances determined from absorption line analyses of the Sun's atmosphere."
+ Predietious of solar neutrino fluxisl are sensitive to the rates of kev nuclear reactions., Predictions of solar neutrino fluxes are sensitive to the rates of key nuclear reactions.
+ Soar models jiwe plaved a imavor role not ouly in uiderstaucine thie Sun. but also as a test beuch for coustituitive plivsies of stellar models in general.," Solar models have played a mayor role not only in understanding the Sun, but also as a test bench for constituitive physics of stellar models in general."
+ The further devolopuieut of SSAIs both miproved treatmens of the physics and more accurate input paraueters as well as tests of the model against new solar observations— are thus of fundamental iniportance to stellar asroplivsics., The further development of SSMs – both improved treatments of the physics and more accurate input parameters – as well as tests of the model against new solar observations are thus of fundamental importance to stellar astrophysics.
+ Over the last decade. the adveut of 3D lycdrocdvuamic inodels o| uenr-urface solar couvection (Stein&Nord-meveretal. 2001)... a more careful selection of spectral lines (0.8. AllendePrietoetal. 20013) zd. iu some cases. relaxation of the assumption of local thermodynamic equilibiunu iu liue forniion (Asplundctal.2001:Bergemannetal.2010) have resulted. iu a thorough reassessinent of the photspheric solar abuudances.," Over the last decade, the advent of 3D hydrodynamic models of near-surface solar convection \citep{stein:1998,asplund:2000,freytag:2002,wedemeyer:2004}, a more careful selection of spectral lines (e.g. \citealp{allende:2001}) ) and, in some cases, relaxation of the assumption of local thermodynamic equilibrium in line formation \citep{asplund:2004,bergemann:2010} have resulted in a thorough reassessment of the photospheric solar abundances."
+ The 108 recent and complete revision by Aspluudetal.(200 Ον hereafter ACGSSO9. shows a reduction in the alnidauces of the volatile CNO elements aud Ne (~0.1- 0.15 dex) with respect to older conrpilations of soar alnidances (Crevesse&Noels1993:CirevesseSauval1998.hereafter CGS98).. : wowell as snaller changes for alnidit refractory clemeuts. ee. οἱ aud Fe.," The most recent and complete revision by \citet{agss09}, hereafter AGSS09, shows a reduction in the abundances of the volatile CNO elements and Ne $\sim$ 0.1-0.15 dex) with respect to older compilations of solar abundances \citep[hereafter
+ GS98]{gn93,gs98}, as well as smaller changes for abundant refractory elements, e.g. Si and Fe."
+ One sinale fiere-ofanerit for comparisons ix the photospheric toal uectal to hwdrogen fraction. (Z/N).. of 0.018. for ΑΟ versus 0.023. KxCSIs.," One simple figure-of-merit for comparisons is the photospheric total metal to hydrogen fraction, $(Z/X)_\odot$ of $\sim$ 0.018 for AGSS09 versus $\sim$ 0.023 forGS98."
+ The ACGSSOO results. iowever. have not gone τιichalleuged.," The AGSS09 results, however, have not gone unchallenged."
+ Based on 3D solar atmosphere models computed with the COSBOLD code (Frevtagetal.2002).. Caffauctal.(2010.audreferencesherein) have derived best-value solar CNO abundances hat are svsteimaticallv üeher than those in ACGSSO9 and associated uncertaiuties that are larger. so that the ranges are formally consistent with CS98 values.," Based on 3D solar atmosphere models computed with the CO5BOLD code \citep{freytag:2002}, \citet[and references there
+ in]{caffau:2010} have derived best-value solar CNO abundances that are systematically higher than those in AGSS09 and associated uncertainties that are larger, so that the ranges are formally consistent with GS98 values."
+ The differences. between the Wo groups που to arise not Toni the 3D ντοναιuical atinosphere models but roni the selection of spectral lines aud. assunuptious iu he line formation calculatious.," The differences between the two groups seem to arise not from the 3D hydrodynamical atmosphere models, but from the selection of spectral lines and assumptions in the line formation calculations."
+ Nuclear reaction rates are another muportait input cor SSAIs., Nuclear reaction rates are another important input for SSMs.
+ A very broad effort to critically evauate pp chain aud CN evele rates was undertaken iu 1997-8. culiuinatius in Solar Fuxion I Adelbergereta.(1998. SEI. ," A very broad effort to critically evaluate pp chain and CN cycle rates was undertaken in 1997-8, culminating in Solar Fusion I \citet[hereafter SFI]{sfi}. ."
+Since SEI a ereat deal of acditional experiauental aud theoretical work has been done. e.g. IN(p.33/0 ," Since SFI a great deal of additional experimental and theoretical work has been done, e.g. $^{14}{\rm N(p,\gamma)^{15}O}$ "
+The gamma-ray burst (GRB) 0505254 is a bright. brief M of gamma-ray racliation detected bv the Burst Alert Telescope (BAT) on 25 Alay2005 00:02:53 (UT) 2005).,The gamma-ray burst (GRB) 050525A is a bright brief flash of gamma-ray radiation detected by the Burst Alert Telescope (BAT) on 25 May 2005 00:02:53 (UT) \citep{b05}.
+. It showed two peaks. will a duration of 79)=8.3£0.58 in the 15-350 keV band (Alarkward(etal.2005;Cummings2005)..," It showed two peaks, with a duration of $T_{90}=8.8\pm0.5 \,{\rm s}$ in the 15-350 keV band \citep{m05,c05b}."
+" Filing to the Band spectral model vields a low-energy photon index of a=1.02:0.1 and a peak energy of E,=(9t4 keV 2005)..", Fitting to the Band spectral model yields a low-energy photon index of $\alpha=1.0\pm 0.1$ and a peak energy of $E_p=79\pm4$ keV \citep{c05b}.
+ The fluencein the 15-350 keV bandis (2.02:0.1)xLOeresem7.," The fluence in the 15-350 keV band is $(2.0\pm0.1)\times10^{-5} \,{\rm ergs\,\,\rm cm^{-2}}$."
+" This burst was i"" detected on-board instruments such as iTEGRAL and MUNWind. withpeak fluxes of~alDy3.2xotherLO“eresem?s! (Gotzetal.2005) and ~8.7xI0Pergsem7s1 (Golenetskiieta2005).. respectively."," This burst was also detected by other on-board instruments such as INTEGRAL and Konus-Wind, withpeak fluxes of $\sim 3.2\times10^{-6}\,{\rm
+ergs\,\,cm^{-2}\,\,s^{-1}}$ \citep{g05} and $\sim 8.7\times10^{-6}\,{\rm
+ergs\,\,cm^{-2}\,\,s^{-1}}$ \citep{g05b}, respectively."
+" The fluences were >1.2x10ergsem7 in 20-200 keV (lor the 12 integration (ime) bv INTEGRAL and ~7.8x10ergsem? by ]xonus-Wind MNin20-1000 ""nikeV (for the 11.5 second integration time). respectivelv."," The fluences were $\ge
+1.2\times10^{-5}\,{\rm ergs\,\,cm^{-2}}$ in 20-200 keV (for the 12 second integration time) by INTEGRAL and $\sim 7.8\times10^{-5}\,{\rm
+ergs\,\,cm^{-2}}$ by Konus-Wind in 20-1000 keV (for the 11.5 second integration time), respectively."
+" The a peak energy. £j,=54.1x1.rkeV by Ixonus-Wind 2005).."," The time-integrated spectrum shows a peak energy $E_p =
+84.1 \pm 1.7 {\rm keV}$ by Konus-Wind \citep{g05b}. ."
+ About six minutes later. the ROTSE III telescope in Namibia was able to," About six minutes later, the ROTSE III telescope in Namibia was able to"
+an carly De star in the SAIC.,an early Be star in the SMC.
+ With an N-ray. positional uncertainty of only 10 areseconcs. we conclude that this Be star is the optical counterpart to the X-ray pulsar.," With an X-ray positional uncertainty of only 10 arcseconds, we conclude that this Be star is the optical counterpart to the X-ray pulsar."
+ tN J0051.8-7231 was discovered in observations., RX J0051.8-7231 was discovered in observations.
+ The X-ray error circle included the bright SAIC star AV 11 Figure 1(0)). which was suggested as the optical counterpar Mubweiler LOST. Wang Wu 1992).," The X-ray error circle included the bright SMC star AV 111 Figure 1(d)), which was suggested as the optical counterpart (Bruhweiler 1987, Wang Wu 1992)."
+ The source displays time variability. as demonstrate »v an increase in luminosity by more than a factor of ten in l vear. and bv a factor of ~5 in 5 days (Ixahabka Yjetsch 1996).," The source displays time variability, as demonstrated by an increase in luminosity by more than a factor of ten in $\sim$ 1 year, and by a factor of $\sim$ 5 in 5 days (Kahabka Pietsch 1996)."
+ Phe X-ray luminosity. time-variabilitv aux arch spectrum Leck Ixababka Pietsch to suggest a De/N-ray. binary nature for the SOULCE.," The X-ray luminosity, time-variability and hard spectrum led Kahabka Pietsch to suggest a Be/X-ray binary nature for the source."
+ Israel ct al. (, Israel et al. (
+1995) detected pulsations from the X-ray source. with a period of 8.9 seconds.,"1995) detected pulsations from the X-ray source, with a period of 8.9 seconds."
+ They subsequently obtained optical CCD images of the field (Israel et al 1997). and concluded that both AVILI and a second star (Star 1) within their error circle displaved Lla activity.," They subsequently obtained optical CCD images of the field (Israel et al 1997), and concluded that both AV111 and a second star (Star 1) within their error circle displayed $\alpha$ activity."
+ In fact. both AVIII and Star 1 do not lie within the error circle derived by WKahabka Pictseh. which contains only one object. visible in the DSS image. star 2.," In fact, both AV111 and Star 1 do not lie within the error circle derived by Kahabka Pietsch, which contains only one object visible in the DSS image, star 2."
+ On 1998 February 2 weobtained a medium resolution Lo spectrum of this object. which is shown in Figure (ο).," On 1998 February 2 we obtained a medium resolution $\alpha$ spectrum of this object, which is shown in Figure 3(c)."
+ Ho is in emission with (lla) = -3.0-1.0A., $\alpha$ is in emission with $\alpha$ ) = $\pm$ 1.0.
+. The error circle of Kahabka Pietsch scems to favour the identification of star 2 as the optical counterpart to RX J0051.8-7231. though in the light of the Ha activity detected by. Israel et al. (," The error circle of Kahabka Pietsch seems to favour the identification of star 2 as the optical counterpart to RX J0051.8-7231, though in the light of the $\alpha$ activity detected by Israel et al. ("
+1997) we cannot rule out star 1. whilst ΑΝΤΕ appears a less likely. candidate.,"1997) we cannot rule out star 1, whilst AV111 appears a less likely candidate."
+ This source appears in a number of catalogues based: onEinstein observations of the SAIC (Inoue. Ixovanma Tanaka 1983: Brubweiler et al," This source appears in a number of catalogues based on observations of the SMC (Inoue, Koyama Tanaka 1983; Bruhweiler et al."
+ LOST: Wang Wu 1992). ane was detected by (IXDP1996).," 1987; Wang Wu 1992), and was detected by (KP1996)."
+ Ixahabka and Pietsch (IXD1996) rejected an X-ray. binary nature for the source on the grounds of the lack of time variability (on a timescale of hours), Kahabka and Pietsch (KP1996) rejected an X-ray binary nature for the source on the grounds of the lack of time variability (on a timescale of hours).
+ Observations mace with the 2XTE satellite on 1998 January 20 detected: a 59 second. pulsar. with a positiona uncertainty of £10 areminutes. consistent with the position of RX J0054.9-7226 (Marshall Lochner 1998).," Observations made with the $RXTE$ satellite on 1998 January 20 detected a 59 second pulsar, with a positional uncertainty of $\pm$ 10 arcminutes, consistent with the position of RX J0054.9-7226 (Marshall Lochner 1998)."
+ Subsequent observations with SeLN reduced the positional uncertainty to a radius of 50 areseconds. confirming the identification with RA .0054.9-7226. and refining the pulse period to 58.969 seconds. (Santangelo Cusumano 1998).," Subsequent observations with $SAX$ reduced the positional uncertainty to a radius of 50 arcseconds, confirming the identification with RX J0054.9-7226, and refining the pulse period to 58.969 seconds (Santangelo Cusumano 1998)."
+ The uncertainty in X-ray position was further reduced to a 10 arcsecond radius from the analysis of archival ROSAT data by Israel (1998)., The uncertainty in X-ray position was further reduced to a 10 arcsecond radius from the analysis of archival ROSAT data by Israel (1998).
+ We obtained BVR and Ho images of the field on the night of 1996 October 5., We obtained $BVR$ and $\alpha$ images of the field on the night of 1996 October 5.
+ An emission—colour diagram for the field revealed four objects within the X-ray. error circle. all identified. as carly tvpe stars by their 2V. colours.," An $emission-colour$ diagram for the field revealed four objects within the X-ray error circle, all identified as early type stars by their $B-V$ colours."
+ Of these. only one shows strong Ho emission. with a It-//a value of 0.49 (the object indicated in Figure 1(0)).," Of these, only one shows strong $\alpha$ emission, with a ${H\alpha}$ value of 0.49 (the object indicated in Figure 1(e))."
+ We obtained spectra of Object 1 on the nights of 1998 February 4 and 5., We obtained spectra of Object 1 on the nights of 1998 February 4 and 5.
+ The medium resolution Hla. spectrum obtained is shown in Figure 3(d)., The medium resolution $\alpha$ spectrum obtained is shown in Figure 3(d).
+ Phe Ho line shows strong emission. with an equivalent width of (lla) = -25 and a radial velocity of 137 28km s.+. consistent with SAIC membership.," The $\alpha$ line shows strong emission, with an equivalent width of $\alpha$ ) = -25 $\pm$ 2, and a radial velocity of 137 $\pm$ 28km $^{-1}$, consistent with SMC membership."
+ The low resolution spectrum in Figure + shows 113 also clearly in emission. with (II) = -2.001.0 A.," The low resolution spectrum in Figure \ref{ma810_low}
+ shows $\beta$ also clearly in emission, with $\beta$ ) = $\pm$ 1.0."
+. This A-ray source was discovered inPOSAL pointed observations in. 1991. October., This X-ray source was discovered in pointed observations in 1991 October.
+ Observations approximately 6 months later fled to detect the source. suggesting à transient nature (Ixahabka. Pietsch 1996).," Observations approximately 6 months later failed to detect the source, suggesting a transient nature (Kahabka Pietsch 1996)."
+ IXDP1996 claim that the source is most likely associated with a 15-16th magnitude Be star., KP1996 claim that the source is most likely associated with a 15-16th magnitude Be star.
+ To the authors’ knowledge. no observations of this star have previously been published.," To the authors' knowledge, no observations of this star have previously been published."
+ We only obtained fe: and Ho images of the field of WN J0101.0-7321 on 1995 October 1., We only obtained $R_{C}$ and $\alpha$ images of the field of RX J0101.0-7321 on 1995 October 1.
+ X plot of the measured llo magnitudes against the measured Ae banc magnitudes shows a clear linear relationship exists between Ro and llo: the scatter at magnitudes 1t 7 15.5 is mostly due to uncertainties in the Ho. magnitudes., A plot of the measured $\alpha$ magnitudes against the measured $R_{C}$ -band magnitudes shows a clear linear relationship exists between R and $\alpha$; the scatter at magnitudes R $>$ 15.5 is mostly due to uncertainties in the $\alpha$ magnitudes.
+ Three points show Lla excesses which appear to be much greater than the local scatter of points., Three points show $\alpha$ excesses which appear to be much greater than the local scatter of points.
+ Objects 2 and 3 each lic a low aveminutes from the X-ray. position. whilst the error associated with this," Objects 2 and 3 each lie a few arcminutes from the X-ray position, whilst the error associated with this"
+magnetorotational instabilitv to regenerate radial [field from vertical field. while magnetic reconnection provides dissipation.,"magnetorotational instability to regenerate radial field from vertical field, while magnetic reconnection provides dissipation."
+ In effect. they wrote down a non-linear dynamical svstem for the three field. components. which could be thought of as either mean or RATS values.," In effect, they wrote down a non-linear dynamical system for the three field components, which could be thought of as either mean or RMS values."
+ The present model is ultimately based on a similar idea. although it treats the Maxwell stress covariantlv as à single tensor object.," The present model is ultimately based on a similar idea, although it treats the Maxwell stress covariantly as a single tensor object."
+ Phe Parker instability (magnetic buoyancey) does not play a role here. and the magnetorotational instability is an outcome of the model rather than an ingredient.," The Parker instability (magnetic buoyancy) does not play a role here, and the magnetorotational instability is an outcome of the model rather than an ingredient."
+ In recent work on the cynamics of eccentric discs | questioned the use of a viscous model of the turbulent stress in an accretion disc (Ogilvie 2001)., In recent work on the dynamics of eccentric discs I questioned the use of a viscous model of the turbulent stress in an accretion disc (Ogilvie 2001).
+ Phe assumption that the stress is linearly and instantaneously related to the rate of strain typically leads to an instability whereby an initially circular disc develops eccentricity on short radial scales (see also Lyubarskij. Postnov Prokhorov 1994): in fact. this is essentially the same ellect as the viscous overstability of axisvmametric waves discovered by Ixato (1978).," The assumption that the stress is linearly and instantaneously related to the rate of strain typically leads to an instability whereby an initially circular disc develops eccentricity on short radial length-scales (see also Lyubarskij, Postnov Prokhorov 1994); in fact, this is essentially the same effect as the viscous overstability of axisymmetric waves discovered by Kato (1978)."
+" In order to take account. of the non-zero relaxation time 7 of the turbulence. E proposed the use of a viscoelastie mocel for the Maxwell stress tensor. where fi ds the ellective shear viscosity. ja, the elfective bulk viscosity. and is a time-derivative operator such that DAL; is equal to the left-hand side of equation (43))."," In order to take account of the non-zero relaxation time $\tau$ of the turbulence, I proposed the use of a viscoelastic model for the Maxwell stress tensor, where $\mu$ is the effective shear viscosity, $\mu_{\rm b}$ the effective bulk viscosity, and $\sD$ is a time-derivative operator such that $\sD M_{ij}$ is equal to the left-hand side of equation \ref{compressible_mij}) )."
+ This model is a compressible version. of the upper-convecteck Maxwell Duid of non-Newtonian Luicl dyvnanics. and is motivated by the fact that the Maxwell tensor satisfies the equation DAL;;=0 exactly in ideal MIID.," This model is a compressible version of the upper-convected Maxwell fluid of non-Newtonian fluid dynamics, and is motivated by the fact that the Maxwell tensor satisfies the equation $\sD M_{ij}=0$ exactly in ideal MHD."
+ The analogy between viscoelasticity anc ΑΔΗ) has been developed. in some detail by Ogilvie Proctor (2002). who show that the maenetorotational instability is also found. in viscoclastic Couette flow.," The analogy between viscoelasticity and MHD has been developed in some detail by Ogilvie Proctor (2002), who show that the magnetorotational instability is also found in viscoelastic Couette flow."
+ Phe introduction of the relaxation term. with 7 being comparable to the dynamical time-scale €.7. has a profound influence on the dynamics of eccentric dises and tends to remove the viscous overstability.," The introduction of the relaxation term, with $\tau$ being comparable to the dynamical time-scale $\Omega^{-1}$, has a profound influence on the dynamics of eccentric discs and tends to remove the viscous overstability."
+ The viscoclastic model was intended. to introduce the ellect of the relaxation time in the simplest. possible wav. and &eneralizecs the conventional viscous model by the introduction of a single new parameter. the Weissenberg number O7.," The viscoelastic model was intended to introduce the effect of the relaxation time in the simplest possible way, and generalizes the conventional viscous model by the introduction of a single new parameter, the Weissenberg number $\Omega\tau$."
+ lt is not a true magnetorotational moclel because it neglects the Revnolds stress and relies on an cllective viscosity to generate a non-zero. Maxwell stress., It is not a true magnetorotational model because it neglects the Reynolds stress and relies on an effective viscosity to generate a non-zero Maxwell stress.
+ In contrast. the model proposed. in the present paper does not require any artificial source terms because the stresses arise through an instability of the trivial state and through a cooperative non-linear interaction between the ltevnolds ancl Maxwell stresses.," In contrast, the model proposed in the present paper does not require any artificial source terms because the stresses arise through an instability of the trivial state and through a cooperative non-linear interaction between the Reynolds and Maxwell stresses."
+ “Phe cost. paid. for this is the need for approximateA wice as many terms. variables and. parameters in the model.," The cost paid for this is the need for approximately twice as many terms, variables and parameters in the model."
+" Nevertheless. the viscoelastic model can be compared closely with equation (43)) if one identifies 7P=(6%1Ο1Ο.(ΛΙp? and understands that the C', term. in acting as a source for Adj). is analogous to the effective viscosity."," Nevertheless, the viscoelastic model can be compared closely with equation \ref{compressible_mij}) ) if one identifies $\tau^{-1}=(C_3+C_5)\Omega_z(M/p)^{1/2}$ and understands that the $C_4$ term, in acting as a source for $M_{ij}$, is analogous to the effective viscosity."
+ The behaviour of the present model in the shearing sheet. as set out in the various parts of Section 5.. agrees at least qualitatively with the results of existing numerical simulations.," The behaviour of the present model in the shearing sheet, as set out in the various parts of Section \ref{Non-linear outcome in the
+ shearing sheet}, agrees at least qualitatively with the results of existing numerical simulations."
+ For a non-rotating shear Low without a magnetic field. one may compare with the simulations bv. Pumir (1996). which demonstrate. the development of statistically steady. homogeneous ancl anisotropic turbulence.," For a non-rotating shear flow without a magnetic field, one may compare with the simulations by Pumir (1996), which demonstrate the development of statistically steady, homogeneous and anisotropic turbulence."
+" X. detailed comparison with the mean stresses obtained in numerical simulations. preferably in the limit of horizontally extended shearing boxes L,.LyX»L.. would test the accuracy of the model and constrain the values of the parameters."," A detailed comparison with the mean stresses obtained in numerical simulations, preferably in the limit of horizontally extended shearing boxes $L_x, L_y\gg L_z$, would test the accuracy of the model and constrain the values of the parameters."
+ Rotating shear Lows without a magnetic field have been stuclied numerically by Hawlev. Balbus Winters (1999).," Rotating shear flows without a magnetic field have been studied numerically by Hawley, Balbus Winters (1999)."
+ 'Phey investigated rotation laws of the form ©κrJu inding that Keplerian shear Lows (q=3/2) are hyvdrodsnamically stable while constant-angular momentum shear Lows (q=2) are non-lincarly unstable and become turbulent.," They investigated rotation laws of the form $\Omega\propto r^{-q}$, finding that Keplerian shear flows $q=3/2$ ) are hydrodynamically stable while constant-angular momentum shear flows $q=2$ ) are non-linearly unstable and become turbulent."
+ They. tried to locate the critical. value of q at which hyvdrodynamic turbulence could just be sustained. and found it to be close to 1.95. although this may depend to some extent on the resolution of the numerical method and the initial conditions adopted.," They tried to locate the critical value of $q$ at which hydrodynamic turbulence could just be sustained, and found it to be close to $1.95$, although this may depend to some extent on the resolution of the numerical method and the initial conditions adopted."
+ These results are in qualitative agreement with Fig., These results are in qualitative agreement with Fig.
+ 4. which also suggests that hydrodynamic turbulence can be sustained only in Hows that are BRavleigh-unstable or only just Itavleigh-stable.," 4, which also suggests that hydrodynamic turbulence can be sustained only in flows that are Rayleigh-unstable or only just Rayleigh-stable."
+ The present model prediets the development of steady ALD turbulence in a Weplerian shear Low. much as seen in numerical simulations by Llawley ct al. (," The present model predicts the development of steady MHD turbulence in a Keplerian shear flow, much as seen in numerical simulations by Hawley et al. ("
+1995). although it is more clirectly applicable to caleulations without an imposed magnetic fux. such as those of Brandenburg. et al. (,"1995), although it is more directly applicable to calculations without an imposed magnetic flux, such as those of Brandenburg et al. ("
+1995) and Stone et al. (,1995) and Stone et al. (
+1996).,1996).
+ Phe initial magnetic field in the latter simulations is well ordered. although of zero mean. and. gives rise to an initial phase of exponential erowth that is not captured in the model.," The initial magnetic field in the latter simulations is well ordered, although of zero mean, and gives rise to an initial phase of exponential growth that is not captured in the model."
+ Nevertheless. the properties of the saturated. stresses appear to be in. good qualitative agreement.," Nevertheless, the properties of the saturated stresses appear to be in good qualitative agreement."
+" (X detailed numerical. comparison would be inappropriate. because unfortunately there is no Consensus as to the ""correct, absolute or relative magnitudes of the mean stress components in a Weplerian shear How."," A detailed numerical comparison would be inappropriate, because unfortunately there is no consensus as to the `correct' absolute or relative magnitudes of the mean stress components in a Keplerian shear flow."
+ For example. Brandenburg et al. (," For example, Brandenburg et al. ("
+1995) and Stone et al. (,1995) and Stone et al. (
+1996). although adopting identical initial conditions and numerical resolutions. obtain widely. dillering quantitative results: in particular. they disagree on whether the turbulent kinetic energy or the turbulent magnetic energy is greater.,"1996), although adopting identical initial conditions and numerical resolutions, obtain widely differing quantitative results; in particular, they disagree on whether the turbulent kinetic energy or the turbulent magnetic energy is greater."
+ Lt is unclear whether this discrepancy is attributable to the dillerence in boundary conditions or to a lack of numerical convergenco., It is unclear whether this discrepancy is attributable to the difference in boundary conditions or to a lack of numerical convergence.
+ Abramowicz et al. (, Abramowicz et al. (
+1996) have simulated. the magnetorotational instability. in non-Weplerian shear Hows. having in mind the application to the inner parts— of accretion ¢lises around black holes.,"1996) have simulated the magnetorotational instability in non-Keplerian shear flows, having in mind the application to the inner parts of accretion discs around black holes."
+ They found a non-linear relation between the shear stress ancl the shear rate. at fixed angular velocity. in good agreement with Fig.," They found a non-linear relation between the shear stress and the shear rate, at fixed angular velocity, in good agreement with Fig."
+ 6., 6.
+ Interestingly. the present model also allows for sustained AMLLD turbulence in the case of negative Rosshy number. AO<0. which corresponds to a situation in. which the angular velocity increases outwards.," Interestingly, the present model also allows for sustained MHD turbulence in the case of negative Rossby number, $A/\Omega<0$, which corresponds to a situation in which the angular velocity increases outwards."
+ This regime may appear astrophysically improbable. but may be relevant to the boundary laver between a star ancl a disc., This regime may appear astrophysically improbable but may be relevant to the boundary layer between a star and a disc.
+ While, While
+filaments at too high temperatures compared with the RT simulation.,filaments at too high temperatures compared with the RT simulation.
+ These effects may be quantified using the temperature—density distributions extracted from the simulations., These effects may be quantified using the temperature--density distributions extracted from the simulations.
+" The temperature-density distribution obtained with closely resembles that found using with radiativetransfer at z>3.0, as shown in Fig. 23,,"," The temperature–density distribution obtained with closely resembles that found using with radiativetransfer at $z\ge3.0$, as shown in Fig. \ref{fig:rho_Tdist_multiz_4models},"
+ with reasonably good agreement down to z—2., with reasonably good agreement down to $z=2$.
+" The distributions overlap very closely, although the hydrodynamics computation shows a somewhat tighter and flatter distribution in temperatures, especially towards z—2."," The distributions overlap very closely, although the hydrodynamics computation shows a somewhat tighter and flatter distribution in temperatures, especially towards $z=2$."
+" While the temperatures in the optically thin simulation with boosted hheating rate agree well with the RT simulations for p/(p)>3 over 3.0€z<3.5, the temperatures are appreciably lower in structures with p/(p)«1."," While the temperatures in the optically thin simulation with boosted heating rate agree well with the RT simulations for $\rho/\langle\rho\rangle > 3$ over $3.0\le z\le 3.5$, the temperatures are appreciably lower in structures with $\rho/\langle\rho\rangle <
+1$."
+" At z=2, good agreement is found between the RT and boosted hheating rate simulations for underdense structures, but the latter produces too high temperatures for overdense structures."," At $z=2$, good agreement is found between the RT and boosted heating rate simulations for underdense structures, but the latter produces too high temperatures for overdense structures."
+" The distribution from the optically thin simulation agrees well with the others at z—4, but once rreionization commences, the temperatures are systematically too low."," The distribution from the optically thin simulation agrees well with the others at $z=4$, but once reionization commences, the temperatures are systematically too low."
+" The temperature distributions of the with radiative transfer and the runs are very similar, as shown in Fig. 24,,"," The temperature distributions of the with radiative transfer and the runs are very similar, as shown in Fig. \ref{fig:Tdist_Models},"
+ which may be expected from the overall agreement in the density-temperaturedistrubtions., which may be expected from the overall agreement in the density-temperaturedistrubtions.
+" By z=2.0, however, the simulation systematically underestimates the temperatures by 3—5x10° K. Most of the volume is too cold in the simulation with the boosted hheating rate for 3.0<z 3.5, although a tail of warmer"," By $z=2.0$, however, the simulation systematically underestimates the temperatures by $3-5\times10^3$ K. Most of the volume is too cold in the simulation with the boosted heating rate for $3.0\le z \le3.5$ , although a tail of warmer"
+axisviunetrie briehtuess profile Bi}. the absorption cross section in eq. (1)),"axisymmetric brightness profile $B(\theta')$, the absorption cross section in eq. \ref{eq:absgal}) )"
+ should be convolved with BO’)., should be convolved with $B(\theta')$.
+ As a special case. we observe that if 0 — for all Incoming ravs. Le. k=z. we recover the absorption cross section in eq. (3))," As a special case, we observe that if $\theta'=0$ for all incoming rays, i.e. $\hb{k}=\hb{z}$, we recover the absorption cross section in eq. \ref{eq:absstar}) )"
+ for a poiut-like illuninatius source., for a point-like illuminating source.
+ The cross section for PAT ciission polarized along a direction w is proportional to (2?) where |]? and |} are vibrational states of the cluitting molecule.," The cross section for PAH emission polarized along a direction $\hb{w}$ is proportional to \citep{leger_88}
+ where $|j\rangle$ and $|k\rangle$ are vibrational states of the emitting molecule."
+" We shall call Fl aud Fe the relative enisson cross sections for du-plane anc out-of-plane modes respectively,", We shall call $F_w^\parallel$ and $F_w^\perp$ the relative emission cross sections for in-plane and out-of-plane modes respectively.
+ For out-of-plane modes. the only nou-vauishing terms iu eq. (6))," For out-of-plane modes, the only non-vanishing terms in eq. \ref{eq:emission}) )"
+ involve the component of d perpendicular to the molecule plane. aud For in-plane modes. the matrix elements in eq. (6))," involve the component of $\b{d}$ perpendicular to the molecule plane, and For in-plane modes, the matrix elements in eq. \ref{eq:emission}) )"
+ are uon-zero only for the components of d in the molecule plane: for a rapidlv-piuuiug disk molecule. rotational invariance with respect to the angle of proper rotation o vields By averaging over the precession auele c for the Cluission process Independently from the absorption (as discussed in §2.2)). the emission cross sections for of-plaue modes with polarization along u aud v are where CO.3) has been defined in eq. 0).," are non-zero only for the components of $\b{d}$ in the molecule plane; for a rapidly-spinning disk molecule, rotational invariance with respect to the angle of proper rotation $\phi$ yields By averaging over the precession angle $\psi$ for the emission process independently from the absorption (as discussed in \ref{sec:abs}) ), the emission cross sections for out-of-plane modes with polarization along $\hb{u}$ and $\hb{v}$ are where $C(\theta,\beta)$ has been defined in eq. \ref{eq:absstar}) )."
+ We have made use of the fact that the oricutation of the angular momentum J is coustaut caving the IR cussion burst. as discussed at the begiuniug of 8??..," We have made use of the fact that the orientation of the angular momentum $\bJ$ is constant during the IR emission burst, as discussed at the beginning of \ref{sec:model}."
+ The corresponding euissiou cross sections for in-plane modes with polarizationalong u aud v cau be computed frou their out-of-plane counterparts by using eq. (8)):, The corresponding emission cross sections for in-plane modes with polarizationalong $\hb{u}$ and $\hb{v}$ can be computed from their out-of-plane counterparts by using eq. \ref{eq:em0perp}) ):
+ By averaging over the azimuthal angle νο of the angular 1ionmentun distribution. we simplify eqs. (93)-(10)).," By averaging over the azimuthal angle $\varphi$ of the angular momentum distribution, we simplify eqs. \ref{eq:em1}) \ref{eq:em2}) )"
+ aud obtain and the corresponding y-averaged cross sectious for plane modes may then be derived from eq. (8))., and obtain and the corresponding $\varphi$ -averaged cross sections for in-plane modes may then be derived from eq. \ref{eq:em0perp}) ).
+ Iun the following. we assune that the augular moneutuni J is randomly oriented m space. but im principle eqs. (9)3- (103) ," In the following, we assume that the angular momentum $\bJ$ is randomly oriented in space, but in principle eqs. \ref{eq:em1}) \ref{eq:em2}) )"
+and eqs. (113)-(12)), and eqs. \ref{eq:em3}) \ref{eq:em4}) )
+" could be convolved with aux aneular distribution for J. οσοι, if the emitting molecules are partially aligned by the interstellar maeuetic field."," could be convolved with any angular distribution for $\bJ$ e.g., if the emitting molecules are partially aligned by the interstellar magnetic field."
+" The probabilitv that a sviuuetrie erain rotates with angular momentun ο= |J| aud nutation angle 3 (between aj and J) may be written (e.g.7) NI .(15hwhere fG/ is the probability that the angular monmentun C|J.| dJ|. Ej; is the grain rotational cucrev and Ti, is the “internal alignment?"," The probability that a symmetric grain rotates with angular momentum $J\equiv|\bJ|$ and nutation angle $\beta$ (between $\ba$ and $\bJ$ ) may be written \citep[e.g.,][]{lazarian&draine_97}
+ J ,where $f(J)\ud J$ is the probability that the angular momentum $\in [J,J+\ud J]$ , $E_{\rm rot}$ is the grain rotational energy and $\Tia$ is the “internal alignment”"
+functions computed for transverse velocities aligned with the C-D ancl A-B axes respectively.,functions computed for transverse velocities aligned with the C-D and A-B axes respectively.
+ Fig., Fig.
+ 11. also contains a kev corresponding to the intervals between event types shown in Fig. 10.., \ref{caust_prob} also contains a key corresponding to the intervals between event types shown in Fig. \ref{scenarios_real}.
+ These intervals are defined below., These intervals are defined below.
+ Firstly. we assume that the 1999 peak was à /—(ee caustic Crossing.," Firstly, we assume that the 1999 peak was a $-ve$ caustic crossing."
+ 1) interval 1-2: The dashed. lines show cumulative probabilities for the separation of two adjacent caustic crossing LIALEs where the first is a |re caustic crossing., $i)$ $interval$ 1-2: The dashed lines show cumulative probabilities for the separation of two adjacent caustic crossing HMEs where the first is a $+ve$ caustic crossing.
+ Secondly. the 1999. peak is interpreted. as a cusp event Following a |re caustic crossing.," Secondly, the 1999 peak is interpreted as a cusp event following a $+ve$ caustic crossing."
+ H) interval 1-2: Phe dot-dashed line shows the distribution of caustic separations in à double horned event where a cusp lies between the two caustics., $ii)$ $interval$ 1-3: The dot-dashed line shows the distribution of caustic separations in a double horned event where a cusp lies between the two caustics.
+ In this case the typical separation of caustics is larger due to the fact that the cusp is generally formed. from independent. caustics. so that the chance of a cusp Lying inside a diamond or another cusp is higher if the separation (and therefore area) is larger.," In this case the typical separation of caustics is larger due to the fact that the cusp is generally formed from independent caustics, so that the chance of a cusp lying inside a diamond or another cusp is higher if the separation (and therefore area) is larger."
+ Hi) interval 1-4: The dotted lines represent the probability for the separation between a cusp and the (immediately) preceding |re caustic crossing., $iii)$ $interval$ 1-4: The dotted lines represent the probability for the separation between a cusp and the (immediately) preceding $+ve$ caustic crossing.
+ iv) interval 4-3: Phe solid lines represent the probability for the separation between a cusp event (that has followed a |re caustic crossing) ancl the subsequent caustic crossing., $iv)$ $interval$ 4-3: The solid lines represent the probability for the separation between a cusp event (that has followed a $+ve$ caustic crossing) and the subsequent caustic crossing.
+ As expected 7/7) and £e) are very similar., As expected $iii)$ and $iv)$ are very similar.
+ The separation of an inferred. |ee caustic crossing between the 1997 and 1998 seasons and the 1999 image € peak is consistent with Fig., The separation of an inferred $+ve$ caustic crossing between the 1997 and 1998 seasons and the 1999 image C peak is consistent with Fig.
+ TL whether the 1999 peak is interpreted as α re caustic crossing or a cusp event., \ref{caust_prob} whether the 1999 peak is interpreted as a $-ve$ caustic crossing or a cusp event.
+ Llowever if the 1999 peak is interpreted as a cusp. we expect another caustic crossing o follow.," However if the 1999 peak is interpreted as a cusp, we expect another caustic crossing to follow."
+ Fig., Fig.
+ Ll shows that in this case. the separation otween the cusp peak and the subsequent caustic is likely ο be z5 vears (although may be significantly longer).," \ref{caust_prob} shows that in this case, the separation between the cusp peak and the subsequent caustic is likely to be $\ga5$ years (although may be significantly longer)."
+ These statisties for the separation of caustic crossings hat surround a cusp does not make use of the information hat the observed cusp event followed the caustic crossing by 500 davs., These statistics for the separation of caustic crossings that surround a cusp does not make use of the information that the observed cusp event followed the caustic crossing by $\sim 500$ days.
+ An alternative way of analysing the separation oween the cusp event peak and the next caustic crossing is to calculate the ratio of the time between the cusp event rcs ancl last caustic crossing and the time between the cusp event peak and next caustic crossing., An alternative way of analysing the separation between the cusp event peak and the next caustic crossing is to calculate the ratio of the time between the cusp event peak and last caustic crossing and the time between the cusp event peak and next caustic crossing.
+ Phe probability of this ratio is plotted in Fig. 12.., The probability of this ratio is plotted in Fig. \ref{caust_ratio}.
+ Phe typical ratio is 1 and we expect it to be between 4+I and ~4. vielding a most likely arrival time for the next event of ~500 clays. and an upper limit of ~2000 days (~90%).," The typical ratio is 1 and we expect it to be between $\sim \frac{1}{4}$ and $\sim4$, yielding a most likely arrival time for the next event of $\sim 500$ days, and an upper limit of $\sim 2000$ days $\sim90\%$ )."
+ Also plotted in Fig., Also plotted in Fig.
+ 19 is the probability of the ratio for a random position (dashed lines)., \ref{caust_ratio} is the probability of the ratio for a random position (dashed lines).
+ The agreement of the curves illustrates the independence of the positions of the cusps and unassociated CAUSLICS., The agreement of the curves illustrates the independence of the positions of the cusps and unassociated caustics.
+ In adclition to the remarkable peak observed during L999 in the image € light-curve. the OGLE data also shows image A brightening over the entire season. with the most rapid variations occurring in the latter observations.," In addition to the remarkable peak observed during 1999 in the image C light-curve, the OGLE data also shows image A brightening over the entire season, with the most rapid variations occurring in the latter observations."
+ We have applied the triggering function (as described in Sec. 4.1.3)), We have applied the triggering function (as described in Sec. \ref{trigger_Im_C}) )
+ to the image A lieht-curve., to the image A light-curve.
+ On the 30th of October 1999 monitoring by OGLE (OGLE web page) showed rises in image A. of 181-155 maes/vear (we used observations on the 20th of October. 30th of October and 9th of November).," On the 30th of October 1999 monitoring by OGLE (OGLE web page) showed rises in image A, of 1.41-1.88 mags/year (we used observations on the 20th of October, 30th of October and 9th of November)."
+ Fig., Fig.
+ 13. shows the resulting triggering functions 7. £ and fee.," \ref{obs_trigger_A} shows the resulting triggering functions $F_{+}$, $F_{-}$ and $F_{C}$."
+ Functions are shown assuming that the source size estimate 5 is correct., Functions are shown assuming that the source size estimate $S$ is correct.
+" The solid. dotted: ancl dashed. lines correspond to ""n(PHEAM jdPLNΕλ0. L5GPsFATos) respectively."," The solid, dotted and dashed lines correspond to $F_{C}(P|T_{obs}\pm\Delta T_{obs})$, $F_{-}(P|T_{obs}\pm\Delta T_{obs})$, $F_{+}(P|T_{obs}\pm\Delta T_{obs})$ respectively."
+ We find that the results are similar to those obtained for the June image € trigger., We find that the results are similar to those obtained for the June image C trigger.
+ The observed trigger precedes a caustic crossing LIAL ~so% of the time., The observed trigger precedes a caustic crossing HME $\sim 80\%$ of the time.
+ The event peak is most likely to occur 13 months following such a trigger if the event was a cusp. ~2 months later if it was ao οὐ caustic and ~1 month later if it was a |ee caustic crossing.," The event peak is most likely to occur $\sim 1-3$ months following such a trigger if the event was a cusp, $\sim 2$ months later if it was a $-ve$ caustic and $\sim 1$ month later if it was a $+ve$ caustic crossing."
+ Having observed a derivative in the quoted range means that a |re caustic crossing is very unlikely anca οὐ caustic crossing is the most likely option.," Having observed a derivative in the quoted range means that a $+ve$ caustic crossing is very unlikely, and a $-ve$ caustic crossing is the most likely option."
+ Unfortunately. these results predict that an event will occur in image A between the 1999 and 2000 observing seasons.," Unfortunately, these results predict that an event will occur in image A between the 1999 and 2000 observing seasons."
+ Lf the impending event is assumed to be a οὐ caustic crossing.," If the impending event is assumed to be a $-ve$ caustic crossing,"
+The values of f and g atv =Av. the inner boundary ofour integration. aro g=2o wd f2=τοπAq.3 where ois a parameter which specifies the relative aout of mmass loading in the solution.,"The values of $f$ and $g$ at $x=\Delta x$, the inner boundary ofour integration, are $g = \sqrt{2/\phi}$ and $f = 1/(2 \pi \, \Delta x^{2} g^{3})$, where $\phi$ is a parameter which specifies the relative amount of mass loading in the solution."
+ / is sot to zero at this out., $h$ is set to zero at this point.
+ Equations δ and 9 are inteerated to the radius i which mass loading begius., Equations \ref{eq:ode1_ab_fe} and \ref{eq:ode2_ab_fe} are integrated to the radius at which mass loading begins.
+ From this poiut onwards we then inteerate Equations 10--12 instead., From this point onwards we then integrate Equations \ref{eq:ode1_ab}- \ref{eq:ode3_ab} instead.
+ The Mach iuuber of the wind rapidly decreases fron its initial value of infinity aud both the deusitv and velocity of the wind respond to the mass addition., The Mach number of the wind rapidly decreases from its initial value of infinity and both the density and velocity of the wind respond to the mass addition.
+ Iutegratiou proceeds uutil he contact discoutinuity (CD) is reached., Integration proceeds until the contact discontinuity (CD) is reached.
+ This occurs when the flow velocity is equal to the coordinate velocity c= dRfdt). which is where goer—04)=2c/(5|A).," This occurs when the flow velocity is equal to the coordinate velocity $v=dR/dt$ ), which is where $g(x=x_{cd}) = 2x/(5+\lambda)$."
+" As previously discussed. this ean sometimes occur the specified position of the iuuer shock (= .0;,) is reached. ia which case the mass-loaded flow connects directly to the CD."," As previously discussed, this can sometimes occur the specified position of the inner shock $x = x_{is}$ ) is reached, in which case the mass-loaded flow connects directly to the CD."
+ Tf an iuncr shock is prescut. its velocity with respect to the bubble ceuter is gs=(5|A). aud the Mach umber of the preshock flow with respect to the shock is where the subscript indicates the preshock values.," If an inner shock is present, its velocity with respect to the bubble center is $g_{s} = 2x_{is}/(5+\lambda)$, and the Mach number of the preshock flow with respect to the shock is where the subscript `1' indicates the preshock values."
+" Because (unlike for the evaporative case) AM, is nof necessarily large. we camnot assiune that the strong shock juup conditions apply aud therefore must use the full Bonkine-Iusoniot relations."," Because (unlike for the evaporative case) $M_{1}$ is not necessarily large, we cannot assume that the strong shock jump conditions apply and therefore must use the full Rankine-Hugoniot relations."
+ The postshock values are then: Even if the flow velocity has not couvereed to the coordinate velocity. an mner shock may not exist.," The postshock values are then: Even if the flow velocity has not converged to the coordinate velocity, an inner shock may not exist."
+ This can occur if the mass loading interior to the proposed position of the inner shock is strong. resulting in A4 being less than unity.," This can occur if the mass loading interior to the proposed position of the inner shock is strong, resulting in $M_{1}$ being less than unity."
+ Iu such a scenario. the jump conditions are uot applied aud the integration proceeds as before.," In such a scenario, the jump conditions are not applied and the integration proceeds as before."
+ At the CD. we aust also satisfy conservation of momentum.," At the CD, we must also satisfy conservation of momentum."
+ We take as a lueasure of the ablated mass to the mass in the swept- shell., We take as a measure of the ablated mass to the mass in the swept-up shell.
+ It follows that at the contact discoutiunuity. The rate of mmassloss frou the star is where which differs from the definition iu the coucuctively driven case.," It follows that at the contact discontinuity, The rate of massloss from the star is where which differs from the definition in the conductively driven case."
+ The mass loading parameter o. which is à measure of the ratio of ablated mass to wind mass is given by ο= M/E.," The mass loading parameter $\phi$, which is a measure of the ratio of ablated mass to wind mass is given by $\phi = 
+\dot{M}_{c}/\Xi$ ."
+ The dimensionless Eqs. 10-, The dimensionless Eqs. \ref{eq:ode1_ab}-
+"19 απο invariant τσ the following trausfoxiiation. which we shall call a normalization: These can be combined to obtain the normalizatious for the mass. and for the kinetic and internal energies of the bubble: The full equation for the mass in the bubble is: The total mass lost by0 the star curing its lifetime is The deeree of 1iass loading iu the bubble. ὃν, is defined as the ratio Ay,ΕΕ "," \ref{eq:ode3_ab} are invariant under the following transformation, which we shall call a normalization: These can be combined to obtain the normalizations for the mass, and for the kinetic and internal energies of the bubble: The full equation for the mass in the bubble is: The total mass lost by the star during its lifetime is The degree of mass loading in the bubble, $\Phi_{b}$, is defined as the ratio $M_{b}/M_{W}$ ."
+"The miass-swept up iuto the shell. A. the ratio of swept-up mass to bubble mass. Αλί, the kinetic energv of the shell. AE. aud the pd work ou the shocked ambient gas. are all identical to the correspouding forms in PDII."," The mass-swept up into the shell, $M_{sh}$, the ratio of swept-up mass to bubble mass, $M_{sh}/M_{b}$, the kinetic energy of the shell, $KE_{sh}$, and the $pdV$ work on the shocked ambient gas, are all identical to the corresponding forms in PDH."
+ Tf an inner shock exists. three main parameters determine the form of a given similarity solution (A. o and wie £c).," If an inner shock exists, three main parameters determine the form of a given similarity solution $\lambda$, $\phi$ and $x_{is}/x_{cd}$ )."
+ However. the ratio of the radius of the onset of mass loading to the radius of the iuner shock is au additional parameter unique to the ablative solutions. aud its value Go) nav also influence the form of the solution.," However, the ratio of the radius of the onset of mass loading to the radius of the inner shock is an additional parameter unique to the ablative solutions, and its value $x_{ml}$ ) may also influence the form of the solution."
+ We fiud that if the mass loading of the bubble is minor. ο] has relatively little infiuence. and vice-versa (since pxà\ and BSAS2 for nost of the solutious presented. most of the mass-loadingo occurs close to the muininnuu niass-loading radius).," We find that if the mass loading of the bubble is minor, $x_{ml}$ has relatively little influence, and vice-versa (since $\dot{\rho} \propto r^{\lambda}$ and $-3 \ltsimm \lambda \ltsimm -2$ for most of the solutions presented, most of the mass-loading occurs close to the minimum mass-loading radius)."
+" Alternatively. if an iuuer shock does not exist. cjufru aud eu, Lave no moeauimg as we have defined them."," Alternatively, if an inner shock does not exist, $x_{is}/x_{cd}$ and $x_{ml}$ have no meaning as we have defined them."
+" Iu this case we report the ratio of the onset radius of mass loading to the radius of the coutact discoutinuitv as the parameter 0,4.", In this case we report the ratio of the onset radius of mass loading to the radius of the contact discontinuity as the parameter $x_{ml}$ .
+ The ratio of mass pick-up frou the clumps to the mass swept-up from the iuterclunip mecdinm(which is related to the value of 0) is obtained only after a particular solution has been found., The ratio of mass pick-up from the clumps to the mass swept-up from the interclump medium(which is related to the value of $\theta$ ) is obtained only after a particular solution has been found.
+ For bubbles whose evolution, For bubbles whose evolution
+"The concentration of a galaxy is determined by locating the radii of the circular apertures that contain (10,2) and (ros) of the total Fux of the galaxy. ancl then taking their ratio: The asvometry parameter (A) is a measure of how svmametric (or asymmetric) a galaxy is with respect to the total Hus of the galaxy.","The concentration of a galaxy is determined by locating the radii of the circular apertures that contain $r_{0.2}$ ) and $r_{0.8}$ ) of the total flux of the galaxy, and then taking their ratio: The asymmetry parameter (A) is a measure of how symmetric (or asymmetric) a galaxy is with respect to the total flux of the galaxy."
+ Galactic asvnimetsrics are predominantly caused by star formation. mainlv in uncderlving structure such as spiral arms.," Galactic asymmetries are predominantly caused by star formation, mainly in underlying structure such as spiral arms."
+ Asvnunetry is also a strong indicator of merger history. as demonstrated by 7..," Asymmetry is also a strong indicator of merger history, as demonstrated by \citet{Conselice2003}."
+ Phe asymmetry. parameter initially involves locating the centre of a galaxy image. which is defined as the elobal minimum of asvmmetry.," The asymmetry parameter initially involves locating the centre of a galaxy image, which is defined as the global minimum of asymmetry."
+ The centre was obtained using the centering algorithm specilied by 2.. which involves an iterative trial-and-error process.," The centre was obtained using the centering algorithm specified by \citet{Conselice2000b}, which involves an iterative trial-and-error process."
+ An initial guess for the centre of the galaxy is mace and the asvnimectry for the specified: pixel. and the surrounding eight pixels is determined.," An initial guess for the centre of the galaxy is made and the asymmetry for the specified pixel, and the surrounding eight pixels is determined."
+ The pixel corresponding to the lowest asvmametry is then redefined to be the centre., The pixel corresponding to the lowest asymmetry is then redefined to be the centre.
+ In order to locate the global asymmetry. minimum. which corresponds to the centre. it is required that the process be repeated until the central pixel is also the pixel with the lowest value of asymmetry.," In order to locate the global asymmetry minimum, which corresponds to the centre, it is required that the process be repeated until the central pixel is also the pixel with the lowest value of asymmetry."
+ The asymmetry itself is calculated by rotating the image 180 about the defined. central point and. subtracting the rotated image from the initial image. in order to determine the absolute value of the residuals.," The asymmetry itself is calculated by rotating the image $\degrees$ about the defined central point and subtracting the rotated image from the initial image, in order to determine the absolute value of the residuals."
+ Following this. the resicluals are divided by the quantity corresponding to the summation of all the pixels in the galaxy.," Following this, the residuals are divided by the quantity corresponding to the summation of all the pixels in the galaxy."
+ Subsequently. a similar method. is then applied to a small section. taken from a corner of the image which contains background only.," Subsequently, a similar method is then applied to a small section, taken from a corner of the image which contains background only."
+ This quantity is then multiplied. by a scale. factor so that it is comparable to the size of the galaxy image., This quantity is then multiplied by a scale factor so that it is comparable to the size of the galaxy image.
+ This quantity is divided by the summation of the total llux of the original galaxy image., This quantity is divided by the summation of the total flux of the original galaxy image.
+ This is done in order to determine the asvmmetry of the background so that it may be discounted from the galactic asymmetry., This is done in order to determine the asymmetry of the background so that it may be discounted from the galactic asymmetry.
+ The calculation for the asymmetry is as follows: where is the image Iux. is the rotated image ux. Bis the background.5 and is the rotated background.," The calculation for the asymmetry is as follows: where is the image flux, is the rotated image flux, is the background, and is the rotated background."
+d The measure of high-spatial frequency clumpiness (8) is directlv related to a galaxys star formation anc high frequeney spatial power., The measure of high-spatial frequency clumpiness (S) is directly related to a galaxy's star formation and high frequency spatial power.
+ The clumpiness is defined as the patchiness of the high frequency light. distribution within a galaxy (7).., The clumpiness is defined as the patchiness of the high frequency light distribution within a galaxy \citep{Conselice2003}.
+ It is computed by subtracting a degracded-resolution version of the galaxv image from the original galaxy image., It is computed by subtracting a degraded-resolution version of the galaxy image from the original galaxy image.
+ As with the asvmimetry parameter the method is repeated on a section of skv background in order to isolate the fraction of the signal that is clue to the galaxy., As with the asymmetry parameter the method is repeated on a section of sky background in order to isolate the fraction of the signal that is due to the galaxy.
+ Again. a corner of the image containing background only is selected. replicated. ancl smoothed. so that an identical. procedure may be carried out on both the background and the galaxy image.," Again, a corner of the image containing background only is selected, replicated, and smoothed, so that an identical procedure may be carried out on both the background and the galaxy image."
+ As the central bulge component is a prominent [eature in the simulated. galaxies. it contributes to a [large portion of the high frequency. structure and hence. clumpiness of the simulated galaxies.," As the central bulge component is a prominent feature in the simulated galaxies, it contributes to a large portion of the high frequency structure and hence clumpiness of the simulated galaxies."
+ As it is already a known issue that simulated: galaxies tend. to have large. central bulge components. with respect to their observed. counterparts. it owas decided that the standard. central region specilied bv Conselice (2003) would be extended to include the total central bulge component of cach galaxy.," As it is already a known issue that simulated galaxies tend to have large central bulge components, with respect to their observed counterparts, it was decided that the standard central region specified by Conselice (2003) would be extended to include the total central bulge component of each galaxy."
+ Furthermore. the inferred. total lux. of the galaxies would also exclude the central region. in order to avoid any bias due to the varying sizes of the bulge components of the galaxies.," Furthermore, the inferred total flux of the galaxies would also exclude the central region, in order to avoid any bias due to the varying sizes of the bulge components of the galaxies."
+ The size of the central bulee component was determined from the light wolile of the oogalaxy byBH consideringIn the radius at which the profile changed from Πα} = 1/4. to exponential decay.," The size of the central bulge component was determined from the light profile of the galaxy by considering the radius at which the profile changed from $I$ ) ${=}$ $^{1/4}$, to exponential decay."
+ The image was then examined to ensure that the correct area had been removed. prior to the computation of the elumpiness parameter.," The image was then examined to ensure that the correct area had been removed, prior to the computation of the clumpiness parameter."
+ The caleulation for the elumpiness is as follows: where {is the image lux. is the smoothed. image lux.D ids the background. and is the smoothed background.," The calculation for the clumpiness is as follows: where is the image flux, is the smoothed image flux, is the background, and is the smoothed background."
+ Further details on the derivation of. the clumpiness parameter are provided in Appendix A. The CAS parameters form a well-established: svstem tha has been used on occasion to analvse galaxy simulations. although not in any systematic manner.," Further details on the derivation of the clumpiness parameter are provided in Appendix A. The CAS parameters form a well-established system that has been used on occasion to analyse galaxy simulations, although not in any systematic manner."
+ For example. (?7) applied the svsteni to dark. matter. merger-remnan simulations. and (?7)— extended: this to the analysis of equal-mass. eas-rich disc. merger remnants.," For example, \citep{Conselice2006} applied the system to dark matter merger-remnant simulations, and \citep{Lotz2008} extended this to the analysis of equal-mass gas-rich disc merger remnants."
+ Our. unique contribution to this area is. for the first time. an application to à suite of simulations realised within a fully cosmologica framework.," Our unique contribution to this area is, for the first time, an application to a suite of simulations realised within a fully cosmological framework."
+ Such a comparison should. enable. underlving physical. dillerences to. be revealed: anc allow for scaling relations to be tested using simulations in a new aux unique manner., Such a comparison should enable underlying physical differences to be revealed and allow for scaling relations to be tested using simulations in a new and unique manner.
+ The main purpose of this paper is to consider the morphological details of the galaxies produce in simulations. and uncover new information regarding the cillerenees between observed. and. simulated: galaxies.," The main purpose of this paper is to consider the morphological details of the galaxies produced in simulations, and uncover new information regarding the differences between observed and simulated galaxies."
+ Το avoid simply re-identifving known issues and biases. we have mocified slightly the underlying CAS parameters. as explained now.," To avoid simply re-identifying known issues and biases, we have modified slightly the underlying CAS parameters, as explained now."
+ The overproduction. of the centrally-concentratecl bulge component of most simulated. disc. galaxies (e.g... Stinson et al.," The overproduction of the centrally-concentrated bulge component of most simulated disc galaxies (e.g., Stinson et al."
+ 2010. and references. therein: ος.," 2010, and references therein; c.f.,"
+ Governato et al., Governato et al.
+ 2010). is à long-standing problem: associated: with cosmological simulations of galaxy formation.," 2010), is a long-standing problem associated with cosmological simulations of galaxy formation."
+ By definition. the concentration parameter is à proxy for the dominance of the central region. simulation or not. although consideration of the surface brightness profiles suggested. that the basic," By definition, the concentration parameter is a proxy for the dominance of the central region, simulation or not, although consideration of the surface brightness profiles suggested that the basic"
+cluster mass within elliptical galaxies (e.g... AIST. Ixundu et al.,"cluster mass within elliptical galaxies (e.g., M87, Kundu et al."
+ 1999)., 1999).
+ Conversely. Vesperini (2000) emphasized that. in the case of à gaussian initial mass function. the racial &eraciient of the average cluster mass is weak ancl consistent with the observations.," Conversely, Vesperini (2000) emphasized that, in the case of a gaussian initial mass function, the radial gradient of the average cluster mass is weak and consistent with the observations."
+ Equivalently. the evolved: mass ane number density profiles match cach other. as shown by the top panels of Fig.," Equivalently, the evolved mass and number density profiles match each other, as shown by the top panels of Fig."
+ 2., 2.
+ The marked. evolution of the number density. profile with respect to the mass density. profile in the case of a power-law mass spectrum arises because. as Vesperini (1998. 2000. 2001). we have assumed that clusters are orbiting at constant galactocentric distance.," The marked evolution of the number density profile with respect to the mass density profile in the case of a power-law mass spectrum arises because, as Vesperini (1998, 2000, 2001), we have assumed that clusters are orbiting at constant galactocentric distance."
+ Should a substantial racial mixing take place. a more racially uniform mass function may emerge.," Should a substantial radial mixing take place, a more radially uniform mass function may emerge."
+ We have thus tested whether our results are significantly allectecl if the orbital eccentricity is e=0.5., We have thus tested whether our results are significantly affected if the orbital eccentricity is $e=0.5$.
+ According to Baumegardt Makino (2003).J) the lifetime of a cluster on an orbit with eccentricitv e is decreased by a actor (1.0) with respect toa cluster on a circular orbit with a radius similar to the apogalactic radius of the eccentric orbit.," According to Baumgardt Makino (2003), the lifetime of a cluster on an orbit with eccentricity $e$ is decreased by a factor $(1-e)$ with respect to a cluster on a circular orbit with a radius similar to the apogalactic radius of the eccentric orbit."
+ Considering the two extreme (and kev) cases. namely. he gaussian C] ancl power-law PE] initial cluster mass clistributions. combined with a coreless D.77 initial number density profile. we have run additional simulations in which he quantity £4. (equation 3)) is halved (νου. multiplied we le).," Considering the two extreme (and key) cases, namely, the gaussian [G] and power-law [PL3] initial cluster mass distributions, combined with a coreless $D^{-3.5}$ initial number density profile, we have run additional simulations in which the quantity $F_{cw}$ (equation \ref{eq:Fcw}) ) is halved (i.e., multiplied by $1-e$ )."
+ The corresponding evolved: mass and number density. profiles are in remarkable agreement with those derived. under the assumption of circular orbits., The corresponding evolved mass and number density profiles are in remarkable agreement with those derived under the assumption of circular orbits.
+ Phe only significant dillerence is the destruction of all the clusters confined within kkpe from the Galactic centre if e=0.5., The only significant difference is the destruction of all the clusters confined within kpc from the Galactic centre if $e=0.5$.
+ The comparison between the predicted profiles and the Old Lalo clistributions leads to (considering AfogD= 0.2): As for the gaussian mass function. the incomplete eamma function is smaller than that derived. for. circular orbits (lable 4).," The comparison between the predicted profiles and the Old Halo distributions leads to (considering $\Delta logD
+= 0.2$ ): As for the gaussian mass function, the incomplete gamma function is smaller than that derived for circular orbits (Table 4)."
+ Yet. this effect is mostly driven by the first bin. this being located on the edge of the region of complete cluster destruction (Le. D x 1.2kpc).," Yet, this effect is mostly driven by the first bin, this being located on the edge of the region of complete cluster destruction (i.e., D $\simeq$ 1.2kpc)."
+ We note that the power-law initial mass spectrum PL3] is again rejected by the poor agreement between the predicted. aud observed. number density profiles. (that is. the increase in the mean cluster. mass with decreasing. galactocentric distance is much larger than is observed).," We note that the power-law initial mass spectrum [PL3] is again rejected by the poor agreement between the predicted and observed number density profiles (that is, the increase in the mean cluster mass with decreasing galactocentric distance is much larger than is observed)."
+ This result. is not unexpected., This result is not unexpected.
+ The evaporation rate of a cluster with a given mass depends on its orbit. especially its pericentre (sce e.&.. Baumeardt 1998).," The evaporation rate of a cluster with a given mass depends on its orbit, especially its pericentre (see e.g., Baumgardt 1998)."
+ Considering a svstem of clusters extending up to 40kkpe from the Galactic centre. if e-0.5. the range of perigalactic distances is c 1-13.3 kpe.," Considering a system of clusters extending up to kpc from the Galactic centre, if e=0.5, the range of perigalactic distances is $\simeq$ 1-13.3 kpc."
+ The lack of consistency. between the slopes of the evolved. mass and number density profiles in the PL3] case illustrates that such a range of perigalactie distances is still too large to erase radial variations in the mean cluster mass., The lack of consistency between the slopes of the evolved mass and number density profiles in the [PL3] case illustrates that such a range of perigalactic distances is still too large to erase radial variations in the mean cluster mass.
+ As demonstrated by Pall Zhang (2001). a narrow clistribution of pericentres and. thus. disruption rates almost independent of the cluster galactocentric distance. can be achieved if the initial velocity distribution. shows some radial. anisotropyand the radial anisotropy is increasing outward.," As demonstrated by Fall Zhang (2001), a narrow distribution of pericentres and, thus, disruption rates almost independent of the cluster galactocentric distance, can be achieved if the initial velocity distribution shows some radial anisotropy the radial anisotropy is increasing outward."
+ Should such conditions be satisfied. the mean cluster masses. in the inner (2<Skkpe) and outer (27Skkpe) groups of clusters are similar.," Should such conditions be satisfied, the mean cluster masses in the inner $D < 5$ kpc) and outer $D >5$ kpc) groups of clusters are similar."
+" However. that result is obtained in the case of an initial Schechter mass spectrum (ie. dNdimxme""mcamM. ) Le. a mass distribution steeper than a power-law a=2."," However, that result is obtained in the case of an initial Schechter mass spectrum (i.e., $dN/dm \propto m^{-2} e^{-m/5 \times
+10^6\,M_{\odot}}$ ), i.e., a mass distribution steeper than a power-law $\alpha = -2$."
+ In the case of an initial power-law mass spectrum. the evolved. mass distribution fails to reproduce the high mass-regime of the present-day distribution (Fall Zhang 2001. their Fig.," In the case of an initial power-law mass spectrum, the evolved mass distribution fails to reproduce the high mass-regime of the present-day distribution (Fall Zhang 2001, their Fig."
+ 3)., 3).
+ A racially uniform cluster mass function can thus be achieved even if the initial mass distribution is steadily rising toward low-mass. although such a solution requires appropriate tuning.," A radially uniform cluster mass function can thus be achieved even if the initial mass distribution is steadily rising toward low-mass, although such a solution requires appropriate tuning."
+ The initial distribution in mass of the Galactic halo elobular clusters has so far remained a poorlv-known function., The initial distribution in mass of the Galactic halo globular clusters has so far remained a poorly-known function.
+ This is due to the fact that both à gaussian initial mass function (i.c.. à two-index. power-law mass spectrum) ancl a power-law initial mass spectrum evolve into the presently observed bell-shaped cluster mass function.," This is due to the fact that both a gaussian initial mass function (i.e., a two-index power-law mass spectrum) and a power-law initial mass spectrum evolve into the presently observed bell-shaped cluster mass function."
+ As a result. the study of the temporal evolution of the mass function/spectrum only does not enable one to assess how large was the contribution of low-mass (say a few thousand solar masses) objects to the initial population of clusters.," As a result, the study of the temporal evolution of the mass function/spectrum only does not enable one to assess how large was the contribution of low-mass (say a few thousand solar masses) objects to the initial population of clusters."
+ In this paper. we have proposed a new method for shedding light on this issue which consists in comparing globular cluster system density profiles with their modelled counterparts. as a function of mass density as well as a function of number density.," In this paper, we have proposed a new method for shedding light on this issue which consists in comparing globular cluster system density profiles with their modelled counterparts, as a function of mass density as well as a function of number density."
+ We assume. as in previous studies of globular cluster svstem dynamical evolution. that the cluster mass range and the cluster mass spectrum are initially independent. of their ealactocentric distance.," We assume, as in previous studies of globular cluster system dynamical evolution, that the cluster mass range and the cluster mass spectrum are initially independent of their galactocentric distance."
+ This is equivalent to assuming that the initial mass and number density. profiles are identical in shape., This is equivalent to assuming that the initial mass and number density profiles are identical in shape.
+ On the other hand. Ίο present mass and number density profiles of the Old llalo are alike as well (see Section 2).," On the other hand, the present mass and number density profiles of the Old Halo are alike as well (see Section 2)."
+ The new constraints on the initial elobular cluster mass function we have derived arise [rom combining this widespread. assumption and this observational fact., The new constraints on the initial globular cluster mass function we have derived arise from combining this widespread assumption and this observational fact.
+ The clusters most. vulnerable (ο evaporation and disruption are the low-mass ones as well as those located at small galactocentric distance., The clusters most vulnerable to evaporation and disruption are the low-mass ones as well as those located at small galactocentric distance.
+ In other words. the disruption rate of globular clusters depends on their initial distribution in space as well as on their initial distribution in mass.," In other words, the disruption rate of globular clusters depends on their initial distribution in space as well as on their initial distribution in mass."
+ The key. point is that the mass and number density. profiles show contrasting temporal evolutions., The key point is that the mass and number density profiles show contrasting temporal evolutions.
+ While the evolution of the latter is heavily driven by the initial cluster mass μα»ectrum. the former remains almost unchanged during evolution for a Llubble-time. providing always that the Galactic potential remains smooth anc was slowly. varving.," While the evolution of the latter is heavily driven by the initial cluster mass spectrum, the former remains almost unchanged during evolution for a Hubble-time, providing always that the Galactic potential remains smooth and was slowly varying."
+ Llenee. the robustness of the mass density. profile provides us with an immediate estimate of the initial steepness of the spatial cistribution.," Hence, the robustness of the mass density profile provides us with an immediate estimate of the initial steepness of the spatial distribution."
+ This can then be evolved. for various initial cluster mass spectra and the resulting number density profiles compared to the observed. one in order to discriminate which cluster niass spectrum provides the best match between the data and the mocol., This can then be evolved for various initial cluster mass spectra and the resulting number density profiles compared to the observed one in order to discriminate which cluster mass spectrum provides the best match between the data and the model.
+ In order to test this idea. we have evolved. various putative elobular cluster systems characterized by dlilferent combinations of initial number density. profiles (i.e... how clusters are distributed in space) and initial mass spectra," In order to test this idea, we have evolved various putative globular cluster systems characterized by different combinations of initial number density profiles (i.e., how clusters are distributed in space) and initial mass spectra"
+Different simulations were performed adding to each set-up spectra a grid of svuthetic spectra with broad gaussian components of various widths ancl amplitudes centered on Ho.,Different simulations were performed adding to each set-up spectra a grid of synthetic spectra with broad gaussian components of various widths and amplitudes centered on $\alpha$.
+ We then applied our template and extraction analvsis to deduce the following limits., We then applied our template and extraction analysis to deduce the following limits.
+ For the CAILA and OSN spectra. broad components equivalent to or higher of the total blended flix in Io -[NII|] were recovered.," For the CAHA and OSN spectra, broad components equivalent to or higher of the total blended flux in $\alpha+$ [NII] were recovered."
+ Using medians of AGN blended flux and redshlfit (his transforms into a detection limit in Ha. broad luminosity of 3.5x10 eres |., Using medians of AGN blended flux and redshfit this transforms into a detection limit in $\alpha$ broad luminosity of $3.5\times10^{38}$ ergs $^{-1}$.
+ We find slightly higher fraction (201) lor the NOT and SPM spectra. equivalent to a detection limit in huninosity of 4.0x105 eres +.," We find slightly higher fraction ) for the NOT and SPM spectra, equivalent to a detection limit in luminosity of $4.0\times10^{38}$ ergs $^{-1}$."
+ Only three BLAGNs in the HES97 sample have a luminosity lower than these limits., Only three BLAGNs in the HFS97 sample have a luminosity lower than these limits.
+ Obviously. the lack of BLAGNs encountered in our samples cannot be explain by a higher detection limits in our samples.," Obviously, the lack of BLAGNs encountered in our samples cannot be explain by a higher detection limits in our samples."
+ There seem to be no obvious observational biases or differences in reduction and analysis methods capable of reproducing the lack of BLAGNs in CGs as compared to lower density environments., There seem to be no obvious observational biases or differences in reduction and analysis methods capable of reproducing the lack of BLAGNs in CGs as compared to lower density environments.
+ It is consequently reasonable to conclude that this phenomenon must be related to the environment (tvpical of CGs., It is consequently reasonable to conclude that this phenomenon must be related to the environment typical of CGs.
+ In the unification model for AGNs. a torus of matter is assumed to be responsible for hiding the BLR from the observer.," In the unification model for AGNs, a torus of matter is assumed to be responsible for hiding the BLR from the observer."
+ To be consistent with our analvsis. this mechanism should be much more efficient in CGs.," To be consistent with our analysis, this mechanism should be much more efficient in CGs."
+ However. this assumption goes against the evidence of tidal stripping: in CGs the infall of gas in the disk seem {ο be stopped. generally. diminishing the amount of star lormation (Cozioletal.1998.2000:Verdes-Montenegro2001:delaRosaetal.2007;Durbala 2008).," However, this assumption goes against the evidence of tidal stripping: in CGs the infall of gas in the disk seem to be stopped, generally diminishing the amount of star formation \citep{coz98,coz00,ver01,ros07,dur08}."
+.. At the same time. the number of earlv-tvpe galaxies in CGs is observed to increase.," At the same time, the number of early-type galaxies in CGs is observed to increase."
+ Therefore. a possible reason why no BLRs appear in AGNs in CGs may be because any amount of gas that has reached the center was consumed into stars. building larger bulges (deCarvalho&Coziol1999).," Therefore, a possible reason why no BLRs appear in AGNs in CGs may be because any amount of gas that has reached the center was consumed into stars, building larger bulges \citep{car99}."
+. Alternatively. the bulges of galaxies in CGs may have grown without gas. through dry mergers (Coziol&Plauchu-Fravn2007).," Alternatively, the bulges of galaxies in CGs may have grown without gas, through dry mergers \citep{coz07}."
+. The fact that the average Iuminositv of the AGNs in CGs is low is another argument in favor of the dissolution hypothesis for the BLR., The fact that the average luminosity of the AGNs in CGs is low is another argument in favor of the dissolution hypothesis for the BLR.
+ According to recent results obtained by reverberation mapping. the size of the BLR. in ACNs is correlated to the optical Iuminosity at citepkas05..," According to recent results obtained by reverberation mapping, the size of the BLR in AGNs is correlated to the optical luminosity at \\citep{kas05}."
+ Ho is consequently possible to imagine the size of the DLE shrinking almost to zero al some low threshold huninosity (Elitzur&Shlosman2006)., It is consequently possible to imagine the size of the BLR shrinking almost to zero at some low threshold luminosity \citep{eli06}.
+. The huninositv at iin our samples range [rom log(AL4(5100A ))) 40.7 to 43.1 (in units of erg 1): comparing with data of Petersonetal.(2004) we are in the lower Iuminosity. part of their distribution.," The luminosity at in our samples range from $\lambda$ $_{\lambda}$ )) 40.7 to 43.1 (in units of erg $^{-1}$ ); comparing with data of \citet{pet04} we are in the lower luminosity part of their distribution,"
+The left. middle and right panel of Figure 1. plots (hie results for the case that the projection was done onto the xv. vz. aud zx plane. respectively.,"The left, middle and right panel of Figure \ref{fig:theory} plots the results for the case that the projection was done onto the xy, yz, and zx plane, respectively."
+ The errors include both of (he sample variance as well as the Poisson noise., The errors include both of the sample variance as well as the Poisson noise.
+ The sample variance is caleulated as one standard deviation among the 1000 bootstrap resamples., The sample variance is calculated as one standard deviation among the $1000$ bootstrap resamples.
+ As one can see. the distributions has a sharp peak al the first bin 0°<©10°. as in the three dimensional case. which proves (hat. projections onto the two dimensional planes do not diminish the correlations between the major axes of matter and satellite galaxy distributions.," As one can see, the distributions has a sharp peak at the first bin $0^{o}\le\phi\le 10^{o}$, as in the three dimensional case, which proves that projections onto the two dimensional planes do not diminish the correlations between the major axes of matter and satellite galaxy distributions."
+" In the following section. we compare (his numerical prediction with the observational result aud test the hypothesis that the observational result is consistent with the numerical prediction based on à ACDM cosmology,"," In the following section, we compare this numerical prediction with the observational result and test the hypothesis that the observational result is consistent with the numerical prediction based on a $\Lambda$ CDM cosmology."
+ Performing two dimensional analysis of the weak lensing shear maps constructed with hieh-quality Subaru/Supreme-cam (Miyazakietal.2002).. Ogurietal.(2010) have measured the (wo dimensional projected shapes and orientations of dark matter distributions in the galaxy clusters [rom the Local Cluster Substructure Survey (Okabeetal.2010).," Performing two dimensional analysis of the weak lensing shear maps constructed with high-quality Subaru/Supreme-cam \citep{miyazaki-etal02}, \citet{Oguri-etal10} have measured the two dimensional projected shapes and orientations of dark matter distributions in the galaxy clusters from the Local Cluster Substructure Survey \citep{okabe-etal10}."
+. Using a subsample of 18 galaxy. clusters for which the weak lensing shear maps are well fitted to the elliptical density profiles (Jing&Suto2002).. thev determined the longest axis of the clark malter distribution and measured its position angle for each cluster (see Table 1 in Oeuri et al.," Using a subsample of $18$ galaxy clusters for which the weak lensing shear maps are well fitted to the elliptical density profiles \citep{JS02}, they determined the longest axis of the dark matter distribution and measured its position angle for each cluster (see Table 1 in Oguri et al."
+ 2010)., 2010).
+ selecting those member galaxies whose r-band magnitudes are brighter than 22 mag. thev also fitted the member galaxy distribution (confined to the same region) in each cluster to an elliplical power law profile and determined (he positions anele of the longest axis οἱ ealaxy distributions (see Table 2 in Oguri οἱ al.," Selecting those member galaxies whose r-band magnitudes are brighter than $22$ mag, they also fitted the member galaxy distribution (confined to the same region) in each cluster to an elliptical power law profile and determined the positions angle of the longest axis of galaxy distributions (see Table 2 in Oguri et al."
+ 2010)., 2010).
+ According to Ogurietal.(2010).. the fitting of the dark matter and member galaxy distribution in each cluster was done at (he square region 20!x20. around cluster's center chosen to be the location of the brightest cluster galaxy. where 20° corresponds to approximately 2-3 Mpe for the selected clusters in a redshift range of 0.1<z«0.3.," According to \citet{Oguri-etal10}, the fitting of the dark matter and member galaxy distribution in each cluster was done at the square region $20^{\prime}\times 20^{\prime}$, around cluster's center chosen to be the location of the brightest cluster galaxy, where $20^{\prime}$ corresponds to approximately $2$ $3$ Mpc for the selected clusters in a redshift range of $0.1< z < 0.3$."
+ For a detailed description of the (wo dimensional fitting of weak lensing shear field and the cluster sample. see Ogurietal.(2010).," For a detailed description of the two dimensional fitting of weak lensing shear field and the cluster sample, see \citet{Oguri-etal10}."
+. By comparing the position angles of dark matter and galaxy distributions in each cluster. they found only very weak correlations between (he two.," By comparing the position angles of dark matter and galaxy distributions in each cluster, they found only very weak correlations between the two."
+ They considered several possible svslemalies such as galaxy bias. dilution effect. fitting area and field-galaxy contamination but found no svstematies significant enough to explain the observed trend.," They considered several possible systematics such as galaxy bias, dilution effect, fitting area and field-galaxy contamination but found no systematics significant enough to explain the observed trend."
+ Although Oguri(2010) asserted that the ellipticity distribution of dark matter distributions in the observed galaxy clusters agrees well with the theoretical prediction based on a ACDM, Although \citet{Oguri-etal10} asserted that the ellipticity distribution of dark matter distributions in the observed galaxy clusters agrees well with the theoretical prediction based on a $\Lambda$ CDM
+"can write F(r)ccr(?*9*!—y(s)ocs(?*9, where F(r) is the flux density in the image domain, V(s) is the Fourier transform of the flux density, s is u-v distance and, p and q are the power- indices of the density (nc r?) and temperature structure (Tοr ), respectively.","can write $F(r) \propto r^{-(p+q)+1} \rightarrow V(s) \propto s^{(p+q)-3}$, where $F(r)$ is the flux density in the image domain, $V(s)$ is the Fourier transform of the flux density, $s$ is $u$ $v$ distance and, $p$ and $q$ are the power-law indices of the density $n \propto r^{-p}$ ) and temperature structure $T \propto r^{-q}$ ), respectively."
+ The least-squares fit (see Fig. 2)), The least-squares fit (see Fig. \ref{Fklamall15}) )
+" to both the 1 mm and 3 mm data yields a slope of —1.01 + 0.02, corresponding to p*q~ 2.0."," to both the 1 mm and 3 mm data yields a slope of $-$ 1.01 $\pm$ 0.02, corresponding to $p + q \sim$ 2.0."
+" Points starting at 40 KA, which corresponds to an angular scale of 6"" (0.1 pc), were excluded from the fit because the signal-to-noise ratio gets very low."," Points starting at 40 $\lambda$, which corresponds to an angular scale of $\sim$ (0.1 pc), were excluded from the fit because the signal-to-noise ratio gets very low."
+" If the index g= 0.4 (?),, the density profile indexis 1.6."," If the index $q =$ 0.4 , the density profile indexis 1.6."
+" We estimate a mean volume gas density of the core (n— , where R is the core radius) of 7x10? using the mass EISderived from the dust observations at 1 mm (13 Mo)) and the geometric mean of the deconvolved linear size as the core radius (R=J(FWHMmq/2)x(FWHMiyi,/2) = 0.04 pc)."," We estimate a mean volume gas density of the core $n = \frac{M_g}{\frac{4}{3}\pi R^3\mu_{H_2} m_H}$ , where $R$ is the core radius) of $7 \times 10^5$ using the mass derived from the dust observations at 1 mm (13 ) and the geometric mean of the deconvolved linear size as the core radius $R = \sqrt{(FWHM_{\rm maj}/2)\times (FWHM_{\rm nin}/2)}$ = 0.04 pc)."
+ The scaling factor of 15 corresponds to a spectral index o= 3.0 + 0.4 and we thus obtain an opacity spectral index B= 1.0 + 0.4., The scaling factor of 15 corresponds to a spectral index $\alpha =$ 3.0 $\pm$ 0.4 and we thus obtain an opacity spectral index $\beta =$ 1.0 $\pm$ 0.4.
+" The derived dust opacity index corresponds, taking the upper limit, to a “normal” interstellar dust material value."," The derived dust opacity index corresponds, taking the upper limit, to a “normal” interstellar dust material value."
+" Nonetheless, if taking the lower limit, the dust opacity index goes down to a value that is usually found toward disks??),, which has been ascribed to grain growth."," Nonetheless, if taking the lower limit, the dust opacity index goes down to a value that is usually found toward disks, which has been ascribed to grain growth."
+" Alternatively, the presence of winds/outflows (see Sect. ?2))"," Alternatively, the presence of winds/outflows (see Sect. \ref{g11:diout}) )"
+ could mimic a low value of 8(?).., could mimic a low value of $\beta$.
+ Here we adopt a 6=1.6 in the rest of the paper., Here we adopt a $\beta = 1.6$ in the rest of the paper.
+" Spectroscopic parameters of the detected C*4S 2 > 1 and CH30H 2,—1, and 5,— 4, transitions toward G11.11P1 are listed in Table 3..", Spectroscopic parameters of the detected $^{34}$ S 2 $\rightarrow$ 1 and $_3$ OH $_{k} \rightarrow$$_{k}$ and $_{k} \rightarrow$ $_{k}$ transitions toward G11.11P1 are listed in Table \ref{Teup}. .
+" The interferometric integrated intensity maps, overlaid on their corresponding continuum emission, along with sample spectra are presented in Figs. 3-—5.."," The interferometric integrated intensity maps, overlaid on their corresponding continuum emission, along with sample spectra are presented in Figs. \ref{Fc34s}- \ref{F1mm}."
+" The integrated intensity map of the C?S 2 — 1 line, is shown in the panel of Fig. 3.."," The integrated intensity map of the $^{34}$ S 2 $\rightarrow$ 1 line, is shown in the panel of Fig. \ref{Fc34s}."
+" This emission is spatially coincident with the 3 mm continuum; moreover, some weaker emission (at 30 level) extends along the E-W direction."," This emission is spatially coincident with the 3 mm continuum; moreover, some weaker emission (at $\sigma$ level) extends along the E-W direction."
+ Observed parameters toward the peak position in the integrated intensity map are presented in Table 4.., Observed parameters toward the peak position in the integrated intensity map are presented in Table \ref{Ttrans}.
+" The virial mass of a core with a power-law density distribution is My;=Kkic2R/G where kj=(5--2p)/(3— p), oy is the three-dimensional root-mean- velocity which is related to the FWHM line width, AV, through c2=(8/81n2)AV?, R is the core radius and G is the Gravitational constant."," The virial mass of a core with a power-law density distribution is $M_{\rm vir} = k_1 \sigma_{\rm v}^2R/G$ where $k_1 = (5-2p)/(3-p)$ , $\sigma_{\rm v}$ is the three-dimensional root-mean-square velocity which is related to the FWHM line width, $\Delta{\rm V}$, through $\sigma_{\rm v}^2 = (3/8{\rm ln2})\Delta{\rm V^2}$, $R$ is the core radius and $G$ is the Gravitational constant."
+" For p=1.6 (see Sect. 3.1.3)),"," For $p = 1.6$ (see Sect. \ref{g11:dentemp}) ),"
+" the virial mass can be expressed as A virial mass of ~135 is calculated when using kms! (see Table 4)) of the C?S optically thin line and R= 0.05 pc, which corresponds the geometric mean of the deconvolved linear size."," the virial mass can be expressed as A virial mass of $\sim$ 135 is calculated when using $\Delta{\rm V} = 4.1~\kms$ (see Table \ref{Ttrans}) ) of the $^{34}$ S optically thin line and $R =$ 0.05 pc, which corresponds the geometric mean of the deconvolved linear size."
+" The virial parameter for a spherical cloud is defined as ay,=Myir/Mtot, where Mi is the total mass."," The virial parameter for a spherical cloud is defined as $\alpha_{\rm vir} \equiv M_{\rm vir}/M_{\rm tot}$, where $M_{\rm tot}$ is the total mass."
+" If we assume that Mi=Με37Mo, then ayir= 3.6."," If we assume that $M_{\rm tot}= M_{\rm g} = 37~\msun$, then $\alpha_{\rm vir} = 3.6$ ."
+" It is worth noting that the C*4S line width could be affected by unbound motions, e.g., outflows, partially broadening the line andthereby increasing the virial mass estimate significantly and the virial parameter as well."," It is worth noting that the $^{34}$ S line width could be affected by unbound motions, e.g., outflows, partially broadening the line andthereby increasing the virial mass estimate significantly and the virial parameter as well."
+responsible for the observed. variations in these rings particle size distributons.,responsible for the observed variations in these rings' particle size distributons.
+ For exaniple. the overall clvuamical environments the rings inhabit may influence how efficiently parlicles of different sizes are generated and how quickly particles are lost to the svstem by collisions with nearby rings and moons.," For example, the overall dynamical environments the rings inhabit may influence how efficiently particles of different sizes are generated and how quickly particles are lost to the system by collisions with nearby rings and moons."
+ In this context. the compact F-rine structures might correspond (o regions where larger particles are produced more rapicly (han other parts of the ring. perhaps again due to differences in the (vpical impact speeds into and between particles.," In this context, the compact F-ring structures might correspond to regions where larger particles are produced more rapidly than other parts of the ring, perhaps again due to differences in the typical impact speeds into and between particles."
+ Future detailed comparisons of these spectral data wilh predictions derived from various models of the revelant processes should vield useful new insights into the dynamics and particle properties of these dusty rings., Future detailed comparisons of these spectral data with predictions derived from various models of the revelant processes should yield useful new insights into the dynamics and particle properties of these dusty rings.
+ We wish to acknowledge NASA. the Cassini project and the VIMS team lor providing the data that made (his analysis possible.," We wish to acknowledge NASA, the Cassini project and the VIMS team for providing the data that made this analysis possible."
+ We also thank IL. Throop and J. N. Cuzzi for their detailed and insightful reviews of this manuscript. and L. Esposito and C. Murray. [or the useful discussions.," We also thank H. Throop and J. N. Cuzzi for their detailed and insightful reviews of this manuscript, and L. Esposito and C. Murray for the useful discussions."
+ This work was supported in part bv a Cassini Data Analvsis Program grant NNNO9AETAGC., This work was supported in part by a Cassini Data Analysis Program grant NNX09AE74G.
+Kepler operation which spaus the interval 2009 Jie 20 to September 16., operation which spans the interval 2009 June 20 to September 16.
+ Data gaps correspoud to removed artifacts aud monthly data downloads., Data gaps correspond to removed artifacts and monthly data downloads.
+ The quicsceut V baud magnitude is reported conteuiporaucouslv by the American Association of Variable Star Observers to be Vix 19., The quiescent V band magnitude is reported contemporaneously by the American Association of Variable Star Observers to be V $\simeq$ 19.
+" While we male no attempt at a color correction. these data indicate thatNepler short cadence data achieves relative photometric accuracy upon V = 19 sources, in agreenmient with Calllaud et al. ("," While we make no attempt at a color correction, these data indicate that short cadence data achieves relative photometric accuracy upon V = 19 sources, in agreement with Gilliland et al. ("
+2010).,2010).
+ The data reveal seven cousecutive dwarf nova outbursts over the quarter of varying magnitude. duration and iuterval. with a comparatively short duty cevele ou the order of 10 davs.," The data reveal seven consecutive dwarf nova outbursts over the quarter of varying magnitude, duration and interval, with a comparatively short duty cycle on the order of 10 days."
+ Oue outburst is a superoutburst. confirming V331 Lyr as aun SU UMXa class object. first reported by Nato (1993).," One outburst is a superoutburst, confirming V334 Lyr as an SU UMa class object, first reported by Kato (1993)."
+ Discrete. Fourier Transforms over two short data samples sugecst we are detecting superluunps duriug superoutburst., Discrete Fourier Transforms over two short data samples suggest we are detecting superhumps during superoutburst.
+ Data extracted from 22.155.0062.155.025 reveal a peak at what we will initially refer to as the orbital period. == 204737 + (0.00001 hours. although we argue for an alternative potential origin of this signal iu Sec. 3..," Data extracted from 2,455,006--2,455,025 reveal a peak at what we will initially refer to as the orbital period, = 2.05737 $\pm$ 0.00001 hours, although we argue for an alternative potential origin of this signal in Sec. \ref{discussion}. ."
+ The error is a lo estimate based upon 1.000 Monte Carlo trials where the transform has been recalculated after recasting cach point iu the photometric curve according to the inverse normal cumulative fiction and a random πας generator.," The error is a $\sigma$ estimate based upon 1,000 Monte Carlo trials where the transform has been recalculated after recasting each point in the photometric curve according to the inverse normal cumulative function and a random number generator."
+ After 1.0000 trials all period ieasurciments occured within the ranec 2.057332.05739 hours.," After 1,000 trials all period measurements occurred within the range 2.05733–2.05739 hours."
+ A second epoch was extracted during the superoutburst over range 22.155.0582.155.080.," A second epoch was extracted during the superoutburst over range 2,455,058--2,455,080."
+ While the power spectrum reveals a stronger but less coherent signal compared to the earlier epoch. the best period was determined by identical method to be == 2.2023 + 0.0001 hours and identified as a superhliunp.," While the power spectrum reveals a stronger but less coherent signal compared to the earlier epoch, the best period was determined by identical method to be = 2.2023 $\pm$ 0.0001 hours and identified as a superhump."
+ The fractional excess of the superlauup period over the orbital period is often characterized by εἰ= “Poi LD ϱ, The fractional excess of the superhump period over the orbital period is often characterized by $\epsilon^+ =$ $-$.
+" fwehaccidentificdthetwosiqualsaccarately, Ομ = 7.0LL x"," If wehave identified the two signals accurately, The V344 Lyr $\epsilon^+$ = 7.044 $\pm$."
+ Consequently this binary would occur at the extreme cud of the εἰ distribution of the SU UMa class of stars tabulated by Patterson (2005)., Consequently this binary would occur at the extreme end of the $\epsilon^+$ distribution of the SU UMa class of stars tabulated by Patterson (2005).
+ Neither the putative orbital nor superhuup periods are permanent features of the O2 light curve., Neither the putative orbital nor superhump periods are permanent features of the Q2 light curve.
+ We serformed a dvuamic Fourier power analysis of V3LL Lyr bv taking 5 dav scemeuts stepped by 0.5 days and calculating Discrete Fourier traustorms on each nue slice., We performed a dynamic Fourier power analysis of V344 Lyr by taking 5 day segments stepped by 0.5 days and calculating Discrete Fourier transforms on each time slice.
+ Power spectra are provided in Figure 2.., Power spectra are provided in Figure \ref{fig2}.
+ The somewhat surprising fiudiug is that the quicscent photometric variability found iu the post-uperoutburst data at 22.155.075 docs not have the sae nature as he quiescent signal measured at 22.155.010. Post-superoutburst.," The somewhat surprising finding is that the quiescent photometric variability found in the post-superoutburst data at 2,455,075 does not have the same nature as the quiescent signal measured at 2,455,010. Post-superoutburst,"
+ and throughout the following norual outburst. V31L4 Lyr continues to be dominated bv superhuups.," and throughout the following normal outburst, V344 Lyr continues to be dominated by superhumps."
+ The TTI model strugeles ecnerally to predict superhunips within quiescence or normal outbursts and so theAcpler data provide us with an inmediate challenge., The TTI model struggles generally to predict superhumps within quiescence or normal outbursts and so the data provide us with an immediate challenge.
+ The detection of post-superoutburst superhuups iu a CV ds rare. but not unique.," The detection of post-superoutburst superhumps in a CV is rare, but not unique."
+ Two iuecnbers of the SU. UMa. class. V1150. Oxi and ER UMa. have also revealed ΠΕ within quiesceuce or lormal outburst (Patterson et al.," Two members of the SU UMa class, V1159 Ori and ER UMa have also revealed superhumps within quiescence or normal outburst (Patterson et al."
+ 1995: Cao ct al., 1995; Gao et al.
+ 1999)., 1999).
+ The WZ See star EG Cuc has also revealed superlunps up to 90 davs after the decay from superoutburst (Patterson et al., The WZ Sge star EG Cnc has also revealed superhumps up to 90 days after the decay from superoutburst (Patterson et al.
+ 1998)., 1998).
+ Ποια (2001) has argued that quiesceut aud normal outburst superhumips are possible iu extreme lass ratio objects where qoS0.07 (q=Ab/MA: M» is the colupanion star mass. Ay is the white dwarf mass).," Hellier (2001) has argued that quiescent and normal outburst superhumps are possible in extreme mass ratio objects where $q \lesssim 0.07$ $q = M_2/M_1$; $M_2$ is the companion star mass, $M_1$ is the white dwarf mass)."
+ Weanedio is possible because augular 1101001tn1u transfer by tidal dissipation is comparatively cfiicicut in this reguue: both cooling aud heating frouts can propagate even when the outer edge of the accretion disk exceeds the 3:1 resonance radius., Such a scenario is possible because angular momentum transfer by tidal dissipation is comparatively efficient in this regime; both cooling and heating fronts can propagate even when the outer edge of the accretion disk exceeds the 3:1 resonance radius.
+ Although uo coufiniied orbital period for this object exists currently. we ask iu this section whether V31L Lvr fits this same low-1ass ratio profile.," Although no confirmed orbital period for this object exists currently, we ask in this section whether V344 Lyr fits this same low-mass ratio profile."
+ Tn order to proceed we must employ a small series of clupirical relations., In order to proceed we must employ a small series of empirical relations.
+ Frou Patterson (2005) we extract the orbital and superlauup periods measured from a sample of SU UMa stars., From Patterson (2005) we extract the orbital and superhump periods measured from a sample of SU UMa stars.
+ The distribution of e! within this siuuple is plotted over in Figure 3.., The distribution of $\epsilon^+$ within this sample is plotted over in Figure \ref{fig3}.
+ Standard SU UMa objects are represeuted κ΄ ciele svinbols. while the three previous quiesceut superliup examples are represented by square svuibols.," Standard SU UMa objects are represented by circle symbols, while the three previous quiescent superhump examples are represented by square symbols."
+ Based upon the empirical fitting of superlinnrp sources with relatively secure measured mass ratios. Patterson (2005) provides the relation e!=4(0.18|0.294): we xovide the predicted g on the right-hand vertical axis of Figure 3..," Based upon the empirical fitting of superhump sources with relatively secure measured mass ratios, Patterson (2005) provides the relation $\epsilon^+ = q\left(0.18 + 0.29q\right)$; we provide the predicted $q$ on the right-hand vertical axis of Figure \ref{fig3}."
+ The curve overlaved upou the distribution is a heoretical εἰ prediction using the e! 4 relation and the assumptionshat i) M4 = usun.. aud i) the orbital period-compauion star deusity relation aud zero-age main sequence (ZAMS) companion star mmass-raclius relation parameterized from the data of," The curve overlayed upon the distribution is a theoretical $\epsilon^+$ prediction using the $\epsilon^+$ $q$ relation and the assumptionsthat i) $M_1$ = , and ii) the orbital period-companion star density relation and zero-age main sequence (ZAMS) companion star mass-radius relation parameterized from the data of"
+We inteerate over f to obtain the 1-D likelihood for the mass of the deflector.,We integrate over $f$ to obtain the 1-D likelihood for the mass of the deflector.
+ This vields the most probable mass of the Talo deflector. given with lo error bars: More statistics is obviously required to constrain the mass of halo deflectors.," This yields the most probable mass of the Halo deflector, given with $1\sigma$ error bars: More statistics is obviously required to constrain the mass of halo deflectors."
+ We studied a wide range of disk-halo models., We studied a wide range of disk-halo models.
+ Table 5 sunnmnarzes their characteristics., Table \ref{desc_model} summarizes their characteristics.
+ The rotation velocity at the Sun predicted by these models is always within he observational range: πι.)ze200ian/s CMeryifiek 1992). or Vr«(cR.j=220415kin/s (Binney aud Tremaine 1987).," The rotation velocity at the Sun predicted by these models is always within the observational range: $V_{\rm Tot}({\rm
+R_\odot})\simeq 200 {\rm \:km/s}$ (Merrifield 1992), or $V_{\rm Tot}({\rm
+R_\odot})=220 \pm 15 {\rm \:km/s}$ (Binney and Tremaine 1987)."
+ The optical depth is independent of the mass of he deflectors: the event rate is given assunines that al deflectors in the halo have a single mass equal to 1M.," The optical depth is independent of the mass of the deflectors; the event rate is given assuming that all deflectors in the halo have a single mass equal to $1{\rm
+\;M_\odot}$."
+ To obtain the predicted. value of the eveut rate for other uasses. one ouly needs to scale by aay. the integral of he mass depeudence of the event rate times fa;(AL). the rormadlized mass distribution: This suuplifics to (MfM.)HM? for oa Dirac distribution peaked at AL.," To obtain the predicted value of the event rate for other masses, one only needs to scale by $\eta_M$, the integral of the mass dependence of the event rate times $f_M(M)$, the normalized mass distribution: This simplifies to $(M/{\rm \;M_\odot})^{-1/2}$ for a Dirac distribution peaked at $M$."
+ Model Lis the “standard” halo model: au isotropic aud isothermal spherical halo. with a mass distribution οἼνοιι in spherical coordinates by (Caldwell aud. Coulson 1981): where BR.=5 kpc is the Valo vcore radius? and RS=δι kpe is the distance from the Sun to the Galactic Center.," Model 1 is the “standard” halo model: an isotropic and isothermal spherical halo, with a mass distribution given in spherical coordinates by (Caldwell and Coulson 1981): where $R_c=5$ kpc is the Halo “core radius” and $R_\odot = 8.5$ kpc is the distance from the Sun to the Galactic Center."
+ Model 2a is the equivalent power-law model (Evans. 1993).," Model 2a is the equivalent power-law model (Evans, 1993)."
+ Model 2b has a maximal disk (Ny=LOOAL.per) and a very light halo. model 5 imtermeciate disk aud halo.," Model 2b has a maximal disk $\Sigma_0=100{\rm
+\;M}_\odot/{\rm pc}^2$ ) and a very light halo, model 5 intermediate disk and halo."
+ Model 3 has a flattened EG halo (axis ratio q=0.71) and model La decreasing rotation curve (.)=0.2 where is proportional to the logarithin of the asviuptotic slope)., Model 3 has a flattened E6 halo (axis ratio $q=0.71$ ) and model 4 a decreasing rotation curve $\beta=0.2$ where $\beta$ is proportional to the logarithm of the asymptotic slope).
+ Models 2.5 are all power-law halo models. with self-consistent mass aud velocity distributions.," Models 2–5 are all power-law halo models, with self-consistent mass and velocity distributions."
+ For each of these. we derive the expected uunboer of events versus the mass AL of the objects iu the ITalo (assuniug a Dirac lass distribution) and compare with the observations (cf figure &)).," For each of these, we derive the expected number of events versus the mass $M$ of the objects in the Halo (assuming a Dirac mass distribution) and compare with the observations (cf figure \ref{nbevt}) )."
+ Whatever the halo model consicered. the sole event we observed. if caused by a deflector in the halo of our Galaxy. corresponds to at least of the total optical depth expected for f=1.," Whatever the halo model considered, the sole event we observed, if caused by a deflector in the halo of our Galaxy, corresponds to at least of the total optical depth expected for $f=1$."
+ This is a very small number statistics. however.," This is a very small number statistics, however."
+ The very lone time-scale of the observed event suggests hat it could show measurable distortions im its light curve due to the motion of the Earth around the Sun. (the xuwallax effect: Gould 1992). provided that the Einstein radius projected onto the plane of the Earth is not mach arecry than the Earth orbital radius.," The very long time-scale of the observed event suggests that it could show measurable distortions in its light curve due to the motion of the Earth around the Sun, (the parallax effect: Gould 1992), provided that the Einstein radius projected onto the plane of the Earth is not much larger than the Earth orbital radius."
+ The first detection of parallax in a eravitational iicrolensing event was observed by Alcock et al. (, The first detection of parallax in a gravitational microlensing event was observed by Alcock et al. (
+1995).,1995).
+" The natural parameter olneasure the streneth of parallax is the semi major axis of the Earth orbit. Ro. in units of the projected Einsteiu radius: bu=BRotlo cMHpgswhoee.=Da/D.. with Dy he distance from the observer to the deflector and D, the distance to the source."," The natural parameter to measure the strength of parallax is the semi major axis of the Earth orbit, $R_O$, in units of the projected Einstein radius: $\delta u=R_O(1-x)/R_E$ where $x=D_d/D_s$, with $D_d$ the distance from the observer to the deflector and $D_s$ the distance to the source."
+ No evidence for distortion due to parallax is detectable ou the lieht curve. iuplviug either a very massive defiector with a very large Eiustein radius. or a deflector near the source.," No evidence for distortion due to parallax is detectable on the light curve, implying either a very massive deflector with a very large Einstein radius, or a deflector near the source."
+ Because our results for the standard aud bleuded, Because our results for the standard and blended
+too contaminated.,too contaminated.
+ Instead. we must first correct for backerouncl objects.," Instead, we must first correct for background objects."
+ As noted above. the GC svstems of the dEs are very centrally concentrated.," As noted above, the GC systems of the dEs are very centrally concentrated."
+ Thus we can use the outer regions of the dE images as a fiducial backerouncl., Thus we can use the outer regions of the dE images as a fiducial background.
+" We defined a backerouncl sample in the radial range l.l—1.25’ (corresponding to 5.46.2 kpe at a distance of 17 Alpe). then subtracted the resulting : GCLF in 0.1 mag bins from the central 30"" sample with the appropriate areal correction factor."," We defined a background sample in the radial range $1.1-1.25 \arcmin$ (corresponding to 5.4–6.2 kpc at a distance of 17 Mpc), then subtracted the resulting $z$ GCLF in 0.1 mag bins from the central $30 \arcsec$ sample with the appropriate areal correction factor."
+ We confined the fit to candidates with 0.7«g—2<1.25 because of the small color range of the dEs.," We confined the fit to candidates with $0.7 < g-z < 
+1.25$ because of the small color range of the dEs."
+ The resulting fit (performed as above) gave a turnover of M.=—8.14zx:0.14. compared to the weighted mean of M.=—8.19 for the &Es.," The resulting fit (performed as above) gave a turnover of $M_z = -8.14\pm0.14$, compared to the weighted mean of $M_z = -8.19$ for the gEs."
+ These are consistent., These are consistent.
+ Since ihe 2 GCLFs are being used. the overall color differences. between the GC! systems should have little effect on the peak locations.," Since the $z$ GCLFs are being used, the overall color differences between the GC systems should have little effect on the peak locations."
+ The dispersion of the dEs (c.= 0.74) is less than that of the eEs (o.=1.03): tis is probably partially due to the smaller color range of the dEs.," The dispersion of the dEs $\sigma_z = 0.74$ ) is less than that of the gEs $\sigma_z 
+= 1.03$ ); this is probably partially due to the smaller color range of the dEs."
+ It also indicates (hat the range of galaxy distances is probably nol sienilieant. consistent with the projected appearance of many of (he dEs near the Virgo cluster core (Dinggeli 1937).," It also indicates that the range of galaxy distances is probably not significant, consistent with the projected appearance of many of the dEs near the Virgo cluster core (Binggeli 1987)."
+ That the peak of the GCLF appears to be the same for both &Es ancl dEs is perhaps puzzling., That the peak of the GCLF appears to be the same for both gEs and dEs is perhaps puzzling.
+ Even if both galaxy (tvpes had similar primordial GCLESs. analvtic calculations ancl numerical simulations of GC evolution in the tidal field of the dEs suggest Chat dvnamical friction should destrov high-mass GCs within several Gvr (Oh Lin 2000. Lotz 2001).," Even if both galaxy types had similar primordial GCLFs, analytic calculations and numerical simulations of GC evolution in the tidal field of the dEs suggest that dynamical friction should destroy high-mass GCs within several Gyr (Oh Lin 2000, Lotz 2001)."
+ Tidal shocks (which tend to destroy low-mass GCs) are more significant in massive galaxies. so 1t is unlikely that the loss of high-mass GCs in dEs could be balanced by a corresponding amount of low-mass GC destruction (o preserve a constant. turnover.," Tidal shocks (which tend to destroy low-mass GCs) are more significant in massive galaxies, so it is unlikely that the loss of high-mass GCs in dEs could be balanced by a corresponding amount of low-mass GC destruction to preserve a constant turnover."
+ variations in the GCLF between dIEs and gEs are another possibility. but. cannot be constrained αἱ present.," variations in the GCLF between dEs and gEs are another possibility, but cannot be constrained at present."
+ If we see no evidence in the dE GCLF turnover lor dynamical friction. does this provide significant evidence against this popular scenario for forming dE nuclei?," If we see no evidence in the dE GCLF turnover for dynamical friction, does this provide significant evidence against this popular scenario for forming dE nuclei?"
+ In order study the relationship between the nuclei and GC's of our sample galaxies. we performed photometry and size measuremenis on all of the dE nuclei.," In order study the relationship between the nuclei and GCs of our sample galaxies, we performed photometry and size measurements on all of the dE nuclei."
+ The procedures were identical to. those described above lor the GCs. except that aperture corrections were derived through analytic integration of the derived Ning profile. since many of the nuclei were large compared to (vpical GC sizes.," The procedures were identical to those described above for the GCs, except that aperture corrections were derived through analytic integration of the derived King profile, since many of the nuclei were large compared to typical GC sizes."
+ The properties of the nuclei are listed in Table 3., The properties of the nuclei are listed in Table 3.
+ In Figure 9 we plot parent galaxy Mp vs. M. of the nucleus., In Figure 9 we plot parent galaxy $M_B$ vs. $M_z$ of the nucleus.
+ Svinbol size is proportional to nuclear size., Symbol size is proportional to nuclear size.
+ The overplotted dashed line represents Monte Carlo simulations of nuclear formation through dynamical friction of GCs from Lotz aL(2001)., The overplotted dashed line represents Monte Carlo simulations of nuclear formation through dynamical friction of GCs from Lotz (2001).
+ These assume 5 Gyr of orbital decay (implicitly assuming the GC's are ~5 Gyr old): we have converted their My (o Af. using stellar populations models of Maraston (2005)., These assume 5 Gyr of orbital decay (implicitly assuming the GCs are $\sim 5$ Gyr old); we have converted their $M_V$ to $M_z$ using stellar populations models of Maraston (2005).
+ Lotz [οι that the expected dvnamical friction was inconsistent with observations. since (he observed nuclei were fainter," Lotz found that the expected dynamical friction was inconsistent with observations, since the observed nuclei were fainter"
+lt ijs important to note that equation It ijs an approximation and as such will not provide an exact conversion [rom the observable η to 5.,It is important to note that equation \ref{eq:soft} is an approximation and as such will not provide an exact conversion from the observable $\eta$ to $S$.
+ In. particular. if aris still undergoing reionization equation 14— will underestimate the true softness parameter: not. all. the helium in the IGM will be doubly ionized and equation 13 will no longer hold.," In particular, if is still undergoing reionization equation \ref{eq:soft} will underestimate the true softness parameter; not all the helium in the IGM will be doubly ionized and equation \ref{eq:Hefrac} will no longer hold."
+ Similarlv. i£ is still being reionized. the assumption of ionization equilibrium. will no longer hold and equation 14 will overestimate the true softness parameter.," Similarly, if is still being reionized, the assumption of ionization equilibrium will no longer hold and equation \ref{eq:soft} will overestimate the true softness parameter."
+ There is also a weak temperature dependence in this relation depending on the assumed: values for the recombination cocllicicnts Shulletal. 20049)., There is also a weak temperature dependence in this relation depending on the assumed values for the recombination coefficients \citealt{Shull04}) ).
+ Llowever. we expect this approximation is accurate to within 5 per cent for the redshift range we consider.," However, we expect this approximation is accurate to within $5$ per cent for the redshift range we consider."
+ Equation 14 will be adopted throughout this paper to convert from η to 5., Equation \ref{eq:soft} will be adopted throughout this paper to convert from $\eta$ to $S$.
+ 1n their recent study of the helium reionization history. ZO4b present and forest spectra of the quasar lll 2347—4342 obtained using the Very. Large Telescope (VLT) andFUSE.," In their recent study of the helium reionization history, Z04b present and forest spectra of the quasar $\rm HE$ $2347-4342$ obtained using the Very Large Telescope (VLT) and."
+ They analyse the spectra. by. fitting Gaussian profiles to the lines in the VET spectrum., They analyse the spectra by fitting Gaussian profiles to the lines in the VLT spectrum.
+ Once the components are identified. counterpart. lines are defined. in theFUSE spectrum. constrained. to have the same redshift. and line width as the LLilines: only the column density is allowed to vary.," Once the components are identified, counterpart lines are defined in the spectrum, constrained to have the same redshift and line width as the lines; only the column density is allowed to vary."
+" In the event that optical depth is too high to produce reasonable values of η, additional sub-components are added with a line width μμ=27 kms"," In the event that optical depth is too high to produce reasonable values of $\eta$, additional sub-components are added with a line width $b_{\rm HeII} = 27$ $\rm
+kms^{-1}$."
+" The assumption of equal line widths is based on the analysis of a sub-sample of unblended and lines by ZO4b: they find a line width ratio of £=bgqi/bg,0.95250.12.", The assumption of equal line widths is based on the analysis of a sub-sample of unblended and lines by Z04b; they find a line width ratio of $\xi = b_{\rm HeII}/b_{\rm HI} = 0.95 \pm 0.12$.
+ In this way ZOLb have estimated the column density ratio in the redshift range 2.0<22.9., In this way Z04b have estimated the column density ratio in the redshift range $2.0<z<2.9$.
+ ZO4b detect large variations in the column density ratio on small scales (As& 0.001) and have interpreted these as spatial Uuctuations in the shape of the UVB., Z04b detect large variations in the column density ratio on small scales $\Delta z \simeq 0.001$ ) and have interpreted these as spatial fluctuations in the shape of the UVB.
+ Shullctal.(2004) who used an independent pixel-hby-pixel optical depth technique to analyse the same observational data come to similar conclusions., \cite{Shull04} who used an independent pixel-by-pixel optical depth technique to analyse the same observational data come to similar conclusions.
+ However. as pointed. out. by Lechner&Reimers(2004) and Shulletal.(2004)... values of yeLO and apz400 should be considered with some caution and are probably not real.," However, as pointed out by \cite{FechnerReimers04} and \cite{Shull04}, values of $\eta<10$ and $\eta>400$ should be considered with some caution and are probably not real."
+ Γον are likely to be caused by uncertainties introduced by background subtraction. the low signal-to-noise of the spectra. line saturation and. blending with higher order Lyman series lines and metals.," They are likely to be caused by uncertainties introduced by background subtraction, the low signal-to-noise of the spectra, line saturation and blending with higher order Lyman series lines and metals."
+ We have performed the same analysis as Zü4b on our artificial spectra., We have performed the same analysis as Z04b on our artificial spectra.
+ Phis enables. us to assess if there is à substantial contribution from errors in the line fitting method to the observed variations in 7 and also to determine how well the softness parameter can be recovered., This enables us to assess if there is a substantial contribution from errors in the line fitting method to the observed variations in $\eta$ and also to determine how well the softness parameter can be recovered.
+ We have analvsed. our artificial spectra using an automated: version of the Voigt. profile fitting package ΝΟΕ (CarswellaL. /www.ast.cam.ac.uk/-rfc/vplit.html). moclifiect to simultaneously fit and. absorption lines.," We have analysed our artificial spectra using an automated version of the Voigt profile fitting package VPFIT (Carswell, $\sim$ rfc/vpfit.html), modified to simultaneously fit and absorption lines."
+ “Phe automated Voigt profile fitting. procedure. is run on 50 random lines of sight taken from outputs produced by the 15200 simulation., The automated Voigt profile fitting procedure is run on $50$ random lines of sight taken from outputs produced by the $15-200$ simulation.
+ For each line which is identified. a corresponding line is itted at the same redshift.," For each line which is identified, a corresponding line is fitted at the same redshift."
+ Additional lines are added where no corresponding can be identified., Additional lines are added where no corresponding can be identified.
+ Purbulent line xoadening is assumed. such that the ratio of line widths is fixed to be £=ὑμωιέθηι1.0. consistent with the ZOLb result.," Turbulent line broadening is assumed, such that the ratio of line widths is fixed to be $\xi =
+b_{\rm HeII}/b_{\rm HI} = 1.0$ , consistent with the Z04b result."
+ This should be a reasonable assumption: Llubble xoadening is expected to be the dominant. contribution to he width of absorption lines in the forest. (Weinbergctal.1998))., This should be a reasonable assumption; Hubble broadening is expected to be the dominant contribution to the width of absorption lines in the forest \citealt{Weinberg98}) ).
+. Note. however. that aclelitional motions in the IGM velocity field from convergent lows and shocks may produce a departure. from pure urbulent. broadening for some of the lines.," Note, however, that additional motions in the IGM velocity field from convergent flows and shocks may produce a departure from pure turbulent broadening for some of the lines."
+" The maximum column clensity allowed. in the fitting procedure is fixed at 107""(0 cm27.", The maximum column density allowed in the fitting procedure is fixed at $10^{19}$ $\rm cm^{-2}$.
+ Similarly.m. the upper limit. to the width. of the itted line profiles is set at SO kris," Similarly, the upper limit to the width of the fitted line profiles is set at $80$ $\rm kms^{-1}$."
+ The left. panel of figure. 3. shows a scatter. plot of the column density ratio g we have obtained from our artificial spectra at 2=24., The left panel of figure \ref{fig:compare} shows a scatter plot of the column density ratio $\eta$ we have obtained from our artificial spectra at $z=2.4$.
+ A uniform background: was assumed., A uniform background was assumed.
+ When determining y we only use line pairs which have well defined. profile fits: any Line with an error in the column density greater than a factor of two is rejected., When determining $\eta$ we only use line pairs which have well defined profile fits; any line with an error in the column density greater than a factor of two is rejected.
+ The and column densities are tightly correlated with a moderate scatter., The and column densities are tightly correlated with a moderate scatter.
+ In the right panel of figure 3. we compare the softness parameter calculated from the column density ratio using equation 14. to the actual. value. used. το produce. the artificial spectra at 2«z3., In the right panel of figure \ref{fig:compare} we compare the softness parameter calculated from the column density ratio using equation \ref{eq:soft} to the actual value used to produce the artificial spectra at $2<z<3$.
+ Again we assume a spatially uniform UVB., Again we assume a spatially uniform UVB.
+ For the range of values under consideration (100«S 350). line fitting recovers the actual value of re softness parameter. well.," For the range of values under consideration $100<S<350$ ), line fitting recovers the actual value of the softness parameter well."
+ I£ we take the median column clensity ratio (open squares) the recovered: value is on average 7 per cent less than the actual value used. to produce the artificial spectra., If we take the median column density ratio (open squares) the recovered value is on average $7$ per cent less than the actual value used to produce the artificial spectra.
+ Following Zü4b. we also plot the logarithmic mean of the softness parameter log(5)=Mσσ)n (lled circles with error bars).," Following Z04b, we also plot the logarithmic mean of the softness parameter $\log(\bar{S}) =
+[\sum\log(S_{\rm i})]/n$ (filled circles with error bars)."
+ The uncertainty on the logarithmic mean. log(AS)=STlog(Sp]2nis computed by propagating the uncertainties on the and column densities.," The uncertainty on the logarithmic mean, $\log(\Delta \bar{S}) =
+[\sum[\log(\Delta S_{\rm i})]^{2}]^{1/2}/n$, is computed by propagating the uncertainties on the and column densities."
+ The median appears to provide the better measure of the softness parameter: it is not weighted so heavily in favour of the lower column density lines which are more commonly identified by the fitting package., The median appears to provide the better measure of the softness parameter; it is not weighted so heavily in favour of the lower column density lines which are more commonly identified by the fitting package.
+ As might be expected. we find fitting Voigt profiles to only and absorption features will not recover the softness parameter accurately once the lines are predominantly saturated.," As might be expected, we find fitting Voigt profiles to only and absorption features will not recover the softness parameter accurately once the lines are predominantly saturated."
+ In this case. only higher order absorption lines or an accurate determination of the etfective optical depth from a large statistical sample of absorption spectra would enable the recovery of S.," In this case, only higher order absorption lines or an accurate determination of the effective optical depth from a large statistical sample of absorption spectra would enable the recovery of $S$."
+ Overall. we conclude that the line fitting method of ZOLb will recover the spectral shape of the UVB along the line-of-sight they consider well. although one should keep in mind that the number of observed spectra is very small. and systematic uncertainties in the line fitting procedure means one should view the extreme values of i with a certain amount of scepticism (krechner&Reimers2004:Shull 2004)).," Overall, we conclude that the line fitting method of Z04b will recover the spectral shape of the UVB along the line-of-sight they consider well, although one should keep in mind that the number of observed spectra is very small, and systematic uncertainties in the line fitting procedure means one should view the extreme values of $\eta$ with a certain amount of scepticism \citealt{FechnerReimers04,Shull04}) )."
+ Lt is also clear that the magnitude of the observed scatter in the column densityratio is much larger than our, It is also clear that the magnitude of the observed scatter in the column densityratio is much larger than our
+accretion disk aud the BER. aud possibly t10 torus itself (seee.c.οειν,"accretion disk and the BLR, and possibly the torus itself \citep[see e.g.][]{elvis00,nic00,es06,elvis06}."
+ Tn any case. we would like to point out that our simplification in terms of two separate materials. one compact. roiehlv at ape from the BIT aud. Comptou-thick (the torus) aud the other close to the BIT aud coustituted by Conmptou-thick and/or thin clouds. cau rexoduce the observed plenomenoloey.," In any case, we would like to point out that our simplification in terms of two separate materials, one compact, roughly at a pc from the BH and Compton-thick (the torus) and the other close to the BH and constituted by Compton-thick and/or thin clouds, can reproduce the observed phenomenology."
+ We have presented a monitoring campaign of the Sevtert 2 galaxy. NGC 7582.," We have presented a monitoring campaign of the Seyfert 2 galaxy, NGC 7582."
+ The dramatic spectral variability observed during the Lobservatious is best explained by changes of the absorbing column deusitv of a material close to the N-ray primary source., The dramatic spectral variability observed during the 4 observations is best explained by changes of the absorbing column density of a material close to the X-ray primary source.
+ (ναι the significant variation between a new observation axl the first one. separated by only 20 hours. its cistaice can be estimated to be a few «10D Clu. i.c. roughly consistent wit itιο BLR.," Given the significant variation between a new observation and the first one, separated by only 20 hours, its distance can be estimated to be a few $\times10^{15}$ cm, i.e. roughly consistent with the BLR."
+ Together with this material. the presewe of a Coupton-thick material «i larecr scale. likely the ‘torus envisaged in the Unification models. Is recmired m order to accout for the Comptou refleclon component and the neutral iron ka chussion line. whose fluxes appear constant iu vears.," Together with this material, the presence of a Compton-thick material on larger scale, likely the `torus' envisaged in the Unification models, is required in order to account for the Compton reflection component and the neutral iron $\alpha$ emission line, whose fluxes appear constant in years."
+ Ou the to» of that. a third. Compton-thiu material interce)ots the li16 of sigwt and can be associated to the dust lane observe in optical aud N-vav iages.," On the top of that, a third, Compton-thin material intercepts the line of sight and can be associated to the dust lane observed in optical and X-ray images."
+ Tn the last wears. a nuniber of sources have shown a Seonierv for the absorbers which is necessarily more couples. than what ecucrally assunied in snpe Unification Moxels;," In the last years, a number of sources have shown a geometry for the absorbers which is necessarily more complex than what generally assumed in simple Unification Models."
+ NCC 7582 is a clear-cut example. where three neutral absorberemitter regions must be present on very different scales.," NGC 7582 is a clear-cut example, where three neutral absorber/emitter regions must be present on very different scales."
+ We propose that the scenario adopted for this source may be the rule rather than the exception., We propose that the scenario adopted for this source may be the rule rather than the exception.
+ The Unification Moccel should therefore be modified in order to account for all the observational evidence collected so fax., The Unification Model should therefore be modified in order to account for all the observational evidence collected so far.
+ The new scenario is based ou the model oxeseuted by 2s: to the ubiquitous pe-scale torus. an exteude .Comptou-thin material. likely associated o galactic cdist lanes. has to be added.," The new scenario is based on the model presented by \citet{matt00b}: to the ubiquitous pc-scale torus, an extended, Compton-thin material, likely associated to galactic dust lanes, has to be added."
+ Moreover. we sugeest the presence of another material. mace © Comptothick and/or Compton-thin clouds. ycated rouohlv at the BLR. whose preseuce cau re unvelled «ilv when the torus does not intercept ie hne of sight and the material is patchy. cading to observable absorption variability on VAjort timescales.," Moreover, we suggest the presence of another material, made of Compton-thick and/or Compton-thin clouds, located roughly at the BLR, whose presence can be unveiled only when the torus does not intercept the line of sight and the material is patchy, leading to observable absorption variability on short timescales."
+ SB aud GM acknowledge financial support from ASI (eraut 1/023/05/0)., SB and GM acknowledge financial support from ASI (grant 1/023/05/0).
+ We would like to thauk the anonvinous referee for helpful sugeestions., We would like to thank the anonymous referee for helpful suggestions.
+ We would like to thauk the anonvinous referee for helpful sugeestions.O, We would like to thank the anonymous referee for helpful suggestions.
+ We would like to thauk the anonvinous referee for helpful sugeestions.OO, We would like to thank the anonymous referee for helpful suggestions.
+With the exception of sources VLA 8. VLA 123. and VLA 19. all sources were found to be uuresolved. 0.< 072.,"With the exception of sources VLA 8, VLA 13, and VLA 19, all sources were found to be unresolved, $\theta_s \le 0$ 2."
+ There are several uechanisius that can produce colmpact centimeter radio sources m regions of star formation., There are several mechanisms that can produce compact centimeter radio sources in regions of star formation.
+ Iu regions of low mass star formation. thermal jets and evrosvuchrotrou emitters can be present.," In regions of low mass star formation, thermal jets and gyrosynchrotron emitters can be present."
+ Iu regions of high mass star formation we also have strong ionizius radiation available and. iu addition to the two mechauisuis. preseut in low mnmiass star-forming regions. we cau lave ionized stellar winds. ultracompact TTT regions. and radio proplvds.," In regions of high mass star formation we also have strong ionizing radiation available and, in addition to the two mechanisms present in low mass star-forming regions, we can have ionized stellar winds, ultracompact HII regions, and radio proplyds."
+ It is possible to distinguish between these various possibilities with high aneular resolution. inultifrequeucey observations made with the required scusitivity.," It is possible to distinguish between these various possibilities with high angular resolution, multifrequency observations made with the required sensitivity."
+ Let us first consider the subset of 15 sources that have IR counterparts., Let us first consider the subset of 15 sources that have IR counterparts.
+ It is possible that these objects are ultracompact TTT (UCTIIID reeglons. centered arounc vouns πμνο stars.," It is possible that these objects are ultracompact HII (UCHII) regions, centered around young massive stars."
+ The radio cussion cau then be explained as optically thin frec-free radiation. aud the familiar photoionization equilibrium equations (c.g.?7) can be used to determine physical parameters of the svstem.," The radio emission can then be explained as optically thin free-free radiation, and the familiar photoionization equilibrium equations \citep[e.g.][]{ost1974} can be used to determine physical parameters of the system."
+" However. ΟΠΗ regions are nof expected to exhibit short-period time variability. as ds seen du 5 of the 15 racdio/IR SOULCCS,"," However, UCHII regions are not expected to exhibit short-period time variability, as is seen in 8 of the 15 radio/IR sources."
+ The remaining 7 sources can be considered as ΤΙΤΠ candidates., The remaining 7 sources can be considered as UCHII candidates.
+ Amoug this group are tle two brightest radio sources; VLA-18 aud VLÀ-I1. as well as the brightest IR source. VLA-15 = IRS1a.," Among this group are the two brightest radio sources, VLA-18 and VLA-14, as well as the brightest IR source, VLA-15 = IRS1a."
+ With an assimned electron temperature of 10: Is and a clistance of 600 pe we can estimate that a BO ZAMS star would be required to maintain tle ionization iu cach of these UCTIII regions (??)..," With an assumed electron temperature of $10^4$ K and a distance of 600 pc we can estimate that a B0 ZAMS star would be required to maintain the ionization in each of these UCHII regions \citep{pan1973,tho1984}."
+ As for the sources that show time variability. these are more likely to correspond to voung stellar objects (YSOs) with active magnetospheres.," As for the sources that show time variability, these are more likely to correspond to young stellar objects (YSOs) with active magnetospheres."
+ The evrosvuchnrotron radiation frou such objects is expected to show substantial variation over timescales of a few days., The gyrosynchrotron radiation from such objects is expected to show substantial variation over timescales of a few days.
+ One of the brightest racdio/IR sources that nav be placed iu this category 19 VLA-7 = IRS 2a., One of the brightest radio/IR sources that may be placed in this category is VLA-7 = IRS 2a.
+ The SED of this object has been examined from ucar-IR to müllineter wavelengths. aud it iudicates that the star is surrounded by approximately 0.1 AL. of circwstellar material (?)..," The SED of this object has been examined from near-IR to millimeter wavelengths, and it indicates that the star is surrounded by approximately 0.1 $_{\sun}$ of circumstellar material \citep{val1994}."
+ This fiudiug is consistent with our tentative classification of VLAÀ-T as a YSO. since may stars remain shrouded in their natal cuvelopes all the way to the main sequence.," This finding is consistent with our tentative classification of VLA-7 as a YSO, since many stars remain shrouded in their natal envelopes all the way to the main sequence."
+ It is important to note that this classification is at odds with the results of ? and ?.. who label this source as a deeply embedded OD star that is powering the ionization of the W10 III region.," It is important to note that this classification is at odds with the results of \citet{smi1985} and \citet{val1991}, who label this source as a deeply embedded OB star that is powering the ionization of the W40 HII region."
+ Based ou the data preseuted here. the radio sources VLÀ-L1. 15. aud LS may be equally valid candidates for supporting the TWIT region.," Based on the data presented here, the radio sources VLA-14, 15, and 18 may be equally valid candidates for supporting the HII region."
+ Of the five radio sources that do not have IR counterparts. VLAÀ-17 and 19 are of particular interest.," Of the five radio sources that do not have IR counterparts, VLA-17 and 19 are of particular interest."
+ VWLA-L9 is offset by 775 (0.02 pc) to the north of VLA-I18. aud VLA-17 is ouly (77 (0.002 pc) away to the northwest.," VLA-19 is offset by $7\farcs 5$ (0.02 pc) to the north of VLA-18, and VLA-17 is only $0\farcs 7$ (0.002 pc) away to the northwest."
+ These radio sources are both relatively weak. do uot vary with time on short periods. and have mildly uou-circular radio contours.," These radio sources are both relatively weak, do not vary with time on short periods, and have mildly non-circular radio contours."
+ All of this evidence sugeests that these wo radio sources may correspond to shock frouts roni a thermal jet interacting with the ambieut interstellar miedimm., All of this evidence suggests that these two radio sources may correspond to shock fronts from a thermal jet interacting with the ambient interstellar medium.
+ Based on its position and xoxinitv. VLA-18 = IRSIb could well be the star which powers this radio jet.," Based on its position and proximity, VLA-18 = IRS1b could well be the star which powers this radio jet."
+ Iun order to unambiguously determine the rature of these compact radio sources. further information ix required.," In order to unambiguously determine the nature of these compact radio sources, further information is required."
+ A good indicator of CCUM regions is the spectral index at radio requeucies., A good indicator of UCHII regions is the spectral index at radio frequencies.
+ Free-free cinissiou from an UCTIIT region should have a very flat radio spectrum. so a spectral index more uesative than -0.1 is a clear indication that the object is not an UCU region (?)..," Free-free emission from an UCHII region should have a very flat radio spectrum, so a spectral index more negative than -0.1 is a clear indication that the object is not an UCHII region \citep{rod1993}."
+ This measurement could be made using the EVLA to observe WO with simular angular resolution at 6 cin. for example.," This measurement could be made using the EVLA to observe W40 with similar angular resolution at 6 cm, for example."
+ Another important observation to further our understanding of the W10 cluster will be to eather A-vav data to determine the hiseh-energv characteristics of the dominant radio/IR. sources., Another important observation to further our understanding of the W40 cluster will be to gather X-ray data to determine the high-energy characteristics of the dominant radio/IR sources.
+ Civrosviichrotron cutters should produce detectable N-rav clussion. with measurable tine-variabilitv signatures.," Gyrosynchrotron emitters should produce detectable X-ray emission, with measurable time-variability signatures."
+ A combination of high aneular resolution radio observations at lonecr wavelengths plus AN-rav measurements with NAIAL or Chandra would enable us to definitively determine the astrophysical nature of thesesources., A combination of high angular resolution radio observations at longer wavelengths plus X-ray measurements with XMM or Chandra would enable us to definitively determine the astrophysical nature of thesesources.
+ À first Chandra study of WLO as currently iu preparation (2)., A first Chandra study of W40 is currently in preparation \citep{kuh2010}. .
+(see c.g. Vignali et al.,(see e.g. Vignali et al.
+ 2000)., 2000).
+ Sensitive hard X-ray surveys are therefore powerful ools to select aree samples of ACN less biased against absorption aud extinction., Sensitive hard X-ray surveys are therefore powerful tools to select large samples of AGN less biased against absorption and extinction.
+ Om approacl1 consists in takine acvantage of the laree field of view aud good sensitivity of tιο BeppoSAN MECS iustrument (Boclla et al., Our approach consists in taking advantage of the large field of view and good sensitivity of the BeppoSAX MECS instrument (Boella et al.
+ 1997a.b) to survev eus to huudrοςls of square degrees a fluxes το10«‘ere (Fiore. et al.," 1997a,b) to survey tens to hundreds of square degrees at fluxes $\gs5-10\times10^{-14}$ (Fiore et al."
+ 20002). and using higher scusitivity NMM-Newton aiu| Chadra observations o extend tli» survey down to ~1jH Tji several dee?JD (at his Hux the majorivy of the hard NRB is resolved in sources).," 2000a), and using higher sensitivity XMM-Newton and Chandra observations to extend the survey down to $\sim10^{-14}$ on several $^2$ (at this flux the majority of the hard XRB is resolved in sources)."
+ TI10 LCSTIts on the optical identification of a sample of faiit Chandra sources discoverec over the first 0.11 deg? ave been pullished by Fiore et al. (, The results on the optical identification of a sample of faint Chandra sources discovered over the first 0.14 $^2$ have been published by Fiore et al. (
+2X005).,2000b).
+" This approach is colmplemecutary to deep xicil beam survevs (~0.1 deg?, sce oe, Mushotzky et al."," This approach is complementary to deep pencil beam surveys $\sim0.1$ $^2$, see e.g. Mushotzky et al."
+ 2000. Toruschemieier e al.," 2000, Hornschemeier et al."
+ 2000). aud we cover a different portio Lot the redshiftIuuinositv pane.," 2000), and we cover a different portion of the redshift–luminosity plane."
+ Our purpose is to study cosmic source poptlatious at fluxes where a reasonably large fractio1 of he hard XRD is resolved at he BeppoSAN flux Μι). but where the X-ray tax is high enoug ito provide X-ray s)octrali formation in higher seusiivity observations.," Our purpose is to study cosmic source populations at fluxes where a reasonably large fraction of the hard XRB is resolved at the BeppoSAX flux limit), but where the X-ray flux is high enough to provide X-ray spectral information in higher sensitivity follow-up observations."
+ This would alow tιο determination of the distribution of absorbing colunus in the sources making the uud NRB. providing strong constraints on ACN svuthesis models for the NRB.," This would allow the determination of the distribution of absorbing columns in the sources making the hard XRB, providing strong constraints on AGN synthesis models for the XRB."
+" Furthermore. aree area survevs allow tre search for previously “raro AGN, like red quΠωνα OY OOther miünoritv ACN popuations (Tsim Elvis 1999) alc quantify their fractional coitiibution to the ACN family"," Furthermore, large area surveys allow the search for previously `rare' AGN, like `red' quasars or other minority AGN populations (Kim Elvis 1999) and quantify their fractional contribution to the AGN family."
+" Finally. at our N-raw flux inits the optical counterpars are bright enough to allw relatively ligh quality optical spectroscopy. 1seful to investigate the plysics of t10 πο]ος,"," Finally, at our X-ray flux limits the optical counterparts are bright enough to allow relatively high quality optical spectroscopy, useful to investigate the physics of the sources."
+ The Hieh Euerev Large Arca Survey (HELLAS) has σοι performed iu the keV baud because: a) this is the band closest to the imiaxinnuu of the NRB energv deusitv which is reachable with the cient inaeime N-rav telescopes. aud b) the BeppoSAX MECS Point Spread Function (PSF) ercathy improves with enerev. xovidiugc» a 95 error radius of 1 in the hard baud (Fiore et al.," The High Energy Large Area Survey (HELLAS) has been performed in the 4.5-10 keV band because: a) this is the band closest to the maximum of the XRB energy density which is reachable with the current imaging X-ray telescopes, and b) the BeppoSAX MECS Point Spread Function (PSF) greatly improves with energy, providing a 95 error radius of $'$ in the hard band (Fiore et al."
+ 2000a). allowine optical+ identification+p. of the N-rav- sources.," 2000a), allowing optical identification of the X-ray sources."
+ About 80* dee?J ofHd skv have been surveved so far using 112 DeppoSAXN MECS fielda at |b]>20 deg. (, About 80 $^2$ of sky have been surveyed so far using 142 BeppoSAX MECS fielda at $|b|>20$ deg. (
+Fields ceutered on bright extended sotrees and bright Galactic sources were excluded from the survey. as well as fleds close to LAIC. SAIC and M32.),"Fields centered on bright extended sources and bright Galactic sources were excluded from the survey, as well as fields close to LMC, SMC and M33.)"
+ A robust detection algoritlin has been used in coadded MECSI|MECS2MECS:?) (or MECS2MECS3 for observations performed after f1ο failure of the MECSI uut) 15-10 keV images., A robust detection algorithm has been used in coadded MECS1+MECS2+MECS3 (or MECS2+MECS3 for observations performed after the failure of the MECS1 unit) 4.5-10 keV images.
+" The method cousists iu first cowolving the ταν image with a swavelkt fiction. to sinooth the nuage8 and icreasc| contrast, and then iu runniic a standard slide-cell detection method ou the smoothed image. to locate cot excesses above the local ckerouud."," The method consists in first convolving the X-ray image with a wavelet function, to smooth the image and increase contrast, and then in running a standard slide-cell detection method on the smoothed image, to locate count excesses above the local background."
+ The statistics o ταach candidate detection is theji aceuratelv studied au the final net counts are estimated from the original (1u-noothed) image to preserve Poissou statistics., The statistics of each candidate detection is then accurately studied and the final net counts are estimated from the original (un-smoothed) image to preserve Poisson statistics.
+ The lxiwkeround is calculated usine source-free boxes near tje source region., The background is calculated using source-free boxes near the source region.
+ The detecnπι has heen rua several t11ο or each feld. chaugiug the size of the wavelet function. to a) take iuto acσοιit the variation of tle MECS PSF with the offaxis angle. and b) to detect efficieitly sources with variale extension.," The detection has been run several time for each field, changing the size of the wavelet function, to a) take into account the variation of the MECS PSF with the offaxis angle, and b) to detect efficiently sources with variable extension."
+ The quality of the detection has always, The quality of the detection has always
+systtems. while it plays a very substantial role in close binaries.,"tems, while it plays a very substantial role in close binaries."
+resolution in order to detect new svstems. since al such distances (he pairs wilh semi-major orbital axes [rom 10 to 100 AU are hard to detect both spectroscopically and visually.,"resolution in order to detect new systems, since at such distances the pairs with semi-major orbital axes from 10 to 100 AU are hard to detect both spectroscopically and visually."
+" Using the two following criteria. yo>0.26""/vear and [m/IHI]<—1. we tried to rule out the thick and thin dise stars."," Using the two following criteria, $\mu>0.26\ ''/$ year and $\mathrm{[m/H]}<-1$, we tried to rule out the thick and thin disc stars."
+ In case of the thick dise though. these limits are not stringent and some objects may. belong to the metal-weak tail of the thick dise 2005)..," In case of the thick disc though, these limits are not stringent and some objects may belong to the metal-weak tail of the thick disc \citep{arifyanto}."
+ To separate the halo stars [rom the dise stus in our sample. we used a formal method. described in the appendix of Grether&Lineweaver(2007).," To separate the halo stars from the disc stars in our sample, we used a formal method, described in the appendix of \citet{grether_lineweaver_2007}."
+. The equations establishing the probability that a star belongs to the Chin cise (2444). the Chick disc (Pu) or the halo Gas ) are where The input parameters for these equations are adopted from Robinetal.(2003).. and presented in Table 1..," The equations establishing the probability that a star belongs to the thin disc $P_{\rm thin}$ ), the thick disc $P_{\rm thick}$ ) or the halo $P_{\rm halo}$ ) are where The input parameters for these equations are adopted from \citet{robin_2003}, and presented in Table \ref{populations_properties}."
+ We rejected three out of 221 svstems: (wo double stars G99-48 and, We rejected three out of 221 systems: two double stars G99-48 and
+Chromospherically Active Binaries (CAB) are the class of detached binary systems with spectral types later than FE characterized by a strong chromospheric. transition region. and coronal activity.,"Chromospherically Active Binaries (CAB) are the class of detached binary systems with spectral types later than F characterized by a strong chromospheric, transition region, and coronal activity."
+" Enhanced. emission. cores of Ca LL Η and ly. and sometimes in the Balmer £4, linc are primary indicators of the chromospheric activity and often accompanied by photometric variability due to starspots."," Enhanced emission cores of Ca II H and K, and sometimes in the Balmer $H_{\alpha}$ line are primary indicators of the chromospheric activity and often accompanied by photometric variability due to starspots."
+ The first. published. catalog of CAB by Strassmeicr ct al. (, The first published catalog of CAB by Strassmeier et al. (
+1988) contained the classical RS CVn systems defined hy Lall (1976) and the BY Draconis-twpe binaries defined. by opp Γοκο (1977).,1988) contained the classical RS CVn systems defined by Hall (1976) and the BY Draconis-type binaries defined by Bopp Fekel (1977).
+ The 168 CAB in the first catalog have been increased to 206 in the second catalog by Strassmeicr et al. (, The 168 CAB in the first catalog have been increased to 206 in the second catalog by Strassmeier et al. (
+1993).,1993).
+ Ehe third catalog has not vet been published but Eker privately continued to collect CAB., The third catalog has not yet been published but Eker privately continued to collect CAB.
+ The unpublished list of Eker. which became the main database for this study. nowadays contains about 280 CAL.," The unpublished list of Eker, which became the main database for this study, nowadays contains about 280 CAB."
+ At the first attempt. Eker (1992). was rather unsuccessful in breaking up the 146 CAB sample into kinematically distinct sub. samples.," At the first attempt, Eker (1992), was rather unsuccessful in breaking up the 146 CAB sample into kinematically distinct sub samples."
+ Containing spectral tvpes from Foto M. ancl luminosity classes from V. to Il. the already heterogeneous CAB sample was found. to be heterogeneous also in the sense that the kinematically vounger and older systems exist among the evolved. binaries with at least one component of à giant or a sub giant. and among the un-evolved main sequence binaries.," Containing spectral types from F to M and luminosity classes from V to II, the already heterogeneous CAB sample was found to be heterogeneous also in the sense that the kinematically younger and older systems exist among the evolved binaries with at least one component of a giant or a sub giant, and among the un-evolved main sequence binaries."
+ The llipparcos data. (Derryman ct al., The Hipparcos data (Perryman et al.
+ 1997). was not available to Eker (1992)., 1997) was not available to Eker (1992).
+ llowever. Aslan et al. (," However, Aslan et al. ("
+1999). who used the Hipparcos proper motions ancl parallaxes of 178 CAB.,"1999), who used the Hipparcos proper motions and parallaxes of 178 CAB,"
+ER I radio galaxies (defined by edge-darkened radio structure) have lower radio power than FR II galaxies (defined by edge-brightened radio structure due to compact Jet-terminating hot spots) (Fanaroff&Riley1974)..,FR I radio galaxies (defined by edge-darkened radio structure) have lower radio power than FR II galaxies (defined by edge-brightened radio structure due to compact jet-terminating hot spots) \citep{1974MNRAS.167P..31F}.
+ Relativistic| jets are observed in many radio-loud active galactic nuclei (AGNs)., Relativistic jets are observed in many radio-loud active galactic nuclei (AGNs).
+ In the frame of unification schemes of radio-loud AGNs. FR I radio galaxies are believed to be misaligned BL Lacertae (BL Lac) objects. and FR II radio galaxies correspond to misaligned radio quasars (seeUrry&Padovani1995.fora review)..," In the frame of unification schemes of radio-loud AGNs, FR I radio galaxies are believed to be misaligned BL Lacertae (BL Lac) objects, and FR II radio galaxies correspond to misaligned radio quasars \citep*[see][for a review]{1995PASP..107..803U}."
+ Most BL Lac objects have featureless optical and ultraviolet continuum spectra. and only a small fraction of BL Lae objects show very weak broad emission lines. while quasars usually have strong broad-line emission.," Most BL Lac objects have featureless optical and ultraviolet continuum spectra, and only a small fraction of BL Lac objects show very weak broad emission lines, while quasars usually have strong broad-line emission."
+ The broad emission lines of quasars are produced by distant gas clouds in broad-line regions (BLR). which are photo-tonized by the optical/UV continua radiated from the accretion disks surrounding massive black holes.," The broad emission lines of quasars are produced by distant gas clouds in broad-line regions (BLR), which are photo-ionized by the optical/UV continua radiated from the accretion disks surrounding massive black holes."
+ The difference of the broad-line emission between radio-loud quasars and BL Lac objects may be attributed to their different central engines (e.g..Cavaliere&D'Elia2002:Cao 2003)..," The difference of the broad-line emission between radio-loud quasars and BL Lac objects may be attributed to their different central engines \citep*[e.g.,][]{2002ApJ...571..226C,2002ApJ...570L..13C,2003ApJ...599..147C}."
+ The unified scheme of BL Lac objects and FR I radio galaxies have been extensively explored by many previous authors with different approaches. such as the comparisons of the spectral energy distributions (SEDs) in. different wavebands (e.g..Owenetal.1996;Capetti&Lee 2001). the radio morphology. (e.g..Marchaetal.2005).. Padova," The unified scheme of BL Lac objects and FR I radio galaxies have been extensively explored by many previous authors with different approaches, such as the comparisons of the spectral energy distributions (SEDs) in different wavebands \citep*[e.g.,][]{1996AJ....111...53O,2000MNRAS.318..493C,2001ApJ...548..244B}, the radio morphology, \citep*[e.g.,][]{2005MNRAS.361..469M}."
+ni&Urry(1992) derived the radio LFs of flat spectrum radio quasars (FSRQs) and FR II galaxies from a sample of radio-loud AGNs.," \citet{1992ApJ...387..449P}
+ derived the radio LFs of flat spectrum radio quasars (FSRQs) and FR II galaxies from a sample of radio-loud AGNs."
+ They considered a two-component model. in which the total luminosity is the sum of an unbeamed part and a beamed jet luminosity. (," They considered a two-component model, in which the total luminosity is the sum of an unbeamed part and a beamed jet luminosity. ,"
+SSRQs). which strengthens the unification of FR II galaxies and radio quasars (seePadovani&Urry1992.forthe details).," which strengthens the unification of FR II galaxies and radio quasars \citep*[see][for the
+details]{1992ApJ...387..449P}."
+. Similar analyses were carried out on the relation between FR I galaxies and BL Lae objects (Padovani&Urry1991:Padovani1995).. which is also consistent with the unification of FR Is and BL Lac objects. Marchàetal.2005:Landt&Bignall 2008).. Landt&," Similar analyses were carried out on the relation between FR I galaxies and BL Lac objects \citep{1991ApJ...368..373P,1995PASP..107..803U}, which is also consistent with the unification of FR Is and BL Lac objects. \citep*[e.g.,][]{2005MNRAS.361..469M,2008MNRAS.391..967L}."
+Bignall(2008) found that a considerable number of BL Lac objects can be identified with the relativistically beamed counterparts of FR II radio galaxies in a sample of BL Lac objects selected from the Deep X-ray Radio Blazar Survey (DXRBS)., \citet{2008MNRAS.391..967L} found that a considerable number of BL Lac objects can be identified with the relativistically beamed counterparts of FR II radio galaxies in a sample of BL Lac objects selected from the Deep X-ray Radio Blazar Survey (DXRBS).
+ Ledlow&Owen(1996) found that FR I and FR II radio galaxies can be clearly divided in the host galaxy optical lummosity—-radio luminosity (Mr-Lraw) plane. by a dividing line showing that radio power is proportional to the optical luminosity of the host galaxy.," \citet{1996AJ....112....9L} found that FR I and FR II radio galaxies can be clearly divided in the host galaxy optical luminosity–radio luminosity $M_R$ $L_{\rm Rad}$ ) plane, by a dividing line showing that radio power is proportional to the optical luminosity of the host galaxy."
+ What causes the FR ΕΕΚ IL division is still unclear. 2007).," What causes the FR I/FR II division is still unclear, ."
+ Ghisellini&Celotti(2001)used the optical luminosity of, \citet{2001A&A...379L...1G} used the optical luminosity of
+of degrees of freedom (in this case. I8 data points minus 2 litted poaranmeers.),"of degrees of freedom (in this case, 18 data points minus 2 fitted parameters.)"
+" Vhis fit vields Puerequator=0.513190.00007 clays. and Prog,=WILY days."," This fit yields $P_{\mbox{equator}}=0.51319\pm 0.00007$ days, and $P_{\mbox{beat}}=109\pm 9$ days."
+ Lhe parameters Zgaior 2d Zo ave strongly correlated., The parameters $P_{\mbox{equator}}$ and $P_{\mbox{beat}}$ are strongly correlated.
+ “Phe dillerential rotation parameters for the 1995 December dataset. obtained by Donati&Col-lierCameron (1997).. lie within the inner contour. close to the major axis of the error ellipse. indicating excellent agreement between the two methods.," The differential rotation parameters for the 1995 December dataset, obtained by \scite{donati97doppler}, lie within the inner contour, close to the major axis of the error ellipse, indicating excellent agreement between the two methods."
+" A re-analysis of the 1996 December data using the sheared-image method. of Donatietal.(2000). gives Loquator=0.51332+0.00012 days. and Z4,=115415 days. again in good agreement."," A re-analysis of the 1996 December data using the sheared-image method of \scite{donati2000rxj} gives $P_{\mbox{equator}}=0.51332\pm 0.00012$ days, and $P_{\mbox{beat}}=115\pm 15$ days, again in good agreement."
+ The scatter about the fit to the data shown in Fig., The scatter about the fit to the data shown in Fig.
+ 5 is considerably ereater than the sizes of the errors on the data points.," \ref{fig:diffrot}
+ is considerably greater than the sizes of the errors on the data points."
+ Lhe measured. periods of some of the best-observed spots deviate from the mean relation by as much as 0.005 dav. giving an accumulated error of order an hour over the 6-dav span of the data.," The measured periods of some of the best-observed spots deviate from the mean relation by as much as 0.005 day, giving an accumulated error of order an hour over the 6-day span of the data."
+ This is too great to be ascribed to any plausible source of svstematic error. such as irregularly-shaped spot. &roups spanning a range of LIatitudes.," This is too great to be ascribed to any plausible source of systematic error, such as irregularly-shaped spot groups spanning a range of latitudes."
+ We are forced. to conclude that the scatter in rotation rates [or spots at any given latitude must be intrinsic to the spots themselves., We are forced to conclude that the scatter in rotation rates for spots at any given latitude must be intrinsic to the spots themselves.
+ lig., Fig.
+ 7 may provide a clue to the physical origin of this scatter., \ref{fig:resid} may provide a clue to the physical origin of this scatter.
+ Here we plot the residuals in rotational velocity measured relative to the mean fit. as a function of spot size.," Here we plot the residuals in rotational velocity measured relative to the mean fit, as a function of spot size."
+ The scatter about the mean relation. appears to be considerably: ercater for the small spots than for the larger spots., The scatter about the mean relation appears to be considerably greater for the small spots than for the larger spots.
+ This size-depencent scatter in the rotation rates of individual spots suggests that the smaller. spots may be more susceptible to random bulleting bv supergranular Lows (sce. ee Llowarel.Gilman& 1984)).," This size-dependent scatter in the rotation rates of individual spots suggests that the smaller spots may be more susceptible to random buffeting by supergranular flows (see, e.g. \pcite{howard84}) )."
+ Phere does not. however. appear to be any significant dilference between the average rotation rate of the large spots (i.e. those having Mo0.003 kms D. and that of the smaller spots.," There does not, however, appear to be any significant difference between the average rotation rate of the large spots (i.e. those having $W>0.003$ km $^{-1}$ ), and that of the smaller spots."
+ The results. of this investigation suggest that spectral subtraction. least-squares cleconvolution and. matched-filter analysis provide a powerful diagnostic method. for identifying large numbers of individual starspot signatures in trailed spectrograms of AB Doraclus and. by implication. other similar stars.," The results of this investigation suggest that spectral subtraction, least-squares deconvolution and matched-filter analysis provide a powerful diagnostic method for identifying large numbers of individual starspot signatures in trailed spectrograms of AB Doradus and, by implication, other similar stars."
+ The matchec-llterl analysis allows the positions ancl rotation periods.l of individual spots to be determined to high. precision. ancl so. provides a powerful alternative to methods involving regularised image reconstruction.," The matched-filter analysis allows the positions and rotation periods of individual spots to be determined to high precision, and so provides a powerful alternative to methods involving regularised image reconstruction."
+ We have successfully measured the rotation periods of individual spots at latitudes ranging from the stellar equator up to a latitude of 55°., We have successfully measured the rotation periods of individual spots at latitudes ranging from the stellar equator up to a latitude of $^{\circ}$ .
+ The dilferential rotation law fitted to the larger spots shows excellent agreement with the mean relation obtained by Donati&CollierCameron(1997)., The differential rotation law fitted to the larger spots shows excellent agreement with the mean relation obtained by \scite{donati97doppler}.
+. Phe rate of shear in the surface rotation speed as a function of απο appears to be very. similar to the Sun. with the equator pulling one full turn ahead of the pole every 110 claws OL SO.," The rate of shear in the surface rotation speed as a function of latitude appears to be very similar to the Sun, with the equator pulling one full turn ahead of the pole every 110 days or so."
+ The rotation rates of the smaller spots at the low and intermediate latitudes studied here appear to have a larger intrinsic scatter than the large spots., The rotation rates of the smaller spots at the low and intermediate latitudes studied here appear to have a larger intrinsic scatter than the large spots.
+ Further observations of similar or better precision would. be desirable to verily whether this size dependence is present., Further observations of similar or better precision would be desirable to verify whether this size dependence is present.
+ H£ confirmed. this result could give a first glimpse of the turbulent sub-surface velocity field in à late-tvpe star other than the Sun.," If confirmed, this result could give a first glimpse of the turbulent sub-surface velocity field in a late-type star other than the Sun."
+ We thank Ixeith Lorne for helpful conversations ancl useful comments on an early version of the manuscript., We thank Keith Horne for helpful conversations and useful comments on an early version of the manuscript.
+ An anonvmous referee mace several useful suggestions [for enhancing the clarity of the paper. and for improving the treatment of limb darkening.," An anonymous referee made several useful suggestions for enhancing the clarity of the paper, and for improving the treatment of limb darkening."
+ ACC acknowledges the support of a PPARC Senior Research Fellowship during the course of this work., ACC acknowledges the support of a PPARC Senior Research Fellowship during the course of this work.
+ This paper is based on observations nace using the 3.9-m Anelo-AXustralian Telescope., This paper is based on observations made using the 3.9-m Anglo-Australian Telescope.
+ Phe project made use of support software and data analysis facilities provided by the Starlink Project which is runby CCLRC on behalf of PPARC., The project made use of support software and data analysis facilities provided by the Starlink Project which is runby CCLRC on behalf of PPARC.
+Since the sample is absolute magnitude lnuited with My<6.5 at 100 pc. ic. well below the πιο huninosity. it shows no bias against metalaxich aud super metal-rich (SAMIR) stars (Crenon 1987)).,"Since the sample is absolute magnitude limited with $M_{\mbox{v}} < 6.5$ at 100 pc, i.e. well below the turn-off luminosity, it shows no bias against metal-rich and super metal-rich (SMR) stars (Grenon \cite{grenon2}) )."
+ It is also properanotion limited aud lence tangential velocity limited. ie. dans *owith distance d iu [pc].," It is also proper-motion limited and hence tangential velocity limited, i.e. $V_{t} > 0.85\cdot d$ km $^{-1}$ with distance $d$ in [pc]."
+ Our sky coverage à<[107 meludes im particulara complete south ealactic cap defined bv bcΠΡ, Our sky coverage $\delta < +10\degr$ includes in particulara complete south galactic cap defined by $b < -53\degr$.
+" For the WIPPARCOS inissiou separation and for the obteution of physical parameters frou, plotometiry. all progrguuue stars south of 6=|10° were observed at the Swiss telescope at La Silla. from 1981 on."," For the HIPPARCOS mission preparation and for the obtention of physical parameters from photometry, all programme stars south of $\delta=+10\degr$ were observed at the Swiss telescope at La Silla, from 1981 on."
+ A total of 9135 measurements were accummilated for 5113 different stars., A total of 39435 measurements were accumulated for 5443 different stars.
+" Calibrations of the Geneva plotometric svstem aud sushbsequent revisions by Crenon (1977)) were used to derive T,4p Me and [AZ]. with precisions varving with the temperature and gravitv ranges."," Calibrations of the Geneva photometric system and susbsequent revisions by Grenon \cite{grenon1}) ) were used to derive $T_{\small{\mbox{eff}}}$, $M_{v}$ and $[M/H]$, with precisions varying with the temperature and gravity ranges."
+ For late G dwarfs. the internal errors are of 20- LOTS on Τομ. 0.03-0.05 dex on [UIT] aud 0.15 on AY. for single stars.," For late G dwarfs, the internal errors are of 20-40 K on $T_{\small{\mbox{eff}}}$, 0.03-0.05 dex on $[M/H]$ and 0.15 on $M_{v}$, for single stars."
+ Stars later than K3V have no AL{TE estimate., Stars later than K3V have no $[M/H]$ estimate.
+ Radial velocities. were obtained for all programuuc stars bv Crenon with CORAVEL at Taute-Proveuce Observatory in the declination zone 6= -20° to 10° and as part of a dedicated ESO Kev Progranune (Alavor et al. 1998)), Radial velocities were obtained for all programme stars by Grenon with CORAVEL at Haute-Provence Observatory in the declination zone $\delta=$ $\degr$ to $\degr$ and as part of a dedicated ESO Key Programme (Mayor et al. \cite{mayor2}) )
+ south of 8= -207., south of $\delta=$ $\degr$.
+ A few velocities were taken frou literature. im particular from the Carney ct al. (19913).," A few velocities were taken from literature, in particular from the Carney et al. \cite{carney}) )"
+ and from the Darbier-Drossat et al. (199 1)), and from the Barbier-Brossat et al. \cite{barbier}) )
+ catalogues., catalogues.
+ Dinaries were ideutified either from the radial velocity changes or the comparison between trigonometric and photometric parallaxcs. or detected by TIPPARCOS For the preseut study. stars were kept if their HIPPARCOS aud photometric parallaxes were larger thui LO ias.," Binaries were identified either from the radial velocity changes or the comparison between trigonometric and photometric parallaxes, or detected by HIPPARCOS For the present study, stars were kept if their HIPPARCOS and photometric parallaxes were larger than 10 mas."
+" Moreover. unucertai parallaxes with standard errors >5r, amas were rejected."," Moreover, uncertain parallaxes with standard errors $ >
+5$ mas were rejected."
+ Iu our final sample. space velocities are available for 11123 stars iui overall metallicities for 2619 of thoai.," In our final sample, space velocities are available for 4143 stars and overall metallicities for 2619 of them."
+ The old disc stars of our sample show larecly positive ean « niofions. confirmune the trend already appareut in the earlier results recalled iu the Introduction.," The old disc stars of our sample show largely positive mean $\overline{u}$ motions, confirming the trend already apparent in the earlier results recalled in the Introduction."
+" At given ecco30 hans +. there is no bias in favour of juward or outward motions when stars are observed iu the south galactic cap or towards the auti-rotation where ΈτVu?que,"," At given $v < -30$ km $^{-1}$, there is no bias in favour of inward or outward motions when stars are observed in the south galactic cap or towards the anti-rotation where $V_{t}=\sqrt{u^{2}+w^{2}}$."
+ With iucreasine distauces the fraction of small [ü| decreases in favour of large 01. 0/— velocities.," With increasing distances the fraction of small $|u|$ decreases in favour of large $u+$, $u-$ velocities."
+ When old dise stars are selected in the z range from 10 to LO uas. the unbalance between & and s becomes strickine. see Fie. 1.," When old disc stars are selected in the $\pi$ range from 10 to 40 mas, the unbalance between $u+$ and $u-$ becomes stricking, see Fig. \ref{fig:Ubh},"
+ and reaches up to [50 lan ! towards the auti-rotation., and reaches up to $+50$ km $^{-1}$ towards the anti-rotation.
+ When subsamples are complete according to e components (c.g. all stars with e<50 lins * are retained). the 4 anomaly is correctly estimated with values within the rauge |29+2 kins + wat.," When subsamples are complete according to $v$ components (e.g. all stars with $v < -50$ km $^{-1}$ are retained), the $u$ anomaly is correctly estimated with values within the range $+29 \pm 2$ km $^{-1}$ w.r.t."
+ the Sun. Le. (19 +£9)kms towort.," the Sun, i.e. $+19 \pm 9$ km $^{-1}$ w.r.t."
+ the Galactic Center (GC) after correction for local motions inferred from Iuijkeu Tremaine (19911)., the Galactic Center (GC) after correction for local motions inferred from Kuijken Tremaine \cite{KuijTre}) ).
+ Stars with metallicities atypical for the voung local disc. io. with [AL/I7|20.25 or <(0.30. forma the best test sample to investigate larec scale —perturbations of stellar orbits.," Stars with metallicities atypical for the young local disc, i.e. with $[M/H] >
+0.25$ or $ < -0.30$, form the best test sample to investigate large scale perturbations of stellar orbits."
+ The metal poor group should contain stars born inside aud outside the solar orbit. whereas the iietalaich eroup should consist of stars born uniquolv within the solar orbit.," The metal poor group should contain stars born inside and outside the solar orbit, whereas the metal-rich group should consist of stars born uniquely within the solar orbit."
+ It is therefore well suited to est the orbital anomalies iu the immer disc., It is therefore well suited to test the orbital anomalies in the inner disc.
+ Its (a.0) plane is displaved i- Fie. 2..," Its $(u,v)$ plane is displayed in Fig. \ref{fig:planUV}."
+ Tf we consider the (o.0) plane for all stars with kuown [A£/H|. the area >10 kimi t and 130c<30 ans *1 contains not only metal-rich stars but a sample of old dise stars with metallicities covering the whole rauge of LUIT| observed in the old disc population from ~0.6 dex to [0.6 dex aud showing a peak at solar metallicity (Fig. 3)).," If we consider the $(u,v)$ plane for all stars with known $[M/H]$ , the area $u > 10$ km $^{-1}$ and $ -130 < v < -30$ km $^{-1}$ contains not only metal-rich stars but a sample of old disc stars with metallicities covering the whole range of $[M/H]$ observed in the old disc population from $\sim -0.6$ dex to $+0.6$ dex and showing a peak at solar metallicity (Fig. \ref{fig:mhhsg}) )."
+ The [MIT| distribution is essenutiallv the same, The $[M/H]$ distribution is essentially the same
+"in terms of brenisstralhluns enussion. uormalized to 0.5 MeV in oder to compare it to ICS. is ~2&105 3, which is 20/100 times the umuber of fast electrons needed for ICS. assuming a single power-law distribution with à=3.","in terms of bremsstrahlung emission, normalized to $0.5$ MeV in oder to compare it to ICS, is $\sim2\times 10^6$ $^{-3}$ , which is $-$ 100 times the number of fast electrons needed for ICS, assuming a single power-law distribution with $\delta=3$."
+ The ratio likewise holds for the cucrey content of fast electrons., The ratio likewise holds for the energy content of fast electrons.
+ The iudex of the clectron distribution function must lie in the ranee 6=2 3 to account for the continu s-rav enüssion di terms of ICS., The index of the electron distribution function must lie in the range $\delta=2-$ 3 to account for the continuum $\gamma$ -ray emission in terms of ICS.
+ We conclude our discussion of the continu s-ray eveuts bv sugeesting that ICS of photospheric photons on relativistic electrous can account for the observed emission., We conclude our discussion of the continuum $\gamma$ -ray events by suggesting that ICS of photospheric photons on relativistic electrons can account for the observed emission.
+ It is energeticallv more favorable than brenisstrahlung by a factor of 20 (isotropic) to 100 (pancake anisotropy)., It is energetically more favorable than bremsstrahlung by a factor of 20 (isotropic) to 100 (pancake anisotropy).
+ Au intriguing example of a coronal IEXNR source for which ICS may have plaved a role is that reported byὃ, An intriguing example of a coronal HXR source for which ICS may have played a role is that reported by.
+ν A powerful flare occurred in AR 10150 when it was 10 behind the limb on 2002 October 27. as viewed from Earth.," A powerful flare occurred in AR 10180 when it was $40^\circ$ behind the limb on 2002 October 27, as viewed from Earth."
+ The flare was observed directly frou. Maus by the Gamuna-Rav Spectrometer on board theOdvssev inission., The flare was observed directly from Mars by the Gamma-Ray Spectrometer on board the mission.
+ Ou this basis it was estimated to be colmparable in SKR. class to the flares ou 2003 October 28 and 2005 January 20 ciseussed iu the previous section., On this basis it was estimated to be comparable in SXR class to the flares on 2003 October 28 and 2005 January 20 discussed in the previous section.
+ A lavee. diffuse IINR source was observed by RITESSI.," A large, diffuse HXR source was observed by RHESSI."
+" Both the footpoiuts and the thermal flare loops were occulted for this event - iu fact. the source was high iu the corona: the radial occultation height is 3007+65"" and the source ceutroid was 580"" above lah. implying a source height of ~0.3 0.1 νι."," Both the footpoints and the thermal flare loops were occulted for this event - in fact, the source was high in the corona: the radial occultation height is $300''\pm 65''$ and the source centroid was $\approx 80''$ above limb, implying a source height of $\sim 0.3-$ 0.4 $_\odot$."
+ The IIXR spectrum was observed up to 60 keV. The photon spectral iudex was relatively hard. chaugiug svstematically frou 5 during the exponential decay phase of the event. niplviug an electron spectral iudex 6=2.5 3.," The HXR spectrum was observed up to 60 keV. The photon spectral index was relatively hard, changing systematically from $-$ 3 during the exponential decay phase of the event, implying an electron spectral index $\delta=2.5-$ 3."
+ Based ou Fig., Based on Fig.
+ 3 in7.. in which the IIXR. spectra is preseuted for a inue during the decay phase. we estimate the photon fux at 30 keV to be O1 photons tem? 1.," 3 in, in which the HXR spectrum is presented for a time during the decay phase, we estimate the photon flux at 30 keV to be 0.1 photons $^{-1}$ $^{-2}$ $^{-1}$."
+ When interpreted in terms of nou-thermal brensstralluus cluission. the authors estimate that the instantaneous nuuber density of electrons >10 keV required ~10° cmnον or ~104 of the total electron umber deusity or un ambicut plasima denstv of ~LO* °.," When interpreted in terms of non-thermal bremsstrahlung emission, the authors estimate that the instantaneous number density of electrons $>10$ keV required $\sim 10^7$ $^{-3}$, or $\sim10$ of the total electron number density for an ambient plasma density of $\sim 10^8$ $^{-3}$."
+" The ower limit to the total euergv in non-thermal electrous ix 2<10°? ores,", The lower limit to the total energy in non-thermal electrons is $2\times 10^{29}$ ergs.
+ Is ICS a viable alternative for this event?, Is ICS a viable alternative for this event?
+ For the case of ultra-relativistie electrons the photon spectral index of a=3.1 nuplies an clectrom spectral iudex is 6=5.2., For the case of ultra-relativistic electrons the photon spectral index of $\alpha= 3.1$ implies an electron spectral index is $\delta=5.2$.
+ When normalized to 0.5 MeV. a lavee uunber of energetie electrons is required. ~ὃνLOT emi).> containing 1075 oes.," When normalized to 0.5 MeV, a large number of energetic electrons is required, $\sim 3\times 10^7$ $^{-3}$, containing $10^{31}$ ergs."
+ We note. however. that the electrous involved in IC in this case have cucreics Z20 MeV. The πα Srequired nuuber deusitv of electrous with energies 220 MeV is far more modest. Hag~LO 7. containing ~2&1075 eres.," We note, however, that the electrons involved in ICS in this case have energies $\gtrsim 20$ MeV. The minimum required number density of electrons with energies $>20$ MeV is far more modest, $n_{min}\sim 10$ $^{-3}$, containing $\sim2\times10^{26}$ ergs."
+ Au anisotropic electron distribution brings tljose estimates down vet further., An anisotropic electron distribution brings these estimates down yet further.
+ Of course. ICS depoeuls entirely on the details of the electron distribution fiuction.," Of course, ICS depends entirely on the details of the electron distribution function."
+ If the electron distribution is described by a single power law (6= 5.2) that extends below ~2 MeV. then it becomes cherectically unfavorable compared to brenisstraliluug.," If the electron distribution is described by a single power law $\delta=5.2$ ) that extends below $\sim 2$ MeV, then it becomes energetically unfavorable compared to bremsstrahlung."
+ Iu the context of a double power law with 4;-=5.2. the break energy - or low energy. cutoff - would need to be a few MeV in order for ICS to be relevant.," In the context of a double power law with $\delta_U=5.2$, the break energy - or low energy cutoff - would need to be a few MeV in order for ICS to be relevant."
+ Consideration of an anisotropic electron distribution cau bring the cutoff down to £yz500 keV. We now consider an alternative to ICS iu the relativistic reguune aud iustead consider a nmüldlv relativistic distribution of electrons., Consideration of an anisotropic electron distribution can bring the cutoff down to $E_b\lesssim 500$ keV. We now consider an alternative to ICS in the relativistic regime and instead consider a mildly relativistic distribution of electrons.
+ Tu 822 we showed that the up-scattered photon spectrum resulting frou iildlv relativistic electrons is significautlv steeper thu that from the classic relation (6|13/2 at simall values of εοει., In 2 we showed that the up-scattered photon spectrum resulting from mildly relativistic electrons is significantly steeper than that from the classic relation $(\delta+1)/2$ at small values of $\epsilon_2/\epsilon_1$.
+ Iu the mnülkdlv relativistic reguue. however. mono-cherectic photons are no longer appropriate as an approximation for the incident photon spectrum.," In the mildly relativistic regime, however, mono-energetic photons are no longer appropriate as an approximation for the incident photon spectrum."
+ Therefore. as a crude estimate. we use the simulated flare spectrum produced by the CUIANTI software package based on an M2 flare in the EUV/SXNR. energy range (0.1 to 20 keV. see the left panel of Fig. 11))," Therefore, as a crude estimate, we use the simulated flare spectrum produced by the CHIANTI software package based on an M2 flare in the EUV/SXR energy range (0.1 to 20 keV, see the left panel of Fig. \ref{fig:mrel_chianti}) )"
+ aud scale the photon nuniber deusity up to represent a large N-class flare., and scale the photon number density up to represent a large X-class flare.
+ The EUV/SXR spectrum declines rapidly with euergv above 12 keV aud so the redistribution of photons by dowu-scatter can be neglected for the ITXR euerev range in question., The EUV/SXR spectrum declines rapidly with energy above $1-2$ keV and so the redistribution of photons by down-scatter can be neglected for the HXR energy range in question.
+ After some experimentation. we find that a double power-law electron energy distribution with an upper spectral iudex of é;-=3.8 produces au upescattered IINR spectrum with the observed photon spectral index of a=3.1 (Fig. LL.," After some experimentation, we find that a double power-law electron energy distribution with an upper spectral index of $\delta_U=3.8$ produces an up-scattered HXR spectrum with the observed photon spectral index of $\alpha=3.1$ (Fig. \ref{fig:mrel_chianti},"
+ blue solid. curve)., blue solid curve).
+ The electron. distribution is assumed to be isotropic at all energies., The electron distribution is assumed to be isotropic at all energies.
+ The spectrum has a break energy at 300 keV with a flatter spectral iudex at lower energies (dg=1.5)., The spectrum has a break energy at 300 keV with a flatter spectral index at lower energies $\delta_L=1.5$ ).
+ Au ambicut deusitv of 105 cin? las been assiuued aud he παπιου density of energetic electrons with ὅτι=3.8 in 3«107 7. normalized to 0.5 MeV. coutaining E1075 ergs - a factor zz50 less than that inuplied by the xenisstrahluue interpretation.," An ambient density of $10^8$ $^{-3}$ has been assumed and the number density of energetic electrons with $\delta_U=3.8$ is $3\times 10^3$ $^{-3}$, normalized to 0.5 MeV, containing $\approx 4\times 10^{27}$ ergs - a factor $\approx 50$ less than that implied by the bremsstrahlung interpretation."
+ The electron distribution vields EIB broiisstraliluung shown as the red solid curve in Fie. LL., The electron distribution yields EIB bremsstrahlung shown as the red solid curve in Fig. \ref{fig:mrel_chianti}.
+ The EIB photon spectrin is siguificautlv Hatter than the ICS spectrum. aud becomes comparable ο ICS at photon energies of 260 keV. The dashed lines show the ICS and EIB cussion from a coue-hbeam (half widtha of 107) distribution of electrons.," The EIB photon spectrum is significantly flatter than the ICS spectrum, and becomes comparable to ICS at photon energies of $\gtrsim 60$ keV. The dashed lines show the ICS and EIB emission from a cone-beam (half width of $10^\circ$ ) distribution of electrons."
+ In this case. ICS clearly exceeds EIB by a large factor aud the estimates of the required number density of fast electrous and their enerev content is reduced by a corresponding factor.," In this case, ICS clearly exceeds EIB by a large factor and the estimates of the required number density of fast electrons and their energy content is reduced by a corresponding factor."
+ We acknowledge that the photon field resulting from a flaring active region is no longer wolbapproxiuated by an infinite plane (as was emploved for the plotospheric photou field)., We acknowledge that the photon field resulting from a flaring active region is no longer well-approximated by an infinite plane (as was employed for the photospheric photon field).
+ ITowever. since ICS is dominated by large- scattering the error introduced is not laree.," However, since ICS is dominated by large-angle scattering the error introduced is not large."
+ We conclude that ICS from uuldly relativistic clectrous may make a significant contribution to the observed IIXR enission iu this case., We conclude that ICS from mildly relativistic electrons may make a significant contribution to the observed HXR emission in this case.
+ Again. a broken power-law electrou spectruni is required with a break energy of 300 keV. Another example of an intriguing coronal JINR source is the lunh-occulted flare on 2007 December 31(€2).," Again, a broken power-law electron spectrum is required with a break energy of 300 keV. Another example of an intriguing coronal HXR source is the limb-occulted flare on 2007 December 31."
+. The flare was z12° over the lamb as observed by RIESSI., The flare was $\approx 12^\circ$ over the limb as observed by RHESSI.
+ However. comparing direct observations by the N-rav Spectrometer ou board the Alercury MESSENGER spacecraft with GOES observations showedthat it was an M2 SNR flare.," However, comparing direct observations by the X-ray Spectrometer on board the Mercury MESSENGER spacecraft with GOES observations showedthat it was an M2 SXR flare."
+" The flare shared mauyv attributes with the ""Masuda flare(?).. a flare that showed a IIXR source above the thermal SNR loops."," The flare shared many attributes with the “Masuda"" flare, a flare that showed a HXR source above the thermal SXR loops."
+ The coronal IIXR spectrum of the 2007 December 31 event is clearly nou-thermal with a photon spectral iudex L2 between 80 keV. Iu order for the thin-tarect broiisstraliluug, The coronal HXR spectrum of the 2007 December 31 event is clearly non-thermal with a photon spectral index 4.2 between $-$ 80 keV. In order for the thin-target bremsstrahlung
+companion and interstellar extinction. we cannot adjust the companion spectrum in any. wax to fit these LR data.,"companion and interstellar extinction, we cannot adjust the companion spectrum in any way to fit these IR data."
+ While these were not simultaneous with the optical data. we clo nol expect the continuum brightness of the companion to vary from epoch to epoch.," While these were not simultaneous with the optical data, we do not expect the continuum brightness of the companion to vary from epoch to epoch."
+ Small variations due to ellipsoidal modulations will be present as a function of orbital phase. but these are less than £10% (Shahbazetal.1994)..," Small variations due to ellipsoidal modulations will be present as a function of orbital phase, but these are less than $\pm10\,\%$ \citep{Shahbaz:1994a}."
+. Aluno&Alauerhan(2006) interpreted. the mid-Lt excess seen in. W404 €vg as the tail of a Shakura-Sunvaev accretion disc., \citet{Muno:2006a} interpreted the mid-IR excess seen in V404 Cyg as the tail of a Shakura-Sunyaev accretion disc.
+ This model. however. predicts a UV. excess which is inconsistent with our{1 observations. exceeding our F330W measurement by an order of magnitude and the F250W by two orders.," This model, however, predicts a UV excess which is inconsistent with our observations, exceeding our F330W measurement by an order of magnitude and the F250W by two orders."
+ During quiescence instead most of the disc is expected to be in a cool state. Kh. with a more isothermal temperature distribution (see IHynesctal.2005 and references therein).," During quiescence instead most of the disc is expected to be in a cool state, K, with a more isothermal temperature distribution (see \citealt{Hynes:2005a} and references therein)."
+ The contamination from this material will then be most pronounced in the mid-LH as it is cooler than the companion star. but no UV excess will be present.," The contamination from this material will then be most pronounced in the mid-IR as it is cooler than the companion star, but no UV excess will be present."
+ We use a solid line in Fig., We use a solid line in Fig.
+ 3. to show a moclel inclucling a 2000lxIx. blackbody in addition to the companion star spectrum discussed above., \ref{OptSEDFig} to show a model including a K blackbody in addition to the companion star spectrum discussed above.
+ This can adequately fit both the near-It to UV SED and theSpifzer near-H1 data (it is less successful at jm)., This can adequately fit both the near-IR to UV SED and the near-IR data (it is less successful at $\mu$ m).
+" Phe blackbody fux is consistent with a disc of radius ~9«10H cem tilted at an inclination of 56"" (Shahbazetal.1994). and at a distance of kkpe.", The blackbody flux is consistent with a disc of radius $\sim9\times10^{11}$ cm tilted at an inclination of $^{\circ}$ \citep{Shahbaz:1994a} and at a distance of kpc.
+ For the system parameters discussed above. this corresponds to about ppercent of the radius of the Roche lobe of the black hole.," For the system parameters discussed above, this corresponds to about percent of the radius of the Roche lobe of the black hole."
+ This is à very. plausible size for a quiescent accretion disc., This is a very plausible size for a quiescent accretion disc.
+ We note that while Muno&Mauerhan(2006). preferred a Shakura-Sunvaev disc model for the Ht excess in sspecilicallv. in other systems they invoked a much cooler circum-binary. disc.," We note that while \citet{Muno:2006a} preferred a Shakura-Sunyaev disc model for the IR excess in specifically, in other systems they invoked a much cooler circum-binary disc."
+ Such a disc would remain a possibility in the light of our UV observations as it would not contribute to the UV flux., Such a disc would remain a possibility in the light of our UV observations as it would not contribute to the UV flux.
+ The difficulty in fitting the Εμ point supports a more controversial interpretation as proposed by Galloetal.(2007).., The difficulty in fitting the $\mu$ m point supports a more controversial interpretation as proposed by \citet{Gallo:2007a}.
+ The flux in this band is rather close to an extrapolation of the flat radio spectrum raising the possibility that the LR excess could be due to svnchrotron jet emission., The flux in this band is rather close to an extrapolation of the flat radio spectrum raising the possibility that the IR excess could be due to synchrotron jet emission.
+ The major dilliculty in quantifving this is that the radio is known to be strongly. variable (Section 5)) and was not observed simultaneously with the ji point., The major difficulty in quantifying this is that the radio is known to be strongly variable (Section \ref{RadioSection}) ) and was not observed simultaneously with the $\mu$ m point.
+ We will return to this topic in Section 6.3.., We will return to this topic in Section \ref{JetSection}.
+ Our data also provide tenuous evidence for a UV. excess., Our data also provide tenuous evidence for a UV excess.
+ One possible explanation for this is rec-leaks in the UV bandpasses (Chiaberge&Sirianni2007).., One possible explanation for this is red-leaks in the UV bandpasses \citep{Chiaberge:2007a}.
+ Although this may make a [arge contribution in F250W. the ellect is not expected for the F330W filter.," Although this may make a large contribution in F250W, the effect is not expected for the F330W filter."
+ The estimates for recd-Lleaks from Chiaberge&Sirianni(2007) are of no use as they only consider unreddened: stars. so we perform our own calculations.," The estimates for red-leaks from \citet{Chiaberge:2007a} are of no use as they only consider unreddened stars, so we perform our own calculations."
+ In addition to recd-leaks. because the spectrum is so steep in the UV it is hard to gauge the expected. [ux in à broad-band without performing synthetic photometry.," In addition to red-leaks, because the spectrum is so steep in the UV it is hard to gauge the expected flux in a broad-band without performing synthetic photometry."
+ To better assess the significance of any UV excess. we jorefore use (Laidleretal.2005) to estimate count-rates in cach bandpass for our model reddened. spectrum.," To better assess the significance of any UV excess, we therefore use \citep{Laidler:2005a} to estimate count-rates in each bandpass for our model reddened spectrum."
+ We consider two spectra. one with the full spectral coverage. and the other truncated above tto cdisentangle in-banc stellar Ες from the rec-leak.," We consider two spectra, one with the full spectral coverage, and the other truncated above to disentangle in-band stellar flux from the red-leak."
+ We find that the rec-leak is negligible in the F330W bandpass. as expected. and that the in-band stellar (lux can account for approximately half of the count-rate seen in this band: this is a little more than a 36 excess.," We find that the red-leak is negligible in the F330W bandpass, as expected, and that the in-band stellar flux can account for approximately half of the count-rate seen in this band; this is a little more than a $\,\sigma$ excess."
+ The P250W band appears to have a large contribution from red-Ieak. about half. of the Hus in the band.," The F250W band appears to have a large contribution from red-leak, about half of the flux in the band."
+ After subtraction of the companion star. both in-band and red-Ieak. the residual excess is 0.40 ss .so is not a significant detection of disc light in this banc.," After subtraction of the companion star, both in-band and red-leak, the residual excess is $0.36\pm0.40$ $^{-1}$, so is not a significant detection of disc light in this band."
+ We therefore appear to have a significant detection of a UV excess in the F330W band only. but even this is at the mercy. of uncertainties in the companion spectral classification. (which do become a concern in the UV) and recdening.," We therefore appear to have a significant detection of a UV excess in the F330W band only, but even this is at the mercy of uncertainties in the companion spectral classification (which do become a concern in the UV) and reddening."
+ We therefore choose not to pursue the spectrum of the UV exeess further., We therefore choose not to pursue the spectrum of the UV excess further.
+ We extracted spectra for both observations with CIAO v3.2 ancl analysed them with NSPEC v12.2., We extracted spectra for both observations with CIAO v3.2 and analysed them with XSPEC v12.2.
+ We used standard CLAO techniques to select response files according to the CCD temperature., We used standard CIAO techniques to select response files according to the CCD temperature.
+ We considered. many. spectral models within NSPEC. all with interstellar absorption. and report results from fitting with power-law. bremsstrahlung. and Ravimond-Smith models here.," We considered many spectral models within XSPEC, all with interstellar absorption, and report results from fitting with power-law, bremsstrahlung, and Raymond-Smith models here."
+ Blackbody. fits to the data were so poor that we neglect. them. for. discussion., Blackbody fits to the data were so poor that we neglect them for discussion.
+ Due to the lower count rate in the 2003 observation. we were prompted to try using CASLL statistics in. tandem. with the Churazov approximation for chi-squared statistics (Cash. 1979: Churazov. 1996).," Due to the lower count rate in the 2003 observation, we were prompted to try using CASH statistics in tandem with the Churazov approximation for chi-squared statistics (Cash, 1979; Churazov, 1996)."
+ For each set of statistics to, For each set of statistics to
+clusters like Ál689. since thev have a relatively wide internal velocity spread aud they extend to large radii: this iucreases the chance for iuterlopers. which can have a large effect on the derived projected velocity dispersion. especially at larec radi where the number density of cluster members is low.,"clusters like A1689, since they have a relatively wide internal velocity spread and they extend to large radii; this increases the chance for interlopers, which can have a large effect on the derived projected velocity dispersion, especially at large radii where the number density of cluster members is low."
+ Wojtak Lokas (2007) showed that without the proper removal of iuterlopers he inferred. parameters of the mass distribution in the cluster are strongly biased towards higher mass aud ower concentration., Wojtak Lokas (2007) showed that without the proper removal of interlopers the inferred parameters of the mass distribution in the cluster are strongly biased towards higher mass and lower concentration.
+ See also Wojtak et ((2007) or mteresting comparison between different methods of interlopers removal. which unfortunately do not iuclude he D99 caustics approach used in this paper.," See also Wojtak et (2007) for interesting comparison between different methods of interlopers removal, which unfortunately do not include the D99 caustics approach used in this paper."
+ To define the caustics we cuploved the techiuique Xoneered by Diaferio (1999. hereafter D99) based ou a imultidineusional adaptive kernel method (Silverman 1986: Pisani 1993: μαι 1996).," To define the caustics we employed the technique pioneered by Diaferio (1999, hereafter D99) based on a multidimensional adaptive kernel method (Silverman 1986; Pisani 1993; Pisani 1996)."
+" A siuunuuwv of the application of this techuique can be found in Diaterio. Geller, Rines (2005): it has previously been applied to several clusters (Reiscnegecr et 22000: Biviano Carardi 2003: Diaferio. Geller. Rines 2005: Rines Diaterio 2006)."," A summary of the application of this technique can be found in Diaferio, Geller, Rines (2005); it has previously been applied to several clusters (Reisenegger et 2000; Biviano Girardi 2003; Diaferio, Geller, Rines 2005; Rines Diaferio 2006)."
+" Briefly a redshift-space diagram is constructed. ic. the line of sight velocity ο is plotted versus the projected distance A frou the cluster center. and for a well defined cluster these points should be distributed in a characteristic ""truupet shape. the boundaries of which are termed caustics (saiser 1987: Regos Geller 1989)."," Briefly, a redshift-space diagram is constructed, i.e., the line of sight velocity $v$ is plotted versus the projected distance $R$ from the cluster center, and for a well defined cluster these points should be distributed in a characteristic “trumpet” shape, the boundaries of which are termed caustics (Kaiser 1987; Regos Geller 1989)."
+ The D99 procedure locates the caustics aud ctermunes the radial dependence of their amplitude (iu uits of velocity). which is related to the escape velocity and thus depends ou the mass profile.," The D99 procedure locates the caustics and determines the radial dependence of their amplitude (in units of velocity), which is related to the escape velocity and thus depends on the mass profile."
+ Galaxies that are inside the caustics can then be considered to be cluster members., Galaxies that are inside the caustics can then be considered to be cluster members.
+ Of course. some of these galaxies might be interlopers. but their uuniber is typically a few percent and has little effect ou dynamical analyses (A. Diaterio. private conmuuuication).," Of course, some of these galaxies might be interlopers, but their number is typically a few percent and has little effect on dynamical analyses (A. Diaferio, private communication)."
+ To apply this procedure we determined the threshold & that defines the caustic location through ΤΠ.ο)=6 (D99).," To apply this procedure we determined the threshold $\kappa$ that defines the caustic location through $f_q(R,v)=\kappa$ (D99)."
+" Tere f,2.e) is the galaxy density distribution in the redshift-space diagram. sinoothed with au adaptive kernel. and ce is the peculiar velocity (neasured with respect to the clusters mean redshitt)."," Here $f_q(R,v)$ is the galaxy density distribution in the redshift-space diagram, smoothed with an adaptive kernel, and $v$ is the peculiar velocity (measured with respect to the cluster's mean redshift)."
+ The parameter q sets the scaling between the quantities R aud ο within the smoothing procedure., The parameter $q$ sets the scaling between the quantities $R$ and $v$ within the smoothing procedure.
+ We used 4=27. close to the value usually used. 25 (D99: Raines et 22003).," We used $q=27$, close to the value usually used, $25$ (D99; Rines et 2003)."
+ We note that different values of gin the 1050 range have little effect on the results (9091., We note that different values of $q$ in the $10-50$ range have little effect on the results (D99).
+" The parameter & was chosen by iuninüzius the quantity Ste.(65)=[lezaepLicP. where CeciPASmRJU(""RyoΜο.otl ds the nean value of the square of the caustic amplitude ACR) within CR; (the mean projected radius of the cluster ucnbers) oR)=fFíUR.e)de. aud is the one-dimensional velocity dispersion of the cluster(025 meubers."," The parameter $\kappa$ was chosen by minimizing the quantity $S(\kappa,\langle R\rangle)=|\langle v_{esc}^2\rangle_{\kappa,\langle
+R\rangle}-4 \langle v^2\rangle|^2$, where $\langle
+v_{esc}^2\rangle_{\kappa,\langle R\rangle}= \int_0^{\langle
+R\rangle}A^2(R)\phi(R)dR/\int_0^{\langle R\rangle}\phi(R)dR$ is the mean value of the square of the caustic amplitude $A(R)$ within $\langle R\rangle$ (the mean projected radius of the cluster members), $\phi(R)=\int f_q(R,v)dv$, and $\langle v^2\rangle$ is the one-dimensional velocity dispersion of the cluster members."
+ The uncertainty of Ais proportional to the inverse of he galaxy πο density within the caustics. and was estimated as AGRO)ZAGR)=κinaxLAUR.ey}. where he maxim is found along the ¢-axis at cach /& (D99).," The uncertainty of $A$ is proportional to the inverse of the galaxy number density within the caustics, and was estimated as $\delta A(R)/A(R) =
+\kappa/ \max \left\{f_q(R,v)\right\}$, where the maximum is found along the $v$ -axis at each $R$ (D99)."
+ Figure 2 shows the velocity-space diagram aud radial xofle of the derived velocity caustics., Figure \ref{velocity space diagram} shows the velocity-space diagram and radial profile of the derived velocity caustics.
+ There are a otal of πο galaxies ideutified as cluster members vine between τηthese caustics., There are a total of $476_{-43}^{+27}$ galaxies identified as cluster members lying between these caustics.
+ The velocity-space diagram clearly looks qualitatively differcut inside aud outside the caustics. iudicating that the caustic-fiudine procedure has identified a /phvsicall-1ueaniugful boundary.," The velocity-space diagram clearly looks qualitatively different inside and outside the caustics, indicating that the caustic-finding procedure has identified a physically-meaningful boundary."
+ The density of the ealaxies that lie between the caustics shows a smooth decline with distance from the ceuter. indicating that the cluster has a simple monolithic ανασα. structure.," The density of the galaxies that lie between the caustics shows a smooth decline with distance from the center, indicating that the cluster has a simple monolithic dynamical structure."
+ The radius at which the caustics iicet at zero velocity represents an effective lias radius for the cluster. which plysically includes cluster ealaxies both withiu the virial radius aud others that may be currently intalling outo the cluster (see tle discussion i 7).," The radius at which the caustics meet at zero velocity represents an effective maximum radius for the cluster, which physically includes cluster galaxies both within the virial radius and others that may be currently infalling onto the cluster (see the discussion in \ref{Discussion}) )."
+ Note that the two caustics are by definition svuunetric about e=0 in this method. which figure 2 clearly shows is a valid assuniptiou for this cluster.," Note that the two caustics are by definition symmetric about $v=0$ in this method, which figure \ref{velocity space diagram}
+ clearly shows is a valid assumption for this cluster."
+ ILwius defined cluster imenibership based on the location of the caustics. we can estimate the projected velocity dispersion profile restricted to cluster iienibers.," Having defined cluster membership based on the location of the caustics, we can estimate the projected velocity dispersion profile restricted to cluster members."
+ We divided the cluster members iuto velocity bins of size G00 kui/sec. sufficiently large. compared to the redshift neasurement uncertainty of ~200 laun/sec.," We divided the cluster members into velocity bins of size 600 km/sec, sufficiently large compared to the redshift measurement uncertainty of $\sim 200$ km/sec."
+ Dittercut ye eave very simular results., Different binning gave very similar results.
+ We fitted a Caussian to he velocity histogram to calculate the overall projected velocity dispersion (figure 3)): it provided a reasonable fit with a reduced. 4? (hereafter 42) of 1.7 per degree of freedom., We fitted a Gaussian to the velocity histogram to calculate the overall projected velocity dispersion (figure \ref{velocity histogram}) ); it provided a reasonable fit with a reduced $\chi^2$ (hereafter $\chi^2_r$ ) of 1.7 per degree of freedom.
+ If we ignore the caustics and simply bin all ealaxies within the maximal velocity range |e|«1000 au/sec. the fit to a Cassia is significautlv worse. \2=2.5. with a nore asvunnetrie distribution and a somewhat larger velocity dispersion.," If we ignore the caustics and simply bin all galaxies within the maximal velocity range $|v|<4000$ km/sec, the fit to a Gaussian is significantly worse, $\chi^2_r=2.5$, with a more asymmetric distribution and a somewhat larger velocity dispersion."
+ It is iuportaut to nention that we did not use the results of the Gaussian Bt in our analysis., It is important to mention that we did not use the results of the Gaussian fit in our analysis.
+ We show this fit to argue in favor of the real plivsical iieaniug of the caustics as we lave neasured them. aud demoustrate that without excludiug unrelated iuterlopers with large velocities. cvnamically derived quantities of the cluster may be significantly skewed. particularly at large projected radi.," We show this fit to argue in favor of the real physical meaning of the caustics as we have measured them, and demonstrate that without excluding unrelated interlopers with large velocities, dynamically derived quantities of the cluster may be significantly skewed, particularly at large projected radii."
+ Note that we fitted a Caussian for simplicity. although the real distribution aay be shebthy non-Gaussian (sec. e.g. Iazautzidis. \Mlagorrian Moore 2001: Sauchis. Lokas Abunon 20014: Diemiand. Moore Stadel 2001: Wojtak et 22005).," Note that we fitted a Gaussian for simplicity, although the real distribution may be slightly non-Gaussian (see, e.g., Kazantzidis, Magorrian Moore 2004; Sanchis, Lokas Mamon 2004; Diemand, Moore Stadel 2004; Wojtak et 2005)."
+ We next estimated the radial profile of the projected velocity dispersion from our sample of cluster 11enibers, We next estimated the radial profile of the projected velocity dispersion from our sample of cluster members
+Because it is Likely that more interlopers will lave volocities that are ercater than their hosts (e.e@.. Zartisky White 1991). we determined σι for the host galaxies bv fitting Gaussians plus constant offsets to 23 different velocity difference histograms: (0) velocity differences taken to be the absolute value. |de]|. of the measured difference. (i) ueeative velocity differences. and (i) positive velocity differcuces.,"Because it is likely that more interlopers will have velocities that are greater than their hosts (e.g., Zartisky White 1994), we determined $\sigma_v$ for the host galaxies by fitting Gaussians plus constant offsets to 3 different velocity difference histograms: (i) velocity differences taken to be the absolute value, $|dv|$, of the measured difference, (ii) negative velocity differences, and (iii) positive velocity differences."
+ We define the velocity difference to be de=ChosrCus 80 that negative values of de correspond to satellites which are more distaut in velocity space.," We define the velocity difference to be $dv \equiv v_{\rm host} - v_{\rm sat}$, so that negative values of $dv$ correspond to satellites which are more distant in velocity space."
+ In all cases. the bestfitting velocity dispersions are in very good agreement amongst the 3 histograms.," In all cases, the best–fitting velocity dispersions are in very good agreement amongst the 3 histograms."
+" Iu αποοποια, we find a clear difference in the munber of interlopers."," In addition, we find a clear difference in the number of interlopers."
+ In the case of the full sample of 5089 hosts. the interloper fraction is (314: for hostsatellite pairs with de«0. (204 for hostsatellite pairs with de> 0. aud (272X on average (ie. fitting to the distribution of |de|).," In the case of the full sample of 809 hosts, the interloper fraction is $(31\pm 3)$ for host--satellite pairs with $dv < 0$, $(20 \pm 3)$ for host–satellite pairs with $dv > 0$ , and $(27\pm 2)$ on average (i.e., fitting to the distribution of $|dv|$ )."
+ For the hosts with spiral morphologies. the mean iuterloper fraction is (33+: while for hosts with clliptical and SO morphologics the interloper fraction is much lower. (Ll1)..," For the hosts with spiral morphologies, the mean interloper fraction is $(33\pm 3)$ while for hosts with elliptical and S0 morphologies the interloper fraction is much lower, $(14\pm 4)$."
+" Lastly. for all of our suüuples of host galaxies. we find that 96, is independent of the radius at which it is determined."," Lastly, for all of our samples of host galaxies, we find that $\sigma_v$ is independent of the radius at which it is determined."
+" Results for o6, as a function of median host. huninosity are shown in panel d of ll. 2. and 3."," Results for $\sigma_v$ as a function of median host luminosity are shown in panel d of 1, 2, and 3."
+ In the case of the full host sample aud the elliptical/S0 host sample. the relationship between velocity dispersion aud median host hunuiuositv is fit best bv linear relations. aud these are shown by the solid lines in panel d of 11 aud .," In the case of the full host sample and the elliptical/S0 host sample, the relationship between velocity dispersion and median host luminosity is fit best by linear relations, and these are shown by the solid lines in panel d of 1 and 2."
+" For full host sample we fud σι=(31ESJLIL),|un+18)the hu Ἡ, and for the clliptical/S0 sample we Bud o,=(134το|(11922) kin"," For the full host sample we find $\sigma_v = (31\pm 8)L/L^\ast + (141\pm 18)$ km $^{-1}$, and for the elliptical/S0 sample we find $\sigma_v = (43\pm 7)L/L^\ast
+ + (149\pm 22)$ km $^{-1}$."
+" This is somewhat different frou the results of SDSS02 who ford a,xVL."," This is somewhat different from the results of SDSS02 who found $\sigma_v \propto
+\sqrt{L}$."
+ We note. however. that our data are consistent with such a relationship. aud the bestfitting function of the form σεXVL is shown by the dashed lines in these figures.," We note, however, that our data are consistent with such a relationship, and the best–fitting function of the form $\sigma_v \propto \sqrt{L}$ is shown by the dashed lines in these figures."
+ Strikinely differcut frou: these results. however. is the relationship of velocity dispersion and median host huninosity for the spiral hosts.," Strikingly different from these results, however, is the relationship of velocity dispersion and median host luminosity for the spiral hosts."
+" From panel d of 33. it is clear that σι for the spiral hosts is independent of host Iuniunositv. and we fud e,=180d9 kins | for these objecta."," From panel d of 3, it is clear that $\sigma_v$ for the spiral hosts is independent of host luminosity, and we find $\sigma_v = 189\pm 9$ km $^{-1}$ for these objects."
+ To determine the masses of the dark matter halos which suround our host galaxies. we adopt the following mass estimator: (ee. Binney Tremaine. 1987).," To determine the masses of the dark matter halos which surround our host galaxies, we adopt the following mass estimator: (e.g., Binney Tremaine, 1987)."
+ Tere r ds a Ὁ-cmieusioual radius. 402) is the 11ean square racial velocity of the satellites. p(r) is the nuuber deusitv of satellites. and ο) is a ineasure of the anisotropy in the velocity dispersion of the satellites: Although other methods of obtaining dynamical masses using satellite ealaxies have been adopted in the literature (sec. e.g. Dahlcall Tremaine 1982: Zaritsky White 1991: Zavitsky ct 11997). this is the method adopted by SDSSU2 and. therefore. we adopt it as well.," Here $r$ is a 3-dimensional radius, $\left< v^2_r \right>$ is the mean square radial velocity of the satellites, $\rho(r)$ is the number density of satellites, and $\beta$ is a measure of the anisotropy in the velocity dispersion of the satellites: Although other methods of obtaining dynamical masses using satellite galaxies have been adopted in the literature (see, e.g., Bahcall Tremaine 1982; Zaritsky White 1994; Zaritsky et 1997), this is the method adopted by SDSS02 and, therefore, we adopt it as well."
+ SDSS02 have used the GIF simulation. which iucorporates senianalytic ealaxy formation within a laree cosmological N simulation (c.g.. Ikauffinaun et 11999). to evaluate this mass estimator.," SDSS02 have used the GIF simulation, which incorporates semi--analytic galaxy formation within a large cosmological N--body simulation (e.g., Kauffmann et 1999), to evaluate this mass estimator."
+ In particular. SDSS02 find that the velocity anisotropy of the satellite galaxies in the CIF siuulation is small G.c.. ο) is consistent with zero at the c level). aud that the mean square line of sight velocity dispersion. 92. is consisteut with the mean square radial velocity dispersion. £025. at the σ level.," In particular, SDSS02 find that the velocity anisotropy of the satellite galaxies in the GIF simulation is small (i.e., $\beta$ is consistent with zero at the $\sigma$ level), and that the mean square line of sight velocity dispersion, $\sigma_v^2$, is consistent with the mean square radial velocity dispersion, $\left< v^2_r \right>$, at the $\sigma$ level."
+ Combining this with the fact that the line of sight velocity dispersion is observed to be independent of radius. the mass estimator used by SDSSO2 reduces to: We compute p(r) for the satellites in our sample aud find p(r)xor7117996 foy the satellites suroundiug the hosts in the full z;uuple. p(r)2.110009. fo the satellites surrounding elliptical and SO hosts. and pir)xkc2279 for the satellites surrounding spiral hosts.," Combining this with the fact that the line of sight velocity dispersion is observed to be independent of radius, the mass estimator used by SDSS02 reduces to: We compute $\rho(r)$ for the satellites in our sample and find $\rho(r) \propto r^{-2.11\pm 0.06}$ for the satellites surrounding the hosts in the full sample, $\rho(r) \propto r^{-2.11\pm 0.09}$ for the satellites surrounding elliptical and S0 hosts, and $\rho(r) \propto r^{-2.2 \pm 0.1}$ for the satellites surrounding spiral hosts."
+ These are all consistent with the results of SDSS02 who find p(r) for their sample., These are all consistent with the results of SDSS02 who find $\rho(r) \propto r^{-2.1}$ for their sample.
+" Taving obtained the nuuber deusitv of satellites as a ""unction of radius. we now use the values of σι from 11. 2. aud 3 to determine the masstolight ratios of the tost galaxies."," Having obtained the number density of satellites as a function of radius, we now use the values of $\sigma_v$ from 1, 2, and 3 to determine the mass–to–light ratios of the host galaxies."
+ For consistency with SDSS0U2. we adopt a fiducial projected radius of 2607.1 kpc.," For consistency with SDSS02, we adopt a fiducial projected radius of $260h^{-1}$ kpc."
+" Shown in 11 are the results for(ALL),, iun units of PhALL/L. asa unction of median host luminosity.", Shown in 4 are the results for $(M_{260}^{\rm dyn} /L)_{b_J}$ in units of $h~M_\odot/L_\odot$ as a function of median host luminosity.
+ The top panel shows results for the full sample of host ealaxics. the iuddle xuel shows results for the ellipticalaud SO hosts. aud he bottom panel shows results for tle spiral hosts.," The top panel shows results for the full sample of host galaxies, the middle panel shows results for the ellipticaland S0 hosts, and the bottom panel shows results for the spiral hosts."
+" From EL thou. we find that(ALL),, for the 509 rosts in our full sample is fairly constant for hosts with Lz2L aud has a value of (19341D ALLE..."," From 4, then, we find that $(M_{260}^{\rm dyn}/L)_{b_J}$ for the 809 hosts in our full sample is fairly constant for hosts with $L \gs 2L^\ast$ and has a value of $(193\pm 14) h~M_\odot/L_\odot$ ."
+ For hosts, For hosts
+the data.,the data.
+ The cluster mass function is primarily determined by the cluster magnitudes. while (he radius distribution is constrained by (he sharpness measurements.," The cluster mass function is primarily determined by the cluster magnitudes, while the radius distribution is constrained by the sharpness measurements."
+ Thus we can accuratelv determine a mass function slope of a=—1.5040.07 and a distribution of core radii centered on ο.=1.53 40.15pc., Thus we can accurately determine a mass function slope of $\alpha = -1.50 \pm 0.07$ and a distribution of core radii centered on $r_c = 1.53 \pm 0.15$ pc.
+ Less certain were our measurements of (he extinction distribution and cluster formation history. as both affect the colors.," Less certain were our measurements of the extinction distribution and cluster formation history, as both affect the colors."
+ These variables will be more strongly constrainted in our future work. which will include galaxies observed through three broadband filters.," These variables will be more strongly constrainted in our future work, which will include galaxies observed through three broadband filters."
+ Nevertheless. we were able to measure a mean extinction that was consistent with that measured from (he Cepheics. giving some indication (hat even that poorlv-constrained. parameter was indeed measured correctly.," Nevertheless, we were able to measure a mean extinction that was consistent with that measured from the Cepheids, giving some indication that even that poorly-constrained parameter was indeed measured correctly."
+ Support for this work was provided bv NASA through erant number AR-09196.01 [rom the Space Telescope Science Institue., Support for this work was provided by NASA through grant number AR-09196.01 from the Space Telescope Science Institute.
+ All of the clata presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (NLAST)., All of the data presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (MAST).
+ STScl is operated by the Association of Universities for Research in Astronomy. Inc.. under NASA contract NAS5-26555.," STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555."
+matrix using weights proportional to the backerouud subtracted couuts from each uode.,matrix using weights proportional to the background subtracted counts from each node.
+ We have ignored the effects due to the secular decline in the ACIS quantum elliciency since the observations ol NGC LOGS took place early iu the mission. auc the effects ou the derived values of the columu density are negligible (less than 105«)) compared to the statistical error.," We have ignored the effects due to the secular decline in the ACIS quantum efficiency since the observations of NGC 1068 took place early in the mission, and the effects on the derived values of the column density are negligible (less than $10$ ) compared to the statistical error."
+ Prior to performing the spectral analysis with version 11.1.0 (Arnaud.1996)... we rebinnecd the source spectra so that there were at least 15 counts per bin (for sources with >100 counts) or at least LO counts per bin (for sources with <100 counts). tius allowing use of the 47 statistic.," Prior to performing the spectral analysis with version 11.1.0 \citep{arn96}, we rebinned the source spectra so that there were at least 15 counts per bin (for sources with $\geq 100$ counts) or at least 10 counts per bin (for sources with $<100$ counts), thus allowing use of the $\chi^{2}$ statistic."
+ The results of the spectral analysis are giveu in Table 2.., The results of the spectral analysis are given in Table \ref{tbl-2}.
+ Light curves were extracted iu tlie 0. 1-5 keV bail for each of the 21 uou-nuclear sources witli sullicient. counts for spectral analysis (see Table 2))., Light curves were extracted in the $0.4$ $5$ keV band for each of the 21 non-nuclear sources with sufficient counts for spectral analysis (see Table \ref{tbl-2}) ).
+ The low count rates for most of the sources preclude an analysis of the light curves based on the 47 statistic. which requires the data to be biuued.," The low count rates for most of the sources preclude an analysis of the light curves based on the $\chi^{2}$ statistic, which requires the data to be binned."
+ To assess the probability of variability. we have. instead. compared the distribution of event arrival times for each source with that of a large source-[ree region of the $3 chip (105257 eveuts). using the Ixolmogorov-Smirnov (4X —9) test (see e.g.. Press et al.," To assess the probability of variability, we have, instead, compared the distribution of event arrival times for each source with that of a large source-free region of the S3 chip (10557 events), using the Kolmogorov-Smirnov $K$ $S$ ) test (see e.g., Press et al."
+ 1992. p. 623).," 1992, p. 623)."
+ We find a total of 81 compact sources on chip $8 (Table 1)). of which 66 are located within the 25.0 B-magnitude (are 7 isophote of the galactie disk of NGC. LOGS (de Vaucouleurs et al.," We find a total of 84 compact sources on chip S3 (Table \ref{tbl-1}) ), of which 66 are located within the 25.0 B-magnitude (arc $^{-2}$ isophote of the galactic disk of NGC 1068 (de Vaucouleurs et al."
+ 1991)., 1991).
+ Fig., Fig.
+ 2. shows the locations of the 81 compact sources superposed ou the Digitized Sky Survey inge., \ref{fig2} shows the locations of the 84 compact sources superposed on the Digitized Sky Survey image.
+ As cau be seen in this figure. approximately half of the sources are projected onto the bright inner disk region of NGC LOGS.," As can be seen in this figure, approximately half of the sources are projected onto the bright inner disk region of NGC 1068."
+ In order to quantify the expected number of background sources in the field. we have performed an identical analysis on an. ACIS-S observation of 3C. 17 (obsid 2129). which has a similar exposure time (11.6 ks) to our observation of NGC LOGS.," In order to quantify the expected number of background sources in the field, we have performed an identical analysis on an ACIS-S observation of 3C 47 (obsid 2129), which has a similar exposure time $41.6$ ks) to our observation of NGC 1068."
+ We have cousidered only regions >1! [from 3C 17 as there is bright exteuded X-ray emission surrounding[n] this source., We have considered only regions $\ge 1^{\prime}$ from 3C 47 as there is bright extended X-ray emission surrounding this source.
+ Regions aloug the reaclout trail were also excluded from our aualysis., Regions along the readout trail were also excluded from our analysis.
+ We find a mean surface density of 0.37 sources (are  iu this observation. which is almost equal to the mean surface deusity of sources -[zz0.15 sources (arc NEP 7] outside the ae25.0 B-magnitude⋅ (are∣ ? 3.isophote of the galactic disk of NGC 1068.," We find a mean surface density of $0.37$ sources (arc $^{-2}$ in this observation, which is almost equal to the mean surface density of sources $\approx 0.45$ sources (arc $^{-2}$ ] outside the $25.0$ B-magnitude (arc $^{-2}$ isophote of the galactic disk of NGC 1068."
+ This suggestsMOD that most of the latter sources are background objects uurelated to NGC 1068., This suggests that most of the latter sources are background objects unrelated to NGC 1068.
+ From the mean surface density of sources in the 3C 17 field. we would expect. in our observation. approximately 12 backgrouud sources projected outo the ealactic disk of NGC LOGS. whereas we detect 66 sources.," From the mean surface density of sources in the 3C 47 field, we would expect, in our observation, approximately $12$ background sources projected onto the galactic disk of NGC 1068, whereas we detect 66 sources."
+ This uuumber of backgrouud sources is a couservative estimate.," This number of background sources is a conservative estimate,"
+at oc 20 Cllz but the proportion of such objects to be detected in very high frequency surveys remains unknown.,at $\ge$ 20 GHz – but the proportion of such objects to be detected in very high frequency surveys remains unknown.
+ The first ventures at higher frequencies were to search or contaminants to CAIB experiments. with the Rvle Telescope survey at 15 Cillz leading the way.," The first ventures at higher frequencies were to search for contaminants to CMB experiments, with the Ryle Telescope survey at 15 GHz leading the way."
+ Tavlor et al. (, Taylor et al. (
+2001) found 66 sources in 63 above 20 niv.,2001) found 66 sources in 63 above 20 mJy.
+ The continuation of this came to be the Cambridge 9€ survey (Waldram et al., The continuation of this came to be the Cambridge 9C survey (Waldram et al.
+ 2003: Bolton et al., 2003; Bolton et al.
+ 2004: Waldram ct al., 2004; Waldram et al.
+ 2009)., 2009).
+ Some 465 sources were catalogued above 25 mv in 520deg., Some 465 sources were catalogued above 25 mJy in 520.
+. Examination of a subset of these sources Bolton et al., Examination of a subset of these sources (Bolton et al.
+ 2004) revealed that (a) at a Bux limit of 25 mJv LS per cent show spectra that peak at 5 CGllz or above: and (b) at GO mJv 27 per cent peak at 5 CGllz or above. and increasing the Dux limit increases the fraction of rising-spectrum sources. (," 2004) revealed that (a) at a flux limit of 25 mJy 18 per cent show spectra that peak at 5 GHz or above; and (b) at 60 mJy 27 per cent peak at 5 GHz or above, and increasing the flux limit increases the fraction of rising-spectrum sources. ("
+Phe statistical dependence of radio spectrum. on Llusx-clensitw level has been. discovered relatively regularly since. 1970: see e.g. Ixellermann Wall 1987.),The statistical dependence of radio spectrum on flux-density level has been discovered relatively regularly since 1970; see e.g. Kellermann Wall 1987.)
+ Colours and the stellar nature of optical counterparts indicate that between a third ancl a half of the spectrum sources are QSOs., Colours and the stellar nature of optical counterparts indicate that between a third and a half of the peaked-spectrum sources are QSOs.
+ The major contribution to surveys and samples at these higher frequencies has come from NT20C. the large-area 20-CGllz survey with the Australia Telescope Compact. Array (Ricci οἱ al 2004: Sacer et al.," The major contribution to surveys and samples at these higher frequencies has come from AT20G, the large-area 20-GHz survey with the Australia Telescope Compact Array (Ricci et al 2004; Sadler et al."
+ 2006: Massarcdi et al., 2006; Massardi et al.
+ 2008)., 2008).
+" ὃν 2007 the survey had covered 90"" po<15"" ccataloguing 4400 sources down to 50 mJ: another 1500 are expected in. the final⋅ survey zone 15 ""mi18<O°.", By 2007 the survey had covered $-90$ $ < \delta < -15$ cataloguing 4400 sources down to $\sim$ 50 mJy; another 1500 are expected in the final survey zone $-15$ $ < \delta < 0$.
+" ""Phese catalogues will render incorrect the statements made at the start of cach of these papers about our poor knowledge of the high-frequency racio sky.", These catalogues will render incorrect the statements made at the start of each of these papers about our poor knowledge of the high-frequency radio sky.
+ Massardi et al., Massardi et al.
+ studied the 320 brightest sources [rom the 2007 version of the catalogue., studied the 320 brightest sources from the 2007 version of the catalogue.
+ The major conclusions do not differ dramatically from those from 9€5 (a) most sources do not show power-law spectra: (b) the spectral complexity ensures that it is impossible to select a Iow-[requencey sample which will constitute a complete high-frequency. sample: (0) spectral steepening is common and correlates with redshift. perhaps due to changing rest frame frequeney: (d) 77 per cent of the brightest-source. sample are QSOs. 19 per cent are galaxies. and the remainder are blank on the SuperCosmos Sky Survey (SSS: see wlawrocacuk/sss/) scanned versions of the UINSTU plates.," The major conclusions do not differ dramatically from those from 9C: (a) most sources do not show power-law spectra; (b) the spectral complexity ensures that it is impossible to select a low-frequency sample which will constitute a complete high-frequency sample; (c) spectral steepening is common and correlates with redshift, perhaps due to changing rest frame frequency; (d) 77 per cent of the brightest-source sample are QSOs, 19 per cent are galaxies, and the remainder are blank on the SuperCosmos Sky Survey (SSS; see ) scanned versions of the UKSTU plates."
+ The galaxies have a median 2) = 17.7. the QSOs 18.6 these medians lie far above the plate limits.," The galaxies have a median $B_J$ = 17.7, the QSOs 18.6 – these medians lie far above the plate limits."
+ Further characterization of the source population at high frequencies has come from multi-frequency ux-density measurements of sources detected in the WALADP (Wilkinson Microwave Anisotropy Probe) CAIB surveys., Further characterization of the source population at high frequencies has come from multi-frequency flux-density measurements of sources detected in the WMAP (Wilkinson Microwave Anisotropy Probe) CMB surveys.
+ WMADP has produced. complete samples of the brightest. radio sources over the sky to a Husx-density limit of about 2 Jy at each of its primary frequencies 23. 33. 43. 61 and 94 CllIz (Llinshaw et al.," WMAP has produced complete samples of the brightest radio sources over the sky to a flux-density limit of about 2 Jy at each of its primary frequencies 23, 33, 43, 61 and 94 GHz (Hinshaw et al."
+ 2007: Lóppez-C'aniego et al., 2007; Lóppez-Caniego et al.
+ 2007: Wright et al., 2007; Wright et al.
+ 2009)., 2009).
+ With measurements at 16 and 33 GlIz using the VSA and ΑΔΗ. Davies et al. (," With measurements at 16 and 33 GHz using the VSA and AMI, Davies et al. ("
+2009) and Franzen et al. (,2009) and Franzen et al. (
+"2009) found he following for a sample complete to 1.1 Jy a 3 Gllz: (1) he proportion of Dlat-spectrum objects is very high: 93 per cent have spectra with' aT3lcg,>0.5: (2) 44 per cent have OLLI3dGHz70.0. i.e. rising spectra: (3) the variable-source bias (à selection elfect which mimics an Eddington bias: see Wall et al.","2009) found the following for a sample complete to 1.1 Jy at 33 GHz: (1) the proportion of flat-spectrum objects is very high: 93 per cent have spectra with $\alpha^{34}_{14~\rm{GHz}} > -0.5$; (2) 44 per cent have $\alpha^{34}_{14~\rm{GHz}} > 0.0$, i.e. rising spectra; (3) the variable-source bias (a selection effect which mimics an Eddington bias; see Wall et al."
+ 2005: Wall 2007) is in clear evidence from comparing repeated WALAP ancl [AME 33-CGLHIz Dux. densities: (3) on imeseales of —1.5vr. 20 per cent of sources varied by more han 20 per cent at 33 GHz with a high degree of correlation »etween variation at 16 and 33 Gllz: and (4) variable sources jwe a marginally Hatter spectrum ία=0.06+£0.05) than non-variable sources (à=0.13+0.04).," 2005; Wall 2007) is in clear evidence from comparing repeated WMAP and AMI 33-GHz flux densities; (3) on timescales of $\sim$ 1.5yr, 20 per cent of sources varied by more than 20 per cent at 33 GHz with a high degree of correlation between variation at 16 and 33 GHz; and (4) variable sources have a marginally flatter spectrum $\overline{\alpha} = -0.06\pm0.05$ ) than non-variable sources $\overline{\alpha} = -0.13\pm0.04$ )."
+ Sadler et al. (, Sadler et al. (
+2008) attempted to characterize the populations to be found at vet higher frequencies.,2008) attempted to characterize the populations to be found at yet higher frequencies.
+ With flux densities measured. for 130 sources at 95 GlIz and 20 6112 with the ATCA. the resultant distribution of spectral indices ogether with source counts [rom ΑΕθέ was used to predict he source count at 95 CGllz above a level of LOO mJy.," With flux densities measured for 130 sources at 95 GHz and 20 GHz with the ATCA, the resultant distribution of spectral indices together with source counts from AT20G was used to predict the source count at 95 GHz above a level of 100 mJy."
+ This wrecliction is lower than that from models of luminosity unctions and evolution (De Zotti οἱ al., This prediction is lower than that from models of luminosity functions and evolution (De Zotti et al.
+ 2005)., 2005).
+ The authors conclude that this is due to the fact that the majority of sources with Lat or rising spectra over the range 5 to 20 61112 show a spectral turnover between 20 and 95 Cllz., The authors conclude that this is due to the fact that the majority of sources with flat or rising spectra over the range 5 to 20 GHz show a spectral turnover between 20 and 95 GHz.
+ lt ds the purpose of this paper to describe a. limite survey at mv levels at 43.4 Gllz with the VLA. and to use this to construct a source count at this frequeney.," It is the purpose of this paper to describe a limited survey at mJy levels at 43.4 GHz with the VLA, and to use this to construct a source count at this frequency."
+ We wan to see if there is evidence for emergent new populations tha night allect future CALB experiments such as Planck., We want to see if there is evidence for emergent new populations that might affect future CMB experiments such as Planck.
+ Io addition we wish to track the space density of the dilleren classes of object which emerge as survey [recqueney is raised. and for this. à source count is an essential complement to radio and optical spectral data.," In addition we wish to track the space density of the different classes of object which emerge as survey frequency is raised, and for this, a source count is an essential complement to radio and optical spectral data."
+ Descriptions of this space density forcombined populations of Ilat-spectrum. objects have been obtained (e.g De Zotti ct al., Descriptions of this space density for populations of flat-spectrum objects have been obtained (e.g De Zotti et al.
+ 2005: Wall ct al., 2005; Wall et al.
+ 2005: Ricci et al., 2005; Ricci et al.
+ 2006). but not for the individual (e.g. GPS) populations.," 2006), but not for the individual (e.g. GPS) populations."
+ With regard to à source count at 43 Cllz. the WALAP all-sky survey at 43 Gllz defined the high-Dlux-density. end of the count as well as it can ever be done.," With regard to a source count at 43 GHz, the WMAP all-sky survey at 43 GHz defined the high-flux-density end of the count as well as it can ever be done."
+ To extend the 43- count to lower Dux levels. Waldram et al. (," To extend the 43-GHz count to lower flux levels, Waldram et al. ("
+2007) used their 15-Cllz 9€ survey plus spectral information to predict source counts at this frequency (as well as at several other frequencies).,2007) used their 15-GHz 9C survey plus spectral information to predict source counts at this frequency (as well as at several other frequencies).
+ At a somewhat lower frequency. Mason οἱ al. (," At a somewhat lower frequency, Mason et al. ("
+2009) measured [ux densities at 31 Gllz for 3165 NVSS sources using the Green Bank 100m and the Owens Valley 40m telescopes.,2009) measured flux densities at 31 GHz for 3165 NVSS sources using the Green Bank 100m and the Owens Valley 40m telescopes.
+ From these measurements they projected an integral source count at 31 GllIz in the range 1 to 4 mv. and found that the surface densities implied account for 21+7 per cent of the amplitude of the power detected in excess of intrinsic anisotropy by the Cosmic Background. Imager (CDI) at (m2000.," From these measurements they projected an integral source count at 31 GHz in the range 1 to 4 mJy, and found that the surface densities implied account for $21 \pm 7$ per cent of the amplitude of the power detected in excess of intrinsic anisotropy by the Cosmic Background Imager (CBI) at $\ell > 2000$."
+ We return to these results in 5. and 6., We return to these results in \ref{scest} and \ref{disc}.
+ Surveving at [frequencies as high as 43 Cllz is not straightforward. as this pilot study will demonstrate.," Surveying at frequencies as high as 43 GHz is not straightforward, as this pilot study will demonstrate."
+ Hence this paper does not follow the conventional survey. analysis. catalogue. source-count ete.," Hence this paper does not follow the conventional survey, analysis, catalogue, source-count etc."
+ route — the data do not lend themselves to such a procedure., route – the data do not lend themselves to such a procedure.
+ We thus present a road map for orientation., We thus present a road map for orientation.
+stopped being active. a galaxy could possibly re-acquire cold gas. producing cyclical phases of AGN feedback and star formation.,"stopped being active, a galaxy could possibly re-acquire cold gas, producing cyclical phases of AGN feedback and star formation."
+ Our galaxies. showing at best only very weak AGN activity. may currently be in the renewed star-forming phase. (," Our galaxies, showing at best only very weak AGN activity, may currently be in the renewed star-forming phase. ("
+In contrast. in more massive ΕΤΟΣ. M.10!!M... the accretion may be permanently disrupted.),"In contrast, in more massive ETGs, $M_*>10^{11} M_{\odot}$, the accretion may be permanently disrupted.)"
+ Large UV sizes reffig:fwhm)) and similarities with. XUV-disk phenomenon (Thilkeretal.2007:Donovan2009) point towards this scenario. but we currently cannot strongly favor IGM aceretion over minor mergers as the main source of gas in our sample.," Large UV sizes \\ref{fig:fwhm}) ) and similarities with XUV-disk phenomenon \citep{thilker,donovan} point towards this scenario, but we currently cannot strongly favor IGM accretion over minor mergers as the main source of gas in our sample."
+ In closing we turn to the nature of the UV emission of the compactcentral. source., In closing we turn to the nature of the UV emission of the compact source.
+ While the compactness of UV emission suggests an AGN. ACS/SBC resolution cannot outright rule out a highly nucleated stellar bulge.," While the compactness of UV emission suggests an AGN, ACS/SBC resolution cannot outright rule out a highly nucleated stellar bulge."
+ We see a correlation between the strength of the central source and velocity dispersion that may (through accretion rate being a function of σ. Martinetal. 2007)) suggest an AGN. but again does not exclude stellar origin.," We see a correlation between the strength of the central source and velocity dispersion that may (through accretion rate being a function of $\sigma$; \citealt{martin}) ) suggest an AGN, but again does not exclude stellar origin."
+ While emission lines are too weak (by selection) to allow secure individual BPT classification. averaging the measurements results in strong and weak central sources landing in the AGN (Seyfert) part of the BPT diagram. and those with no central source in the SF/AGN composite region.," While emission lines are too weak (by selection) to allow secure individual BPT classification, averaging the measurements results in strong and weak central sources landing in the AGN (Seyfert) part of the BPT diagram, and those with no central source in the SF/AGN composite region."
+ Finally. galaxies with no extended UV. where we can measure central UV color have FUV—NUV« 0.4. which is more consistent with a blue continuum source (Agiierosetal.2005) than an old stellar population.," Finally, galaxies with no extended UV, where we can measure central UV color have ${\rm
+ FUV}-{\rm NUV}<0.4$ , which is more consistent with a blue continuum source \citep{agueros} than an old stellar population."
+ Early-type galaxies with strong far-UV excess are dominated by SF. except when they are red in the NUV—;. in which case the FUV comes from a compact source. possibly an optically weak AGN.," Early-type galaxies with strong far-UV excess are dominated by SF, except when they are red in the ${\rm NUV}-r$, in which case the FUV comes from a compact source, possibly an optically weak AGN."
+ Extended SF around ETGs has recently been found in someindividual nearby galaxies. but these have stellar masses one to two orders of magnitude smaller (Thilkeretal.2010;Donovan2009).," Extended SF around ETGs has recently been found in some nearby galaxies, but these have stellar masses one to two orders of magnitude smaller \citep{thilker10,donovan}."
+. We extended SF in galaxies up to M.~I0!!M..., We extended SF in galaxies up to $M_*\sim 10^{11} M_{\odot}$.
+" However. strong UV excess Is rare in yet more massive ETGs—possible consequence of AGN feedback and a reflection of an ""E-E dichotomy."," However, strong UV excess is rare in yet more massive ETGs–possible consequence of AGN feedback and a reflection of an `E-E' dichotomy."
+ In the majority of cases the SF gas is probably external in origin (minor mergers or IGM accretion). making these ETGs ‘rejuvenated’.," In the majority of cases the SF gas is probably external in origin (minor mergers or IGM accretion), making these ETGs `rejuvenated'."
+ Most common UV morphology are rings. which likely result from resonances (from yet undetected bars or other disturbances).," Most common UV morphology are rings, which likely result from resonances (from yet undetected bars or other disturbances)."
+ A smaller fraction of UV-excess ETGs may be more typical green-valley galaxies. where we see the declining original SF. activity.," A smaller fraction of UV-excess ETGs may be more typical green-valley galaxies, where we see the declining original SF activity."
+ Similar conclusions regarding the heterogeneity of evolutionary paths of transiting galaxies (but not necessarily ETGs) were recently reached by Cortese&Hughes(2009) based on HI content of nearby (<25 Mpc) galaxies., Similar conclusions regarding the heterogeneity of evolutionary paths of transiting galaxies (but not necessarily ETGs) were recently reached by \citet{cortese} based on HI content of nearby $<25$ Mpc) galaxies.
+ We thank Stefano Zibetti for help with ADAPTSMOOTH and Ray Lucas for help with ACS data., We thank Stefano Zibetti for help with ADAPTSMOOTH and Ray Lucas for help with ACS data.
+ Based on observations made with the NASA/ESA Hubble Space Telescope and supported with NASA grant HST-GO-11158.03., Based on observations made with the NASA/ESA Hubble Space Telescope and supported with NASA grant HST-GO-11158.03.
+The wavelength resolution is AA ~0.06 ((FWHM).,The wavelength resolution is $\Delta\lambda\sim$ 0.06 (FWHM).
+ The wavelength uncertainty in the calibration is a few bbelow 30 aand about 0.020 iin the region above 30Α., The wavelength uncertainty in the calibration is a few below 30 and about 0.020 in the region above 30.
+. The spectra are background subtracted., The spectra are background subtracted.
+ The statistical errors in the line fluxes include errors from the background., The statistical errors in the line fluxes include errors from the background.
+ For further instrumental details see also Brinkman et al. (, For further instrumental details see also Brinkman et al. (
+2000).,2000).
+ Later (on October 22. 2000). the spectrum of Procyon was observed by XMM-Newton using the RGS and EPIC-MOS.," Later (on October 22, 2000), the spectrum of Procyon was observed by XMM-Newton using the RGS and EPIC-MOS."
+" The total observing time was ~ 107 ksee: however. due to large solar flare activity at the end of the observation. we removed 16.7 ksee of data. leaving a total of 90.5 ksec of ""good"" data."," The total observing time was $\approx$ 107 ksec; however, due to large solar flare activity at the end of the observation, we removed 16.7 ksec of data, leaving a total of 90.5 ksec of “good” data."
+ In XMM-Newton three telescopes focus X-rays onto three EPIC cameras (two MOS and one pn)., In XMM-Newton three telescopes focus X-rays onto three EPIC cameras (two MOS and one pn).
+ About half of the photons in the beams of two telescopes (Turner et al., About half of the photons in the beams of two telescopes (Turner et al.
+ 2001) are diffracted by sets of reflecting gratings and are then focussed onto the RGS detectors., 2001) are diffracted by sets of reflecting gratings and are then focussed onto the RGS detectors.
+ The RGS spectral resolution is AA ~0.07A.. with a maximum effective area of about 140 em? around 15A.," The RGS spectral resolution is $\Delta\lambda\sim$ 0.07, with a maximum effective area of about 140 $^2$ around 15."
+. The wavelength uncertainty is 7-8mA.., The wavelength uncertainty is 7-8.
+ The first spectral order has been selected by means of the energy resolution of the individual CCDs., The first spectral order has been selected by means of the energy resolution of the individual CCDs.
+ For further details see den Herder et al. (, For further details see den Herder et al. (
+2001).,2001).
+ The data were processed by the XMM-Newton SAS using the calibration of February 2001., The data were processed by the XMM-Newton SAS using the calibration of February 2001.
+ The RGS cover the range, The RGS cover the range
+" As we see, in general if only one mode dominates in pulsar profile, then we expect statistically D (double) mean profile for the O-mode and 5S (single) one for the X-mode."," As we see, in general if only one mode dominates in pulsar profile, then we expect statistically $D$ (double) mean profile for the O-mode and $S$ (single) one for the X-mode."
+" Of course, we suppose here that there is nonzero natural width of the radiative domains that efficiently smooths the hollow cone corresponding to the narrow X-mode."," Of course, we suppose here that there is nonzero natural width of the radiative domains that efficiently smooths the hollow cone corresponding to the narrow X-mode."
+" Nevertheless, as is shown on Fig. 19,,"," Nevertheless, as is shown on Fig. \ref{CCmod},"
+ the X-mode can in principal produce D profile and O-mode can produce S ones., the X-mode can in principal produce $D$ profile and O-mode can produce $S$ ones.
+"region. On the other hand, if there are two modes with the similar intensity, the triple (T) and multiple (M) profiles can be obtained as well."," On the other hand, if there are two modes with the similar intensity, the triple (T) and multiple (M) profiles can be obtained as well."
+" So, this picture is in agreement with Rankin’s morphological classification (Rankin 1983)."," So, this picture is in agreement with Rankin's morphological classification (Rankin 1983)."
+" Our investigation provides an additional information to morphological structure, i.e., the information of the type of the mode that forms mainly the mean profile."," Our investigation provides an additional information to morphological structure, i.e., the information of the type of the mode that forms mainly the mean profile."
+" As was demonstrated, only the correlation between p.a. derivative and V signs provides us this chance."," As was demonstrated, only the correlation between $p.a.$ derivative and $V$ signs provides us this chance."
+" Finally, the determination of the angular dimension of the core and conal beams depends greatly on the emission radius (i.e., again, on the generation mechanism)."," Finally, the determination of the angular dimension of the core and conal beams depends greatly on the emission radius (i.e., again, on the generation mechanism)."
+" So the detailed analysis is also beyond the scope of current investigation (one can find empirical results on this in, e.g., Maciesiak Gil 2011)."," So the detailed analysis is also beyond the scope of current investigation (one can find empirical results on this in, e.g., Maciesiak Gil 2011)."
+" Thus, the approach developed in this paper allows the predictions of the theory of radio emission to be quantitatively compared with the observational data."," Thus, the approach developed in this paper allows the predictions of the theory of radio emission to be quantitatively compared with the observational data."
+" Moreover, based on the observational data, one can reject a number of models in which the shapes of the profiles of the position angle and the degree of circular polarization never realized in practice are obtained."," Moreover, based on the observational data, one can reject a number of models in which the shapes of the profiles of the position angle and the degree of circular polarization never realized in practice are obtained."
+" Thus, it becomes possible to solve the inverse problem, i.e., to determine the parameters of the outflowing plasma and the magnetic field structure."," Thus, it becomes possible to solve the inverse problem, i.e., to determine the parameters of the outflowing plasma and the magnetic field structure."
+ We will be glad to any collaboration with observers for investigating the observational features of the mean profiles on the base of our theory., We will be glad to any collaboration with observers for investigating the observational features of the mean profiles on the base of our theory.
+ The source code of our calculations (in Mathematica 7 or C) is available under request., The source code of our calculations (in Mathematica 7 or C) is available under request.
+ We thank Prof. A.V. Gurevich and Ya., We thank Prof. A.V. Gurevich and Ya.
+"N. Istomin for his interest and support, and J. Dyks, A. Jessner, P. Jaroenjittichai, M. Kramer, V.V. Kocharovsky, D. Mitra, M.V. Popov, R. Shcherbakov, B. Rudak, and H.-G. Wang for useful discussions.","N. Istomin for his interest and support, and J. Dyks, A. Jessner, P. Jaroenjittichai, M. Kramer, V.V. Kocharovsky, D. Mitra, M.V. Popov, R. Shcherbakov, B. Rudak, and H.-G. Wang for useful discussions."
+ We also thank A. Spitkovsky for valuable opportunity and assistance in dealing with his numerical solution and for outstanding discussions., We also thank A. Spitkovsky for valuable opportunity and assistance in dealing with his numerical solution and for outstanding discussions.
+ This work was partially supported by Russian Foundation for Basic Research (Grant no., This work was partially supported by Russian Foundation for Basic Research (Grant no.
+ 11-02-01021)., 11-02-01021).
+JaLamMeters. as shown iu Tabe 7..,"parameters, as shown in Table \ref{tab:t07}."
+ The top row shows a reasonable estimate for the uncertainties inTay. while giving even wider latitude in .Mog aid Av.," The top row shows a reasonable estimate for the uncertainties in, while giving even wider latitude in ${\Delta}$ and ${\Delta}$."
+ The variations are determined using0 he model aud data [rom star 2210 in NGC 6752. wlich has a that is near the meclian valte Or our sample.," The variations are determined using the model and data from star 2240 in NGC 6752, which has a that is near the median value for our sample."
+ Since the setsitivity to the parameters is uot symametri€. we report in the tabe he worse result of either inereasing OL ¢ecreaslug a glvel parameter.," Since the sensitivity to the parameters is not symmetric, we report in the table the worse result of either increasing or decreasing a given parameter."
+ Oily the 2001 data are usecl in this test since the spectrun has bot1 a [ul COMpLeLaeut of lines ane a hieh S/N. The resulS are not surprisug: plays the bigges role i1 the error budget. wih he low E. P. liues of Ti I ald Cr I being the most seusitive. while the iconized lines are little ated.," Only the 2001 data are used in this test since the spectrum has both a full complement of lines and a high S/N. The results are not surprising: plays the biggest role in the error budget, with the low E. P. lines of Ti I and Cr I being the most sensitive, while the ionized lines are little affected."
+ Except for Ti I ai| E LIL. au increase in leads to an inc‘ease D [abundatce/H].," Except for Ti II and Eu II, an increase in leads to an increase in [abundance/H]."
+ Unceralnties in lead to ouM7 iall errors for eveu the pressure-sensitive lous with some elements ireasing with (AL C:1l. L aud Cr) aud others decreasiug (Ti II. Fe. d. La. aud Eu).," Uncertainties in lead to only small errors for even the pressure-sensitive ions with some elements increasing with (Al, Ca, Ti I, and Cr) and others decreasing (Ti II, Fe, Ni, La, and Eu)."
+ Final certainties due to a'e quie suall. which is cousistent with the many weak lines employed in our analysis.," Finally, uncertainties due to are quite small, which is consistent with the many weak lines employed in our analysis."
+ lucreasiug cdesaturates the liue thereby increasing tle derived abundance. wl1chi s what we fixl in our results. wl1 [AI/Fe] having zero dependence within the uncertainties.," Increasing desaturates the line thereby increasing the derived abundance, which is what we find in our results, with [Al/Fe] having zero dependence within the uncertainties."
+" Of course. the hree parameter""s uot uucorrelatecd. so variations in one necessarily lead to variaious in the others fas cau be seeli he formula forei. whieh is a function of Tay) aud the seusitivities reported i Table 7 can be be uodified to reflect this in the total uucertainty due to the atijospherie models."," Of course, the three parameters are not uncorrelated, so variations in one necessarily lead to variations in the others (as can be seen by the formula for, which is a function of ) and the sensitivities reported in Table \ref{tab:t07} can be be modified to reflect this in the total uncertainty due to the atmospheric models."
+ Since the focus oL this studs is primarily on aluminum. we note that theinodel iucertaiuties introduce ouly ~ 0.1 dex uncertainty in the [Al/Fe] ratio.," Since the focus of this study is primarily on aluminum, we note that themodel uncertainties introduce only ${\sim}$ 0.1 dex uncertainty in the [Al/Fe] ratio."
+ In Tab ed we present the abundaices Of LL stars in NGC 6752 aud 10 in M80., In Table \ref{tab:t08} we present the abundances of 11 stars in NGC 6752 and 10 in M80.
+ These were accej»ted as cluster members after deteWide that their derived [Fe/H] ratios were consistent with the anicipated metallicities of each cluster. around —1.5to—1.6 dex.," These were accepted as cluster members after determining that their derived [Fe/H] ratios were consistent with the anticipated metallicities of each cluster, around $-1.5~{\rm to}~-1.6$ dex."
+ We estimate that tle abunudances are uncertain by about 0.15 ¢ex after accountiug for line-to-line scatter. uncertainties 1li oscilator st'eugthlis and atmospheric parameters. aud iustrumenta errors.," We estimate that the abundances are uncertain by about 0.15 dex after accounting for line-to-line scatter, uncertainties in oscillator strengths and atmospheric parameters, and instrumental errors."
+ The abuudances for tle NGC' 6702 «lata are averaged together wieu available by giving the second-epochli results twice tle welgit of the first to account for he facor of two improvement ir resolution., The abundances for the NGC 6752 data are averaged together when available by giving the second-epoch results twice the weight of the first to account for the factor of two improvement in resolution.
+ With the exception of star 3589. the abundances or NGC 6722 from the the two epocls were cetermiued with a sinee moclel for each star as given in Taο 6..," With the exception of star 3589, the abundances for NGC 6752 from the the two epochs were determined with a single model for each star as given in Table \ref{tab:t06}."
+ For star 3589. we used the separately cle‘ived models given in Table 6 to derive abundalCes for eac‘th epoch before forming tie welglited average.," For star 3589, we used the separately derived models given in Table \ref{tab:t06}
+ to derive abundances for each epoch before forming the weighted average."
+ Of course. ALSO was observed curing OLly 1re second epoch aud the results were derive with the modeS listed in Table 6..," Of course, M80 was observed during only the second epoch and the results were derived with the models listed in Table \ref{tab:t06}."
+ The ir OllaUL@LLCes asstume a solar valte of log e=7.52. -€and the rest of tle abunuclances are on the Aiders&Cirevesse(1989). scale.," The iron abundances assume a solar value of log ${\epsilon}~=~7.52$, and the rest of the abundances are on the \citet{AG89} scale."
+ The aluminum aud laulanui abundances are not corrected for hyperfit espitting (hís). but the euroxui data are.," The aluminum and lanthanum abundances are not corrected for hyperfine splitting (hfs), but the europium data are."
+ TIe hfs line list for 4e A6615 line was provided tο US yw C. Sueden (2002. private communication ancl we used the blends driver in MOOG o cleteriije the abundances.," The hfs line list for the ${\lambda}6645$ line was provided to us by C. Sneden (2002, private communication) and we used the blends driver in MOOG to determine the abundances."
+ The :ΠΠ abuaudaices are derived {rom only A6696A. since this feature is strouger than the other ια of the doublet. ancl is offen visibe even when AGOUS is uo.," The aluminum abundances are derived from only ${\lambda}6696$, since this feature is stronger than the other half of the doublet, and is often visible even when ${\lambda}6698$ is not."
+Most [A/Fe] values were determined from equivalei width cata. exceyyt,"Most [Al/Fe] values were determined from equivalent width data, except"
+we reformulate the reconstruction as a convex optimisation problem and solve it iteratively with an algorithm based on Hybrid Steepest Descent (HSD)(?)..,we reformulate the reconstruction as a convex optimisation problem and solve it iteratively with an algorithm based on Hybrid Steepest Descent (HSD)\citep{Yamada}.
+" We define the multiresolution support M, which is determined by the set of detected significant coefficients after hypothesis testing: We formulate the reconstruction problem as a convex constrained minimization problem: where 4 denotes the IUWT synthesis operator."," We define the multiresolution support $\mathcal{M}$, which is determined by the set of detected significant coefficients after hypothesis testing: We formulate the reconstruction problem as a convex constrained minimization problem: where $\mathbf{\Phi}$ denotes the IUWT synthesis operator."
+" This problem is solved with the following iterative scheme: the image is initialised by X)=0, and the iteration scheme is, for n=0 to Nmax— 1:"," This problem is solved with the following iterative scheme: the image is initialised by $\mathbf{X}^{(0)} = 0$, and the iteration scheme is, for $n=0$ to $N_{\max}-1$ :"
+o expected at high resolutions.,to expected at high resolutions.
+ So the simulations only place lower limis OI a eastrement of he decay timescale of the magnetic flux threadiug the black hole event horizon., So the simulations only place lower limits on a measurement of the decay timescale of the magnetic flux threading the black hole event horizon.
+ A convergent solution implies the result is indepeudent of auy microscopic resistivity 1jodel and onv deJens ipon the self-consistently generated turbuleut resistivity., A convergent solution implies the result is independent of any microscopic resistivity model and only depends upon the self-consistently generated turbulent resistivity.
+ A non-convergent result woulc impy the uicroscopic resistivity controls the clissipation aud must be specified., A non-convergent result would imply the microscopic resistivity controls the dissipation and must be specified.
+ Fo| such a uon-convergent Case. we cau estimate the expected resistivity (Uzdeusky&Λοιπον2011).. such as cOne in the jext section.," For such a non-convergent case, we can estimate the expected resistivity \citep{um11}, such as done in the next section."
+" Figure 1 shows the initial aud. quasi-steady contour plots for V=RA, lor the axisvuujetric 256x1281 resolution aud black hole with spiu «—0.99.", Figure 1 shows the initial and quasi-steady contour plots for $\Psi=R A_\phi$ for the axisymmetric $256\times 128\times 1$ resolution and black hole with spin $a=0.99$.
+ The initial dipole field has collapse to essentially a split monopole field with most of the magnetic [lux passing through the οσοι iustead of reconnecting near the equator., The initial dipole field has collapsed to essentially a split monopole field with most of the magnetic flux passing through the horizon instead of reconnecting near the equator.
+ This was predicted by Lyutikov(2011b)., This was predicted by \cite{2011arXiv1104.1091L}.
+. All other models at different spins show qualitatively similar final states. except the 3D simulatious with resolution 256x12816 or 256x12832 show slightly more closed field structures at the equator beyouc the event horizon.," All other models at different spins show qualitatively similar final states, except the 3D simulations with resolution $256\times 128\times 16$ or $256\times
+128\times 32$ show slightly more closed field structures at the equator beyond the event horizon."
+ Figure 2 shows the value of magnetic flux dj vs. time (/) for all black hole spius aud botl the force-[ree electrodyuaimics equatious of motion (racliatively ellicient regime) and the MHD equations of motion (radiatively inellicient regime)., Figure 2 shows the value of magnetic flux $\Phi_{\rm EM}$ vs. time $t$ ) for all black hole spins and both the force-free electrodynamics equations of motion (radiatively efficient regime) and the MHD equations of motion (radiatively inefficient regime).
+ At /=0. the black hole rotation begins to drive radial currents into the magnetosphere.," At $t=0$, the black hole rotation begins to drive radial currents into the magnetosphere."
+ Then. during the first /~2067M/c'*. some portion ol the dipole fiekl is absorbed by the black hole.," Then, during the first $t\sim 20GM/c^3$, some portion of the dipole field is absorbed by the black hole."
+ For bigh black hole spin rales. the auount of lost jagnetic [Tux is fouud to correspond to losing those field lines that do uot cross the Allvénn surface (R~OpeczzOy/(2c). for field rotation frequency Qy-) at late times.," For high black hole spin rates, the amount of lost magnetic flux is found to correspond to losing those field lines that do not cross the Alfvénn surface $R\sim \Omega_F/c\approx \Omega_H/(2c)$, for field rotation frequency $\Omega_F$ ) at late times."
+ This means that those field ines that eventually ypen-up to infinity are prevented from beiug absorbed by the black hole by he poloidal currents driven into the maguetosphere by the rotation of the |slack hole., This means that those field lines that eventually open-up to infinity are prevented from being absorbed by the black hole by the poloidal currents driven into the magnetosphere by the rotation of the black hole.
+ For low black ole spin rates. such as for a=0. the early loss of magnetic flux is limited by the eqratorial toroidal current sheet that forms over a titnescale of /AL25G/c. which correspouds to tie light-crossiug . ] ∐∐≺↵⊳∖∢∙⋜↕↥≺↵∩∖⇁≺↵↥⋅⋜↕∖∖↽⋜↕∖⊽↩↥↩∐∑≟↕∐∩↥⋋≈−≻↓⋅↙∪⋀⊔↙∕∣≬⋡−↥∩↕⋅≺↵↥≺↵∢∙↕⋅∩⋯⋜↕∑≟∐≺↲⊔∢∙↓⋅⋜↕≺∐⋜↕↕i−⋅⊾ dp ⋅ ⋅ion from the decay of a lipole Geld for a vacuum black hole (Batungarte&Shapiro2003).," For low black hole spin rates, such as for $a=0$, the early loss of magnetic flux is limited by the equatorial toroidal current sheet that forms over a timescale of $t\sim 25GM/c^3$, which corresponds to the light-crossing timescale over a wavelength of $\lambda\approx 24.7GM/c^2$ for electromagnetic radiation from the decay of a dipole field for a vacuum black hole \citep{bs03}."
+. This meaus that the decay by vacuum radiatiou is avoided by toroidal current sheet formation due to collapse of the dipole field in the presence ofz plasiua., This means that the decay by vacuum radiation is avoided by toroidal current sheet formation due to collapse of the dipole field in the presence of a plasma.
+ Overall. decay of the dipole field is avoided by poloidal aud toroidal currents that spontaneously form due to black hole rotation aud collapse of the dipole field iu tle preseuce of a plasma.," Overall, decay of the dipole field is avoided by poloidal and toroidal currents that spontaneously form due to black hole rotation and collapse of the dipole field in the presence of a plasma."
+ For the MHD solutions. over the first /~50M/c*. the solution then settles into an state of a roughly constant magnetic dissipation uear tlie black hole.," For the MHD solutions, over the first $t\sim 50GM/c^3$, the solution then settles into an quasi-steady state of a roughly constant magnetic dissipation near the black hole."
+ The dissipation creates plasma pressure within the current laver that supports tlie laver against magnetic reconnection., The dissipation creates plasma pressure within the current layer that supports the layer against magnetic reconnection.
+ The dlecays are fit well by an exponential decay. where the decay timescale for the 256x128 resolution [or the MHD models with «={0.0.1.0.5.0.9.0.99} is 7~[100.110.120.100.CM/5001. respectively.," The decays are fit well by an exponential decay, where the decay timescale for the $256\times 128$ resolution for the MHD models with $a=\{0,0.1,0.5,0.9,0.99\}$ is $\tau\sim
+\{100,110,120,400,500\}GM/c^3$, respectively."
+ As expected for the discontinuity resolved at the numerical grid scale. the decay timescale is found to be directly proportional to the resolution (i.e. 7x{ΝεNo.NS ]).," As expected for the discontinuity resolved at the numerical grid scale, the decay timescale is found to be directly proportional to the resolution (i.e. $\tau\propto \{N_r,N_\theta,N_\phi\}$ )."
+ Otherwise identical 3D models, Otherwise identical 3D models
+abundant low velocity dispersion svstems in the redshift survey with convergence peaks.,abundant low velocity dispersion systems in the redshift survey with convergence peaks.
+ Our definitions are similar to those adopted by ILamana et al. (, Our definitions are similar to those adopted by Hamana et al. (
+2004) in their simulations of (he efficiency and completeness of weak lensing survevs for massive cluster identification.,2004) in their simulations of the efficiency and completeness of weak lensing surveys for massive cluster identification.
+ Schirmer et al. (, Schirmer et al. (
+2007) and Gavazzi Soucail (2007) develop uniform procedures [or evaluating the coincidence of concentrations of galaxies wilh convergence map peaks.,2007) and Gavazzi Soucail (2007) develop uniform procedures for evaluating the coincidence of concentrations of galaxies with convergence map peaks.
+ The approaches of both Schirmer et al. (, The approaches of both Schirmer et al. (
+2007) ancl Gavazzi Soucail (2007) are based on complete photometric survevs of (their fields.,2007) and Gavazzi Soucail (2007) are based on complete photometric surveys of their fields.
+ Thev evaluate the significance of cluster candidates internal to their survevs., They evaluate the significance of cluster candidates internal to their surveys.
+ Geller et al. (, Geller et al. (
+2010) develop an analogous approach for evaluating the coincidence of concentrations of galaxies in a foreground redshift survey with convergence map peaks.,2010) develop an analogous approach for evaluating the coincidence of concentrations of galaxies in a foreground redshift survey with convergence map peaks.
+ Here we follow the Geller et al. (, Here we follow the Geller et al. (
+2010) approach.,2010) approach.
+ Because (he redshilt survey is substantially complete we can sample the survey. itself to evaluate the significance of galaxy condensations along the line-ol-sight to lensing peaks., Because the redshift survey is substantially complete we can sample the survey itself to evaluate the significance of galaxy condensations along the line-of-sight to lensing peaks.
+ The sensitivitv plot (Figure 9)) suggests that. we need not look deeper than z~0.75 to identify clusters corresponding to the six weak lensing peaks wilh »>3.7 in the Subaru weak lensing map., The sensitivity plot (Figure \ref{fig:sensitivity.ps}) ) suggests that we need not look deeper than $z \sim 0.75$ to identify clusters corresponding to the six weak lensing peaks with $\nu > 3.7$ in the Subaru weak lensing map.
+ The most distant svstems must be very rich and they should be obvious in (he photometric data., The most distant systems must be very rich and they should be obvious in the photometric data.
+ We use the Re-band counts in the image (o check for such distant svslens., We use the $_C$ -band counts in the image to check for such distant systems.
+ We examine the faint Subaru Re band counts (Mivazaki et al., We examine the faint Subaru $_C$ band counts (Miyazaki et al.
+ 2007) in the range 2]«Re<23 by binning the data in 0.3 bins and smoothing the map of counts with a 1.5 Gaussian (the same smoothing aud pixel scale as the Subaru weak lensing map)., 2007) in the range $21 < R_C < 23$ by binning the data in $^{\prime}$ bins and smoothing the map of counts with a $^{\prime}$ Gaussian (the same smoothing and pixel scale as the Subaru weak lensing map).
+ The range in apparent magnitude just overlaps the rj;<21.3 magnitude limit of our main redshift survev sample., The range in apparent magnitude just overlaps the $r_{petro} < 21.3$ magnitude limit of our main redshift survey sample.
+ The binning and smoothing leneths are essentially the same as those for the weak lensing map., The binning and smoothing lengths are essentially the same as those for the weak lensing map.
+ We define σς as the standard deviation of the pixel value in the map of counts and we note peaks that rise Z264 above the median as corroborating evidence for svslenms of galaxies along the line-oF-sight toward weak lensing peaks and/or as evidence [or svstenms bevond (he reach of the redshift survey., We define $\sigma_S$ as the standard deviation of the pixel value in the map of counts and we note peaks that rise $\gtrsim 2\sigma_S$ above the median as corroborating evidence for systems of galaxies along the line-of-sight toward weak lensing peaks and/or as evidence for systems beyond the reach of the redshift survey.
+ Table 3. indicates (he lensing peak directions where there is a count excess., Table \ref{tbl:VDisp} indicates the lensing peak directions where there is a count excess.
+ Probes through the redshift survey centered on each galaxy in the survey contain many ol (he selection effects which impact a probe toward a lensing peak., Probes through the redshift survey centered on each galaxy in the survey contain many of the selection effects which impact a probe toward a lensing peak.
+ Thus comparison of these probes with similar probes toward the peaks should provide a measure of the significance of cluster candidates in (he redshift survey., Thus comparison of these probes with similar probes toward the peaks should provide a measure of the significance of cluster candidates in the redshift survey.
+ We have demonstrated bv experiment (hat centering of probes on galaxies or on random points in the region has no effect on the results of our, We have demonstrated by experiment that centering of probes on galaxies or on random points in the region has no effect on the results of our
+1998).,.
+. Thev reside at cistances of 2—8 kpe typically (Egan 1993).. and most are located roughly toward the Galactic centre.," They reside at distances of $2-8$ kpc typically \citep{egan98,carey98}, and most are located roughly toward the Galactic centre."
+ The IRDCs have unusual physical characteristics compared (to nearby examples of molecular clouds., The IRDCs have unusual physical characteristics compared to nearby examples of molecular clouds.
+ Egan found that IRDCs appear to have properties intermediate between massive molecular cloud cores (nzz104—105 7: Tz230—100 IK) and isolated dark clouds (ηz107—10/em Tsz8—15 Ix). having densities (wpical of molecular cloud cores. but temperatures (tvpical of isolated dark clouds.," \citet{egan98} found that IRDCs appear to have properties intermediate between massive molecular cloud cores $n\approx 10^4-10^6$ $^{-3}$ ; $T\approx 30-100$ K) and isolated dark clouds $n\approx 10^2-10^4 {\rm cm}^{-3}$; $T\approx 8-15$ K), having densities typical of molecular cloud cores, but temperatures typical of isolated dark clouds."
+ Many URDC! filaments appear to be at a relatively early phase of evolution. with fewer signposts of star formation (han nearby molecular clouds.," Many IRDC filaments appear to be at a relatively early phase of evolution, with fewer signposts of star formation than nearby molecular clouds."
+ The IRDCs and nearby molecular clouds share the common feature that a large fraction of both populations are filamentary. but the IRDC's are olten less fragmented.," The IRDCs and nearby molecular clouds share the common feature that a large fraction of both populations are filamentary, but the IRDCs are often less fragmented."
+" Relatively ""prisüne ΠΟ filaments have simpler density structures than. lvagmented. star-forming filaments. which makes them good candidates for detailed comparison with theoretical models."," Relatively “pristine” IRDC filaments have simpler density structures than fragmented, star-forming filaments, which makes them good candidates for detailed comparison with theoretical models."
+ One drawback is (hat ΗλCs are more distant than many other filaments and are therefore less well resolved., One drawback is that IRDCs are more distant than many other filaments and are therefore less well resolved.
+ The GI1.11-0.12 cloud studied in this paper is a good example of a filamentary IC. which was observed by SCUBA in 3504 emission Careyetal.(2000).. following its discovery bv the AISN survey.," The G11.11-0.12 cloud studied in this paper is a good example of a filamentary IRDC, which was observed by SCUBA in $\micron$ emission \citet{carey00}, following its discovery by the MSX survey."
+ Johnstoneelal.—(2003) (herealter Paper I) fit a non-magnetic equilibrium filament model to the radial structure of (vo regions of GI1.11-0.12., \citet{johnstone03} (hereafter Paper I) fit a non-magnetic equilibrium filament model to the radial structure of two regions of G11.11-0.12.
+ This paper presents a more thorough analvsis. which compares the SCUBA &50;a emission with 3 existing models of filamentary clouds in detail.," This paper presents a more thorough analysis, which compares the SCUBA $\micron$ emission with 3 existing models of filamentary clouds in detail."
+ These models are (he pressure-truncated non- Ostriker (Ostriker1964). model. a generalized Stodólkiewicez(1963) (hereafter GS) helical field model. and the Fiege Pudritz (hereafter FP) helical field model.," These models are the pressure-truncated non-magnetic Ostriker \citep{ostriker64} model, a generalized \citet{stod63} (hereafter GS) helical field model, and the Fiege Pudritz (hereafter FP) helical field model."
+ We discuss these models in the next section., We discuss these models in the next section.
+ Note that the Ostriker model is the zero magnetic field limit of both the GS and EP models. but this limit is sulliientlv important to warrant separale treatment.," Note that the Ostriker model is the zero magnetic field limit of both the GS and FP models, but this limit is sufficiently important to warrant separate treatment."
+ This paper uses a novel genetic algorithim-based. technique (to explore the parameter space of each model completely and to fit solutions to the data., This paper uses a novel genetic algorithm-based technique to explore the parameter space of each model completely and to fit solutions to the data.
+ This is the first astronomical research application for a new multü-object genetic algorithm (“Ferret”) developed by one of the authors (J.F.)., This is the first astronomical research application for a new multi-object genetic algorithm (“Ferret”) developed by one of the authors (J.F.).
+ The technique is discussed in Section 5.., The technique is discussed in Section \ref{sec:GAs}.
+ The goal of this paper is a thorough comparison of the GI11.11-0.12 IRDC with 3 different. theoretical models of filamentary molecular clouds., The goal of this paper is a thorough comparison of the G11.11-0.12 IRDC with 3 different theoretical models of filamentary molecular clouds.
+ These models have different magnetic field configurations.or no magnetic field at all in the case of the Ostriker mocel.," These models have different magnetic field configurations,or no magnetic field at all in the case of the Ostriker model."
+Having cletermined the properties of the flare source region. we cau compute its obscuring elfect ou,"Having determined the properties of the flare source region, we can compute its obscuring effect on"
+Data were reduced using the and software packages.,Data were reduced using the and software packages.
+ Spectra were calibrated using HIPE and then exported to CLASS format with the shortest possible pre- (typically ~40 seconds)., Spectra were calibrated using HIPE and then exported to CLASS format with the shortest possible pre-integration (typically $\sim40$ seconds).
+ For each scan we combined the four sub-bands in each polarization to create a GGHz spectrum., For each scan we combined the four sub-bands in each polarization to create a GHz spectrum.
+ From this spectrum we computed the underlying continuum using the line-free channels; we then subtracted first-order baselines from individual sub-bands., From this spectrum we computed the underlying continuum using the line-free channels; we then subtracted first-order baselines from individual sub-bands.
+ The baseline-subtracted sub-bands were again combined and the continuum level added., The baseline-subtracted sub-bands were again combined and the continuum level added.
+ This results in a noise-weighted GGHz spectrum for each scan., This results in a noise-weighted GHz spectrum for each scan.
+" These spectra were inspected for remaining baseline instabilities, and scans with distorted baselines were omitted."," These spectra were inspected for remaining baseline instabilities, and scans with distorted baselines were omitted."
+ For each line we inspected the co-added result in both polarizations (H and V) separately., For each line we inspected the co-added result in both polarizations (H and V) separately.
+ The continuum level was found to agree better than for each frequency., The continuum level was found to agree better than for each frequency.
+" The integrated line intensities agree within without significant differences of the line profiles (although the latter comparison is limited by the signal-to-noise ratio, in particular for the GGHz observations)."," The integrated line intensities agree within without significant differences of the line profiles (although the latter comparison is limited by the signal-to-noise ratio, in particular for the GHz observations)."
+" In the following we therefore use the noise-weighted average of both polarizations, which yields effective on-source integration times for the three lines are 750, 2250, and 1300 seconds for theo-H5O(110-101).,p-H20(111-0o9),, and Ilines, respectively."," In the following we therefore use the noise-weighted average of both polarizations, which yields effective on-source integration times for the three lines are 750, 2250, and 1300 seconds for the, and lines, respectively."
+" Since HIFI's calibration is still preliminary, we used the theoretical predictions based on the Herschel's expected surface accuracy and the geometrical aperture size to convert the antenna temperatures to flux density (Kramer 2006))."," Since HIFI's calibration is still preliminary, we used the theoretical predictions based on the Herschel's expected surface accuracy and the geometrical aperture size to convert the antenna temperatures to flux density (Kramer \cite{kramer06}) )."
+" This yields 462 KK!, 467 KK!, and 470 KK""! at GGHz, GGHz, and GGHz, respectively."," This yields 462 $^{-1}$, 467 $^{-1}$, and 470 $^{-1}$ at GHz, GHz, and GHz, respectively."
+" The final spectra are shown at a velocity resolution of 20, 10, and 25 iin reflines (top)."," The final spectra are shown at a velocity resolution of 20, 10, and 25 in \\ref{lines} (top)."
+" All three water lines have been detected with high significance, demonstrating that faint (few mK) broad lines can be observed with HIFI."," All three water lines have been detected with high significance, demonstrating that faint (few mK) broad lines can be observed with HIFI."
+ The lline is detected in emission and shows a double-peaked line profile., The line is detected in emission and shows a double-peaked line profile.
+" Within noise uncertainty, no emission (or absorption) is detected on the systemic velocity of M82 (visa=225 kms))."," Within noise uncertainty, no emission (or absorption) is detected on the systemic velocity of M82 $_{\rm LSR}=225$ )."
+" Both components are well fit by Gaussian profiles, with"," Both components are well fit by Gaussian profiles, with"
+Ro/Co~ OUbs ffes)., 640 ).
+A)This defines the rising πανnescale of the colision-induced clectromaguetic flash., This defines the rising timescale of the collision-induced electromagnetic flash.
+" The αιαον deposit rate during this stage is E~E/ty),9«107eresμιμα dyes).This is ~L8«107 nues the Edcdinetoi luminosity of the Jovian planet. LadaALστ95«lOeretiny (where e. G. any. στ IxcGinare the seed of light. exavitatiouals constant. xoton nass. and Thomson cross section. respectively)."," The energy deposit rate during this stage is $\dot E \sim E/\tau_{sh} \sim 9\times 10^{37}~{\rm erg~s^{-1}}
+(m_\oplus m_J/r_J) (r_\oplus /c_{s})$ .This is $\sim 1.8\times 10^3$ times the Eddington luminosity of the Jovian planet, $L_{\rm Edd} = 4\pi cG m_pM/\sigma_T \sim
+5\times 10^{34} {\rm
+erg~s^{-1}} m_J$ (where $c$, $G$, $m_p$, $\sigma_T$ are the speed of light, gravitational constant, proton mass, and Thomson cross section, respectively)."
+ About one half of tie collisional energy is converted to internal energv. of wich only a small fraction is racliated xoniptly: while a large fraction is sunk as latent heat.," About one half of the collisional energy is converted to internal energy, of which only a small fraction is radiated promptly, while a large fraction is sunk as latent heat."
+ The other half of the energy is initially in kinetic form. driving a radiatiou-supported expanding cuvelope aud inducing oscillation aud convection within the eiaut planet interior.," The other half of the energy is initially in kinetic form, driving a radiation-supported expanding envelope and inducing oscillation and convection within the giant planet interior."
+ This kinetic energy will be thermalized later aud radiated away over a longer period of time along with the lateut heat. as an IR afterglow (see below).," This kinetic energy will be thermalized later and radiated away over a longer period of time along with the latent heat, as an IR afterglow (see below)."
+ In the prompt phase. photons must be trapped bv the high opacity iu the reeious where the heat is initially deposited. aud the peak bolometric huuinositv is — ying ABEL cap. where L.=3.8&10°?eres tis the solar huuinositw. ys Lis a factor correcting for less than optimal emissiou ecolctry and other radiative ineficiencics.," In the prompt phase, photons must be trapped by the high opacity in the regions where the heat is initially deposited, and the peak bolometric luminosity is = m_J 13 m_J, where $L_\odot = 3.8\times 10^{33}~{\rm erg~s^{-1}}$ is the solar luminosity, $\eta \leq 1$ is a factor correcting for less than optimal emission geometry and other radiative inefficiencies."
+ Since Ex Lg. o Lis likely achieved.," Since $\dot E
+\gg L_{\rm Edd}$ , $\eta \sim 1$ is likely achieved."
+ The spectrua is esseutiallv thevinal from a hot with radius ~Ry. with a nonthermal hard tail caused by Comptonization of the thermal photons by the non-thermal clectrous accelerated from the shocks.," The spectrum is essentially thermal from a hot with radius $\sim R_2$, with a non-thermal hard tail caused by Comptonization of the thermal photons by the non-thermal electrons accelerated from the shocks."
+ The typical thermal temperature at the peak time 15 LISIORamE nOI/2.€ which is insensitive to both the planet mass aud the nukuown radiation⋅⋅ efficiency⋅⋅ (Le.⋅ X(gHti1/1 ).," The typical thermal temperature at the peak time is 10^5 , which is insensitive to both the planet mass and the unknown radiation efficiency (i.e. $\propto \eta^{1/4} m_J^{1/4}$ )."
+ The . . ↻↸∖⋜∐↘↽↕↥⋅↸∖≺∣⋯∖∐↸⊳⋅↖↽↕↴∖↴⊼∕∣↗∕⊽⊽∶∩∙∩∖⋅⋅↽↣⊥∩⊥⊽↾∐∑∣∕↓‘ny⊥⊥⊥⊐ri Gvaveleuegth ~150A and euergv ~28 eV).," The peak frequency is $\nu_{pk}
+= 6.6\times 10^{15} ~{\rm Hz}~ \eta^{1/4} m_J^{1/4}
+r_\oplus^{-1/2}$ (wavelength $\sim 450 {\rm \AA}$ and energy $\sim 28$ eV)."
+ This is in the extreme ultraviolet (EUV) baud., This is in the extreme ultraviolet (EUV) band.
+" The peak flux is where P. k. aud c are the Plauck's.Boltzuium's aud Stefan-Doltzuiuun coustauts. respectively,"," The peak flux is (pk) = )^2 =60, where $h$, $k$, and $\sigma$ are the Planck's,Boltzmann's and Stefan-Boltzmann constants, respectively."
+ The hard. nou-thermal tail is expected to extend into the soft-N-ray baud., The hard non-thermal tail is expected to extend into the soft-X-ray band.
+" For the thermal spectrmu. the flax around v could be estimated as F(v)~(R/D)PBAT\v. where D,OP) is the Plauk function. Ris the emission radius. aud D is the distance to Earth."," For the thermal spectrum, the flux around $\nu$ could be estimated as $F(\nu) \sim (R/D)^2 B_\nu(T) \nu$, where $B_\nu(T)$ is the Plank function, $R$ is the emission radius, and $D$ is the distance to Earth."
+" One can define a ratioftv)=F(v.p)/PF(v.s) as the fiux contrast between the planet and the star at a certain frequency i, where F(r.p) aud Εν.) are the fluxes of the planet and the star in the band around v. respectively,"," One can define a ratio$f(\nu) \equiv
+{F(\nu,p)}/{F(\nu,*)}$ as the flux contrast between the planet and the star at a certain frequency $\nu$, where $F(\nu,p)$ and $F(\nu,*)$ are the fluxes of the planet and the star in the band around $\nu$, respectively."
+ At the flash peak time. the EUV fare can greatly out-shine a main-sequence host star later than Bh. or a white dwarf host star.," At the flash peak time, the EUV flare can greatly out-shine a main-sequence host star later than B5, or a white dwarf host star."
+ In the optical and IR bands. high contrasts are also achievable for late-type host stars.," In the optical and IR bands, high contrasts are also achievable for late-type host stars."
+" For example. for a C2 (solar-type) star. the U. V. aud I bands contrasts at the peak time are F(U.pk)~0.2. F(V.pk)0.02. and f(Lpk)~0.008. respectively,"," For example, for a G2 (solar-type) star, the U, V, and I bands contrasts at the peak time are $f({\rm U, pk})\sim 0.2$, $f({\rm V,pk}) \sim 0.02$, and $f({\rm
+I,pk}) \sim 0.008$, respectively."
+ Even higher contrasts are achievable when a cooler star is adopted., Even higher contrasts are achievable when a cooler star is adopted.
+" For f=0.01. the flare is detectable with photometric monitoring of eround-based telescopes (οιο, Borucki et al."," For $f\geq 0.01$, the flare is detectable with photometric monitoring of ground-based telescopes (e.g. Borucki et al."
+ 2001)., 2001).
+ Fixing a particular baud. a lighteurve could be depicted.," Fixing a particular band, a lightcurve could be depicted."
+ The rising liehteurve is very steep. with a time scale of ~Tuv10 minutes.," The rising lightcurve is very steep, with a time scale of $\sim \tau_{sh}
+\sim 10$ minutes."
+ The decay time scale after the peal is likely not very long. based ou existing observational data of other related phenomena. c.g. Comet Shoemaker-Levy ϱ vs. Jupiter collision (Orton et al.," The decay time scale after the peak is likely not very long, based on existing observational data of other related phenomena, e.g. Comet Shoemaker-Levy 9 vs. Jupiter collision (Orton et al."
+ 1995). ταν bursts in accreting neutrou stars (Lewin. van Paradij Taam 1993) aud soft gamma-ray οσα flares from maguctars (Feroci et al.," 1995), X-ray bursts in accreting neutron stars (Lewin, van Paradijs Taam 1993) and soft gamma-ray giant flares from magnetars (Feroci et al."
+ 2001)., 2001).
+ The prompt radiation pressure would drive a fraction of the iupactiug materials outwards. so hat initially the Edcdingtou-liuited photosphere expands and then cools.," The prompt radiation pressure would drive a fraction of the impacting materials outwards, so that initially the Eddington-limited photosphere expands and then cools."
+ As the radiation pressure eradually drops. he photosphere contracts and the materials falls back outo the Jovian planet surface.," As the radiation pressure gradually drops, the photosphere contracts and the materials falls back onto the Jovian planet surface."
+ This fallback refreshes he surface heating. leading to continually refreshed clectromaguetic flashes with descending maguitudes.," This fallback refreshes the surface heating, leading to continually refreshed electromagnetic flashes with descending magnitudes."
+ A jest guess ds that the prompt flash will follow a decaying ail extending up to several 10's of the rising time scale., A best guess is that the prompt flash will follow a decaying tail extending up to several 10's of the rising time scale.
+" The typical timescale for the prompt flash during which he eiiissiou huuinosity is above. say 50% of £,,, could be estimated with. the averaged bolometric. Iuninosity lastLOL ajay."," The typical timescale for the prompt flash during which the emission luminosity is above, say $50\%$ of $L_{pk}$ could be estimated _1 10 2 /c_s), with the averaged bolometric luminosity $\bar L_{bol}^{flash} \sim 10 
+L_\odot \eta m_J$ ."
+" This is the first important time scaleLf"" in the problems.", This is the first important time scale in the problem.
+ The total radiated energy curving this prompt flash stage j EfeLiτι--ὃὁν1003107.ere. which: is: much less than the total impact energw. ie. <E~6x107 eres.," The total radiated energy during this prompt flash stage is $E^{flash} \sim \bar L_{bol}^{flash} \tau_1 \sim 3 \times
+10^{38}~{\rm erg}$, which is much less than the total impact energy, i.e. $\ll E \sim 6\times 10^{40}$ ergs."
+ So the bulk of the collisiou-deposited energy isnof radiated prompth., So the bulk of the collision-deposited energy is radiated promptly.
+ Rather. it is mixed iuto deep inside of the plauct and stored as heat.," Rather, it is mixed into deep inside of the planet and stored as heat."
+ It is released during auch lousger time scale through a decaviug thermalradiation afterelow., It is released during a much longer time scale through a decaying thermal-radiation afterglow.
+ Model simulations reveal that at later times a Jovian planet follows an ορΊσα bolometric linunosity law (Black 1980: Sammon et al., Model simulations reveal that at later times a Jovian planet follows an empirical bolometric luminosity law (Black 1980; Saumon et al.
+" 1996). Le. L/L.~10Pe177, and Tage~οBKett 977, where fg is the age in unit of Myr. andthe relations are valid for tg> 1."," 1996), i.e., $L/L_\odot
+\sim 10^{-5} m_J^{2.35} t_6^{-1.22}$, and $T_{\rm eff} \sim 860 ~{\rm
+K}~ m_J^{0.51} t_6^{-0.28}$ , where $t_6$ is the age in unit of Myr, andthe relations are valid for $t_6 > 1$ ."
+ Notice that the luminosity decaving index 1.227 1. so that the total cuerey released duiug fo606 vy is nesligible compared with the euerev released during the carly epoch.," Notice that the luminosity decaying index $\sim 1.22 > 1$ , so that the total energy released during $t> 10^6$ yr is negligible compared with the energy released during the early epoch."
+" At eulier times. generally we expect a huninosity decaviug law £=Ly(t/ty)"" with"," At earlier times, generally we expect a luminosity decaying law $L=L_0 (t/t_0)^{-a}$ with"
+2006).,.
+. Note that the true gas velocities are likely to be higher because we are measuring only the motion along the line-of-sight., Note that the true gas velocities are likely to be higher because we are measuring only the motion along the line-of-sight.
+ Our multi-wavelength observations indicate that $G1120-S1 does not have a significant AGN component (Table 1)): our 1.4 GHz and 24jzm measurements show that the galaxy falls on the infrared-radio relation for local star-forming galaxies, Our multi-wavelength observations indicate that SG1120-S1 does not have a significant AGN component (Table \ref{tab:sour}) ): our 1.4 GHz and $24\mu$ m measurements show that the galaxy falls on the infrared-radio relation for local star-forming galaxies.
+" SG1120-S1 is X-ray detected but with (Yunan X-ray 2001)..luminosity of ~1032ergs~!, any embedded AGN must be Compton thick2011)."," SG1120-S1 is X-ray detected but with an X-ray luminosity of $\sim 10^{42} \ergsec$, any embedded AGN must be Compton thick."
+". For comparison, most U/LIRGs in (Nardinithis merger stage with similar Lrr are starburst or composite systems"," For comparison, most U/LIRGs in this merger stage with similar $L_{IR}$ are starburst or composite systems."
+ Star formation in U/LIRGs also tends (Yuanto be concentrated2010).. in the central kpc which corresponds to less than one spaxel in our IFU maps., Star formation in U/LIRGs also tends to be concentrated in the central kpc which corresponds to less than one spaxel in our IFU maps.
+ 'The lines in the core spaxels are emitted primarily by gas that is participating in the outflow and their line ratios appear to be dominated by shock excitation as is commonly seen in starburst galaxies2010)., The lines in the core spaxels are emitted primarily by gas that is participating in the outflow and their line ratios appear to be dominated by shock excitation as is commonly seen in starburst galaxies.
+" We conclude that SG1120-S1(Sharp is not AGN dominated, thus the strong gas outflows detected in both the IFU and single slit spectroscopy are driven primarily by star formation."," We conclude that SG1120-S1 is not AGN dominated, thus the strong gas outflows detected in both the IFU and single slit spectroscopy are driven primarily by star formation."
+required in the observational samples.,required in the observational samples.
+" This is done in order to expand the dynamic range in the plots, but none of our results depends on this choice."," This is done in order to expand the dynamic range in the plots, but none of our results depends on this choice."
+ The for each bin is calculated., The for each bin is calculated.
+" This is given by the ratio of mass in metals to total cold gas mass, normalised to the solar metalabundance?:: Zeoiq9.0+log(Mcoa,z /Meoia/0.02)."," This is given by the ratio of mass in metals to total cold gas mass, normalised to the solar metal: $Z_{\textnormal{cold}}=9.0+\textnormal{log}(M_{\textnormal{cold,Z}}/M_{\textnormal{cold}}/0.02)$ ."
+" The M,-Z relation for our z=0 model sample is shown in Fig. 7..", The $M_{*}$ $Z$ relation for our $z=0$ model sample is shown in Fig. \ref{fig:L-Gals_ave_z=0}.
+" We see a positive correlation between stellar mass and metallicity, as well as a segregation of this relation by SFR."," We see a positive correlation between stellar mass and metallicity, as well as a segregation of this relation by SFR."
+" Interestingly, this segregation is similar to that seen in Sample T2, with a reversal in the SFR-dependence from low to high masses."," Interestingly, this segregation is similar to that seen in Sample T2, with a reversal in the SFR-dependence from low to high masses."
+" Another way of stating this result is that there is a ‘turnover’ at high stellar masses in the M,-Z relation for galaxies.", Another way of stating this result is that there is a `turnover' at high stellar masses in the $M_{*}$ $Z$ relation for galaxies.
+" Such a feature is hinted at in Sample T2, but we note that the number of real massive galaxies with emission lines that are strong enough to measure metallicity and that are not dominated by AGN emissionis quite small."," Such a feature is hinted at in Sample T2, but we note that the number of real massive galaxies with emission lines that are strong enough to measure metallicity and that are not dominated by AGN emissionis quite small."
+ This may explain the, This may explain the
+The aim of this paper is to investigate the scatter evolution of supermassive black hole for dry. merger driven erowth.,The aim of this paper is to investigate the scatter evolution of supermassive black hole for dry merger driven growth.
+ Despite the interesting insights we have neglected another important physical growth mechanism. e.g. the accretion of gas onto black holes.," Despite the interesting insights we have neglected another important physical growth mechanism, e.g. the accretion of gas onto black holes."
+ The importance of an additional &rowth mechanism is also confirmed. by the comparison of the black hole mass function resulting from merging only with observations at 2=0: (2)) (227223). (11). (77)). 773). ? (7272))). (273). ?.. 2..," The importance of an additional growth mechanism is also confirmed by the comparison of the black hole mass function resulting from merging only with observations at $z=0$ \citealp{Soltan}) \citealp{Springel05, DiMatteo05, Johansson09, Hopkins08a}) \citealp{Johansson09}) \citealp{Khochfar06, Khochfar09}) \citealp{Hasinger06, Ueda05}) \citet{Hasinger06} \citealp{Khochfar08, Dekel08, Naab09_2, Birnboim07}) \citealp{Dekel06, Khochfar09}) \citet{Johansson09_2}, \citet{Johansson09_2}."
+radio light curves would be poorer and a value of 2 close to 6 (indeed. close to the limiting value 2=5 to keep the expansion self-similar) be too low.,"radio light curves would be poorer and a value of $n$ close to 6 (indeed, close to the limiting value $n = 5$ to keep the expansion self-similar) be too low."
+ Baron et al. (1995)).," Baron et al. \cite{Baron1995}) ),"
+ for instance. fitted ejecta density profiles with -10 for 11993J. based on non-local thermodynamic equilibrium (NLTE) synthetic spectra.," for instance, fitted ejecta density profiles with $n \sim 10$ for 1993J, based on non-local thermodynamic equilibrium (NLTE) synthetic spectra."
+ The wide fractional shell reported in Mareaide et al. (2009a)), The wide fractional shell reported in Marcaide et al. \cite{Marcaide2009}) )
+ and Paper I therefore remains., and Paper I therefore remains.
+. In any case. we notice that it is difficult to discern whether 5 takes the value of 2 or smaller. since there is a tight coupling between the different variables that affect the evolution of the radiation produced by the circumstellar interaction of the expanding supernova shock.," In any case, we notice that it is difficult to discern whether $s$ takes the value of 2 or smaller, since there is a tight coupling between the different variables that affect the evolution of the radiation produced by the circumstellar interaction of the expanding supernova shock."
+ Our software satisfactorily fits the radio data assuming s=2. but it also assumes that the evolution of the magnetic-field energy density (as well as the acceleration efficiency for the electrons) is proportional to the specific kinetic energy of the shock.," Our software satisfactorily fits the radio data assuming $s = 2$, but it also assumes that the evolution of the magnetic-field energy density (as well as the acceleration efficiency for the electrons) is proportional to the specific kinetic energy of the shock."
+ A change in any of these (or other) assumptions of the model may affect our conclusions about the radial density profile of the CSM., A change in any of these (or other) assumptions of the model may affect our conclusions about the radial density profile of the CSM.
+ We note however that the conclusions extracted in. the following subsections are independent of the real value of 5., We note however that the conclusions extracted in the following subsections are independent of the real value of $s$.
+ Weiler et al. (2007)), Weiler et al. \cite{Weiler2007}) )
+ found an enhancement of the flux-density decay rate after day ~3100., found an enhancement of the flux-density decay rate after day $\sim$ 3100.
+ These authors fitted the enhanced decay rate using an exponential factor with an e-folding time of «1100 days. but noted that the same data can also be fitted (although worse) using a power-law decay withB= —2.7.," These authors fitted the enhanced decay rate using an exponential factor with an e-folding time of $\sim$ 1100 days, but noted that the same data can also be fitted (although worse) using a power-law decay with $\beta =$ $-$ 2.7."
+ This decay takes place at all frequencies at the same time. leaving the spectral index of the supernova unaltered.," This decay takes place at all frequencies at the same time, leaving the spectral index of the supernova unaltered."
+ These authors interpreted the exponential-like decay as being produced by a steeper. and fine-tuned. radial density profile of the CSM. beginning on day ~3100 after the explosion.," These authors interpreted the exponential-like decay as being produced by a steeper, and fine-tuned, radial density profile of the CSM, beginning on day $\sim$ 3100 after the explosion."
+ In this section. we show that even an extreme drop in the density of the CSM cannot explain the observed fall-off in the flux density at all frequencies.," In this section, we show that even an extreme drop in the density of the CSM cannot explain the observed fall-off in the flux density at all frequencies."
+ If all the electrons accelerated during the supernova expansion (and not only those that have just been accelerated after being affected by the shock. as it can be derived from Sect.," If all the electrons accelerated during the supernova expansion (and not only those that have just been accelerated after being affected by the shock, as it can be derived from Sect."
+ 5.3 of Weiler et al. 2007)), 5.3 of Weiler et al. \cite{Weiler2007}) )
+ contribute to the radio emission. a change in the CSM density profile is not enough to explain such a rapid decay of the flux density.," contribute to the radio emission, a change in the CSM density profile is not enough to explain such a rapid decay of the flux density."
+ Cooling effects. and eventually the escape of electrons from the emitting region. can have important effects on the evolution of the supernova flux density at these late epochs. which we now consider in more detail.," Cooling effects, and eventually the escape of electrons from the emitting region, can have important effects on the evolution of the supernova flux density at these late epochs, which we now consider in more detail."
+ We can obtain an upper bound to the flux-density decay rate after day 3100 by assuming that the CSM density profile approaches zero (or negligible values) from day 3100 onwards., We can obtain an upper bound to the flux-density decay rate after day 3100 by assuming that the CSM density profile approaches zero (or negligible values) from day 3100 onwards.
+ In that case. the evolution of the expanding structure will no longer be self-similar.," In that case, the evolution of the expanding structure will no longer be self-similar."
+ However. we can still derive. the evolution of the expanding structure from the velocity fields given in Chevalier (19822)): the radial velocity of the shocked gas between the contact discontinuity and the forward shock is approximately equal to the expansion velocity of the contact discontinuity.," However, we can still derive the evolution of the expanding structure from the velocity fields given in Chevalier \cite{Chevalier1982a}) ): the radial velocity of the shocked gas between the contact discontinuity and the forward shock is approximately equal to the expansion velocity of the contact discontinuity."
+ If the supernova further expands into a negligibly dense CSM. the absolute (r.e.. not fractional!)," If the supernova further expands into a negligibly dense CSM, the absolute (i.e., not fractional!)"
+ width of the shocked CSM region (i.e.. from the contact discontinuity to the forward shock) will remain constant. and the expansion velocity of the forward shock will be roughly equal to the expansion velocity of the contact discontinuity.," width of the shocked CSM region (i.e., from the contact discontinuity to the forward shock) will remain constant, and the expansion velocity of the forward shock will be roughly equal to the expansion velocity of the contact discontinuity."
+ Using these assumptions. we can compute an upper bound to the flux-density decay corresponding only to electron cooling (Le.. turning off the electron-escape term).," Using these assumptions, we can compute an upper bound to the flux-density decay corresponding only to electron cooling (i.e., turning off the electron-escape term)."
+ Thus. from day 3100 onwards. we ensure that the total number of relativistic electrons inside the shell remains constant.," Thus, from day 3100 onwards, we ensure that the total number of relativistic electrons inside the shell remains constant."
+ We also assume the ratio of the particle energy density to the magnetic-field energy density to be constant (as in the Chevalier model)., We also assume the ratio of the particle energy density to the magnetic-field energy density to be constant (as in the Chevalier model).
+ In Fig. 4((, In Fig. \ref{OnlyCooling}( (
+a). we show the resulting simulated flux densities at late epochs superimposed on the observations reported in Weiler et al. (2007)).,"a), we show the resulting simulated flux densities at late epochs superimposed on the observations reported in Weiler et al. \cite{Weiler2007}) )."
+ We note that the flux-density decay rate predicted by RAMSES describes remarkably well the enhanced flux-density decay rate reported in Weiler et al. (2007)).," We note that the flux-density decay rate predicted by RAMSES describes remarkably well the enhanced flux-density decay rate reported in Weiler et al. \cite{Weiler2007}) ),"
+ although with systematic slightly higher (around 1056)) flux densities., although with systematic slightly higher (around ) flux densities.
+ These systematics would be lower if we were to shift the initial epoch of negligible CSM to earlier dates by 50-100 days., These systematics would be lower if we were to shift the initial epoch of negligible CSM to earlier dates by $-$ 100 days.
+ In other words.curves. provided the CSM density becomes negligible from day 3000-3100 onwards.," In other words, provided the CSM density becomes negligible from day $-$ 3100 onwards."
+ A steep boundary in the CSM may be unrealistic. but it is conceivable that a rapid drop in the density is produced by the peculiarities in the onset of the stellar wind.," A steep boundary in the CSM may be unrealistic, but it is conceivable that a rapid drop in the density is produced by the peculiarities in the onset of the stellar wind."
+ For a smoother density drop. cooling effects cannot describe. alone. the observed flux-density decay rate.," For a smoother density drop, cooling effects cannot describe, alone, the observed flux-density decay rate."
+ In that case. other contributions to the flux-density decay rate must be invoked.," In that case, other contributions to the flux-density decay rate must be invoked."
+ The escape of the electrons from the radiating region is a natural way to enhance the flux-density decay rate., The escape of the electrons from the radiating region is a natural way to enhance the flux-density decay rate.
+ For instance. if the CSM structure index. s. were changed from 2 to. say. 4 on day 3100 (ever disregarding the self-similar. standard interaction model scenario for radio supernovae) a mean lifetime of the electrons of about 1500 days (together with electron-cooling effects) could also explain the observations.," For instance, if the CSM structure index, $s$, were changed from 2 to, say, 4 on day 3100 (even disregarding the self-similar, standard interaction model scenario for radio supernovae), a mean lifetime of the electrons of about 1500 days (together with electron-cooling effects) could also explain the observations."
+ We note that the mean lifetime of the electrons affects the electron evolution during the entire expansion. not only after day 3100.," We note that the mean lifetime of the electrons affects the electron evolution during the entire expansion, not only after day 3100."
+ In the case of a negligible CSM after day 3100. adding electron escape with a relatively long mean lifetime (around 2500 days) decreases the flux densities predicted by the model without increasing so much the enhanced decay rate at late epochs. resulting in a better fit to the data (see Fig. 4((," In the case of a negligible CSM after day 3100, adding electron escape with a relatively long mean lifetime (around 2500 days) decreases the flux densities predicted by the model without increasing so much the enhanced decay rate at late epochs, resulting in a better fit to the data (see Fig. \ref{OnlyCooling}( ("
+b)).,b)).
+ This actually corresponds to our final choice for the simultaneous modelling of the radio light curves and the expansion curve of 11993J (see Sect., This actually corresponds to our final choice for the simultaneous modelling of the radio light curves and the expansion curve of 1993J (see Sect.
+ 3 and Fig. 1))., \ref{III} and Fig. \ref{SIMULFIT}) ).
+ It is difficult to identify a model with the right combination of both factors (electron escape and enhanced radial density profile of the CSM). since both quantities are completely coupled.," It is difficult to identify a model with the right combination of both factors (electron escape and enhanced radial density profile of the CSM), since both quantities are completely coupled."
+ Therefore. the only clear conclusions we can reach at," Therefore, the only clear conclusions we can reach at"
+"may be reached sequentially, starting from the edges and ending at the midplane.","may be reached sequentially, starting from the edges and ending at the midplane."
+ Our standard procedure does not allow for this kind of gradual evolution because the 1D code settles all dust particles at their local terminal speeds regardless of their location., Our standard procedure does not allow for this kind of gradual evolution because the 1D code settles all dust particles at their local terminal speeds regardless of their location.
+ In this sense our standard procedure is too blunt because it does not allow for slower settling at the stirred-up edges and faster settling at the quiescent midplane., In this sense our standard procedure is too blunt because it does not allow for slower settling at the stirred-up edges and faster settling at the quiescent midplane.
+" True marginally stable states should have midplane dust-to-gas ratios even higher than the maximum ones displayed in Figures 2,, 5,, 6,, and 9.."," True marginally stable states should have midplane dust-to-gas ratios even higher than the maximum ones displayed in Figures \ref{fig:solarmoney}, , \ref{fig:solar19101930}, \ref{fig:mrmoney}, and \ref{fig:mr2122}."
+"4 'To remedy this shortcoming, we modify our procedure by applying a weighting function 0«W(z)<1 to each dust particle's settling velocity."," To remedy this shortcoming, we modify our procedure by applying a weighting function $0 < W(z) \le 1$ to each dust particle's settling velocity."
+" Starting with a KH-stable state near the end of our standard sequence of iterations, we continue to let dust particles drift to the midplane in the 1D code, not at vye1(z) but at the weighted velocity W(z)vre(z)."," Starting with a KH-stable state near the end of our standard sequence of iterations, we continue to let dust particles drift to the midplane in the 1D code, not at $\vrel(z)$ but at the weighted velocity $W (z)
+\vrel(z)$."
+" We use either a Fermi function described by two parameters zso and Zw, or a Gaussian described by a single parameter zy."," We use either a Fermi function described by two parameters $\zhalf$ and $\zw$ , or a Gaussian described by a single parameter $\zw$."
+ The choice of weighting function is somewhat arbitrary; it depends on the shape of the dust profile to be settled and is made case-by-case according to considerations outlined below., The choice of weighting function is somewhat arbitrary; it depends on the shape of the dust profile to be settled and is made case-by-case according to considerations outlined below.
+ The intent of the weighting function is to slow the steepening of density gradients at the dust layer's edges and thereby prevent the edges from destabilizing the entire layer., The intent of the weighting function is to slow the steepening of density gradients at the dust layer's edges and thereby prevent the edges from destabilizing the entire layer.
+ We start with the KH-stable profile in iteration #118 of our solar metallicity sequence (black solid curve in top left panel of Figure 10))., We start with the KH-stable profile in iteration 18 of our solar metallicity sequence (black solid curve in top left panel of Figure \ref{fig:solarwin}) ).
+" The dust layer is characterized by a “core” from z=0 to z©0.52max over which (1) is fairly constant, and an “edge” from z&0.52max to Z=Ζπιακ over which the dust content drops to zero."," The dust layer is characterized by a “core” from $z = 0$ to $z \approx 0.5 \zmax$ over which $\langle \mu \rangle$ is fairly constant, and an “edge” from $z \approx 0.5 \zmax$ to $z =
+\zmax$ over which the dust content drops to zero."
+" Because the instabilities that threaten to disrupt the layer originate in the edge and not the core, we seek a weighting function W(z) that slows the downward drift of dust in the edge but not in the core."," Because the instabilities that threaten to disrupt the layer originate in the edge and not the core, we seek a weighting function $W(z)$ that slows the downward drift of dust in the edge but not in the core."
+ At the same time W(z) should not be so strongly weighted toward the midplane that the core disconnects from the edge and opens a local gap in dust content., At the same time $W(z)$ should not be so strongly weighted toward the midplane that the core disconnects from the edge and opens a local gap in dust content.
+ We find upon experimentation that a Gaussian does not have enough flexibility to meet these requirements for this particular iteration., We find upon experimentation that a Gaussian does not have enough flexibility to meet these requirements for this particular iteration.
+" However a Fermi function with z59=0.005.H approximately to the boundary of the core—and, —correspondingz,=0.05259 seems to work well (blue dashed curve in top left panel of Figure 10))."," However a Fermi function with $\zhalf =
+0.005\Hg$ —corresponding approximately to the boundary of the core—and $\zw = 0.05\zhalf$ seems to work well (blue dashed curve in top left panel of Figure \ref{fig:solarwin}) )."
+ We use this weighting function to settle the dust profile until its midplane µο increases by to a value of 2.9 (red dot-dashed curve)., We use this weighting function to settle the dust profile until its midplane $\mu_0$ increases by to a value of 2.9 (red dot-dashed curve).
+ This settled layer remains KH stable (top middle and right panels of Figure 10))—unlike the layer settled without the weighting function (bottom panels of Figure In the 5))., This settled layer remains KH stable (top middle and right panels of Figure \ref{fig:solarwin}) )—unlike the layer settled without the weighting function (bottom panels of Figure \ref{fig:solar19101930}) ).
+"new profile settled with our modified procedure, the distinction between the core and the edge is no longer so sharp."," In the new profile settled with our modified procedure, the distinction between the core and the edge is no longer so sharp."
+" Thus to settle this new profile even further, a simple Gaussian for the weighting function seems to suffice."," Thus to settle this new profile even further, a simple Gaussian for the weighting function seems to suffice."
+" Choosing z,,=zz0.252max, we attempt to increase the midplane 0.0025jig yetH, again by30%,, but find the resultant profile to be KH unstable (bottom panels of Figure 10))."," Choosing $\zw = 0.0025 \Hg
+\approx 0.25 \zmax$, we attempt to increase the midplane $\mu_0$ yet again by, but find the resultant profile to be KH unstable (bottom panels of Figure \ref{fig:solarwin}) )."
+ Similar results are obtained for the metal-rich case (Figure 11))., Similar results are obtained for the metal-rich case (Figure \ref{fig:mrwin}) ).
+" Using Gaussian weighting functions we are able to push the midplane dust-to-gas ratio uo to a new record of 26.4, which is greater than the value attained using our unweighted standard procedure."," Using Gaussian weighting functions we are able to push the midplane dust-to-gas ratio $\mu_0$ to a new record of 26.4, which is greater than the value attained using our unweighted standard procedure."
+" 'To form rocky planets and gas giant cores, dust must amass in a circumstellar disk."," To form rocky planets and gas giant cores, dust must amass in a circumstellar disk."
+" In the classic scenario for forming planetesimals, dust settles vertically toward the midplane into an ever thinner and denser layer that eventually becomes gravitationally unstable (Safronov1969;Goldreich&Ward 1973)."," In the classic scenario for forming planetesimals, dust settles vertically toward the midplane into an ever thinner and denser layer that eventually becomes gravitationally unstable \citep{safronov69, goldreichward73}."
+". Toomre's criterion for gravitational instability (GI) is satisfied for midplane dust-to-gas ratios (pa/pg)o=Ho2HOo,Toomre, Where [40,Toomre©34 for a minimum-mass nebula orbiting a solar-mass star (equation 10; note that μρ,Τοοπιτο 1$ nearly independent of disk radius)."," Toomre's criterion for gravitational instability (GI) is satisfied for midplane dust-to-gas ratios $(\rhod/\rhog)_0 \equiv \mu_0 \gtrsim
+\mu_{\rm 0,Toomre}$, where $\mu_{\rm 0,Toomre} \approx 34$ for a minimum-mass nebula orbiting a solar-mass star (equation \ref{eqn:muToomre}; note that $\mu_{\rm 0,Toomre}$ is nearly independent of disk radius)."
+" For comparison, in a disk of well-mixed dust and gas at solar abundance, µο70.015 (Lodders"," For comparison, in a disk of well-mixed dust and gas at solar abundance, $\mu_0 \approx 0.015$ \citep{lodders03}."
+ Whether dust can accumulate until its density 2003).increases by more than three orders of magnitude depends on how turbulent the ambient gas is., Whether dust can accumulate until its density increases by more than three orders of magnitude depends on how turbulent the ambient gas is.
+" Evensupposing that gas in certain regions of the disk is not intrinsically turbulent because it istoo weakly ionized to support the magnetorotational(e.g., instability), the dust itself may excite turbulence in gas by a"," Evensupposing that gas in certain regions of the disk is not intrinsically turbulent (e.g., because it istoo weakly ionized to support the magnetorotational instability), the dust itself may excite turbulence in gas by a"
+The radial velocities were measured by Gaussian fits to he whole profile. using the SPLOT proerauuue of IRAE. This averages over anv real or instrumental assviuuetry.,"The radial velocities were measured by Gaussian fits to the whole profile, using the SPLOT programme of IRAF, This averages over any real or instrumental assymmetry."
+ Saturated mes were measured by fitting to the wines., Saturated lines were measured by fitting to the wings.
+ We rave checked that the radial velocities determuned in this way closely agree with the line centroid radial velocities., We have checked that the radial velocities determined in this way closely agree with the line centroid radial velocities.
+ It is because of our method that we were able to use emissiou ines. whose centres are saturated.," It is because of our method that we were able to use emission lines, whose centres are saturated."
+ The nature of our IUE spectra does not enable us to do auvthiug more precise., The nature of our IUE spectra does not enable us to do anything more precise.
+" Let us note however that the IIST/STIS spectra of Sugeest sole assvinuctry of the resonance enissiou line profiles, with line ceutres beime shifted by apparent absorption on the low waveleneth side (Young 2006. private communication as well as the MAST archive)."," Let us note however that the HST/STIS spectra of suggest some assymmetry of the resonance emission line profiles, with line centres being shifted by apparent absorption on the low wavelength side (Young 2006, private communication as well as the MAST archive)."
+ The errors. iudicated in the radial veocitv tables. correspond to Loa errors of the respective mean values.," The errors, indicated in the radial veocity tables, correspond to 1$\sigma$ errors of the respective mean values."
+ Thev clo not indicate systematic errors. such as for lustaunce those of errors in the absolute wavelcneth scale.," They do not indicate systematic errors, such as for instance those of errors in the absolute wavelength scale."
+ Πονονο the mean values are based on data of caiffereut rausitions. so their σ values mcelude that due to uncertzntv of laboratory wawveleugths.," However the mean values are based on data of different transitions, so their $\sigma$ values include that due to uncertainty of laboratory wavelengths."
+ The other systematic errors due to the absolute wavelength scale. have been estimated bv comparing our radial velocity estimates from differcut spectra of tle same target. taken at the same epoch.," The other systematic errors due to the absolute wavelength scale, have been estimated by comparing our radial velocity estimates from different spectra of the same target, taken at the same epoch."
+ In particular we found 11 pairs of IUE spectra (1 forAud.. 10 for Dra)) taken on the sale day or one day ater.," In particular we found 14 pairs of IUE spectra (4 for, 10 for ) taken on the same day or one day later."
+ The offset in radial velocity for these pairs is jetween Av~07μις1. with an offset or LO pairs of 2-8 . a ποσα value of around dnas! and a lean of ~6.7Eολων5. ," The offset in radial velocity for these pairs is between $\Delta v \sim 0-20 km\,s^{-1}$, with an offset for 10 pairs of 2-8 $^{-1}$, a median value of around 5 km $^{-1}$ and a mean of $\sim 6.7 \pm 1.3 km
+s^{-1}$."
+There is no cdiffereuce in the offset between saturated aud nou-saturated lines. when we compare a spectrum on which a particular line is saturated and another spectrum on which it is not saturated.," There is no difference in the offset between saturated and non-saturated lines, when we compare a spectrum on which a particular line is saturated and another spectrum on which it is not saturated."
+ The corresponding offset for 2 HIST/STIS spectra of taken ou the same dav is oulv 1.7 1 , The corresponding offset for 2 HST/STIS spectra of taken on the same day is only 1.7 $^{-1}$ .
+Our measured radial velocity meaus and their 10 errors are tabulated iu table 1 forPer.. andIba. in table 2 for andl in table 3 forDra.," Our measured radial velocity means and their 1 $\sigma$ errors are tabulated in table 1 for, and, in table 2 for and in table 3 for."
+. The phases of the tables are from the ephemeris of Mikolajewska (20033). with respect to inferior conjunction of the elaut.," The phases of the tables are from the ephemeris of ajewska \cite{Mik03}) ), with respect to inferior conjunction of the giant."
+ The last 8 Hues in table 1 for aud the ast 2 lines in table 3 for are based on IST/STIS spectra., The last 8 lines in table 1 for and the last 2 lines in table 3 for are based on HST/STIS spectra.
+ The wavelength of the 1610 lue used. is a mean wavelength of the fine structure. even by Cleee et al (1999)).," The wavelength of the 1640 line used, is a mean wavelength of the fine structure, given by Clegg et al \cite{Cl99}) )."
+ Let us note that this mean is fairly insensitive to the exact plivsical conditions of line ‘oration: ranges of electron density from 10! to Lo? ? and temperatures from 10 000 to 30 00018 eive for case B ine formation a shift of 0.007 oor 1.3 ‘au 1, Let us note that this mean is fairly insensitive to the exact physical conditions of line formation; ranges of electron density from $^4$ to $^9$ $^{-3}$ and temperatures from 10 000 to 30 000K give for case B line formation a shift of 0.007 or 1.3 km $^{-1}$.
+ Some radial velocities iu the tables. includiug meaus alc the radial velocity of the 1610 line. plotted against orbital phase. are shown in fig 1. forMus.. while radial velocity differences with πιο] less scatter are shown iu fies.," Some radial velocities in the tables, including means and the radial velocity of the 1640 line, plotted against orbital phase, are shown in fig 1, for, while radial velocity differences with much less scatter are shown in figs."
+ 2 and 3 for the same svinbiotic., 2 and 3 for the same symbiotic.
+ Silay figures shown. ave {ν 5 and 6 for aud 7. 8 alc 9 forDra.," Similar figures shown, are 4, 5 and 6 for and 7, 8 and 9 for."
+. The difference erapls ave with respect to the resonance doublet ofIv. as measurements of the strone lines belonging to this doublet could be made for all the spectra studied by us.," The difference graphs are with respect to the resonance doublet of, as measurements of the strong lines belonging to this doublet could be made for all the spectra studied by us."
+ Shifts. such as those of the intercombination lines relative to1v. depend ou the shift of itself.," Shifts, such as those of the intercombination lines relative to, depend on the shift of itself."
+ The values for ave not plotted. but hei beliaviour can be seen from the tables.," The values for are not plotted, but their behaviour can be seen from the tables."
+ Tn the graphs mentioned. the open circles are of fits of the line centres to a Gaussian for IUE observations i quiescence. while in the cases of IUE observations of iwl Dra. the filled circles ave for IUE observationsiu outburst.," In the graphs mentioned, the open circles are of fits of the line centres to a Gaussian for IUE observations in quiescence, while in the cases of IUE observations of and , the filled circles are for IUE observationsin outburst."
+ Open and filled triangles denote corresponding, Open and filled triangles denote corresponding
+Over the past years. CMB polarization has been. measured by several experiments and is being measured by the Planck surveyor (22222222222).,"Over the past years, CMB polarization has been measured by several experiments and is being measured by the Planck surveyor \citep{DASI:data,DASI:instrument,DASI:I,DASI:II,DASI:III,DASI:3yr,QUaD1,QUaD2,QUaD:instrument,QUaD_improved, Planck_bluebook}."
+ CMB polarization pattern may be considered as the sum of gradient-like E mode and curl-like B mode (??).," CMB polarization pattern may be considered as the sum of gradient-like E mode and curl-like B mode \citep{Seljak-Zaldarriaga:Polarization,Kamionkowski:Flm}."
+ In the standard model. B mode polarization is not produced by scalar perturbation. but solely by tensor perturbation.," In the standard model, B mode polarization is not produced by scalar perturbation, but solely by tensor perturbation."
+ Therefore. measurement of B mode polarization makes it possible to probe the universe on the energy scale at inflationary period (22222).," Therefore, measurement of B mode polarization makes it possible to probe the universe on the energy scale at inflationary period \citep{Kamionkowski:Flm,Seljak-Zaldarriaga:Polarization,Modern_Cosmology,Inflation,Foundations_Cosmology}."
+ In most inflationary models. tensor-to-scalar ratio r is much smaller than one. and the WMAP 7 year data imposes an upper bound on +<0.36 at 95% confidence level (??)..," In most inflationary models, tensor-to-scalar ratio $r$ is much smaller than one, and the WMAP 7 year data imposes an upper bound on $r<0.36$ at $95\%$ confidence level \citep{WMAP7:powerspectra,WMAP7:Cosmology}."
+ Besides instrumet noise. there are complications. which limits detectability of tensor perturbation.," Besides instrument noise, there are complications, which limits detectability of tensor perturbation."
+ Imperfection 1 removing foregrounc and gravitational lensing impose:PA observational limit on tensor-sealar-ratio: p.~107+ and ré3x107 respectively (??)..," Imperfection in removing foreground and gravitational lensing imposes observational limit on tensor-scalar-ratio: $r\sim10^{-4}$ and $r\sim 3\times10^{-5}$ respectively \citep{B_lensing,B_mode_limit_foreground}."
+ Due to the nature of the observatio or heavy foreground contamination. reliable of estimatio on CMB polarizatio signal is not available over a whole sky.," Due to the nature of the observation or heavy foreground contamination, reliable of estimation on CMB polarization signal is not available over a whole sky."
+ Incomplete sky coverage leads to E/B mixing. andvery significantly limit our capacity to measure tensor perturbatio as well (?)..," Incomplete sky coverage leads to E/B mixing, andvery significantly limit our capacity to measure tensor perturbation as well \citep{Bunn:EB-Separation}."
+ Therefore. there have been various efforts to understand and reduce E/B mixing (??????)..," Therefore, there have been various efforts to understand and reduce E/B mixing \citep{Kim:optimization,Kim:measuring_a2lm,Bunn:EB-Separation,EB_incomplete_sky,EB_harmonic,Smith:pseudo_EB}."
+ It is best to implement E/B decomposition in map space. since diffuse foregrounds and point sources are well-localized in map space. and their spatial information are known relatively better than other properties.," It is best to implement E/B decomposition in map space, since diffuse foregrounds and point sources are well-localized in map space, and their spatial information are known relatively better than other properties."
+ In this paper. we investigate E/B decomposition in pixel space.," In this paper, we investigate E/B decomposition in pixel space."
+ Our investigation shows that E/B mixing is highly localized in pixel space., Our investigation shows that E/B mixing is highly localized in pixel space.
+ Therefore. we may reduce E/B mixing effectively by excluding the ambiguous pixels.," Therefore, we may reduce E/B mixing effectively by excluding the ambiguous pixels."
+ We have applied our method to simulated maps partially blocked by a foreground (diffuse + point source) mask., We have applied our method to simulated maps partially blocked by a foreground (diffuse + point source) mask.
+ After excluding ambiguous pixels. we find that leakage power in retained pixels (sky fraction ~ 0.48) is smaller than primordial (i.e. unlensed) B mode power spectrum of tensor-to-scalar ratio r~1x107 atwide range of multipoles (50€/ 2000).," After excluding ambiguous pixels, we find that leakage power in retained pixels (sky fraction $\sim 0.48$ ) is smaller than primordial (i.e. unlensed) B mode power spectrum of tensor-to-scalar ratio $r\sim 1\times10^{-3}$ atwide range of multipoles $50\la l \la 2000$ )."
+ The outline of this paper is as follows., The outline of this paper is as follows.
+ In Sec. 2..," In Sec. \ref{Stokes},"
+ we discuss all-sky analysis of CMB polarization., we discuss all-sky analysis of CMB polarization.
+ In See. 3," In Sec. \ref{pixel_filter},"
+ we derive E/B decomposition in pixel space., we derive E/B decomposition in pixel space.
+ In Sec. 4..," In Sec. \ref{cutsky},"
+ we discuss the application to cut sky. and the method to diagnose ambiguous pixels.," we discuss the application to cut sky, and the method to diagnose ambiguous pixels."
+ In Sec., In Sec.
+ 5 and 6.. we present our simulation result.," \ref{simulation} and \ref{scale}, we present our simulation result."
+ In Section. 7.. we summarize our investigation.," In Section \ref{Discussion}, we summarize our investigation."
+" In Appendix AppendixA:.. we discuss error analysis of pseudo C, estimation. and show interpixel noise correlation may be neglected."," In Appendix \ref{noise}, we discuss error analysis of pseudo $C_l$ estimation, and show interpixel noise correlation may be neglected."
+ The state of polarization is described by Stokes parameter (??)..," The state of polarization is described by Stokes parameter \citep{Kraus:Radio_Astronomy,Tools_Radio_Astronomy}."
+ Since Thompson scattering does not generate circular polarization. Stokes parameter Q and U are sufficient to describe CMB polarization (?)..," Since Thompson scattering does not generate circular polarization, Stokes parameter Q and U are sufficient to describe CMB polarization \citep{Modern_Cosmology}. ."
+" Stokes parameter Q and U transform under rotation of an angle w on the plane perpendicular to direction fi (??):: Therefore. all-sky Stokes parameters may be decomposed into spin £2 spherical harmonics (2). as follows: where the decomposition coefficients dsj, are obtained by: Though the quantity shown in Eq."," Stokes parameter $Q$ and $U$ transform under rotation of an angle $\psi$ on the plane perpendicular to direction $\mathbf {\hat n}$ \citep{Seljak-Zaldarriaga:Polarization,Zaldarriaga:Polarization_Exp}: Therefore, all-sky Stokes parameters may be decomposed into spin $\pm2$ spherical harmonics \citep{Seljak-Zaldarriaga:Polarization} as follows: where the decomposition coefficients $a_{\pm2,lm}$ are obtained by: Though the quantity shown in Eq."
+ 2 has direct association with physical observables (1.8. Stokes parameters). rotational variance leads to computational complication.," \ref{Q_lm+iU_lm} has direct association with physical observables (i.e. Stokes parameters), rotational variance leads to computational complication."
+" Therefore. two real scalar quantities. termed ""E and B! mode. areoften built out of QU)+(LGA) (22): where J and ὅ refer to lowering and raising operator respectively (?).."," Therefore, two real scalar quantities, termed `E' and `B' mode, areoften built out of $Q(\mathbf {\hat n})\pm i U(\mathbf {\hat n})$ \citep{Kamionkowski:Flm,Seljak-Zaldarriaga:Polarization}: : where $\lo$ and $\ra$ refer to lowering and raising operator respectively \citep{Seljak-Zaldarriaga:Polarization}. ."
+ The explicit expression of ὅ and ὅ are given, The explicit expression of $\lo$ and $\ra$ are given
+distant QSOs.,distant QSO's.
+ The values of ?IL/IL inferred from these clouds are consistent with each other and average to ?2IL/IL73-10? (O'Mearaοἱal.2001).., The values of $\hh$ /H inferred from these clouds are consistent with each other and average to $\hh/H\sim 3 \cdot 10^{-5}$ \citep{omeara}.
+ Because the deuterium vield in the big bang is a sensitive Function of the entropy per barvon (or alternatively barvon to photon ratio 7). the observed 21L/IE places important constraints on conditions during the epoch of BBN. and gives a determination of the barvon rest mass contribution to closure (scaled by the Hubble parameter in units of LOO kms.!Mpe.!). OA?zz0.02.," Because the deuterium yield in the big bang is a sensitive function of the entropy per baryon (or alternatively baryon to photon ratio $\eta$ ), the observed $\hh$ /H places important constraints on conditions during the epoch of BBN, and gives a determination of the baryon rest mass contribution to closure (scaled by the Hubble parameter in units of 100 ${\rm km s ^{-1} Mpc^{-1}}$ ), $\Omega_b
+h^2\approx 0.02$."
+ Estimates of the present day interstellar medium values of ?IL/IE in different directions from the UV Lyman à. transition are also consistent with each other and average to about Πα21.5410. (MeCullough.1992).. in accord with the idea that stellar processing results in a net destruction of deuterium.," Estimates of the present day interstellar medium values of ${^2}{\rm
+H/H}$ in different directions from the UV Lyman $\alpha$ transition are also consistent with each other and average to about ${^2}{\rm
+H/H}\approx 1.5\cdot 10^{-5}$ \citep{mcullough}, in accord with the idea that stellar processing results in a net destruction of deuterium."
+ There have also been a number of careful radio searches for the hxperfime (transition in deuterium., There have also been a number of careful radio searches for the hyperfine transition in deuterium.
+ So far these have only vielded upper bounds about a [actor of four higher than the abundance measured by UV satellites., So far these have only yielded upper bounds about a factor of four higher than the abundance measured by UV satellites.
+ Nonethless. radio searches [or pockets of 7LL in supernova remnants may represent a promising avenue [or detecting evidence of relativistic flows.," Nonethless, radio searches for pockets of ${^2}{\rm H}$ in supernova remnants may represent a promising avenue for detecting evidence of relativistic flows."
+ Ultimately whether or not deuterium [rom relativistic flows is detectable depends on the interaction of the flow with the ISM., Ultimately whether or not deuterium from relativistic flows is detectable depends on the interaction of the flow with the ISM.
+ We do not caleulate this here., We do not calculate this here.
+" The evolution of a hot plasma is governed by the initial radius Jy of the plasma. the initial temperature of (he plasma Zi. and (he entropy per barvon s/f, in units of Doltzinann's constant."," The evolution of a hot plasma is governed by the initial radius $R_0$ of the plasma, the initial temperature of the plasma $T_0$, and the entropy per baryon $s/k_b$ in units of Boltzmann's constant."
+ Initially the plasma accelerates under the influence of thermal pressure., Initially the plasma accelerates under the influence of thermal pressure.
+ For optically thick flows the acceleration ceases when the thermal energy of the fireball is converted to kinetic energv., For optically thick flows the acceleration ceases when the thermal energy of the fireball is converted to kinetic energy.
+5. Enerev5. conservation 5gives the final Lorentz factor 5;if of the fireball as llere £i is the total initial energy. (thermal plus rest mass) in the fireball ancl VW is the mass of the barvons confined to the fireball (in this and in what follows we set —¢= 1)., Energy conservation gives the final Lorentz factor $\gamma_f$ of the fireball as Here $E_{0}$ is the total initial energy (thermal plus rest mass) in the fireball and $M$ is the mass of the baryons confined to the fireball (in this and in what follows we set $\hbar=c=1$ ).
+" In this paper we will Locus on the case where the initial plasma is barvon poor. so 5,> 10."," In this paper we will focus on the case where the initial plasma is baryon poor, so $\gamma_f > 10$ ."
+" Non-relativistic or mildly relativistic outflows wilh modest entropy per barvon (s/h,~ 100) can give rise (o interesting heavy element svnthesis (Meverοἱal.1992:WoosleyetTakahashi.Witt.&Janka 1994).. but are relatively uninteresting as far as A<& nuclei are concerned."," Non-relativistic or mildly relativistic outflows with modest entropy per baryon $s/k_b \sim 100$ ) can give rise to interesting heavy element synthesis \citep{meyer92,wilsonandmathews,takahashi}, but are relatively uninteresting as far as $<$ 8 nuclei are concerned."
+ In terms of 5; and the initial temperature. the entropy per barvon in the fireballls," In terms of $\gamma_f$ and the initial temperature, the entropy per baryon in the fireballis"
+factor 10 reduction in black hole mass will only reduce the lower and upper limits on /? by and increase ng by250%.,factor 10 reduction in black hole mass will only reduce the lower and upper limits on $R$ by and increase $n_{\rm H}$ by.
+". We now consider how the physical constraints we can place on qe properties of the eclipsing cloud in NGC 3227 compare with le constraints expected for gas in various regions of the AGN,", We now consider how the physical constraints we can place on the properties of the eclipsing cloud in NGC 3227 compare with the constraints expected for gas in various regions of the AGN.
+ Tost obviously. the location of the cloud (roughly 10-100 light-days from the central X-ray source) 1s consistent with an origin in ye BLR of NGC 3227. the size of which is estimated from optical reverberation studies to be between 10-20 light-days (Wandel et al.," Most obviously, the location of the cloud (roughly 10-100 light-days from the central X-ray source) is consistent with an origin in the BLR of NGC 3227, the size of which is estimated from optical reverberation studies to be between 10-20 light-days (Wandel et al."
+" 1999, Winge et al."," 1999, Winge et al."
+ 1995. Salamanca et al.," 1995, Salamanca et al."
+ 1994). although note iat since these estimates refer to the region of the BLR which emits most of the Balmer lines. the full extent of the BLR may be significantly larger.," 1994), although note that since these estimates refer to the region of the BLR which emits most of the Balmer lines, the full extent of the BLR may be significantly larger."
+ Similarly. the density we estimate for the cloud is also consistent with the lower limit on densities of BLR clouds estimated from optical line ratios (e.g. see Krolik 1999 or discussion) and rules out an origin in the lower density. more distant. narrow line region (NLR) or large-scale molecular torus envisioned by the standard unification model.," Similarly, the density we estimate for the cloud is also consistent with the lower limit on densities of BLR clouds estimated from optical line ratios (e.g. see Krolik 1999 for discussion) and rules out an origin in the lower density, more distant, narrow line region (NLR) or large-scale molecular torus envisioned by the standard unification model."
+ This last constraint is jxarticularly strong. since even in the absence of any lower bound on ionisation parameter. the density ofthe cloud can only increase with assumed orbital radius. due to the well constrained total column density and the strong dynamical constraint on cloud size imposed by the cloud crossing time.," This last constraint is particularly strong, since even in the absence of any lower bound on ionisation parameter, the density of the cloud can only increase with assumed orbital radius, due to the well constrained total column density and the strong dynamical constraint on cloud size imposed by the cloud crossing time."
+ By the same token. the cloud cannot be very dense. as according to our equation 2. the density scales only with the square root of distance from the central X-ray source. so if the assumed density were significantly larger than 107 ? the cloud would be located much further from the X-ray source than is consistent with any known component of the AGN (the same reasoning can be used to strongly constrain eloud size).," By the same token, the cloud cannot be very dense, as according to our equation 2, the density scales only with the square root of distance from the central X-ray source, so if the assumed density were significantly larger than $10^{8}$ $^{-3}$ the cloud would be located much further from the X-ray source than is consistent with any known component of the AGN (the same reasoning can be used to strongly constrain cloud size)."
+ Our strong constraint on cloud density is in conflict with the estimate of much denser BLR clouds (ray~10‘tom 7) in the luminous Seyfert | NGC 5548. by Ferland et al. (," Our strong constraint on cloud density is in conflict with the estimate of much denser BLR clouds $n_{\rm
+H}\sim10^{11}$ $^{-3}$ ) in the luminous Seyfert 1 NGC 5548, by Ferland et al. ("
+1992). who required a large hydrogen number density to prevent the clouds becoming too heavily ionised in the ugh radiation density environment implied by the relatively small BLR size estimated from reverberation mapping.,"1992), who required a large hydrogen number density to prevent the clouds becoming too heavily ionised in the high radiation density environment implied by the relatively small BLR size estimated from reverberation mapping."
+ However. we note qat although NGC 3227 has a similar sized BLR to NGC 5548 (e.g. see Peterson et al.," However, we note that although NGC 3227 has a similar sized BLR to NGC 5548 (e.g. see Peterson et al."
+ 2002). it has a much lower X-ray luminosity (1077 erg + versus 510% erg lj. so that the radiation density in the BLR of NGC 3227 is likely to be much lower than in NGC 5548.," 2002), it has a much lower X-ray luminosity $10^{42}$ erg $^{-1}$ versus $5\times10^{43}$ erg $^{-1}$ ), so that the radiation density in the BLR of NGC 3227 is likely to be much lower than in NGC 5548."
+ In fact. if for comparison. we convert from our ionisation parameter £ to the equivalent formulation ( used by Ferland et al. (," In fact, if for comparison, we convert from our ionisation parameter $\xi$ to the equivalent formulation $U$ used by Ferland et al. ("
+1992) (e.g. see George et al.,1992) (e.g. see George et al.
+ [998a for conversion factors). we find (/«0.1 in NGC 3227. consistent with the value estimated for NGC 5548 and the BLR in other AGN (see Ferland et al.," 1998a for conversion factors), we find $U<0.1$ in NGC 3227, consistent with the value estimated for NGC 5548 and the BLR in other AGN (see Ferland et al."
+" 1992, Peterson 1997. and references therein)."," 1992, Peterson 1997, and references therein)."
+ The size inferred for the eclipsing cloud is quite large. implying that only of order 100 clouds are required to produce he —10€ covering fraction estimated for the BLR (Peterson 1997. Krolik 1999).," The size inferred for the eclipsing cloud is quite large, implying that only of order $\sim100$ clouds are required to produce the $\sim10$ covering fraction estimated for the BLR (Peterson 1997, Krolik 1999)."
+ This small number of clouds conflicts with the arge number (=107) required to produce the smooth broad line ofiles observed in type | AGN (e.g. Atwood. Baldwin Carswell 1982). assuming that the cloud intrinsic line widths are small. hermal widths only.," This small number of clouds conflicts with the large number $>10^{4}$ ) required to produce the smooth broad line profiles observed in type 1 AGN (e.g. Atwood, Baldwin Carswell 1982), assuming that the cloud intrinsic line widths are small, thermal widths only."
+ This latter assumption might be relaxed if he cloud line widths are intrinsically broadened. e.g. if the clouds are dynamically unstable. which might be expected given the large cloud size.," This latter assumption might be relaxed if the cloud line widths are intrinsically broadened, e.g. if the clouds are dynamically unstable, which might be expected given the large cloud size."
+ Note also that our cloud size estimates can be used © constrain the size of the primary X-ray emitting region to be ess than ~1 light-day. or around 170 Sehwartzsehild radii for ü5-I0 M. black hole (assuming the primary X-ray emitting region is totally covered. as suggested by the lack of variability in the unabsorbed spectral component).," Note also that our cloud size estimates can be used to constrain the size of the primary X-ray emitting region to be less than $\sim1$ light-day, or around 170 Schwartzschild radii for a $5\cdot
+10^{7}$ $_{\odot}$ black hole (assuming the primary X-ray emitting region is totally covered, as suggested by the lack of variability in the unabsorbed spectral component)."
+ When considering the implications of these results for models of the BLR in AGN in general. one must bear in mind the caveat that strong observational selection effects almost certainly exist in any detection of BLR clouds by X-ray absorption.," When considering the implications of these results for models of the BLR in AGN in general, one must bear in mind the caveat that strong observational selection effects almost certainly exist in any detection of BLR clouds by X-ray absorption."
+ For example. eclipses of the X-ray emitting region by much smaller clouds. of the higher densities envisaged by Ferland et al. (," For example, eclipses of the X-ray emitting region by much smaller clouds, of the higher densities envisaged by Ferland et al. ("
+1992). would only be noticable if the X-ray emitting region was itself very compact (smaller than hundreds of light-seconds or less than one Schwartzschild radius for a 5-10* M. black hole).,"1992), would only be noticable if the X-ray emitting region was itself very compact (smaller than hundreds of light-seconds or less than one Schwartzschild radius for a $5\cdot10^{7}$ $_{\odot}$ black hole)."
+ Furthermore. the duration of such events would be much shorter than the observed ~100 d (perhaps only a few thousand seconds). so would only be detectable with continuous monitoring. which we do not have.," Furthermore, the duration of such events would be much shorter than the observed $\sim100$ d (perhaps only a few thousand seconds), so would only be detectable with continuous monitoring, which we do not have."
+ Therefore we cannot rule out the existence of such small clouds in NGC 3227., Therefore we cannot rule out the existence of such small clouds in NGC 3227.
+ Other selection effects may also be at work., Other selection effects may also be at work.
+ For example. the relatively low radiation density in the BLR of NGC 3227 may permit the existence of large. relatively low-density clouds which might not exist in other AG," For example, the relatively low radiation density in the BLR of NGC 3227 may permit the existence of large, relatively low-density clouds which might not exist in other AGN."
+ The presence of such clouds may also account for the smaller amplitude column-density variations reported in NGC 3227 by George et al. (, The presence of such clouds may also account for the smaller amplitude column-density variations reported in NGC 3227 by George et al. (
+1998b).,1998b).
+ Alternatively. the fact that column variations are apparently common in NGC 3227 might be explained if the BLR has a planar geometry and we are viewing the X-ray source through the BLR plane.," Alternatively, the fact that column variations are apparently common in NGC 3227 might be explained if the BLR has a planar geometry and we are viewing the X-ray source through the BLR plane."
+ Even taking into account the above caveuts. our results imply that a population of large BLR clouds exists in at least some AGN.," Even taking into account the above caveats, our results imply that a population of large BLR clouds exists in at least some AGN."
+ Taking the estimate of much smaller. higher density clouds in NGC 5548 (Ferland et al.," Taking the estimate of much smaller, higher density clouds in NGC 5548 (Ferland et al."
+ 1992) at face value. it," 1992) at face value, it"
+Ongoing surveys of the CDES with the Spitzer. infra-reck telescope will when complete. provide accurate. dust-independent measurements of the SERs in our redshift sample.,"Ongoing surveys of the CDFS with the Spitzer infra-red telescope will, when complete, provide accurate, dust-independent measurements of the SFRs in our redshift sample."
+ The ongoing SILVDISS survey of bright. ο>Smad. sub-pim galaxies will provide a measurement of the clustering of these probable progenitors of the ERGs (Dunlop 2005: Mortier et al.," The ongoing SHADES survey of bright, $F(850\rm \mu m)>8 mJy$, sub-mm galaxies will provide a measurement of the clustering of these probable progenitors of the ERGs (Dunlop 2005; Mortier et al."
+ 2005)., 2005).
+ We have planned further ancl more extensive spectroscopic surveys. including in the near-Lt. of high redshift galaxy clusters and large-scale structure.," We have planned further and more extensive spectroscopic surveys, including in the near-IR, of high redshift galaxy clusters and large-scale structure."
+" This paper is based on observations obtained at the Ciemini Observatory, which is operated. by the Association of Universities for Research in. Astronomy on behalf of the Gemini partnership."," This paper is based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy on behalf of the Gemini partnership."
+ Εμίν paper also makes use of publically available GOODS data. from the Advanced Camera for Surveys on the NASA Llubble Space Telescope and from the Infrared Spectrometer and Array Camera on the Very. Large Telescope operated by the European Southern Observatory., This paper also makes use of publically available GOODS data from the Advanced Camera for Surveys on the NASA Hubble Space Telescope and from the Infrared Spectrometer and Array Camera on the Very Large Telescope operated by the European Southern Observatory.
+ We thank the GOODS team for providing reduced. data products., We thank the GOODS team for providing reduced data products.
+ JSD and NDIt acknowledge. PPARC funding., JSD and NDR acknowledge PPARC funding.
+ NDR. acknowledges funding from the Universidad: Nacional Autonoma cde Mexico., NDR acknowledges funding from the Universidad Nacional Autonoma de Mexico.
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+selection elfects include Galactic extinction. ancl crowded star-lields.,selection effects include Galactic extinction and crowded star-fields.
+ “Phere is also the problem of host galaxy extinction. although this is expected to be a less of problem given that at least some of the afterglows are aobserved oll-set from. the hosts.," There is also the problem of host galaxy extinction, although this is expected to be a less of a problem given that at least some of the afterglows are observed off-set from the hosts."
+ From the current. optical follow-up attempts of triggered. bursts. it is clear that a significant fraction of SGRBs that are apparently on-axis have not been observed in the optical at all.," From the current optical follow-up attempts of triggered bursts, it is clear that a significant fraction of SGRBs that are apparently on-axis have not been observed in the optical at all."
+ Εις can partly be attributed to the much greater. distances of theSwift bursts. compared to a GW triggered search.," This can partly be attributed to the much greater distances of the bursts, compared to a GW triggered search."
+ We have based our calculations on a sensitivity distance of ~200 Alpe (non-coincidence rate). but note that GAY wave sources with smaller angles of inclination will be detected. at greater distances.," We have based our calculations on a sensitivity distance of $\sim 200$ Mpc (non-coincidence rate), but note that GW wave sources with smaller angles of inclination will be detected at greater distances."
+ Preliminary studies suggest that this bias can increase the median distance to ~269 MMpe lor binaries with inclination angles less than25., Preliminary studies suggest that this bias can increase the median distance to $\sim 269$ Mpc for binaries with inclination angles less than.
+. A numerical study is xanned to investigate how this bias will ellect the rate of coincident detections., A numerical study is planned to investigate how this bias will effect the rate of coincident detections.
+ Another important issue for the joint searches is the arge errors in the GW source localisation. which can extenc o some tens of degrees.," Another important issue for the joint searches is the large errors in the GW source localisation, which can extend to some tens of degrees."
+ This is à significant problem eiven he first observed. GRB optical afterglows required: smial error boxes of arc-minute size., This is a significant problem given the first observed GRB optical afterglows required small error boxes of arc-minute size.
+ To counter this. the CAV rigecred search strategv will use the estimated. horizon distance of the detector network to reduce the number of »»tential host galaxies. as opposed. to a ‘blind’ error. box hat extends to cosmological distances.," To counter this, the GW triggered search strategy will use the estimated horizon distance of the detector network to reduce the number of potential host galaxies, as opposed to a `blind' error box that extends to cosmological distances."
+ Εις technique was demonstrated in the recent LIGO/Virgo 86: the telescope rolling strategv usec a galaxy weighting taking into account the mass ancl the distance of cataloguccl galaxics., This technique was demonstrated in the recent LIGO/Virgo S6; the telescope pointing strategy used a galaxy weighting taking into account the mass and the distance of catalogued galaxies.
+ A detailed: description ancl analysis of the S6 coincidence searches will be published as a LIGO Scientific Collaboration paper., A detailed description and analysis of the S6 coincidence searches will be published as a LIGO Scientific Collaboration paper.
+ Another problem that manifests with large coincidence error boxes is the increasing chance of detecting false coincident optical transients., Another problem that manifests with large coincidence error boxes is the increasing chance of detecting false coincident optical transients.
+ False coincident sources may include supernovac. [are stars. variable active galactic nuclei and even Earth orbiting space debris.," False coincident sources may include supernovae, flare stars, variable active galactic nuclei and even Earth orbiting space debris."
+ Fortunately. some of these sources can be excluded. in the analysis because of the sensitivity distance of GW searches and the expectation that the strongest GAY sources. will be associated with catalogued host galaxies.," Fortunately, some of these sources can be excluded in the analysis because of the sensitivity distance of GW searches and the expectation that the strongest GW sources will be associated with catalogued host galaxies."
+ Another possibility [or optimizing the coincidence search is to data mine archived images from very wide field optical surveys. such as," Another possibility for optimizing the coincidence search is to data mine archived images from very wide field optical surveys, such as"
+The presence of a logarithm in the Graaisformation between fux ancl magnitude produces an asvnmmeltric distribution of points around the mean color. mainly in the region of faint objects.,"The presence of a logarithm in the transformation between flux and magnitude produces an asymmetric distribution of points around the mean color, mainly in the region of faint objects."
+" We obtain (hiis behavior by defining (wo curves with n, the level over the mean.", We obtain this behavior by defining two curves with $n_{\sigma}$ the level over the mean.
+ As can be seen from Eq. 46..," As can be seen from Eq. \ref{eq:curve},"
+ for a given narrow-band magnitude. σος) depends on the properties of the filters encoded in 6. on the exposure (mes in the (wo bands (/x aid fp). on the color of the skv background (Q4) and on the relative brightness of the sky when compared with the brightness of the objects (q).," for a given narrow-band magnitude, $\sigma[\mathcal{Q}_\mathrm{O}]$ depends on the properties of the filters encoded in $b$, on the exposure times in the two bands $t_{\narrowband}$ and $t_{\broadband}$ ), on the color of the sky background $\mathcal{Q}_{\mathrm{S}}$ ) and on the relative brightness of the sky when compared with the brightness of the objects $q$ )."
+ In Fig., In Fig.
+" 2. we show how the selection curve changes (with a fixed value of n,=3) with the exposure time and sky. background in the {wo bancs for the sample filler set CCAFOSS200 (the dilferent filter properties are detailed in sect. 6))", \ref{fig:seleccion2} we show how the selection curve changes (with a fixed value of $n_\sigma$ =3) with the exposure time and sky background in the two bands for the sample filter set CAFOS8200 (the different filter properties are detailed in Sect. \ref{sec:examples}) )
+ The increase of exposure tme or the reduction of the skv background (effectively increasing the sky-background magnitude) reduces the slope of the selection curve., The increase of exposure time or the reduction of the sky background (effectively increasing the sky-background magnitude) reduces the slope of the selection curve.
+ ]solines represent. curves joining points wilh an equal given property., Isolines represent curves joining points with an equal given property.
+ Lines with equal EW correspond with lines of equal color in (he magnitucde-color diagram. but lines with equal continuum or line (lux are not direetIv related with neither the color nor (he magnitude of anv of the bands.," Lines with equal EW correspond with lines of equal color in the magnitude-color diagram, but lines with equal continuum or line flux are not directly related with neither the color nor the magnitude of any of the bands."
+ We can compute the isolines in magnitude-color space starting wilh Eq. 9.., We can compute the isolines in magnitude-color space starting with Eq. \ref{eq:flujo2}.
+ If we assume a constant line flux. fj. the magnitudes of the objects as a function of f; and EW only are:," If we assume a constant line flux $\fluxl$, the magnitudes of the objects as a function of $\fluxl$ and EW only are:"
+the οz0 case. the phase velocity vanishes [R(u)—0].,"the $\Omega_F \neq 0$ case, the phase velocity vanishes $\Re(\omega) \rightarrow 0$ ]."
+ This has two consequences., This has two consequences.
+" First. a larger Traction of the heavier lavers momentum is available to help destabilize the flow (ancl so. the laver remains unstable for larger Ii, j)."," First, a larger fraction of the heavier layer's momentum is available to help destabilize the flow (and so, the layer remains unstable for larger $\Ri_{\eff}$ )."
+ And second. a vanishingly small w means a small A4 (eq. [20]]).," And second, a vanishingly small $\omega$ means a small $K_1$ (eq. \ref{k1_eq}] ]),"
+ so that vertical pressure gradients in the gas laver are very small., so that vertical pressure gradients in the gas layer are very small.
+ So. as the cust/gas mixture flows through the sinuous focus of the perturbed midplane laver. only the centrifugal force in the vertical direction (VT0z) and gravity (02.0z) are involved in the force balacing. while gas pressure gradients are negleeible.," So, as the dust/gas mixture flows through the sinuous focus of the perturbed midplane layer, only the centrifugal force in the vertical direction $V^2 k_y^2 \, \delta z$ ) and gravity $\Omega_K^2 \, \delta z$ ) are involved in the force balacing, while gas pressure gradients are neglegible."
+ In contrast. for the ο=0 case. Rie) is non-vanishing at the onset of instability. vielding a finite A4 value.," In contrast, for the $\Omega_F = 0$ case, $\Re(\omega)$ is non-vanishing at the onset of instability, yielding a finite $K_1$ value."
+ This implies non-neglegible vertical pressure gradients. which introduce a dependence on(he dust-to-gas density ratio in the force balance.," This implies non-neglegible vertical pressure gradients, which introduce a dependence onthe dust-to-gas density ratio in the force balance."
+ Small or vanishing Wie) is possible only in the οςz0 case. when the Coriolis forces (rom ry) are able to partly of fully compensate lor azimuthal pressure [orces.," Small or vanishing $\Re(\omega)$ is possible only in the $\Omega_F
+\neq 0$ case, when the Coriolis forces (from $v_{1x}$ ) are able to partly of fully compensate for azimuthal pressure forces."
+" We can allernatively write the instabilitv condition for O740 in terms of more basic physical parameters by considering equation (26)) in the large and small ji limits: Notice that the minimum stable laver thickness approaches a constant value as Xj/Y, increases. as shown by the dotted lines of Figure 2.."," We can alternatively write the instability condition for $\Omega_F \neq 0$ in terms of more basic physical parameters by considering equation \ref{rich_eq}) ) in the large and small $\mu$ limits: Notice that the minimum stable layer thickness approaches a constant value as $\Sigma_d / \Sigma_g$ increases, as shown by the dotted lines of Figure \ref{sdh_fig}."
+ This stands in contrast with the instability criterion for the Ὡς=0 case (see equation [25] and Figure 2)). for which (Ha/Leninoxδα)| as the surface density increases.," This stands in contrast with the instability criterion for the $\Omega_F = 0$ case (see equation\ref{instb_eq}] ] and Figure \ref{sdh_fig}) ), for which $(H_d/H_g)_{crit} \propto (\Sigma_d / \Sigma_g)^{-1}$ as the surface density increases."
+" Thus. for large pr. the instability criterion is Qnοκ lor Qj40. compared to GabyLL70,<Coamarh’y for Qe=0."," Thus, for large $\mu$, the instability criterion is $\Omega_K <
+v_{0,max} k_y$ for $\Omega_F \neq 0$, compared to $(\mu k_y H_d)^{1/2} \Omega_K < v_{0,max} k_y$ for $\Omega_F = 0$."
+ While we have concentrated on the long-wavelength limit in the discussion of this section. we note that the general dispersion relation of equation (23)) can easily be solved in the lint ke]=0.," While we have concentrated on the long-wavelength limit in the discussion of this section, we note that the general dispersion relation of equation \ref{disprel_eq}) ) can easily be solved in the limit $|\omega| \rightarrow 0$."
+" When 0,40. an unstable solution will appear when Εάν passes below the value with growth rate x(Bi,;—RiggJQe."," When $\Omega_F \neq 0$, an unstable solution will appear when $\Ri_{\eff}$ passes below the value with growth rate $\propto (\Ri_{crit}-\Ri_{\eff})\Omega_F$."
+ Thus. the necessary criterion for instability in (he lone-waveleneth lait is quite similar to a more general cvilerion for a niarginallv-unstable mode to appear.," Thus, the necessary criterion for instability in the long-wavelength limit is quite similar to a more general criterion for a marginally-unstable mode to appear."
+ Since lanh(y)/\> 1. Hi; is in practice slightly larger than the value deduced in the long-wavelength limit.," Since $\tanh(\chi)/\chi > 1$ , $\Ri_{crit}$ is in practice slightly larger than the value deduced in the long-wavelength limit."
+lk denotes pixels corresponding (o a certain sidelobe. then the overall sidelobe response is SAαμ-GP)TL/N.,"$k$ denotes pixels corresponding to a certain sidelobe, then the overall sidelobe response is $\sum_k(G_k^A-G_k^B)T_k/N$."
+" The overall sidelobe response should contribute a pseudo component to the dilferential Doppler signal The spacecraft LOS vectors and n, are all constant vectors in spacecraft. coordinates.", The overall sidelobe response should contribute a pseudo component to the differential Doppler signal The spacecraft LOS vectors and $\bf{n}_{_k}$ are all constant vectors in spacecraft coordinates.
+ Since G ave normalized gains. it can also be (reated as constant.," Since $G$ are normalized gains, it can also be treated as constant."
+" Therefore we have two constant vectors: n?Tw=So,Gin,/N and Ant—$5,GPn/N."," Therefore we have two constant vectors: $\Delta\bf{n}_{_A}^*=\rm\it\sum_k G_k^A\, \bf{n}_{_k}/\it{N}$ and $\Delta\bf{n}_{_B}^*=\rm\it\sum_k G_k^B\, \bf{n}_{_k}/\it{N}$."
+" Since the sidelobe gain is much smaller than main lobe gain. n? and An, are (wo small constant vectors. and By comparing Eq."," Since the sidelobe gain is much smaller than the main lobe gain, $\Delta\bf{n}_{_A}^*$ and $\Delta\bf{n}_{_B}^*$ are two small constant vectors, and By comparing Eq."
+ 2 and Eq., 2 and Eq.
+ 4. we can see that the entire sidelobe Doppler pickup can be exactly described by a small constant deviation to the differential spacecralt LOS vector.," 4, we can see that the entire sidelobe Doppler pickup can be exactly described by a small constant deviation to the differential spacecraft LOS vector."
+ sinilarlv. the effect of the sidelobe response uncertainty on the differenüal Doppler dipole is also equivalent to introducing a small differential LOS error An*.," Similarly, the effect of the sidelobe response uncertainty on the differential Doppler dipole is also equivalent to introducing a small differential LOS error $\Delta\bf{n}^*$."
+ 1t is important to notice that. even if the LOS is absolutely accurate. Chis pseudo-dipole signal from. the sidelobe response uncertainty can still exist.," It is important to notice that, even if the LOS is absolutely accurate, this pseudo-dipole signal from the sidelobe response uncertainty can still exist."
+ The amplitude of the equivalent LOS error induced by sidelobe Doppler pickup can be estimated by An;|/In4arcsin(||). which are ~50—15 uesjectablefor different bands.," The amplitude of the equivalent LOS error induced by sidelobe Doppler pickup can be estimated by $\arcsin(|\Delta\bf{n}_{_A}^*|/|\bf{n}_{_A}|)$, which are $\sim 50' - 75'$ for different bands."
+ The WMAP team believes that the Galactic sidelobe pickup is except for the Ix-band. the reason is largelv that the Galactic emission is much weaker in other bands (han in the K-band (Barnesοἱal.2003).," The WMAP team believes that the Galactic sidelobe pickup is neglectable except for the K-band, the reason is largely that the Galactic emission is much weaker in other bands than in the K-band \citep{barnes03}."
+. The case is sienilicantly different lor the Doppler signals because thev have almost the same amplitucle in all bancls., The case is significantly different for the Doppler signals because they have almost the same amplitude in all bands.
+ Moreover. the Doppler signal is as strong as several mlx all over the skv: however. the Galactic emission is strong onlv for the low Galactic latitude ancl (he strength decreases rapidly for higher latitude.," Moreover, the Doppler signal is as strong as several mK all over the sky; however, the Galactic emission is strong only for the low Galactic latitude and the strength decreases rapidly for higher latitude."
+ Thus the overall power of the Doppler signal is about 10—40 times higher than the overall Galactic foreground emission power. ancl the effect of sidelobe Doppler signal pickup is at least 10 times stronger than the sidelobe Galactic pickup.," Thus the overall power of the Doppler signal is about $\sim 10-40$ times higher than the overall Galactic foreground emission power, and the effect of sidelobe Doppler signal pickup is at least 10 times stronger than the sidelobe Galactic pickup."
+ The WAIAP antenna sidelobe gain patterns are estimated by ground-based measurements and in-flight lunar measurements., The WMAP antenna sidelobe gain patterns are estimated by ground-based measurements and in-flight lunar measurements.
+ For the Ix-band. the WALAP first-vear in-flight) gain measurenents are 6056 svstematically brighter than ground-based measurements.," For the K-band, the WMAP first-year in-flight gain measurements are $60\%$ systematically brighter than ground-based measurements."
+ The WAIAP (eam scaled up the ground-based measurements by 30% and scaled down the Iunar results, The WMAP team scaled up the ground-based measurements by $30\%$ and scaled down the lunar results
+coordinate svstem.,coordinate system.
+ Here.5 is the flux of the leused image. so that ο=Sy/detAy. as usual. where detdy is the Jacobian at the origin 0=0.," Here,$S$ is the flux of the lensed image, so that $S=S_0/\det\A_0$, as usual, where $\det\A_0$ is the Jacobian at the origin $\theta=0$."
+ The origin of the coordinates in the source plane is the inage of the origin in the leus pluie as mapped with the leus equation., The origin of the coordinates in the source plane is the image of the origin in the lens plane as mapped with the lens equation.
+ In particular. this does not coincide with the center-ofleht of the source. which is given by Mouj[/ Sy. or 03)) II(0))e," In particular, this does not coincide with the center-of-light of the source, which is given by $\bar\beta\equiv {\rm Mom}[\beta]/S_0$ , or ) )."
+ Expanding the iutegrand. we note that terms linear in 0 vanish. due to (30)).," Expanding the integrand, we note that terms linear in $\theta$ vanish, due to \ref{eq:center}) )."
+ Define the second-order brightuess. moments of. the image. iu. the form+ HIQO)): ET(0)). we obtain for the source centroid shitt OQ.ezva) έν Mae," Defining the second-order brightness moments of the image in the form ); ), we obtain for the source centroid shift Q_2 Q_0 Q_2^*."
+(QU We now write these equations in a more compact form: for this. we define the matrix G by GT=(GS.G1.G4.G3). where the T denotes the transpose of the mati.," We now write these equations in a more compact form; for this, we define the matrix $\tens{G}$ by $\tens{G}^{\rm
+T}=(G_3^*,G_1^*,G_1,G_3)$, where the `T' denotes the transpose of the matrix."
+ Then.ΙΡ. whereWere the| coefficieutseoohcieη of: —B=(0.55.53.04) aresey givenovn byM," Then, where the coefficients of $\tens{B}=(b_1,b_2,b_3,b_4)$ are given by."
+Et The ceutroid shift iu the source plane is thus eiven bv the product of the derivatives of the reduced. shear (expressed by Gyand 63) and the area of the image. which is proportional to Qu and Qe.," The centroid shift in the source plane is thus given by the product of the derivatives of the reduced shear (expressed by $G_1$and $G_3$ ) and the area of the image, which is proportional to $Q_0$ and $Q_2$."
+ Of course. since the reduced shear aud itsderivatives arenot directly observable. the centroid. shiftMu in unobservable as well.," Of course, since the reduced shear and itsderivatives arenot directly observable, the centroid shift in unobservable as well."
+" To ect au orcder-of-.magnitude estimate of J. we assuiie that thesource has a linear angular size O,. cousider the reduced shear to be of order unity. aud let Ου be the angular scale ouwhich the reduced shearvaries."," To get an order-of-magnitude estimate of $\bar\beta$, we assume that thesource has a linear angular size $\Theta_{\rm s}$, consider the reduced shear to be of order unity, and let $\Theta_{\rm c}$ be the angular scale onwhich the reduced shearvaries."
+ Then.Fuit.," Then,."
+" Next we consider the secoud-order brightuess moments of the source. defined as ο=Moi?n/syMoi]""|S and Q)-— Mou[(ncanj/Syο1” A£."," Next we consider the second-order brightness moments of the source, defined as $Q_2^{\rm s}={\rm Mom}[(\beta-\bar\beta)^2]/S_0
+={\rm Mom}[\beta^2]/S_0-\bar\beta^2$ and $Q_0^{\rm s}={\rm Mom}[(\beta-\bar\beta)(\beta-\bar\beta)^*]/S_0={\rm Mom}[\beta\beta^*]/S_0-\bar\beta\bar\beta^*$ ."
+ By defining the briehtuess° moments of the image? through? EI(QO)): ET(0)) . κ =QeT E 2gQuE | PPSτὸ ES T1 ggMT boetar.. = g Qs|{lL.ga Qu- GQ)- CAT?," By defining the third-order brightness moments of the image through ); ) we obtain = Q_2 - 2g Q_0 + g^2 Q_2^* T_3 T_1 T_1^* T_3^* , = -g^* Q_2+(1+g )Q_0-g Q_2^* T_3"
+D appears able to destabilize a circular. coplanar svstem.,"B appears able to destabilize a circular, coplanar system."
+" Next we relax (he requirement (hat the planets began on coplanar orbils and ran simulations with initial mutual inclinations V,=3°. 10. and 30. and with αν=700. ου=O. ig=30°."," Next we relax the requirement that the planets began on coplanar orbits and ran simulations with initial mutual inclinations $\Psi_0 =
+3^\circ$ , $10^\circ$ , and $30^\circ$ , and with $a_B = 700$, $e_B = 0$, $i_B = 30^\circ$."
+ The two planets began with their current. best fit semi-major axes and masses. but on circular orbits. ancl one inclination was set to 3°. 107 or 30° with the masses held constant.," The two planets began with their current best fit semi-major axes and masses, but on circular orbits, and one inclination was set to $3^\circ$, $10^\circ$ or $30^\circ$ with the masses held constant."
+ These svstems were integrated for 1 Gyr., These systems were integrated for 1 Gyr.
+ In each of these cases. shown in Fig. 2..," In each of these cases, shown in Fig. \ref{fig:kozainonzero},"
+ the initial value of V is maintained for the duration of the simulation., the initial value of $\Psi$ is maintained for the duration of the simulation.
+ The widths of the bration increase wilh V because (he interactions among the planets are driving a secular oscillation., The widths of the libration increase with $\Psi_0$ because the interactions among the planets are driving a secular oscillation.
+ It therefore seems (hat even if (he planets began wilh a nonzero relative inclination. D is unable to piunp it to the range shown in Fig. 1..," It therefore seems that even if the planets began with a nonzero relative inclination, B is unable to pump it to the range shown in Fig. \ref{fig:secular}."
+ These simulations also demonstrate that plausible orbits of B will not destabilize the observed svstem., These simulations also demonstrate that plausible orbits of B will not destabilize the observed system.
+ These simulations indicate that it is unlikely that B could have twisted the orbits into the mutually inclined svstem we see today., These simulations indicate that it is unlikely that B could have twisted the orbits into the mutually inclined system we see today.
+" However. it could have destabilized the svstem. which we will see in the next section is a process which can lift coplanar orbits to VS30""."," However, it could have destabilized the system, which we will see in the next section is a process which can lift coplanar orbits to $\Psi \gsim 30^\circ$."
+ Our second hvpothesis considers the possibility that the planetary svstem [ormed in an unstable configuration independent of D. and that encounters between planets ultimately ejected an original companion leaving a svstem with biel mutual inclinations.," Our second hypothesis considers the possibility that the planetary system formed in an unstable configuration independent of B, and that encounters between planets ultimately ejected an original companion leaving a system with high mutual inclinations."
+ Such impulsive interactions could drive inclinations to laree values. perhaps as large as 607 (Alarzari Weidenschiling 2002: Chatterjee 2008: Ravinond 2010).," Such impulsive interactions could drive inclinations to large values, perhaps as large as $60^\circ$ (Marzari Weidenschilling 2002; Chatterjee 2008; Raymond 2010)."
+ We considered 41.000 different initial configuration of the svslenm. one or two additional planets with masses in the range 1. 15 Jap:," We considered 41,000 different initial configuration of the system, one or two additional planets with masses in the range 1 – 15 $_{\textrm{Jup}}$."
+ At (he end of this section. we summarize the results of all these simulations. but initially we focus on one subset.," At the end of this section we summarize the results of all these simulations, but initially we focus on one subset."
+ We completed 5.000 simulations that began with three 10. 15 Myup (Chambersmass objects (uniformly. distributed in mass) separated by 45 mutual Lill radii 1996). with e< 0.05.7<1 and 0.75<a«4 AU.," We completed 5,000 simulations that began with three 10 – 15 $_{\textrm{Jup}}$ mass objects (uniformly distributed in mass) separated by 4–5 mutual Hill radii (Chambers 1996), with $e <
+0.05$, $i < 1^\circ$ and $0.75 < a < 4$ AU."
+ We integrated these cases for 100 vears with Merenry’s hvbrid integrator. conserving energy to 1 part in 104.," We integrated these cases for $10^6$ years with Mercury's hybrid integrator, conserving energy to 1 part in $10^4$."
+ About of cases failecl lo conserve energy al this level and were thrown out., About of cases failed to conserve energy at this level and were thrown out.
+ These ranges are somewhat arbitrary but Follow the recent study by Bayxmond (2010). which considered smaller mass planets ab larger distances.," These ranges are somewhat arbitrary but follow the recent study by Raymond (2010), which considered smaller mass planets at larger distances."
+ They found that a systemconsistingofthree 2. jaup planets could. alter removal of one planet aud settling into a stable configuration. end up with V>30 ," They found that a systemconsistingofthree 3 $_{\textrm{Jup}}$ planets could, after removal of one planet and settling into a stable configuration, end up with $\Psi > 30^\circ$ "
+"with the standard deviation given by expressions: The probability that the amplitude is greater than a certain chosen value is giveu by the formula with standard deviation of this amplitude This test was originally introduced by TawleyPeebles(L975) and substantially modified by Codlowski(1991). for the case when talking iuto account higher Fourier modes: Iu our case (all Nj,; are equal) it leads to forunas for the Aj; coefficients: aud with the standard deviation aud Tf we analvsed Fourier iiodes separately. probability that the amplitude is greater than a certain chosen value is giveu by the formula: When we analysed first and second Fourier modes together we obtain andl The value of cocficient Ay, gives us the direction of departure from isotropy.","with the standard deviation given by expressions: The probability that the amplitude is greater than a certain chosen value is given by the formula with standard deviation of this amplitude This test was originally introduced by \citet{h4} and substantially modified by \citet{g3} for the case when taking into account higher Fourier modes: In our case (all $N_{0,k}$ are equal) it leads to formulas for the $\Delta_{ij}$ coefficients: and with the standard deviation and If we analysed Fourier modes separately, probability that the amplitude is greater than a certain chosen value is given by the formula: When we analysed first and second Fourier modes together we obtain and The value of coefficient $\Delta_{11}$ gives us the direction of departure from isotropy."
+ If A44<0 then the excess of the eroup with position angles near 90 degrees (parallel to Local Supercluster plane) is observed. while for Ayy70 the excess siat Me loop with position aueles perpendicular to Local Supercluster plane is observed.," If, $\Delta_{11}<0$ then the excess of the group with position angles near $90$ degrees (parallel to Local Supercluster plane) is observed, while for $\Delta_{11}>0$ the excess of the group with position angles perpendicular to Local Supercluster plane is observed."
+ The isotropy the resultant cdistributious of the angles o was investigated using Ixolmogorov- Suuiznovtest (Ik-S test)., The isotropy the resultant distributions of the angles $\phi$ was investigated using Kolmogorov- Smirnovtest (K-S test).
+ We divided our sample according to distance D between group centres., We divided our sample according to distance $D$ between group centres.
+ We assumed that the theoretical. random distribution contains the same muuber of objects as the observed oue.," We assumed that the theoretical, random distribution contains the same number of objects as the observed one."
+ Our null lwpothesis ZZ) is that the distribution is randoni ouc., Our null hypothesis $H_0$ is that the distribution is random one.
+ In order to reject the Πρ hypothesis. the value of observed statistics A should be greater than Αι ," In order to reject the $H_0$ hypothesis, the value of observed statistics $\lambda$ should be greater than $\lambda_{cr}$ ."
+At the significance level a=0.05 the value A...= 1.358., At the significance level $\alpha=0.05$ the value $\lambda_{cr} = 1.358$ .
+"distribution in the number density of early-type galaxies, in better agreement with the data.","distribution in the number density of early-type galaxies, in better agreement with the data."
+" A very preliminary test to improve on the flatness of the predicted size-mass relation would be to simply remove from the mock catalog the large and less massive galaxies (Rez3 kpc and Mstar~10 ΜΟ)), and the compact and massive ones (ReS1 kpc and Mstar~10 Mo))."," A very preliminary test to improve on the flatness of the predicted size-mass relation would be to simply remove from the mock catalog the large and less massive galaxies $\gtrsim 3$ kpc and $\sim 10^{10}$ ), and the compact and massive ones $\lesssim 1$ kpc and $\sim 10^{11}$ )."
+" One simple way to do this would be to assume that those galaxies which are strong outliers with respect to the observed local size-mass relation in Figure 3,, should be considered as irregulars and/or disk-dominated galaxies instead of spheroids (e.g., ?,, and references therein)."," One simple way to do this would be to assume that those galaxies which are strong outliers with respect to the observed local size-mass relation in Figure \ref{fig|SizeMassRelation}, should be considered as irregulars and/or disk-dominated galaxies instead of spheroids (e.g., \citealt{Hopkins09transf}, and references therein)."
+" The region in the pplane within the dotted lines in Figure 3 roughly defines the “region of acceptance"" at a ~3c level.", The region in the plane within the dotted lines in Figure \ref{fig|SizeMassRelation} roughly defines the “region of acceptance” at a $\sim 3 \sigma$ level.
+ All the mock galaxies which lie below and above those lines can be safely considered as outliers., All the mock galaxies which lie below and above those lines can be safely considered as outliers.
+" Recomputing the predicted ERF and SMF only considering the subsample of galaxies within the dotted lines, yields the red long-dashed lines in Figures 4empha,b,c,d.."," Recomputing the predicted ERF and SMF only considering the subsample of galaxies within the dotted lines, yields the red long-dashed lines in Figures \ref{fig|Models}."
+" It is clear that, irrespective of the actual physical basis motivating such drastic cut in the mock catalog of early-type galaxies, this is not a satisfactory solution, as it severely further underpredicts the number density of massive galaxies, without significantly reducing the bump around Mia1.5x10!°Mo,, and also worsens the match to the ERF."," It is clear that, irrespective of the actual physical basis motivating such drastic cut in the mock catalog of early-type galaxies, this is not a satisfactory solution, as it severely further underpredicts the number density of massive galaxies, without significantly reducing the bump around $\sim 1.5\times 10^{10}$, and also worsens the match to the ERF."
+ We have checked that a significant fraction of the outlier galaxies are satellites., We have checked that a significant fraction of the outlier galaxies are satellites.
+" Therefore, one possible way to remove the outlier galaxies might be to simply assume a higher rate of mergers with the central galaxy in the same parent halo."," Therefore, one possible way to remove the outlier galaxies might be to simply assume a higher rate of mergers with the central galaxy in the same parent halo."
+ Figure 5 shows the model predictions modified by assuming that all satellite galaxies with stellar mass below Metar=3X10'°Mo hhave merged with their respective centrals., Figure \ref{fig|ModelsImproved} shows the model predictions modified by assuming that all satellite galaxies with stellar mass below $=3\times 10^{10}$ have merged with their respective centrals.
+" It is apparent how this hypothetical increase in the merger rate of the less massive spheroids yields a considerably better agreement with the observed SMF, lowering the number density at low masses, and increasing it at higher masses."," It is apparent how this hypothetical increase in the merger rate of the less massive spheroids yields a considerably better agreement with the observed SMF, lowering the number density at low masses, and increasing it at higher masses."
+" Although the simple improvement proposed here for this model yields outputs in better agreement with the data, it requires a valid physical explanation and several"," Although the simple improvement proposed here for this model yields outputs in better agreement with the data, it requires a valid physical explanation and several"
+where Vua is the radial velocity measurement for each stars.,where $V_{{\rm rad}}$ is the radial velocity measurement for each stars.
+ The least-squares fit to this function ⋅∏↸∖∐∖↴⋜↧∖−↖⇁⋜↧↕∏↸∖↕∪↥⋅↸∖⋯⊳∐∣∣⋖⊡∶↴∙⊾∙−≽⋜↧⋜∪∙⋅, The least-squares fit to this function yields a $\chi^2$ value for each $\theta$ (Fig.
+ ⋅≻ . ⋅⊀, \ref{fig2}a a).
+" The: distribution: euclosed by μμ2|,1 corresponds to lo for a normal distribution (Pressetal.1992)."," The distribution enclosed by $\chi^2_{\rm min}+1$ corresponds to $1 \,
+\sigma$ for a normal distribution \citep{1992nrfa.book.....P}."
+ We use this to obtain the error in the position angle which. minimizes∙∙∙D7., We use this to obtain the error in the position angle which minimizes$\chi^2$.
+ Wo ⇁⋅find a positionee angle. of. 35.!55Is degrees. with. a velocity eradient of 1623xlanss !kpe |. aud a zeropoiut of [5.11c0.381ans!.," We find a position angle of $-35^{+18}_{-23}$ degrees with a velocity gradient of $16 \pm 3$ $^{-1}$ $^{-1}$, and a zeropoint of $45.11\pm 0.38{\rm \, km\, s^{-1}}$."
+ We obtain a reduced 4? of 3.89 for our 18 stars with2 degrees of freedoun., We obtain a reduced $\chi^2$ of $3.89$ for our 18 stars with2 degrees of freedom.
+ Following Walkerctal.(2008).. we estimate the significance of the velocity. gradient using 10° Monte Carlo realisations.," Following \citet{2008ApJ...688L..75W}, we estimate the significance of the velocity gradient using $10^5$ Monte Carlo realisations."
+ Ta each realisation. we sample the velocity aud spatial distributions independently.," In each realisation, we sample the velocity and spatial distributions independently."
+ Thus we scramble the correlation between velocity aud position. while maiutainimg the original velocity distribution and spatial positions.," Thus we scramble the correlation between velocity and position, while maintaining the original velocity distribution and spatial positions."
+ This is valid as long as the pliase space distribution function of the stars is separable (which implies that the velocity dispersion is independent of radius)., This is valid as long as the phase space distribution function of the stars is separable (which implies that the velocity dispersion is independent of radius).
+ We determine the significance of the velocity eradieut by computing the fraction of realisations that fail to produce a \? as low as the one calculated from the real data., We determine the significance of the velocity gradient by computing the fraction of realisations that fail to produce a $\chi^2$ as low as the one calculated from the real data.
+ We find a sienificance of the velocity eradicut of 78 per ceut (1.230)., We find a significance of the velocity gradient of 78 per cent $1.23 \sigma$ ).
+ In the previous section we found tentative evidence for a velocity eradicut in the Iercules dSphli., In the previous section we found tentative evidence for a velocity gradient in the Hercules dSph.
+ The peak-to-peak difference of ~10s within a radius of less than 1kkpe could be interpreted as the effect of Galactic tides (Reacetal. 2006b)..," The peak-to-peak difference of $\sim 10
+\, {\rm km \, s^{-1}}$ within a radius of less than kpc could be interpreted as the effect of Galactic tides \citep{2006MNRAS.367..387R}."
+ Iuterestiugh Martin—(2009) recentlyestimated the orbit of Hercules base ou the assumption that its elongation is tidally induced. and predicted a velocity eradieut of a least -Thaustkpe*1 which. is. eonsisteut. with. our observed eradicut.," Interestingly, \citet{Martin2009}
+ recentlyestimated the orbit of Hercules based on the assumption that its elongation is tidally induced, and predicted a velocity gradient of at least $7\, {\rm km \, s^{-1}} {\rm kpc}^{-1}$ which is consistent with our observed gradient."
+ Additionally. in Adénctal.(2009) we found that the spatial distribution of the Hercules stars is asvinmetric at ~0.5 kpe.," Additionally, in \citet{Aden2009} we found that the spatial distribution of the Hercules stars is asymmetric at $\sim$ 0.5 kpc."
+ There are three significant outlicr stars to the South. but no corresponding members at this distance im either the Northern or Western fields (Fig. 3)).," There are three significant outlier stars to the South, but no corresponding members at this distance in either the Northern or Western fields (Fig. \ref{fig3}) )."
+ These three stars are unambiguously identified as RGB stars from Strónuueren photometry., These three stars are unambiguously identified as RGB stars from Strömmgren photometry.
+ We now consider the possibility that the velocity eradient aud the positional outlicrs are evidence that Uereules is beiug tically disrupted. aud use this information to obtain a second mass estimate for the svsteun.," We now consider the possibility that the velocity gradient and the positional outliers are evidence that Hercules is being tidally disrupted, and use this information to obtain a second mass estimate for the system."
+ The tidal radius of a dSph depeuds on the potential of the host galaxy. the potential of the dSph. the orbit of the dSph within the host ealaxy aud the orbit of the stars within the dSph (e.g.Readetal.2006c).. ," The tidal radius of a dSph depends on the potential of the host galaxy, the potential of the dSph, the orbit of the dSph within the host galaxy and the orbit of the stars within the dSph \citep[e.g.][]{2006MNRAS.366..429R}. ."
+We paraueterise the \ilky Way potential using the default model iu Johustonetal.(2005).. analysed in the Calactic plane: aud the Hercules potential using a generalised Ieruquist profile (IHeruquist1990):: where AM is the total mass. (£y the scaleleneth. and 5 the ceutral logarithmic cusp slope.," We parameterise the Milky Way potential using the default model in \citet{2005ApJ...619..800J}, analysed in the Galactic plane; and the Hercules potential using a generalised Hernquist profile \citep{1990ApJ...356..359H}: where $M$ is the total mass, $r_s$ the scalelength, and $\gamma$ the central logarithmic cusp slope."
+ We consider the ranges 0.3<r.«3kkpe aud 0«5 1l., We consider the ranges $0.3<r_s<3$ kpc and $0<\gamma<1$ .
+ Our results are not sensitive to these choices., Our results are not sensitive to these choices.
+ The orbit of the Hercules dSphli is uuxnown but it currently Hes at a distance of 132 kkpe from the ealactic centre. aud has a hehoceutric velocity of 15.2E1.09]aunsbt. which iuplies a 8ealactocentric radial velocity of 14521.09Ίαν+ (seeEq.(5)iuCourteau&vaucenBereh 1999)..," The orbit of the Hercules dSph is unknown but it currently lies at a distance of $132$ kpc from the galactic centre, and has a heliocentric velocity of $45.2 \pm 1.09 {\rm \, km\, s^{-1}}$, which implies a galactocentric radial velocity of $145 \pm 1.09 {\rm \, km\,
+ s^{-1}}$ \citep[see Eq. (5) in][]{1999AJ....118..337C}. ."
+ Using tle above potential model for the Milky: Way aud setting the tanecutialo velocity component for Tercules to zero. ↑∐↕↴∖↴∶↴∙⊾↕↖⇁↸∖↴∖↴∏↴∖↴⋜↧∐∐∐↕∐∐∐⊔⋜⋯∪↸⊳↸∖∐⊓⋅↸∖↕≯∪↥⋅↕↕↸∖↥⋅↸⊳↿∏↸∖↴∖↴∪↕⋟ ∣⋮∣∣≱⋯↕⊔∶↕≺∖∖≺∖," Using the above potential model for the Milky Way and setting the tangential velocity component for Hercules to zero, this gives us a minimum apocentre for Hercules of $r_{a,\mathrm{min}} = 188.5$ kpc."
+∖∙⋅↱⊐↨↘↽↨↘↴⋉⊳∙↽∕∏⋯↴∖↴∙↖↖⇁↸∖↸⊳∪∐↴∖↴↕≼∐∖↥⋅⋜∪∪∩∖∐⊓⋅↸∖ ↕⊰↸∖⋜∥↧↸∖↑⋜↧↕∙⊔∩∩⊓↸⊳⋝∱∎⋯≼⊓↕⋯↑↻∐∪∪⋯↸∖↑∏↸⊳ ↕⋟↸∖⋜↧⊓∐⋅↸∖↴∖↴⋜∐⋅↸∖↑⋅," Thus, we consider apocentre and pericentre ranges of $188.5 < r_a < 600$ kpc and $5< r_p<132$ kpc, respectively."
+↖↴≻↕↸⊳⋜↧∐∙↖⇁↴∖↴↸∖↸∖∐↴⋝↸∖∙↖↽∪∐≺↧↑∐↸∖↥⋅↸∖↥⋅∪∶↴⋁↥⋅⋯∐∖ ↑↕≼↧⋜↧⊔⋅⋜∥∐∏↴∖↴∙⊺↴∪↻↥⋅∪↸⊳↸∖↸∖≼↧∙↖↖⇁↸∖⋯⋜∐↘↽↸∖↑∐↸∖⋜↧↴∖∷∖↴∏∐∏≻↑↕∪∐, \citet{2006MNRAS.366..429R} find that photometric features are typically seen beyond the retrograde tidal radius.
+ ↑∐⋜↧↑↑↕∐∖⋯↕↑∐↸∖↥⋅↴∖↴↕∐⊟∶↴⋁∙∶≩↕∐≺∐↸⊳⋜⋔∖↑↕∐∖↕∪↸⊳⋜↕↑↕∪∐ ∪↕≯↑↕∐," To proceed, we make the assumption that the outliers in Fig."
+∖↥⋅↸∖⊓⋅∪∶↴∙⊾↥⋅⋜∥∐∖↑↕≼↧⋜↧↕↥⋅⋜⊔∐∏↴∖↴∪↕≯∐↸∖↥⋅↸⊳∏↕↸∖↴∖↴∙↕∙↸∖∙ ∣⋮⊤∿↓≺∖∖⋅↱⊐⋯⋉⊳∙↕↕⋟↖↖↽↸∖⋜↧↴∖↴↴∖↴⋯⋯∖↕≯↿∐⋅↑∐↸," \ref{fig3} indicate the location of the retrograde tidal radius of Hercules, i.e. $r_t \sim
+485$ pc."
+∖↥⋅↑∐⋜↧↑↑∐↸∖↑↕≼↧⋜↧↕ ↥⋅⋜⊔∐∏↴∖↴∪↕⋟∐↸∖↥⋅↸⊳∏↕↸∖↴∖↴↕↴∖↴↴∖↴↸∖↑⋜↧↑↑∐↸∖↻↸∖↥⋅↕↸⊳↸∖↕↓⊓⋅↸∖∪↕≯↕↑↴∖↴ orbit. we can solve Eq. (," If we assume further that the tidal radius of Hercules is set at the pericentre of its orbit, we can solve Eq. ("
+7) of Readetal.(20060). to calculate the mass. Mig. of Hercules as a function of its orbital periceutre c. for our assuued rages of ry. 5 aud rc.,"7) of \citet{2006MNRAS.366..429R} to calculate the mass, $M_{\mathrm{tid}}$, of Hercules as a function of its orbital pericentre $r_p$ , for our assumed ranges of $r_s$ , $\gamma$ and $r_a$ ."
+" We calculateboth its mass within ry (Maa, ) anc its mass within ppe (Mua iss} which can then be compared with our masμα estimate based on the spherical Jeans equatio-CMaST,133 qna)."," We calculateboth its mass within $r_t$ $M_{\mathrm{tid},r_t}$ ) and its mass within pc $M_\mathrm{tid,433}$ ) which can then be compared with our mass estimate based on the spherical Jeans equation$M_{\mathrm{SJ},433}$ )."
+ The results of this calculation are neiven, The results of this calculation are given
+accretion disc (Fabian&Rees1979:Milgrom1010:Abell&Margon 1979).,"accretion disc \citep{FabianRees:1979,
+Milgrom:1979,AbellMargon:1979}."
+". Phe ""moving"" lines are also seen in the Chandra X-ray spectra (Marshallctal.2002)", The “moving” lines are also seen in the Chandra X-ray spectra \citep{MarshallEt:2002}.
+" Another major periodicity found is the so-called. 7nodding"" of the disc (Ilxatzetal.1982) "," Another major periodicity found is the so-called “nodding"" of the disc \citep{KatzEt:1982}. ."
+This is a 6.28 day period. ancl is attributed to the donor star's passage through the nodal line of the titled disc which causes tidal deformation of the disc., This is a $6.28$ day period and is attributed to the donor star's passage through the nodal line of the titled disc which causes tidal deformation of the disc.
+ In addition to the jets there is evidence of a very powerful outflow from the disc with a velocity of approximately. 1500 !(seeBegelman.King.&Pringle2006.andreferencestherein).. elements, In addition to the jets there is evidence of a very powerful outflow from the disc with a velocity of approximately $1500$ \citep[see][and references therein]{BegelmanEt:2006}.
+ Figurela of Barnesetal.(2006) shows the major of 4433. which include the compact object. the Roche filling donor star. the aceretion disc. the precessing relativistic jets and the outflow in a form of extend disc or disc wind.," Figure 1a of \cite{BarnesEt:2006} shows the major elements of 433, which include the compact object, the Roche filling donor star, the accretion disc, the precessing relativistic jets and the outflow in a form of extend disc or disc wind."
+ The precessional ancl orbital periods have provided a very satisfactory model of the motion of thejets (Ixatzetal.1982:Collins&CGarasi1994) and the photometric variations (CGoranskietal.L99S).," The precessional and orbital periods have provided a very satisfactory model of the motion of the jets \citep{KatzEt:1982,CollinsGarasi:1994} and the photometric variations \citep{GoranskiiEt:1998}."
+. However. there are still many unanswered questions about the system.," However, there are still many unanswered questions about the system."
+ These include. the formation process for the jets and the disc precession mechanism (Wijers&Pringle1999).," These include, the formation process for the jets and the disc precession mechanism \citep{WijersPringle:1999}."
+ In Table we eive a summary of the main parameters we adopted for 4433., In Table \ref{HerX1:table:SS433Parameters} we give a summary of the main parameters we adopted for 433.
+ We take the basic hyverodsnamies equations governing the viscous gaseous material of the accretion disc., We take the basic hydrodynamics equations governing the viscous gaseous material of the accretion disc.
+ It is assumed that the disc sel-eravity is negligible in comparison to the binary components tidal field., It is assumed that the disc self-gravity is negligible in comparison to the binary component's tidal field.
+ The equations of continuity and motion may be written as where in the above is the material derivative operator (Batchelor1970)...MM p. v and P are the fluid. density. velocity. ancl pressure respectively.," The equations of continuity and motion may be written as where in the above is the material derivative operator \citep{Batchelor:1970}, , $\rho$, $\mathbf{v}$ and $P$ are the fluid density, velocity and pressure values respectively."
+ The gravitational potential is 9. the viscous force is fy and the force due to the radiation on the [uid is Ενν.," The gravitational potential is $\Phi$, the viscous force is $\mathbf{f_{v}}$ and the force due to the radiation on the fluid is $\mathbf{f_{irr}}$."
+ Phe equation for the rate of change of thermal energy per unit mass can be expressed as We solve the basic equations in section 3. using a SPILL (Lucy1977:Cingold&Monaghan.1977) computercode first developed by Murray (1996. 1998)," The equation for the rate of change of thermal energy per unit mass can be expressed as We solve the basic equations in section \ref{sec:Basic Equations} using a SPH \citep{Lucy:1977, GingoldMonaghan:1977} computercode first developed by Murray (1996, 1998)."
+ SPLL uses smoothed quantities defined by the equation here fr) is an arbitrary function. Wo isa smoothing kernel. fis the smoothing length and the integration is over the entire space.," SPH uses smoothed quantities defined by the equation here $f(\mathbf{r})$ is an arbitrary function, $W$ isa smoothing kernel, $h$ is the smoothing length and the integration is over the entire space."
+ A Monte-Carlo. representation of (5) is used for cach particle position where mj is the mass of particle j. p; is the density of the particle j and the sum is over all. particles N.. We use à variable smoothing length. such that cach particle. { παν an associated smoothing length h;.," A Monte-Carlo representation of $(\ref{eq:3})$ is used for each particle position where $m_j$ is the mass of particle $j$, $\rho_j$ is the density of the particle $j$ and the sum is over all particles N. We use a variable smoothing length such that each particle, $i$, has an associated smoothing length $h_i$."
+ The kernel is svnunetrized with respect to. particle pairs in order to CODserve monendum., The kernel is symmetrized with respect to particle pairs in order to conserve momentum.
+ The equation of motion for each particle ὁ is given by which gives the force on particle 7 as a summationof the pressure. gravity. viscousand irradiation forces. Ep;. Fe.;. Fuji.) and E;445 respectively.," The equation of motion for each particle $i$ is given by which gives the force on particle $i$ as a summationof the pressure, gravity, viscousand irradiation forces, $\mathbf{F}_{P,i}$ , $\mathbf{F}_{G,i}$ , $\mathbf{F}_{visc,i}$ and $\mathbf{F}_{irr,i}$ respectively."
+ Lor the pressure gradient we make use of the fact that and the pressure is given by Similarlv the SPILL energy. equation is, For the pressure gradient we make use of the fact that and the pressure is given by Similarly the SPH energy equation is
+eveuly over the central part of Arp 220. but to be concentrated in two knots which have a radius zz D pe each (2) (ic. the size of a UCD).,"evenly over the central part of Arp 220, but to be concentrated in two knots which have a radius $\approx$ 50 pc each \citep{lonsdale2006a} (i.e. the size of a UCD)."
+ This nuplies that iudeed a major part of the star formation in the central part of Arp 220 takes place within these two knots., This implies that indeed a major part of the star formation in the central part of Arp 220 takes place within these two knots.
+" This would iuplv projected star formation densities of a few 10OAL,vrobpe7 2.in the knots.", This would imply projected star formation densities of a few $10^{-3} \ {\rm M}_{\odot} \ {\rm yr}^{-1} \ {\rm pc}^{-2}$ in the knots.
+ SNRs calculated from Equation. (30)) for a constant SFR of 100ML.xr| are shown as functions of the ligh-imass slope of the IMF in Figure &..," SNRs calculated from Equation \ref{eq:SNR2}) ) for a constant SFR of $100 \ {\rm M}_{\odot} \, {\rm yr}^{-1}$ are shown as functions of the high-mass slope of the IMF in Figure \ref{fig:SNrate}."
+ The two curves correspond to differeut times at which the star burst was initialized. but the expected number of SN per vear (i6. the SNR) is low in anv case.," The two curves correspond to different times at which the star burst was initialized, but the expected number of SN per year (i.e. the SNR) is low in any case."
+ The umber of SN that occur within oue vear can therefore ditter substantially from the caleulated SNR. as the frequency of SN. over such a time span obevs low-unniber statistics.," The number of SN that occur within one year can therefore differ substantially from the calculated SNR, as the frequency of SN over such a time span obeys low-number statistics."
+ Thus. the probability for a certain ΠΠ of SN to happen within one vear is quantified by the Poisson distribution function.," Thus, the probability for a certain number of SN to happen within one year is quantified by the Poisson distribution function."
+ Now consider the case that the star-burst iu Arp 220 already lasts for more than LO Myr., Now consider the case that the star-burst in Arp 220 already lasts for more than 40 Myr.
+ This imuples mua=8Me. so that the nunuber of SNII per vear is maxiuized for the given SER.," This implies $m_{\rm low} = 8 \ {\rm M}_{\odot}$, so that the number of SNII per year is maximized for the given SFR."
+" The expectation value for the ο[ταῖς is then about one per vear if the ΤΝΤ was canonical (1.6. a3= 2.3). but about two per vear for a top-licavy IMF with lSa,=2. where the SN-rate is only a weak function of a; (cf"," The expectation value for the SNII-rate is then about one per year if the IMF was canonical (i.e. $\alpha_3=2.3$ ), but about two per year for a top-heavy IMF with $1 \lesssim \alpha_3 \lesssim 2$, where the SN-rate is only a weak function of $\alpha_3$ (cf."
+ Figure 8))., Figure \ref{fig:SNrate}) ).
+ Thus. the probability to actually observe our new SNII iu a given vear (2) d$ then abou two per cent if the IME is canonical. but abou 12 per cout for 150522.D4 ," Thus, the probability to actually observe four new SNII in a given year \citep{lonsdale2006a} is then about two per cent if the IMF is canonical, but about 12 per cent for $1 \lesssim \alpha_3 \lesssim 2$."
+"A anore elaborate discussion of the SNIT rate iu Arp 220 is obtained by taking iuto account that stars in a galaxy form in star-clusters of different Lüasses, ποσο μιας Of the IAIF depeuds on the mass of the star-cluster for low-mass star-clusters."," A more elaborate discussion of the SNII rate in Arp 220 is obtained by taking into account that stars in a galaxy form in star-clusters of different masses, since $m_{\rm max}$ of the IMF depends on the mass of the star-cluster for low-mass star-clusters."
+ This implies that the iuteerated galactic IME GGIME) of all star-clusters in Arp 220 combined is not equal to the IME in its star-chisters., This implies that the integrated galactic IMF (IGIMF) of all star-clusters in Arp 220 combined is not equal to the IMF in its star-clusters.
+ This IGIME is even bx wherem is the iitial stellar mass. Af.) is the initial stellar mass of a star cluster. τω ds the minima mass of newly formed star-clusters. Moclauax(SFR) is the SFR-cdependenut maximal mass of newly formed star-clisters. 1) ds the IME aud £440AM44) is the stir-cluster nias function (??73..," This IGIMF is given by where$m$ is the initial stellar mass, $M_{\rm ecl}$ is the initial stellar mass of a star cluster, $M_{\rm ecl,min}$ is the minimum mass of newly formed star-clusters, $M_{\rm ecl,max}{\rm (SFR)}$ is the SFR-dependent maximum mass of newly formed star-clusters, $\xi (m)$ is the IMF and $\xi_{\rm ecl}(M_{\rm ecl})$ is the star-cluster mass function \citep{weidner2005a,weidner2010a}. ."
+ The IGCIME can be parametrized by a imulti-power law.," The IGIMF can be parametrized by a multi-power law,"
+ ↕≻∐≼⊲↖↕≻≼↧⋜↧↑⋜↧,DnC+D data only.
+∪↕∐⋅↖↴↽∕∏∐∖↻↥⋅↕⋯⋜∐⋅∙↖↽↴⋝↸∖⋜∐⊔↸⊳∪↥⋅↥⋅↸∖↸⊳⊓∪∐∐⋜↧↴∖↴ been applied to the radio map restricted to its contour., The primary beam correction has been applied to the radio map restricted to its contour.
+ Such a conservative choice was adopted to avoid any problem towards the otherwise noisy edees of the map., Such a conservative choice was adopted to avoid any problem towards the otherwise noisy edges of the map.
+ Remarkably. Fig.," Remarkably, Fig."
+ 1 shows that C'veuus N-3 is located just at the οσο of au extended radio feature clearly visible towards the South and South-West direction., \ref{cd+d} shows that Cygnus X-3 is located just at the edge of an extended radio feature clearly visible towards the South and South-West direction.
+ Despite some deep negative artifacts. ikelv due to missing short baselines aud perlaps to the subtraction procedure as well. the reality of the extended emission is bevond doubt thanks to confirmation by independent radio instrmucuts (see below).," Despite some deep negative artifacts, likely due to missing short baselines and perhaps to the subtraction procedure as well, the reality of the extended emission is beyond doubt thanks to confirmation by independent radio instruments (see below)."
+" Iu addition. two clumps of extended emission are also present to the North aud East from the microquasar. respectively, but well separated by several arc-minuutes."," In addition, two clumps of extended emission are also present to the North and East from the microquasar, respectively, but well separated by several arc-minutes."
+ Despite the morphological suggestionOO frou Fig. 1..," Despite the morphological suggestion from Fig. \ref{cd+d},"
+ it is not possible to conclude whether or not the exteuded Cluission is plivsicallv associated to Cyenus N-.., it is not possible to conclude whether or not the extended emission is physically associated to Cygnus X-3.
+ While unrelated. extended euussion from the Calactic Plane cannot be ruled out. we could also be seeing traces of a disturbed lobe with au cussion plume extending sidewavs or as long as 5 pe at the Cyguus. A-3 distance.," While unrelated extended emission from the Galactic Plane cannot be ruled out, we could also be seeing traces of a disturbed lobe with an emission plume extending sideways for as long as 5 pc at the Cygnus X-3 distance."
+ Silay scaled-up effects are very often seeu in the large scale lobes and plumes of radio galaxies. where the interaction of the jet pointing towards the observer with the iutergalactic uediun induces turus in the flow direction. changing it Nwoas πιο as 907.," Similar scaled-up effects are very often seen in the large scale lobes and plumes of radio galaxies, where the interaction of the jet pointing towards the observer with the intergalactic medium induces turns in the flow direction changing it by as much as $90^{\circ}$."
+ A representative example of this shenomenon is found. for iustauce. iu the large lobes of 3€ 261 (Laraetal.2001)) where the most remarkable feature is the transition between a well collimated narrow jet at distances from the core below 80 pe. to a conical-shaped wide jet. with a large opening augle.," A representative example of this phenomenon is found, for instance, in the large lobes of 3C 264 \cite{l2004}) ) where the most remarkable feature is the transition between a well collimated narrow jet at distances from the core below 80 pc, to a conical-shaped wide jet, with a large opening angle."
+ The possibility than an analog plivsical xocess could take place iu Cyeuus X-3 is a scenario worth exploring based on this evidence., The possibility than an analog physical process could take place in Cygnus X-3 is a scenario worth exploring based on this evidence.
+ Moreover. as in 3€ 261. we also find iu our map a more distant patch of extended emission right at the expected position anele of the counterjct.," Moreover, as in 3C 264, we also find in our map a more distant patch of extended emission right at the expected position angle of the counterjet."
+ The resulting map obtained after combining the full DuCDD coufiguration data at 6 cinis prescuted in Fig. 2.., The resulting map obtained after combining the full DnC+D+B configuration data at 6 cm is presented in Fig. \ref{cd+d+b}.
+ It is part of a composite image containiug both radio (this work) aux near infrared data (Martíctal.2006)., It is part of a composite image containing both radio (this work) and near infrared data \cite{m2006}) ).
+ This noutage is meluded to render a global aud cuhanced view of the «different source conrponents present iu the complex Cveuus N-3 field (e.g. extended emission. conipact cores. galactic and extragalactic jets. etc.)," This montage is included to render a global and enhanced view of the different source components present in the complex Cygnus X-3 field (e.g. extended emission, compact cores, galactic and extragalactic jets, etc.)"
+ which otherwise would be difücult to visualize simultaneously., which otherwise would be difficult to visualize simultaneously.
+ As quoted above. the Cyeuus. N-3 variable radio core was removed to avoid variability artifacts during the CLEANing process.," As quoted above, the Cygnus X-3 variable radio core was removed to avoid variability artifacts during the CLEANing process."
+ A compact point source. with the average flux density level. was later added later at its location for illustration. purposes.," A compact point source, with the average flux density level, was later added later at its location for illustration purposes."
+ At this point. The svuthesized beam was practically a circular Gaussian with 175 augular size aud the map ruis noise 9.5 ji dv ," At this point, The synthesized beam was practically a circular Gaussian with 5 angular size and the map rms noise 9.5 $\mu$ Jy $^{-1}$."
+Correction for primary beam respouse was also applied to the radio data up to a conservative response value of for appropriate physical mecasurcuicuts., Correction for primary beam response was also applied to the radio data up to a conservative response value of for appropriate physical measurements.
+ As a final cosmetic step. the extended radio euissiou next to Cyenus X-3 aud whose nature is still uuclear has been slightly sioothed for better display.," As a final cosmetic step, the extended radio emission next to Cygnus X-3 and whose nature is still unclear has been slightly smoothed for better display."
+ The radio map in Fig., The radio map in Fig.
+ 2. is also deop enough to reveal a a triple radio source with a compact core and two aligned lobefhot spots componucuts previously uuknown., \ref{cd+d+b} is also deep enough to reveal a a triple radio source with a compact core and two aligned lobe/hot spots components previously unknown.
+" The location of the proposed core ds at coordinates oqogou=207327 3872940802 and qoos=|10755/317 025072, with a 6 cin flux density of 0.27—0.06 mJy."," The location of the proposed core is at coordinates $\alpha_{\rm J2000.0}= 20^h 32^m $ $\pm$ 02 and $\delta_{\rm J2000.0}= +40^{\circ} 55^{\prime}$ $\pm$ 2, with a 6 cm flux density of $0.27 \pm 0.06$ mJy."
+ Its aneular distance your Cvenus X-3 is 3/06., Its angular distance from Cygnus X-3 is 06.
+ The overall morphology of this triple source cousistent with a FR II radio galaxy. as shown iu Fie. 3. ," The overall morphology of this triple source consistent with a FR II radio galaxy, as shown in Fig. \ref{triple}. ."
+The erey scale in Fie., The grey scale in Fig.
+ 3. comes from previous near mfrared JVs- observations (Martíetal. 2006)).where a point- counterpart with magnitude As17.240.1 is coincident with the compac core Within astrometric error.," \ref{triple} comes from previous near infrared $Ks$ -band observations \cite{m2006}) ),where a point-like counterpart with magnitude $Ks=17.2 \pm 0.1$ is coincident with the compact core within astrometric error."
+The underlving mechanisms which govern the large range in low-frequency racio Luminosity observed. in. racdio-LIoud. active galactic nuclei (AGN) are currently unknown.,The underlying mechanisms which govern the large range in low-frequency radio luminosity observed in radio-loud active galactic nuclei (AGN) are currently unknown.
+ Over, Over
+The cemography of local galaxies suggests that most galaxies host a quiescent supermassive black hole (SMBDBLL) at the present time and the properties of the black hole are correlated. with those of the host spheroid.,The demography of local galaxies suggests that most galaxies host a quiescent supermassive black hole (SMBH) at the present time and the properties of the black hole are correlated with those of the host spheroid.
+ In. particular. recent observational evidence points to the existence of a tight correlation between the mass of the central black hole and the velocity. dispersion of the host spheroic (2???) in nearby galaxies.," In particular, recent observational evidence points to the existence of a tight correlation between the mass of the central black hole and the velocity dispersion of the host spheroid \citep{tremaine02,ferrarese2000,gebhardt2000b} in nearby galaxies."
+ This correlation strongly suggests. coeval erowth of the black hole and the stellar component via likely regulaion of the gas supply in galactic nuclei (72777)," This correlation strongly suggests coeval growth of the black hole and the stellar component via likely regulation of the gas supply in galactic nuclei \citep{silk98,kauffmann00,king03,thompson05}."
+ 1)ack hole growth is believed. to be powered. by σας aceretion (2) and accreting black holes are detected. as opticalv bright quasars., Black hole growth is believed to be powered by gas accretion \citep{lyndenbell69} and accreting black holes are detected as optically bright quasars.
+ These optically bright quasars appear to exist out to the highest. redshifts probed at the present time., These optically bright quasars appear to exist out to the highest redshifts probed at the present time.
+ Therefore. the mass build-up of SMDlIS is likely o have commenced at extremely high redshifts (2> 10).," Therefore, the mass build-up of SMBHs is likely to have commenced at extremely high redshifts $z > 10$ )."
+ In fact. optically bright quasars have now been detected al 2>6 (e.g. ??)) in the Sloan Digital Sky Survey (SDSS).," In fact, optically bright quasars have now been detected at $z > 6$ (e.g., \citealt{fan04,fan06}) ) in the Sloan Digital Sky Survey (SDSS)."
+ llosts of high redshift quasars are often strong sources of dust emission (????).. suggesting that quasars were already in place in massive galaxies at a time when galaxies were undergoing vigourous star formation.," Hosts of high redshift quasars are often strong sources of dust emission \citep{omont01,cox02,carilli02,walter03}, suggesting that quasars were already in place in massive galaxies at a time when galaxies were undergoing vigourous star formation."
+ The growth spurts of SALBLIs are also seen in the X-ray waveband., The growth spurts of SMBHs are also seen in the X-ray waveband.
+ The integratec emission. from. these X-ray quasars generates the cosmic X-rav background (XRB). and its spectrum suggests tha most. black-hole erowth is obscured in optical wavelengths (22%," The integrated emission from these X-ray quasars generates the cosmic X-ray background (XRB), and its spectrum suggests that most black-hole growth is obscured in optical wavelengths \citep{fabian99,mushot00,hasinger01,barger03,barger05,worsley05}."
+" There exist. examples of obscured black-hole erowth in the form of ""Evpe-2. quasars. but their detectec numbers are fewer than expected from models of the NIU."," There exist examples of obscured black-hole growth in the form of `Type-2' quasars, but their detected numbers are fewer than expected from models of the XRB."
+ However. there is recent tantalizing evidence from infra-rec (LR) studies that clust-obscurecl accretion is ubiquitous (?)..," However, there is recent tantalizing evidence from infra-red (IR) studies that dust-obscured accretion is ubiquitous \citep{martinez05}."
+ Current work suggests that while SALBIIs might spend mos of their lifetime in an optically dim phase. the bulk of mass erowth occurs in the short-lived quasar stages.," Current work suggests that while SMBHs might spend most of their lifetime in an optically dim phase, the bulk of mass growth occurs in the short-lived quasar stages."
+ The assembly of BIL mass in the Universe. has been tracked using optical quasar activity., The assembly of BH mass in the Universe has been tracked using optical quasar activity.
+ “Phe current, The current
+not bias in mass ancl concentration.,not bias in mass and concentration.
+ We discuss the results of these authors in the context of our findings in Section 6.., We discuss the results of these authors in the context of our findings in Section \ref{sec:conclusion}.
+ For cach cluster we produce surface mass density maps bv projecting the 105. MMpe. box along one side and interpolating to a regular grid of 20007 pixels using a triangular shaped clouds (TSC) algorithm (2)...," For each cluster we produce surface mass density maps by projecting the $10\,h^{-1}$ Mpc box along one side and interpolating to a regular grid of $^2$ pixels using a triangular shaped clouds (TSC) algorithm \citep{Hockney_Eastwood_1988}."
+ Ht was checked that other interpolation algorithms. such as cloud-in-cells or smoothed. particle hvdrodynamics. (SPILL) interpolation. cid not vield significantly dilferent results.," It was checked that other interpolation algorithms, such as cloud-in-cells or smoothed particle hydrodynamics (SPH) interpolation, did not yield significantly different results."
+ At the adopted resolution. 1 pixel length corresponds to 5.! kpe. equivalent to the gravitational softening length of the AIS. which makes it the highest sensible resolution.," At the adopted resolution, 1 pixel length corresponds to $5\,h^{-1}$ kpc, equivalent to the gravitational softening length of the MS, which makes it the highest sensible resolution."
+ To further boost our statistics. ancl explore the role of viewing angle for individual objects. we repeated: this procedure for each. simulated: cluster five times. cach time selecting a random orientation. and thus. produce. five different maps/ for each. cluster.," To further boost our statistics, and explore the role of viewing angle for individual objects, we repeated this procedure for each simulated cluster five times, each time selecting a random orientation and thus produce five different maps for each cluster."
+ In the DIES. observations will be mace of ~20.000clusters with masses in excess of 2.JOMAL..," In the DES, observations will be made of $\sim20,000$clusters with masses in excess of $2\times 10^{14} M_{\odot}$."
+ The total number of projections we produce is comparable to this survey., The total number of projections we produce is comparable to this survey.
+ For the most massive clusters (with. Moog210:077 A7.) in our sample. we have electec to boost our statistics even further by. producing maps corresponding to 50 dilferent random projections.," For the most massive clusters (with $M_{200} \ge 10^{14.8} M_\odot$ ) in our sample, we have elected to boost our statistics even further by producing maps corresponding to 50 different random projections."
+ The logic behind this is that i) such massive svstems are the mos commonly-tareetted in. present WL observational. studies (and therefore are the systems for which we can presently make comparisons between theory ancl observations): ane iij) very high mass (LOM AL.)clusters are rare even in a 5005+ Alpe volume., The logic behind this is that i) such massive systems are the most commonly-targetted in present WL observational studies (and therefore are the systems for which we can presently make comparisons between theory and observations); and ii) very high mass $\sim 10^{15} M_\odot$ )clusters are rare even in a $500 h^{-1}$ Mpc volume.
+ Convergence (8) maps are then computed according to with “oy as above. and the shear maps formed according to equations and (2).," Convergence $\kappa$ ) maps are then computed according to with $\Sigma_{\rm cr}$ as above, and the shear maps formed according to equations and ."
+". To solve these equations. we adopt the same EET method as ?.. computing 2 and α are Fourier transforms of 7, and 6. and A are the whereappropriate wave vectors."," To solve these equations, we adopt the same FFT method as \citet{Clowe_et_al_2004}, computing where $\tilde{\gamma}$ and $\tilde{\kappa}$ are Fourier transforms of $\gamma$ and $\kappa$, and $\hat{k}$ are the appropriate wave vectors."
+ As an example. Fig.," As an example, Fig."
+ 1 shows three projections of one high-mass cluster with the corresponding shear fields as tick marks., \ref{projections} shows three projections of one high-mass cluster with the corresponding shear fields as tick marks.
+ Ht is apparent that dillerent. projections of the same cluster can look very dillerent and lead to clillerent shear signals., It is apparent that different projections of the same cluster can look very different and lead to different shear signals.
+ With shear maps for all the projections generated. an arrav of background galaxies. extending bv default to a projected distance of ra=15! from the cluster. centre. was then simulated for each individual projection.," With shear maps for all the projections generated, an array of background galaxies, extending by default to a projected distance of $r_\text{out} = 15^\prime$ from the cluster centre, was then simulated for each individual projection."
+ This limit in projected. distance excludes potential numerical artifacts near the edges of the shear field: it also matches the typical lield-of-view. of observations (e.g. see also 2)). where for the most massive clusters the shear signal is detected out to roughly this radius.," This limit in projected distance excludes potential numerical artifacts near the edges of the shear field; it also matches the typical field-of-view of observations (e.g. \citealt{Clowe_Schneider_2001}; see also \citealt{Hoekstra_2003}) ), where for the most massive clusters the shear signal is detected out to roughly this radius."
+ Although galaxies usedin real weak lensing surveys have a range of redshifts. we consider only the simplified case of a uniform: source redshift of z=1.07.," Although galaxies usedin real weak lensing surveys have a range of redshifts, we consider only the simplified case of a uniform source redshift of $z = 1.0$."
+ This is a value typical for observational studies such as CEIITLS (7) and COSMOS (7). and has also been used in other theoretical work concerning the accuracy of weak lensing (e.g.. ?:: 2: 7: 7)).," This is a value typical for observational studies such as CFHTLS \citep{Ilbert_et_al_2006} and COSMOS \citep{Ilbert_et_al_2009} and has also been used in other theoretical work concerning the accuracy of weak lensing (e.g., \citealt{Clowe_et_al_2004}; \citealt{Oguri_Hamana_2011}; \citealt{Gruen_et_al_2011}; \citealt{Becker_Kravtsov_2011}) )."
+ In real surveys the dispersion in source redshifts and its finite sampling (7) has to be considered. as an additional error source., In real surveys the dispersion in source redshifts and its finite sampling \citep{Hoekstra_et_al_2011} has to be considered as an additional error source.
+ We note that for this choice of lens and. source redshift. the critical surface mass density X4=3.33:102AJ. Alpe?.," We note that for this choice of lens and source redshift, the critical surface mass density $\Sigma_\text{cr} = 3.33 \cdot 10^{15} M_\odot / $ $^2$."
+" To exclude the strong lensing regime in the cluster centre. we excised. the central region. within +rin. With a default inner limit racdius 7j,=30 aresee. a value tvpica of the extent of strongly. lensed ares in observational stuclics (e.g... 2))."," To exclude the strong lensing regime in the cluster centre, we excised the central region within $r = r_\text{in}$ , with a default inner limit radius $r_\text{in} = 30$ arcsec, a value typical of the extent of strongly lensed arcs in observational studies (e.g., \citealt{Hennawi_et_al_2008}) )."
+ We consider the effect of changing the values of both rj and ra in appendix D3: while the exact amount of bias and scatter depends somewhat on the exact choice of ra and rua (see also 2)) our overall conclusions are unalfectec by their exact choice., We consider the effect of changing the values of both $r_\text{in}$ and $r_\text{out}$ in appendix \ref{sec:cuttest}; while the exact amount of bias and scatter depends somewhat on the exact choice of $r_\text{in}$ and $r_\text{out}$ (see also \citealt{Becker_Kravtsov_2011}) ) our overall conclusions are unaffected by their exact choice.
+ The background. galaxy ellipticities. ον were drawn randomly from ao Gaussian. distribution. with stanclare deviation σ=0.2 per component as in ?: note that this choice of σ is also similar to that used by ?.., The background galaxy ellipticities $\epsilon_s$ were drawn randomly from a Gaussian distribution with standard deviation $\sigma = 0.2$ per component as in \citet{Geiger1998}; note that this choice of $\sigma$ is also similar to that used by \citet{Hoekstra_et_al_2011}.
+ The elect of increasing the value of σ is shown in appendix D2: it docs not alfect our results significantly., The effect of increasing the value of $\sigma$ is shown in appendix \ref{sec:sigmatest}; ; it does not affect our results significantly.
+ Phe positions of galaxies were allocated randomly. with average unlensed density [ron which shapes can be measured accurately n=30 ? and accounting for Poisson noise.," The positions of galaxies were allocated randomly, with average unlensed density from which shapes can be measured accurately $n = 30$ $^{-2}$, and accounting for Poisson noise."
+ Εις is comparable to background. densities achievable with current &rouncd-based pointed observations of galaxy clusters (e.g. 2)) and similar to what will be achieved. with upcoming large surveys such as the and that. from the Large. Svnoptie Survey Telescopeος).," This is comparable to background densities achievable with current ground-based pointed observations of galaxy clusters (e.g., \citealt{Dahle_2006}) ) and similar to what will be achieved with upcoming large surveys such as the and that from the Large Synoptic Survey Telescope."
+ The distortion of these simulated background. galaxies by the cluster-inducecl shear field was then computed according to to. vield. the lensecl background: galaxy ellipticities ο., The distortion of these simulated background galaxies by the cluster-induced shear field was then computed according to to yield the lensed background galaxy ellipticities $\epsilon$ .
+ To account. for number count. depletion as described. in Section 2.. the lensecl source density was adjusted: locally according to equations and(πι.," To account for number count depletion as described in Section \ref{sec-lensingtheory}, the lensed source density was adjusted locally according to equations and."
+ At this stage we have the final svnthetie catalogues of LIensed ealaxies., At this stage we have the final synthetic catalogues of lensed galaxies.
+ We derive the best-fit mass. Mago. ane concentration. c. by fitting projected NEW mass distributions (7: ?:: 2)) toour mock WL data.," We derive the best-fit mass, $M_{200}$ , and concentration, $c$ , by fitting projected NFW mass distributions \citealt{Navarro_Frenk_White_1995}; ; \citealt{Navarro_Frenk_White_1996}; \citealt{Navarro_Frenk_White_1997}) ) toour mock WL data."
+ The NEW profile. which has been shown to reproduce the 3D sphericallv-averaged mass density profilesof dark matter haloes spanning a wide range of masses. is described by thetwo parameterform:," The NFW profile, which has been shown to reproduce the 3D spherically-averaged mass density profilesof dark matter haloes spanning a wide range of masses, is described by thetwo parameterform:"
+Large Magellanie Cloud (the coverage is shown in Fig. 1).,Large Magellanic Cloud (the coverage is shown in Fig. \ref{selection}) ).
+ Their johoetometrie catalogue contains 24 million objects. so that each of he individually analysed. fields contains some 107 stars.," Their photometric catalogue contains 24 million objects, so that each of the individually analysed fields contains some $10^4$ stars."
+ For the synthetic colour magnitude diagrams the isochrones of the Padova group were used (2).., For the synthetic colour magnitude diagrams the isochrones of the Padova group were used \citep{girardi-etal2002}.
+ The temporal resolution of the star formation story is given by age bins of 0.3 dex size for ages younger than 100 Myr and bins of 0.2 dex for older ages., The temporal resolution of the star formation history is given by age bins of 0.3 dex size for ages younger than 100 Myr and bins of 0.2 dex for older ages.
+ Solutions for the star ormation history were obtained for four metallicities. Z=0.001. Z=0.0025 (interpolated and used only for ages larger than 100 Tyr). Z=0.004 and Z=0.008.," Solutions for the star formation history were obtained for four metallicities, $Z=0.001$, $Z=0.0025$ (interpolated and used only for ages larger than 100 Myr), $Z=0.004$ and $Z=0.008$."
+ As their photometry does not reach the main sequence turnoff yoint for the old population. the extraction of the early star ormation history was difficult., As their photometry does not reach the main sequence turnoff point for the old population the extraction of the early star formation history was difficult.
+ Therefore ? restricted StarFISH to tit only a single age bin covering all ages older than 4 Gyr in the bar region., Therefore \citet{harris+zaritsky2009} restricted StarFISH to fit only a single age bin covering all ages older than 4 Gyr in the bar region.
+ Within the bar region they used the typical star formation ustory from solutions for the star formation history derived using HST data (2.. 2 and ?.. which widely agree with each other).," Within the bar region they used the typical star formation history from solutions for the star formation history derived using HST data \citealp{olsen1999}, \citealp{holtzman-etal1999} and \citealp{smeckerhane-etal2002}, which widely agree with each other)."
+ Figure 7 shows the star formation history. derived by ?..," Figure \ref{sfh-hz} shows the star formation history, derived by \citet{harris+zaritsky2009}."
+ The star formation history of the fields within the regions of ?. (open circles in Fig. 13) , The star formation history of the fields within the regions of \citet{massey2002} (open circles in Fig. \ref{selection}) )
+in which star clusters have been observed is shown as the solid line. with the grey area being its uncertainty.," in which star clusters have been observed is shown as the solid line, with the grey area being its uncertainty."
+ This partial star formation history follows the total star formation history (for the whole area of the Large Magellanic Cloud. dotted line) at about half the star formation rate.," This partial star formation history follows the total star formation history (for the whole area of the Large Magellanic Cloud, dotted line) at about half the star formation rate."
+ In Figure 8 we summarise the solutions for the star formation history of the Large Magellanic Cloud derived from. colour-magnitude diagrams. the most massive star cluster or the total star cluster mass.," In Figure \ref{sfh-comparison} we summarise the solutions for the star formation history of the Large Magellanic Cloud derived from colour-magnitude diagrams, the most massive star cluster or the total star cluster mass."
+ As before in Fig., As before in Fig.
+ 7. the dashed line within the grey shaded area gives the star formation history with uncertainty derived from colour-magnitude diagrams within the ?. fields., \ref{sfh-hz} the dashed line within the grey shaded area gives the star formation history with uncertainty derived from colour-magnitude diagrams within the \citet{massey2002} fields.
+ The thick solid lines give the upper and lower limits for the star formation history derived from the most massive star clusters (Sec. ??) , The thick solid lines give the upper and lower limits for the star formation history derived from the most massive star clusters (Sec. \ref{sec_mostmassive}) )
+and the dotted line gives the amount of star formation in star clusters (Sec. 22)., and the dotted line gives the amount of star formation in star clusters (Sec. \ref{sec_total}) ).
+" The comparison of the AZ, estimate for the star formation history with the result from the CMD shows good agreement for ages up to zz| Gyr. choosing the AZ,4, solution which includes R136."," The comparison of the $M_\mathrm{max}$ estimate for the star formation history with the result from the CMD shows good agreement for ages up to $\approx 1$ Gyr, choosing the $M_\mathrm{max}$ solution which includes R136."
+ The absolute value for the star formation rate. derived from the most massive clusters. is at the level of the CMD star formation rate within the ? fields.," The absolute value for the star formation rate, derived from the most massive clusters, is at the level of the CMD star formation rate within the \citet{massey2002} fields."
+ However. as the most massive method is intended to give the star formation rategalaxy. and the spatial coverage of the star clusters contains most of the area recently active in star formation. it is perhaps more appropriate to compare to the galaxy-wide star formation rate.," However, as the most massive cluster-method is intended to give the star formation rate, and the spatial coverage of the star clusters contains most of the area recently active in star formation, it is perhaps more appropriate to compare to the galaxy-wide star formation rate."
+ In this ease the star formation rate would be underestimated by a factor of 2 (compared to the dotted line in Fig. 7)), In this case the star formation rate would be underestimated by a factor of 2 (compared to the dotted line in Fig. \ref{sfh-hz}) )
+ and the normalisation in eq., and the normalisation in eq.
+ 5 would need adjusting., \ref{sfrdeltat} would need adjusting.
+ As long as the spatial coverage is incomplete it is unfortunately impossible to disentangle inappropriate normalisation and effects of spatial incompleteness., As long as the spatial coverage is incomplete it is unfortunately impossible to disentangle inappropriate normalisation and effects of spatial incompleteness.
+ Compared to the CMD solution the most massive cluster solution shows an offset of the peaks in star formation., Compared to the CMD solution the most massive cluster solution shows an offset of the peaks in star formation.
+ This is perhaps accounted for by the moving averaging window. although different isochrone sets could also account for this.," This is perhaps accounted for by the moving averaging window, although different isochrone sets could also account for this."
+ If the ages of the star clusters could be more accurately determined the moving window could be reduced. by at the same time including not only the most massive star cluster but also the second. third ete.," If the ages of the star clusters could be more accurately determined the moving window could be reduced, by at the same time including not only the most massive star cluster but also the second, third etc."
+ most massive and thus keeping the sample of data points used large enough for the statistically necessary averaging., most massive and thus keeping the sample of data points used large enough for the statistically necessary averaging.
+ An investigation in this direction is bevond the scope of this work., An investigation in this direction is beyond the scope of this work.
+" We turn now to a the comparison of the star formation rate derived from the total mass with the CMD and Ad, solution.", We turn now to a the comparison of the star formation rate derived from the total mass with the CMD and $M_\mathrm{tot}$ solution.
+ For the very youngest ages (<< 10-20 Myr) the Au star formation history shows an increase of the star formation rate with age. in contrast to the of the CMD solution.," For the very youngest ages $< 10$ $20$ Myr) the $M_\mathrm{tot}$ star formation history shows an increase of the star formation rate with age, in contrast to the of the CMD solution."
+ This is because very young clusters are in a class of objects not selected by and thus the used data set is incomplete for very young ages., This is because very young clusters are in a class of objects not selected by \citet{hunter-etal2003} and thus the used data set is incomplete for very young ages.
+" Within the age range where we can assume that the A44, solution is based on a complete data set (both in object selection and ower mass completeness). from = 10-20 Myr to zz 200 Myr. the Mo solution shows the same structure as the CMD and Minas solution. but at a lower star formation rate."," Within the age range where we can assume that the $M_\mathrm{tot}$ solution is based on a complete data set (both in object selection and lower mass completeness), from $\approx 10$ $20$ Myr to $\approx$ 200 Myr, the $M_\mathrm{tot}$ solution shows the same structure as the CMD and $M_\mathrm{max}$ solution, but at a lower star formation rate."
+ The fraction of star ormation in star clusters. i.e. the ratio between the CMD curve and the Adj curve. appears for the whole age range to be at a level.," The fraction of star formation in star clusters, i.e. the ratio between the CMD curve and the $M_\mathrm{tot}$ curve, appears for the whole age range to be at a level."
+ This is caused by either the formation of only a Traction of stars in star clusters. or. when assuming that all stars orm in eusters. by the dissolution of star clusters caused by the oesumably violent transition from their gas-embedded to the gas-ree state.," This is caused by either the formation of only a fraction of stars in star clusters, or, when assuming that all stars form in custers, by the dissolution of star clusters caused by the presumably violent transition from their gas-embedded to the gas-free state."
+ The second explanation could in principle be detected in the comparison. it should lead to a Mq star formation rate hat is identical to the CMD star formation rate at very young ages (the embedded cluster phase).," The second explanation could in principle be detected in the comparison, it should lead to a $M_\mathrm{tot}$ star formation rate that is identical to the CMD star formation rate at very young ages (the embedded cluster phase)."
+ However. as the data set is not complete at the youngest ages a distinction between the alternative explanations is not possible.," However, as the data set is not complete at the youngest ages a distinction between the alternative explanations is not possible."
+ For ages older than | Gyr the number of detected star clusters is very small. and their distribution shows many gaps.," For ages older than 1 Gyr the number of detected star clusters is very small, and their distribution shows many gaps."
+ Thus no large amount of information about the star formation history can be extracted from the star clusters., Thus no large amount of information about the star formation history can be extracted from the star clusters.
+ The Mj method gives an upper and lower limit for the star formation history. either assuming no star cluster formation in the gaps (lower bold line in Fig. 8» ," The $M_\mathrm{max}$ method gives an upper and lower limit for the star formation history, either assuming no star cluster formation in the gaps (lower bold line in Fig. \ref{sfh-comparison}) )"
+or assuming the mass of the detection limit as the upper limit for the mass of a cluster that could have formed in an gap (upper line)., or assuming the mass of the detection limit as the upper limit for the mass of a cluster that could have formed in an gap (upper line).
+ As the mass associated with the detection limit increases with increasing age. the upper limit for the star formation history shows also an increase of the star formation rate with age.," As the mass associated with the detection limit increases with increasing age, the upper limit for the star formation history shows also an increase of the star formation rate with age."
+ The, The
+ , 
+"fluxes and the high energy cutoff, depend strongly on {ν.","fluxes and the high energy cutoff, depend strongly on $t_v$."
+" If the mechanism responsible of the X-ray flares is the same as the prompt emission (IS) one could expect a similar variability time, thus we initially assume for flares that 4,~ 100 ms."," If the mechanism responsible of the X-ray flares is the same as the prompt emission (IS) one could expect a similar variability time, thus we initially assume for flares that $t_v \sim$ 100 ms."
+" However, given our ignorance on the average interval between consecutive shells when the engine is reactivated, that these shells could be emitted close to the central engine or at larger distances, and that the Lorentz factor Γ could be smaller for X-ray flares in comparison with the prompt emission, in the following we also discuss the case of a larger t,."," However, given our ignorance on the average interval between consecutive shells when the engine is reactivated, that these shells could be emitted close to the central engine or at larger distances, and that the Lorentz factor $\Gamma$ could be smaller for X-ray flares in comparison with the prompt emission, in the following we also discuss the case of a larger $t_v$."
+ The other relevant quantities are the luminosity L of the flare and the average flare duration ty., The other relevant quantities are the luminosity $L$ of the flare and the average flare duration $t_f$.
+" In order to determine the typical t; and L of the flares, we select a sample of X-ray flares with known redshift and published light curves, taking place at times shorter than ~1000 s, when the IS mechanism can still be active."," In order to determine the typical $t_f$ and $L$ of the flares, we select a sample of X-ray flares with known redshift and published light curves, taking place at times shorter than $\approx 1000$ s, when the IS mechanism can still be active."
+" At longer times other mechanisms can be responsible for the flare emission, e.g. inhomogeneities in the external medium To estimate the high energy counterpart of X-ray flares, we repeat the calculations made in Sect."," At longer times other mechanisms can be responsible for the flare emission, e.g. inhomogeneities in the external medium To estimate the high energy counterpart of X-ray flares, we repeat the calculations made in Sect."
+ 1. by changing only the luminosity from L=10°?ergs! to L=2x10?ergs!., \ref{prompt} by changing only the luminosity from $L=10^{52}~erg s^{-1}$ to $L=2 \times 10^{49}~erg~s^{-1}$.
+ As noted by Fanetal.(2007) it is difficult to predict the expected SSC emission as we do not have a good estimate of the typical Lorentz factor of the electrons accelerated at the shock to produce the X-rays., As noted by \citet{fan07} it is difficult to predict the expected SSC emission as we do not have a good estimate of the typical Lorentz factor of the electrons accelerated at the shock to produce the X-rays.
+" However we can give some constraints on the Lorentz factor I"" of the colliding shells.", However we can give some constraints on the Lorentz factor $\Gamma$ of the colliding shells.
+" If we assume that the X-ray flares are produced by accelerated electrons which cool very fast, the condition v;«v,, introduces an upper limit to the fireball Lorentz factor I."," If we assume that the X-ray flares are produced by accelerated electrons which cool very fast, the condition $\nu_c < \nu_m$ introduces an upper limit to the fireball Lorentz factor $\Gamma$ ."
+" In particular, for typical flare luminosity and temporal variability {ν--100 ms, and for e;=0.45, eg=0.1 and p=2.5, we find Γx340."," In particular, for typical flare luminosity and temporal variability $t_v$ =100 ms, and for $\epsilon_e=0.45$, $\epsilon_B=0.1$ and $p=2.5$, we find $\Gamma \lesssim 340$."
+" In addition, in order to have SSC GeV flare emission we have to require that the spectral cutoff energy due topair production is"," In addition, in order to have SSC GeV flare emission we have to require that the spectral cutoff energy due topair production is"
+"5Η MIT. respectively,","341 MHz, respectively."
+ The observations at the high requencies using the Effelsbere aud the Lovell telescope rave been described in I&karastereiou ct al. (, The observations at the high frequencies using the Effelsberg and the Lovell telescope have been described in Karastergiou et al. (
+2001. 2002).,"2001, 2002)."
+ The bandwidths used at these frequencies were 500 NUIT (Effelsbere) aud 32 MIIz (Lovell)., The bandwidths used at these frequencies were 500 MHz (Effelsberg) and 32 MHz (Lovell).
+ At the CAIRT we used vandwidths of 16 MIIz at each frequency., At the GMRT we used bandwidths of 16 MHz at each frequency.
+ The ceutral requenucies for the lowest receiving band differed for he two sources observed. as these signals were recorded w the same data logeine svsteni (ie. signals frou voth frequencies were added together after square-law detection. aud the pulsars dispersion delay was used to separate the two data streams during off-line analysis).," The central frequencies for the lowest receiving band differed for the two sources observed, as these signals were recorded by the same data logging system (i.e. signals from both frequencies were added together after square-law detection, and the pulsar's dispersion delay was used to separate the two data streams during off-line analysis)."
+ The lower observing frequency was chosen such that the dispersion delay relative to the 626 MITz resulted in well separated pulses., The lower observing frequency was chosen such that the dispersion delay relative to the 626 MHz resulted in well separated pulses.
+ Data from all telescopes were stored for off-line processing., Data from all telescopes were stored for off-line processing.
+ Using observing parameters sunuuarized in Table 1.. we observed 2513 pulses (30 mun) of PSR D0329151 and [375 pulses (86 min) of PSR D1133]|16.," Using observing parameters summarized in Table \ref{obstab}, we observed 2513 pulses (30 min) of PSR B0329+54 and 4375 pulses (86 min) of PSR B1133+16."
+ All data elescopes were converted to a ΟΜΙΛΟΙ EPN format and time-aliened following the procedures detailed in INarastereiou et al. (, All data telescopes were converted to a common EPN format and time-aligned following the procedures detailed in Karastergiou et al. (
+2001).,2001).
+ The effective resolutions of our ime series were 711.0 ys for PSR Bo329|5L aud 1159.85 yes or PSR B1133|16. respectively.," The effective resolutions of our time series were 711.0 $\mu$ s for PSR B0329+54 and 1159.8 $\mu$ s for PSR B1133+16, respectively."
+ Some mecasured pulses. in particulary at the low frequencies of the GAIRT. were affected by radio interfercuce.," Some measured pulses, in particular at the low frequencies of the GMRT, were affected by radio interference."
+ Such pulses. 251 for PSR D0329|5[ and 11 for PSR B1133|16. were excluded from he subsequent analysis.," Such pulses, 251 for PSR B0329+54 and 41 for PSR B1133+16, were excluded from the subsequent analysis."
+" Tn Effelsbere and Jodrell Bank. the pulses were compared to a signal of a calibrated noise diode. which itself was compared to the strength of known fiux calibrators observed durug poiuting observations before. after and dunus the observations. respectively,"," In Effelsberg and Jodrell Bank the pulses were compared to a signal of a calibrated noise diode, which itself was compared to the strength of known flux calibrators observed during pointing observations before, after and during the observations, respectively."
+ Flux densities of these point sources. such as NGCTO27 and 3CLs. were obtained from Pene et al.," Flux densities of these point sources, such as NGC7027 and 3C48, were obtained from Peng et al."
+ and Ott ct al. (1991).," \nocite{pkkw00} and Ott et al. \nocite{owq+94},"
+. resulting in an estimated uncertaintv iu the lucasurements of about., resulting in an estimated uncertainty in the measurements of about.
+. At the CAIRT. known flux calibrators such as BCLIF. 30295 aud 3C286 were observed before and after every pulsar observation.," At the GMRT, known flux calibrators – such as 3C147, 3C295 and 3C286 – were observed before and after every pulsar observation."
+ Unlike the pulsar observations. the calibration sources were observed separately at two frequencies. ie. one frequency at a time. using the suunied signal frou only the auteuuas at the selected frequency;," Unlike the pulsar observations, the calibration sources were observed separately at two frequencies, i.e. one frequency at a time, using the summed signal from only the antennas at the selected frequency."
+ From these observations. au effective. calibration scale was computed for cach of the two frequencies.," From these observations, an effective calibration scale was computed for each of the two frequencies."
+" Using the nearest calibration source observation. this scale was applied to the ""onu-off deflection of the pulsu. separately for cach frequency of observation. i order to flux calibrate the pulsar signal at both frequencies."," Using the nearest calibration source observation, this scale was applied to the “on-off” deflection of the pulsar, separately for each frequency of observation, in order to flux calibrate the pulsar signal at both frequencies."
+ As we will demonstrate later. the corresponding flux densities of the average profiles are consistent with what is known from the literature. eiving us coufideuce in the accuracy of our sinele pulse measurements.," As we will demonstrate later, the corresponding flux densities of the average profiles are consistent with what is known from the literature, giving us confidence in the accuracy of our single pulse measurements."
+ Before we can use flux density data for a computation of the radio spectrün. we must studyv the possible imupact of ISS.," Before we can use flux density data for a computation of the radio spectrum, we must study the possible impact of ISS."
+ Iu particular. we have to estimate the iuflueuce of diffractive scintillation which occurs on timescales of minutes," In particular, we have to estimate the influence of diffractive scintillation which occurs on timescales of minutes."
+ Up to frequencies of about 1 GIIz. the sources are usually in the strong scintillation regime. aud the de- bandwidth scales as (e.g. Cordes Rickett where s;=l.l for à Rkohuogorov turbulence ↴∖↴⋉∖↸⊳⊓↴," Up to frequencies of about 1 GHz, the sources are usually in the strong scintillation regime, and the de-correlation bandwidth scales as (e.g. Cordes Rickett \nocite{cr98} where $\beta=4.4$ for a Kolmogorov turbulence spectrum."
+"∐⊔∙ The modulation timescale.iniescale. Afpysss isds expected to increase with frequencyfiCV as ∆∤↕↔↔∖⊼∕↓⊳−⋖↸∖∙∶↴∙⊾∙≼∶∏↻↑⋜↧⊥≝⋟⋂⋅↱↗∙≼⊲∪↥⋅≺∐∖↴∖↴∙∖↽↕⊰↕↸⊳↨↘↽↸∖↑↑↓∩∩≺∖⋟∙D , : ↽∖ ↕"," The modulation timescale, $\Delta t_{\rm ISS}$ is expected to increase with frequency as $\Delta t_{\rm ISS} \propto \nu^{1.2}$ (e.g. Gupta 1995, Cordes Rickett 1998)."
+↕≯∐↸∖↴∖↴↸⊳↕∐↑↕∐⋜↧↑↕∪∐↴∖↴↕∐↑↕⋯↸∖⋜⋯≼↧↕∐↕⋟↥⋅↸∖≺⋯∖∐↸⊳↖↽⋜∐⋅↸∖ | ↴∖↴↿∏−⊔↸⊳↕↸∖∐↑↕⋅↖⇁↖↖↽↸∖∐↴∖↴⋜⋯∏≻↕↸∖≼↧↑↕∐∖∐↑∐↸∖⋯∪≼↧∏↕⋜↧↑↕∪∐↕∐≼∐∖⊼∙∣⊔∙ defined as is unity in the stroug (diffractive) scintillation regime if Sis he measured fiux density aud «52 its mean value.," If the scintillations in time and in frequency are sufficiently well sampled then the modulation index, $m$ , defined as is unity in the strong (diffractive) scintillation regime if $S$ is the measured flux density and $<\!S\!>$ its mean value."
+ However. if the scintillation αποποτα is siguificautlv less than the observing bandwidth. the sciutles will wash out’ and the measured modulation iudex will be wach lower.," However, if the scintillation bandwidth is significantly less than the observing bandwidth, the scintles will `wash out' and the measured modulation index will be much lower."
+ Iu the weak scintillation regime. modulation decreases while timescales are expected to decrease with frequency as well Atha;xÓÉU7.," In the weak scintillation regime, modulation decreases while timescales are expected to decrease with frequency as well, $\Delta
+t_{\rm ISS} \propto \nu^{-0.5}$."
+ For our sources. scintillation parameters have been measured by Malofeev et al. (," For our sources, scintillation parameters have been measured by Malofeev et al. ("
+1996) and both sources are fouud to be in the weals scintillation rogue at [850 MITIz. with modulation mdices measured by Malofeev. et al.,"1996) \nocite{mss+96} and both sources are found to be in the weak scintillation regime at 4850 MHz, with modulation indices measured by Malofeev et al."
+ of only an~0.12 for D032915L aud m)c0.22 for PSR D1132116., of only $m\sim0.12$ for B0329+54 and $m\sim0.22$ for PSR B1133+16.
+ Timescales expected for weak scintillation from the ratio of Fresnel scale aud velocity for both our pulsus are listed in Table 1.., Timescales expected for weak scintillation from the ratio of Fresnel scale and velocity for both our pulsars are listed in Table \ref{obstab}.
+" Pulsar BO329|5[s critical frequency. v,.. marking the ransition frou the strong to weak scintillation regine. is estimated to be around 3 GIIz. cousistent with the results w Malofeev et al. ("," Pulsar B0329+54's critical frequency, $\nu_c$, marking the transition from the strong to weak scintillation regime, is estimated to be around 3 GHz, consistent with the results by Malofeev et al. ("
+1996).,1996).
+ This places PSR Bo32915Ll in he strone ISS reguue at all our frequencies below £850 Miz., This places PSR B0329+54 in the strong ISS regime at all our frequencies below 4850 MHz.
+ For these frequencies. we estimate the expected scintillation parameters by using measurements made by Bhat et al. (," For these frequencies, we estimate the expected scintillation parameters by using measurements made by Bhat et al. ("
+1999b) at 327 MIIz. aud Ioucdratiey et al. (,"1999b) at 327 MHz, and Kondratiev et al. ("
+2001) and refereuces therein.,2001) and references therein.
+ We fit a power laws to data available iu a frequency range from 102 MIIz to 1120 AIIz aud derive he values Liste in Table 1.., We fit a power laws to data available in a frequency range from 102 MHz to 1420 MHz and derive the values listed in Table \ref{obstab}.
+ At 238 iuxa 626 MIIz. our observing banchwidths are 11uch larger tha- the expected de-correlation bandwidths. as displaved iu Table 1..," At 238 and 626 MHz, our observing bandwidths are much larger than the expected de-correlation bandwidths, as displayed in Table \ref{obstab}."
+ The observing bandwidth at 1112 MIIz is abou+ twice the estimated de-correlation bandwidth. so that sole flux density modulation can be expected.," The observing bandwidth at 1412 MHz is about twice the estimated de-correlation bandwidth, so that some flux density modulation can be expected."
+ Inudeed. all these expectations are confined in Fie.," Indeed, all these expectations are confirmed in Fig."
+ d. where we showthe measured equivalent continuum flux density for all four frequencies., \ref{flux0329} where we showthe measured equivalent continuum flux density for all four frequencies.
+ We will discuss details in Sect., We will discuss details in Sect.
+ 3.2., 3.2.
+83 covers part of the thir«d short-on orbital cvele. where he LECS intensity is severely reduced. while SI covers uost of the final short-on evele. which may also include ow-state enussion.,"S3 covers part of the third short-on orbital cycle, where the LECS intensity is severely reduced, while S4 covers most of the final short-on cycle, which may also include low-state emission."
+ The spectra were rebiuned to oversample the fu lwidth jialf masta (FWIAD of the cuerey resolution bv a actor 3 aud to have additionally a 100àninnun of 20 counts ver bin to allow use of the 47 statistic., The spectra were rebinned to oversample the full width half maximum (FWHM) of the energy resolution by a factor 3 and to have additionally a minimum of 20 counts per bin to allow use of the $\chi^2$ statistic.
+ Additionally a systematic error Wwvas added to the uncertaimties., Additionally a systematic error was added to the uncertainties.
+ Events were selected iu the energv ranges 0.)0 keV (LECS) and 1.510 keV (MECS) where the iistruneut responses are well deteined au sufficiei counts obtained., Events were selected in the energy ranges 0.1–4.0 keV (LECS) and 1.8–10 keV (MECS) where the instrument responses are well determined and sufficient counts obtained.
+ The photoclectvic absorpticVL CYOSS SCCious of Alorisson MeCanunon L983)) audt1ο solar abudances of Aixors Cirovesse 1989)) are 1sed throughout., The photoelectric absorption cross sections of Morisson McCammon \cite{m:83}) ) and the solar abundances of Anders Grevesse \cite{a:89}) ) are used throughout.
+ A factor was mcelhuded in t1ο spectral fitting to allow for nonualization uncertainties between the iustrmucuts., A factor was included in the spectral fitting to allow for normalization uncertainties between the instruments.
+ This Was COistrained to be within the usial range of 0.8.1.0 (LECSΠΟ) during al fitting., This was constrained to be within the usual range of 0.8–1.0 (LECS/MECS) during all fitting.
+ luitlally. all 45 Spectra were fit with tjo. absorbed power-sv and Dlackhocv continuum τοσοher with 2 broad Cien enudsskn features at —1 keV axd 6.1 keV as used by O97.," Initially, all 5 spectra were fit with the absorbed power-law and blackbody continuum together with 2 broad Gaussian emission features at $\sim$ 1 keV and $\sim$ 6.4 keV as used by O97."
+ The ASCA Solid-Sate huaeiug Spectrometer (SIS) resits of Milia οίng (19913) indicate that the ~1 keV feature may be beter modeled as two narrow lines at at 0.91. keV aud |6 keV. at least in the low-state.," The ASCA Solid-State Imaging Spectrometer (SIS) results of Mihara Soong \cite{m:94}) ) indicate that the $\sim$ 1 keV feature may be better modeled as two narrow lines at at 0.91 keV and 1.06 keV, at least in the low-state."
+ Siuce the full-scidth half masiuuun (FWIIM) euerev resolution of the LECS is 0.2 keV : lkeV.these features are not well resolved in twe LEC'S a so the broad feature was replaced by two narrow lines witli the energies fixed at the alove values in all subsequeut fipA," Since the full-width half maximum (FWHM) energy resolution of the LECS is 0.2 keV at 1 keV, these features are not well resolved in the LECS and so the broad feature was replaced by two narrow lines with the energies fixed at the above values in all subsequent fits."
+ We refor to this as the vstandard” model., We refer to this as the “standard” model.
+ The fit resuts preseuted in Table 1 indicate that this model adequately describes the ALL. SI. 82 and Sl spectra. but not the $3 spectrun where a 47D of GET383 forJ 83Me degrees of eq.freedom (dof). is obtained.," The fit results presented in Table \ref{tab:ave_fits} indicate that this model adequately describes the M1, S1, S2 and S4 spectra, but not the S3 spectrum where a $\chi^{2}$ of 383 for 83 degrees of freedom (dof) is obtained."
+" The fit quality ""ofthe $2 spectruiu is somewhat worse than to the MI. S1. and Sl spectra."," The fit quality of the S2 spectrum is somewhat worse than to the M1, S1, and S4 spectra."
+ This may be because of unresolved varlability which was not excluded from the accumulation (sce the upper aud middle pancls of Fig. 2))., This may be because of unresolved variability which was not excluded from the accumulation (see the upper and middle panels of Fig. \ref{fig:lc}) ).
+ The S1 specti1 niis siguificautly harder aud the CCpuvaleut width (EW) of the Fe-Ix line is higher than the other spectra (Table 1))., The S4 spectrum is significantly harder and the equivalent width (EW) of the Fe-K line is higher than the other spectra (Table \ref{tab:ave_fits}) ).
+ There is no comparable change in the EW of the Fe-L liaes., There is no comparable change in the EW of the Fe-L lines.
+ The intense Fe-I&. feature is learlv visible in the Slsectrua shown in Fie. 3.., The intense Fe-K feature is clearly visible in the S4 spectrum shown in Fig. \ref{fig:spectra}.
+" There is no siguificant change in Fe-I& line energy. between the 4.vectra,", There is no significant change in Fe-K line energy between the spectra.
+ The equivalent blackbody radii are ~200 kin for oeiervals MI. Sl aud $2. decreasing to ~150 kin durius," The equivalent blackbody radii are $\sim$ 200 km for intervals M1, S1 and S2, decreasing to $\sim$ 150 km during"
+if a spatially varving or time dependent underlving equilibrium distribution is chosen since this results in a more complex algorithm |13.14]..,"if a spatially varying or time dependent underlying equilibrium distribution is chosen since this results in a more complex algorithm \cite{thomas1,thomas2}."
+ Particle methods typically solve the Doltzmann equation by applving a splitting in which molecular motion is simulated as a series of collisionless advection and collision steps of length Af., Particle methods typically solve the Boltzmann equation by applying a splitting in which molecular motion is simulated as a series of collisionless advection and collision steps of length $\Delta t$.
+" In such a seheme. the collisionless advection step integrates bv simply advecting particles for a timestep A’. while the collision step integrates by changing the distribution by an amount [Icoll(r.c.1),A "," In such a scheme, the collisionless advection step integrates by simply advecting particles for a timestep $\Delta t$, while the collision step integrates by changing the distribution by an amount $\dfcoll({\bf r},{\bf c},t)\Delta t$."
+Spatial diseretization is introduced by treating collisions as spatially homogeneous within (small) computational cells of volume 12μ., Spatial discretization is introduced by treating collisions as spatially homogeneous within (small) computational cells of volume ${\cal V}_{cell}$.
+ Our approach retains (this basic structure: the particular form of these steps can be summarized as lollows., Our approach retains this basic structure; the particular form of these steps can be summarized as follows.
+ lt can be easily verifiel (hat when the underlying equilibrium distribution is not a function of space or (nme. as is (the case here. the advecüon step for deviational particles is identical to that of physical particles fie. [is also governed by equation (5)) during the advection step].," It can be easily verified that when the underlying equilibrium distribution is not a function of space or time, as is the case here, the advection step for deviational particles is identical to that of physical particles [i.e. $f^d$ is also governed by equation \ref{advection}) ) during the advection step]."
+ Boundary condition implementation. however. differs somewhat because the mass flux (ο boundaries is now split into a deviational contribution and an equilibrium," Boundary condition implementation, however, differs somewhat because the mass flux to boundaries is now split into a deviational contribution and an equilibrium"
+main results obtained by our chemical model for phase I and phase II.,main results obtained by our chemical model for phase I and phase II.
+ In Fig.9. (left panels) we present the fractional abundances (with respect to hydrogen nuclei) of some molecular species of interest CCO. Ν.Ο. Hy.NH3..N2H™.. and HCO)) together with the aabundance —ratio as a function of time for the central core (CC model) and the western/eastern cores (WC/EC model) during phase IL. First of all. we note that the fractional abundances obtained for the central core (top panel) and the western/eastern cores (middle panel) do not show significant differences because they only differ in the core size and the final density reached at the end of this phase.," In Fig.\ref{modelfigs}~ (left panels) we present the fractional abundances (with respect to hydrogen nuclei) of some molecular species of interest CO, $_2$, $_{3}^{+}$, and ) together with the abundance ratio as a function of time for the central core (CC model) and the western/eastern cores (WC/EC model) during phase I. First of all, we note that the fractional abundances obtained for the central core (top panel) and the western/eastern cores (middle panel) do not show significant differences because they only differ in the core size and the final density reached at the end of this phase."
+ As can be seen in Fig. 9..," As can be seen in Fig. \ref{modelfigs},"
+ the fractional abundance of CO.ΝΟ.NH3.. NH... and uncreases with time (with à more pronounced increase for the case of CO.No... and NH3)). reaching its maximum value at times 7=2.5x10° vr.," the fractional abundance of CO, and increases with time (with a more pronounced increase for the case of CO, and ), reaching its maximum value at times $t\simeq2.5\times10^6$ yr."
+ After this time. the fractional abundances start to decrease because molecules freeze out onto dust grains due to the low temperature (T=12 K) and high density qa~10°— em™)) achieved in these cores.," After this time, the fractional abundances start to decrease because molecules freeze out onto dust grains due to the low temperature $T\simeq12$ K) and high density $n\simeq10^{5}-10^{6}$ ) achieved in these cores."
+ On the other hand. the ion kkeeps more or less a constant abundance until =2.5x10° yr. and then its abundance rises moderately.," On the other hand, the ion keeps more or less a constant abundance until $t\simeq2.5\times10^6$ yr, and then its abundance rises moderately."
+ The primary formation route for iis through the reaction between H and N>. which also produces Η».," The primary formation route for is through the reaction between H and $_2$, which also produces $_2$."
+ ecan then react with many atoms and molecules such as C. O. H:O or CO to form Ν..," can then react with many atoms and molecules such as C, O, $_2$ O or CO to form $_2$."
+ Hence it is therefore clear that the abundance in the gas phase of these removal agents will be critical in determining the abundance of ppresent., Hence it is therefore clear that the abundance in the gas phase of these removal agents will be critical in determining the abundance of present.
+ At high densitites. but before the effects of freeze-out dominate. all the molecular species that are important in the removal of aand will be high.," At high densitites, but before the effects of freeze-out dominate, all the molecular species that are important in the removal of and will be high."
+ In particular. for both models the main formation route for iis initially TRNQ—NHHS ++ H. and then the chemical reaction ++ N27 ++ H becomes important. while iis mainly formed via dissociative recombination of NH.," In particular, for both models the main formation route for is initially + $_{2}^{+}$ $\to$ + H, and then the chemical reaction + $\to$ + H becomes important, while is mainly formed via dissociative recombination of $_{4}^{+}$."
+ The latter is formed via consecutive hydrogenation of N (via reactions with molecular hydrogen)., The latter is formed via consecutive hydrogenation of $^{+}$ (via reactions with molecular hydrogen).
+ During phase I the aabundance ratio behaves in the same way for both the central core and the western/eastern cores models. andonly during the latest time steps it depends on the adopted density 10° ffor the western/eastern cores and 1=10° ffor the central core).," During phase I the abundance ratio behaves in the same way for both the central core and the western/eastern cores models, andonly during the latest time steps it depends on the adopted density $n\simeq10^5$ for the western/eastern cores and $n\simeq10^6$ for the central core)."
+ In both cases. the rratio starts from low values and 1t reaches values around 1000 and even higher for t>2.5x10° yr because CO depletion favors the formation of aagainst iim core centers at high densities.," In both cases, the ratio starts from low values and it reaches values around 1000 and even higher for $t>2.5\times10^6$ yr because CO depletion favors the formation of against in core centers at high densities."
+ This suggests that long-lived starless cores are associated with high values of the aabundance ratio., This suggests that long-lived starless cores are associated with high values of the abundance ratio.
+ Figure 9. (right panels) shows the fractional abundances (with respect to hydrogen nuclei) as a function of time of CO. No. Hi.NH3..NH.. and derived from our chemical model during phase II for the central core (top panel) and the western/eastern cores (middle panel).," Figure \ref{modelfigs}~ (right panels) shows the fractional abundances (with respect to hydrogen nuclei) as a function of time of CO, $_2$, $_{3}^{+}$, and derived from our chemical model during phase II for the central core (top panel) and the western/eastern cores (middle panel)."
+ The fractional abundances in the central core clearly have a different behavior compared to that of the western and eastern cores as a consequence of the different density. and especially the different temperature reached at the end of phase II.," The fractional abundances in the central core clearly have a different behavior compared to that of the western and eastern cores as a consequence of the different density, and especially the different temperature reached at the end of phase II."
+" In the central core. the fractional abundance of CO. N».... and iis constant during the initial stages. and once the temperature increases, their fractional abundances rise significantly."," In the central core, the fractional abundance of CO, $_2$, and is constant during the initial stages, and once the temperature increases, their fractional abundances rise significantly."
+ On the other hand. the fractional abundance of the molecular ton ddecreases with time.," On the other hand, the fractional abundance of the molecular ion decreases with time."
+ This ts of course due to the fact that iis highly reactive at high densities., This is of course due to the fact that is highly reactive at high densities.
+ The abrupt increase of the fractional abundance takes place at the same time for all molecules. around ¢=6.1x10° yr (corresponding to a temperature of 7521 K).," The abrupt increase of the fractional abundance takes place at the same time for all molecules, around $t\simeq6.1\times10^3$ yr (corresponding to a temperature of $T\simeq21$ K)."
+ This time corresponds to the time when a fraction of weakly bound species evaporates from the grain mantle (see Table | in Viti et al., This time corresponds to the time when a fraction of weakly bound species evaporates from the grain mantle (see Table 1 in Viti et al.
+ 2004)., 2004).
+ The fractional abundance of rreaches its maximum value 4.6x107! (corresponding to a ecolumn density of ~1.810 emp at r=17x107 yr., The fractional abundance of reaches its maximum value $4.6\times10^{-11}$ (corresponding to a column density of $\sim1.8\times10^{12}$ ) at $t\simeq1.7\times10^4$ yr.
+" The rapid destruction of ((by the chemical reaction ++ CO ++ Ns) occurs att~2.3x107 yr. with a variation of the fractional abundance of one oder of magnitude. being the ecolumn density around ~8x10!""ο. consistent. with the value estimated from the observations."," The rapid destruction of (by the chemical reaction + CO $\to$ + $_2$ ) occurs at $t\simeq2.3\times10^4$ yr, with a variation of the fractional abundance of one oder of magnitude, being the column density around $\sim8\times10^{11}$, consistent with the value estimated from the observations."
+ On the other hand. fractional abundance is constant at ~6x1075.," On the other hand, fractional abundance is constant at $\sim6\times10^{-10}$."
+ Assuming we estimated an ecolumn density of ||240.~1.3x10em7?.. in good agreement with the estimation from the observational data (see Table 2)).," Assuming ]=40, we estimated an column density of $\sim1.3\times10^{12}$, in good agreement with the estimation from the observational data (see Table \ref{coldens}) )."
+ This chemical behaviour could be explained in terms of CO desorption from grain mantles. which causes a substantial destruction of ffavoring the formation of HCO™.. being this reaction the main removal mechanism of ffor ‘standard’ CO abundances of |CO/H;]]21077 Jorgensen 2004).," This chemical behaviour could be explained in terms of CO desorption from grain mantles, which causes a substantial destruction of favoring the formation of , being this reaction the main removal mechanism of for `standard' CO abundances of $\simeq10^{-4}$ \citep{jorgenssen2004,lee2004}."
+ In fact. as can be seen in," In fact, as can be seen in"
+diagrams in Figure 4..,diagrams in Figure \ref{fig:VHcolmag}.
+ In these CALDs the red supergiant positions are consistent with ages of 25 Myr or 440 Myr for NGCISIS and 40 Myr or 254440 Myr [or Νας1505., In these CMDs the red supergiant positions are consistent with ages of 25 Myr or 40 Myr for NGC1818 and 40 Myr or 40 Myr for NGC1805.
+ To compare the observations and the cdillerent age simulations we first look at the shape of the observed. and expected CMDs., To compare the observations and the different age simulations we first look at the shape of the observed and expected CMDs.
+ Ehe reddening also alfects the position of the CMD at bright magnitudes. and so we have compared simulations ancl observations for 25. 40. LO&225 and 254440 Avr and E(B-V) between 0 and 1.," The reddening also affects the position of the CMD at bright magnitudes, and so we have compared simulations and observations for 25, 40, 25 and 40 Myr and E(B-V) between 0 and 1."
+ We find that only the 25 Myr and 25&440 Myr simulations with E(D-V)—0.075. describe wel the observed. shape.," We find that only the 25 Myr and 40 Myr simulations with E(B-V)=0.075, describe well the observed shape."
+ To distinguish. between these two simulations we look at the observed and expected numbers of bright sws in 0.5 mae bins down the main sequence., To distinguish between these two simulations we look at the observed and expected numbers of bright stars in 0.5 mag bins down the main sequence.
+ ‘Table 9 gives 116 observed. ancl expected. numbers of the brightest stars for the 25 and 254440 Myr simulations for 12 V«17., Table \ref{tab:brightnum} gives the observed and expected numbers of the brightest stars for the 25 and 40 Myr simulations for $<$ $<$ 17.
+ Reeall that the simulations and observations were forced to have he same number of stars in the range 1T«V «19 used for normalisation., Recall that the simulations and observations were forced to have the same number of stars in the range $<$ $<$ 19 used for normalisation.
+ The observed. numbers here include the De stars., The observed numbers here include the Be stars.
+ The simulations do not contain 3o stars. but the Be phenomenon just moves stars recwards in the colour-magnitucle diagram. and so the numbers per magnitude bin remain the same.," The simulations do not contain Be stars, but the Be phenomenon just moves stars redwards in the colour-magnitude diagram, and so the numbers per magnitude bin remain the same."
+ The comparison of the observed and simulated numbers is complicated. by the fact that the number of blue, The comparison of the observed and simulated numbers is complicated by the fact that the number of blue
+"defines the metric potentials y(n,x) and yP(y,x).","defines the metric potentials $\psi^{(1)}(\eta,{\bf x})$ and $\chi^{(1)}(\eta,{\bf x})$."
+ Below we use the convention ao=1 for today’s scale factor., Below we use the convention $a_{0}=1$ for today's scale factor.
+ We use conformal time 7 and the traceless differential operator D;;—0;0;—1óij on a spatially flat background., We use conformal time $\eta$ and the traceless differential operator $D_{ij}=\partial_{i}\partial_{j}-\frac{1}{3}\delta_{ij}\Delta$ on a spatially flat background.
+ The geometrical quantities of interest are the local expansion rate and the spatial curvature., The geometrical quantities of interest are the local expansion rate and the spatial curvature.
+ The former follows from the expansion tensor and reads where ()’ stands for the derivative with respect to conformal time., The former follows from the expansion tensor and reads where $()^{\prime}$ stands for the derivative with respect to conformal time.
+" Calculating the spatial Ricci curvature from the above metric yields By the covariant conservation of the energy momentum tensor, y“ is related to the matter density contrast by with ὅ(x) denoting a constant of integration."," Calculating the spatial Ricci curvature from the above metric yields By the covariant conservation of the energy momentum tensor, $\psi^{\left(1\right)}$ is related to the matter density contrast by with $\bar{\zeta}\left({\bf x}\right)$ denoting a constant of integration."
+" This constant plays no role in the following, because 0 and & involve only time derivatives of y))."," This constant plays no role in the following, because $\theta$ and ${\cal R}$ involve only time derivatives of $\psi^{\left(1\right)}$."
+" For dust and a cosmological constant, Einstein's equations give the well-known relation at first order in perturbationtheory."," For dust and a cosmological constant, Einstein's equations give the well–known relation at first order in perturbationtheory."
+" For the ACDM model the solution reads [see, e.g. Greenetal. (2005)]] where δρ(x) is the density perturbation today."," For the $\Lambda$ CDM model the solution reads [see, e.g. \citet{green2005}] ] where $\delta_{0}\left(\bx\right)$ is the density perturbation today."
+ D(a) is the growth factor given by and 2F is a hypergeometric function., $D\left(a\right)$ is the growth factor given by and $_{2}F_{1}$ is a hypergeometric function.
+ In the following we denote today’s value of the growth factor by Do=D(1)., In the following we denote today's value of the growth factor by $D_0 \equiv D\left(1\right)$.
+" Plugging this solution into (10)) and using (13)), we find the local expansion rate expressed in terms of the growth rate From (11)) we find the local spatial curvature From these quantities we can define local © functions, demonstrating that the importance of curvature effects grows proportional to the formation of structures."," Plugging this solution into\ref{eq:theta-pert}) ) and using \ref{eq:def-psi}) ), we find the local expansion rate expressed in terms of the growth rate From \ref{eq:R-pert}) ) we find the local spatial curvature From these quantities we can define local $\Omega$ functions, demonstrating that the importance of curvature effects grows proportional to the formation of structures."
+" A remarkable property is that >)O;(a,x)=1 holds not only for the FLRW background, but also at the level of perturbations."," A remarkable property is that $\sum \Omega_i(a,{\bf x}) = 1$ holds not only for the FLRW background, but also at the level of perturbations."
+" For linear perturbations the kinematic backreaction term does not play any role, but becomes important as soon as quadratic terms are considered."," For linear perturbations the kinematic backreaction term does not play any role, but becomes important as soon as quadratic terms are considered."
+ Let us now compare these local quantities with the domain-averaged expansion rate and the spatial curvature (Li& 2008).., Let us now compare these local quantities with the domain-averaged expansion rate and the spatial curvature \citep{li:scale}.
+" From the definition of the average () in (2)) we find that, in principle, fluctuations in the volume element dg have to be taken into account."," From the definition of the average $\average{}$ in \ref{eq:Def-Mittel}) ) we find that, in principle, fluctuations in the volume element $d\mu_{g}$ have to be taken into account."
+" Writing du,=Jdx with the functional determinant J=a?−sy), the average over the perturbed hypersurface agrees(1 with an average over an unperturbed Euclidean domain if we restrict our attention to linear perturbations."," Writing ${\rm d}\mu_{g} = J \rmd^{3}x$ with the functional determinant $J=a^{3}\left(1-3\psi^{\left(1\right)}\right)$, the average over the perturbed hypersurface agrees with an average over an unperturbed Euclidean domain if we restrict our attention to linear perturbations."
+" We express domain-averaged quantities in terms of the volume scale factor ap, because we assume that the measured redshift in an inhomogeneous universe is related to the average scale factor."," We express domain-averaged quantities in terms of the volume scale factor $a_{\cal D}$, because we assume that the measured redshift in an inhomogeneous universe is related to the average scale factor."
+" This has been advocated by Rüsánen(2009),, where the relation has been established."," This has been advocated by \cite{rasanen:light}, where the relation has been established."
+ Note that in principle one would have to introduce averaging on some larger scale than D to connect this background average on some domain $$ to the redshift., Note that in principle one would have to introduce averaging on some larger scale than $\cD$ to connect this background average on some domain $\CB$ to the redshift.
+" For the sake of simplicity, and because we content ourselves with small redshifts, we use the same domain 4)."," For the sake of simplicity, and because we content ourselves with small redshifts, we use the same domain $\cD$ ."
+ This limits the validity of the result to small redshifts., This limits the validity of the result to small redshifts.
+" In order to relate a and ag, we start from To first order this relation gives"," In order to relate $a$ and $a_\cD$ , we start from To first order this relation gives"
+Ad=17555+0'05.,$\Delta\delta=1\farcs555\pm0\farcs05$.
+ These values are consistent with the rough offset of 9771 east and 1606 north given by Quirrenbachοἱal.(2001) for their AO imaging and in reasonable agreement with the offset of 975402 east and 175402 north they. found from their lower-resolution NIRC imaging., These values are consistent with the rough offset of 7 east and 6 north given by \citet{qui01} for their AO imaging and in reasonable agreement with the offset of $9\farcs5\pm0\farcs2$ east and $1\farcs5\pm0\farcs2$ north they found from their lower-resolution NIRC imaging.
+ In order to use this offset with the FCS coordinates for AO-GS. we convert it from the offset [rom the center of light of the binary. to the offset from the primary (easternmost) component of the binary.," In order to use this offset with the FGS coordinates for AO-GS, we convert it from the offset from the center of light of the binary to the offset from the primary (easternmost) component of the binary."
+" The magnitude difference of the components at A"" is 0.ΕΕ mag. and we use the accurate separation of 1147 and position angle of ddetermined from the FCS trausfer-mode observations."," The magnitude difference of the components at $K'$ is 0.44 mag, and we use the accurate separation of 147 and position angle of determined from the FGS transfer-mode observations."
+ The offset of the stellar object from ihe AO-GS primary is then 97750+0702 east and 17520+0705 north., The offset of the stellar object from the AO-GS primary is then $9\farcs750\pm0\farcs02$ east and $1\farcs520\pm0\farcs05$ north.
+ Adding this offset to the position of the AO-GS primary found above gives the position of the stellar object near 2294 as 1H4069LE1459+00017.7-347112479140705.," Adding this offset to the position of the AO-GS primary found above gives the position of the stellar object near 294 as $14^{\rm h} 06^{\rm m} 44\fs1459 \pm 0\fs0017, +34\arcdeg 11\arcmin 24\farcs91 \pm 0\farcs05$."
+ The position of the unresolved [lat-spectrum radio core of 2294 has been given bv AMeCarthyetal.(1990).. from their observations at 6 em with the Very Large Array (VLA) in the A-array configuration.," The position of the unresolved flat-spectrum radio core of 294 has been given by \citet{mcC90}, from their observations at 6 cm with the Very Large Array (VLA) in the A-array configuration."
+ Their position for the core is 14043450602502005.Q'05 in the reference frame defined by the VLA B1950 astrometric calibrators.," Their position for the core is $14^{\rm h} 04^{\rm m} 34\fs060 \pm 0\fs005, +34\arcdeg 25\arcmin 40\farcs00 \pm 0\farcs05$ in the reference frame defined by the VLA B1950 astrometric calibrators."
+ Conversion ol these coordinates to the ICRS (J2000) reference frame was kindly carried out [or us by Jim Ulvestad., Conversion of these coordinates to the ICRS (J2000) reference frame was kindly carried out for us by Jim Ulvestad.
+ The result is 145064074+(0:005.2-24711247950705.," The result is $14^{\rm h} 06^{\rm m} 44\fs074 \pm 0\fs005, +34\arcdeg 11\arcmin 24\farcs95 \pm 0\farcs05$."
+ These values. are very close 0o those given bv Stockton.Canalizo.&Ridgeway(1999).. who used a different procedure. bul Chev should be more accurate: uncertainties introduced by the conversion should be <1 mas (J. Ulvestacd. private communication).," These values are very close to those given by \citet{sto99}, who used a different procedure, but they should be more accurate; uncertainties introduced by the conversion should be $<1$ mas (J. Ulvestad, private communication)."
+" With our precise mapping of the AO imaging onto the ICRS reference frame. we can now locate the position of the radio core on the AO A""-band image."," With our precise mapping of the AO imaging onto the ICRS reference frame, we can now locate the position of the radio core on the AO $K'$ -band image."
+ It lies 0789+0707 west and Q'042:0707 north of the stellar object. as shown in Fig. 3..," It lies $0\farcs89\pm0\farcs07$ west and $0\farcs04\pm0\farcs07$ north of the stellar object, as shown in Fig. \ref{ao_core}."
+" The radio core. which presumably marks (he nucleus. lies very. close to one of the brightness peaks of what have called ""the western component"" of 2294."," The radio core, which presumably marks the nucleus, lies very close to one of the brightness peaks of what \citet{qui01} have called “the western component” of 294."
+ On the other hand. this position of the radio core is clearly inconsistent with its being identified with the stellar object. as was suggested by Quirrenbachetal.(2001).," On the other hand, this position of the radio core is clearly inconsistent with its being identified with the stellar object, as was suggested by \citet{qui01}."
+. One of the most striking features of Fig., One of the most striking features of Fig.
+ 1 is the difference in appearance of the various AO images., \ref{aoimagefig} is the difference in appearance of the various AO images.
+ All show a stellar component and a diffuse component about (to (he west of it. but there is little detail (hat convincingly reproduces from image to image.," All show a stellar component and a diffuse component about to the west of it, but there is little detail that convincingly reproduces from image to image."
+" To be sure. panel C is a rather noisy image. and panel D (Quirrenbachetal.2001) is an IH-band image (the other 3 are all A""). so we should probably concentrate on comparing the"," To be sure, panel $C$ is a rather noisy image, and panel $D$ \citep{qui01} is an $H$ -band image (the other 3 are all $K'$ ), so we should probably concentrate on comparing the"
+into the corona (Parnell 2010b) where current sheets between the emerging and pre-existing coronal field were shown to be threaded. by numerous separators. a feature ol complex magnetic fields suggested by Albright (1999).,"into the corona (Parnell 2010b) where current sheets between the emerging and pre-existing coronal field were shown to be threaded by numerous separators, a feature of complex magnetic fields suggested by Albright (1999)."
+ Parnell (2010b) found the regions of hiehest integrated parallel electric field to be at or in association with the separators. so arguing (hiat reconnection occurs at (he separators themselves.," Parnell (2010b) found the regions of highest integrated parallel electric field to be at or in association with the separators, so arguing that reconnection occurs at the separators themselves."
+ Separator reconnection also appears to be a kev process al the magnetopause where magnetic nulls or even clusters of nulls appear in the northern and southern polar cusp regions (Dorelli 2007)., Separator reconnection also appears to be a key process at the magnetopause where magnetic nulls or even clusters of nulls appear in the northern and southern polar cusp regions (Dorelli 2007).
+ These nulls are linked bv à separator along the davside maguelopause which will therefore be involved in reconnection with the interplanetary magnetic field., These nulls are linked by a separator along the dayside magnetopause which will therefore be involved in reconnection with the interplanetary magnetic field.
+ Indeed reconnection in such a configuration has been implied by observations from the Cluster spacecralt (Xiao 2007)., Indeed reconnection in such a configuration has been implied by observations from the Cluster spacecraft (Xiao 2007).
+ These studies partially motivate the present work where we want to investigate how and why multüple separators can be involved in reconnection., These studies partially motivate the present work where we want to investigate how and why multiple separators can be involved in reconnection.
+" Although separator configurations have been analvsed for some time (e.g. Chance 1992: Craig 1999). (here remain several unanswered basic questions surrounding the nature of separator reconnection. particularly involving the wav magnetic flix evolves. i.e, exactly how the changes in field line connectivity occur."," Although separator configurations have been analysed for some time (e.g. Chance 1992; Craig 1999), there remain several unanswered basic questions surrounding the nature of separator reconnection, particularly involving the way magnetic flux evolves, i.e. exactly how the changes in field line connectivity occur."
+ There have been suggestions that the reconnection involves a cut and paste of field lines at the separator itself in a manner similar to the 2D null case (e.g. Lau Finn 1990: Priest Titov 1996: Longcope 2001)., There have been suggestions that the reconnection involves a `cut and paste' of field lines at the separator itself in a manner similar to the 2D null case (e.g. Lau Finn 1990; Priest Titov 1996; Longcope 2001).
+ However. previous detailed investigations into the nature of individual 3D reconnection events both at and away from null points have shown the flux evolution to be quite distinct from the 2D picture (ILornig Priest 2003: Pontin 2004. 2005) and. in the light of these models. such a simple flux evolution would be surprising.," However, previous detailed investigations into the nature of individual 3D reconnection events both at and away from null points have shown the flux evolution to be quite distinct from the 2D picture (Hornig Priest 2003; Pontin 2004, 2005) and, in the light of these models, such a simple flux evolution would be surprising."
+ Indeed a number of kev features of 3D reconnection are now known which show many ol its features are quite distinct to the 2D case., Indeed a number of key features of 3D reconnection are now known which show many of its features are quite distinct to the 2D case.
+ Accordingly it would appear useful to re-examine (he fundamental nature of separator reconnection and consider whether it has, Accordingly it would appear useful to re-examine the fundamental nature of separator reconnection and consider whether it has
+Lyman-r emission.,$\alpha$ emission.
+ The distribution of active regions on the solar disk prior. during. and after the passage of the sunspot groups are presented in the bottom of the figure.," The distribution of active regions on the solar disk prior, during, and after the passage of the sunspot groups are presented in the bottom of the figure."
+ Note that during the transit of sunspots on the solar disk the total solar irradiance reduces significantly as expected., Note that during the transit of sunspots on the solar disk the total solar irradiance reduces significantly as expected.
+ The transit of the active regions from their appearance on 2011-07-26 to the maximum on is observed as a progressive increase of the filling factors from the limb to the center rings., The transit of the active regions from their appearance on 2011-07-26 to the maximum on 2011-08-01 is observed as a progressive increase of the filling factors from the limb to the center rings.
+ At the same time. the TSI decreases progressively. reaching its minimum value around 12:00 when three (3) large active regions in the northern hemisphere are near the disk center.," At the same time, the TSI decreases progressively, reaching its minimum value around 2011-08-01 12:00 when three (3) large active regions in the northern hemisphere are near the disk center."
+ A gradual decrease of the area covered by the active regions/sunspots and the concomitant increase of the TSI occurs from the 2011-08-01 to 2011-08-10., A gradual decrease of the area covered by the active regions/sunspots and the concomitant increase of the TSI occurs from the 2011-08-01 to 2011-08-10.
+ However. as the pattern of the Lyman-w emission illustrates. the presence of active regions produce different patterns in different bands of the spectrum.," However, as the pattern of the $\alpha$ emission illustrates, the presence of active regions produce different patterns in different bands of the spectrum."
+ In contrast to the decrease of the TSI during the passage of the active regions. a clear enhancement of the Lyman-a emission is observed from 2011-07-26 to 2011-08-10.," In contrast to the decrease of the TSI during the passage of the active regions, a clear enhancement of the $\alpha$ emission is observed from 2011-07-26 to 2011-08-10."
+ A large enhancement of the Lyman-a emission is also observed from 2011-07-11 to 2011-07-26 without a concomitant decrease of the TSI., A large enhancement of the $\alpha$ emission is also observed from 2011-07-11 to 2011-07-26 without a concomitant decrease of the TSI.
+ In this way. the feature extraction procedure is able to identify the presence and the distribution of the active regions/sunspots that are needed to compute the evolution of the solar irradiance employing a neural network.," In this way, the feature extraction procedure is able to identify the presence and the distribution of the active regions/sunspots that are needed to compute the evolution of the solar irradiance employing a neural network."
+ As discussed in Sect., As discussed in Sect.
+ 3. the coefficients of the network are obtained by comparing the output of the model and observations of the solar irradiance made by instruments onboard of SORCE spacecraft.," 3, the coefficients of the network are obtained by comparing the output of the model and observations of the solar irradiance made by instruments onboard of SORCE spacecraft."
+ Here we focus on two regions of the spectra that are covered by SORCE instruments., Here we focus on two regions of the spectra that are covered by SORCE instruments.
+information available at the time of our analysis. OGLE-TR-LOb.,"information available at the time of our analysis. OGLE-TR-10b,"
+ Llib. 182b. 211b. and SWISEPS-04b and 11b were removed due to the host stars being of unknown spectral tvpe.," 111b, 182b, 211b, and SWEEPS-04b and 11b were removed due to the host stars being of unknown spectral type."
+ GJ 436b. was not included: because it has the only Al type host. and. will have dillerent high energy. emisison characteristics to the FGIx hosts prevalent in the rest of the sample.," GJ 436b was not included because it has the only M type host, and will have different high energy emisison characteristics to the FGK hosts prevalent in the rest of the sample."
+ LLDI49026b. was not included because its host star is not on the main sequence., HD149026b was not included because its host star is not on the main sequence.
+ In Figure 1. we plot mass against orbital period for our sample after 2.. and surface gravity vs period after ?..," In Figure \ref{correlations} we plot mass against orbital period for our sample after \cite{mazeh}, and surface gravity vs period after \cite{pjw}."
+ 3oth plots show greater scatter than the original suggested correlations. especially to the upper right. but the mass-period. distribution still shows a sharp cut-off to the lower left.," Both plots show greater scatter than the original suggested correlations, especially to the upper right, but the mass-period distribution still shows a sharp cut-off to the lower left."
+ Phe lower left of the surface eravitv-period distribution does not show a sharp cutoll. but this area remains relatively sparsely populated. despite strong selection effects in transit surveys for planets in this region.," The lower left of the surface gravity-period distribution does not show a sharp cutoff, but this area remains relatively sparsely populated, despite strong selection effects in transit surveys for planets in this region."
+ The upper right of both of these diagrams also remain sparsely populated., The upper right of both of these diagrams also remain sparsely populated.
+ “Phis is unlikely to be related to evaporation and instead may be a feature of planetary migration., This is unlikely to be related to evaporation and instead may be a feature of planetary migration.
+ Selection ellects in transit surveys may also be a factor. since they select against long orbital periods and high surface eravities.," Selection effects in transit surveys may also be a factor, since they select against long orbital periods and high surface gravities."
+ ]t has been shown bv ? that Jeans escape of material from hot Jupiters cannot drive significant mass loss., It has been shown by \cite{lammer} that Jeans escape of material from hot Jupiters cannot drive significant mass loss.
+ Instead they suggest that energy-limited escape of particles due to absorption of EUV/X-rav. photons is the most-likely. ciriver of mass loss in these svstenis., Instead they suggest that energy-limited escape of particles due to absorption of EUV/X-ray photons is the most-likely driver of mass loss in these systems.
+" In order to estimate the evaporation rates of exoplanets we adopt the model of 2.. who used the ratio of planetary potential energy. ancl the time-averagecl power incident on the planet in order to predict mass loss rates. given hy where m ds mass loss rate. £, is the stars X-ray/EUVY luminosity. Ze, is the planetary radius. AZ, is the planetary mass. G is the gravitational constant. ancl is the time-averaged orbital clistance."," In order to estimate the evaporation rates of exoplanets we adopt the model of \cite{lecav}, who used the ratio of planetary potential energy and the time-averaged power incident on the planet in order to predict mass loss rates, given by where $\dot{m}$ is mass loss rate, $L_{x}$ is the star's X-ray/EUV luminosity, $R_{p}$ is the planetary radius, $M_{p}$ is the planetary mass, G is the gravitational constant, and $\langle$ $\rangle$ is the time-averaged orbital distance."
+ This simple model assumes energy-Iimited evaporation. which may overestimate real mass loss rates.," This simple model assumes energy-limited evaporation, which may overestimate real mass loss rates."
+ It also assumes that a test particle must be moved to infinity to escape the eravitational potential of a planet. (Sect.4.3of2).. which ? suggest may underestimate mass loss by up to when a planet is close to filling its Roche lobe.," It also assumes that a test particle must be moved to infinity to escape the gravitational potential of a planet \citep[Sect.\,4.3 of][]{lecav}, which \cite{rochelobeeffect} suggest may underestimate mass loss by up to when a planet is close to filling its Roche lobe."
+ Further. we neglect any evolution of the planet radius. assuming the current value for each planet.," Further, we neglect any evolution of the planet radius, assuming the current value for each planet."
+ Since unperturbed gaseous. planets contract slowly with time. this will tend to underestimate mass loss rates.," Since unperturbed gaseous planets contract slowly with time, this will tend to underestimate mass loss rates."
+ X-rayοV. emission from stars is believed to be driven by a dvnamo action that decreases with time as the star spins down due to magnetic braking (2).., X-ray/EUV emission from stars is believed to be driven by a dynamo action that decreases with time as the star spins down due to magnetic braking \citep{Skumanic}.
+ N-rav/EUN emission is seen to increase with stellar rotation rate. up to à maximum of Ly fbi. c 10. where it saturates (?2)..," X-ray/EUV emission is seen to increase with stellar rotation rate, up to a maximum of $_x$ $_{bol}$ $\simeq$ $^{-3}$, where it saturates \citep{vil,vilwil}."
+ We expect therefore. all known exoplanets to have sulfered à. period of high and approximately constant Narav/EEUV irradiation in the past. when their host stars were voung and rapidly rotating.," We expect therefore all known exoplanets to have suffered a period of high and approximately constant X-ray/EUV irradiation in the past, when their host stars were young and rapidly rotating."
+ This saturated: phase is likely to dominate their irracliation history., This saturated phase is likely to dominate their irradiation history.
+ The duration of the saturated. phase ofX-ray emission depends on spectral type. but is not well known.," The duration of the saturated phase ofX-ray emission depends on spectral type, but is not well known."
+" For E tvpe stars a saturation time of 0.1 Gyr is used. taken from ? who used optical anc X-ray telescope data of five star-lorming regions of varving age (p Ophiuchi. the Orion Nebula Cluster, NGC 2264. Chamacleon 1. and. ἡ Chamacleontis) and two voung clusters (the Pleiades and NGC 2516)."," For F type stars a saturation time of 0.1 Gyr is used, taken from \cite{fstarsat} who used optical and X-ray telescope data of five star-forming regions of varying age $\rho$ Ophiuchi, the Orion Nebula Cluster, NGC 2264, Chamaeleon I, and $\eta$ Chamaeleontis) and two young clusters (the Pleiades and NGC 2516)."
+ This is their upper limit for the saturation period of stars with the mass of E stars., This is their upper limit for the saturation period of stars with the mass of F stars.
+ For G type stars we estimate a saturation time of 0.2 Car using the results of ?.. who used the ASCA and ROSAT X-ray satellites to probe the coronae of a sample of nine solar-like CG stars. with ages between 70 Myr and 9 Civr.," For G type stars we estimate a saturation time of 0.2 Gyr using the results of \cite{gstarsat}, who used the ASCA and ROSAT X-ray satellites to probe the coronae of a sample of nine solar-like G stars, with ages between 70 Myr and 9 Gyr."
+ The saturation timescale for Ix type stars is estimated rom the observation (e.g...2) that Ix stars are saturatec in he Pleiades (ageo£0.1Gyr.7) but not in the Hades (ageof~0.6Gyr. 2)...," The saturation timescale for K type stars is estimated from the observation \citep[e.g., ][]{xraysatwheatley} that K stars are saturated in the Pleiades \citep[age of 0.1 Gyr,][]{pleiades} but not in the Hyades \citep[age of $\sim$0.6 Gyr,][]{hyades}."
+ We choose a value of 0.35 Gyr. but adjusting his value to the extremes does not alleet our conclusions (sce Section 3.3)).," We choose a value of 0.35 Gyr, but adjusting this value to the extremes does not affect our conclusions (see Section \ref{effect}) )."
+ For each of the exoplanets in our sample. we have calculated the total incident energy over the planet surface or the complete saturation interval of its host. star.," For each of the exoplanets in our sample, we have calculated the total incident energy over the planet surface for the complete saturation interval of its host star."
+ The stellar parameters required are. [listed in. Table 1 and Mdanetaryv properties in Table 2.., The stellar parameters required are listed in Table \ref{table2} and planetary properties in Table \ref{table1}.
+ The saturation duration is based. only on the spectral type of the star., The saturation duration is based only on the spectral type of the star.
+ Phe X-ray saturation level used can be found by using the saturation evel for stars of a given mass (inTable3of?).. Four stars in this work are slightlv more massive than those studied, The X-ray saturation level used can be found by using the saturation level for stars of a given mass \citep[in Table 3 of][]{pizzolato}.. Four stars in this work are slightly more massive than those studied
+found the barvon fraction inside μι of9O%.. which agrees well with our method at higher masses.,"found the baryon fraction inside $r_{500}$ of, which agrees well with our method at higher masses."
+" However. Crainetal.(2006) also "" (his fraction still holds at lower masses. and that the barvon fraction is still at the virial radius (again with no feedback). while we find a higher fraction in both of these instances."," However, \citet{CrainEFJMNP06z} also find this fraction still holds at lower masses, and that the baryon fraction is still at the virial radius (again with no feedback), while we find a higher fraction in both of these instances."
+ Note that if we had instead started with an initial state consistent with a barvon fraction. then we would require a lower amount of feedback to reproduce observed gas fractions.," Note that if we had instead started with an initial state consistent with a baryon fraction, then we would require a lower amount of feedback to reproduce observed gas fractions."
+ It is useful in this context to compare with recent hydrodyenamical simulations which include feedback., It is useful in this context to compare with recent hydrodynamical simulations which include feedback.
+" The left panel of Fig. relfig:sims “= eas lo total mass inside μμ: the solid line shows the best-fit, power law relation al 2='0 found from adaptive refinement simulations by lhravisovetal.(20060)."," The left panel of $.$ \\ref{fig:sims}
+ relates gas to total mass inside $r_{500}$; the solid line shows the best-fit power law relation at $z=0$ found from adaptive refinement simulations by \citet{KravtsovVN06}."
+ With no feedbac we [ind a similar slope (slightly larger (han unity). but [or a given. halo mass there is a higher gas fraction than in the simulations.," With no feedback we find a similar slope (slightly larger than unity), but for a given halo mass there is a higher gas fraction than in the simulations."
+ Adding feedback reduces this discrepancy al higher masses. but at lower masses our maximun feedback case instead has lower gas fractions.," Adding feedback reduces this discrepancy at higher masses, but at lower masses our maximum feedback case instead has lower gas fractions."
+ As Ixravisovetal.(2006) did not attempt to include AGN feedback it is nol surprising that an intermediate amount of feedback would provide the best match., As \citet{KravtsovVN06} did not attempt to include AGN feedback it is not surprising that an intermediate amount of feedback would provide the best match.
+ We do not match the simulations when it comes to the total mass-temperature relation. however.," We do not match the simulations when it comes to the total mass-temperature relation, however."
+ For Equ. (9)). Avavtso," For Eqn. \ref{eqn:mtpowlaw}) ),"
+vetal.(2006). [ind A=3.8520.19 anda =1.524c0.07. while al.(2006) with an SPII simulation find A=4.47£0.19 and à=1.76£0.07.," \citet{KravtsovVN06} find $A=3.85\pm 0.19$ and $\alpha=1.524\pm 0.07$, while \citet{KaydSABLPST06z} with an SPH simulation find $A=4.47\pm 0.19$ and $\alpha=1.76\pm 0.07$."
+ We obtain a similar slope (a= 1.62). but more importantly we find a different normalization— that is. we find hotter temperatures al a given mass. as would be expected because we included a higher level of feedback. The near future will see a number of survevs that select clusters of galaxies via (heir sunvaev-Zelclovich (3Z) decrement. which is proportional to the gas pressure in the cluster integrated along the line of sight.," We obtain a similar slope $\alpha=1.62$ ), but more importantly we find a different normalization— that is, we find hotter temperatures at a given mass, as would be expected because we included a higher level of feedback, The near future will see a number of surveys that select clusters of galaxies via their Sunyaev-Zel'dovich (SZ) decrement, which is proportional to the gas pressure in the cluster integrated along the line of sight."
+ Currently the Sunyaev-Zeldovich. Array (8ZÀ:Carlstrometal.2000) and the Are Minute Microkelvin Imager (AMI:Ixneiss]etal.2001) are equipped to perform such a search on tens of square degrees on the sky., Currently the Sunyaev-Zel'dovich Array \citep[SZA;][]{Carlstrom:00} and the Arc Minute Microkelvin Imager \citep[AMI;][]{Kneissl:01} are equipped to perform such a search on tens of square degrees on the sky.
+" IIowever. the Atacama. Cosmology Telescope (ACT). the South Pole Telescope (SPT). APEX. and ultimately the ""n survevor will scan ol square degrees on "" skv in the radio (Ruhletal.GiistenPMNNetal.ni>Clavel&Tauber 2005)."," However, the Atacama Cosmology Telescope (ACT), the South Pole Telescope (SPT), APEX, and ultimately the Planck surveyor will scan thousands of square degrees on the sky in the radio \citep{Ruhl:04,Fowler:04,Gusten:06,Planck}."
+. These surveys will detect ne of clusters: for SPT will scan 4.000lhe deg? on the skv and observe ol the order of 6.000. clusters of galaxies (this is for a sensitivitv of about L.5mJv at the 4-sigma detection threshold with a 1 arcmin beam operating at 150 Gllz. and assumes ay= 0.75).," These surveys will detect thousands of clusters; for example, the SPT will scan 4,000 $deg^2$ on the sky and observe of the order of $6,000$ clusters of galaxies (this is for a flux sensitivity of about $1.5$ mJy at the 4-sigma detection threshold with a 1 arcmin beam operating at 150 GHz, and assumes $\sigma_8 = 0.75$ )."
+" This translates to a limiting mass of Mg,zz10Η!M, if one assumes the clusters ave in hydrostatic equilibrium.", This translates to a limiting mass of $M_{\rm lim}\approx 10^{14.2}h^{-1}M_\odot$ if one assumes the clusters are in hydrostatic equilibrium.
+ Sebealetal.(2006). found a similar limit for a complete cluster sample from ACT., \citet{SehgalTHB06z} found a similar limit for a complete cluster sample from ACT.
+ The redshift distribution of clusters is very sensitive to the amplitude and growth of, The redshift distribution of clusters is very sensitive to the amplitude and growth of
+images from the 6-m BTA telescope.,images from the 6-m BTA telescope.
+ Additionally. the HII region abundances and Ha rotation curves are being obtained on the William Herschel Telescope (WHT). Isaac Newton Telescope (INT) telescopes on La Palma and 6-m Russian BTA telescope. respectively.," Additionally, the HII region abundances and $\alpha$ rotation curves are being obtained on the William Herschel Telescope (WHT), Isaac Newton Telescope (INT) telescopes on La Palma and 6-m Russian BTA telescope, respectively."
+ This paper is organised as follows., This paper is organised as follows.
+ In Section 2. we describe the design and the properties of the galaxy sample., In Section \ref{sec:sample} we describe the design and the properties of the galaxy sample.
+ The main science drivers for FIGGS are described in Section.3.., The main science drivers for FIGGS are described in \ref{sec:science}.
+ The GMRT observations are described in Section 4 and the results of the survey are presented and discussed in Section 5.., The GMRT observations are described in Section \ref{sec:obs} and the results of the survey are presented and discussed in Section \ref{sec:result}.
+ The Faint Irregular. Galaxies GMRT Survey FIGGS. is a large observing program aimed at providing a comprehensive and statistically robust characterisation of the neutral ISM properties of faint. gas rich. dwarf irregular galaxies using the Giant Metrewave Radio Telescope (GMRT).," The Faint Irregular Galaxies GMRT Survey $-$ FIGGS, is a large observing program aimed at providing a comprehensive and statistically robust characterisation of the neutral ISM properties of faint, gas rich, dwarf irregular galaxies using the Giant Metrewave Radio Telescope (GMRT)."
+ The FIGGS sample forms a subsample of the Karachentsev et al.(2004) catalog of galaxies within  10 Mpc., The FIGGS sample forms a subsample of the Karachentsev et al.(2004) catalog of galaxies within $\sim$ 10 Mpc.
+ Specifically. the FIGGS sample consists of 65 faint dwarf irregular (dIrrygalaxies with: The sample choice was dictated by a balance between achieving the scientific. goals described in Section.3 and the practical limitations of the observing time.," Specifically, the FIGGS sample consists of 65 faint dwarf irregular (dIrr)galaxies with: The sample choice was dictated by a balance between achieving the scientific goals described in \ref{sec:science} and the practical limitations of the observing time."
+ We note that the above mentioned criterion on the optical B band major axis was not strictly followed in few cases., We note that the above mentioned criterion on the optical B band major axis was not strictly followed in few cases.
+ Some unusual. very faint dwarf galaxies. with optical B band major axis < | aremin were still included in our sample. as they are interesting cases to study in detail in HI.," Some unusual, very faint dwarf galaxies, with optical B band major axis $<$ 1 arcmin were still included in our sample, as they are interesting cases to study in detail in HI."
+ Further. for some of the galaxies in the FIGGS sample. fresh estimates of the distance (obtained after our observations were complete) imply absolute magnitudes slightly larger than the cut off above.," Further, for some of the galaxies in the FIGGS sample, fresh estimates of the distance (obtained after our observations were complete) imply absolute magnitudes slightly larger than the cut off above."
+ These galaxies have however been retained in the sample., These galaxies have however been retained in the sample.
+ Some properties (mainly derived from. optical observations) of galaxies in the FIGGS sample are listed in Table I.., Some properties (mainly derived from optical observations) of galaxies in the FIGGS sample are listed in Table \ref{tab:figgs_sample}.
+ The columns are as follows: Columntl) the galaxy name. &(3) the equatorial coordinates (2000). Column¢(+) the absolute blue magnitude (corrected for galactic extinction). Column(S) the Holmberg diameter in aremin. Column(6) the (B-V) colour. Columnt7) the distance in Mpc. Columnt8) the method used to measure the distance from the tip of the red giant branch (rgb). from membership in a group with known distance (erp). from the Tully-Fisher relation (t£). and from the Hubble flow (hy.," The columns are as follows: Column(1) the galaxy name, (3) the equatorial coordinates (J2000), Column(4) the absolute blue magnitude (corrected for galactic extinction), Column(5) the Holmberg diameter in arcmin, Column(6) the (B-V) colour, Column(7) the distance in Mpc, Column(8) the method used to measure the distance $-$ from the tip of the red giant branch (rgb), from membership in a group with known distance (grp), from the Tully-Fisher relation (tf), and from the Hubble flow (h)."
+ Columnt(9) gives the group membership of the galaxy. ColumntclO) the inclination determined from optical photometry (and assuming an intrinsic thickness. 450.) and Columntl1) the reference for the (B-V) colour. and/or revised distance.," Column(9) gives the group membership of the galaxy, Column(10) the inclination determined from optical photometry (and assuming an intrinsic thickness, $_o$ =0.2) and Column(11) the reference for the (B-V) colour, and/or revised distance."
+ The data presented in the Table |. (except for the colour) is taken from Karachentsevetal.(2004). catalog. except that revised distances have been adopted. if available.," The data presented in the Table \ref{tab:figgs_sample} (except for the colour) is taken from \cite{kk04} catalog, except that revised distances have been adopted, if available."
+ As can be seen from the Table |. tip of the red giant branch (rgb) distances (which are generally accurate to ~ 10569) are available for most of the galaxies in our sample.," As can be seen from the Table \ref{tab:figgs_sample}, tip of the red giant branch (rgb) distances (which are generally accurate to $\sim$ ) are available for most of the galaxies in our sample."
+ Figure | shows the histogram of the absolute blue magnitude (Mz). distance. HI mass. and HI mass to light ratio (Mai La) for the FIGGS sample. while Figure 2. compares the distributions/ of gas fraction. luminosity and dynamical mass of the FIGGS galaxies with that of existing samples of galaxies with HI aperture synthesis observations.," Figure \ref{fig:hist_figgs} shows the histogram of the absolute blue magnitude ${\rm{_B}}$ ), distance, HI mass, and HI mass to light ratio ${\rm{_{HI}/L_B}}$ ) for the FIGGS sample, while Figure \ref{fig:gasfrac} compares the distributions of gas fraction, luminosity and dynamical mass of the FIGGS galaxies with that of existing samples of galaxies with HI aperture synthesis observations."
+ The gas fraction and the dynamical masses for the FIGGS sample have been derived from the GMRT observations., The gas fraction and the dynamical masses for the FIGGS sample have been derived from the GMRT observations.
+ The FIGGS sample has a median Mis19 and a median HI mass ~35107 M.. while spanning range of more than 100 in stellar light. gas mass and dynamical mass. and more than + in gus fraction.," The FIGGS sample has a median ${\rm{M_B}} \sim -13$ and a median HI mass $\sim 3 \times 10^7$ $_\odot$, while spanning range of more than 100 in stellar light, gas mass and dynamical mass, and more than 4 in gas fraction."
+ It can also be clearly seen that by focusing on fainter. lower mass galaxies than those observed in previous HI imaging studies. FIGGS bridges the transition to rotation dominated low mass spirals and provides a substantially extended baseline in mass and luminosity space for a comparative study of galaxy properties.," It can also be clearly seen that by focusing on fainter, lower mass galaxies than those observed in previous HI imaging studies, FIGGS bridges the transition to rotation dominated low mass spirals and provides a substantially extended baseline in mass and luminosity space for a comparative study of galaxy properties."
+ The aim of FIGGS is to provide a large multi-wavelength database for a systematically selected sample of extremely faint dwarf irregular galaxies., The aim of FIGGS is to provide a large multi-wavelength database for a systematically selected sample of extremely faint dwarf irregular galaxies.
+ As mentioned in Section l.. such a database could be used to address a diverse range of astrophysical questions.," As mentioned in Section \ref{sec:intro}, such a database could be used to address a diverse range of astrophysical questions."
+ Rather than attempting to enumerate all of these. in this section. we deseribe in some detail a couple of key science drivers for the FIGGS survey.," Rather than attempting to enumerate all of these, in this section, we describe in some detail a couple of key science drivers for the FIGGS survey."
+ One of the main goals of FIGGS is to use the HI interferometric images in conjunction with the optical data to study the interplay between the neutral ISM and star formation in the faintest. lowest mass. gas rich dIrr galaxies.," One of the main goals of FIGGS is to use the HI interferometric images in conjunction with the optical data to study the interplay between the neutral ISM and star formation in the faintest, lowest mass, gas rich dIrr galaxies."
+ The gravitational binding energy for very faint dwarf irregular galaxies is not much larger than the energy output from a few supernovae: consequently star formation in such galaxies could have a profound effect on the morphology and kinematics of the ISM of these systems., The gravitational binding energy for very faint dwarf irregular galaxies is not much larger than the energy output from a few supernovae; consequently star formation in such galaxies could have a profound effect on the morphology and kinematics of the ISM of these systems.
+ The FIGGS data will enable us to study the ISM of most of our sample galaxies at a linear resolution of ~15300 pe Le. comparable to the scales at which energy is injected into the ISM through supernova and stellar winds., The FIGGS data will enable us to study the ISM of most of our sample galaxies at a linear resolution of $\sim15-300$ pc $-$ i.e. comparable to the scales at which energy is injected into the ISM through supernova and stellar winds.
+ FIGGS thus provide a unique opportunity to study the effects of feedback from star formation in low mass. gas rich galaxies. which in turn will allow us to understand the processes driving the evolution of these gulaxies.," FIGGS thus provide a unique opportunity to study the effects of feedback from star formation in low mass, gas rich galaxies, which in turn will allow us to understand the processes driving the evolution of these galaxies."
+ For example. it has been suggested that star formation in. dwarf galaxies occurs only above a constant threshold HI column density of Nyy~107 7 (e.g. Skillman1987:Taylor 19949).," For example, it has been suggested that star formation in dwarf galaxies occurs only above a constant threshold HI column density of ${\rm{_{HI}}}\sim10^{21}$ $^{-2}$ (e.g. \citealp{skillman87,taylor94}) )."
+ Such : threshold could be a consequence of disk dynamics (e.g. related to Toomre's instability criterion: Kennicutt (1989))) or a consequence of some other physical process. e.g. self shielding or thermo-gravitational instability (Schaye (200439).," Such a threshold could be a consequence of disk dynamics (e.g. related to Toomre's instability criterion; \cite{rob89}) ) or a consequence of some other physical process, e.g. self shielding or thermo-gravitational instability \cite{schaye04}) )."
+ A preliminary study of a small subsample of FIGGS (Begumetal. (2006))) suggested that unlike brighter dwarfs. the faintest dwarf galaxies do not show well detined threshold density.," A preliminary study of a small subsample of FIGGS \cite{b06}) ) suggested that unlike brighter dwarfs, the faintest dwarf galaxies do not show well defined threshold density."
+ A detailed comparison of Ha and UV images with HI column density maps for the FIGGS sample will allow us to definitively answer the issue of the existence of a threshold density in the faintest galaxies and also to check whether the recipes for star formation derived from larger galaxies (Kennicutt (1989))) continue to be valid at this mass regime., A detailed comparison of $\alpha$ and UV images with HI column density maps for the FIGGS sample will allow us to definitively answer the issue of the existence of a threshold density in the faintest galaxies and also to check whether the recipes for star formation derived from larger galaxies \cite{rob89}) ) continue to be valid at this mass regime.
+ These are critical issues in hierarchical galaxy formation models., These are critical issues in hierarchical galaxy formation models.
+ The second major aim of this survey is to study the relation between dark and baryonic matter in the smallest known star. forming galaxies., The second major aim of this survey is to study the relation between dark and baryonic matter in the smallest known star forming galaxies.
+ According to several models of galaxy formation and evolution. the first burst of star formation in dwarf galaxies below a," According to several models of galaxy formation and evolution, the first burst of star formation in dwarf galaxies below a"
+so that The above expression in the round brackets is the denominator of «yh.ΛΑ.0).,"so that The above expression in the round brackets is the denominator of $\omega_\mathbf{x}(h,k,\Omega)$."
+ Further simplification vields, Further simplification yields
+a set ol objects in two dimensions (hat have a radial distribution given by equation 1. (,a set of objects in two dimensions that have a radial distribution given by equation 1. (
+3) Calculate the £ and (e of each object as seen from the Sun al R.=Skpc. (,"3) Calculate the $\ell$ and $v$ of each object as seen from the Sun at $\rsol = 8\,\mathrm{kpc}$. ("
+4) Generate a random Gaussian deviate lor each object from a distribution with o=σι and zero mean. and adel it (to each ο. (,"4) Generate a random Gaussian deviate for each object from a distribution with $\sigma \, = \, \sigma_v$ and zero mean, and add it to each $v$. ("
+5) Compare the resulting set of random (6.0) values to the catalog ol actual molecular clouds. using a two-dimensional Ixolmogorov-Smirnov test 1992).. (,"5) Compare the resulting set of random $(\ell, v)$ values to the catalog of actual molecular clouds, using a two-dimensional Kolmogorov-Smirnov test \citep{press92}. ("
+6) Go back to step 1. varving parameters (o maximize the similarity between (he random set of objects and the observed catalog.,"6) Go back to step 1, varying parameters to maximize the similarity between the random set of objects and the observed catalog."
+" This procedure results in the parameters OQ,=215kms!. Ro=LLTkpc. a=2.3. and c,=vkmsFH."," This procedure results in the parameters $\Theta_0 \, = \, 215 \, \kms$, $R_\mathrm{s} \, = \, 1.7 \, \mathrm{kpc}$, $\alpha = 2.3$, and $\sigma_v \,=\, 7.7\,\kms$."
+" This value of O, should not be taken to be a measure of the Suns rotational velocity.", This value of $\Theta_0$ should not be taken to be a measure of the Sun's rotational velocity.
+" If the rotation curve were parameterized as O(2)&Oy+(Η/Μ.)O44.... then O, cannot be constrained bv fitting to objects in the inner Galaxy: galactic kinematics looks (he same wilh a solic-body term added to the rotation curve."," If the rotation curve were parameterized as $\Theta(R) \,\approx\, \Theta_0 \,+\, (R/\rsol)\Theta_1 \,+\,\ldots\,$, then $\Theta_1$ cannot be constrained by fitting to objects in the inner Galaxy: galactic kinematics looks the same with a solid-body term added to the rotation curve."
+ These data do not constrain a verv well either. since our cutoff at 6=20° eliminates much of the molecular hole at Rox3.5kpe.," These data do not constrain $\alpha$ very well either, since our cutoff at $\ell\, = \, 20\de $ eliminates much of the molecular hole at $R\,\approx\,3.5\,\mathrm{kpc}$."
+ The shape of the inner edge of this hole is parameterized by a. and almost all values of a between 2 and 3 fit well.," The shape of the inner edge of this hole is parameterized by $\alpha$, and almost all values of $\alpha$ between 2 and 3 fit well."
+ The quick drop-oll of themolecular ring to larger radi. Ru=Lakpe. is (vpical of fits (o first quadrant molecular line survevs (Burton&Solomonetal. 1979).," The quick drop-off of themolecular ring to larger radii, $R_\mathrm{s} \, = \, 1.7 \, \mathrm{kpc}$, is typical of fits to first quadrant molecular line surveys \citep{burton78,solomon79}."
+". The value of σι=7.7kims! is consistent with the value o,=—7.80.6jskms+ determined: forJ moderately-large↽ clouds near the Sun by- Brand (1989).."," The value of $\sigma_v \,=\, 7.7\,\kms$ is consistent with the value $\sigma_v \,=\, 7.8^{+0.6}_{-0.5}\,\kms$ determined for moderately-large clouds near the Sun by \citet{stark89b}."
+ The distribution of [.Ne| values for the Monte Carlo objects is plotted in green in figure 2. and their median value of |Nel is T.81220.05kms. !.," The distribution of $\dv$ values for the Monte Carlo objects is plotted in green in figure 2, and their median value of $|\Delta v|$ is $7.81\pm0.05\,\kms$ ."
+. This value is robust. in the sense that essentially any azimuthallv-symmetrie distribution of random objects in the ealactic plane that we have tried will produce a median value of νο that is greater than Tkms.!.," This value is robust, in the sense that essentially any azimuthally-symmetric distribution of random objects in the galactic plane that we have tried will produce a median value of $|\Delta v|$ that is greater than $7\,\kms$."
+ The small molecular clouds are more concentrated to spiral arms (han random objects. and the GAICSs are even more concentrated than the small clouds.," The small molecular clouds are more concentrated to spiral arms than random objects, and the GMCs are even more concentrated than the small clouds."
+ lt is unfortunate that we cannot observe the Milkv. Way [rom the outside. and see the spatial relationship between the molecular clouds and the spiral arms directly.," It is unfortunate that we cannot observe the Milky Way from the outside, and see the spatial relationship between the molecular clouds and the spiral arms directly."
+ Interferometric observations of CO in other galaxies having well-defined spiral structure shows that the CO enission is concentrated (ο the spiral arms (Murgiaetal.2005)., Interferometric observations of CO in other galaxies having well-defined spiral structure shows that the CO emission is concentrated to the spiral arms \citep{murgia05}.
+. In the Milky Was. most ol the CO emission is from GMCs.," In the Milky Way, most of the CO emission is from GMCs."
+ This is illustrated in figure 3., This is illustrated in figure 3.
+ IIere we have taken the calalog subset discussed in Stark&Lee(2005).. where the cloucls were selected to have distances. and huther restricted. (hat subset to clouds with 6> 20°.," Here we have taken the catalog subset discussed in \citet{stark05c}, where the clouds were selected to have well-determined distances, and further restricted that subset to clouds with $\ell\,>\,20\de$ ."
+ Figure 3 shows the distributionof virial masses lor (hese clouds., Figure 3 shows the distributionof virial masses for these clouds.
+ The vast majority of the CO-emitting, The vast majority of the CO-emitting
+"much higher temperatures in the inner part; e.g. at 2000 AU from the region Bl, diffusion raises the temperature from 150 K in a low-density model, where diffusion is not effective, to 450 K in the high-density model for G10.47--0.03.","much higher temperatures in the inner part; e.g. at 2000 AU from the region B1, diffusion raises the temperature from 150 K in a low-density model, where diffusion is not effective, to 450 K in the high-density model for G10.47+0.03."
+ This mechanism therefore seems indispensable to account for large masses of hot gas., This mechanism therefore seems indispensable to account for large masses of hot gas.
+" Additionally, by absorbing stellar UV radiation the dust protects the molecules from dissociation."," Additionally, by absorbing stellar UV radiation the dust protects the molecules from dissociation."
+" In general, heating is more efficient if the heating source is inside a density condensation (?) than if itis external, as could be the case for G34.26+0.15 (??).."," In general, heating is more efficient if the heating source is inside a density condensation \citep{Kaufman98} than if it is external, as could be the case for G34.26+0.15 \citep{Watt99, Mookerjea07}."
+" In our models, the diffusion mechanism is accomplished by dense spherical clumps around heating sources."," In our models, the diffusion mechanism is accomplished by dense spherical clumps around heating sources."
+" It is however not clear how many heating sources there are: The stars in theHir regions are sufficient to reproduce the observed line strength, as the presented models demonstrate."," It is however not clear how many heating sources there are: The stars in the regions are sufficient to reproduce the observed line strength, as the presented models demonstrate."
+" But a good fit can also be achieved when lower-luminosity stars are added to the models according to the Initial Luminosity Function, without changing the stars in the regions."," But a good fit can also be achieved when lower-luminosity stars are added to the models according to the Initial Luminosity Function, without changing the stars in the regions."
+ So on the basis of these data it seems not possible to distinguish between the scenarios of heating only by the stars ionizing the regions and heating by regions plus embedded sources which do not emit in the radio regime., So on the basis of these data it seems not possible to distinguish between the scenarios of heating only by the stars ionizing the regions and heating by regions plus embedded sources which do not emit in the radio regime.
+" We may speculate that there are much more stars present than detected from the free-free radiation, since the observed sources are clusters in formation."," We may speculate that there are much more stars present than detected from the free-free radiation, since the observed sources are clusters in formation."
+" Also, there could be more extended (faint) line emission (and the absorption coming from a larger column) than in the models, which would point to additional heating sources."," Also, there could be more extended (faint) line emission (and the absorption coming from a larger column) than in the models, which would point to additional heating sources."
+" Furthermore, the simple models predict strong in the (optically thick) submm rotational transitions within the v? 1 state of HCN and H?CN, which is however not observed with the SMA, presumably due to clumpiness (Rolffs et al."," Furthermore, the simple models predict strong in the (optically thick) submm rotational transitions within the $v_2$ =1 state of HCN and $^{13}$ CN, which is however not observed with the SMA, presumably due to clumpiness (Rolffs et al."
+" 2011, in prep.)."," 2011, in prep.)."
+" While not excluding the possibility of additional heating sources, we could show that heating by the stars in the hypercompactHu regions is sufficient to reproduce the observations."," While not excluding the possibility of additional heating sources, we could show that heating by the stars in the hypercompact regions is sufficient to reproduce the observations."
+" Hence, our modeling results suggest that it is unlikely that the main heating is due to embedded objects without ionized gas."," Hence, our modeling results suggest that it is unlikely that the main heating is due to embedded objects without ionized gas."
+" Given that the sources are dense condensations deeply embedded in giant molecular clouds and have high column densities, diffusion of radiation, which is necessary to increase the heating, is reasonable even for more inhomogeneous and asymmetric structures than modeled."," Given that the sources are dense condensations deeply embedded in giant molecular clouds and have high column densities, diffusion of radiation, which is necessary to increase the heating, is reasonable even for more inhomogeneous and asymmetric structures than modeled."
+28. 2009 (UT).,"28, 2009 (UT)."
+ Conditions on the latter night were clear with a stable seeing of 076 at A-band., Conditions on the latter night were clear with a stable seeing of $\farcs$ 6 at $K$ -band.
+ The SpeX prism mode provides 0.75-2.5 jm continuous spectroscopy with resolution z120 for the 075 slit employed (dispersion across the chip is 20-30A !'), The SpeX prism mode provides 0.75–2.5 $\micron$ continuous spectroscopy with resolution $\approx 120$ for the $\farcs$ 5 slit employed (dispersion across the chip is 20–30 $^{-1}$ ).
+ Both components were observed separately. with the slit oriented. north-south. roughly aligned with the parallactic angle and perpendicular to the separation axis.," Both components were observed separately, with the slit oriented north-south, roughly aligned with the parallactic angle and perpendicular to the separation axis."
+ For the primary. eight exposures of 90s each were obtained at an average airmass of 2.33. while guiding on spillover light from the slit.," For the primary, eight exposures of 90s each were obtained at an average airmass of 2.33, while guiding on spillover light from the slit."
+ For the secondary. eight exposures of 150s each were obtained at an average armass of 2.34. while guiding on the primary off-slit.," For the secondary, eight exposures of 150s each were obtained at an average airmass of 2.34, while guiding on the primary off-slit."
+ For both sources. the AO V star HD 8977 was observed immediately before the target for telluric and flux calibration while the quartz and Ar are lamps were observed for flat field and wavelength calibration. respectively.," For both sources, the A0 V star HD 8977 was observed immediately before the target for telluric and flux calibration while the quartz and Ar arc lamps were observed for flat field and wavelength calibration, respectively."
+ Data were reduced using the SpeXtool package. version 3.4 (2?) using standard settings: see ? for details.," Data were reduced using the SpeXtool package, version 3.4 \citep{Vacca2003, Cushing2004} using standard settings; see \citet{Burgasser2007d} for details."
+ Figure 2. shows the spectra of the two components of 4445AB; signal-to-noise at the 7/7AN flux peaks was 100—150 and 25-35 for the A and B components. respectively.," Figure \ref{Fig: nirspec} shows the spectra of the two components of AB; signal-to-noise at the $JHK$ flux peaks was 100–150 and 25–35 for the A and B components, respectively."
+ Both spectra. show the characteristic infrared (NIR) features of late-type M and L dwarfs 222): steep red optical slopes (0.8—1.0 j/m) from the pressure broadened wing of the 0.77 jim doublet: molecular absorption bands arising from (1.4 and 1.9 jm). CO (2.3 yam). and FeH (0.99 jim); and an overall red spectral energy distribution (SED). consistent with the photometric colors.," Both spectra show the characteristic near-infrared (NIR) features of late-type M and L dwarfs \citep[e.g.,][]{Reid2001a, McLean2003,
+ Cushing2005}: steep red optical slopes (0.8–1.0 $\micron$ ) from the pressure broadened wing of the 0.77 $\micron$ doublet; molecular absorption bands arising from (1.4 and 1.9 $\micron$ ), CO (2.3 $\micron$ ), and FeH (0.99 $\micron$ ); and an overall red spectral energy distribution (SED), consistent with the photometric colors."
+ 4445A also exhibits additional absorption features in the 0.8—1.2 j/m region arising from TiO. VO. FeH. and unresolved and lines. all typical for a late-type M dwarf.," A also exhibits additional absorption features in the 0.8–1.2 $\micron$ region arising from TiO, VO, FeH, and unresolved and lines, all typical for a late-type M dwarf."
+ The corresponding region in the spectrum of 4445B ts considerably smoother. albeit more noisy. suggesting that many of these gaseous species have condensed out (e.g..??)..," The corresponding region in the spectrum of B is considerably smoother, albeit more noisy, suggesting that many of these gaseous species have condensed out \citep[e.g.,][]{Tsuji1998, Ackerman2001}."
+ The appearance of weak absorption at 1.15 sm and the very red SED of the NIR spectrum all indicate that 4445B ts a mid- to late-type L dwarf with relatively thick condensate clouds at the photosphere., The appearance of weak absorption at 1.15 $\micron$ and the very red SED of the NIR spectrum all indicate that B is a mid- to late-type L dwarf with relatively thick condensate clouds at the photosphere.
+ To determine spectral types we compared our NIR spectra of 4445AB with 463 spectra of 439 M7 or later dwarfs from the SpeX Prism Spectral., To determine spectral types we compared our NIR spectra of AB with 463 spectra of 439 M7 or later dwarfs from the SpeX Prism Spectral.
+Libraries?.. All templates wre chosen to have median S/N > 10 and could not be binaries. giants. subdwarfs. or spectral classifications that were peculiar or uncertain.," All templates wre chosen to have median S/N $>$ 10 and could not be binaries, giants, subdwarfs, or spectral classifications that were peculiar or uncertain."
+ Best matches were determined by finding the minimum deviation between component spectra and templates in the 0.95—1.35. 1.45-1.80. and 2.00-2.35 jim regions {1.δ.. avoiding telluric bands). following the procedure of ? with no pixel weighting.," Best matches were determined by finding the minimum deviation between component spectra and templates in the 0.95--1.35, 1.45–1.80, and 2.00–2.35 $\micron$ regions (i.e., avoiding telluric bands), following the procedure of \citet{Cushing2008} with no pixel weighting."
+ The two best matching templates to both components of 4445AB are shown in Figure 3.., The two best matching templates to both components of AB are shown in Figure \ref{Fig: nirspec_temp}. .
+ For 4445A. these are the optically classified M8 dwarf 2MASS 4105172730 3348593 (?) and LO dwarf DENIS-P 253450 (?). for 4445B these are the optically classified LS dwarf 2MASSW 272937 (?) and L7 dwarf 2MASS 3421292 (?)..," For A, these are the optically classified M8 dwarf 2MASS $-$ 3348593 \citep{Cruz2003} and L0 dwarf DENIS–P $-$ 253450 \citep{Phan-Bao2008}; for B these are the optically classified L5 dwarf 2MASSW $-$ 272937 \citep{Gizis2002} and L7 dwarf 2MASS $-$ 3421292 \citep{Kirkpatrick2008}."
+ Note that despite the differences in optical type. these spectra provide equivalently good fits—the two fits were different only by 1.75 for the primary and l|.]o for the secondary based on the F-test which gauges whether two different fits to data are significantly distinct based on the ratio of values and degrees of freedom (?)-—to the components of 4445AB. a reflection of the discrepancies between optical and NIR spectral morphologies for late- M and L dwarfs (e.g.22).," Note that despite the differences in optical type, these spectra provide equivalently good fits—the two fits were different only by $\sigma$ for the primary and $\sigma$ for the secondary based on the F-test which gauges whether two different fits to data are significantly distinct based on the ratio of values and degrees of freedom \citep{Burgasser2010a}- —to the components of AB, a reflection of the discrepancies between optical and NIR spectral morphologies for late-type M and L dwarfs \citep[e.g.,][]{Geballe2002, Kirkpatrick2005}."
+ A 4? weighted mean of all the SpeX templates (e.g..2?)— indicates component types of M9.0+0.5 for 4445A and L6+1 for 4445B. We also derived classifications using a suite of spectral indices and spectral index/spectral type relations from ?.. 9.. 2.. ?.. and ?..," A $\chi^2$ weighted mean of all the SpeX templates \citep[e.g.,][]{Burgasser2010a} indicates component types of $\pm$ 0.5 for A and $\pm$ 1 for B. We also derived classifications using a suite of spectral indices and spectral index/spectral type relations from \citet{Tokunaga1999}, \citet{Reid2001a}, \citet{Geballe2002}, \citet{Burgasser2006b}, and \citet{Burgasser2007c}."
+ Table 2. shows the measured values and the inferred spectral subtypes of 4445A and4445B for each of the indices.," Table \ref{Tab:
+ indices} shows the measured values and the inferred spectral subtypes of A andB for each of the indices."
+ The mean and scatter from these indices yieldsclassifications of LO.S+1.0 for 4445A and L7.0+1.5 for 4445B. These are consistent with. but less precise than. the types inferred from spectral. template matching. so we adopt the latter for our subsequent analysis.," The mean and scatter from these indices yieldsclassifications of $\pm$ 1.0 for A and $\pm$ 1.5 for B. These are consistent with, but less precise than, the types inferred from spectral template matching, so we adopt the latter for our subsequent analysis."
+is altributed to the asvinmetrie diit of nearby stars. discussed in more detail in the next paragraph.,"is attributed to the asymmetric drift of nearby stars, discussed in more detail in the next paragraph."
+" Generally. there are (three clilferent reasons of physical and technical kind for our estimation of Vy to be biased: As long as a relatively small local area of the Galaxy around the Sun is considered (r<< Ry. where Jt, is (he ealactocentric distance of (he Sun). it is appropriate to expand the svstemic velocity field of stars in a Taylor series over the galactocentric evlindrical coordinates (p.0.z)."," Generally, there are three different reasons of physical and technical kind for our estimation of $V_Y$ to be biased: As long as a relatively small local area of the Galaxy around the Sun is considered $r << R_0$ , where $R_0$ is the galactocentric distance of the Sun), it is appropriate to expand the systemic velocity field of stars in a Taylor series over the galactocentric cylindrical coordinates $(\rho, \theta, z)$."
+ Assuming that the local field is planar. that is. all svstemic motions are in the galactic plane. the velocity vector o(p.8.2)=atp.8)-op.0. z). where the former component (0s radial wilh respect to the Galactic center. and the latter. à is tangential and orthogonal to the former.," Assuming that the local field is planar, that is, all systemic motions are in the galactic plane, the velocity vector $\vec{v}(\rho, \theta, z)=\vec{a}(\rho,
+\theta)+\vec{\omega}(\rho, \theta, z)$ , where the former component $\vec{a}$ is radial with respect to the Galactic center, and the latter, $\vec{\omega}$ is tangential and orthogonal to the former."
+ Note that the radial component is assumed to be independent of z. (hat is. that there is no radial svstemic motion depending on the distance from the plane.," Note that the radial component is assumed to be independent of $z$, that is, that there is no radial systemic motion depending on the distance from the plane."
+ Retaining ouly first-degree terms in the corresponding Tavlor expansion. one can write where the subscript Q denotes Che corresponding parameters al the Sun's location (po.85. 24).," Retaining only first-degree terms in the corresponding Taylor expansion, one can write where the subscript 0 denotes the corresponding parameters at the Sun's location $(\rho_0, \theta_0, z_0)$ ."
+ We left the dependence of the rotational velocity & on z in ils generic form. since on physical grounds. (his dependence is expected to be sanmetric around (he plane. and can not be represented by a simple linear term.," We left the dependence of the rotational velocity $\omega$ on $z$ in its generic form, since on physical grounds, this dependence is expected to be symmetric around the plane, and can not be represented by a simple linear term."
+ Retaining only terms to O(—-). these model relations can be rewritten moreconveniently in theheliocentric coordinates (r.(6) introducedin Appendix. A:," Retaining only terms to $O(\frac{r}{\rho_0})$, these model relations can be rewritten moreconveniently in theheliocentric coordinates $(r, \ell, b)$ introducedin Appendix A:"
+The possibility of the central X-ray cinissiou being larecly a svuchrotron nebula. as cousidere bv ISISS. is ruled out.,"The possibility of the central X-ray emission being largely a synchrotron nebula, as considered by KKS, is ruled out."
+ Furthermore. since the absorbing columns toward the X-rav οιτας regions are approxiuatoelv equal to the tota coluun densities of hvdrogen nuclei in those directions. the N-ravs. probably originate behliuk most of the atomic and molecular sas iu the Sagittarius arm.," Furthermore, since the absorbing columns toward the X-ray emitting regions are approximately equal to the total column densities of hydrogen nuclei in those directions, the X-rays probably originate behind most of the atomic and molecular gas in the Sagittarius arm."
+ The results therefore are compatible with W51C being a single. isotherma SNR at the far side of the Sagittarius aru.," The results therefore are compatible with W51C being a single, isothermal SNR at the far side of the Sagittarius arm."
+" At a distance of 6 kpc. the linear extent of NIC is 88&66 pe. which makes it one of the largest SNRs iu the Galaxy (οιοι, Strom 1996)."," At a distance of 6 kpc, the linear extent of W51C is $88\times 66$ pc, which makes it one of the largest SNRs in the Galaxy (e.g., Strom 1996)."
+ We have derived the parameters of the central region: au elliptical (22« 15’) area surrounding regions ANN and NS., We have derived the parameters of the central region; an elliptical $23'\times 15'$ ) area surrounding regions XN and XS.
+" The root mean. voluuc-averaged. square density a.=(fidV/V)~0.58 7. the N-vay enüttiug mass Afy=~290 AL... and the thermal cnerev Ey,&0.39<1075 eres,"," The root mean, volume-averaged, square density $\bar n_e\equiv \left( \int n_e^2 dV/V \right)^{1/2}\simeq 0.58$ $^{-3}$, the X-ray emitting mass $M_X\simeq 290~M_\odot$ , and the thermal energy $E_{\rm th}\simeq 0.39\times 10^{51}$ ergs."
+" For comparison. WINS derived η2%0.29 3 Aly~1900AL... and Ey,~2.6«10° eres for the whole (50ς 38’) SNR."," For comparison, KKS derived $\bar n_e\simeq 0.29$ $^{-3}$, $M_X\simeq 1900~M_\odot$, and $E_{\rm th}\simeq 2.6\times 10^{51}$ ergs for the whole $50'\times 38'$ ) SNR."
+ If the N-rav cutting eas is chuupy with a volume flline factor of f£. the true electron deusitv η. would be ereater by a factor of FO? whereas the N-vay cutting mass would be zualler by a factor of £172.," If the X-ray emitting gas is clumpy with a volume filling factor of $f$ , the true electron density $n_e$ would be greater by a factor of $f^{-1/2}$, whereas the X-ray emitting mass would be smaller by a factor of $f^{1/2}$."
+ Thermal energy. could be greater (f L7) or smaller (£57) depending ou whether the iuterchunp region is filled for the SNR to be isobaric or not., Thermal energy could be greater $f^{-1/2}$ ) or smaller $f^{1/2}$ ) depending on whether the interclump region is filled for the SNR to be isobaric or not.
+ Frou the nuage (Fie., From the image (Fig.
+" 1). we estimate that the volumuec filling factor of the N-rav cinitting gas f&OL ‘or the whole SNR. so that »,20.16 cm”. Aly£1200AL... aud Ey,2LLx1075 eres or Ey,21.6«109 eres."," 1), we estimate that the volumne filling factor of the X-ray emitting gas $f\simlt 0.4$ for the whole SNR, so that $n_e\simgt 0.46$ $^{-3}$, $M_X\simlt 1200~M_\odot$, and $E_{\rm th}\simgt 4.1\times 10^{51}$ ergs or $E_{\rm th}\simlt 1.6\times 10^{51}$ ergs."
+" ISISS estimated the age of he SNR at ~23«104 sys,", KKS estimated the age of the SNR at $\sim 3\times 10^4$ yrs.
+ As a single SNR. W51C has a composite nuorphologv: a very thick (~13’ or 23 pc) outer radio and X-ray shell aud au A-ravbright ceutral region.," As a single SNR, W51C has a composite morphology: a very thick $\sim 13'$ or 23 pc) outer radio and X-ray shell and an X-ray–bright central region."
+ There are about 15 other composite SNRs. correspoucding to 25% of N-ravdetected Galactic SNRs (Rho&Petre1905).," There are about 15 other composite SNRs, corresponding to $\sim 25$ of X-ray–detected Galactic SNRs \citep{rho98}."
+ Several models have been suggested. which iuclude the evaporation cloud model of White&Long and the heat conduction model of Sheltonetal. (19993.," Several models have been suggested, which include the evaporation cloud model of \citet{whi91} and the heat conduction model of \citet{she99}."
+.. The former assunes small deuse clouds evaporating within the hot iuterior. while the latter considers oulv heat conduction within the hot interior.," The former assumes small dense clouds evaporating within the hot interior, while the latter considers only heat conduction within the hot interior."
+ Observatiouallv. many of these SNRs me known to be interacting with large molecular clouds (see Rho Petre 1998 and references therein).," Observationally, many of these SNRs are known to be interacting with large molecular clouds (see Rho Petre 1998 and references therein)."
+ It has been shown that W51C is interacting with a molecular cloud along the western boundary of region XN (soo&Moou1997h).. and some. if uot all. of the N-ray euüttiug eas im the central area could have been originated from this molecular cloud.," It has been shown that W51C is interacting with a molecular cloud along the western boundary of region XN \citep{koo97b}, and some, if not all, of the X-ray emitting gas in the central area could have been originated from this molecular cloud."
+ Since W51C is located iu the Sagittarius ari and is interacting with a molecular cloud. is likely to be a type II SNR.," Since W51C is located in the Sagittarius arm and is interacting with a molecular cloud, is likely to be a type II SNR."
+ According to our results. however. most of the soft N-rav enüssion in the ceutral region is thermal. certainly not consistent with a bright pulsar wind uchula.," According to our results, however, most of the soft X-ray emission in the central region is thermal, certainly not consistent with a bright pulsar wind nebula."
+ Among hard X-rav sources. TIN? and TN3-west are wuideutified and could be candidates for a weak pulsar nebula.," Among hard X-ray sources, HX2 and HX3-west are unidentified and could be candidates for a weak pulsar nebula."
+ Their spectra can be fitted by a power-law model ↖∏↑↓↕⇁∶−≻∙⋅↱⊐∩↴⋃⋡⊐∩⋜⋯≼↧↓∙∣↴∖≺∪↥⋅∿−∙∣⋟∙⋅ ⊀⋃⊳yr2. ↽↼⊥⊳⋃ ⊀↼ respectively. although they can be fitted equally well by a hermal emission model.," Their spectra can be fitted by a power-law model with $\Gamma=2.59^{+0.27}_{-0.29}$ and $1.7^{+1.0}_{-\infty}$ (or $\simlt 2.7$ ), respectively, although they can be fitted equally well by a thermal emission model."
+ These plotou indices fall into the rauge (1.62.5) of the indices of known pulsar nebulae (Seward&Wang1988)., These photon indices fall into the range (1.6–2.5) of the indices of known pulsar nebulae \citep{sew88}.
+. If we extraploate their spectra to low euergies by the same power lav. them N-rav hDpnuuimosities iu the band (0.2| keV) would be 0. 2.107! ores assunune d=6 kpe.," If we extraploate their spectra to low energies by the same power law, their X-ray luminosities in the band (0.2–4 keV) would be $\sim$ $\times 10^{34}$ ergs $^{-1}$ assuming $d=6$ kpc."
+ This is ZOLL of Crab nebula aud is not unreasonable as a luminosity of a ~3«105 ves old pulsar (Sew&Wang 1988)., This is $\simlt 0.1$ of Crab nebula and is not unreasonable as a luminosity of a $\sim 3\times 10^4$ yrs old pulsar \citep{sew88}.
+. Recent liscoveries also indicate that many voung pulsars and pulsar wine nebulac do not have high bhuninosities. (Olbertetal.2001:IIughes 2001).," Recent discoveries also indicate that many young pulsars and pulsar wind nebulae do not have high luminosities \citep{olb01, hug01}."
+. μοι spatia resolution is needed in order to explore the uature of these sources aud to look for a faint pulsa iux its nebula., High spatial resolution is needed in order to explore the nature of these sources and to look for a faint pulsar and its nebula.
+ We have detected hard X-ray sources coincident with several compact II IT regions in W5l., We have detected hard X-ray sources coincident with several compact H II regions in W51.
+ The hard image (Fie., The hard image (Fig.
+ 2c) implies that W51À welt be the brightest lard N-ray source among he W51 star-forming regious. but the source is located at the edge of the CIS fold and ifs properties could not be obtained.," 2c) implies that W51A might be the brightest hard X-ray source among the W51 star-forming regions, but the source is located at the edge of the GIS field and its properties could not be obtained."
+ Among heW51D sources. two compact IT II reeious. 0.3 and 0.3. are coincident with the tightest hard N-rav source UNL.," Among theW51B sources, two compact H II regions, $-$0.3 and $-$ 0.3, are coincident with the brightest hard X-ray source HX1."
+ The brightest radio contimuni source in WB51ID. 0.3. Is also detected as IENS-cast. but is much. fainter.," The brightest radio continuum source in W51B, $-$ 0.3, is also detected as HX3-east, but is much fainter."
+ These X-rays uuelit be from early-type stars, These X-rays might be from early-type stars
+explanation for the scaling of σι.,explanation for the scaling of $\sigma_{l}$.
+" To this eud. we insert spatially uncorrelated scatter with a).=0.5 ou the scale of 100pe aud compute σι as a function of averaging scale for log-noriual deusity distributions of various widths o, aud for a 1D. 2D and 3D configuration of Πω,"," To this end, we insert spatially uncorrelated scatter with $\sigma{}_{l_{\rm min}}=0.5$ on the scale of $100$pc and compute $\sigma_{l}$ as a function of averaging scale for log-normal density distributions of various widths $\sigma_\rho$ and for a 1D, 2D and 3D configuration of $\H2$."
+" Specifically. for cach scale 7 we first generate (1a) independent density values (d is the assumed dimensionality of the eas configuration) drawn from a loe-norimal distribution with width a, ind the same muuber of correspoudius star formation rate density values according to equation (D2))."," Specifically, for each scale $l$ we first generate $(l/l_{\rm min})^d$ independent density values $d$ is the assumed dimensionality of the gas configuration) drawn from a log-normal distribution with width $\sigma_\rho$ and the same number of corresponding star formation rate density values according to equation \ref{eq:SFRAlt}) )."
+" We then compute the average Ho density and SER over the 1ἶμμιο ποσοκο, ", We then compute the average $\H2$ density and SFR over the $(l/l_{\rm min})^d$ “cells”.
+This pair of spatially averaged Που density aud star formation rate density constitutes a data point inthe relation aud the scatter of the relation is computed frou 1000 data points eenerated in this wav., This pair of spatially averaged $\H2$ density and star formation rate density constitutes a data point in the relation and the scatter of the relation is computed from 1000 data points generated in this way.
+ Fig., Fig.
+" Br shows that if 5, is sufficicutly small (<<1). 6j scales as VU as anticipated."," \ref{fig:scaleScatterTest} shows that if $\sigma_\rho$ is sufficiently small $\ll{}1$ ), $\sigma_{l}$ scales as $V^{-0.5}$ as anticipated."
+" Supersonic turbulence sinulatious. however. predict =lub|024423172 with M~dh10 aud 3~0.250,5 (McI&ee&Ostriker2007).. "," Supersonic turbulence simulations, however, predict $\sigma_\rho=\ln(1+\beta^2\mathcal{M}^2)^{1/2}$ with $\mathcal{M}\sim{}5-10$ and $\beta\sim{}0.25-0.5$ \citep{2007ARA&A..45..565M}. ."
+Thus. συ%1.5 is a more reasonable assumption.," Thus, $\sigma_\rho{}\approx{}1.5$ is a more reasonable assumption."
+ In this case. σι scales approxiuately x70 M the gas is arvanece iu a disk.," In this case, $\sigma_{l}$ scales approximately $\propto{}l^{-0.5}$ if the gas is arranged in a disk."
+The uncertainties in the red. and blue spectra. should ο statistically independent.,The uncertainties in the red and blue spectra should be statistically independent.
+ Several features are immediately. apparent from. these spectra., Several features are immediately apparent from these spectra.
+ The emission. line spectrum is. dominated. by broad Balmer lines (lla. Lb. Le and 1190) and narrow forbidden oxvgen lines OH]A3727 anc OLWJAL959/5007).," The emission line spectrum is dominated by broad Balmer lines $\alpha$ , $\beta$ , $\gamma$ and $\delta$ ) and narrow forbidden oxygen lines $\lambda
+3727$ and $\lambda 4959/5007$ )."
+ The emission line spectrum. which is clearly related to the Seyfert activity. is examined in more detail below (Section 2.1.2).," The emission line spectrum, which is clearly related to the Seyfert activity, is examined in more detail below (Section 2.1.2)."
+ The spectrum between παπα mimay be severely alfected bv incorrec subtraction of atmospheric water features. and so the line-like features in this region of the spectrum should be treated with caution.," The spectrum between and may be severely affected by incorrect subtraction of atmospheric water features, and so the line-like features in this region of the spectrum should be treated with caution."
+ The presence of the doublet in absorption (near .)) shows there to be a non-neelieible fraction of starlight from the host galaxy contributing to this spectrum., The presence of the doublet in absorption (near ) shows there to be a non-negligible fraction of starlight from the host galaxy contributing to this spectrum.
+ We now examine this stellar component., We now examine this stellar component.
+ For the blue data we extracted the spectrum of the host galaxy (which appears on the long slit frames as a clearly defined. “fuzz surrounding the bright nuclear emission.), For the blue data we extracted the spectrum of the host galaxy (which appears on the long slit frames as a clearly defined `fuzz' surrounding the bright nuclear emission.)
+ The spectra. extracted from the individual long-slit’ frames were again ¢o-acded to form a single spectrum., The spectra extracted from the individual long-slit frames were again co-added to form a single spectrum.
+ Figure 2a shows this co-added: spectrum., Figure 2a shows this co-added spectrum.
+ The absence of any. broad jalmer lines from this spectrum is verification that we have successfully isolated the stellar component., The absence of any broad Balmer lines from this spectrum is verification that we have successfully isolated the stellar component.
+ Note that there narrow OL) and ΟΠΗ] line emission. present. in this spectrum., Note that there narrow [OII] and [OIII] line emission present in this spectrum.
+ We attribute this to the narrow line region (NLA) possessing significant spatial extent (and. therefore contaminating the spectrum of the host galaxy.), We attribute this to the narrow line region (NLR) possessing significant spatial extent (and therefore contaminating the spectrum of the host galaxy.)
+ The stellar spectrum reveals several absorption features., The stellar spectrum reveals several absorption features.
+ The weak LL? absorption indicates the presence of a young stellar population., The weak $\beta$ absorption indicates the presence of a young stellar population.
+ An older stellar population give rise to the Ale D.Fer. G-band and absorption features.," An older stellar population give rise to the Mg B, G-band and absorption features."
+ Assuming the stellar spectrum to be spatially uniform. we can subtract the spectrum of the host galaxy from the total nuclear spectrum in order to isolate the spectrum of the active nucleus.," Assuming the stellar spectrum to be spatially uniform, we can subtract the spectrum of the host galaxy from the total nuclear spectrum in order to isolate the spectrum of the active nucleus."
+" In. practice. we progressively subtract more of the stellar spectrum from the total spectrum until the doublet feature at 40004, vanishes."," In practice, we progressively subtract more of the stellar spectrum from the total spectrum until the doublet feature at $\sim 4000$ vanishes."
+ The resulting dilference spectrum is taken to be the spectrum of the active nucleus., The resulting difference spectrum is taken to be the spectrum of the active nucleus.
+ This is shown in ligure 2b., This is shown in Figure 2b.
+ A similar procedure was not performed. for the red (FORS) data because the poorer spatial resolution of our FORS data makes isolation of the host galaxy spectrum more cillicult., A similar procedure was not performed for the red (FORS) data because the poorer spatial resolution of our FORS data makes isolation of the host galaxy spectrum more difficult.
+ This would make subtraction of the stellar component rather subjective ancl consequently diminish its value., This would make subtraction of the stellar component rather subjective and consequently diminish its value.
+ As is typical for Sevfert 1 nuclei. the spectrum consists of a strong non-stellar continuum. broad. Balmer lines and narrow permitted and forbidden lines.," As is typical for Seyfert 1 nuclei, the spectrum consists of a strong non-stellar continuum, broad Balmer lines and narrow permitted and forbidden lines."
+ Both the red and blue spectra were visually examined for known prominent lines., Both the red and blue spectra were visually examined for known prominent lines.
+ All such identified lines were characterized by fitting a single Gaussian profile whilst modeling the local continuum as a power-law., All such identified lines were characterized by fitting a single Gaussian profile whilst modeling the local continuum as a power-law.
+" ""This procedure was performed on the galaxy-subtracted spectrum. with the exception of those few lines that were identified. in the FORS data (for the reasons given above)."," This procedure was performed on the galaxy-subtracted spectrum, with the exception of those few lines that were identified in the FORS data (for the reasons given above)."
+ Since our stellar spectrum is contaminated with forbidden oxveen line emission. (presumably from an extended. NLR). we note that these oxygen lines will be suppressed by ~LO per cent in the galaxy-subtracted AGN spectrum.," Since our stellar spectrum is contaminated with forbidden oxygen line emission (presumably from an extended NLR), we note that these oxygen lines will be suppressed by $\sim 10$ per cent in the galaxy-subtracted AGN spectrum."
+ The single Gaussian parameterization. 1s a (visually) good fit to all of the emission. lines except 117., The single Gaussian parameterization is a (visually) good fit to all of the emission lines except $\beta$.
+ Three Gaussian components are required to. properly describe this line: a) A narrow line component at the svstemic velocity of the galaxy (delined as the velocity of the OLLI] emission line region) with FWIIMZτοῦkms and lux £043)= +., Three Gaussian components are required to properly describe this line: a) A narrow line component at the systemic velocity of the galaxy (defined as the velocity of the [OIII] emission line region) with $\approxlt 700\kmps$ and flux $F_{\rm n}(H\beta)=9.6\times 10^{-15}\ergpcmsqps$ .
+ b) X broad line component blueshifted hy 200+G60kms relative to. the systemic velocity with FWIIM-2400+200kms and Εν 1).=63.10Mereem7s +.," b) A broad line component blueshifted by $200\pm 60\kmps$ relative to the systemic velocity with $\sim 2400\pm 200\kmps$ and flux $F_{\rm
+b}(H\beta)=6.3\times 10^{-14}\ergpcmsqps$ ."
+ C) A verv broad. line component redshifted by 500+120kms relative to the svstemic velocity. with ENIM7100+700kms and lux {11)=4.610οσοιο7s 1.," c) A very broad line component redshifted by $500\pm 120\kmps$ relative to the systemic velocity with $\sim 7100\pm 700\kmps$ and flux $F_{\rm
+vb}(H\beta)=4.6\times 10^{-14}\ergpcmsqps$ ."
+ The resulting line identifications. wavelengths. line widths and total line fluxes for all of the identified lines are given in Table 1.," The resulting line identifications, wavelengths, line widths and total line fluxes for all of the identified lines are given in Table 1."
+ The errors quoted in this table (aud those for the 11 components above) include both statistical errors and an estimate of any systematic errors resultine from the waveleneth/Uux calibration., The errors quoted in this table (and those for the $\beta$ components above) include both statistical errors and an estimate of any systematic errors resulting from the wavelength/flux calibration.
+ The statistical errors are derived from. v7 fitting of the Gaussian. models to the data. assuming that the pixel-to-pixel variation is due to random Gaussian noise.," The statistical errors are derived from $\chi^2$ fitting of the Gaussian models to the data, assuming that the pixel-to-pixel variation is due to random Gaussian noise."
+ The optical spectrum clearly shows the cLlleet of dust extinetion: the continuum flux declines towards the blue end of the spectrum and the Balmer decrements are large., The optical spectrum clearly shows the effect of dust extinction: the continuum flux declines towards the blue end of the spectrum and the Balmer decrements are large.
+ The line-of-sight extinction. and a comparison of this extinction to the line-of-sight’ X-ray absorption will be addressed. in Section 3.," The line-of-sight extinction, and a comparison of this extinction to the line-of-sight X-ray absorption will be addressed in Section 3."
+ Vhe ultraviolet data were taken by LUE on 1994 July 2324 (i.e. three days prior to the observation described below.), The ultraviolet data were taken by IUE on 1994 July 23--24 (i.e. three days prior to the observation described below.)
+ Due to excessive scattered. solar light. we present only the data taken with the short. wavelength. primary (SWI) camera.," Due to excessive scattered solar light, we present only the data taken with the short wavelength primary (SWP) camera."
+ These short-wavelength data were taken during one full shift on 1994 Jub 24 (ellective exposure time S hours)., These short-wavelength data were taken during one full shift on 1994 July 24 (effective exposure time 8 hours).
+ Data reduction was performed using the software package., Data reduction was performed using the software package.
+ The resulting SWP spectrum is shown in Fig., The resulting SWP spectrum is shown in Fig.
+ 3., 3.
+ The source is found to be weak in the UV compared with other wavebands (apart fromradio)., The source is found to be weak in the UV compared with other wavebands (apart fromradio).
+ Thehypothesis that this UV spectrum is Hat and featureless can be rejected with a high level of confidence: such a hypothesis gives X7=153 for, Thehypothesis that this UV spectrum is flat and featureless can be rejected with a high level of confidence: such a hypothesis gives $\chi^2=153$ for
+the interplay between cooling and heating processes becomes dominant.,the interplay between cooling and heating processes becomes dominant.
+ The relevant time scale is therefore that reported in Eq. (159., The relevant time scale is therefore that reported in Eq. \ref{eq:tserio}) ).
+ After detailed analysis of extensive numerical modeling of decaying flaring loops. we have checked that the decay time from Eq. (1))," After detailed analysis of extensive numerical modeling of decaying flaring loops, we have checked that the decay time from Eq. \ref{eq:tserio}) )"
+ should be computed with a correction factor: where ὁ=1.3 to better fit the decay time measured from numerical models., should be computed with a correction factor: where $\phi \approx 1.3$ to better fit the decay time measured from numerical models.
+ Hydrodynamic simulations confirm that the time required to reach full equilibrium scales as the loop cooling time (τς). and. as shownfor instance in Fig.," Hydrodynamic simulations confirm that the time required to reach full equilibrium scales as the loop cooling time $\tau_s$ ), and, as shownfor instance in Fig."
+ 5. (see also Jakimiec et al., \ref{fig:mod_lin} (see also Jakimiec et al.
+ 1992). the time to reach flare steady-state equilibrium ts: We have verified that Eq.(6)) holds for loops of different lengths.," 1992), the time to reach flare steady-state equilibrium is: We have verified that \ref{eq:tequil}) ) holds for loops of different lengths."
+ Fig., Fig.
+ 6 shows the ratio of the time required to reach of the pressure equilibrium value to the time computed from Eq. (6)), \ref{fig:tequil} shows the ratio of the time required to reach of the pressure equilibrium value to the time computed from Eq. \ref{eq:tequil}) )
+ for simulations of flaring loops with three different loop lengths. and three different heating rates appropriate to reach the temperature of 10. 20 and 30 MK. respectively.," for simulations of flaring loops with three different loop lengths, and three different heating rates appropriate to reach the temperature of 10, 20 and 30 MK, respectively."
+ The agreement is within10%., The agreement is within.
+. For!2 fey. the density asymptotically approaches the equilibrium value: where à=1.4x107 (c.g.s.," For $t \geq t_{eq}$ , the density asymptotically approaches the equilibrium value: where $a = 1.4 \times 10^3$ (c.g.s."
+ units). or as obtained from the loop sealing laws (Rosner et al.," units), or as obtained from the loop scaling laws (Rosner et al."
+ 1978) and the plasma equation of state., 1978) and the plasma equation of state.
+ If the heat pulse stops before the loop reaches equilibrium conditions. the flare maximum density ts lower than the value at equilibrium.," If the heat pulse stops before the loop reaches equilibrium conditions, the flare maximum density is lower than the value at equilibrium."
+ Fig., Fig.
+ 5. shows that. after the initial impulsive evaporation on a time scale given by Eq.(2)). the later progressive pressure growth can be approximated with a linear trend.," \ref{fig:mod_lin} shows that, after the initial impulsive evaporation on a time scale given by \ref{eq:tsound}) ), the later progressive pressure growth can be approximated with a linear trend."
+ Since the temperature is almost constant in this phase (Fig. 3).," Since the temperature is almost constant in this phase (Fig. \ref{fig:mod_evol}) ),"
+ we can as well approximate that the density increases linearly for most of the time., we can as well approximate that the density increases linearly for most of the time.
+ We can then estimate the value of the maximum density at the loop apex as: where fa; is the time at which the density maximum occurs. which can be measured directly from observations.," We can then estimate the value of the maximum density at the loop apex as: where $t_{M}$ is the time at which the density maximum occurs, which can be measured directly from observations."
+ Since it Is reasonable to assume that the volume of the flaring region does not change much. at least on the relatively short time scale of the rise phase. the time of the naximum emission measure Is à good proxy for fay.," Since it is reasonable to assume that the volume of the flaring region does not change much, at least on the relatively short time scale of the rise phase, the time of the maximum emission measure is a good proxy for $t_M$."
+ Phase III ranges between the time at which the heat pulse ends and the time of the density maximum., Phase III ranges between the time at which the heat pulse ends and the time of the density maximum.
+ Fig., Fig.
+ 3. shows that the start time of the temperature decay marks well the end of the heat pulse., \ref{fig:mod_evol} shows that the start time of the temperature decay marks well the end of the heat pulse.
+ In many stellar flares. the temperature e“‘olution is not well resolved. and we may take the time of the temperature maximum as indicative of the end of the heat pulse.," In many stellar flares, the temperature evolution is not well resolved, and we may take the time of the temperature maximum as indicative of the end of the heat pulse."
+ The time of the density maximum is also the time at which the decay path crosses the locus of the equilibrium loops (QSS curve)., The time of the density maximum is also the time at which the decay path crosses the locus of the equilibrium loops (QSS curve).
+ As reported in Sec. 2.1..," As reported in Sec. \ref{sec:fl_ev},"
+ Cargill Klimchuk (2004) remarked that. at this time. the radiative cooling time Is exactly equal to the conductive cooling time.," Cargill Klimchuk (2004) remarked that, at this time, the radiative cooling time is exactly equal to the conductive cooling time."
+ By equating the time scales in Eqs. (4)), By equating the time scales in Eqs. \ref{eq:trad}) )
+ and (3)). we can then derive the temperature T3; at which the flare maximum density occurs: or If a value for L is already available tous and we are able to measure 73. we can derive 73; from Eq.(9)). in alternative to Eq.(8)).," and \ref{eq:tcon}) ), we can then derive the temperature $T_{M}$ at which the flare maximum density occurs: or If a value for $L$ is already available tous and we are able to measure $T_M$ , we can derive $n_M$ from \ref{eq:trc}) ), in alternative to \ref{eq:nmax}) )."
+"similar recursion relations (see, e.g., ?)), for which a recent implementation is available (?)..","similar recursion relations (see, e.g., \citealt{kostelec-rockmore-2008}) ), for which a recent implementation is available \citep{prezeau-reinecke-2010}."
+" Many packages allow the user to supply an array containing precomputed ccoefficients to the SHT routines, which can speed up the computation by a factor of up to «2."," Many packages allow the user to supply an array containing precomputed coefficients to the SHT routines, which can speed up the computation by a factor of up to $\approx 2$."
+" ddoes not provide such an option; this may seem surprising at first, but it was left out deliberately for the following reasons: uuses an adapted version of the well-known library (?),, ported to the C language (with a few minor adjustments) by the author."," does not provide such an option; this may seem surprising at first, but it was left out deliberately for the following reasons: uses an adapted version of the well-known library \citep{swarztrauber-1982}, ported to the C language (with a few minor adjustments) by the author."
+" This package performs well for FFTs of lengths N whose prime decomposition only contains the factors 2, 3, and 5."," This package performs well for FFTs of lengths $N$ whose prime decomposition only contains the factors 2, 3, and 5."
+" For prime N, its complexity degrades toO(N?),, which should be avoided, so for N containing large prime factors, Bluestein's algorithm (?) is employed, which allows replacing an FFT of length N by a convolution (conveniently implemented as a multiplication in Fourier space) of any length >2N—|."," For prime $N$, its complexity degrades to, which should be avoided, so for $N$ containing large prime factors, Bluestein's algorithm \citep{bluestein-1968} is employed, which allows replacing an FFT of length $N$ by a convolution (conveniently implemented as a multiplication in Fourier space) of any length $\ge2N-1$."
+" By appropriate choice of the new length, the desired complexity can be reestablished, at the cost of a higher prefactor."," By appropriate choice of the new length, the desired complexity can be reestablished, at the cost of a higher prefactor."
+" Obviously it would have been possible to use the very popular and efficient library (?) for the Fourier transforms, but since this package is fairly heavyweight, this would conflict with libpsht'ss completeness or compactness goals."," Obviously it would have been possible to use the very popular and efficient library \citep{frigo-johnson-2005}
+ for the Fourier transforms, but since this package is fairly heavyweight, this would conflict with s completeness or compactness goals."
+" Even though with the Bluestein modification will be somewhat slower thanFFTW,, the efficiency goal is not really compromised: libpsht'ss algorithms have an overall complexity ofO(N»/2,,),, so that the FFT part with its total complexity of will (for nontrivial problem sizes) never dominate the total CPU time."," Even though with the Bluestein modification will be somewhat slower than, the efficiency goal is not really compromised: s algorithms have an overall complexity of, so that the FFT part with its total complexity of will (for nontrivial problem sizes) never dominate the total CPU time."
+" As «5p and (6) and the complexity of already suggest, a code implementing SHTs on an iso-latitude grid will contain at least three nested loops which iterate on y, m, and l, respectively."," As \ref{syn_used}) ) and \ref{ana_used}) ) and the complexity of already suggest, a code implementing SHTs on an iso-latitude grid will contain at least three nested loops which iterate on $y$, $m$, and $l$, respectively."
+" The overall speed and/or memory consumption of the final implementation depend sensitively on the order in which these loops are nested, since different hierarchies are not equally well suited for precomputation of common quantities used in the innermost loop."," The overall speed and/or memory consumption of the final implementation depend sensitively on the order in which these loops are nested, since different hierarchies are not equally well suited for precomputation of common quantities used in the innermost loop."
+" The choice made for the computation of the,4j,(9)., tto obtain them by /-recursion (see refylmcalc)) strongly favors placing the loop over / innermost; so only the order of the y and m loops remains to be determined."," The choice made for the computation of the, to obtain them by $l$ -recursion (see \\ref{ylmcalc}) ) strongly favors placing the loop over $l$ innermost; so only the order of the $y$ and $m$ loops remains to be determined."
+" From the point of view of speed optimisation, arguably the best approach would be to iterate over m in the outermost loop, then over the ring number y, and finally over /."," From the point of view of speed optimisation, arguably the best approach would be to iterate over $m$ in the outermost loop, then over the ring number $y$, and finally over $l$."
+" This allows precomputing the Απ and B, quantities (see refylmcalc)) exactly once and in small chunks, for an additional memory cost of onlyO(/max)-."," This allows precomputing the $A_{lm}$ and $B_{lm}$ quantities (see \\ref{ylmcalc}) ) exactly once and in small chunks, for an additional memory cost of only."
+" However, this method requires additional sstorage for the intermediate products Fi, or G,,,, which is comparable to the size of the input and output data."," However, this method requires additional storage for the intermediate products $F_{m,y}$ or $G_{m,y}$, which is comparable to the size of the input and output data."
+" This was considered too wasteful, since it prohibits performing SHTs whose input and output data already occupy most of the available computer memory."," This was considered too wasteful, since it prohibits performing SHTs whose input and output data already occupy most of the available computer memory."
+" Unfortunately, switching the order of the two outermost loops does not improve the situation: now, one is presented with the choice of either precomputing the full two-dimensional Aj, and B,,, arrays, which is not acceptable for the same reason that prohibits storing the entire F,,, and μην arrays, or recomputing their values over and over, which consumes too much CPU time."," Unfortunately, switching the order of the two outermost loops does not improve the situation: now, one is presented with the choice of either precomputing the full two-dimensional $A_{lm}$ and $B_{lm}$ arrays, which is not acceptable for the same reason that prohibits storing the entire $F_{m,y}$ and $G_{m,y}$ arrays, or recomputing their values over and over, which consumes too much CPU time."
+" To solve this dilemma, it was decided not to perform the SHTs for a whole map in one go, but rather to subdivide them into a number of partial transforms, each of which only deals with a subset of rings (or ring pairs)"," To solve this dilemma, it was decided not to perform the SHTs for a whole map in one go, but rather to subdivide them into a number of partial transforms, each of which only deals with a subset of rings (or ring pairs)."
+" An appropriate choice for the number of parts will result in near-optimal performance combined with only small memory overhead; the most symmetrical choice of N,/* leads to a time", An appropriate choice for the number of parts will result in near-optimal performance combined with only small memory overhead; the most symmetrical choice of $N^{1/2}_{\vartheta}$ leads to a time
+[rom the top of the crust to neutron drip. 10.7<py0.66. and over the inner crust to the core. 8.013<pyo166. to obtain for piox0.66 ancl for pyo> 8.9.,"from the top of the crust to neutron drip, $10^{-5}<\rho_{12}<0.66$, and over the inner crust to the core, $8.913<\rho_{12}<166$, to obtain for $\rho_{12} \le 0.66$ and for $\rho_{12}\ge 8.9$ ."
+ In these equatlous ZNp is the temperattre at neutron drip (presumed constaut [or 0.60<pio« 8.0). yy=R/Ou. iis=g/101135∢∙∐↕⊳∖−⋅⋜↕∐≺⇂∠∩∡∶↜⋝⊽↴⋜↕∐⇂∠∣∶∶↜⋝⊽↴⋜⊔⋅≺↵↕∐≺↵∐⋯∐∐∩⊳∖∐↓≺↵↠∖⋅ ⋅⋅⋅ in unis of 10eres+. for pie<0.66 and pi»8.0. respectively.," In these equations $T_\ND$ is the temperature at neutron drip (presumed constant for $0.66<\rho_{12}<8.9$ ), $R_{10}=R/10\km$, $g_{14.3}=g/10^{14.3}\cm\second^{-2}$, and $L_{o,35}$ and $L_{i,35}$ are the luminosities, in units of $10^{35}\erg\second^{-1}$, for $\rho_{12}<0.66$ and $\rho_{12}>8.9$, respectively."
+" Both £, aud. L; are signed: they are positive if the flux is clirected outwards aud negative if direced inwards.", Both $L_o$ and $L_i$ are signed: they are positive if the flux is directed outwards and negative if directed inwards.
+ Notice that the coelliclent of Los in equation (22)) is au order of magniude larger tLian the coellicient of Ly.35 in equaticn (23)).," Notice that the coefficient of $L_{o,35}$ in equation \ref{eq:T-integrand-low}) ) is an order of magnitude larger than the coefficient of $L_{i,35}$ in equation \ref{eq:T-integrand-high}) )."
+TIis disparity relects that the inner crust requires a tuich sinaller thermal eracient than the outer crust to carry a [n]given [Iux., This disparity reflects that the inner crust requires a much smaller thermal gradient than the outer crust to carry a given flux.
+" Te» solve for the thermal st""ucture. E also require that the luminosity flowing away οι the crust |eal source Is aud that the core neutrino luminosity balances tje heat conducted into the core. Evaluating equation (22)) at py=10. and equaion (23)) at pio=166. aud using equatious (21)) aud (25)) to replace Lj and L; with Ly and L,. Τουalu an equation for the core temperature. Here T;=μιαs»."," To solve for the thermal structure, I also require that the luminosity flowing away from the crust heat source is and that the core neutrino luminosity balances the heat conducted into the core, Evaluating equation \ref{eq:T-integrand-low}) ) at $\rho_{12}=10^{-5}$ and equation \ref{eq:T-integrand-high}) ) at $\rho_{12}=166$, and using equations \ref{eq:luminosity-sum}) ) and \ref{eq:core-luminosity}) ) to replace $L_o$ and $L_i$ with $L_N$ and $L_\nu$, I obtain an equation for the core temperature, Here $T_\circ=T|_{\rho_{12}=10^{-5}}$."
+" Solving eqtation (26)) for a modified Urea luminosity L,455x aud Lxz1.07x105eres!Gig) gives Tusco=0.31 and Las=0.92Lsa5. Le.. ol the heat geueratec in the crust {ows iuto tlie core."," Solving equation \ref{eq:core-temperature}) ) for a modified Urca luminosity $L_{\nu,35}\approx 5\ee{4} \GTcore^8$ and $L_N\approx
+1.07\ee{36}\erg\second^{-1} (\dot{m}/\dot{m}_E)$ gives $\GTcore=0.34$ and $L_{\nu,35}=0.92L_{N,35}$, i.e., of the heat generated in the crust flows into the core."
+ This compares reasonably well with the uunmerical calculation without core supeΠιοity (Figures { and 6.. panels)).," This compares reasonably well with the numerical calculation without core superfluidity (Figures \ref{fig:pQ}
+ and \ref{fig:pQ-L}, )."
+ Iu that case. Ile preyper temperature at the crust botom is 3.1κ105 Ix. the luminosity flowing out the top of Ije CTst is LLQUOy. and the luuinosity lowing into the core is 0.8.," In that case, the proper temperature at the crust bottom is $3.1\ee{8}\K$ , the luminosity flowing out the top of the crust is $0.04L_N$, and the luminosity flowing into the core is $0.8L_N$."
+ Neutrino emission from the ‘rust |yalances the remainder o| Lw., Neutrino emission from the crust balances the remainder of $L_N$.
+ Ucbstituting Teoreg ΤΕ into edmation (23)). [find that νου=1.1. which is an overestimation ‘the maximum crust temperaure.," Substituting $\GTcore$ and $L_{\nu,35}$ into equation \ref{eq:T-integrand-high}) ), I find that $T_{\ND,9}=1.1$, which is an overestimation of the maximum crust temperature."
+ Thisis a «'onsequeuce of the “two-zone” treatment. which puts €a Lol he neutriuo cooling iu the core.," Thisis a consequence of the “two-zone” treatment, which puts all of the neutrino cooling in the core."
+ $ill. t1ο qualitative features of the numerical solutious are reprocucec.," Still, the qualitative features of the numerical solutions are reproduced."
+ In equation (23)). he iucrease in temperature [rom core to neutron drip (second teri.," In equation\ref{eq:T-integrand-high}) ), the increase in temperature from core to neutron drip (second term,"
+engaged in encounters with b<0.5. while =5 per cent are having encounters with 6«0.2. and 1.2per cent encounters with b«0.1.,"engaged in encounters with $b<0.5$, while $\approx 5$ per cent are having encounters with $b<0.2$, and $1-2$per cent encounters with $b<0.1$."
+ These figures are consistent with the 120 per cent of merging and. interacting galaxies observed: by Oenmler et al. (, These figures are consistent with the $10-20$ per cent of merging and interacting galaxies observed by Oemler et al. (
+1997) in four clusters at zc0.4.,1997) in four clusters at $z \simeq 0.4$.
+ The selfsimilarity of our simulations should ensure that such ellects do not exhibit any trend with redshift. as it is indeed. the case al least for 2.<1.," The self-similarity of our simulations should ensure that such effects do not exhibit any trend with redshift, as it is indeed the case at least for $z < 1$."
+ Our results support the claim that close encounters between galaxies in clusters are fairly frequent., Our results support the claim that close encounters between galaxies in clusters are fairly frequent.
+ Phe of encounters produced by our simulations is lower than the one invoked by Moore ct for galaxy harrassment. as most satellites in our simulations have at most one encounters at b«ol (as shown in the top panel of Figure 7)). and. not several as required by galaxy harrassment.," The of encounters produced by our simulations is lower than the one invoked by Moore et for galaxy harrassment, as most satellites in our simulations have at most one encounters at $b<1$ (as shown in the top panel of Figure \ref{fig:coll1}) ), and not several as required by galaxy harrassment."
+ However. the rate we found. should. be considered as a lower limit. as in real clusters collisions mostly take place in the central. densest region of the cluster. where we do not have enough resolution to preserve the satellite identity.," However, the rate we found should be considered as a lower limit, as in real clusters collisions mostly take place in the central, densest region of the cluster, where we do not have enough resolution to preserve the satellite identity."
+ Further investigation with higher resolution simulations is required. to estimate this ellect., Further investigation with higher resolution simulations is required to estimate this effect.
+ There exist two cillerent approaches to studying the properties of substructure within dark matter halos:(2) we can identify satellites before they merge with the main halo and keep track of their. particle trajectories after merging:(di) we can identify overdensity peaks within the main halo at each output time., There exist two different approaches to studying the properties of substructure within dark matter halos: we can identify satellites before they merge with the main halo and keep track of their particle trajectories after merging; we can identify overdensity peaks within the main halo at each output time.
+ The latter approach has the advantage of identifving physically motivated substructure., The latter approach has the advantage of identifying physically motivated substructure.
+ Moreover. us approach is equivalent to the observational procedure when it is applied to two-dimensional projections.," Moreover, this approach is equivalent to the observational procedure when it is applied to two-dimensional projections."
+ However. iis method is not straightforward if we are interested in the volution of the substructure properties.," However, this method is not straightforward if we are interested in the evolution of the substructure properties."
+ In fact. in order to heck the identity of density peaks identified: at cillerent imes. we need to check whether a substantial fraction. of 1e particles of a density peak at a Later time belonged to a density peak at an earlier time (IxIvpin et al.," In fact, in order to check the identity of density peaks identified at different times, we need to check whether a substantial fraction of the particles of a density peak at a later time belonged to a density peak at an earlier time (Klypin et al."
+ 1997)., 1997).
+ In the present work. we have eight output times Lor each of the nine simulations.," In the present work, we have eight output times for each of the nine simulations."
+ DPherefore. we use the former approach. which is more straightforward in keeping track of each satellite.," Therefore, we use the former approach, which is more straightforward in keeping track of each satellite."
+ Phe relatively large number of output times is the most important feature of our analvsis., The relatively large number of output times is the most important feature of our analysis.
+ In fact. we can consistently compare the satellite orbit evolution with the dynamical friction predictions.," In fact, we can consistently compare the satellite orbit evolution with the dynamical friction predictions."
+ Moreover. we can accurately estimate the evolution of the satellite internal properties as tidal effects ancl collisions disrupt them.," Moreover, we can accurately estimate the evolution of the satellite internal properties as tidal effects and collisions disrupt them."
+ This information provides a quantitative definition of the survival time., This information provides a quantitative definition of the survival time.
+ Our analysis complements Chigna ct al., Our analysis complements Ghigna et al.
+ analysis., analysis.
+ CGhigna ct al., Ghigna et al.
+ identify overdensities in a sinele halo simulation with higher space ancl force resolution than ours., identify overdensities in a single halo simulation with higher space and force resolution than ours.
+ Thus. they can accurately estimate the satellite density. profiles. and the spatial distribution of the satellites within the main halo.," Thus, they can accurately estimate the satellite density profiles, and the spatial distribution of the satellites within the main halo."
+ On the other hand. our work has the advantage of a better statistics provided. by the nine independent. clusters ancl the use of a larger number of output times.," On the other hand, our work has the advantage of a better statistics provided by the nine independent clusters and the use of a larger number of output times."
+ The results of this paper may be summarized as follows: The application of our results to galaxies in clusters requires us to specify a procedure for populating the smaller dark matter satellites with stellar material., The results of this paper may be summarized as follows: The application of our results to galaxies in clusters requires us to specify a procedure for populating the smaller dark matter satellites with stellar material.
+ We discuss an empirical method of assigning galaxies to haloes in Section ?7.., We discuss an empirical method of assigning galaxies to haloes in Section \ref{sec:disc}.
+ Phe results can be summarised as follows: Roughly 50 per cent of galaxies in present-clay clusters should have a dark. matter halo truncated at or. below zz l0r., The results can be summarised as follows: Roughly 50 per cent of galaxies in present-day clusters should have a dark matter halo truncated at or below $\approx 10 r_e$ .
+ Some of these galaxies may themselves have been clisruptecl., Some of these galaxies may themselves have been disrupted.
+with mock CMB skies in which the point source population follows the same distribution as our model.,with mock CMB skies in which the point source population follows the same distribution as our model.
+" We found that although the overall level of correlation is low, W band correlations with actual data are in good agreement with correlations with the mock data."," We found that although the overall level of correlation is low, W band correlations with actual data are in good agreement with correlations with the mock data."
+" In Q band, instead, the degree of correlation between simulations and actual maps is significantly lower than what expected from the mock data."," In Q band, instead, the degree of correlation between simulations and actual maps is significantly lower than what expected from the mock data."
+ In this band foregrounds contamination is significantly higher than at higher frequency and may contribute to decreasing the level of correlation between simulations and real maps., In this band foregrounds contamination is significantly higher than at higher frequency and may contribute to decreasing the level of correlation between simulations and real maps.
+ We used the simulated maps to estimate the unresolved source correction to WMAPS measured spectra., We used the simulated maps to estimate the unresolved source correction to WMAP5 measured spectra.
+" As a general result, we find that unresolved sources in our simulations have a steeper frequency behaviour than the contribution estimated by the WMAP team."," As a general result, we find that unresolved sources in our simulations have a steeper frequency behaviour than the contribution estimated by the WMAP team."
+ This is due to the fact that observation of sources in the 20-90 GHz range with S<1 Jy typically show steeper spectral indexes than the best fit spectral index from the WMAPS catalog., This is due to the fact that observation of sources in the 20–90 GHz range with $S \lesssim 1$ Jy typically show steeper spectral indexes than the best fit spectral index from the WMAP5 catalog.
+" The observed scaling in the simulations corresponds to an effective spectral index a=—0.20~-—0.25, compared to WMAP5 a=0--0.09."," The observed scaling in the simulations corresponds to an effective spectral index $\alpha = -0.20 \sim -0.25$, compared to WMAP5 $\alpha = 0 \sim
+-0.09$."
+" Thus, the unresolved point sources contribution derived by the simulation either agrees with the WMAPS estimates in Q band but is lower by up to ~30—40% at 94 GH, depending on the way we mask bright sources."," Thus, the unresolved point sources contribution derived by the simulation either agrees with the WMAP5 estimates in Q band but is lower by up to $\sim 30-40\%$ at 94 GH, depending on the way we mask bright sources."
+" The corresponding shift in estimates for parameter like ng and w, can reach up to ~0.3— 0.4σ.", The corresponding shift in estimates for parameter like $n_S$ and $\omega_c$ can reach up to $\sim 0.3-0.4\sigma$ .
+" Given a sufficient number of simulations, the approach followed in this work allows in principle to estimate the probability distribution for the flux at an arbitrary frequency of each source from a template low frequency survey."," Given a sufficient number of simulations, the approach followed in this work allows in principle to estimate the probability distribution for the flux at an arbitrary frequency of each source from a template low frequency survey."
+" In turn this would allow to naturally account for the added uncertainty due to undetected sources during the map-making stage, as is done for galactic foregrounds with Gibbs sampling methods."," In turn this would allow to naturally account for the added uncertainty due to undetected sources during the map-making stage, as is done for galactic foregrounds with Gibbs sampling methods."
+" This would require that Gibbs samplers can be made to efficiently work at multipoles £~2000—3000, where point sources will be relevant for Planck and other possibly upcoming high-resolution experiments."," This would require that Gibbs samplers can be made to efficiently work at multipoles $\ell \simeq 2000-3000$, where point sources will be relevant for Planck and other possibly upcoming high-resolution experiments."
+" Even if this were not possible, having an estimate of unresolved source contributions in each pixel would be helpful in defining a sky mask, e.g. by flagging all pixel in which the source contribution is expected to exceed a given threshold even if there is no actual detection of a point source in that region of the map."," Even if this were not possible, having an estimate of unresolved source contributions in each pixel would be helpful in defining a sky mask, e.g. by flagging all pixel in which the source contribution is expected to exceed a given threshold even if there is no actual detection of a point source in that region of the map."
+ EP is and NSF-ADVANCE fellow (AST-0649899)., EP is and NSF–ADVANCE fellow (AST–0649899).
+" LPLC and EP were supported by NASA grant NNX07AH59G and Planck subcontract 1290790 for this work, and would like to thank Caltech for hospitality during this period."," LPLC and EP were supported by NASA grant NNX07AH59G and Planck subcontract 1290790 for this work, and would like to thank Caltech for hospitality during this period."
+" We acknowledge the use of the LAMBDA data archive and of the CosmoMC, PICO and HEALPix packages."," We acknowledge the use of the LAMBDA data archive and of the CosmoMC, PICO and HEALPix packages."
+ We thank the members of the US Planck ADG team for helpful and stimulating discussion., We thank the members of the US Planck ADG team for helpful and stimulating discussion.
+ (~99 per cent) with assigned distances have larger luminosities than would be expected from a B3 star — assuming the observed luminosity arises from a single embedded star.,$\sim$ 99 per cent) with assigned distances have larger luminosities than would be expected from a B3 star — assuming the observed luminosity arises from a single embedded star.
+ Massive stars are known to form exclusively in clusters., Massive stars are known to form exclusively in clusters.
+ However. ? showed that for a realistic initial mass function the most massive member of a cluster is only a couple of spectral classes lower than that determined by assuming the luminosity was due to a single star.," However, \citet{wood1989} showed that for a realistic initial mass function the most massive member of a cluster is only a couple of spectral classes lower than that determined by assuming the luminosity was due to a single star."
+ Taking this into account we find that any RMS source with a luminosity greater than 10* ((Le.. equivalent to a star of spectral type BI or earlier) is likely to contain a massive star.," Taking this into account we find that any RMS source with a luminosity greater than $^4$ (i.e., equivalent to a star of spectral type B1 or earlier) is likely to contain a massive star."
+ Applying this limit to identify a bona fide sample of massive young stars we find approximately two-thirds (193) of our sources have luminosities over 10Isun.. and are therefore luminous enough to host massive young stars.," Applying this limit to identify a bona fide sample of massive young stars we find approximately two-thirds (193) of our sources have luminosities over $^4$, and are therefore luminous enough to host massive young stars."
+ We note that these fractions only apply to the current sample since there is likely to be number of high mass stars that fall below our detection threshold., We note that these fractions only apply to the current sample since there is likely to be number of high mass stars that fall below our detection threshold.
+ In the right panel of reftig:luminosities we present a plot showing the distribution of source luminosities as a function of heliocentric distance along with the MSX llimiting sensitivity., In the right panel of \\ref{fig:luminosities} we present a plot showing the distribution of source luminosities as a function of heliocentric distance along with the MSX limiting sensitivity.
+" The sensitivity has been estimated by first calculating the flux in the MSX (band E) using F,=4.041x10IS, WWnyo (2). where $5, is the MSX ddetection limit (2.7 JJy)."," The sensitivity has been estimated by first calculating the flux in the MSX (band E) using $F_E=4.041\times10^{-14}\,S_{21}$ $^2$ \citealt{cohen2000}) ), where $S_{21}$ is the MSX detection limit $\sim$ Jy)."
+ This value is then multiplied by a factor of 24. calculated from the ratio of bolometric fluxes. determined from SED fits. to the MSX band E fluxes (Mottram et al.," This value is then multiplied by a factor of 24, calculated from the ratio of bolometric fluxes, determined from SED fits, to the MSX band E fluxes (Mottram et al."
+ 2010. submitted).," 2010, submitted)."
+ The fairly homogeneous distribution in heliocentric distance of RMS sources above 10*Isun.. would suggest that our sample of young massive stars is complete for luminosities above this across the Galaxy.," The fairly homogeneous distribution in heliocentric distance of RMS sources above $^4$, would suggest that our sample of young massive stars is complete for luminosities above this across the Galaxy."
+ To properly determine the completeness limit requires modelling of the luminosity function which is beyond the scope of this paper and will be the subject of a subsequent paper (Davies et al., To properly determine the completeness limit requires modelling of the luminosity function which is beyond the scope of this paper and will be the subject of a subsequent paper (Davies et al.
+ 2010. in prep.)," 2010, in prep.),"
+ however. preliminary results are consistent with the completeness level derived from reffis:luminosities .," however, preliminary results are consistent with the completeness level derived from \\ref{fig:luminosities} ."
+ In reftbl:rmsunmuarvwepresentasummaryvo f eachRMS sourcelocatedwithintheGRS lon Galactocentricdistanceandbolometricluminositvtonlvasmallportiono f thetabli strasbg.fr30.79.125.Sjorviahttp {οδόstrasbg.frícgi—bin{gealds{A+Af).," In \\ref{tbl:rms_summary} we present a summary of each RMS source located within the GRS longitude range including positions, associated GRS cloud name, assigned kinematic distance, height above the mid-plane, Galactocentric distance and bolometric luminosity (only a small portion of the table is provided here, the full table is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.125.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/)."
+ In this section we will discuss the Galactic distribution of molecular sus and compare the physical properties of clouds found to be associated with the RMS sampleof young missive stars., In this section we will discuss the Galactic distribution of molecular gas and compare the physical properties of clouds found to be associated with the RMS sampleof young massive stars.
+ To facilitate these discussions we need to calculate a few key parameters currently unavailable in either of the ? or? papers., To facilitate these discussions we need to calculate a few key parameters currently unavailable in either of the \citet{roman2009} or \citet{rathborne2009} papers.
+ In this subsection we will estimate masses and physical radii for every cloud that has an unambiguous distance. and for which an excitation temperature has been calculated by ?..," In this subsection we will estimate masses and physical radii for every cloud that has an unambiguous distance, and for which an excitation temperature has been calculated by \citet{rathborne2009}."
+ The observed radius of each cloud has been estimated using the projected area of the cloud on the sky. A. given by 2. via: To correct for the spatial resolution we usethe following: where 8a is the etfective beam size of the smoothed GRS cube.," The observed radius of each cloud has been estimated using the projected area of the cloud on the sky, $A$, given by \citet{rathborne2009} via: To correct for the spatial resolution we usethe following: where $\theta_{\rm{FWHM}}$ is the effective beam size of the smoothed GRS cube."
+" ? calculated the optical depths and excitation temperatures for all of the clouds by combining ""CO (J21—0) data from the University of Massachusetts-Stony Brook survey (2)) and the ""CO data from the GRS.", \citet{rathborne2009} calculated the optical depths and excitation temperatures for all of the clouds by combining $^{12}$ CO $J$ =1–0) data from the University of Massachusetts-Stony Brook survey \citealt{sanders1986}) ) and the $^{13}$ CO data from the GRS.
+ They calculated typical optical depths and excitation temperatures of «0.13 and KK respectively., They calculated typical optical depths and excitation temperatures of $\sim$ 0.13 and K respectively.
+ We can now use these results to estimate the mass of the individual clouds., We can now use these results to estimate the mass of the individual clouds.
+ Following the method outlined by ? it is possible to estimate the total mass of a cloud using a single optically thin molecular transition and assuming an appropriate abundance ratio., Following the method outlined by \citet{bourke1997} it is possible to estimate the total mass of a cloud using a single optically thin molecular transition and assuming an appropriate abundance ratio.
+" Since only a single transition is used. the mass is determined assuming local thermodynamic equilibrium (LTE) and is referred to hereafter as the LTE mass (M, 44:."," Since only a single transition is used, the mass is determined assuming local thermodynamic equilibrium (LTE) and is referred to hereafter as the LTE mass $M_{\rm{LTE}}$ )."
+" We use the following equation taken from 9.. where the fa, is the mean molecular mass per H» molecule. nj is atomic mass of hydrogen. D is the heliocentric distance. v is in aand © is the source solid angle in aremin?."," We use the following equation taken from \citet{bourke1997}: : where the $\mu_{\rm{m}}$ is the mean molecular mass per $_2$ molecule, $m_{\rm{H}}$ is atomic mass of hydrogen, $D$ is the heliocentric distance, $v$ is in and $\Omega$ is the source solid angle in $^2$."
+ We have assumed an abundance ratio of Hs to CO of 7x10 (221) a mean molecular mass of 1.36 a7). taking account of the H. content of of the gas.," We have assumed an abundance ratio of $_2$ to $^{13}$ CO of $7 \times 10^5$ \citealt{frerking1982,bourke1997}) ) a mean molecular mass of 1.36 $m_{\rm{H}}$ , taking account of the $_{\rm{e}}$ content of of the gas."
+ Distances were takenfrom ? and the Έως from ?.., Distances were takenfrom \citet{roman2009} and the $T_{\rm{ex}}$ from \citet{rathborne2009}. .
+ The [ThydvdO term in 33 is equal to the total integrated intensity which has been ealeulated for each cloud and presented in, The $\int \int T_{\rm{mb}}\ {\rm{d}}v\ {\rm{d}}\Omega$ term in 3 is equal to the total integrated intensity which has been calculated for each cloud and presented in
+respect to suuspot numbers. with a time lag of two to three vears between the peak times in solar cvele 2] (Wagner 1988: Asclavanden 1991. Dronmmud et al.,"respect to sunspot numbers, with a time lag of two to three years between the peak times in solar cycle 21 (Wagner 1988; Aschwanden 1994, Bromund et al."
+ 1995)., 1995).
+ Tlowever. Wilsou (1993) did not find evidence for a tine lag between the maxima of the rates of optical flares and N-vav background flux in solar cycle 22.," However, Wilson (1993) did not find evidence for a time lag between the maxima of the rates of optical flares and X-ray background flux in solar cycle 22."
+ Wheatland Litvinenko (2001) fore that there is an average σαν of aout 6 nouths considering evcles 21 and 22 together by using correlation analysis., Wheatland Litvinenko (2001) found that there is an average delay of about 6 months considering cycles 21 and 22 together by using correlation analysis.
+ For the evele 23. Tan (2010). comxwed the sunspot nuubers with the distributions of the :)pearance rate of solar GOES flares.," For the cycle 23, Tan (2010) compared the sunspot numbers with the distributions of the appearance rate of solar GOES flares."
+ They found iif there is a time lag jetween the maxim value of sunspot nunibors auc the ιαπαι values of the amuual προς of C-class. Me-class. and X-class flares.," They found that there is a time lag between the maximum value of sunspot numbers and the maximum values of the annual numbers of C-class, M-class, and X-class flares."
+" The aΜ��] averaged relative suuspot number reaches its asim m about 2000 and the annual uiuubers of C-class; M-class; ancl N-class flares reach tweir αμαπια in about 2001. 2001. acl 2002 respectively,"," The annual averaged relative sunspot number reaches its maximum in about 2000 and the annual numbers of C-class, M-class, and X-class flares reach their maxima in about 2001, 2001, and 2002 respectively."
+ Tenuuer et al. (, Temmer et al. (
+2003) found that there is a characerislic time lag between flare activity and suuspot activity in the rauge of 10-15 mouths for solar cycles 19. 21. aid 23 by using the uuuber of the flares in each mouth.,"2003) found that there is a characteristic time lag between flare activity and sunspot activity in the range of 10-15 months for solar cycles 19, 21, and 23 by using the number of the flares in each month."
+ Iu order to investigate all kinds of long-term flare activity. the flare ineex calculated by T. Atac aud. A. Ozeuc from Dogaziei University Nauclilli Observatory and the daily soft N-rav flare iudex defined by Autalové (1996) are two mad flare indices in the present research.," In order to investigate all kinds of long-term flare activity, the flare index calculated by T. Atac and A. Ozguc from Bogazici University Kandilli Observatory and the daily soft X-ray flare index defined by $\acute{a}$ (1996) are two main flare indices in the present research."
+ For lustance. Li et al. (," For instance, Li et al. ("
+2010) found that the northeru-hemuspheric flare activity should lead the southerp-hemispheric flare activity for low-frequency colponcuts by usine the former flare index.,2010) found that the northern-hemispheric flare activity should lead the southern-hemispheric flare activity for low-frequency components by using the former flare index.
+ Joshi et al. (, Joshi et al. (
+2001) reported a real north-south asvuunetrv durus solar ήπια and obtained the significant periods of approximately 28.26 davs. 550.3 davs. aud 3.72 vears by using the latter flare iudex.,"2004) reported a real north-south asymmetry during solar minimum and obtained the significant periods of approximately 28.26 days, 550.3 days, and 3.72 years by using the latter flare index."
+ Iu this paper. we focus on the three solar evcles (21. 22. and 23) to investigate the phase relation between sunspot nunibers and flare activity (C-class. M-class. and N-class flares) by usingC» a flare index. which is stuuilar to alové (1996).," In this paper, we focus on the three solar cycles (21, 22, and 23) to investigate the phase relation between sunspot numbers and flare activity (C-class, M-class, and X-class flares) by using a flare index, which is similar to $\acute{a}$ (1996)."
+ The data used in this paper are as follows: 1., The data used in this paper are as follows: 1.
+ The monthly mean northern aud southern sunspot numbers from January 1915 to December 2001 were compiled by Tenuuer et al., The monthly mean northern and southern hemispheric sunspot numbers from January 1945 to December 2004 were compiled by Temmer et al.
+ httpedsweliu:ΜΕυπμμ| AfMT/T35)., $http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/447/735$ ).
+ The| monthly mean northern aud southern henmisphleric: sunspot nunibers frou January 1992 to Deceiiber 2008 can be downloaded from: NOAA's National Geophysical Data Ceuter (NCDC?) (ftp:ftp.ngdeioua.goc/STPILEMISPIIERENUAMBERS/)., The monthly mean northern and southern hemispheric sunspot numbers from January 1992 to December 2008 can be downloaded from NOAA's National Geophysical Data Center (NGDC) $ftp://ftp.ngdc.noaa.gov/STP\protect\newline/SOLAR\_DATA/SUNSPOT\_NUMBERS\protect\newline/HEMISPHERE\_NUMBERS/$ ).
+ There is an overlapping time span frou January 1992 through. December 2001 with the first sunspot ¢lata., There is an overlapping time span from January 1992 through December 2004 with the first sunspot data.
+ The first eroup of data eenerallv matches the second group (for details. sec Teunner et al.," The first group of data generally matches the second group (for details, see Temmer et al."
+ 2006)., 2006).
+ In this paper. we extracted the sunspot nunibers from λίαν 1976 to May. 2008 as our saluples.," In this paper, we extracted the sunspot numbers from May 1976 to May 2008 as our samples."
+ 2., 2.
+ The soft XN-rav flares were downloaded from NGDCC (ftp:/fftp.igdοουSTP/SOLAR_DATSOLARFLARES/FLARES NRA)., The soft X-ray flares were downloaded from NGDC $ftp://ftp.ngdc.noaa.gov/STP/SOLAR\_DATA\protect\newline/SOLAR\_FLARES/FLARES\_XRAY/$ ).
+ Iu this paper. we adoted. the peak fluxes of soft ray observed by GOES ]SA aas the defined flare fiux when one flare occurred.," In this paper, we adopted the peak fluxes of soft X-ray observed by GOES 1-8 as the defined flare flux when one flare occurred."
+ Afterwards. we ste|l monthly peak fluxes of the defined flares (C-class. M-class. or N-class flares) as a flare index.," Afterwards, we summed monthly peak fluxes of the defined flares (C-class, M-class, or X-class flares) as a flare index."
+ The equations are as follows: Iu the above equations. C;. M;. X; indicate all sub- of C-class. Meclass. N-class flares (c.¢.. For C-class flares. the sub-level flares (Cj) are CLO. CI.1. C1.2. 05006 . aud so on.).," The equations are as follows: In the above equations, $C_i$, $M_i$, $X_i$ indicate all sub-levels of C-class, M-class, X-class flares (e.g., For C-class flares, the sub-level flares $C_i$ ) are C1.0, C1.1, C1.2, $\cdot\cdot\cdot\cdot\cdot\cdot$ , and so on.)."
+ 77 is eiven as a float wmuber with one decimal (prior to 1980 April 27. 7/7 is listed as an iuteeer.).," $i$ "" is given as a float number with one decimal (prior to 1980 April 27, $i$ "" is listed as an integer.)."
+" The value of ""7 is 1. L1. 1.2 "," The value of $i$ "" is 1, 1.1, 1.2 $\cdot\cdot\cdot\cdot\cdot\cdot$ , $j$."
+For C-class and M-class flares. the value of 7/7 is no more than 9.9.," For C-class and M-class flares, the value of $i$ "" is no more than 9.9."
+ 5; denotes the mouthly απους of sub-level fares (e.g. 05 denotes the mouthly numbers of Cs. M3 ov X3.). Fe Pay.," $n_i$ denotes the monthly numbers of sub-level flares (e.g., $n_5$ denotes the monthly numbers of $C_5$, $M_5$ or $X_5$ .). $F_C$, $F_M$,"
+ aud Fx indicate the mouthily peak fluxes of C-class. MEclass. aad X-class flares.," and $F_X$ indicate the monthly peak fluxes of C-class, M-class, and X-class flares."
+ This method is the same as the flare iudex (FI) defined by Antalové (1996) aud used by Joshi et al. (, This method is the same as the flare index (FI) defined by $\acute{a}$ (1996) and used by Joshi et al. (
+2001).,2004).
+ The unit of the poak flux is οTs1, The unit of the peak flux is $\cdot$$cm^{-2}$$\cdot$$s^{-1}$.
+ As weak soft X-ray flares can not be detected during times of high activity (Feldiman et al., As weak soft X-ray flares can not be detected during times of high activity (Feldman et al.
+ 1997: Veronig et al., 1997; Veronig et al.
+ 2002). we cliscarded the classification of soft N- flare lower than C1.0 to removethe effect. of the hemisphericvariation ofsoft N-rav background radiation.," 2002), we discarded the classification of soft X-ray flares lower than C1.0 to removethe effect of the variation ofsoft X-ray background radiation."
+ Figure 1 shows the distribution of the mouthiy uuuboers of three types of flares from λίαν 1976 to Alay 2008., Figure 1 shows the distribution of the monthly numbers of three types of flares from May 1976 to May 2008.
+of an astrophvsical svstem.,of an astrophysical system.
+ For example. fitting the interstellar extinction curve requires a distribution of erain sizes.," For example, fitting the interstellar extinction curve requires a distribution of grain sizes."
+ Mathis.ΠΡΙ&Nordsieck(1977.hereafterMIN). used a simple power law size distribution with index q=—3.5. ancl a lower and upper cutolf size. (minv0.05son and eu~0.25jan.," \citet[hereafter MRN]{mat77} used a simple power law size distribution with index $q=-3.5$, and a lower and upper cutoff size, $a_{\rm min} \sim 0.05\,\mu\rm m$ and $a_{\rm max} \sim 0.25\,\mu\rm m$."
+ This distribution has been revised many (mes in (he literature., This distribution has been revised many times in the literature.
+" For example. im.Martin&Hendry(1994). proposed a distribution in the form where e,=0.14jm for silicate grains and a)=0.28jn for graphite grains."," For example, \citet{kim94} proposed a distribution in the form where $a_0=0.14\,\mu\rm m$ for silicate grains and $a_0=0.28\,\mu\rm m$ for graphite grains."
+ The upper cutoff was replaced by an exponential decay. to allow for the presence of large dust grains (e=0.2 jum).," The upper cutoff was replaced by an exponential decay, to allow for the presence of large dust grains $a \gtrsim 0.2 \,\mu \rm m$ )."
+ Although the [aet that the dust has a distribution of grain sizes is well established. the exact form of this distribution is still an open issue.," Although the fact that the dust has a distribution of grain sizes is well established, the exact form of this distribution is still an open issue."
+ As we will demonstrate below. part ol this uncertainty is due to the AI of the problem. which makes (wo models with similar size distributions virtually indistinguishable.," As we will demonstrate below, part of this uncertainty is due to the AI of the problem, which makes two models with similar size distributions virtually indistinguishable."
+ Another source of uncertainty lies on our. vel incomplete knowledge of the details of erain formation and growth., Another source of uncertainty lies on our yet incomplete knowledge of the details of grain formation and growth.
+ To complicate matters even further. the grain size distribution is likely to vary across the wind because different erain sizes have different αντ velocities (Elitzur&Ivezié2001).," To complicate matters even further, the grain size distribution is likely to vary across the wind because different grain sizes have different drift velocities \citep{eli01}."
+. In this section. we study the SED of models wilh a grain size distribution.," In this section, we study the SED of models with a grain size distribution."
+ For simplicity. we adopt a power-law (MIN) grain size distribution funcüon: although the «quantitative details of the results shown below do depend on the adopted size distribution. the fundamental concepts discussed do not.," For simplicity, we adopt a power-law (MRN) grain size distribution function; although the quantitative details of the results shown below do depend on the adopted size distribution, the fundamental concepts discussed do not."
+ Previously. in section 4.. we used the concept of AI to reveal important similarities in the Η SED (and in the extinction SED for optically thin models) for single grain size models.," Previously, in section \ref{single}, we used the concept of AI to reveal important similarities in the IR SED (and in the extinction SED for optically thin models) for single grain size models."
+ In the following. we study what controls the spectral shape of the SED [or models with a size distribution in an attempt to idenüfv (he same sort of similarities.," In the following, we study what controls the spectral shape of the SED for models with a size distribution in an attempt to identify the same sort of similarities."
+ Let us consider an envelope model. consisting of spherical dust grains of the same chemical composition and sizes distributed according to a given size distribution function. dnda.," Let us consider an envelope model, consisting of spherical dust grains of the same chemical composition and sizes distributed according to a given size distribution function, $dn/da$."
+ The envelope IR SED. Fy. will be given bx the sum of the contributions from each graln size where (a) is the wavelength integrated emission for each grain size. a measure of how," The envelope IR SED, ${F}_{\rm IR}$, will be given by the sum of the contributions from each grain size where $f(a)$ is the wavelength integrated emission for each grain size, a measure of how"
+Figure 9 shows the total fluxes of the forward ancl reverse shock emission at the observed frequency v—5 Gllz lor observed times /=1 hr. 1 day. 10 days ancl LOO davs. as a function ol redshift z. together with the 5o sensitivities of the VLA and SNA.,"Figure \ref{fig:magnetized:z} shows the total fluxes of the forward and reverse shock emission at the observed frequency $\nu=5$ GHz for observed times $t=1$ hr, $1$ day, $10$ days and $100$ days, as a function of redshift $z$, together with the $5 \sigma$ sensitivities of the VLA and SKA."
+ Comparing Figures 9 and 1 wesee that the detection of the magnetized GRB afterglow is somewhat easier (han lor the standard case. and even the eurrent VLA could detect such magnetized GRB afterglows up to z~30.," Comparing Figures \ref{fig:magnetized:z} and \ref{fig:standard:z}
+ we see that the detection of the magnetized GRB afterglow is somewhat easier than for the standard case, and even the current VLA could detect such magnetized GRB afterglows up to $z \sim 30$."
+ A small increase in the fIux is eaused by the reverse shock emission at early times /<10 day. due to the higher reverse magnetic field.," A small increase in the flux is caused by the reverse shock emission at early times $t \siml 10$ day, due to the higher reverse magnetic field."
+ Figure 10. shows the light curves of the magnetized GRB afterglows al z=6 and pcc200 ΑΔΗ and at z=13 andy ~100 MlII., Figure \ref{fig:magnetized} shows the light curves of the magnetized GRB afterglows at $z=6$ and $\nu \sim 200$ MHz and at $z=13$ and $\nu \sim 100$ MHz.
+ The56 sensitivities of the VLA. LOFAR and SIVA are also shown.," The$5 \sigma$ sensitivities of the VLA, LOFAR and SKA are also shown."
+ The reverse shock emission al low frequencies ~LOO MlIZ is still below the Forward shock emission. being dimmer al these frequencies than in the stancard case in Figure 2..," The reverse shock emission at low frequencies $\sim 100$ MHz is still below the forward shock emission, being dimmer at these frequencies than in the standard case in Figure \ref{fig:standard}."
+ Since the forward shock emission is similar to the standard case. the detectabilitv is similar to that in the standard case.," Since the forward shock emission is similar to the standard case, the detectability is similar to that in the standard case."
+ We see that LOFAR and SIXÁÀ can detect the magnetized allerglows at low frequencies 100 MlII&E and high redshift z—10. while the VLA maa have difficulty detecting them.," We see that LOFAR and SKA can detect the magnetized afterglows at low frequencies $\sim 100$ MHz and high redshift $z \sim 10$, while the VLA may have difficulty detecting them."
+ From Figure 10.. the peak flux is seen to be about ~1-10 μ.]ν ad the redshifted 21 em frequency. v=1420/(1+2) Alle. for 2=6.," From Figure \ref{fig:magnetized}, the peak flux is seen to be about $\sim 1$ $10$ $\mu$ Jy at the redshifted 21 cm frequency, $\nu=1420/(1+z)$ MHz, for $z \simg 6$."
+ The association of Type Ib/c supernovae (SNe) with GRBs has confirmed (he massive stellar origin of long GRBs (e.g..Zhang&Mészáros2003:2002).," The association of Type Ib/c supernovae (SNe) with GRBs has confirmed the massive stellar origin of long GRBs \citep[e.g.,][]{zhang03,meszaros02}."
+. Onlv. a small fraction of core collapse SNe appear to lead to GRD jets with high Lorentz factor 2o2100 (Bergerοἱal.2003:Soderberg.Frail.&Wieringa 2004)..," Only a small fraction of core collapse SNe appear to lead to GRB jets with high Lorentz factor $\gamma_{0} \simg 100$ \citep{berger03,soderberg04}. ."
+ However. there may be a larger inuuber of SNe accompanied by mildly relativistic ejecta. which could in principle involve more energy (han tvpical GRB jets (AlacFacdven.Woosley.&LegerRuiz 2004).," However, there may be a larger number of SNe accompanied by mildly relativistic ejecta, which could in principle involve more energy than typical GRB jets \citep{macfadyen01,granot04}."
+. A large kinetic energy.10 erg is also suggested in some core collapse SNe. by the optical light curve and spectral fits (Nomotoetal.2003:Poclsiadlowski2004).. and these objects appear in a number of cases not to be associated with detected GRB.," A large kinetic energy $\sim 10^{52}$ erg is also suggested in some core collapse SNe, by the optical light curve and spectral fits \citep{nomoto03,pod04}, and these objects appear in a number of cases not to be associated with detected GRB."
+ Such energetic supernovae (or hvpernovae. IINe) may be powered bv jets that fail to bore their wav through the stellar envelope (IxXhokhlovetal.1999;Mészáros&Rees2001).," Such energetic supernovae (or hypernovae, HNe) may be powered by jets that fail to bore their way through the stellar envelope \citep{khokhlov99,meszaros01}."
+. At high redshifts. stars are expected (o be very massive (Abel.Dryan.&Norman2002;Dromm.Coppi.&Larson2002:OumkaiPalla 2003)... ancl it max be more frequent for the jet to be choked even while the jet could be —10 times more energetic (Ileger 2001)..," At high redshifts, stars are expected to be very massive \citep{abel02,bromm02,omukai03}, and it may be more frequent for the jet to be choked even while the jet could be $\sim 10$ times more energetic \citep{heger03,fryer01}. ."
+ Based on these considerations. an isotropic equivalent energv ofἐνουLO’? ere may be reasonable for high redshift ILNe. and depending on (he energy conversion elliciency. even an," Based on these considerations, an isotropic equivalent energy of$E\sim 10^{53}$ erg may be reasonable for high redshift HNe, and depending on the energy conversion efficiency, even an"
+citedijets)) channels atDO.,) channels at.
+. The preliminary lower CL linits on the conrpactitication scale Me: of a single extra dineusiou sot in these analyses are 1.12 and 1.41 TeV respectively., The preliminary lower CL limits on the compactification scale $M_C$ of a single extra dimension set in these analyses are 1.12 and 1.4 TeV respectively.
+ Both the CDF and ccollaborations pioncered direct searches for the effects of IKaluza-dI&leim gravitons in the Randall-Sundrmn model in Run 2 25.26].," Both the CDF and collaborations pioneered direct searches for the effects of Kaluza-Klein gravitons in the Randall-Sundrum model in Run 2 \cite{RS1,RS2}."
+ As discussed above. the siuplest RS model is fixed by specifving just two parameters: the curvature of the AdS space & and the radius of compactification R.," As discussed above, the simplest RS model is fixed by specifying just two parameters: the curvature of the AdS space $k$ and the radius of compactification $R$."
+ Towever. from the experimental point of view. it is more couvenicut to work with au equivaleut set of paramcters. which correspond to direct observables: AA. tlie mass of the first Raluza-klem excitation of the eraviton. aud a dimensionless parameter k=kfAlp). which governs the coupling of the eravitous to the SM fields aud hence defines the internal width of the II& exavitous.," However, from the experimental point of view, it is more convenient to work with an equivalent set of parameters, which correspond to direct observables: $M_1$, the mass of the first Kaluza-Klein excitation of the graviton, and a dimensionless parameter $\tilde{k} = k/\overline{M}_{\rm Pl}$, which governs the coupling of the gravitons to the SM fields and hence defines the internal width of the KK gravitons."
+ Both the CDF aud CCollaborations used the dilepton and diphoton mass spectra aud searched for a narrow resonance. which would indicate a production of the first excitation of the KK eraviton.," Both the CDF and Collaborations used the dilepton and diphoton mass spectra and searched for a narrow resonance, which would indicate a production of the first excitation of the KK graviton."
+" The best scusitivity comes from the 1 combined ee25 analysis ([25| and 2.5 CDP μμ analysis. which exclude RS eravitous with masses below 0.9. 0.7. and 0.3 TeV for k=0.1. 0.05. and 0.01. respectively,"," The best sensitivity comes from the 1 combined $ee + \gamma\gamma$ analysis \cite{RS1} and 2.5 CDF $\mu\mu$ analysis, which exclude RS gravitons with masses below 0.9, 0.7, and 0.3 TeV for $\tilde{k} = 0.1$, 0.05, and 0.01, respectively."
+" Ultimate sensitivity to laree.1, and RS extra dimeusious will be achieved very soon at the LIIC."," Ultimate sensitivity to large, and RS extra dimensions will be achieved very soon at the LHC."
+ In inost of the cases. just ~1 of data is sufficieut to discover them or severely constrain parameters of these models.," In most of the cases, just $\sim 1$ of data is sufficient to discover them or severely constrain parameters of these models."
+ The LIIC would also allow to search for the IIx excitations of the gauge bosons directly., The LHC would also allow to search for the KK excitations of the gauge bosons directly.
+ Undoubtedly. the most exciting possibility is production of niui-black holes at the LIC |s} possible in theADD and RS models.," Undoubtedly, the most exciting possibility is production of mini-black holes at the LHC \cite{dlgt} possible in theADD and RS models."
+ Consider two partous with the ceuter-ofass energv Vs=Mon colliding ou.," Consider two partons with the center-of-mass energy $\sqrt{\hat s} =
+\MBH$ colliding head-on."
+ If the impact parameter of the collision is less than the Cuigher dimensional) Schwarzschild radius. correspondiug to this energy. a black hole (BIT) with the mass is formed.," If the impact parameter of the collision is less than the (higher dimensional) Schwarzschild radius, corresponding to this energy, a black hole (BH) with the mass is formed."
+ Therefore the total cross section of black hole production in particle collisions cau be estimated from pure ecometrical arguments aud is of order πο. DIT, Therefore the total cross section of black hole production in particle collisions can be estimated from pure geometrical arguments and is of order $\pi R_S^2$.
+ production is expected to be a threshold phenomena aud the onset is expected to happen for a iiniumui black hole mass Ap (ADD) or A (RS)., BH production is expected to be a threshold phenomena and the onset is expected to happen for a minimum black hole mass $\sim M_D$ (ADD) or $\Lambda_\pi$ (RS).
+ The total production cross section above this threshold at the LOC. ranges between 15 ub aud 1 pb for the Planck scale between 1 TeV aud 5 TeV. and varies by less than a factor of two for ο between 1 (RS) aud 7.," The total production cross section above this threshold at the LHC, ranges between 15 nb and 1 pb for the Planck scale between 1 TeV and 5 TeV, and varies by less than a factor of two for $n$ between 1 (RS) and 7."
+ Note that this cross section is comparable with. c.e.. ff production cross section. which result iu ~1 Uz signal event rate at the nominal LITC huninosity.," Note that this cross section is comparable with, e.g., $t\bar t$ production cross section, which result in $\sim 1$ Hz signal event rate at the nominal LHC luminosity."
+ Once7 produced. minifoot black holes quicklyH (~E10.D7° κ) evaporate via. IHEckius- radiation. |27|] with a characteristic temperature of ~100 GeV [s].," Once produced, mini black holes quickly $\sim 10^{-26}$ s) evaporate via Hawking radiation \cite{Hawking} with a characteristic temperature of $\sim 100$ GeV \cite{dlgt}. ."
+ The, The
+environment.,environment.
+ The result of this paper which is relevant al this point is that BIRGs seem {ο have. {ο a high confidence level. more large ancl close companions (D>10 Kpe. d<GOApe) than seviert 1s. and seem to be similar in their environment to Sevlert 2s (cf. (1999a))).," The result of this paper which is relevant at this point is that BIRGs seem to have, to a high confidence level, more large and close companions $D\geq 10$ Kpc, $d\leq 60Kpc$ ) than Seyfert 1s, and seem to be similar in their environment to Seyfert 2s (cf. \citet{dd99a}) )."
+ This statistical result gives support to a scheme that several workers have considered (IIeckmanetal.(1989).. and references therein: Sandersetal.(1988))).," This statistical result gives support to a scheme that several workers have considered \citet{he89}, and references therein; \citet{s88}) )."
+ The scheme is an evolutionary sequence for AGN driven by interaction: where the brace indicates that Svl and Sv2 may be actually the same kind of objects seen in different orientations., The scheme is an evolutionary sequence for AGN driven by interaction: where the brace indicates that Sy1 and Sy2 may be actually the same kind of objects seen in different orientations.
+ There are several lines of (achmittecdly ciretuustantial) evidence that also support (his simple evolutionary path: first. the contribution of thermal emission to the bolometric Iuminosity appears to decrease along (he sequence (Dullzin-LacvanRuano (1996))).," There are several lines of (admittedly circumstantial) evidence that also support this simple evolutionary path: first, the contribution of thermal emission to the bolometric luminosity appears to decrease along the sequence \citet{dd96}) )."
+ Sv1 nuclei have been revealed in several evolved mergers (for instance. see Ralanelli.Marziani.Birkle.&Thiele (1993))).," Sy1 nuclei have been revealed in several evolved mergers (for instance, see \citet{raf93}) )."
+ Second. there are several active galaxies in the DIRG sample.," Second, there are several active galaxies in the BIRG sample."
+ Of 87 galaxies. hhost a Sevfert 2 nucleus (15 objects). but only hhost a Sevlert 1 nucleus (2 objects).," Of 87 galaxies, host a Seyfert 2 nucleus (15 objects), but only host a Seyfert 1 nucleus (2 objects)."
+ There is no statistical difference between the ool active and non-active galaxies. except for a slightly hieher value in the Svl objects.," There is no statistical difference between the of active and non-active galaxies, except for a slightly higher value in the Sy1 objects."
+ The value of F(25jm)/F(60jan) for Sv2s is 220.18 and for Sv1Is zz0.20. compared {ος0.13 for non-aclive galaxies.," The value of $\mu$ $\mu$ m) for Sy2s is $\approx 0.18$ and for Sy1s $\approx 0.20$, compared $\approx 0.13$ for non-active galaxies."
+ F(25jmi)/F(60jm) is larger in Sy1s and Sv2s due to the contribution to the continuum of a non-thermal source (deGrijpetal. (1992)))., $\mu$ $\mu$ m) is larger in Sy1s and Sy2s due to the contribution to the continuum of a non-thermal source \citet{dg92}) ).
+ Estimating in a careful wav the ratio thermal/nonthermal emission for (the DIBRGs is not possible [rom published data., Estimating in a careful way the ratio thermal/nonthermal emission for the BIRGs is not possible from published data.
+ However. ool the Seyfert 2 (11 of 15) show evidence of significant star formation (there is evidence of a cireum-nuclear starburst in ool the star forming Sv2s).," However, of the Seyfert 2 (11 of 15) show evidence of significant star formation (there is evidence of a circum-nuclear starburst in of the star forming Sy2s)."
+ Also the 2 DIRG Sevfert I galaxies show evidence of a starburst (this makes selecting samples of Sevfert 1 and Sevfert 2 rom eeven more improper for environmental studies than selecting them from catalogues!), Also the 2 BIRG Seyfert 1 galaxies show evidence of a circum-nuclear starburst (this makes selecting samples of Seyfert 1 and Seyfert 2 from even more improper for environmental studies than selecting them from catalogues!)
+ The evolutionary sequence outlined above can be understood in three different wavs., The evolutionary sequence outlined above can be understood in three different ways.
+grains being the three species with the least destruction.,grains being the three species with the least destruction.
+" Although Fe grains had a higher survival rate than other dust species, the amount of Fe dust destruction was considerably higher than any of the simulations in our previous work."," Although Fe grains had a higher survival rate than other dust species, the amount of Fe dust destruction was considerably higher than any of the simulations in our previous work."
+" In the most extreme case, only of the Fe dust mass survived."," In the most extreme case, only of the Fe dust mass survived."
+" If we look at dust survival for two other commonly studied grain species for the same extreme case, and of the initial dust mass remains for carbonaceous (C) and silicate (8105) grains, respectively."," If we look at dust survival for two other commonly studied grain species for the same extreme case, and of the initial dust mass remains for carbonaceous (C) and silicate $\subscript{2}$ ) grains, respectively."
+" For example, consider the differences in dust survival across grain species."," For example, consider the differences in dust survival across grain species."
+" The difference, Agsuryival, between most destroyed grains and least destroyed grains ranges from as high as in Simulations 4 and 5 (Al,O3 compared to to as low as in Simulation 12 (again, AlgO3 Si)compared to Si)."," The difference, $\Delta \subscript{survival}$, between most destroyed grains and least destroyed grains ranges from as high as in Simulations 4 and 5 $\subscript{2}$ $\subscript{3}$ compared to Si) to as low as in Simulation 12 (again, $\subscript{2}$ $\subscript{3}$ compared to Si)."
+ This variation stems primarily from the initial distributions in grain radius., This variation stems primarily from the initial distributions in grain radius.
+ A high degree of destruction is observed in the species with a majority of their mass locked up in small grains., A high degree of destruction is observed in the species with a majority of their mass locked up in small grains.
+ Destruction is reduced significantly for those species with appreciable mass in large grains., Destruction is reduced significantly for those species with appreciable mass in large grains.
+" As noted in Paper 1, grains with initial radii less than 0.1 wm are easily obliterated, while larger grains take a considerable duration to be substantially sputtered."," As noted in Paper 1, grains with initial radii less than 0.1 $\mu$ m are easily obliterated, while larger grains take a considerable duration to be substantially sputtered."
+" If the actual size distributions of dust grains in SNRs were greatly disparate from the distributions assumed here, the final dust mass survival fractions might change considerably."," If the actual size distributions of dust grains in SNRs were greatly disparate from the distributions assumed here, the final dust mass survival fractions might change considerably."
+" In an attempt to understand the differences in the physical environments between the twelve simulations, we present phase-space diagrams of density and temperature to illustrate where the dust particles spend their time 6))."," In an attempt to understand the differences in the physical environments between the twelve simulations, we present phase-space diagrams of density and temperature to illustrate where the dust particles spend their time (Figure \ref{phaseplots}) )."
+ We also over-plot erosion rate contours for C (Figuregrains as a function of density and temperature to determine where in phase-space the, We also over-plot erosion rate contours for C grains as a function of density and temperature to determine where in phase-space the
+inclinations of planctesimals are not selfdamped by spiral density waves launched at the disk edge.,inclinations of planetesimals are not self-damped by spiral density waves launched at the disk edge.
+ The eccentricities and inclinations of plauctesimals could also be damped by dynamical friction caused by sinaller. unobserved bodies whose ταο velocities were damped bv collisions (2) or by resonant relaxation (?)..," The eccentricities and inclinations of planetesimals could also be damped by dynamical friction caused by smaller, unobserved bodies whose random velocities were damped by collisions \citep{2004G} or by resonant relaxation \citep{1998T}."
+ Wo asstuuce that neither of these effects contribute sienificautly during the era of iuterest., We assume that neither of these effects contribute significantly during the era of interest.
+ Finally. the ecceutricities aud melinatious of the planctesimals may not be excited in the first place if the planetesimal disk is massive cuough to cause Neptune's eccentricity to quickly precess. lowering the forced eccentricity of the cold classicals.," Finally, the eccentricities and inclinations of the planetesimals may not be excited in the first place if the planetesimal disk is massive enough to cause Neptune's eccentricity to quickly precess, lowering the forced eccentricity of the cold classicals."
+ Tu order for the disk to significantly affect the precession rate of the cold Classicals. its mass must be at least that of Neptune.," In order for the disk to significantly affect the precession rate of the cold classicals, its mass must be at least that of Neptune."
+ A scenario involving a quickly precessiug Neptune is explored iu 7 ancl discussed iu our conrpauiou paper. ?," A scenario involving a quickly precessing Neptune is explored in \citet{2011B} and discussed in our companion paper, \citet{2012D}."
+ We preseut here the results of a set of iuteerations contaiiug Neptune and a collection of test particles designed. to represeut the plauctesimals that vecolme fodav cold classical KDBOs., We present here the results of a set of integrations containing Neptune and a collection of test particles designed to represent the planetesimals that become today's cold classical KBOs.
+ The orbital evolution of Neptuues. is modeled as described in Section 2.2. and the Appendix., The orbital evolution of Neptune is modeled as described in Section \ref{subsec:modelnep} and the Appendix.
+ In Section L1.. we describe row we colputationally model the planetesinals.," In Section \ref{subsec:compmodel}, we describe how we computationally model the planetesimals."
+ In Section L2.. we characterize the secular excitation of he planctesimals under the influence of an inclined aud eccentric Neptune aud place limits ou the eccentricity alc inclination of Neptune.," In Section \ref{subsec:evolve}, we characterize the secular excitation of the planetesimals under the influence of an inclined and eccentric Neptune and place limits on the eccentricity and inclination of Neptune."
+ In Section L.3.. we determine he effect of damping of Neptuue’s eccentricity and inclination on the excitation of the plauctcsimals.," In Section \ref{subsec:damp}, we determine the effect of damping of Neptune's eccentricity and inclination on the excitation of the planetesimals."
+ We demonstrate that there are two regimes for campine Neptune's orbi in the slow-damping regnue. the planctesimals evolve to thei initial free ecceutricity and inclination. set by Neptune's initial eccentricity and oeinclination. aud in the fast-damping reece. their eccentricity and inclination are frozen at the value reached at the damping tine.," We demonstrate that there are two regimes for damping Neptune's orbit – in the slow-damping regime, the planetesimals evolve to their initial free eccentricity and inclination, set by Neptune's initial eccentricity and inclination, and in the fast-damping regime, their eccentricity and inclination are frozen at the value reached at the damping time."
+ We also show that the evolution of the iuclination and ecceutricity can be treated separately., We also show that the evolution of the inclination and eccentricity can be treated separately.
+ Tn Section LL. we consider the effects of Neptunes migration.," In Section \ref{subsec:migrate}, we consider the effects of Neptune's migration."
+ Finally. we preseut two example scenarios iu Section [.5 demonstrating the principles we have derived: a pathological. scenario i which the evolution of the Neptuue's semi-niajor axis. eccentricity. and inclination occur cach on a different. timescale. and a realistic scenario cousistent with preserving the cold population.," Finally, we present two example scenarios in Section \ref{subsec:ex} demonstrating the principles we have derived: a pathological scenario in which the evolution of the Neptune's semi-major axis, eccentricity, and inclination occur each on a different timescale, and a realistic scenario consistent with preserving the cold population."
+ We performed N-body integrations of au initially cold planetesimal disk under the imflucuce of a dvuamically-evolving Neptune using the hvbrid sviuplectic imteerator (7) with an accuracy parameter of 1017. a step size of 200 days and user- forces and velocitics imposing the mieratiou aud damping of Neptune. a complete description of which is provided in Section 2.2. and the Appendix.," We performed N-body integrations of an initially cold planetesimal disk under the influence of a dynamically-evolving Neptune using the hybrid symplectic integrator \citep{1999C} with an accuracy parameter of $10^{-12}$, a step size of 200 days, and user-defined forces and velocities imposing the migration and damping of Neptune, a complete description of which is provided in Section \ref{subsec:modelnep} and the Appendix."
+ We modeled the initial population of planetesimals. which become todav cold classical INBOs. as GOO inassless test particles withs initial & eveuly spaced between LO to 60 AU and initial ο=;/0.," We modeled the initial population of planetesimals, which become today's cold classical KBOs, as 600 massless test particles with initial $a$ evenly spaced between 40 to 60 AU and initial $e = i = 0$."
+ We will use the terms “planctesimals” and “test particles” interchangeably from here forward.," We will use the terms “planetesimals"" and “test particles"" interchangeably from here forward."
+ Although simulations that deliver KBOs iuto the classical region via scattering require ~10i plauctesimals in cach run. our niniediate goal of understanding the couditious wader which the cold classical population can be retained can be achieved witli a παο uuuber. allowing us to explore a wider rauge of orbital conditions.," Although simulations that deliver KBOs into the classical region via scattering require $\sim 10^{4}$ planetesimals in each run, our immediate goal of understanding the conditions under which the cold classical population can be retained can be achieved with a smaller number, allowing us to explore a wider range of orbital conditions."
+ lutegratious probing solar system histories απο typically yun for { Cr in order to test the stability of the system uuder current solar svsteim conditions., Integrations probing solar system histories are typically run for 4 Gyr in order to test the stability of the system under current solar system conditions.
+ ILlowever. the dynamical histories of Neptune we test are not dependent ou the loung-teriu survival of the planctesimals. because we have chosen observational criteria that are independent of the loug-terii evolutiou of the belt under the current configuration of the solar system (Sectou 3.1)).," However, the dynamical histories of Neptune we test are not dependent on the long-term survival of the planetesimals, because we have chosen observational criteria that are independent of the long-term evolution of the belt under the current configuration of the solar system (Secton \ref{sec:obs}) )."
+ To save computation time. our inteerationus probe onlv the period in which Neptune nuderegocs planetesinaledriven migration and/or ecceutricitv— damping aud/or inclination— damping.," To save computation time, our integrations probe only the period in which Neptune undergoes planetesimal-driven migration and/or eccentricity damping and/or inclination damping."
+— Iu scenarios of the carly solar svsteii in which Neptune is scattered onto au eccentric aud/or inclined orbit. Neptune can potentially disrupt au population of planctesimals above the coufined ecceutricities and iuclinations we observe (Secion 3.1)) iu the cold classical Iuper belt today.," In scenarios of the early solar system in which Neptune is scattered onto an eccentric and/or inclined orbit, Neptune can potentially disrupt an population of planetesimals above the confined eccentricities and inclinations we observe (Section \ref{sec:obs}) ) in the cold classical Kuiper belt today."
+ Neptune| forces the plauctesimals to nudereo secular evolution (Section 2.2)). in which they oscillate through bieh ecceuricities and inclinations.," Neptune forces the planetesimals to undergo secular evolution (Section \ref{subsec:sec}) ), in which they oscillate through high eccentricities and inclinations."
+ The rate of secular evolution depeuds ou the location of the Xauetesimual relative to Neptune (Fie. 1))., The rate of secular evolution depends on the location of the planetesimal relative to Neptune (Fig. \ref{fig:freqs}) ).
+ Fig 6 shows snapshots of the secular evolutiou of test particles uuder he influence of Neptune when the planet is stationary at 20 AU (top). migrates from 20 to 30 AU (uiddle). aud is stationary at 30 AU (bottom).," Fig \ref{fig:excite} shows snapshots of the secular evolution of test particles under the influence of Neptune when the planet is stationary at 20 AU (top), migrates from 20 to 30 AU (middle), and is stationary at 30 AU (bottom)."
+ Iu their secular evolution. he plauetesiuals reach maxima ο=2forecq aud {= Pitorecd ," In their secular evolution, the planetesimals reach maxima $e = 2 \eforced$ and $i = 2 \iforced$ ."
+In the middle pauel the evolution transitions roni mnatchine the top panel (early suapshots. loft) to natching the bottom panel (late snapshots. right).," In the middle panel, the evolution transitions from matching the top panel (early snapshots, left) to matching the bottom panel (late snapshots, right)."
+ Thus. in the case of mueration. we can cstimate the secular evolution using a constaut aa.," Thus, in the case of migration, we can estimate the secular evolution using a constant $a_N$."
+ When ey or fy has not vet damped. we can estimate the secular evolution as if ay were constaut at the location where Neptune spent most tine.," When $e_N$ or $i_N$ has not yet damped, we can estimate the secular evolution as if $a_N$ were constant at the location where Neptune spent most time."
+ Because we have chosen a migration law for which n decreases with time. planctesimals behave as ify has abwavs had the value αν).," Because we have chosen a migration law for which $\frac{\dot{a}_N}{a_N}$ decreases with time, planetesimals behave as if $a_N$ has always had the value $a_N(t)$ ."
+ We will discuss nueration in detail iu Section LL., We will discuss migration in detail in Section \ref{subsec:migrate}.
+ Tt Neptune's eccentricity aud inclination are small (Fig. T))," If Neptune's eccentricity and inclination are small (Fig. \ref{fig:obey}) ),"
+" the test particles remain below the observational lit ofο<0.1.76"" indefinitely."," the test particles remain below the observational limit of $e< 0.1, i< 6^\circ$ indefinitely."
+ Since. to first order. Heorecd 18 Independent of a. the planctcsimal’s position relative to Neptune. the planetesimals will remain below (L7 dE <3°.," Since, to first order, $\iforced$ is independent of $\alpha$ , the planetesimal's position relative to Neptune, the planetesimals will remain below $i < 6^\circ$ if $i_N < 3^\circ$."
+ The forced eccentricity does depeud on a (Fig. tj).," The forced eccentricity does depend on $\alpha$ (Fig. \ref{fig:freqs}) ),"
+ decreasing with the planctesimal’s distance frou Neptune., decreasing with the planetesimal's distance from Neptune.
+ Thus to satisfv our conservative criteria established in Section 3.1.. which requires planuctesimals with 12.5<a«I5 AU to remain at €<0.1. Neptune's ecceutricitv must stav below 0.09 at 20 AU or 0.06 at 30 AU.," Thus to satisfy our conservative criteria established in Section \ref{sec:obs}, which requires planetesimals with $42.5 < a < 45$ AU to remain at $e < 0.1$, Neptune's eccentricity must stay below 0.09 at 20 AU or 0.06 at 30 AU."
+ Tlowever.if ον aud ὃν damp due to dvuamical friction with the plauetesinals. the initial values of Neptune's cecentricity aud inclination can be higher.," However,if $e_N$ and $i_N$ damp due to dynamical friction with the planetesimals, the initial values of Neptune's eccentricity and inclination can be higher."
+ Iuthe next subsection. we discuss the effects of damping.," Inthe next subsection, we discuss the effects of damping."
+"llot. massive stars are not expected. a priori, to exhibit magnetic fields as they lack the convective envelopes necessary for driving a dynamo.","Hot, massive stars are not expected, a priori, to exhibit magnetic fields as they lack the convective envelopes necessary for driving a dynamo."
+ However. Ap/Bp stars (chemically peculiar A ancl B type stars) are well known to possess organized magnetic fields with surface streneths up to tens of kG (e.g. Borra Landstreet 1979).," However, Ap/Bp stars (chemically peculiar A and B type stars) are well known to possess organized magnetic fields with surface strengths up to tens of kG (e.g. Borra Landstreet 1979)."
+ The hottest of the Bp stars are the so-called. heliume-strong stars - early B-type stars on the main sequence with enhanced ancl typically variable helium. lines., The hottest of the Bp stars are the so-called helium-strong stars - early B-type stars on the main sequence with enhanced and typically variable helium lines.
+ Some helium-strong stars show additional emission and. variability in Balmer lines. variable photometric brightness and color. variable UV resonance lines. and non-thermal radio emission (Pedersen Vhomsen 1977: Walborn 1982: Shore Brown 1990: Leone Umana 1993).," Some helium-strong stars show additional emission and variability in Balmer lines, variable photometric brightness and color, variable UV resonance lines, and non-thermal radio emission (Pedersen Thomsen 1977; Walborn 1982; Shore Brown 1990; Leone Umana 1993)."
+ AL of these quantities vary with a single period. interpreted to be the rotational period of the star.," All of these quantities vary with a single period, interpreted to be the rotational period of the star."
+ The Balmer line (specifically Ho) and. UV line variability point to the presence of a stellar wind coupled with a magnetic field to form a circumstellar magnetosphere (Shore Brown 1990)., The Balmer line (specifically $\alpha$ ) and UV line variability point to the presence of a stellar wind coupled with a magnetic field to form a circumstellar magnetosphere (Shore Brown 1990).
+ Following their discovery of a magnetic field in the D2Vp star o Ori I (Landstreet Borra 1978). Borra," Following their discovery of a magnetic field in the B2Vp star $\sigma$ Ori E (Landstreet Borra 1978), Borra"
+the 610.3nnm abundance corrections for ALI)<4.0 are approximately constant for fixed stellar parameters.,the nm abundance corrections for $A\rm(Li)<4.0$ are approximately constant for fixed stellar parameters.
+ As described in Barklemetal.(2003).. including hydrogen collisions strengthens the collisional coupling between singly ionized lithium. and neutral lithium. especially through the charge transfer reaction involving H and Lill in the 3s state and its inverse reaction.," As described in \citet{Barklem03b}, including hydrogen collisions strengthens the collisional coupling between singly ionized lithium and neutral lithium, especially through the charge transfer reaction involving $\rm H$ and I in the 3s state and its inverse reaction."
+ The number populations of low excited states increase when charge transfer is included. while the line source functions are barely affected.," The number populations of low excited states increase when charge transfer is included, while the line source functions are barely affected."
+ In turn. the abundance corrections become generally lower. lless positive or more negative.," In turn, the abundance corrections become generally lower, less positive or more negative."
+ Notice that the added collisions do not always have a thermalising effect. but can contribute to over-population of levels compared to LTE.," Notice that the added collisions do not always have a thermalising effect, but can contribute to over-population of levels compared to LTE."
+ The effect is illustrated in 11. for a selected metal-poor and solar-metallicity dwarf and giant.," The effect is illustrated in 1, for a selected metal-poor and solar-metallicity dwarf and giant."
+ As seen in Fig., As seen in Fig.
+ |. the change in A(Li;aigi—Αν. when including charge transfer is 20.12 ddex for the metal-poor giant and —0.06 ddex for the metal-poor dwarf.," 1, the change in $A\rm(Li)_{\rm non-LTE}-\it A\rm(Li)_{\rm LTE}$ when including charge transfer is $-0.12$ dex for the metal-poor giant and $-0.06$ dex for the metal-poor dwarf."
+ While the charge transfer reaction has a significant influence on the statistical equilibrium of lithium. including bound-bound (and bound-free) transitions due to collisions with neutral hydrogen has no influence at all.," While the charge transfer reaction has a significant influence on the statistical equilibrium of lithium, including bound-bound (and bound-free) transitions due to collisions with neutral hydrogen has no influence at all."
+ However. this is not true if one relies on the classical Drawin recipe for estimates of the cross-section for those collisions.," However, this is not true if one relies on the classical Drawin recipe for estimates of the cross-section for those collisions."
+ For comparison. Fig.," For comparison, Fig."
+ | also shows the results obtained wher neglecting charge transfer reactions but including excitation and ionization by neutral hydrogen according to Drawin's recipe. with the rate formula given by Lambert(1993) ane with the scaling factor Sy=1.0.," 1 also shows the results obtained when neglecting charge transfer reactions but including excitation and ionization by neutral hydrogen according to Drawin's recipe, with the rate formula given by \citet{Lambert93} and with the scaling factor $S_{\rm H}=1.0$."
+ Generally. the Drawn cross-sections are much higher than the quantum mechanical calculations by Belyaev Barklem.," Generally, the Drawin cross-sections are much higher than the quantum mechanical calculations by Belyaev Barklem."
+ When adopting the higher classical rates. the ground state is more populated anc the abundance corrections become lower. by ~—O0.05 ddex. compared to the results obtained with quantum mechanical rates.," When adopting the higher classical rates, the ground state is more populated and the abundance corrections become lower, by $\sim-0.05$ dex, compared to the results obtained with quantum mechanical rates."
+ We have compared our results for the lithium resonance line to Carlssonetal.(1994) abundance corrections for dwarts and giants with [Fe/H]=[—3.0.0.0] to Pavlenko&Magazzu(1996) results for dwarfs and subgiants with solar metallicity and to Takeda&Kawanomoto(2005) corrections for dwarfs," We have compared our results for the lithium resonance line to \citet{Carlsson94} abundance corrections for dwarfs and giants with $\rm[Fe/H]=[-3.0,0.0]$ to \citet{Pavlenko96} results for dwarfs and subgiants with solar metallicity and to \citet{Takeda05} corrections for dwarfs"
+obtain unacceptable numbers of incorrect. matches.,obtain unacceptable numbers of incorrect matches.
+ In. the Bayesian scheme. the positional uncertainty is determined self consistently from the X-ray and optical data themselves.," In the Bayesian scheme, the positional uncertainty is determined self consistently from the X-ray and optical data themselves."
+ In addition. the simple scheme picks only the nearest cataloged optical counterpart.," In addition, the simple scheme picks only the nearest cataloged optical counterpart."
+ If there 1s a brighter optical source (which has a lower probability of being there by chance) slightly further away and the X-ray positional error is large. the Bayesian scheme may assign this to be the true optical counterpart.," If there is a brighter optical source (which has a lower probability of being there by chance) slightly further away and the X-ray positional error is large, the Bayesian scheme may assign this to be the true optical counterpart."
+ To further test the reliability of our Bayesian technique. we ran Monte Carlo simulations of the optical and X-ray datasets to model the problem of matching the X-ray point sources from the XBoóttes Survey to the NDWFS optical catalog.," To further test the reliability of our Bayesian technique, we ran Monte Carlo simulations of the optical and X-ray datasets to model the problem of matching the X-ray point sources from the XBoöttes Survey to the NDWFS optical catalog."
+ The Monte Carlo results provide confidence in our overall approach and a basis for interpreting the results for the real NDWFS/XBoéttes data., The Monte Carlo results provide confidence in our overall approach and a basis for interpreting the results for the real NDWFS/XBoöttes data.
+ In particular. we use the Monte Carlo tests to understand the completeness of the identifications (hereafter IDs).," In particular, we use the Monte Carlo tests to understand the completeness of the identifications (hereafter IDs)."
+ We modeled the X-ray positional uncertainties by a Gaussian (Eqn. A9)), We modeled the X-ray positional uncertainties by a Gaussian (Eqn. \ref{eqn:psfmod}) )
+" with a constant. 7,2075. the PSF width at the pointing center. 7=075. and the quadratic growth of the PSF width as one moves off-axis. 05002470."," with a constant, $\sigma_{a}=0\farcs5$, the PSF width at the pointing center, $\sigma_0=0\farcs5$, and the quadratic growth of the PSF width as one moves off-axis, $\sigma_{600}=4\farcs0$."
+ We produced synthetic background catalogs by taking the XBoóttes catalog and generating new optical catalogs for the regions around each X-ray source., We produced synthetic background catalogs by taking the XBoöttes catalog and generating new optical catalogs for the regions around each X-ray source.
+ We used broken power-law models for the number counts. empirically normalized to match the observed number counts of the NDWFS field.," We used broken power-law models for the number counts, empirically normalized to match the observed number counts of the NDWFS field."
+ The total number of sources was normalized to match the ~[0° observed sources: in the matching regions used for the X-ray sources., The total number of sources was normalized to match the $\sim 10^5$ observed sources in the matching regions used for the X-ray sources.
+ The X- sources were modeled with three regions. ©=0.4 for 15«in<20. à=0 for 20«m23. and à=—0.3 for 23«m26. based on the magnitudes of the most probable optical IDs for the X-ray sources in the central d;<40070 regions of the ACIS fields.," The X-ray sources were modeled with three regions, $\alpha=0.4$ for $15 < m < 20$, $\alpha=0$ for $20 < m < 23$, and $\alpha=-0.3$ for $23 < m < 26$, based on the magnitudes of the most probable optical IDs for the X-ray sources in the central $d_k < 400\farcs0$ regions of the ACIS fields."
+ The background optical source positions. were distributed randomly in the matching region for each X-ray source., The background optical source positions were distributed randomly in the matching region for each X-ray source.
+ To test our ability to recognize X-ray sources genuinely lacking optical counterparts. we assumed that a fraction (1— of the X-ray sources have no optical counterpart.," To test our ability to recognize X-ray sources genuinely lacking optical counterparts, we assumed that a fraction $1-f$ ) of the X-ray sources have no optical counterpart."
+" For the f£)sources with counterparts. the position of the counterpart was randomly drawn based on the “PSF model"" using the observed X-ray counts and off-axis distance"," For the sources with counterparts, the position of the counterpart was randomly drawn based on the “PSF model” using the observed X-ray counts and off-axis distance"
+It is well known that the concentration parameter of halos al a given mass. e(M). has a broad log-normal distribution. and is correlated with the halo formation un m in such a wav that halos of earlier formation have a higher concentration (Jing2000:--Bul-al.2001:Jing&Suto2002:WechsleretZhaoοἱ 2003a.,"It is well known that the concentration parameter of halos at a given mass, $c(M)$, has a broad log-normal distribution, and is correlated with the halo formation epoch $z_{\rm form}$ in such a way that halos of earlier formation have a higher concentration \citep{jing00,bullock01,js02,Wechsler02,zhao03a,zhao03b}."
+b).. the age dependence of halo bias should vield a concentration dependence of halo bias., Therefore the age dependence of halo bias should yield a concentration dependence of halo bias.
+ such dependence is indeed observed in recent N-bocly simulations., Such dependence is indeed observed in recent $N$ -body simulations.
+" Wechsler found that. for masses A<AL,. halos with higher concentrations are more strongly clustered. as expected from the age dependence aud (he correlation between the age ancl concentration of dark halos."," \citet{Wechsler06} found that, for masses $M<M_\ast$, halos with higher concentrations are more strongly clustered, as expected from the age dependence and the correlation between the age and concentration of dark halos."
+" More Wechsleretal.(2006) found reversed concentration dependence for M>©M,.Interestingly. in the sense (hat halos with higher concentrations are actually less biased."," More interestingly, \citet{Wechsler06} found reversed concentration dependence for $M>M_\ast$, in the sense that halos with higher concentrations are actually less biased."
+ This result wasnot seen by Gaoetal.(2005) who explored the age dependence only for AfxLOM..., This result wasnot seen by \citet{gao05} who explored the age dependence only for $M\le 10M_\ast$.
+ More recently. Wetzeletal.(2006) examined both the age dependence ancl concentration dependence of halo clustering for massive halos.," More recently, \citet{white06} examined both the age dependence and concentration dependence of halo clustering for massive halos."
+ While they found concentration dependence similar to Chat found by Wechsleretal.(2006).. thev did nol detect anv significant dependence on formation epoch.," While they found concentration dependence similar to that found by \citet{Wechsler06}, , they did not detect any significant dependence on formation epoch."
+ Note. however. (hat (2006) used halos identified at different redshifts to increase the dvnamical range probed by their simulations. while both Gaoetal.(2005) ancl Wetzeletal.(2006) used. halos identified at the same Gime.," Note, however, that \citet{Wechsler06} used halos identified at different redshifts to increase the dynamical range probed by their simulations, while both \citet{gao05} and \citet{white06}
+ used halos identified at the same time."
+ It is unclear how (o make a detailed comparison between the different results., It is unclear how to make a detailed comparison between the different results.
+ The environmental dependence of halo Formation aud structure has important implications not only for improving the PS theory but also for improving semi-analvtical models of galaxy formation based on PS merger trees and current halo occupation distribution (10D) mocleV., The environmental dependence of halo formation and structure has important implications not only for improving the PS theory but also for improving semi-analytical models of galaxy formation based on PS merger trees and current halo occupation distribution (HOD) models.
+ In thisLeller. we use a large set of cosmological simulations. each with 1024* particles. to investigate in detail (he concentration and formation epoch dependence of halo clustering.," In this, we use a large set of cosmological simulations, each with $1024^3$ particles, to investigate in detail the concentration and formation epoch dependence of halo clustering."
+" The model considered here is a canonical spatiallv-flat Cold Dark MM (CDM) model with the parameter O,,=0.263. the constant O4=0.732. MN(the IIubble constant --Uemf=0.71. ancl (he baryon density. parameterMN€;=0.045."," The model considered here is a canonical spatially-flat Cold Dark Matter (CDM) model with the density parameter $\Omega_m=0.268$, the cosmological constant $\Omega_\Lambda=0.732$, the Hubble constant $h=0.71$, and the baryon density parameter $\Omega_b=0.045$."
+ The clensily fielcl is to be Gaussian wilh a seale-invariant. power spectrum xA., The primordial density field is assumed to be Gaussian with a scale-invariant power spectrum $\propto k$.
+ For the linear spectrum. we adopt the fitting transfer ΠΙΟ of Eisenstein&Iu(1998)... ancl the normalization is sel bv σε= 0.85. where σς is the present linear RAIS density [Iuctuation within a sphere of radius 8h IMpe.," For the linear spectrum, we adopt the fitting transfer function of \citet{eh98}, and the normalization is set by $\sigma_8=0.85$ , where $\sigma_8$ is the present linear RMS density fluctuation within a sphere of radius $8\mpchi$ ."
+ We use an upgraded. version of the Particle-Particle-Particle-Mesh (DPM) code of to simulate structure formation in the universe.," We use an upgraded version of the Particle-Particle-Particle-Mesh $\pppm$ ) code of \citet{js98,js02} to simulate structure formation in the universe."
+ Thecode hasnow, Thecode hasnow
+"We set where L, is the value of L, at the coronal base. so that {νι increases in. proportion to the eross-sectional radius of the flux tube in our model.","We set where $L_{\perp \sun}$ is the value of $L_\perp$ at the coronal base, so that $L_\perp$ increases in proportion to the cross-sectional radius of the flux tube in our model."
+" In addition to linear damping. KAWs at Kjpy~| undergo nonlinear damping through the ""stochastic heating"" of protons (2??).."," In addition to linear damping, KAWs at $k_\perp \rho_{\rm p} \sim 1$ undergo nonlinear damping through the “stochastic heating” of protons \citep{mcchesney87,chen01,johnson01}."
+ In stochastic proton heating. AW/KAW fluctuations at Kip~| cause proton orbits in the plane perpendicular to By to become stochastic. and the protons are subsequently energized by the time-varying electrostatic potential.," In stochastic proton heating, AW/KAW fluctuations at $k_\perp \rho_{\rm p} \sim 1$ cause proton orbits in the plane perpendicular to $\bm{B}_0$ to become stochastic, and the protons are subsequently energized by the time-varying electrostatic potential."
+" Using numerical simulations of test particles interacting with a spectrum of randomly phased AWs and KAWs. ? found that the effective damping rate of KAWs atk, Pp~ Land pipXI from stochastic proton heating is where c» Is a dimensionless constant. and vip=yDyTip is the proton perpendicular thermal speed."," Using numerical simulations of test particles interacting with a spectrum of randomly phased AWs and KAWs, \cite{chandran10a} found that the effective damping rate of KAWs at $k_\perp \rho_{\rm p} \sim 1$ and $\beta_{\parallel \rm p} \lesssim 1$ from stochastic proton heating is where $c_2$ is a dimensionless constant, and $v_{\perp \rm p} = \sqrt{2 k_{\rm B} T_{\perp \rm p}/m_{\rm p}}$ is the proton perpendicular thermal speed."
+ ? found that /nityc»=0.34 for randomly phased AWs and KAWs. but conjectured that c» is smaller (and hence stochastic heating is more effective) in. strong. AW/KAW turbulence. because much of the dissipation in. strong AW/KAW turbulence occurs within. coherent structures in which the fluctuation amplitudes are larger than. their. rms values (?)..," \cite{chandran10a} found that $c_2
+= 0.34$ for randomly phased AWs and KAWs, but conjectured that $c_2$ is smaller (and hence stochastic heating is more effective) in strong AW/KAW turbulence, because much of the dissipation in strong AW/KAW turbulence occurs within coherent structures in which the fluctuation amplitudes are larger than their rms values \citep{dmitruk04}."
+ In the vicinity of such structures. proton orbits are more stochastic than on average. allowing for more efficient stochastic heating.," In the vicinity of such structures, proton orbits are more stochastic than on average, allowing for more efficient stochastic heating."
+ ? developed a model of ton temperatures in coronal holes based on stochastic heating. and this model matched the observed O? temperature profile when e» was set equal to 0.15.," \cite{chandran10b} developed a model of ion temperatures in coronal holes based on stochastic heating, and this model matched the observed ${\rm O}^{+5}$ temperature profile when $c_2$ was set equal to 0.15."
+ lr our model. we leave c» às a free parameter.," In our model, we leave $c_2$ as a free parameter."
+ The total effective damping rate of KAWs at Kjpp=| is We define [ to be the fraction of the cascade power that is dissipated at ki~1., The total effective damping rate of KAWs at $k_\perp \rho_{\rm p} =1 $ is We define $\Gamma$ to be the fraction of the cascade power that is dissipated at $k_\perp \rho_{\rm p} \sim 1$.
+ This fraction is roughly Yo when VoteX| and roughlypp| when yout>>|., This fraction is roughly $\gamma_{\rm tot} t_{\rm c}$ when $\gamma_{\rm tot} t_{\rm c} \ll 1$ and roughly1 when $\gamma_{\rm tot} t_{\rm c} \gg 1$.
+" To interpolate smoothly between these limits. we set Landau damping and transit-time damping of KAWs on protons contribute to Qj, but not to Qj; (?).."," To interpolate smoothly between these limits, we set Landau damping and transit-time damping of KAWs on protons contribute to $Q_{\parallel \rm p}$ but not to $Q_{\perp \rm
+ p}$ \citep{stix92}."
+ On the other hand. stochastic heating leads primarily to perpendicular proton heating when By«| (y.," On the other hand, stochastic heating leads primarily to perpendicular proton heating when $\beta_{\parallel \rm
+ p} \ll 1$ \citep{chandran10a}."
+ 2 have shown that stochastic heating leads to significant perpendicular proton heating even at βιp [and for simplicity we take stochastic heating to contribute only to regardless of the value of ο.," \cite{johnson01} have shown that stochastic heating leads to significant perpendicular proton heating even at $\beta_{\parallel \rm p} \sim 1$, and for simplicity we take stochastic heating to contribute only to $Q_{\perp \rm p}$, regardless of the value of $\beta_{\parallel \rm p}$."
+ We divide the cascade Ορ.power that is dissipated at AiPpβι~I into three parts — electron heating. perpendicular proton heating. and parallel proton heating — in proportion to the corresponding damping rates. Yo. Ys. and yy.," We divide the cascade power that is dissipated at $k_\perp \rho_{\rm p} \sim 1$ into three parts — electron heating, perpendicular proton heating, and parallel proton heating — in proportion to the corresponding damping rates, $\gamma_{\rm e}$, $\gamma_{\rm s}$, and $\gamma_{\rm p}$."
+ As mentioned previously. we assume that the fraction of the cascade power that is not dissipated at Kjpp—1 cascades to scales «pp. dissipates via interactions. with. electrons. and contributes to ος.," As mentioned previously, we assume that the fraction of the cascade power that is not dissipated at $k_\perp \rho_{\rm p} \sim 1$ cascades to scales $\ll \rho_{\rm p}$, dissipates via interactions with electrons, and contributes to $Q_{\rm e}$ ."
+ This procedure leadsto the relations and We integrate Equations (10)) through (17)) forward in time until a steady state is reached usingthe implicit numerical method described by ?.. We use a logarithmicgrid. in r with grid.points 7; extending from ΓΑ. to 1.2 AU. where i—0.1.2.....N|1 and N=1320.," This procedure leadsto the relations and We integrate Equations \ref{eq:cont1}) ) through \ref{eq:dEwdt}) ) forward in time until a steady state is reached usingthe implicit numerical method described by \cite{hu97}.. We use a logarithmicgrid in $r$ with gridpoints $r_i$ extending from $1 \;R_{\sun}$ to $1.2$ AU, where $i= 0, 1, 2, \dots, N+1$ and $N=1320$."
+ At each time step. we update the variables by integrating the equations only at FENry.," At each time step, we update the variables by integrating the equations only at $r_1, r_2, \dots, r_N$."
+ We then update the variables at ry and νε] using the following boundary conditions., We then update the variables at $r_0$ and $r_{N+1}$ using the following boundary conditions.
+" At r=ry we set n-n..T. TipTpT.. and ‘fy=n.Hy(Ov,πι. where a..T. and v. are constants (see Table 1))."," At $r=r_0$ we set $n= n_{\sun}$, $T_e = T_{\perp \rm p} = T_{\parallel \rm p} =
+T_{\sun}$, and ${\cal E}_{\rm w} = n_{\sun} m_{\rm p} (\delta
+v_{\sun})^2$ , where $n_{\sun}$, $T_{\sun}$, and $\delta v_{\sun}$ are constants (see Table \ref{tab:parameters}) )."
+" We determine U. 41p. and qj at r=ro by linearly extrapolating from the values atr =r, and r=ro."," We determine $U$ , $q_{\perp \rm p}$, and $q_{\parallel \rm p}$ at $r=r_0$ by linearly extrapolating from the values at $r=r_1$ and $r=r_2$ ."
+ We determine U(509) in this way rather than by fixing the value of CU(ro) so that the variables can evolve towards a steady-state transonic solution that passes smoothly through the sonic point., We determine $U(r_0)$ in this way rather than by fixing the value of $U(r_0)$ so that the variables can evolve towards a steady-state transonic solution that passes smoothly through the sonic point.
+" We determine qj and qj, at r—rg by linear extrapolation for the following reason.", We determine $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ at $r=r_0$ by linear extrapolation for the following reason.
+" At radii slightly greater than ry. the values of gi, and gp are determined to a good approximation by neglecting the terms on the left-hand sides of Equations (15)) and (16)). because the collision time Vy! ismuch shorter than the expansion time r/U."," At radii slightly greater than $r_0$ , the values of $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ are determined to a good approximation by neglecting the terms on the left-hand sides of Equations \ref{eq:dqperpdt}) ) and \ref{eq:dqpardt}) ), because the collision time $\nu_{\rm p}^{-1}$ ismuch shorter than the expansion time $r/U$."
+" We call the values of gi) and qj, determined in this way the “collisional values.”", We call the values of $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ determined in this way the “collisional values.”
+" [f boundary conditions were Imposed on qp and qj, at r—7o that were different from the collisional values. then in steady state a boundary layer would develop at r=ry of thickness ~UGVpo. Where Uy=U(ro) and vpo=Vp (vo). and gp and 4p would approach their collisional values at 7~ry| "," If boundary conditions were imposed on $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ at $r=r_0$ that were different from the collisional values, then in steady state a boundary layer would develop at $r=r_0$ of thickness $\sim U_0/\nu_{\rm p0}$, where $U_0 = U(r_0)$ and $\nu_{\rm p0} = \nu_{\rm p}(r_0)$ , and $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ would approach their collisional values at $r\sim r_0 + U_0/\nu_{\rm p0}$."
+"However. by linearly extrapolating the values of gi)Uy/Vpo. and qj, to r—ry. we prevent such unphysical boundary layers from appearing."," However, by linearly extrapolating the values of $q_{\perp \rm p}$ and $q_{\parallel \rm p}$ to $r=r_0$, we prevent such unphysical boundary layers from appearing."
+ At F—ΑΙ we evaluate all variables by linearly extrapolating from their values at 7=ry; and r= ry.For the solutions presented in Sections 4 and 5.7 we used the following initialconditions: f;=TigἩνT.(32R./r)tr/R. 7. U—(655km/s)|l|20(R. Fin9n.Ugac (Ua). P. and qq=Gp/r— 9.," At $r=r_{N+1}$ we evaluate all variables by linearly extrapolating from their values at $r=r_{N-1}$ and $r=r_{N}$ .For the solutions presented in Sections \ref{sec:solution}
+ and \ref{sec:swani_cyclotron} we used the following initialconditions: $T_{\rm e} = T_{\perp \rm p} = T_{\parallel \rm p} =
+T_{\sun} (3-2R_{\sun}/r) (r/R_{\sun})^{-2/7} $ , $U = (655 \mbox{
+ km/s})[1 + 20 (R_{\sun}/r)^{3}]^{-1}$ , $n = n_{\sun} U_0 a_{\sun}/(U
+a)$ ,${\cal E}_{\rm w} = n m_{\rm p} (\delta v_{\sun})^2 $ , and $q_{\perp \rm p} = q_{\parallel \rm p} = 0$ ."
+ In this section. we focus on a steady-state solution toEquations (10))through (17)) in which the model parameters areset equal to the values listed in Table 1..," In this section, we focus on a steady-state solution toEquations \ref{eq:cont1}) )through \ref{eq:dEwdt}) ) in which the model parameters areset equal to the values listed in Table \ref{tab:parameters}. ."
+" Our choice for L,. is motivated by Faraday-rotation measurements along linesof sight passing through the coronaand near-Sun solar wind (2)..", Our choice for $L_{\perp \sun}$ is motivated by Faraday-rotation measurements along linesof sight passing through the coronaand near-Sun solar wind \citep{hollweg10}.
+ After choosing the super-radial expansion factors, After choosing the super-radial expansion factors
+The timescale is chosen such that the sound speed is of order unity.,The timescale is chosen such that the sound speed is of order unity.
+ The sounc speed has been calculated as 3.26⋅−⋅.IO1ems loch.," The sound speed has been calculated as $c_{\rm s} = 3.26 \times 10^4 \,\mbox{cm} \,\mbox{s}^{-1}$ ."
+ This implies. a time. scale Using these units. the velocity of the plasma [iow is now Finally. a mass scale is chosen.," This implies a time scale Using these units, the velocity of the plasma flow is now Finally, a mass scale is chosen."
+ This is set so that the mass density of the plasma is of order unity., This is set so that the mass density of the plasma is of order unity.
+ The mass density has been calculated as po=3.8910Pegem 7.," The mass density has been calculated as $\rho_0 = 3.89 \times 10^{-18} \, \mbox{g} \, \mbox{cm}^{-3}$ ."
+ Using the length scale chosen above. this gives a characteristic mass scale The magnitude of the magnetic field. by=22.8µία. can also be transformed into a dimensionless quantity. by From this calculation onwards. the units of B. E and J are written such that the factors of ἐπ and the speed of light no longer appear.," Using the length scale chosen above, this gives a characteristic mass scale The magnitude of the magnetic field, $B_0 = 22.8\, \mu \mbox{G}$, can also be transformed into a dimensionless quantity, by From this calculation onwards, the units of $\mathbf{B}$, $\mathbf{E}$ and $\mathbf{J}$ are written such that the factors of $4 \pi$ and the speed of light no longer appear."
+ Lt is possible now to transform the conductivities. and the resulting resistivities. into cdimensionless units.," It is possible now to transform the conductivities, and the resulting resistivities, into dimensionless units."
+" The Following values are found: These values can be seen to correspond το those implemented in Paper L Ehe amount of ambipolar resistivity is equivalent to. the. simulation with high. ambipolar resistivity in Paper lL (ambi-high-hr) with magnetic Revnolds number Rey,—28.4. while the amount of Πα resistivity is equivalent to the simulation with moderate Llall resistivity in Paper E (hall-med-br) with magnetic Revnolds number Rey,=284."," The following values are found: These values can be seen to correspond to those implemented in Paper I. The amount of ambipolar resistivity is equivalent to the simulation with high ambipolar resistivity in Paper I (ambi-high-hr) with magnetic Reynolds number $Re_{\rm m} = 28.4$, while the amount of Hall resistivity is equivalent to the simulation with moderate Hall resistivity in Paper I (hall-med-hr) with magnetic Reynolds number $Re_{\rm m} = 284$."
+ As in Paper LL these simulations are carried out on a DD slab erid in the wy-plane.," As in Paper I, these simulations are carried out on a D slab grid in the $xy$ -plane."
+ The initial set-up used. was that of two plasmas llowing anti-parallel side-by-side on a erid of sizer=0.32L] and y=0.L].," The initial set-up used was that of two plasmas flowing anti-parallel side-by-side on a grid of size $x = [0, 32L]$ and $y = [0, L]$."
+ The plasma velocities are eiven. by |2M and — EMin the yecireetion., The plasma velocities are given by $+ \frac{V_0}{2}$ and $-\frac{V_0}{2}$ in the $y$ -direction.
+". rpThe unit. of Mtime used is the sound crossing time. /,=4, the time taken for a signal travelling at the speed. of sound to cross the grid in he y-clireetion."," The unit of time used is the sound crossing time, $t_{\rm s} \equiv \frac{L}{c_{\rm s}}$, the time taken for a signal travelling at the speed of sound to cross the grid in the $y$ -direction."
+ We use periocic boundary. conditions at the high and ow y boundaries., We use periodic boundary conditions at the high and low $y$ boundaries.
+ Since we wish to study not only the initial erowth phase of the instability. but also its subsequent non-inear behaviour we must ensure that waves interacting with he high and low ο boundaries do not rellect back into the domain to influence the ecinamics.," Since we wish to study not only the initial growth phase of the instability, but also its subsequent non-linear behaviour we must ensure that waves interacting with the high and low $x$ boundaries do not reflect back into the domain to influence the dynamics."
+" Previous test simulations or various parameters have shown that a large width of 294, is necessary to ensure this.", Previous test simulations for various parameters have shown that a large width of $32L$ is necessary to ensure this.
+ We use gradient. zero boundary conditions at the high and low ος and z boundaries., We use gradient zero boundary conditions at the high and low $x$ and $z$ boundaries.
+ The grid is chosen therefore to consist of cells. in the .r. jy. and z directions respectively.," The grid is chosen therefore to consist of $6400 \times 200 \times 1$ cells, in the $x$, $y$, and $z$ directions respectively."
+ Εις resolution was chosen on the basis that it reproduces the initial linear. growth of the ideal MIEID. system in Ixeppens, This resolution was chosen on the basis that it reproduces the initial linear growth of the ideal MHD system in \cite{keppens99}.
+ Resolution studies were performed to confirm the resolution as being appropriate for multilluid: MED. as well (see refsee:valiclation-resolution))., Resolution studies were performed to confirm the resolution as being appropriate for multifluid MHD as well (see \\ref{sec:validation-resolution}) ).
+ The plasma velocity field is initiated with a tangential shear laver of width 2e at the interface at16L., The plasma velocity field is initiated with a tangential shear layer of width $2a$ at the interface at.
+ This velocity profile is described by The width of the shear laver is chosen to be fo= 0.05. or approximately 20 grid zones.," This velocity profile is described by The width of the shear layer is chosen to be $\frac{a}{L} =
+0.05$ , or approximately 20 grid zones."
+ The magnetic field is initially set to be uniform and aligned with the plasma Low., The magnetic field is initially set to be uniform and aligned with the plasma flow.
+ The initial backeround for all four [luids in the svsteni is now an exact equilibrium., The initial background for all four fluids in the system is now an exact equilibrium.
+ The initial velocity field. Vi is then augmented with a perturbation given by where Vy is set to. 10το," The initial velocity field, $V_0$ is then augmented with a perturbation given by where $\delta V_0$ is set to $10^{-4}\,V_0$."
+ The wavelength of the »erturbation is set equal to the characteristic length scale. z-o-—L. so that a single wavelength. fits exactly. into he computational domain.," The wavelength of the perturbation is set equal to the characteristic length scale, $\lambda
+= \frac{2 \pi}{k_y} = L$, so that a single wavelength fits exactly into the computational domain."
+ This maximises the possibility of resolving structures that are small relative to. the initial perturbed: wavelength. (Frank.ctal.1996).," This maximises the possibility of resolving structures that are small relative to the initial perturbed wavelength, \citep{frank96}."
+".. The »erturbation attenuation scale is chosen so that it is larger han the shear laver. but small enough so that the instability canbe assumed to interact only minimally with the a xindaries (seePalottietal.2008).. and is set using ""=02 (seeIxeppensetal.1999:Palotti2008)."," The perturbation attenuation scale is chosen so that it is larger than the shear layer, but small enough so that the instability canbe assumed to interact only minimally with the $x$ -boundaries \citep[see][]{palotti08}, and is set using $\frac{\sigma}{L} = 0.2$ \citep[see][]{keppens99, palotti08}."
+. Finally. the wavenumber & is chosen to be 27 in order to maximise the erowth rate of the instability (Ixeppensetal.1999).," Finally, the wavenumber $k_y$ is chosen to be $2 \pi$ in order to maximise the growth rate of the instability \citep{keppens99}."
+. We note that our choice of periodic boundary conditions on the 5g boundaries limits the behaviour of the system to some extent: wavelengths longer than the erid domain can not grow and coalescence of large-scale vortices will not occur (c.g.Batyetal.2003)., We note that our choice of periodic boundary conditions on the $y$ boundaries limits the behaviour of the system to some extent: wavelengths longer than the grid domain can not grow and coalescence of large-scale vortices will not occur \citep[e.g.][]{baty03}.
+.. Furthermore. our choice of gradient zero. boundary. conditions in the vw direction implies that we are studying only the surface. modes. of the instability.," Furthermore, our choice of gradient zero boundary conditions in the $x$ direction implies that we are studying only the surface modes of the instability."
+ “Phis latter point. however. is not a major restriction as the body modes. will not be important. in. for example. the development of the KL instability in the bowshocks created by voung stellar outllows.," This latter point, however, is not a major restriction as the body modes will not be important in, for example, the development of the KH instability in the bowshocks created by young stellar outflows."
+ The set-up described above has been shown in Paper I to be valid for producing the WIL instability under the ideal AILLD approximation through comparisons with previously published. literature., The set-up described above has been shown in Paper I to be valid for producing the KH instability under the ideal MHD approximation through comparisons with previously published literature.
+ Both hyerodynamic and. ideal. ΑΔΗ simulations were run using the IIYDRA code., Both hydrodynamic and ideal MHD simulations were run using the code.
+ C'omparisons between the linear growth of the simulated. hydrodynamic instability and the growth rate calculated analytically for the same wavenumber by Miura.&Pritchett(1982) , Comparisons between the linear growth of the simulated hydrodynamic instability and the growth rate calculated analytically for the same wavenumber by \citet{miura82} 
+In order to model the coronae of E. Tauri stars we need to assume something about the form of the magnetic field.,In order to model the coronae of T Tauri stars we need to assume something about the form of the magnetic field.
+ Observations suggest it is compact and inhomogeneous and may vary not only with time on cach star. but also [fron one star to the next.," Observations suggest it is compact and inhomogeneous and may vary not only with time on each star, but also from one star to the next."
+ To capture this behaviour. we use as examples the field structures of two cillerent main sequence stars. LO. Lva and. AB Dor determined. from Zeeman- imaging (?)..," To capture this behaviour, we use as examples the field structures of two different main sequence stars, LQ Hya and AB Dor determined from Zeeman-Doppler imaging \citep{don03}."
+ Although we cannot be certain whether or not the magnetic field structures extrapolated from surface magnetograms of voung main sequence stars do represent the magnetically confined coronae of TE Tauri stars. they do satisfy the currently available observational constraints.," Although we cannot be certain whether or not the magnetic field structures extrapolated from surface magnetograms of young main sequence stars do represent the magnetically confined coronae of T Tauri stars, they do satisfy the currently available observational constraints."
+ In future it will be possible to use real T Jauri magnetograms derived from. Zeeman-Doppler images obtained using the ESPaDOnS instrument at the Canacla-France-Llawaii telescope. (2).., In future it will be possible to use real T Tauri magnetograms derived from Zeeman-Doppler images obtained using the ESPaDOnS instrument at the Canada-France-Hawaii telescope \citep{don07}.
+ However. in the meantime. the example field. geometries. used. in this letter (see Fig. 11} ," However, in the meantime, the example field geometries used in this letter (see Fig. \ref{coronae}) )"
+capture the essential features of E Fauri coronae., capture the essential features of T Tauri coronae.
+ They reproduce X-ray emission. measures (Als) and corona densities which are tvpical of T Lauri stars (2).., They reproduce X-ray emission measures (EMs) and coronal densities which are typical of T Tauri stars \citep{jar06}.
+ Phe surface field. structures are complex. consistent with polarisation nmieasurements (72). and N-ray emitting plasma is confines within unevenly distributed. magnetic structures close to the stellar surface. giving rise to significant rotationa moclulation of X-ray emission (7)..," The surface field structures are complex, consistent with polarisation measurements \citep{val04} and X-ray emitting plasma is confined within unevenly distributed magnetic structures close to the stellar surface, giving rise to significant rotational modulation of X-ray emission \citep{gre06b}."
+ We extrapolate from surface magnetoerams by assuming that the magnetic field. B. is potential such tha Vo.D—0., We extrapolate from surface magnetograms by assuming that the magnetic field $\bmath{B}$ is potential such that $\nabla \times \bmath{B} = 0$.
+ This condition is satisfied by writing the fiel in terms of a scalar [lux function V. such that B=VW.," This condition is satisfied by writing the field in terms of a scalar flux function $\Psi$, such that $\bmath{B}=-\nabla \Psi$."
+" Thus. in order to ensure that the field. is divergence-free (V.D— (0). V must satisfy Laplace's equation. Vow=0: the solution of which is a linear combination of spherical harmonics. where £P, denote the associated Legendre functions."," Thus, in order to ensure that the field is divergence-free $\nabla.\bmath{B}=0$ ), $\Psi$ must satisfy Laplace's equation, $\nabla^2\Psi=0$; the solution of which is a linear combination of spherical harmonics, where $P_{lm}$ denote the associated Legendre functions."
+" The coellicients er, and 5/5, are determined from the radial field at the stellar surface obtained. from. Zeeman-Doppler maps and also by assuming that at some height 2. above the surface (known as the source surface) the field becomes radial and hence By(fo)=0. emulating the ellect of the corona blowing open field lines to form a stellar wind. (?).."," The coefficients $a_{lm}$ and $b_{lm}$ are determined from the radial field at the stellar surface obtained from Zeeman-Doppler maps and also by assuming that at some height $R_s$ above the surface (known as the source surface) the field becomes radial and hence $B_{\theta}(R_s)=0$, emulating the effect of the corona blowing open field lines to form a stellar wind \citep{alt69}."
+ In order to extrapolate the field we used a mocified: version ofa code originally developed by. 2.., In order to extrapolate the field we used a modified version of a code originally developed by \citet*{van98}.
+ lor à given surface magnetogram we calculate the extent of the closed corona for a specified: κο of stellar parsumeters., For a given surface magnetogram we calculate the extent of the closed corona for a specified set of stellar parameters.
+ As this process has been described in detail bv 2. and. we provide only a brief outline here.," As this process has been described in detail by \citet{jar06} and \citet{gre06a,gre06b} we provide only a brief outline here."
+ We assume that the corona is isothermal and that plasma along field line loops is in hydrostatic equilibrium., We assume that the corona is isothermal and that plasma along field line loops is in hydrostatic equilibrium.
+ The pressure is calculated along the path of field. Lines loops ancl is set to zero for open Ποιά lines ancl for field. lines where. at some point along the loop. the gas pressure exceeds the magnetic pressure.," The pressure is calculated along the path of field lines loops and is set to zero for open field lines and for field lines where, at some point along the loop, the gas pressure exceeds the magnetic pressure."
+ The pressure along a field line scales with the pressure at its foot. point. and we assume that this scales with the magnetic pressure (py=EBG).," The pressure along a field line scales with the pressure at its foot point, and we assume that this scales with the magnetic pressure $p_0=KB_0^2$ )."
+ This technique has been used. successfully to caleulate mean coronal densities and X-ray EMSs for the Sun and other main sequence stars as well as T Tauri stars (7).., This technique has been used successfully to calculate mean coronal densities and X-ray EMs for the Sun and other main sequence stars \citep{jar02} as well as T Tauri stars \citep{jar06}.
+ Phe extent of the corona depends both on the value of dy and also on Ly whieh is determined. directly [rom surface magnetograms ancl varies across the stellar surface., The extent of the corona depends both on the value of $K$ and also on $B_0$ which is determined directly from surface magnetograms and varies across the stellar surface.
+ ? determined the best value of A or à given surface magnetogram which results in the best fit o the X-ray EMs of stars from the COUP database., \citet{jar06} determined the best value of $K$ for a given surface magnetogram which results in the best fit to the X-ray EMs of stars from the COUP database.
+ We have aciopted the same values in this letter., We have adopted the same values in this letter.
+ We note that we make a conservative estimate of the location of the source surface w calculating the largest radial distance at which a dipole ield Line would remain closed. with the same average field strength.," We note that we make a conservative estimate of the location of the source surface by calculating the largest radial distance at which a dipole field line would remain closed, with the same average field strength."
+ The AD Dor-like coronal field has an X-ray of og[M=53.73em (without considering accretion) and a mean EM-weighted coronal density of log»=10.57em consistent with estimates from the modelling of individual lares (?)..," The AB Dor-like coronal field has an X-ray of $\log{EM}=53.73\,{\rm cm}^{-3}$ (without considering accretion) and a mean EM-weighted coronal density of $\log{\bar{n}}=10.57\,{\rm cm}^{-3}$, consistent with estimates from the modelling of individual flares \citep{fav05}."
+ The LO lLlya-like field has a more extended corona and consequently a lower coronal density and. EM. ogEAM=52.6lem.logn—9.59cm7.," The LQ Hya-like field has a more extended corona and consequently a lower coronal density and EM, $\log{EM}=52.61\,{\rm cm}^{-3}, \log{\bar{n}}=9.79\,{\rm cm}^{-3}$."
+ We consider a thin disc anc assume that the disc norma is aligned. with the stellar rotation axis., We consider a thin disc and assume that the disc normal is aligned with the stellar rotation axis.
+ We consider a steady state Llow model and assume that the structure of the magnetic field remains undistorted by the in-falling materia and that the magnetosphere rotates as a solic body., We consider a steady state flow model and assume that the structure of the magnetic field remains undistorted by the in-falling material and that the magnetosphere rotates as a solid body.
+ ? and ? demonstrated that for most stars the closed. Lele region is confined. close to the stellar surface., \citet{jar06} and \citet{gre06a} demonstrated that for most stars the closed field region is confined close to the stellar surface.
+ T Tauri stars have hot coronae. with temperatures in excess of severa tens of mega-Ixclvin. e.g. see the COUP database. ?]].," T Tauri stars have hot coronae, with temperatures in excess of several tens of mega-Kelvin [e.g. see the COUP database, \citet{get05}] ]."
+ Such high temperatures leack to closed. field. lines being openec by the gas pressure. limiting the closed field region to close to the stellar surface.," Such high temperatures lead to closed field lines being opened by the gas pressure, limiting the closed field region to close to the stellar surface."
+ In such cases the inner disc may si in a reservoir of open field. (?).., In such cases the inner disc may sit in a reservoir of open field \citep{gre06a}.
+ This can be seen in Fig., This can be seen in Fig.
+ 1 where the field lines that are carrving accreting gas are open., \ref{coronae} where the field lines that are carrying accreting gas are open.
+ The question of where the disc is truncated. remains a major problem for aceretion models., The question of where the disc is truncated remains a major problem for accretion models.
+ It is still unknown if 10 disc is truncated in the vicinity of the corotation radius. jo assumption of traditional accretion models (e.g ?T)) or whether it extends closer to the stellar surface (e.g. 2)).," It is still unknown if the disc is truncated in the vicinity of the corotation radius, the assumption of traditional accretion models (e.g. \citealt{kon91}) ), or whether it extends closer to the stellar surface (e.g. \citealt{mat05}) )."