diff --git "a/batch_s000045.csv" "b/batch_s000045.csv" new file mode 100644--- /dev/null +++ "b/batch_s000045.csv" @@ -0,0 +1,10385 @@ +source,target + The best-fitting quadratie ephemeris calculated is: and is shown in Fig. 3.., The best-fitting quadratic ephemeris calculated is: and is shown in Fig. \ref{fig:ephem}. + Xdding a cubic term had little ellect on the quality of the fit and is thus not done., Adding a cubic term had little effect on the quality of the fit and is thus not done. + Note that this gives a slightly longer period than Mason (1997). and a smaller value for the period change.," Note that this gives a slightly longer period than Mason (1997), and a smaller value for the period change." + We extractedsource. anc background. spectra. from. the, We extractedsource and background spectra from the +Consequently. the flux associated to the measured photometric quantity can be written as the dot product of the original flux distribution. /(À;) and a weighting function. so that each column of the sensing matrix ® in Eq. (1)),"Consequently, the flux associated to the measured photometric quantity can be written as the dot product of the original flux distribution $f(\lambda_i)$ and a weighting function, so that each column of the sensing matrix $\mathbf{\Phi}$ in Eq. \ref{eq:sensing}) )" + is given bv: The sparsity constraint of (he spectrum is fulfilled automatically when using a principal component decomposition. with the additional advantage of knowing exactly which coefficients ol (he sparse vector a are non-zero.," is given by: The sparsity constraint of the spectrum is fulfilled automatically when using a principal component decomposition, with the additional advantage of knowing exactly which coefficients of the sparse vector $\mathbf{a}$ are non-zero." + Therefore. the solution of the problem given by Eq. (3))," Therefore, the solution of the problem given by Eq. \ref{eq:sensing2}) )" + is simpler than the full CS problem in which the non-zero elements of a have to be identified., is simpler than the full CS problem in which the non-zero elements of $\mathbf{a}$ have to be identified. + Consequently. the solution to Eq. (3))," Consequently, the solution to Eq. \ref{eq:sensing2}) )" + is given by the sparse vector that minimizes the following f5-norm: In other words. welook for the sparse vector a with the first AV elements different [rom zero aud the rest set to zero that minimizes (he square difference between (he photonmetric flix on the g. r aud / fillers aud the ones reconstructed using (hie previous formalism. where the [ας is obtained as a linear combination of ἐν principal components.," is given by the sparse vector that minimizes the following $\ell_2$ -norm: In other words, welook for the sparse vector $\mathbf{a}$ with the first $K$ elements different from zero and the rest set to zero that minimizes the square difference between the photometric flux on the $g$, $r$ and $i$ filters and the ones reconstructed using the previous formalism, where the flux is obtained as a linear combination of $K$ principal components." + We now develop in more detail the steps to be followed., We now develop in more detail the steps to be followed. + Assume that the signal of interest. F(A) ean be written as a linear combination of A (sparsity) PCA basis functions D;(À). so that: The measurement process produces the following linear combinations: where Af is (he number of wavelength points. NV is the number of measurements ancl the ®;; ave the matrix elements of the sensing matrix 9. given in Eq. (3)).," Assume that the signal of interest $F(\lambda)$ can be written as a linear combination of $K$ (sparsity) PCA basis functions $B_i(\lambda)$, so that: The measurement process produces the following linear combinations: where $M$ is the number of wavelength points, $N$ is the number of measurements and the $\Phi_{ij}$ are the matrix elements of the sensing matrix $\mathbf{\Phi}$ , given in Eq. \ref{eq:sensing_matrix_final}) )." + Plugging Eq. (10)), Plugging Eq. \ref{eq:develop}) ) + into Eq. (11)).," into Eq. \ref{eq:sensing_linear_combination}) )," + we end up with:, we end up with: +for some function q;(r). then the cross power spectrum. will be Η we wish to caleulate the cross power of & between two redshifts. then we see by comparing equations (10)) and (15)) that we do have two projections of the necessary form. with q; given by so we therefore obtain (noting that. the two-point statistics of α and agree. c.g. Blandford et al 19911. where (is the angular wavenumber.,"for some function $q_i(r)$, then the cross power spectrum will be If we wish to calculate the cross power of $\kappa$ between two redshifts, then we see by comparing equations \ref{eq:kappa}) ) and \ref{eq:pi}) ) that we do have two projections of the necessary form, with $q_i$ given by so we therefore obtain (noting that the two-point statistics of $\kappa$ and $\gamma$ agree, e.g. Blandford et al 1991), where $\ell$ is the angular wavenumber." + In the case where we wish to examine the power of the shear at one particular redshift. we can simplify this to We can obtain corresponding cross-correlation functions for shears at two clilferent redshifts 24. 22.," In the case where we wish to examine the power of the shear at one particular redshift, we can simplify this to We can obtain corresponding cross-correlation functions for shears at two different redshifts $z_1$ , $z_2$." +" We will use three diferent correlation functions: C, represents (4-7). where ut represents the first shear component of two galaxies. in a coordinate frame where zero position angle lies along the line joining the galaxies (e.g. Bacon et al 2003)."," We will use three different correlation functions: $C_1$ represents $\langle \gamma_1^a \gamma_1^b +\rangle$, where $\gamma_1^{a,b}$ represents the first shear component of two galaxies, in a coordinate frame where zero position angle lies along the line joining the galaxies (e.g. Bacon et al 2003)." + C2 represents (5555‘1 sand C=Cy|€».," $C_2$ represents $\langle \gamma_2^a \gamma_2^b \rangle$, and $C=C_1+C_2$." + Given these definitions. the correlation functions are simple transforms of the power: These are the 3-D. shear correlation functions we have been seeking: we will use these functions to compare our data with various evolving cosmological models.," Given these definitions, the correlation functions are simple transforms of the power: These are the 3-D shear correlation functions we have been seeking; we will use these functions to compare our data with various evolving cosmological models." + ]t is convenient to have a simple means of calculating the shear power spectrum [rom the matter power spectrum: this has been described in the last section., It is convenient to have a simple means of calculating the shear power spectrum from the matter power spectrum; this has been described in the last section. + However. it is also useful to have a means of calculating the matter power spectrum given the shear power spectrum: here we outline how this can be achieved for theoretical models.," However, it is also useful to have a means of calculating the matter power spectrum given the shear power spectrum; here we outline how this can be achieved for theoretical models." + Noisy data ave difficult to invert with this procedure: it would be more convenient to fit models to the data and then use the results of this section to invert these models., Noisy data are difficult to invert with this procedure; it would be more convenient to fit models to the data and then use the results of this section to invert these models. + ln order to show how to invert the shear power spectrum. we start by considering an integral of the form where D(r.r)' is. à smooth. continuous. function.," In order to show how to invert the shear power spectrum, we start by considering an integral of the form where $B(r',r)$ is a smooth, continuous function." +⋅. Carrving out a partial dillerentiation of ο with respect tor. we find This result will be used below in order to untangle the matter power spectrum from its integral projection found in calculating the shear.," Carrying out a partial differentiation of $A$ with respect to $r$, we find This result will be used below in order to untangle the matter power spectrum from its integral projection found in calculating the shear." + We also require the result which can be easily verified directly from equation (12))., We also require the result which can be easily verified directly from equation \ref{eq:fk}) ). + Using these two results repeatedly. upon the shear power spectrum (for). we find that we can calculate the matter power spectrum: Alternatively. we can find a similar means of calculating the matter power spectrum from the cross power spectrum. of shear at two clilferent redshifts: These equations therefore permit us to find the matter power spectrum given a model for the shear power spectrum.," Using these two results repeatedly upon the shear power spectrum $P_{\gamma}(l,r)$, we find that we can calculate the matter power spectrum: Alternatively, we can find a similar means of calculating the matter power spectrum from the cross power spectrum of shear at two different redshifts: These equations therefore permit us to find the matter power spectrum given a model for the shear power spectrum." + Note that upon noisy data. the first and second dilferentials in this equation can lead to unphysical negative power spectra.," Note that upon noisy data, the first and second differentials in this equation can lead to unphysical negative power spectra." + Thus we reiterate that a better approach is to fit a continuous shear model to the data which can then be inverted with these equations., Thus we reiterate that a better approach is to fit a continuous shear model to the data which can then be inverted with these equations. + We conclude this section with a discussion of models for the matter and shear power spectra which allow us to directly examine the redshift evolution of the power spectrum., We conclude this section with a discussion of models for the matter and shear power spectra which allow us to directly examine the redshift evolution of the power spectrum. + Firstly. suppose we have measured the shear cross-correlation function C' between many recshifts. from which we can calculate the shear power spectrum.," Firstly, suppose we have measured the shear cross-correlation function $C$ between many redshifts, from which we can calculate the shear power spectrum." +" We will initially restrict ourselves to power within a redshift shell. 2.2 we will consider £7,» later."," We will initially restrict ourselves to power within a redshift shell, $P_\gamma$; we will consider $P_{\gamma 1 2}$ later." + We will attempt to use our formalism with a simple model for the shear power spectrum: a power law in both theangular and redshift directions. i.c. The correlation function corresponding to this moclel. calculated from equation (20)) is:," We will attempt to use our formalism with a simple model for the shear power spectrum: a power law in both theangular and redshift directions, i.e. The correlation function corresponding to this model, calculated from equation \ref{ctheta}) ) is:" +Viable optical counterparts have been suggested. for ouly five among the26 sources shown iu roffxlun.. all of them probable cataclesiuic variables.,"Viable optical counterparts have been suggested for only five among the 26 sources shown in \\ref{fxlum}, all of them probable cataclysmic variables." + We compare the ratio of X-ray flux to optical flux of these sources with the ratios measured for cataclysuic variables and for RS CVu systems in the Galactic Disk in reffeompa.., We compare the ratio of X-ray flux to optical flux of these sources with the ratios measured for cataclysmic variables and for RS CVn systems in the Galactic Disk in \\ref{fcompa}. + It is seen that the sueeested optical counterparts for the sources ino 66397. and 66752 lead to ratios which are compatible with those of cataclvsiic variables. whereas those i [7 Tuc are too bright in N-ravs. in aerecineut with reffxhun..," It is seen that the suggested optical counterparts for the sources in 6397 and 6752 lead to ratios which are compatible with those of cataclysmic variables, whereas those in 47 Tuc are too bright in X-rays, in agreement with \\ref{fxlum}." + Tf these sources are indeed catachvsndc variables. their excessive N-rayv luminosity needs to be explained: alternatively. the suggestedOO identifications may be chance coiucideuces (as discussed by Verbuut Wasinger 1998).," If these sources are indeed cataclysmic variables, their excessive X-ray luminosity needs to be explained; alternatively, the suggested identifications may be chance coincidences (as discussed by Verbunt Hasinger 1998)." + All sugeested counterparts lead to hnieher X-ray o optical flux ratios than those of RS CVu binaries., All suggested counterparts lead to higher X-ray to optical flux ratios than those of RS CVn binaries. + The accurate positions that we determine for individual sources are valid for separately detected sources in particular., The accurate positions that we determine for individual sources are valid for separately detected sources in particular. + Iu the case of overlapping sources. we do not lave unique solutious.," In the case of overlapping sources, we do not have unique solutions." + Thus. iu the core of 66397 fits with 5 and 6 sources are both acceptable. at simular quality: aud we cannot exclude that more sources contribute to the observed flix. which would invalidate our derived positions.," Thus, in the core of 6397 fits with 5 and 6 sources are both acceptable, at similar quality; and we cannot exclude that more sources contribute to the observed flux, which would invalidate our derived positions." + Binaries mav reside away from the core either because he cluster has uudergouc little mass segresation. or )ocaise a three-body interaction aa close encounter Pa binary with a single star) in the core has expelled the inuv from the core HITut et 11992).," Binaries may reside away from the core either because the cluster has undergone little mass segregation, or because a three-body interaction a close encounter of a binary with a single star) in the core has expelled the binary from the core Hut et 1992)." + In the latter case the binary is expected to be ecceutric nunuediatelv after beiug expelled: tidal forces iav in time circularize he orbit again., In the latter case the binary is expected to be eccentric immediately after being expelled; tidal forces may in time circularize the orbit again. + Such binaries are only a minoritv of the overall binary population of a cluster: however. Nouv observations may preforablv select sucht binaries if tical orces act in them.," Such binaries are only a minority of the overall binary population of a cluster; however, X-ray observations may preferably select such binaries if tidal forces act in them." + Since sources away frou the core can be ore- or background sources. optical identification of thoi is required to settle whether they belong to the cluster or rot.," Since sources away from the core can be fore- or background sources, optical identification of them is required to settle whether they belong to the cluster or not." + Our accurate positions should help iu finding such counterparts., Our accurate positions should help in finding such counterparts. +"ποιο! of 2.8«10τοσα” by suggesting a breaking strain of 10.2,","moment of $2.8\times +10^{41}$ $\cdot $ $^2$ by suggesting a breaking strain of $10^{-2}$." + Our estimation is one order ligher than hat of Owen(2005). and shows that the solidified nountains can provide a bigger quadrupole moment thaw clastic mountains do.," Our estimation is one order higher than that of \cite{Owen2005}, and shows that the solidified mountains can provide a bigger quadrupole moment than elastic mountains do." + Horowitzetal.(2009) have done sole work to show that the breakiug straiu on normal jeutron stars can be as big as LO+., \cite{Horowitz2009} have done some work to show that the breaking strain on normal neutron stars can be as big as $10^{-1}$. + Our result about the Παπ quadrupole moment on the solid quark stars would agree with Owen's if the breaking strain on solid quark stars can also be as high as 10.+., Our result about the maximum quadrupole moment on the solid quark stars would agree with Owen's if the breaking strain on solid quark stars can also be as high as $10^{-1}$. + Tt is evideut that th the dependence on radius aud mass are the sanie iu his work auc that of Owen(2005)., It is evident that both the dependence on radius and mass are the same in this work and that of \cite{Owen2005}. +.. What can we coustrain the equation of state ly he CW observations?, What can we constrain the equation of state by the GW observations? + To compare the theoretical quadrupole moment with the LICO 85 data (Abbottetal. 2010).. we are to provide the relationship between Qo and CAV amplitude fy.," To compare the theoretical quadrupole moment with the LIGO S5 data \citep{Abbott2010}, , we are to provide the relationship between $Q_{22}$ and GW amplitude $h_0$." + Assuniug a density perturbation of 6p=Ro|pz23o2(0. (Ushomirskyetal.2000) aud following the definition 09]of Qo» above. oue comes to where a factor 1/2 before the amplitude of the trace-reversed perturbation is added for fy estimation.," Assuming a density perturbation of $\delta \rho={\rm Re}[\rho_{22}Y_{22}(\theta,\phi)]$ \citep{Ushomirsky2000} and following the definition of $Q_{22}$ above, one comes to where a factor $1/\sqrt{2}$ before the amplitude of the trace-reversed perturbation is added for $h_0$ estimation." + Similar reatiments can also be found in other eravitational wave iteratures (CJaranowskietal.1998)., Similar treatments can also be found in other gravitational wave literatures \citep{Jaranowski1998}. +. The ouly difference is applviug ellipticity € rather than quadrupole moment Qo»., The only difference is applying ellipticity $\epsilon$ rather than quadrupole moment $Q_{22}$. + Note the relationship e=VarLiQoofT... it is casy o prove that they are equivalent.," Note the relationship $\epsilon=\sqrt{8\pi/15}Q_{22}/I_{zz}$, it is easy to prove that they are equivalent." + With our estimation of the maxinuuun quadrupole uonicnt into the equation above. we can obtain the uaxinumn amplitude of the exavitational waves. LIGO i focusing ou the direct detection of eravitational waves.," With our estimation of the maximum quadrupole moment into the equation above, we can obtain the maximum amplitude of the gravitational waves, LIGO is focusing on the direct detection of gravitational waves." + Recent LIGO $5 data has given an upper bound of the fy for 116 known pulsars (Abbottetal.2010)., Recent LIGO S5 data has given an upper bound of the $h_0$ for 116 known pulsars \citep{Abbott2010}. +. The value of the upper bound of the 116 pulsus varies from 107° to 1077., The value of the upper bound of the 116 pulsars varies from $10^{-26}$ to $10^{-25}$. + However. from Eq.(5)). it sees that solid quark stars can have mountains high enough to radiate gravitational waves with amplitudes as high as 10.77.," However, from \ref{h0esti}) ), it seems that solid quark stars can have mountains high enough to radiate gravitational waves with amplitudes as high as $10^{-22}$." + Does this mean that these 116 pulsars cannot be solid quark stars?, Does this mean that these 116 pulsars cannot be solid quark stars? + Should we believe that the mountains ou solid quar stars have the imaxiumnun heights?, Should we believe that the mountains on solid quark stars have the maximum heights? + Iu addition. iu our estimation of Eq.(5)). we assune a distribution of 35». which cau attribute most to the exavitational waves.," In addition, in our estimation of \ref{h0esti}) ), we assume a distribution of $Y_{22}$, which can attribute most to the gravitational waves." + It is surcly quite strange if natural mountains on stars have he exact distribution of 359 im order to produce the ΠαπΙΙ eravitational waves., It is surely quite strange if natural mountains on stars have the exact distribution of $Y_{22}$ in order to produce the maximum gravitational waves. + The orogeny iu the Earth's crust is powered by the uautle convection and thus the engagement of tectonic ates., The orogeny in the Earth's crust is powered by the mantle convection and thus the engagement of tectonic plates. + What could be the force to build mountaius onu solid quark stars?, What could be the force to build mountains on solid quark stars? + Elastic energy develops when a star evolves. and both bulk-iuvariable aud bulk-variable Orces can result in decreases of moment of inertia durimg a star quake (Penge&Xu2008).," Elastic energy develops when a star evolves, and both bulk-invariable and bulk-variable forces can result in decreases of moment of inertia during a star quake \citep{Peng2008}." + This force would be responsible for mountain building too., This force would be responsible for mountain building too. + We cau then estimate the real heights of mountains from the elitch phenomena of pulsars below., We can then estimate the real heights of mountains from the glitch phenomena of pulsars below. + We think elitehes occur if the moment of inertia change suddenly inside a solid quark stars (Zhouetal.2001)., We think glitches occur if the moment of inertia change suddenly inside a solid quark stars \citep{Zhou2004}. +". The pulsus aneular momentum, πιοον the assmuption of sphere. is 3ALR?O/5."," The pulsars' angular momentum, under the assumption of sphere, is $\sim 3MR^2\Omega/5$ ." +" Making a differential of it. we have |RI~5«10""[8(Q/Q)/10""|."," Making a differential of it, we have $\left|\delta R/R\right|\sim +5\times 10^{-7}[\delta(\Omega/\Omega)/10^{-6}]$ ." + Although this consideration could not be the real heights of mountains on pulsars. we think that actually the height of mountains might be of the same order of 62.," Although this consideration could not be the real heights of mountains on pulsars, we think that actually the height of mountains might be of the same order of $\delta R$." + Therefore. replacing 6R with {μι aud inserting this iuto the estimation of Eq.(1)). we have the estimation of amplitude /y below. We draw Fie.," Therefore, replacing $\delta R$ with $H_{\rm +m}$ and inserting this into the estimation of \ref{Q22h0}) ), we have the estimation of amplitude $h_0$ below, We draw Fig." + l to show how fy varies as the mass of solid quark stus changes iu the model of Lai&Xu(2009).. for different elitch amplitude 00/0.," 1 to show how $h_0$ varies as the mass of solid quark stars changes in the model of \cite{Lai2009}, for different glitch amplitude $\delta\Omega/\Omega$." + In the calculations. we choose the numbers of quarks in a quar clusters to be 18 aud the potential Uy to be 50 MeV. We can see that if our estimation of the real height of mountains is valid. it would be natural that LICO still hasn't detect the eravitational waves directly. since the miaxiuuni CAV amplitude preseuted iu Eq.(5)) requires maxim height and a particular Y55-distribution of mountains.," In the calculations, we choose the numbers of quarks in a quark clusters to be $18$ and the potential $U_0$ to be $50$ MeV. We can see that if our estimation of the real height of mountains is valid, it would be natural that LIGO still hasn't detect the gravitational waves directly, since the maximum GW amplitude presented in \ref{h0esti}) ) requires maximum height and a particular $Y_{22}$ -distribution of mountains." + What if the mountains havent the Y55-distribution?, What if the mountains haven't the $Y_{22}$ -distribution? + This will be discussed im the next section., This will be discussed in the next section. + Above we give an approach to the actual amplitude of CAVS using the elitch plenomenon., Above we give an approach to the actual amplitude of GWs using the glitch phenomenon. + However. the auplitude of GAVs from quark stars depends also ou another factor.," However, the amplitude of GWs from quark stars depends also on another factor." + Besides the maxinnun height of mountains. the distribution of mountains plavs a kev role as well.," Besides the maximum height of mountains, the distribution of mountains plays a key role as well." + Iu the previous section. we eive an estimation of the iiaxiumia amplitude of CAs by assinuing a specific distribution described in Eq.," In the previous section, we give an estimation of the maximum amplitude of GWs by assuming a specific distribution described in Eq." +(2).. Such a hypothesis is too strong to approach the physical circtiustances., Such a hypothesis is too strong to approach the physical circumstances. + Iu this section we will discard this assumption aud consider another Likely distribution., In this section we will discard this assumption and consider another likely distribution. + Tt is certainly very cdifüeult to kuow the real distribution of imountaius on a solid quark star., It is certainly very difficult to know the real distribution of mountains on a solid quark star. + Nevertheless. an idea comes out that we can give a random distribution to approach the actual situation.," Nevertheless, an idea comes out that we can give a random distribution to approach the actual situation." + Yet another problem of the definition of random distribution cmerees., Yet another problem of the definition of random distribution emerges. + The imost strict and plysical definition should satisfv the following requirements. (, The most strict and physical definition should satisfy the following requirements. ( +1) The height of mountains should varies frou 0 to the miaxiuun height of mountains. (,1) The height of mountains should varies from 0 to the maximum height of mountains. ( +2) The surface of the star should be coutimmous. sav the fiction 0.ω) which describe the distribution of iiouutaius should be infinitely differentiable. (,"2) The surface of the star should be continuous, say the function $H(\theta,\phi)$ which describe the distribution of mountains should be infinitely differentiable. (" +3) Thethird requireimieut comes uot from imathematics. but from plivsics: the partial derivative of /7(0.0)should not be too high. otherwise a 1nountain which is quite precipitous would tend to fall down.,"3) Thethird requirement comes not from mathematics, but from physics: the partial derivative of $H(\theta,\phi)$should not be too high, otherwise a mountain which is quite precipitous would tend to fall down." + With all the considerations above. we find it is nupossible or at least ναν hard to use the most," With all the considerations above, we find it is impossible or at least very hard to use the most" +Ix8.,K8. + All of tjejr Baliner lines are in enilssiol along with the Ca H aud Ix lines., All of their Balmer lines are in emission along with the Ca H and K lines. + We detect Li I AGTOS line iu the spectrum of idl659 even at OUL ow resolution., We detect Li I $\lambda$ 6708 line in the spectrum of id1659 even at our low resolution. + There is oue M star in our sample (k113) with weak Ha aud [NU] emission., There is one M star in our sample (id113) with weak $H\alpha$ and [NII] emission. + To examine the pessljlitw that KlE13 is a dMe star. we estimate the number of nearby «Me stars from the observed stellar density of 0.081 star/pe% (Henryeal.2002).," To examine the possibility that id113 is a dMe star, we estimate the number of nearby dMe stars from the observed stellar density of 0.084 $\rm pc^{-3}$ \citep{henr02}." +. Roughly 70% of nearby stars are M dwarls (Her'οἱal.2002)., Roughly $70\%$ of nearby stars are M dwarfs \citep{henr02}. +. At our completeness imit (V.c11.9). we cau detect M cdwarfs at clistances of 100 ye.," At our completeness limit $V \simeq 14.9$ ), we can detect M dwarfs at distances of 100 pc." + The number of detectable M dwars towards NGC 6871 is two., The number of detectable M dwarfs towards NGC 6871 is two. + The number of «Me stars alno& M cdwarfs increases monotonically with spectral type (1% around MO and 9056 M5 and later (Joy&Abt 1971)))., The number of dMe stars among M dwarfs increases monotonically with spectral type $4\%$ around M0 and $90\%$ M5 and later \citep{joy74}) ). + The overall ratio of dle stars among M cdwarfs is 175€. whieh yields a imocest probability (~ 31%) of detecting oue in our spectroscopic sample.," The overall ratio of dMe stars among M dwarfs is $17\%$, which yields a modest probability $\simeq 34\%$ ) of detecting one in our spectroscopic sample." + Thus. id113 could be a dMe star.," Thus, id113 could be a dMe star." + High resolution optical spectra of Li E AGTOS would cliscriminate between a field «Me star and a pre-main sequence star., High resolution optical spectra of Li I $\lambda$ 6708 would discriminate between a field dMe star and a pre-main sequence star. + There are 15 emission line stars iu this eeroup., There are 15 emission line stars in this group. + Their spectral tvpes raneeOm from early B to late G. Most are definite emission liue objects., Their spectral types range from early B to late G. Most are definite emission line objects. + More than half of the 15 cluster members (9 altogether) appear to be PMS stars., More than half of the 15 cluster members (9 altogether) appear to be PMS stars. + For 8 stars. we base this ideutification ou their position on the HRD.," For 8 stars, we base this identification on their position on the HRD." + Oue star has [NID] ancl [SII] emission., One star has [NII] and [SII] emission. + The‘e are three B type stars (dS. 14156. i138) amoung the 11 definite emission liue stars.," There are three B type stars (id8, id186, id38) among the 11 definite emission line stars." + They show emission ouly in Πα., They show emission only in $H\alpha$. + Two of them (145 aud id38) were previously known Polcaro 2001).," Two of them (id8 and id38) were previously known \citep[see][]{grig88,bern01}." +.. They could be cassical Be stars., They could be classical Be stars. + The‘e ape two A stars (10492100. id1691) in this subgroup.," There are two A stars (id2139, id1694) in this subgroup." + One (id2139) shows sigus of [NII] and [SII] emission., One (id2139) shows signs of [NII] and [SII] emission. + The Ha emission of id160[ is very weak. but its position on the HRD iudicates that it may be a PMS star.," The $H\alpha$ emission of id1694 is very weak, but its position on the HRD indicates that it may be a PMS star." + There are three F type stars witl clelinite eiilssion featwes among the cluster members., There are three F type stars with definite emission features among the cluster members. + Two (i2616. 12631) are PAIS caudidates with [NI aud [SII] emission.," Two (id2646, id2631) are PMS candidates with [NII] and [SII] emission." + The third one (id12055) shows emission in Ho and lies close to the 1ail secuelce., The third one (id2055) shows emission in $H\alpha$ and lies close to the main sequence. + Three G type members (12621. i¢12138. id26το) may be PMS stars.," Three G type members (id2621, id2138, id2675) may be PMS stars." + All lie well above the main sequence., All lie well above the main sequence. + The spectrum of id2621 coiains [NIL and [SII] emission: we detect Li I AGTOS line in the spectrum of id2138., The spectrum of id2621 contains [NII] and [SII] emission; we detect Li I $\lambda$ 6708 line in the spectrum of id2138. + The third star of this ο'oup does not show any. particular eature beside the weak emission in the core of the Ha line., The third star of this group does not show any particular feature beside the weak emission in the core of the $H\alpha$ line. + Based ou their reddening.Om we identify 16 emission line stars as possible backgrouudfon) objects.," Based on their reddening, we identify 16 emission line stars as possible background objects." +"mode, the LAT observes the entire sky every 3 hours.","mode, the LAT observes the entire sky every 3 hours." +" For individual sources, pprovides nearly uniform sky coverage down to a photon flux of 4x1071? cm""? s! between 1 and 100 GeV, except for sources at low Galactic latitude (|b|€ 10?) where the diffuse emission dominates (Abdoetal.2010a)."," For individual sources, provides nearly uniform sky coverage down to a photon flux of $4 \times 10^{-10}$ $^{-2}$ $^{-1}$ between 1 and 100 GeV, except for sources at low Galactic latitude $|b|\leq 10^\circ$ ) where the diffuse emission dominates \citep{abdo}." +". Throughout the paper, we assume an Ho —71 km s! Mpc!, Qm=0.27, Qa=0.73 cosmological model."," Throughout the paper, we assume an $H_0 = $ 71 km $^{-1}$ $^{-1}$, $\Omega_m = 0.27$, $\Omega_{\Lambda} = 0.73$ cosmological model." +" As of 2007 November, the public catalog of UHECRs released by the collaboration contains 27 events with energiesabove 55 EeV collected at their site in Malargüee, Argentina (Abrahametal.2004)."," As of 2007 November, the public catalog of UHECRs released by the collaboration contains 27 events with energiesabove 55 EeV collected at their site in Malargüee, Argentina \citep{abraham}." +. The LLAT 1FGL catalog consists of 1451 sources characterized in the 100 MeV—100 GeV energy range (Abdoetal.2010a)., The LAT 1FGL catalog consists of 1451 sources characterized in the 100 MeV–100 GeV energy range \citep{abdo}. +. The data were obtained in an all-sky scanning mode during 2008 August-2009 July and represents the most extensive map of the gamma-ray sky (>100 MeV) ever obtained., The data were obtained in an all-sky scanning mode during 2008 August–2009 July and represents the most extensive map of the gamma-ray sky $\geq 100$ MeV) ever obtained. +" The entire catalog includes 689 blazars, two starburst galaxies, two radio galaxies, 56 pulsars, 50 supernova remnants, and 630 unidentified sources (Abdoetal.2010a)."," The entire catalog includes 689 blazars, two starburst galaxies, two radio galaxies, 56 pulsars, 50 supernova remnants, and 630 unidentified sources \citep{abdo}." +". Since we are interested in testing for possible correlations of UHECRs with various classes of gamma-ray emitters in the MeV-GeV energy range without anypriori assumption, we take advantage of the complete 1FGL catalog without any redshift or type discrimination."," Since we are interested in testing for possible correlations of UHECRs with various classes of gamma-ray emitters in the MeV-GeV energy range without any assumption, we take advantage of the complete 1FGL catalog without any redshift or type discrimination." +" To test for a possible cross-correlation between UHECRs and ssources in the 1FGL catalog, we checked whether the individual UHECR positions in the sample are clearly contained within the error circle of individual sources listed in the 1FGL catalog."," To test for a possible cross-correlation between UHECRs and sources in the 1FGL catalog, we checked whether the individual UHECR positions in the sample are clearly contained within the error circle of individual sources listed in the 1FGL catalog." +" Specifically, we count correlations whenever an UHECR event is within a circle of 3.P radius around a particular ssource."," Specifically, we count correlations whenever an UHECR event is within a circle of $3.1\!^\circ$ radius around a particular source." + This radius is in line with the value predicted by conventional models of cosmic ray trajectories that consider the full effect of the Galactic magnetic field (Abrahametal.2008a).., This radius is in line with the value predicted by conventional models of cosmic ray trajectories that consider the full effect of the Galactic magnetic field \citep{abraham5}. +" In order to avoid “multiple” counts, we only consider one match per individual UHECR event."," In order to avoid “multiple” counts, we only consider one match per individual UHECR event." +" Within the 27 UHECR events, we find 12 matches with 1FGL sources."," Within the 27 UHECR events, we find 12 matches with 1FGL sources." +" To examine the likelihood of such correlation, we used the BATSE 4B gamma-ray burst (GRB) catalog (Paciesasetal.1999) in place of a random generator of isotropic sky positions that allows us to generate artificial samples of simulated UHECRs."," To examine the likelihood of such correlation, we used the BATSE 4B gamma-ray burst (GRB) catalog \citep{paciesas} in place of a random generator of isotropic sky positions that allows us to generate artificial samples of simulated UHECRs." + The BATSE 4B catalog consists of 1637 GRB positions localized by the BATSE instrument on board the Compton Gamma Ray Observatory (CGRO) in the period between 1991 April 19 and 1996 August 29., The BATSE 4B catalog consists of 1637 GRB positions localized by the BATSE instrument on board the Compton Gamma Ray Observatory (CGRO) in the period between 1991 April 19 and 1996 August 29. +" For our work, we restricted our analysis to events that would be accessible from the southern site at a declination < 24.8 (Abrahametal.2004)."," For our work, we restricted our analysis to events that would be accessible from the southern site at a declination $<$ $24.8\!^\circ$ \citep{abraham}." +". In addition, we assume that the fraction of exposure is the same accross the declination range as stated in Abrahametal.(2008a)."," In addition, we assume that the fraction of exposure is the same accross the declination range as stated in \citet{abraham5}." +". After the proper cuts were applied, we drew 100 random sets of 27 events from the resulting “southern” BATSE 4B catalog to match the UHECR sample."," After the proper cuts were applied, we drew 100 random sets of 27 events from the resulting “southern” BATSE 4B catalog to match the UHECR sample." + We next proceeded to correlate each of the 100 random sets with the 1451 ssources in the same manner as with the original dataset., We next proceeded to correlate each of the 100 random sets with the 1451 sources in the same manner as with the original dataset. + Figure 1 shows the distribution of matches between the artificial samples of UHECRs constructed from the BATSE 4B catalog and the 1FGL catalog., Figure \ref{figure1} shows the distribution of matches between the artificial samples of UHECRs constructed from the BATSE 4B catalog and the 1FGL catalog. +" In particular, 63% of the artificial samples have 12 or more matches consistent with ppositions."," In particular, $\%$ of the artificial samples have 12 or more matches consistent with positions." + Allowing for larger circles (6° and 8° radii) only increases the number of false positives when compared with the BATSE sample to 71% and 73% respectively., Allowing for larger circles $6^\circ$ and $8^\circ$ radii) only increases the number of false positives when compared with the BATSE sample to $\%$ and $\%$ respectively. +" If the position of events were strongly correlated with ssources, we would expect a greater number of coincidences when compared with the outcomes of randomly-generated artificial sets of UHECR events drawn from the BATSE 4B GRB catalog."," If the position of events were strongly correlated with sources, we would expect a greater number of coincidences when compared with the outcomes of randomly-generated artificial sets of UHECR events drawn from the BATSE 4B GRB catalog." + We find no evidence for such an association and subsequently cannot claim a positive cross-correlation between UHECRs and the 1FGL catalog., We find no evidence for such an association and subsequently cannot claim a positive cross-correlation between UHECRs and the 1FGL catalog. +" In summary, we find no cross-correlation between UHECRs and 1FGL sources that cannot be reproduced by chance alignment."," In summary, we find no cross-correlation between UHECRs and 1FGL sources that cannot be reproduced by chance alignment." +" This differs from the findings reported by Abrahametal. (2009),, using a different AGN sample (see 2007).."," This differs from the findings reported by \citet{abraham3}, , using a different AGN sample \citep[see also][]{abraham2}. ." + Examining the matches between, Examining the matches between +bar causes the scale length. of the disk to increase with time as well.,bar causes the scale length of the disk to increase with time as well. + Measured. with a double exponential fit to the surface density. profile. this increase in the disk scale length is fairly rapid initially. jumping to 2.3 from 1.0 between |—50 and /=110. and slowly increasing to 2.8 hy the end of the simulation.," Measured with a double exponential fit to the surface density profile, this increase in the disk scale length is fairly rapid initially, jumping to 2.3 from 1.0 between $t=50$ and $t=110$, and slowly increasing to 2.8 by the end of the simulation." + This increase more than olfsets the pattern speed slow-clown. with only one comparison still too low at /=110 while two are higher than observed. and the rest are within the accepted ranges.," This increase more than offsets the pattern speed slow-down, with only one comparison still too low at $t=110$ while two are higher than observed, and the rest are within the accepted ranges." + By the end of the simulation three are again too low. only one is too high. and the rest are acceptable (with two at the lower limit of the error bars. anc two well within the measured. range)," By the end of the simulation three are again too low, only one is too high, and the rest are acceptable (with two at the lower limit of the error bars, and two well within the measured range)." + Therefore. for à considerable duration. these long-lived. bar is rotating at speeds. comparable to those observed.," Therefore, for a considerable duration, these long-lived bar is rotating at speeds comparable to those observed." + This reiterates the corotation to bar length ration finding above: both the simulated. and observed. bars are “fast” (ratio less than 1.5).," This reiterates the corotation to bar length ration finding above: both the simulated and observed bars are ""fast"" (ratio less than 1.5)." + Some attention has recently been placed on the identification of cclge-on barred galaxies., Some attention has recently been placed on the identification of edge-on barred galaxies. + MM. Athanassoula Bureau (1999) (ΑΦ). and BAA have studied: the. line-ol-sight. velocity. profile. of peanut-shaped. edge-on galaxies and have shown that thei kinematics are explained by the presence of a bar.," M, Athanassoula Bureau (1999) B), and A have studied the line-of-sight velocity profile of peanut-shaped, edge-on galaxies and have shown that their kinematics are explained by the presence of a bar." + The characteristic plot. of position versus velocity gives a unique figure-eight pattern for barred ealaxies. since some orbits are depleted in these systems.," The characteristic plot of position versus velocity gives a unique figure-eight pattern for barred galaxies, since some orbits are depleted in these systems." + Our simulation results have been similarly plotted., Our simulation results have been similarly plotted. + After initial xw formation. the plot appears similar to the unbarred dots: as the bar develops. the central positions darken while the clensity increases. and the areas immediately to he outside of this central pole suller some depopulation. similar to the observational plots in WaAIAL and DD. With bar buckling. though. comes increased scatter in he plot. removing any forming pattern. and once again caving a plot which would be considered unbarred i£ viewed observationallv.," After initial bar formation, the plot appears similar to the unbarred plots; as the bar develops, the central positions darken while the density increases, and the areas immediately to the outside of this central pole suffer some depopulation, similar to the observational plots in M and B. With bar buckling, though, comes increased scatter in the plot, removing any forming pattern, and once again leaving a plot which would be considered unbarred if viewed observationally." + Plots at /250 and /—76 are shown in Fig., Plots at $t$ =50 and $t$ =76 are shown in Fig. + 12., 12. + The original simulation is composed only of a disc and halo anc was constructed to be in equilibrium. with a JToomre (OQ greater. than 1., The original simulation is composed only of a disc and halo and was constructed to be in equilibrium with a Toomre Q greater than 1. + This generates a higher velocity. dispersion. than would be observed: 0.23. (times maximum velocity) tangential ancl 0.215 radial dispersion at 2.4 scale lengths., This generates a higher velocity dispersion than would be observed: 0.23 (times maximum velocity) tangential and 0.215 radial dispersion at 2.4 scale lengths. + Phe radial dispersion increases to 0.258 by £216., The radial dispersion increases to 0.258 by $t$ =76. + “Phis blurs the Geure-cight pattern in the IK&MM plot., This blurs the figure-eight pattern in the M plot. + Since all the observational results used. cool gas to determine velocities. our results would be much closer to the analytically predicted orbits. and the distinctive figure-cight plot.," Since all the observational results used cool gas to determine velocities, our results would be much closer to the analytically predicted orbits, and the distinctive figure-eight plot." + ‘To counter part of the possible velocity. dispersion problem statec above. a model containing a small. compact bulge was constructed in the lower resolution range (0001 particles). keeping a similar halo and rotation curve (see Section 2 for a more complete description).," To counter part of the possible velocity dispersion problem stated above, a model containing a small, compact bulge was constructed in the lower resolution range (500K particles), keeping a similar halo and rotation curve (see Section 2 for a more complete description)." + The evolution of the disc is similar to the earlier model in that the bar is initiated. a little before. 7-20 (540 Myr): it differs in that the initial bar is about 20 per cent smaller. with dumbbell-shaped inner isophotes when viewed [face-on.," The evolution of the disc is similar to the earlier model in that the bar is initiated a little before $t$ =20 (540 Myr); it differs in that the initial bar is about 20 per cent smaller, with dumbbell-shaped inner isophotes when viewed face-on." + The bar remains smaller than the the previous model. by about the same amount for the duration of the simulation.," The bar remains smaller than the the previous model, by about the same amount for the duration of the simulation." + This is in clirect contrast with the findings of MM who found their bulge model to have a longer bar. although the shape is similar.," This is in direct contrast with the findings of M who found their bulge model to have a longer bar, although the shape is similar." + Llowever. their bulge is more massive in comparison to the disk. and more extended than ours.," However, their bulge is more massive in comparison to the disk, and more extended than ours." + Fig., Fig. + 13 shows /—26 o 250., 13 shows $t$ =26 to 250. + The initial velocity. dispersion (0.146. versus 0.23 above for tangential and 0.18. versus 0.215 for racial) was ower than our original simulation: however. the line-of-sight velocity distribution (LOSVD) plots remained similar.," The initial velocity dispersion (0.146 versus 0.23 above for tangential and 0.18 versus 0.215 for radial) was lower than our original simulation; however, the line-of-sight velocity distribution (LOSVD) plots remained similar." + The veh resolution simulation showed much more detail and ollowecl trends toward the figure-cight shape not seen in he lower resolution plots for the same initial conditions., The high resolution simulation showed much more detail and followed trends toward the figure-eight shape not seen in the lower resolution plots for the same initial conditions. + Creating a simulation of the bulge mass model with higher resolution. with its lowered. velocity dispersion. should vield clearer LOSVD plots than the original high resolution run. therefore we cannot make anv conclusions as to whether or not the inflated velocity dispersion in the original sipiulation is responsible for masking the figure-eight. pattern.," Creating a simulation of the bulge mass model with higher resolution, with its lowered velocity dispersion, should yield clearer LOSVD plots than the original high resolution run, therefore we cannot make any conclusions as to whether or not the inflated velocity dispersion in the original simulation is responsible for masking the figure-eight pattern." + The pattern speed of the bulge model is higher than, The pattern speed of the bulge model is higher than +represent discrete physical structures.,represent discrete physical structures. + Columns (2) and (3) are the coordinates of the sources. and column (4) is the relative positions of the these sources will respect to the strongest component.," Columns (2) and (3) are the coordinates of the sources, and column (4) is the relative positions of the these sources with respect to the strongest component." + Column (5) lists the surface brightnesses of these sources. and column (6) their integrated. [lux densities.," Column (5) lists the surface brightnesses of these sources, and column (6) their integrated flux densities." + Column (7) gives the nominal deconvolved sizes of the Gaussian components at EWIIM as given by JMETT. aud column 8 lists the position angles of the fitted Gaussians.," Column (7) gives the nominal deconvolved sizes of the Gaussian components at FWHM as given by JMFIT, and column 8 lists the position angles of the fitted Gaussians." + The corresponding intrinsic brightness temperatures of (hese compact sources are on the order of 10* to 103 IK. and are listed in column 9.," The corresponding intrinsic brightness temperatures of these compact sources are on the order of $10^7$ to $10^8$ K, and are listed in column 9." + Dased on the VLBA results. the source is composed of (wo dominant structures separated bv ~31 mas.," Based on the VLBA results, the source is composed of two dominant structures separated by $\sim 31$ mas." + The stronger of (hese is consistent will (wo Gaussians (components 1 and 3 in Table 3). while the second dominant source is represented by one Gaussian (component 2 in Table 3).," The stronger of these is consistent with two Gaussians (components 1 and 3 in Table 3), while the second dominant source is represented by one Gaussian (component 2 in Table 3)." + Including a possible faint component to the east (component + in Table 3). gives a total [lux density of 1.7780.109 mJy.," Including a possible faint component to the east (component 4 in Table 3), gives a total flux density of $1.778 \pm +0.109$ mJy." + This value is consistent with the 1.73£0.13 mJv obtained with the VLA FIRST survev (Decker.White.&Ilelfand1995).. and with the 1.316x:0.021 mJv obtained with the VLA ELAIS survey (Ciliegietal.1999).," This value is consistent with the $1.73 \pm 0.13$ mJy obtained with the VLA FIRST survey \citep{BWH95}, and with the $1.816 \pm 0.021$ mJy obtained with the VLA ELAIS survey \citep{CP99}." +. The 1.4 GIIz flux densities measured with the VLA in 1995 and 1997 (Becker.Cilieeietal.1999) and the VLBA in 2007 (this paper) are equal to better than5%... implvine that this source is not highly variable on time scales of vears.," The 1.4 GHz flux densities measured with the VLA in 1995 and 1997 \citep{BWH95,CP99} and the VLBA in 2007 (this paper) are equal to better than, implying that this source is not highly variable on time scales of years." +" We have also svnthesized larger images (2""x 2"") using the VLBA and found no other radio components al >5e level (140 (Jy. +) in the field other than those seen in Figure 4 and listed in Table 3.", We have also synthesized larger images $2'' \times 2''$ ) using the VLBA and found no other radio components at $\ge 5\sigma$ level $140~\mu$ Jy $^{-1}$ ) in the field other than those seen in Figure 4 and listed in Table 3. + We have detected the z=6.12 QSO J14272-3312 at 8.4 GIIz with the VLA A-array and al 1.4 GllIz with the VLBA., We have detected the $z=6.12$ QSO J1427+3312 at 8.4 GHz with the VLA A-array and at 1.4 GHz with the VLBA. + The source is unresolved as seen in the 8.4 GIIz VLA results. and has a steep spectrum wilh a spectral index value of a1=--.1.," The source is unresolved as seen in the 8.4 GHz VLA results, and has a steep spectrum with a spectral index value of $\alpha^{8.4}_{1.4}=-1.1$." + At mas resolution. the VLBA observations show that this QSO is comprised of two dominant continuum components separated by 31 mas (174 pe: Table 3) with a flux density ratio of ~3:1.," At mas resolution, the VLBA observations show that this QSO is comprised of two dominant continuum components separated by 31 mas (174 pc; Table 3) with a flux density ratio of $\sim 3:1$." + The Gaussian fitting suggests that both components are resolved with sizes ol ~4—6 mas (22—34 pe)., The Gaussian fitting suggests that both components are resolved with sizes of $\sim 4-6$ mas $22-34$ pc). + The physical properties observed in this source suggest (hat this high-: QSO could be a Compact Sxiinetric Object (CSO) with two distinct. steep spectra. radio lobes that are confined) by a dense ISM in the host galaxy. (Conway2002).," The physical properties observed in this source suggest that this $z$ QSO could be a Compact Symmetric Object (CSO) with two distinct, steep spectra, radio lobes that are confined by a dense ISM in the host galaxy \citep{CO02}." +.. CSOs are radio sources that have sizes on scales of 1 pe to 1 kpe. and are thought to be verv voung (X10! vi: Reaclheacl (1998))).," CSOs are radio sources that have sizes on scales of 1 pc to 1 kpc, and are thought to be very young $\leq 10^4$ yr; \citet{RTXPWP96,OC98}) )." + Moreover. because CSOs are highly confined," Moreover, because CSOs are highly confined" +Dress Gunn (1973) were the first to propose that a COSELOogical abundance of dark compact objects could be detected. by gravitational lensing of more clistant sources.,Press Gunn (1973) were the first to propose that a cosmological abundance of dark compact objects could be detected by gravitational lensing of more distant sources. + Ifa point mass lies along the observer-source line of sight. the relative motion tween. the lens. source ancl observer produces a change in the magnifications of two lensed images.," If a point mass lies along the observer-source line of sight, the relative motion between the lens, source and observer produces a change in the magnifications of two lensed images." + The presence of foreground. compact objects. is therefore. detected. through a change in the observed. Dux of the background source (termed microlensing)., The presence of foreground compact objects is therefore detected through a change in the observed flux of the background source (termed microlensing). + This οσοι has been used successfully in the search for compact objects in the halo of the Milky Way galaxy (e.g. cock et a., This effect has been used successfully in the search for compact objects in the halo of the Milky Way galaxy (e.g. Alcock et al. + 2000)., 2000). + Also. microlensing «ue to stars in a galaxy at moderate redshift has been observed. in the. gravitationally lensed quasar Q2237|0305 (Lewin et al.," Also, microlensing due to stars in a galaxy at moderate redshift has been observed in the gravitationally lensed quasar Q2237+0305 (Irwin et al." + 1989: Corriganet al., 1989; Corrigan et al. + 1991)., 1991). + The short durations of gamma rav bursts (CBs) oller an alternativo way to study microlensing and hence search [or à cosmological population of compact objects. through observation of repeating bursts.," The short durations of gamma ray bursts (GRBs) offer an alternative way to study microlensing and hence search for a cosmological population of compact objects, through observation of repeating bursts." + In this paper we assume GRBs to be at cosmological distances (Paczynski 1995). and therefore probable sources for gravitational lensing (Lurner. Ostriker Cott 1984).," In this paper we assume GRBs to be at cosmological distances (Paczynski 1995), and therefore probable sources for gravitational lensing (Turner, Ostriker Gott 1984)." + Paczyvuski (1986. LOST) noted that while multiple images of a single GRB cannot be angularly resolved: by present clay detectors. their relative ἅσίαν may be longer than the burst duration so that the lensecl images could be resolved temporally (a pair of lensed GiB images that are produced bv a single mass (Ade) have a relative time delay that is AL~5Oscc(Ade10AL.) (Mao 1992)).," Paczynski (1986, 1987) noted that while multiple images of a single GRB cannot be angularly resolved by present day detectors, their relative delay may be longer than the burst duration so that the lensed images could be resolved temporally (a pair of lensed GRB images that are produced by a single mass $(M_{CO})$ have a relative time delay that is $\Delta t\sim 50 sec\times(M_{CO}/10^{6}M_{\odot})$ (Mao 1992))." + Microlensing of aGkB by a compact object is therefore observed as a GRB that repeats., Microlensing of a GRB by a compact object is therefore observed as a GRB that repeats. + Phe utility. of GRBs to explore cosmological dark matter. as well as the possible inference of properties of the GIU population itself was cliscussed in detail by Blaes Webster (1992).," The utility of GRBs to explore cosmological dark matter, as well as the possible inference of properties of the GRB population itself was discussed in detail by Blaes Webster (1992)." + The short event duration. as well as the transparcney of the universe to ganmuna ravs make GRBs ideal probes of dark matter in the form of compact objects over à wide range of masses.," The short event duration, as well as the transparency of the universe to gamma rays make GRBs ideal probes of dark matter in the form of compact objects over a wide range of masses." + Alicrolensing of existing ancl potential catalogues of GRBs have been used to discuss the cosmological abundance of compact objects., Microlensing of existing and potential catalogues of GRBs have been used to discuss the cosmological abundance of compact objects. + Marani et al. (, Marani et al. ( +1999) use non-detections of lensed images from the BATSE and Ulysses catalogues to set conservative limits on dark compact objects with masses between LO15 and LOAZ...,1999) use non-detections of lensed images from the BATSE and Ulysses catalogues to set conservative limits on dark compact objects with masses between $10^{-16}$ and $10^{-7}M_{\odot}$. + Also a universe proposed. by Gnedin Ostriker (1992) with Meo107A. and O3;=Oco—(015 was ruled out at a confidence level of., Also a universe proposed by Gnedin Ostriker (1992) with $M_{CO}\sim10^{6.5}M_{\odot}$ and $\Omega_{M}=\Omega_{CO}=0.15$ was ruled out at a confidence level of. +. This scenario had. been previously investigated. in detail. using a 3-D lensing code by Mao (1993).," This scenario had been previously investigated in detail, using a 3-D lensing code by Mao (1993)." + He found results that did not depart significantly from those obtained. by a single- approximation., He found results that did not depart significantly from those obtained by a single-screen approximation. + Mao (1992) and Crossman Nowak, Mao (1992) and Grossman Nowak +We lave analysed the distribution of various ‘imematical aspects of our sample.,We have analysed the distribution of various kinematical aspects of our sample. + To ect the information for the overall kinematics of the and to attempt to come to a separation of disk and halo we calculated their orbits., To get the information for the overall kinematics of the and to attempt to come to a separation of disk and halo we calculated their orbits. + The observational parameters 0.0.4.fy.fs.rad Were transformed iuto ἂν}.ZEW (for details see de Boer 1997b).," The observational parameters $\alpha,\delta,d,\mu_{\alpha},\mu_{\delta},v_{\rm rad}$ were transformed into $X,Y,Z,U,V,W$ (for details see de Boer 1997b)." + Also the orbital velocity projected to the ealactic plane. O. and the velocity towards the ealactic ceutre. ®. were calculated.," Also the orbital velocity projected to the galactic plane, $\Theta$, and the velocity towards the galactic centre, $\Phi$, were calculated." + Orbits were calculated over a time span of 10 taking steps of LATIN., Orbits were calculated over a time span of 10 taking steps of Myr. +" Following de Doer (1997) the eccentricity. cece. of the orbit aud the apo- aud perigalactic distances. Ry aud fé were calculated as well as the maxima -dstauce reached. μιας. ad its ealactic radial distance normalised z-extenut. nic. For the valuesIh, for our stars sec Table available in electrouic form."," Following de Boer (1997b) the eccentricity, $ecc$, of the orbit and the apo- and perigalactic distances, $R_{\rm a}$ and $R_{\rm p}$, were calculated as well as the maximum $z$ -distance reached, $z_{\rm max}$, and its galactic radial distance normalised $z$ -extent, $nze$, For the values for our stars see Table \\ref{tablecds} available in electronic form." + The meridional sections of the orbits of ({sec reforhits.fig)) show very different shapes., The meridional sections of the orbits of (see \\ref{orbits.fig}) ) show very different shapes. +" About half of the stars of our sample have perigalactic distances fü,x 3kkpc."," About half of the stars of our sample have perigalactic distances $R_{\rm p} +\leq 3$ kpc." + 65 rroach a perigalactic distance Ryx Lkkpe., 65 reach a perigalactic distance $R_{\rm p}\leq 1$ kpc. + Three stars AAnd. VVir and VVir) reach perigalactic distances less than 0.1 kpe.," Three stars And, Vir and Vir) reach perigalactic distances less than 0.1 kpc." + The other extreme is represented by 5 (CTAAnd. ALCCAG. LLyu. PPer and PPsc) which have perigalactic distances more than 9 Ipc aud always stav bevoud the solar circle.," The other extreme is represented by 5 And, CMi, Lyn, Per and Psc) which have perigalactic distances more than 9 kpc and always stay beyond the solar circle." + Of the 217 663 have boxy orbits auc stay close to the plane., Of the 217 63 have boxy orbits and stay close to the plane. + A subset of 29 inakes only nall excursions in a reforbits.fig top row). 31 have planar but very ecceutric orbits 2. iuiddle row).," A subset of 29 makes only small excursions in $\varpi$ \\ref{orbits.fig} top row), 34 have planar but very eccentric orbits \ref{orbits.fig} middle row)." +" The orbit of itself is not like that of a disk population star (oven although its 444 is only kkpe aud its το= 0.01). because its orbit has ecc=OS with 4,= δρ aud R=9.1 kkpe."," The orbit of itself is not like that of a disk population star (even although its $z_{\rm max}$ is only kpc and its $nze = 0.01$ ), because its orbit has $ecc =0.8$ with $R_{\rm a}=18.4$ kpc and $R_{\rm p}= 2.1$ kpc." + 151 of the sshow orbits resenibliug those of he bottoni row., 154 of the show orbits resembling those of the bottom row. + The shape of these orbits is really chaotic. aud shows movements perpendicular to the plane.," The shape of these orbits is really chaotic, and shows movements perpendicular to the plane." +" These stars partly have orbits geome to vorv small ealactocentric distances and some of them reach very laree apogalactic distances. with as an extreme Ih,Si kkpe (for AAÀqr)."," These stars partly have orbits going to very small galactocentric distances and some of them reach very large apogalactic distances, with as an extreme $R_{\rm p} \simeq 51$ kpc (for Aqr)." + The statistics of the circular componcut O of the velocity is shown in 1tthetalist.fig.., The statistics of the circular component $\Theta$ of the velocity is shown in \\ref{thetahist.fig}. + In nr suuple we have 87 wwitl retrograde orbits., In our sample we have 87 with retrograde orbits. + We compared the characteristics of these stars with the prograde ones., We compared the characteristics of these stars with the prograde ones. + Among the retrograde part of our sample there are no stars with a iietallicity [Fe/TI] >0.9 (see rofanavel upper eft panel)., Among the retrograde part of our sample there are no stars with a metallicity [Fe/H] $>-0.9$ (see \\ref{anavel} upper left panel). + Considering the retrograde eroup separately shows that the peak at high eccentricities in refeccuzehistfig (eft panel) is nearly completely due to the retrogradeLxis.. while the stars with prograde orbits show a flat distribution over the hole metallicity ranee.," Considering the retrograde group separately shows that the peak at high eccentricities in \\ref{eccnzehist.fig} (left panel) is nearly completely due to the retrograde, while the stars with prograde orbits show a flat distribution over the hole metallicity range." + The ecceutzricities ecc of the orbits as well as the values of ize span a large range (see rofeccuzelist.fig))., The eccentricities $ecc$ of the orbits as well as the values of $nze$ span a large range (see \\ref{eccnzehist.fig}) ). + The eccentricity of the orbit of the majority ofour stars is ccc.>0.15., The eccentricity of the orbit of the majority of our stars is $ecc~>0.45$. + The distribution has an absolute maxima at cece=0.9 and a shallow local τή at eeczz0.6., The distribution has an absolute maximum at $ecc= 0.9$ and a shallow local minimum at $ecc\simeq 0.6$. +" The distribution of the uormalised z-exteuts shows a nani at low values neκ0,9, a local nunuuuuni at nie=0.4 followed by another peak at niez0.6."," The distribution of the normalised $z$ -extents shows a maximum at low values $nze<0.2$, a local minimum at $nze =0.4$ followed by another peak at $nze \simeq 0.6$." + There are 100 wwith sies0 12d only 15 mreach niecol., There are 109 with $nze \leq 0.4$ and only 45 reach $nze \geq 1$. + A separation in stars with prograde aud retrograde orbits does not chauge the appearance of the Dic diagram., A separation in stars with prograde and retrograde orbits does not change the appearance of the $nze$ diagram. + Both eroups show Ligh peaks at low values of nie anda oug flat tail., Both groups show high peaks at low values of $nze$ and a long flat tail. + The shape of a stellar orbit gives an indication of the population a star belongs to., The shape of a stellar orbit gives an indication of the population a star belongs to. + stars have orbits simular to those of the younger galactic stars. deine rather circular and staving close to the plane.," stars have orbits similar to those of the younger galactic stars, being rather circular and staying close to the plane." + These are the ones nost likely haviug been bora iu the disk., These are the ones most likely having been born in the disk. + It is. from our data. not easy to discriminate between so-called thin- aud thick-clisk stars. a distinction which is uot well defined auvway.," It is, from our data, not easy to discriminate between so-called thin- and thick-disk stars, a distinction which is not well defined anyway." + The cousists of stars whose orbits are not akin to disk-like galactic rotation., The consists of stars whose orbits are not akin to disk-like galactic rotation. + They most likely have been born outside he disk of the Allkv Wav., They most likely have been born outside the disk of the Milky Way. + Using the above criteria (which were largely set by the distribution of the parameters of our stars as discussed above uxiug aud 5)). one can now roughly sort the stars according to populations.," Using the above criteria (which were largely set by the distribution of the parameters of our stars as discussed above using \\ref{thetahist.fig} and \ref{eccnzehist.fig}) ), one can now roughly sort the stars according to populations." + We define halo stars as those having O-100 oor ecce>OL or nie> 0.L., We define halo stars as those having $\Theta<100$ or $ecc > 0.4$ or $nze>0.4$ . + Each one of the criteria alone, Each one of the criteria alone +textscii] 46548 A and [N textscii] 46584 A transitions must exhibit identical non-thermal motions since both lines are emitted by the same ions.,$]$ $\lambda$ 6548 $\mbox{\AA}$ and $[$ $]$ $\lambda$ 6584 $\mbox{\AA}$ transitions must exhibit identical non-thermal motions since both lines are emitted by the same ions. + The same argument also applies to the [S textscii] 46716 A and [S textscii] 46731 A lines., The same argument also applies to the $[$ $]$ $\lambda$ 6716 $\mbox{\AA}$ and $[$ $]$ $\lambda$ 6731 $\mbox{\AA}$ lines. +" Finally, specific line ratios of the [N textscii]and[S textscii] transitions must agree with standard theoretical models of ionized nebulae."," Finally, specific line ratios of the $[$ $]$ and $[$ $]$ transitions must agree with standard theoretical models of ionized nebulae." +" All conditions taken in consideration, only 30057 emission-line profiles were retained, most of them being associated to the bright, central structure found in the overlapping region between the eastern and western field (see 8 5.2.1)."," All conditions taken in consideration, only 057 emission-line profiles were retained, most of them being associated to the bright, central structure found in the overlapping region between the eastern and western field (see $\S$ 5.2.1)." +" Obviously, the summation of both cubes largely contributes to increase the data quality for duplicated pixels not only by increasing the peak signal for all lines but also by flattening noise fluctuations in empty channels (see Figure 2)."," Obviously, the summation of both cubes largely contributes to increase the data quality for duplicated pixels not only by increasing the peak signal for all lines but also by flattening noise fluctuations in empty channels (see Figure 2)." +" Considering the very large number of emission-line profiles available in our initial data set (more than 0000), we are fully aware that the conditions listed above drastically reduce the size of the retained sample."," Considering the very large number of emission-line profiles available in our initial data set (more than 000), we are fully aware that the conditions listed above drastically reduce the size of the retained sample." +" On the other hand, these conditions assure that the retained profiles are undoubtedly the most reliable available."," On the other hand, these conditions assure that the retained profiles are undoubtedly the most reliable available." + We reiterate that a S/N greater than 6 is required to appropriately conduct an investigation on line ratios., We reiterate that a S/N greater than 6 is required to appropriately conduct an investigation on line ratios. +" However, a *physical detection"" [of a given line] is considered when Z33 (Rola&Pelat 1994).."," However, a “physical detection” [of a given line] is considered when $\gtrsim$ 3 \citep{Rol1994}. ." +" This will be used in a later subsection to investigate notable, although not necessarily reliable, line ratios for particular structures in 11805 (see 8 5.2.2.2)."," This will be used in a later subsection to investigate notable, although not necessarily reliable, line ratios for particular structures in 1805 (see $\S$ 5.2.2.2)." +" The [S[Su]u]A6716A6731 line ratios were computed and used as input values for the Fivel procedure (DeRobertisetal.1987) adapted, for convenience, inIDL."," The $\frac{[\textnormal{S}\,\textsc{ii}]\,\lambda6716}{[\textnormal{S}\,\textsc{ii}]\,\lambda6731}$ line ratios were computed and used as input values for the Fivel procedure \citep{DeR1987} adapted, for convenience, in." +. Assuming a constantelectron temperature of 74400 K throughout, Assuming a constantelectron temperature of 400 K throughout +al progressively longer wavelengths (Meszaros&ReesL997).,at progressively longer wavelengths \citep{Meszaros1997}. +. Ever since the discovery of the first afterglows. these have been modeled succeshilly by combining a model for the blast wave ονπας wilh a svehirotron radiation model. where shock-accelerated particles radiate by interacting with a shock-generated magnetic field (e.g. Wijersetal.1997:&Galama1999:Frailetal.2000:Panaiteseu&Ixirunar. 2002)).," Ever since the discovery of the first afterglows, these have been modeled succesfully by combining a model for the blast wave dynamics with a sychrotron radiation model, where shock-accelerated particles radiate by interacting with a shock-generated magnetic field (e.g. \citealt{Wijers1997, Wijers1999, Frail2000, PK2002}) )." + AnalvGecally tractable solutions for the dvnamies are the self-similar Blandford-MeIxee (BAI. Blandford&Melee1976)). and sedov-von Neumann-Tavlor (ST. Sedov1959:Tavlor1950:VouNeumann 1961)) solutions describing. respectively. (he ulira-relativistic ancl non-relativistic phase of the blast wave evolution.," Analytically tractable solutions for the dynamics are the self-similar Blandford-McKee (BM, \citealt{Blandford1976}) ) and Sedov-von Neumann-Taylor (ST, \citealt{Sedov1959, Taylor1950, vonNeumann1961}) ) solutions describing, respectively, the ultra-relativistic and non-relativistic phase of the blast wave evolution." + At early time lateral spreading of the collimated jet has not vet set in and the outflow is purely radial. while al late times the jet will have become (uly spherical. allowing the application of spherically svmmetrie solutions in both cases.," At early time lateral spreading of the collimated jet has not yet set in and the outflow is purely radial, while at late times the jet will have become truly spherical, allowing the application of spherically symmetric solutions in both cases." + As of vet. no analvtieal solution exists that fully captures the intermediate stage of the blast wave evolution. where the blast wave becomes trausrelativistic. inhomogenous along the shock front (Zhang&Mac-Facdven2009;VanEerten&MacFadven2011) and decollimates.," As of yet, no analytical solution exists that fully captures the intermediate stage of the blast wave evolution, where the blast wave becomes transrelativistic, inhomogenous along the shock front \citep{Zhang2009, vanEerten2011jetspreading} and decollimates." + Early attempts assumed a homogeneous shock front (Rhoads1999). or spherical outflow (παςetal.1999).. while even recently studies (e.g. Granot&Piran 2011)) do not account [or the radial structure of the jet.," Early attempts assumed a homogeneous shock front \citep{Rhoads1999} or spherical outflow \citep{Huang1999}, while even recently studies (e.g. \citealt{Granot2011}) ) do not account for the radial structure of the jet." + The practical implications of the fact that the blast wave evolution is determined by a very small number of variables such that scalings between different explosion energies and circumbiurst density are possible were not fully realized until verv recently., The practical implications of the fact that the blast wave evolution is determined by a very small number of variables such that scalings between different explosion energies and circumburst density are possible were not fully realized until very recently. + The scalings apply in the asvimptolic sell-similar limits. but also in the intermediate regime where the two-dimensional nature of jet decollimation is in full effect.," The scalings apply in the asymptotic self-similar limits, but also in the intermediate regime where the two-dimensional nature of jet decollimation is in full effect." + This made it possible to use only a small set of simulations for different initial jel opening angles as a basis lor a fit code that can be applied to broadband aftereglow data (VanEertenetal.2011).., This made it possible to use only a small set of simulations for different initial jet opening angles as a basis for a simulation-based fit code that can be applied to broadband afterglow data \citep{vanEerten2011boxfit}. + Bul even though a complete recalculation of the dynamics of the blast wave is no longer necessary. (here remained the equations of radiative Urausfer of (a representative number of) ravs (rough (he evolving jet. that have to be solved for each datapoint fit iteration.," But even though a complete recalculation of the dynamics of the blast wave is no longer necessary, there remained the equations of radiative transfer of (a representative number of) rays through the evolving jet, that have to be solved for each datapoint fit iteration." + For a large nuniber of iterations and datapoints this procedure remains computationally expensive and requires the use of a parallel computer., For a large number of iterations and datapoints this procedure remains computationally expensive and requires the use of a parallel computer. + The current study shows (hat the calculation (ime for a given light curve or spectrum can be futher reduced., The current study shows that the calculation time for a given light curve or spectrum can be further reduced. + We demonstrate that scalability between blast waves has straightforward implications for scalability between light curves., We demonstrate that scalability between blast waves has straightforward implications for scalability between light curves. + In 82. we describe how the scaling relations for the clvnamics of blast waves can be used (o scale between light curves as well., In \ref{scalings_section} we describe how the scaling relations for the dynamics of blast waves can be used to scale between light curves as well. + We show that the sealings remain unchanged. between the DM. ancl ST regimes. and in 33 we demonstrate numerically that the sealings also hold in the intermediate regime and [or off-axis observers.," We show that the scalings remain unchanged between the BM and ST regimes, and in \ref{numerics_section} we demonstrate numerically that the scalings also hold in the intermediate regime and for off-axis observers." + We discuss our findings in 84.., We discuss our findings in \ref{discussion_section}. . +"grid along its cardinal directions, necessary to maximise the usage of the cubic simulation volume.","grid along its cardinal directions, necessary to maximise the usage of the cubic simulation volume." + We focus on the thermal state of the gas surrounding the central source before and after the II-front passes., We focus on the thermal state of the gas surrounding the central source before and after the I-front passes. + The evolution of the temperature in slice through the simulation is shown in Figure 1.., The evolution of the temperature in a slice through the simulation is shown in Figure \ref{fig:temperature}. + The role aplayed by dense structures surrounding the source in restraining the growth of the aand ILfronts is apparent in the highly anisotropic temperature patterns., The role played by dense structures surrounding the source in restraining the growth of the and I-fronts is apparent in the highly anisotropic temperature patterns. +" For comparison, the density field at z—2.6 is shown in Figure 2.. ("," For comparison, the density field at $z=2.6$ is shown in Figure \ref{fig:overdensity}. (" +"Since the overdensities change very little over the redshift interval 2.6«z4.6, only the image at z—2.6 is shown.)","Since the overdensities change very little over the redshift interval $2.6 0.," We emphasize here, $\varphi_1$ may not be continuous up to the boundary, i.e. $\varphi_1 \notin C(\overline{Q})$ when $g(x)< f(x)$ for some$x>0$ ." +Alternatively. by performing the O-integral over the 6 function. one obtains where the final expression is obtained by partially integrating twice.,"Alternatively, by performing the $\Theta$ -integral over the $\delta$ function, one obtains ^2(z) ^2(z_c), where the final expression is obtained by partially integrating twice." + Noting that the inlegral over O in (3.3)) is equal to 4/15. the two results agree.," Noting that the integral over $\Theta$ in \ref{eta5}) ) is equal to $4/15$, the two results agree." + An alternative form is , An alternative form is P _c). +The power per unit frequency. (o). is identiliel bywriting the right hand side οἱ οςa ). (3.3)) or (3.3)) as the integral of P(e) over dic.," The power per unit frequency, $P(\omega)$, is identified bywriting the right hand side of \ref{eta5}) ),\ref{eta5a}) ) or \ref{eta5b}) ) as the integral of $P(\omega)$ over $\rmd\omega$ ." + One finds, One finds ^2(z_c) _c). +WolfRavet (WR) stars are evolved massive stars characterized by hieh mass loss ratos (107 and related rotation angle 6 (defined below) rotate sinoothlv about the rotation axis iu a counter-clockwise direction frou A to D (again this continuous rotation for triplets A to D is achieved oulv if source 10 is chosen for the singlet in triplet D)., The triplet position angle $\beta$ and related rotation angle $\theta$ (defined below) rotate smoothly about the rotation axis in a counter-clockwise direction from A to D (again this continuous rotation for triplets A to D is achieved only if source 10 is chosen for the singlet in triplet D). + This rotation direction is the same as the galaxy rotation direction and may sugeest that the triplet position angele is fixed relative to the co-ordinates of the ceutral object., This rotation direction is the same as the galaxy rotation direction and may suggest that the triplet position angle is fixed relative to the co-ordinates of the central object. + The chanec in triplet position angle is thus a measure of the rotation of the central object., The change in triplet position angle is thus a measure of the rotation of the central object. + This possibility is used below to estimate the rotation period ofthe central objec, This possibility is used below to estimate the rotation period of the central object. + Tn all four triplets in Fig., In all four triplets in Fig. + L. the pair orientation auele w (defined by the position angle ou the sky of the line counecting between the two sources 1u cach pair) rotates continuously with the triplet position angle ./ and. im so doing. maintains a siuilar oricutation relative to the siuglet/pair midpoiut/rotation axis plane.," 4, the pair orientation angle $\omega$ (defined by the position angle on the sky of the line connecting between the two sources in each pair) rotates continuously with the triplet position angle $\beta$ and, in so doing, maintains a similar orientation relative to the singlet/pair midpoint/rotation axis plane." + In effect. alb four triplets are similarly structured. differius onlv iu their rotation angles and sizes.," In effect, all four triplets are similarly structured, differing only in their rotation angles and sizes." + To obtain this result bv chance seems extremely unlikely if it is remembered that the sources m each triplet were, To obtain this result by chance seems extremely unlikely if it is remembered that the sources in each triplet were +an H-index below around 7.5 pm. the average profiles show only weak emission features.,"an H-index below around 7.5 pm, the average profiles show only weak emission features." + If the H-index exceeds 7.5 pm. a pronounced asymmetry of Hay and Hog is seen.," If the H-index exceeds 7.5 pm, a pronounced asymmetry of $_{\rm 2V}$ and $_{\rm 2R}$ is seen." + Profiles with a strongly enhanced Hy peak cover around of the area. in agreement with S96.," Profiles with a strongly enhanced $_{\rm 2V}$ peak cover around of the area, in agreement with S96." + If however the area fractions of all profiles with emission signatures (H-index >7.5 pm) are added up. the ratio of profiles in emission to those without is 60:40.," If however the area fractions of all profiles with emission signatures (H-index $>$ 7.5 pm) are added up, the ratio of profiles in emission to those without is 60:40." + Taking the full FOV. the area fraction of profiles with an H-index above 7.5 pm is 64%.," Taking the full FOV, the area fraction of profiles with an H-index above 7.5 pm is 64." +. Another estimate can be made from Fig. 18:, Another estimate can be made from Fig. \ref{fig21}: +: a shock event affects usually three to four profiles. Le. it leads to emission for 60 to 80 seconds afterwards.," a shock event affects usually three to four profiles, i.e. it leads to emission for 60 to 80 seconds afterwards." + If the next shock happens 180 seconds after the first one. the fraction of time spent in emission is around 70/180~40 %.," If the next shock happens 180 seconds after the first one, the fraction of time spent in emission is around $\sim$ 40." +. This definition of the emission from the appearance of any Hay peak then suggests in all estimates that the chromosphere. or more precisely. the core of the Ca II H line spends around half of the time in emission instead of 10%.," This definition of the emission from the appearance of any $_{\rm 2V}$ peak then suggests in all estimates that the chromosphere, or more precisely, the core of the Ca II H line spends around half of the time in emission instead of 10." +. If this emission can by modeled by a static temperature rise. or reflects a temperature rise at all. is another question.," If this emission can by modeled by a static temperature rise, or reflects a temperature rise at all, is another question." + On locations with detected photospheric fields. a quasi-permanent increase of intensity in both emission peaks is present. in addition to similar repetitive bright grains as happen outside fields.," On locations with detected photospheric fields, a quasi-permanent increase of intensity in both emission peaks is present, in addition to similar repetitive bright grains as happen outside fields." + Near to. but still outside strong photospheric magnetic. fields. the emission is generally increased (cf.," Near to, but still outside strong photospheric magnetic fields, the emission is generally increased (cf." + Figs., Figs. + 4. or 16))., \ref{fig5} or \ref{fig19}) ). + Interestingly. the maximum H-index observed in the time series is located outside of magnetic fields. which could fit with the suggestion of ? that collisions between flux concentrations and granules are responsible for the creation of bright grains.," Interestingly, the maximum H-index observed in the time series is located outside of magnetic fields, which could fit with the suggestion of \citet{kalkofen1996} that collisions between flux concentrations and granules are responsible for the creation of bright grains." + We note however that in our case it would be the interaction of a strong uni-polar network element with granulation instead of the weaker mixed-polarity fields suggested by ?.., We note however that in our case it would be the interaction of a strong uni-polar network element with granulation instead of the weaker mixed-polarity fields suggested by \citet{kalkofen1996}. + The areas with least chromospheric emission in our time series are located furthest away (> 10%) from any magnetic fields. in the middle of the observed field of view.," The areas with least chromospheric emission in our time series are located furthest away $>10^{\prime\prime}$ ) from any magnetic fields, in the middle of the observed field of view." + The spatial distribution of emission in our ]-D slit observations would comply very well with a cut through the FOV observed by οςVO7.. if the slit would be placed across one of the field concentrations. visible in their Fig.," The spatial distribution of emission in our 1-D slit observations would comply very well with a cut through the FOV observed by \citet[][V07]{vecchio+etal2007}, if the slit would be placed across one of the field concentrations visible in their Fig." + 2., 2. + The halo of enhanced emission close to the fields found in the present paper would correspond to one high-emission fibril seen in the Ca II nm line by VO7., The halo of enhanced emission close to the fields found in the present paper would correspond to one high-emission fibril seen in the Ca II nm line by V07. + These fibrils are interpreted to reflect the chromospheric magnetic field topology by V07. and enc after around 10 in low-emission dark regions.," These fibrils are interpreted to reflect the chromospheric magnetic field topology by V07, and end after around $10^{\prime\prime}$ in low-emission dark regions." + Like ?.. we do not see an one-to-one correlation ofemissiot in calcium and photospheric fields or Stokes V signal as claimed by ?.. in none of the Figs. 1.. 4..," Like \citet{lites+etal1999}, we do not see an one-to-one correlation of emission in calcium and photospheric fields or Stokes $V$ signal as claimed by \citet{sivaraman+etal2000}, in none of the Figs. \ref{fig2}, \ref{fig5}," + or 16.., or \ref{fig19}. + ? employed data very similar to ours. spectro-polarimetry in nm and spectroscopy in Ca IL H. There are several occurrences of H»y brightenings on locations without any polarization signal above our detection limit of 0.15 of the continuum intensity.," \citet{lites+etal1999} employed data very similar to ours, spectro-polarimetry in nm and spectroscopy in Ca II H. There are several occurrences of $_{\rm 2V}$ brightenings on locations without any polarization signal above our detection limit of 0.15 of the continuum intensity." +" The relation in the other direction is however rather tight: 1f photospheric fields are present. the emission is enhanced and often also affects the Ho, peak as well (Figs."," The relation in the other direction is however rather tight: if photospheric fields are present, the emission is enhanced and often also affects the $_{\rm 2R}$ peak as well (Figs." + 1. and 4))., \ref{fig2} and \ref{fig5}) ). + To quantify the visual impression. we use the scatter plot of integrated unsigned Stokes V signal vs the H-index (Fig. 21)).," To quantify the visual impression, we use the scatter plot of integrated unsigned Stokes $V$ signal vs the H-index (Fig. \ref{fig23}) )." + The scatter plot of the full FOV shows the usual behavior (e.g.22??).. a general increase of chromospheric emission with polarization signal. 1e.. with total magnetic flux.," The scatter plot of the full FOV shows the usual behavior \citep[e.g.][]{skumanich+etal1975,schrijver1987,schrijver+etal1989,reza+etal2007}, a general increase of chromospheric emission with polarization signal, i.e., with total magnetic flux." + To substantiate the claim that emission can occur without fields. we overplotted the values of the quiet region in the middle of the FOV (cf.," To substantiate the claim that emission can occur without fields, we overplotted the values of the quiet region in the middle of the FOV (cf." +Table 2.. 2nd row) separately t light grey.,"Table \ref{tab2}, 2nd row) separately in light grey." + It can be clearly seen that the emission in this part of the FOV covers the as the full FOV. from 6 to around 11 pm. but shows only weak polarization signals.," It can be clearly seen that the emission in this part of the FOV covers the as the full FOV, from 6 to around 11 pm, but shows only weak polarization signals." + We emphasize also again the conclusion of ?. that the presence of magnetic flux influences the minimum H-index. but that the maximum emission seems to be independent of the magnetic flux.," We emphasize also again the conclusion of \citet{reza+etal2007} that the presence of magnetic flux influences the minimum H-index, but that the maximum emission seems to be independent of the magnetic flux." +" This gives another indirect argument that photospheric fields increase the chromospheric emission. but do not actually deliver the main contribution to it,"," This gives another indirect argument that photospheric fields increase the chromospheric emission, but do not actually deliver the main contribution to it." + For the strongest concentration of magnetic flux in. the field of view observed. a stable long-lasting (>1 hr) network element. we find a displacement of 1” between magnetic flux and highest emission. and Jess pronounced photospheric power.," For the strongest concentration of magnetic flux in the field of view observed, a stable long-lasting $>1$ hr) network element, we find a displacement of $^{\prime\prime}$ between magnetic flux and highest emission, and less pronounced photospheric power." + The displacement ts only in one direction along the slit. which we ascribe to the field topology in the FOV.," The displacement is only in one direction along the slit, which we ascribe to the field topology in the FOV." + The strongest field concentration could be connected to one of opposite polarity nearby in the direction of the displacement., The strongest field concentration could be connected to one of opposite polarity nearby in the direction of the displacement. + The chromospheric intensity oscillations show power at all frequencies from 0 to about 10 mHz., The chromospheric intensity oscillations show power at all frequencies from 0 to about 10 mHz. + We do not find a pronounced peak of power at 3 minutes. but a broad distribution over several frequencies.," We do not find a pronounced peak of power at 3 minutes, but a broad distribution over several frequencies." + However. to address the question of heating. the average power spectrum alone is of less interest than the power spectrum of locations with strong chromospheric emission.," However, to address the question of heating, the average power spectrum alone is of less interest than the power spectrum of locations with strong chromospheric emission." + Comparing the spatially resolved chromospheric intensity with chromospheric and photospheric oscillation power. or with the locations of photospheric fields. it can be seen that strong emission in the chromosphere is always related to one of two things (or both): magnetic fields. or high power in the photospheric velocity oscillations.," Comparing the spatially resolved chromospheric intensity with chromospheric and photospheric oscillation power, or with the locations of photospheric fields, it can be seen that strong emission in the chromosphere is always related to one of two things (or both): magnetic fields, or high power in the photospheric velocity oscillations." + These photospheric oscillations are due to isolated small-scale power sources in the frequency range up to the acoustic cutoff frequency of around 5 mHz (cf., These photospheric oscillations are due to isolated small-scale power sources in the frequency range up to the acoustic cutoff frequency of around 5 mHz (cf. + Fig. 7..," Fig. \ref{fig9}," + lowermost panel)., lowermost panel). + This agrees with the finding of ? that the large-scale photospheric 3 mHz oscillations are less important for the generation of H»y bright grains than localized 5 mHz oscillations (seealso?).., This agrees with the finding of \citet{kamio+kurokawa2006} that the large-scale photospheric 3 mHz oscillations are less important for the generation of $_{\rm 2V}$ bright grains than localized 5 mHz oscillations \citep[see also][]{hoekzema+etal2002}. + The result would also be in agreement with both an impulsive excitation of waves. or a stochastic generation by the (random) superposition of large-scale wave patterns. which again would interfere positively only on some locations.," The result would also be in agreement with both an impulsive excitation of waves, or a stochastic generation by the (random) superposition of large-scale wave patterns, which again would interfere positively only on some locations." + The analysis of the phase differences between the oscillations of the Hay peak and the intensities at other wavelengths gives evidence that the acting agent between photosphere and chromosphere are propagating waves with frequencies above 2 mHz.," The analysis of the phase differences between the oscillations of the $_{\rm + 2V}$ peak and the intensities at other wavelengths gives evidence that the acting agent between photosphere and chromosphere are propagating waves with frequencies above 2 mHz." + Below 2 mHz. constant phase shifts," Below 2 mHz, constant phase shifts" +fast timing of black hole candidates in their low states and atoll sources in hard island states.,fast timing of black hole candidates in their low states and atoll sources in hard island states. +a given line is likely to be blended to some extent. resulting in additional uncertainty in an EW analvsis.,"a given line is likely to be blended to some extent, resulting in additional uncertainty in an EW analysis." + We therefore used (he current version of the LTE line analysis program MOOG (Sneclen1973) to fit à svnthetie spectrum toa 10 wavelength region wound the 5782.14 Cul transition., We therefore used the current version of the LTE line analysis program MOOG \citep{moog} to fit a synthetic spectrum to a 10 wavelength region around the 5782.14 Cu I transition. + For those clusters lor which we had sulficient wavelength coverage. a LO spectrum svnthesis was also emploved to derive an abundance from the Cu I line at 5105.5À.," For those clusters for which we had sufficient wavelength coverage, a 10 spectrum synthesis was also employed to derive an abundance from the Cu I line at 5105.5." +. The model atiosphleres and parameters used in the spectrum svntliesis are identical to (hose derived by the Calilornia/Texas eroup., The model atmospheres and parameters used in the spectrum synthesis are identical to those derived by the California/Texas group. + We have adopted those previously determined abundances. particularly Mg and C. since the 5105 region contains a laree number of Mell and Cs lines in the cooler and higher metallicity stars.," We have adopted those previously determined abundances, particularly Mg and C, since the 5105 region contains a large number of MgH and $_2$ lines in the cooler and higher metallicity stars." + Though these molecular lines have been (reated as a simple parameter of fit. in only one case has it been necessary io alter either of these abundances.," Though these molecular lines have been treated as a simple parameter of fit, in only one case has it been necessary to alter either of these abundances." + In the M4 star L2208 the Mg abundance derived from atomic Me lines by. Ivausetal.(1999). did not seem to reproduce the strong Mell features surrounding the 5105 line., In the M4 star L2208 the Mg abundance derived from atomic Mg lines by \citet{Ivans1999} did not seem to reproduce the strong MgH features surrounding the 5105 line. + In this ease. an increase of 0.3 dex (o Ivans et al," In this case, an increase of 0.3 dex to Ivans et al." +is Meg abundance better fit the observed features.,'s Mg abundance better fit the observed features. + This star was problematic in the Ivans el al., This star was problematic in the Ivans et al. + study (see (heir 844.2.2. for a discussion of the anomalies found in (his stars alpha and light odd-elements., study (see their 4.2.2 for a discussion of the anomalies found in this star's alpha and light odd-elements. + The data for this star are of lower resolution and lower S/N than most of the saaple)., The data for this star are of lower resolution and lower S/N than most of the sample). + However. since the Mg abundance serves here only as a parameter of the spectrum svnthesis fit. it does not affect the Cu abundance we derive.," However, since the Mg abundance serves here only as a parameter of the spectrum synthesis fit, it does not affect the Cu abundance we derive." + A line list was prepared bv using MOOG (Sneden1973) to fit lines taken [rom the πο(1993). line list ina region around each επι line to the solar spectrum., A line list was prepared by using MOOG \citep{moog} to fit lines taken from the \citet{KurCD1993} line list in a region around each Cu line to the \citet{Kurucz1984} solar spectrum. + While (he fit was in general quite good with the initial line parameters. gf values For a few Fe lines were modified empirically to produce a better fit.," While the fit was in general quite good with the initial line parameters, $gf$ values for a few Fe lines were modified empirically to produce a better fit." + The Cu I gf , The Cu I $gf$ +Under the assumption that the X-ray spectrum of Cyg X-I is not substantially more complex than what has already been found in theChandra observation. eV-level spectral resolution observations of the system (as those that will be carried out by the satellite) can be used to constrain the dynamics of the disk and establish whether the accretion disk of Cygnus X-] actually precesses. and if so. to determine its dynamics.,"Under the assumption that the X-ray spectrum of Cyg X-1 is not substantially more complex than what has already been found in the observation, eV-level spectral resolution observations of the system (as those that will be carried out by the satellite) can be used to constrain the dynamics of the disk and establish whether the accretion disk of Cygnus X-1 actually precesses, and if so, to determine its dynamics." + Several X-ray binaries present periodic behavior in their light curves on timescales longer than the orbital period., Several X-ray binaries present periodic behavior in their light curves on timescales longer than the orbital period. + Among these systems we can mention Her X-1. SS 433. and LMC X-4.," Among these systems we can mention Her X-1, SS 433, and LMC X-4." + It has been suggested that these long periods correspond to the precession of the accretion disk (e.g. Katz 1973)., It has been suggested that these long periods correspond to the precession of the accretion disk (e.g. Katz 1973). + In the case of SS 433 the precession is directly measured in the jets. so if these are attached to the accretion disk it is reasonable to expect that the disk will also display precession (Katz 1980).," In the case of SS 433 the precession is directly measured in the jets, so if these are attached to the accretion disk it is reasonable to expect that the disk will also display precession (Katz 1980)." + Although there is no reported compelling evidence yet for disk precession in black hole binaries. it is reasonable that the same mechanisms responsible for this phenomenon in neutron binaries will apply.," Although there is no reported compelling evidence yet for disk precession in black hole binaries, it is reasonable that the same mechanisms responsible for this phenomenon in neutron binaries will apply." + The mechanism that produces the precession might be the instability of the response of the disc to the radiation reaction force from the illumination by the central source (e.g. Wijers Pringle 1999. Ogilvie Dubus 2001). or the tidal force of the companion star on a disk which is not coplanar with the binary orbit (Katz 1973. Larwood 1998. Kaufman Bernadó et al.," The mechanism that produces the precession might be the instability of the response of the disc to the radiation reaction force from the illumination by the central source (e.g, Wijers Pringle 1999, Ogilvie Dubus 2001), or the tidal force of the companion star on a disk which is not coplanar with the binary orbit (Katz 1973, Larwood 1998, Kaufman Bernadó et al." + 2002)., 2002). + Uniform disk precession will occur in this case only if the sound crossing time through the disk is considerably shorter than the characteristic precession period induced by the perturbing star., Uniform disk precession will occur in this case only if the sound crossing time through the disk is considerably shorter than the characteristic precession period induced by the perturbing star. + The precession angular velocity is given by (e.g. Romero et al., The precession angular velocity is given by (e.g. Romero et al. +" 2000): 1ο)5:Om-cos. where G is the gravitational constant. 7j, 1s the orbital radius. wy Is the inner disk angular velocity. ϐ is the half-opening angle of the precession cone. and 7 is the mass of the star that exerts the torque upon the disk."," 2000): $ \left|\Omega_{\rm p}\right|\approx \frac{3}{4} +\frac{Gm}{r_{\rm m}^3} \frac{1}{\omega_{\rm d}} \cos\theta, $ where $G$ is the gravitational constant, $r_{\rm m}$ is the orbital radius, $\omega_{\rm d}$ is the inner disk angular velocity, $\theta$ is the half-opening angle of the precession cone, and $m$ is the mass of the star that exerts the torque upon the disk." + The orbital period 7; 1s related with the involved masses and the size of the orbit by Kepler's law: Fa.30ΗMY where M ∣↼is the mass ofκ the accreting. object.," The orbital period $T_{\rm m}$ is related with the involved masses and the size of the orbit by Kepler's law: $ r_{\rm +m}^3=\frac{G(m+M)T_{\rm m}^2}{4\pi^2}, $ where $M$ is the mass of the accreting object." +". The ratio between the orbital and the precessing periods can be related through the disk angular velocity wy—(GM/rjyey where 7,= and #2ry/r,."," The ratio between the orbital and the precessing periods can be related through the disk angular velocity $\omega_{\rm d}=(GM/r_{\rm d}^3)^{1/2}$: where $T_{\rm p}=2\pi/\Omega_{\rm p}$ and $\kappa=r_{\rm d}/r_{\rm +m}$." +" Since &«I. normally Ti/T,<1."," Since $\kappa<1$, normally $T_{\rm m}/T_{\rm p}<1$." + In the 27/Qp)case of Cygnus X-I. Brocksopp et al. (," In the case of Cygnus X-1, Brocksopp et al. (" +1999) have reported multrwavelength evidence for the existence of a period of 142.0+7.1 days.,1999) have reported multiwavelength evidence for the existence of a period of $142.0\pm7.1$ days. + Similar precessing periods have been caleulated by Larwood (1998). Katz (1973. 1980) and Katz et al. (," Similar precessing periods have been calculated by Larwood (1998), Katz (1973, 1980) and Katz et al. (" +1982) for other X-ray binaries on the basis of the same model.,1982) for other X-ray binaries on the basis of the same model. + As shown by Romero et al. (, As shown by Romero et al. ( +2002) in the case of Cygnus X-]. such a period can be obtained from tidally-induced precession for an accretion disk with a size ~4«10!! cm. If the half opening angle of the precession cone is ~157.,"2002) in the case of Cygnus X-1, such a period can be obtained from tidally-induced precession for an accretion disk with a size $\sim 4\times 10^{11}$ cm, if the half opening angle of the precession cone is $\sim 15^{\circ}$." + For a purely wind-fed system. this size might be too large and other mechanism may be in operation to generate the observed timescales.," For a purely wind-fed system, this size might be too large and other mechanism may be in operation to generate the observed timescales." + In particular. radiation-driven precession (Pringle 1996. Maloney Begelman 1997. Ogilvie Dubus 2001). wind-driven warping and precession (Schandl Meyer 1994. Quillen 2001). and spin-spin precession (Bardeen Petterson 1975. Armitage Natarajan 1999) can yield precession periods of several weeks to a few months.," In particular, radiation-driven precession (Pringle 1996, Maloney Begelman 1997, Ogilvie Dubus 2001), wind-driven warping and precession (Schandl Meyer 1994, Quillen 2001), and spin-spin precession (Bardeen Petterson 1975, Armitage Natarajan 1999) can yield precession periods of several weeks to a few months." + In all these mechanisms the observed optical and X-ray modulation points to precession of the disk. whereas the radio variations might originate in the jet.," In all these mechanisms the observed optical and X-ray modulation points to precession of the disk, whereas the radio variations might originate in the jet." + The use of emission lines as a diagnosis tool for the state of binary systems has been proposed in the past (e.g.. for an investigation on supermassive black hole binarity. see Torres et al.," The use of emission lines as a diagnosis tool for the state of binary systems has been proposed in the past (e.g., for an investigation on supermassive black hole binarity, see Torres et al." + 2003. also Gaskell 2003: Zakharov et al.," 2003, also Gaskell 2003; Zakharov et al." + 2004a.b).," 2004a,b)." + Here we show that a similar method can be used to extract information about the precession status of microquasars., Here we show that a similar method can be used to extract information about the precession status of microquasars. + For the case of Cygnus X-1. we shall assume that the time-averaged disk inclination angle is 357. which is in agreement with the fitting of the system's Fe line (Miller et al.," For the case of Cygnus X-1, we shall assume that the time-averaged disk inclination angle is $^\circ$, which is in agreement with the fitting of the system's Fe line (Miller et al." + 2002)., 2002). + Values around this time-averaged inclination angle were also found for other binary systems ( e.g. Her X-1. LMC X-4. SMC X-I. etc.," Values around this time-averaged inclination angle were also found for other binary systems ( e.g., Her X-1, LMC X-4, SMC X-1, etc." + e.g.. as discussed by Larwood 1998).," e.g., as discussed by Larwood 1998)." +" We also assume two extreme cases for the amplitude of the disk precession: (1) the disk inclination angle precesses from to 39"" (very low magnitude of the precession angle): (2) the disk inclination angle precesses from 5° to 65° (large magnitude of the precession angle).", We also assume two extreme cases for the amplitude of the disk precession: (1) the disk inclination angle precesses from $^\circ$ to $^\circ$ (very low magnitude of the precession angle); (2) the disk inclination angle precesses from $^\circ$ to $^\circ$ (large magnitude of the precession angle). + The amplitude of the precession of the inner disk should not be large if it is due to the tidal force of the secondary star and if the disk (especially the outer disk region) develops a significant warp., The amplitude of the precession of the inner disk should not be large if it is due to the tidal force of the secondary star and if the disk (especially the outer disk region) develops a significant warp. + However. if the initial spin direction of the BH ts significantly different from the orbital angular momentum direction. the inner disk. which ts confined to the equatorial plane of the BH (if the spin is high) due to the Bardeen-Peterson effect. may precess around the total angular momentum (dominated by the orbital angular momentum) with an amplitude as large as the initial orbital inclination angle with respect to the BH equatorial plane.," However, if the initial spin direction of the BH is significantly different from the orbital angular momentum direction, the inner disk, which is confined to the equatorial plane of the BH (if the spin is high) due to the Bardeen-Peterson effect, may precess around the total angular momentum (dominated by the orbital angular momentum) with an amplitude as large as the initial orbital inclination angle with respect to the BH equatorial plane." + If we assume that the disk is rigidly precessing around the total angular momentum (dominated by the orbital angular momentum). then the amplitude of the precession may also be around 30”.," If we assume that the disk is rigidly precessing around the total angular momentum (dominated by the orbital angular momentum), then the amplitude of the precession may also be around $^\circ$." + Several calculations on the disk line profiles have been performed., Several calculations on the disk line profiles have been performed. +" We use a ray-tracing technique and elliptic integrals (Rauch Blandford 1994; see also Yu Lu 2000: Lu Yu 2001 and references therein) to follow the trajectories of photons from the observer. keeping track of all coordinates until the photons either intersect the accretion disk plane. disappear below the event horizon. or escape to ""infinity"" (operationally defined to be r21000GM/c7 away from the BH)."," We use a ray-tracing technique and elliptic integrals (Rauch Blandford 1994; see also Yu Lu 2000; Lu Yu 2001 and references therein) to follow the trajectories of photons from the observer, keeping track of all coordinates until the photons either intersect the accretion disk plane, disappear below the event horizon, or escape to ""infinity"" (operationally defined to be $r=1000GM/c^2$ away from the BH)." + We then calculate the redshift factor for a photon (to the observer) emitted from a particular position on the disk., We then calculate the redshift factor for a photon (to the observer) emitted from a particular position on the disk. + The solid angle subtended at the observer by each disk element ts also calculated., The solid angle subtended at the observer by each disk element is also calculated. +" We set the inner radius of the disk to be at the marginally stable orbit (67. for a Schwarzschild black hole or 1.237. for a Kerr black hole with spin ¢/M—0.998. where r,=GM c). and the outer radius at 1607..."," We set the inner radius of the disk to be at the marginally stable orbit $6r_{\rm g}$ for a Schwarzschild black hole or $1.23r_{\rm g}$ for a Kerr black hole with spin $a/M=0.998$, where $r_{\rm g}=GM/c^2$ ), and the outer radius at $160r_{\rm g}$." + We assume that the surface emissivity of line photons follows a power-law. 7%. with g=2.5.," We assume that the surface emissivity of line photons follows a power-law, $r^{-q}$, with $q=2.5$." + Both the power- emissivity law and the size of the disk in Schwarszchild units. are usual assumptions (see. e.g.. Nandra et al.," Both the power-law emissivity law and the size of the disk in Schwarszchild units, are usual assumptions (see, e.g., Nandra et al." + 1997)., 1997). + The BH spin is assumed to be a/M=0.998., The BH spin is assumed to be $a/M=0.998$. + In microquasar systems. both the high frequency quasi-periodic oscillation and relativistic lines suggest a high spin.," In microquasar systems, both the high frequency quasi-periodic oscillation and relativistic lines suggest a high spin." + In any case. we proved that if We were to assume a lower spin. there is not much qualitative difference for the problem we have studied here.," In any case, we proved that if we were to assume a lower spin, there is not much qualitative difference for the problem we have studied here." + With the above assumptions. we sum up all the photons received by the observer. which is emitted from each disk element. and obtain the profile of emergent Fe Ko lines. with different inclination," With the above assumptions, we sum up all the photons received by the observer, which is emitted from each disk element, and obtain the profile of emergent Fe $\alpha$ lines, with different inclination" +"deduce that €,=ἐν and e;=ez; taking into account the symmetry of the equations.",deduce that $C_h=C_v$ and ${\sigma}^2_h = {\sigma}^2_v$ taking into account the symmetry of the equations. +" Moreover. the temperature power spectrum P(À) can be obtained. from the.defait wavelet power spectrum wa.(dA) as follows For the Haar and Mexican wavelets we can calculate: where 47=kp|bs and D. can be obtained from Py. swapping A, and As."," Moreover, the temperature power spectrum $P(k)$ can be obtained from the wavelet power spectrum $w_{\alpha \alpha }(R,R;\vec k)$ as follows For the Haar and Mexican wavelets we can calculate: where $k^2 = k_1^2 + k_2^2$ and $\tilde{\Gamma}_v$ can be obtained from $\tilde{\Gamma}_h$, swapping $k_1$ and $k_2$." + The variance of thedefait. wavelet coellicients for the Haar and Mexican Hat systems. assuming the standard CDAL model. is presented in Figure 1.," The variance of the wavelet coefficients for the Haar and Mexican Hat systems, assuming the standard CDM model, is presented in Figure 1." + As one can see the acoustic peaks can be clearly noticed. being more pronounced for the Mexican Hat. basis.," As one can see the acoustic peaks can be clearly noticed, being more pronounced for the Mexican Hat basis." + This last result is a consequence of being a more localizecl wavelet. svstem., This last result is a consequence of being a more localized wavelet system. + For a more cletailed discussion see Sanz et al., For a more detailed discussion see Sanz et al. + 1905. 1999.," 1998, 1999." + An orthonormal basis of L(30) dilationscan be constructed. from a wavelet co through cvaclic j and translations & 1n addition. a scaling function © can be defined associate to themother wavelet c.," An orthonormal basis of $L^2(\Re )$ can be constructed from a wavelet $\psi$ through dyadic dilations $j$ and translations $k$ In addition, a scaling function $\phi$ can be defined associated to the wavelet $\psi $." + Such a function gives rise to the sO called multiresolution analysis., Such a function gives rise to the so called multiresolution analysis. + multiresolution analysis of L?(8) is defined as a sequence of elosed subspaces V; of ΙΩΝ). j€Z.," A multiresolution analysis of $L^2(\Re )$ is defined as a sequence of closed subspaces $V_j$ of $L^2(\Re )$, $j\in Z$." + Properties can be seen in Ogden (1997), Properties can be seen in Ogden (1997). +" Subspaces V; are generated by dyadic dilations anc translations of the scaling function o (this function forms an orthonormal basis of V5. fO,n6r)=olr 21)."," Subspaces $V_j$ are generated by dyadic dilations and translations of the scaling function $\phi $ (this function forms an orthonormal basis of $V_o$, $\{ \phi_{o,k}(x)=\phi(x-k)\}$ )." +" Aloreover cach Vj can be expressed as the orthogona sum V;—1,1Wya. where Mj4, ds created. from wavelets ojia."," Moreover each $V_j$ can be expressed as the orthogonal sum $V_j=V_{j-1}\oplus W_{j-1}$, where $W_{j-1}$ is created from wavelets $\psi_{j-1,k}$." + Taking into account the properties of the sealing function. together with this last) expression. we can construct at increasing levels of resolution.," Taking into account the properties of the scaling function, together with this last expression, we can construct at increasing levels of resolution." + These are lincar combinations of dilations and translations of a scaling function ©., These are linear combinations of dilations and translations of a scaling function $\phi $. + The dilference between two consecutiveepprozàmealions. Lo. thedelail at the corresponding resolution level. is given by a linear combination of dilations and translations of a wavelet function c.," The difference between two consecutive, i.e. the at the corresponding resolution level, is given by a linear combination of dilations and translations of a wavelet function $\psi $." + The analysis performed in this work assumes equal dilations in the 2 dimensions involved., The analysis performed in this work assumes equal dilations in the 2 dimensions involved. +" At a fixed level of resolution. subspaces in a 2-1) multiresolution analysis are the tensor products of the corresponding one-dimensional ones Vj,=V;acας νο 2-D basis is therefore built by the product of two sealing functions. (epprorimnalion). the. product of wavelet anc scaling functions (horizontal and vertical defails) and the product of two wavelets (diagonal defaits): Horizontal. vertical and clagonaldefeif coellicients represent the variations in these directions relative to a weighted average ata lower resolution level (given by theapprorimalion coellicients)."," At a fixed level of resolution, subspaces in a 2-D multiresolution analysis are the tensor products of the corresponding one-dimensional ones ${\bf V}_{j+1}=V_{j+1}\otimes V_{j+1}$ .The 2-D basis is therefore built by the product of two scaling functions ), the product of wavelet and scaling functions (horizontal and vertical ) and the product of two wavelets (diagonal ): Horizontal, vertical and diagonal coefficients represent the variations in these directions relative to a weighted average at a lower resolution level (given by the coefficients)." +" A discrete orthonormal basis. E,Gr:j.&). can be defined by setting /?—2/ and b—2""Kk in equations (3-5). then (LGPpΟΙOu)=Mad)Opp. where () denotes the scalar ne in £709)."," A discrete orthonormal basis, ${\Gamma}_{\alpha}(\vec{x}; j, \vec{k})$, can be defined by setting $R = 2^{-j}$ and $\vec{b} = 2^{-j}\vec{k}$ in equations (3-5), then $({\Gamma}_{\alpha}(\vec{x}; j, \vec{k}){\Gamma}_{{\alpha}^{\prime}}(\vec{x}; +j^{\prime}, {\vec{k}}^{\prime})) = {\delta}_{\alpha +{\alpha}^{\prime}}{\delta}_{jj^{\prime}}{\delta}_{\vec{k}{\vec{k}}^{\prime}}$, where $()$ denotes the scalar product in $L^2(\Re ^2)$." + Wo we define the discrete wavelet coefficients associated to anydefaif by the equation we can thus reconstruct the image with all the In particular. we get the following expression for the moment of the image that expresses how the energy of the field is. distributed locally at any scale ancldefe.," If we define the discrete wavelet coefficients associated to any by the equation (6) we can thus reconstruct the image with all the In particular, we get the following expression for the second-order moment of the image that expresses how the energy of the field is distributed locally at any scale and." +" Fora finite image. £2,5Rovere. in order to reconstruct it we must add. to equation (19) an. approximation wsGOL,CEE) with LEB)=OrRiehi) PCrRiehe) and ws(he)={απ)Ε.Ε). representing the Ποιά at the lower resolution."," For a finite image, $R_{max}\times R_{max}$, in order to reconstruct it we must add to equation (19) an approximation $w_a(\vec{k}){\Gamma}_a(\vec{x}; \vec{k})$ with ${\Gamma}_a(\vec{x}; \vec{k})\equiv \Phi (x_1; R_{max},k_1)$ $\Phi (x_2; R_{max},k_2)$ and $w_a(\vec{k}) \equiv \int +d\vec{x}\,f(\vec{x})\,{\Gamma}_a(\vec{x},\vec k)$, representing the field at the lower resolution." +" (9) represents the temperature fluctuation field then the variance is given by =(GNT/T))2NS. being AN, the number of pixels."," If $f(\vec x)$ represents the temperature fluctuation field then the variance is given by $<(\Delta T/T)^2>=((\Delta T/T)^2)/N_p$, being $N_p$ the number of pixels." + The orthonormal basis that we are going to use are the standard. Daubechies IN. (Llaar corresponds to N= 1). that has been extensively used in the literature because of their special properties: they are defined in a compact support. have increasing with No and vanishing moments up to order ΑΝ velNN(Daubechies 1988).," The orthonormal basis that we are going to use are the standard Daubechies $N$ (Haar corresponds to $N = 1$ ), that has been extensively used in the literature because of their special properties: they are defined in a compact support, have increasing regularity with $N$ and vanishing moments up to order $N - 1$ (Daubechies 1988)." + On the contrary. the Alexican Lat is not defined in a compact support and it is not appropriate for this multiresolution analysis.," On the contrary, the Mexican Hat wavelet is not defined in a compact support and it is not appropriate for this multiresolution analysis." + For discrete. wavelet. analysis of the CMD we use the Matlab Wavelet ToolboxCMisiti et al., For discrete wavelet analysis of the CMB maps we use the Matlab Wavelet Toolbox (Misiti et al. + 1996)., 1996). +" ---""a""Ehis toolbox is an extensive collection of programs for denoising and compressing signals and 2-D images."," This toolbox is an extensive collection of programs for analyzing, denoising and compressing signals and 2-D images." + Discrete Wavelet decomposition is performed as described above to, Discrete Wavelet decomposition is performed as described above to +parameters to which the BBIL background is sensitive.,parameters to which the BBH background is sensitive. +" Our results show that for M.$ ) and provides the best description of the source spectrum $\chi^{2}$ =255/259)." +" ""Phe amount of extra column density is Ng. —4 107 cm7 while the value of οποίο index is P=1.56: the eross calibration is found to be 1.65023 indicating possible variability between the pointed XMM observation and the IBIS average measurement.", The amount of extra column density is $_{H}$ $\sim$ 4 $\times$ $^{21}$ $^{-2}$ while the value of photon index is $\Gamma$ =1.56; the cross calibration is found to be $^{+0.28}_{-0.22}$ indicating possible variability between the pointed XMM observation and the IBIS average measurement. +" The combined «ΧΑΛΙΛΗΝΓΙΟΛΙ, unfolded spectrum. fitted with this mocde is shown in figure 10 and. described. in table 3.", The combined XMM/INTEGRAL unfolded spectrum fitted with this model is shown in figure 10 and described in table 3. + The spectral. parameters are quite in agreement with what found by Halpern from the NMM data analysis alone but not fully compatible with the spectral parameters reported. by Foleheraiter ct al. (, The spectral parameters are quite in agreement with what found by Halpern from the XMM data analysis alone but not fully compatible with the spectral parameters reported by Folgheraiter et al. ( +1997). although we ect a similar flux in the 0.5-2 keV energy band (~2 £ 7? ο].,1997) although we get a similar flux in the 0.5-2 keV energy band $\sim$ 2 $\times$ $^{-12}$ $^{-2}$ $^{-1}$ ). +" On the other hand. it is worth noting that a thermal mocel does not fit our broad band We conclude that LGA 18538-0102. as already. argued bv Llalpern CGotthell (2010). is unlikely a compact object in the supernova remnant. (92.109 but. could. be a backeround ACN that is coincicentally aligned with the supernova,"," On the other hand, it is worth noting that a thermal model does not fit our broad band We conclude that IGR J18538-0102, as already argued by Halpern Gotthelf (2010), is unlikely a compact object in the supernova remnant G32.1-09 but could be a background AGN that is coincidentally aligned with the supernova." + In this work we have cross-correlated the list of the still unidentified hard N-ray emitters [listed in the 4th. LBIS survey with the archive of all pointings finding a set of 6 objects with archivalLI data., In this work we have cross-correlated the list of the still unidentified hard X-ray emitters listed in the 4th IBIS survey with the archive of all pointings finding a set of 6 objects with archival data. + First. we studied the EPIC images in order to [find in the LBIS error circle the X-ray counterpart(s).," First, we studied the EPIC images in order to find in the IBIS error circle the X-ray counterpart(s)." + In the case where an associated source has been found. the data have then been used together the spectra to study the broad. band slope and investigate the possible nature of the source.," In the case where an associated source has been found, the data have then been used together the spectra to study the broad band slope and investigate the possible nature of the source." + In table 4 à summary of our proposed identifications is In a couple of cases no obvious X-ray counterpart has been found. from the observations. like IGI J1173331.2406 and LOR J17445-2747.," In table 4 a summary of our proposed identifications is In a couple of cases no obvious X-ray counterpart has been found from the observations, like IGR J173331–2406 and IGR J17445-2747." + In the first case. no X-ray source has been detected.," In the first case, no X-ray source has been detected." + This is in perfect agreement with the LBIS survey data where this source has been found to be transient., This is in perfect agreement with the IBIS survey data where this source has been found to be transient. + Extrapolating to the low energies (0.5-10 keV) the spectrum seen by LBIS during the source outburst and comparing it with the NM. upper limit. we found a dynamical range of the order of 3000.," Extrapolating to the low energies (0.5-10 keV) the spectrum seen by IBIS during the source outburst and comparing it with the XMM upper limit, we found a dynamical range of the order of 3000." + Such a high. value strongly suggests that LGR: J17331.2406 could. be. either a transient. black hole in the Galactic bulge or a SEXT although this latter interpretation can be ruled out clue to source location olf the Galactic plane (~ 5r. degrees) as well as its significantly longer. outburst duration. compared. to classical SENXTs., Such a high value strongly suggests that IGR J17331–2406 could be either a transient black hole in the Galactic bulge or a SFXT although this latter interpretation can be ruled out due to source location off the Galactic plane $\sim$ 5 degrees) as well as its significantly longer outburst duration compared to classical SFXTs. + Vhe other case where the NMM observation: does not provide a secure X-ray counterpart is that of LOR J17445-2747., The other case where the XMM observation does not provide a secure X-ray counterpart is that of IGR J17445-2747. + From the imaging analysis we found a faint NAIA source at the border of the IBIS error circle which has also been detected by both ονΧΙΙ and. Chandra., From the imaging analysis we found a faint XMM source at the border of the IBIS error circle which has also been detected by both Swift-XRT and Chandra. + On the other hand. from archival searches we found. an NAIA slew source well inside. the high energy positional uncertainty which has been seen only once by NMM and therefore is again in perfect agreement with the high energy survey data which classified this source extremely variable: we therefore associate LGR J17445-214T 1o NAIAISLI J174429.4-274600 In the remaining four cases we have found a convincing X-rav counterpart in the IBIS error circle for which it has eni possible to search for counterparts in other wavelength xnds and also perform. the spectral data. analysis. in he 0.5110 keV band.," On the other hand, from archival searches we found an XMM slew source well inside the high energy positional uncertainty which has been seen only once by XMM and therefore is again in perfect agreement with the high energy survey data which classified this source extremely variable; we therefore associate IGR J17445-2747 to XMMSL1 J174429.4-274609 In the remaining four cases we have found a convincing X-ray counterpart in the IBIS error circle for which it has been possible to search for counterparts in other wavelength bands and also perform the spectral data analysis in the 0.5–110 keV band." + The spectral parameters obtained ogether with the possible Ht/optical/radio counterpart ound allowec us to investigate on the nature of each source., The spectral parameters obtained together with the possible IR/optical/radio counterpart found allowed us to investigate on the nature of each source. + We conclude that LGR. J15859-5750 is an AGN of intermediate tvpe for which we were able to estimate he amount of the complex. absorption as well as to give constraints on the reflection ancl t1e high. energy. eut-olT., We conclude that IGR J15359-5750 is an AGN of intermediate type for which we were able to estimate the amount of the complex absorption as well as to give constraints on the reflection and the high energy cut-off. + Lor the two ASCA sources AX J1739.3-2923 and ΑΝ J1740.2-2003. we suggest à stronely absorbed. Galactic nature and for both we argue that trev are likely persistent LIAIND systems., For the two ASCA sources AX J1739.3-2923 and AX J1740.2-2903 we suggest a strongly absorbed Galactic nature and for both we argue that they are likely persistent HMXB systems. + More uncertain is the case of LGR. JIS53s-N02 which is spatially coincident with a hot spot. in he supernova remnant (29.109 detected: previously by LOSATL and ASCA., More uncertain is the case of IGR J18538-0102 which is spatially coincident with a hot spot in the supernova remnant G32.1-0.9 detected previously by ROSAT and ASCA. + From the broad band spectral analysis »rformed in this work we can conclude that this object is unrelated: compact object. which happen to coincide with he supernova remnant. probably a background AGN that is coincidentally aligned even if no radio counterpart has ovn found in the more precise NMM error This research has made use of cata obtained from. the SIMDAD database operated at CDS. Strasbourg. France: the Ligh Enerey Astrophysics Science. Archive Research Center (IUEASARC). provided by NASA's Goddard Space Εισαι Center NASA/IPAC Extragalactic Database (NED).," From the broad band spectral analysis performed in this work we can conclude that this object is unrelated compact object which happen to coincide with the supernova remnant, probably a background AGN that is coincidentally aligned even if no radio counterpart has been found in the more precise XMM error This research has made use of data obtained from the SIMBAD database operated at CDS, Strasbourg, France; the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center NASA/IPAC Extragalactic Database (NED)." + We thank the anonymous referee for the very. detailed. and careful review of this paper., We thank the anonymous referee for the very detailed and careful review of this paper. + The authors acknowledge the ASI financial support via ASLINAV erant L/008/07/0., The authors acknowledge the ASI financial support via ASI–INAF grant I/008/07/0. +410% while Larsenοἱal.(2011) claim up to heterogeneity in the initial abundance of 7 A] in the inner solar svstem).,$\pm$ while \citet{lar11} claim up to heterogeneity in the initial abundance of $^{26}$ Al in the inner solar system). + Ho we assume instead that the material from. which Iris formed had the lowest (7 AL/7 Al); in chondrule-containing parent bodies (from (2011))). the ordinary chondrite parent body. ο ΑΙ y=l.63x 10°. then Iris crystallized ab least 1.8 Myr alter CAIs.," If we assume instead that the material from which Iris formed had the lowest $^{26}$ $^{27}$ $_0$ in chondrule-containing parent bodies (from \citet{lar11}) ), the ordinary chondrite parent body, $^{26}$ $^{27}$ $_0$ $\times$ $^{-5}$, then Iris crystallized at least 1.8 Myr after CAIs." + The formation time of Iris compared to ehondrules from other meteorites. as well as a CAI-like Stardust fragment. is shown in Figure 4..," The formation time of Iris compared to chondrules from other meteorites, as well as a CAI-like Stardust fragment, is shown in Figure \ref{almg}." + The probability density curve for Iris is derived from the Monte Carlo simulations discussed above. whereas the curves lor other objects are normalized suns of reported measurements of AD). assuming Gaussian- errors.," The probability density curve for Iris is derived from the Monte Carlo simulations discussed above, whereas the curves for other objects are normalized sums of reported measurements of $^{26}$ $^{27}$ $_0$, assuming Gaussian-distributed errors." + Probability density associated with negative CAD/7* AI) corresponds to ill-clefined Gime since CALI formation. so these densities. as well as (those for 77 half-lives of Al. are not shown. (," Probability density associated with negative $^{26}$ $^{27}$ $_0$ corresponds to ill-defined time since CAI formation, so these densities, as well as those for $>$ 7 half-lives of $^{26}$ Al, are not shown. (" +Consequently. the area under the curves for objects with allowed negalive or very small (ΑΙ Αμ is less than for objects that formed earlier.),"Consequently, the area under the curves for objects with allowed negative or very small $^{26}$ $^{27}$ $_0$ is less than for objects that formed earlier.)" + Since CAIs are believed to have formed when (he Sun was a class 0 or class I protostar. lis formed late in the evolution of the solu nebula. at a time when ~95% of the scattered disk is thought to have eleared (Iersant.Gautheir.&IIwé2001).," Since CAIs are believed to have formed when the Sun was a class 0 or class I protostar, Iris formed late in the evolution of the solar nebula, at a time when $\sim$ of the scattered disk is thought to have cleared \citep{her01}." +". The residual disk material mav have coated Jupiter-lamily comets with a ""late veneer” enriched in inner solar system malerial (Oglioreetal.2010).. ancl this may explain sienifiecant. differences between the Stardust sample. which sampled Wild 2 coma material ejected by jets that entrained near- dust (Bellon2010).. and chondiitie porous interplanetary dust. particles. which probably sample the bulk of Ixuiper Belt comets (Nesvornyetal.2010)."," The residual disk material may have coated Jupiter-family comets with a “late veneer” enriched in inner solar system material \citep{ogl10}, and this may explain significant differences between the Stardust sample, which sampled Wild 2 coma material ejected by jets that entrained near-surface dust \citep{bel10}, and chondritic porous interplanetary dust particles, which probably sample the bulk of Kuiper Belt comets \citep{nes10}." +. These analvses are [or a single object in the Stardust collection., These analyses are for a single object in the Stardust collection. + However. Iris is not unique or even particularly unusual in the suite of particles returned. [rom comet Wild 2.," However, Iris is not unique or even particularly unusual in the suite of particles returned from comet Wild 2." +" Nakamuraetal.(2008) identified four chondrule-like objects in the Stardust samples which are more ""O-enriched than Iris (Figure 2)) and olivines that are less Fe-rich (Forouj. Foos. Foo)."," \citet{nak08} identified four chondrule-like objects in the Stardust samples which are more $^{16}$ O-enriched than Iris (Figure \ref{oiso}) ) and olivines that are less Fe-rich $_{79\text{--}80}$, $_{95}$, $_{91}$ )." + At least one chondrule fragment. was also identified during the Stardust preliminary examination (Zolenskvetal.2006)., At least one chondrule fragment was also identified during the Stardust preliminary examination \citep{zol06}. +. All of these objects show igneous textures similar to Iris. though Iris appears to have formed in a more oxidizing environment.," All of these objects show igneous textures similar to Iris, though Iris appears to have formed in a more oxidizing environment." + Recent measurements bv Joswiaketal.(2011) show two fragments of Fe-rich olivine (Fogssz. Fossa). that also have O isotopic composition close to Iris.," Recent measurements by \citet{jos11} show two fragments of Fe-rich olivine $_{62\text{--}67}$, $_{58\text{--}61}$ ), that also have O isotopic composition close to Iris." + Chondrule-like objects from Stardust show a broad range of isotopic aud mineralogical compositions: a subset of these objects could be geneticallv related to Iris., Chondrule-like objects from Stardust show a broad range of isotopic and mineralogical compositions; a subset of these objects could be genetically related to Iris. +"The received photou density »""(e"") iu the jet comoving faune can be then approximately written as where we use e=(2D)?e42 (there are two Doppler shiftswp due to nurror effects).",The received photon density $\npp(\epspp)$ in the jet comoving frame can be then approximately written as where we use $\epspp=(2\Gam)^2\epsp$ (there are two Doppler shifts due to mirror effects). + The couvenieut form of cross section of photon-photou collision is (Coppi Blandford 1990) where à is the usual ó-fuuctiou., The convenient form of cross section of photon-photon collision is (Coppi Blandford 1990) where $\delta$ is the usual $\delta$ -function. + This approximation is ouly valid for case of isotropic radiation field., This approximation is only valid for case of isotropic radiation field. + The Rsv colmponent is seen bv the blob within the solid anele AQzc2z(lcosP1)x=x/I?., The Rsy component is seen by the blob within the solid angle $\Delta \Omega \approx 2\pi(1-\cos\Gam^{-1})\approx \pi/\Gam^2$. + Although the cross section of photon-plioton interaction holds. the interacting possibility among plotous reduces by a factor of AQ/lz=1/GD)? due to the beaming effects. which effectively reduces the opacity.," Although the cross section of photon-photon interaction holds, the interacting possibility among photons reduces by a factor of $\Delta \Omega/4\pi=1/(2\Gam)^2$ due to the beaming effects, which effectively reduces the opacity." + Thus the pair production optical depth for photon with euergv €» reads here 79.(e.) is We will show the validity of the above approximation in the next subsection., Thus the pair production optical depth for photon with energy $\epsg$ reads here $\tgg^0(\epsg)$ is We will show the validity of the above approximation in the next subsection. +" Supposing that RSC operates eficiently iu οταν loud blazars. we cau eet fi. from the observations From equation (11) we lave πμ... From equation (11) we kuow that the observed 7/4 represents the reflection ratio aud Doppler factor of jet notion as long as the Compton catastrophe does not occi,"," Supposing that RSC operates efficiently in $\gamma$ -ray loud blazars, we can get $\lrsc$ from the observations From equation (11) we have and From equation (11) we know that the observed $\lrsc$ represents the reflection ratio and Doppler factor of jet motion as long as the Compton catastrophe does not occur." + Tf we set fiemlans we got τνD2=(k01 for D10.," If we set $\lrsc\approx \lssc$, we get $\tblr\approx \Gam^{-2}=0.01$ for $\Gam=10$." + This value is the lowest one in the model of Sikora. Boechuan Rees (1991) who suggest τν—0.1~0.01.," This value is the lowest one in the model of Sikora, Begelman Rees (1994) who suggest $\tblr=0.1\sim 0.01$." +" Iu act we can roughly adopt 7,,, as the covering factor which is usually taken to be 0.1 in fitting the broad cussion line x photoionization moclel.", In fact we can roughly adopt $\tblr$ as the covering factor which is usually taken to be 0.1 in fitting the broad emission line by photoionization model. + Iuserting D aud Np [|eqs(1) and (60) intoτὸ. (Eq., Inserting $B$ and $N_T$ [eqs(1) and (6)] into$\tgg^0$ (Eq. +" 10). and letting e;=€,)./P aud 5= “-DAt,,,.. we have the pair xoduction optical depth for €; iu the observers frame where A,=DontSOnhie(DELEh-)~ ΕΞτοςτο4% for a—2.1)."," 10), and letting $\epsg=\eps_{\rm obs}/\cd$ and $s=c\cd \dt$ , we have the pair production optical depth for $\eps_{\rm obs}$ in the observer's frame where $K_{\alpha}=\frac{(3-\al)\st c}{80\pi h \nu_0^2} + \left(\frac{h}{m_ec^2}\right)^{\al-1 \over 2}$ $K_{\alpha}=7.9\times 10^{-18}$ for $\alpha=2.4$ )." +" There are five observational parameters: 14. Msc. a, Afus and fc: aud the uukuown Doppler factor D."," There are five observational parameters: $\nu_{\rm s}$ , $\nu_{\rm rsc}$, $\al$, $\dt$ and $\lrsc$; and the unknown Doppler factor $\cd$." +" For the typical value of pirzuneters; à=2.1 i,—10«10:TIz. and Me=LO.109 Πε. Δι=1 day, and D=10. we have where Diy=D/10."," For the typical value of parameters, $\al=2.4$, $\nu_{\rm s}=4.0\times 10^{14}$ Hz, and $\nu_{\rm rsc}=1.0\times 10^{25}$ Hz, $\dt=1$ day, and $\cd=10$, we have where $\cd_{10}=\cd/10$." + Fieure 1 shows the opacity due to pair production of photous with very high euergv chcountering with the reflected svuchrotron plotous., Figure 1 shows the opacity due to pair production of photons with very high energy encountering with the reflected synchrotron photons. + The equation (15) tells us the constraints on VILE from jet: (1) inaller fie. Le. stronger reflection. will leads to the absorption of TeV plotons.," The equation (15) tells us the constraints on VHE from jet: (1) smaller $\lrsc$, i.e. stronger reflection, will leads to the absorption of TeV photons." + This parameter represents the energv density reflected by the BLR cloud including the bulk relativistic motion., This parameter represents the energy density reflected by the BLR cloud including the bulk relativistic motion. + From this estimation we know that TeV photon will be absorbed by the reflected svuchrotron plotous provided that fp.<1.9. (, From this estimation we know that TeV photon will be absorbed by the reflected synchrotron photons provided that $\lrsc<1.9$. ( +"2) T. Is sensitive to n, and mie. (",2) $\tgg$ is sensitive to $\nu_{\rm s}$ and $\nu_{\rm rsc}$. ( +3) τ.- Is proportional to PD! in contrast to the usual down-linüt (sce Mattox et al 1993. and Doudi Cdiselliui 1995). providing the wpper μπιτ Doppler factor of bulk motion from 7...<1. This is a new coustraint. which is expressed bv the observational quantities.,"3) $\tgg$ is proportional to $\cd^4$, in contrast to the usual down-limit (see Mattox et al 1993, and Dondi Ghisellini 1995), providing the upper limit Doppler factor of bulk motion from $\tgg\leq 1$, This is a new constraint, which is expressed by the observational quantities." + It. leuds us a simple iud efficient way to select TeV candidates from kuown blazars iu teria of their known characteristics., It lends us a simple and efficient way to select TeV candidates from known blazars in term of their known characteristics. + The received photous reflected by BLR iu comoving fune is amisotropic. therefore. the pair opacity should be carefully treated.," The received photons reflected by BLR in comoving frame is anisotropic, therefore, the pair opacity should be carefully treated." + We lave made iportaut o»proxinatious that the BLR is thought to be a plane nürror and treated the photou-photon iuteractiou in au o»proxinate wav., We have made important approximations that the BLR is thought to be a plane mirror and treated the photon-photon interaction in an approximate way. + Now let us show the validity of this o»proxination., Now let us show the validity of this approximation. + We adopt the geometry shown iu Fig lc of Clüsellin Macau (1996)., We adopt the geometry shown in Fig 1c of Ghisellin Madau (1996). + They. show that the energy density of reflected svuchrotron photon strongly depeuds ou the location of cmittine blob., They show that the energy density of reflected synchrotron photon strongly depends on the location of emitting blob. + We should duit that the axiual sviuiuetrv holds iu the reflected svuchrotrou cluission., We should admit that the aximal symmetry holds in the reflected synchrotron emission. +" We approximate the Thonison scattering event by isotropic scattering with cross section o, aud neglect recoil which is a very good approximation when ες 1."," We approximate the Thomson scattering event by isotropic scattering with cross section $\st$ and neglect recoil, which is a very good approximation when $\eps_s \ll 1$ ." + The angular distribution ofreflected svuchrotron cussion is given λεω.fiero)Πο](μιολες (no is aconstaut). the function fiery) determines the angular distribution of reflected svuchrotron photons (κοπή Madau 1996)," The angular distribution ofreflected synchrotron emission is given by $n_{\rm ph}(\eps_s, \mu, r_0)=n_0 f(\mu,r_0)\eps_s^{-q}$ $n_0$ is aconstant), the function $f(\mu, r_0)$ determines the angular distribution of reflected synchrotron photons (Ghisellini Madau 1996)" +This work was partially supported by CAPES (CoordenacGkzmarkmainDBodyEnd1752mainDodyStart1753a00 de Aperleigooamento cle Pessoal de Nivvel Superior).,This work was partially supported by CAPES (Coordena\c{c}\\tikzmark{mainBodyEnd1752}\~\tikzmark{mainBodyStart1753}aoo de Aperfeiçooamento de Pessoal de vel Superior). +Comparison of [Cu/Felsz» to [(Cu/Fe]sio5 indicates that the DID did not in fact affect the Cu feature and ils removal by (his simple means has not introduced any appreciable fitting error: (he spectrum fits for this line are as good as 40.05 dex.,Comparison of $_{5782}$ to $_{5105}$ indicates that the DIB did not in fact affect the Cu feature and its removal by this simple means has not introduced any appreciable fitting error; the spectrum fits for this line are as good as $\pm$ 0.05 dex. + In every case bul (wo (stars L2208 and L4201). the abunclances from the 5782 and 5105 lines agree to 0.1 dex.," In every case but two (stars L2208 and L4201), the abundances from the 5782 and 5105 lines agree to 0.1 dex." + Though the 5105 line was available for all MA stars. it was not possible in every case Lo determine an abundance from it.," Though the 5105 line was available for all M4 stars, it was not possible in every case to determine an abundance from it." + In stars cooler than 4225 Ix. (with the exception of L2406) the spectrum was too crowded by strong Megll lines to achieve a reliable fit of the sviithesis to the observed spectrum., In stars cooler than 4225 K (with the exception of L2406) the spectrum was too crowded by strong MgH lines to achieve a reliable fit of the synthesis to the observed spectrum. + M15 is the most metal poor cluster in Chis sample., M15 is the most metal poor cluster in this sample. + The overall metallicity is so low that the 5782 feature is too weak {ο synthesize even in the cooler stars: all Chat could be derived [rom it were unenlishtening upper limits., The overall metallicity is so low that the 5782 feature is too weak to synthesize even in the cooler stars; all that could be derived from it were unenlightening upper limits. + We have therefore synthesized only the 5105 line in M15 stars., We have therefore synthesized only the 5105 line in M15 stars. + Given the usually g00d agreement between the two lines in other clusters. we are confident that the derived abundance [rom this line acciuratelv reflects the Cu content of a particular star.," Given the usually good agreement between the two lines in other clusters, we are confident that the derived abundance from this line accurately reflects the Cu content of a particular star." + However. even the stronger 5105 line is still quite weak in the M15 spectra. and individual abundances are sUll uncertain.," However, even the stronger 5105 line is still quite weak in the M15 spectra, and individual abundances are still uncertain." + We put larger errors on MI5 abundances. and our average abundance for this clusterhas a high o.," We put larger errors on M15 abundances, and our average abundance for this clusterhas a high $\sigma$." + The most metal rich cluster in the sample. M1. posed a different sort of problem for our analvsis.," The most metal rich cluster in the sample, M71, posed a different sort of problem for our analysis." + The spectrum becomes so crowded with lines Chat it is very difficult to determine the level of the continuum., The spectrum becomes so crowded with lines that it is very difficult to determine the level of the continuum. + Onlv six stars [rom Snedenetal.(1994). vielded credible results., Only six stars from \citet{Sneden1994} yielded credible results. + In the coolest stars. we found that |Cu/Fe]| correlated with Ἐν with the," In the coolest stars, we found that [Cu/Fe] correlated with $_{eff}$ , with the" +systems with edge-on class Il. and one not-edge-on class II object.,"systems with edge-on class II, and one not-edge-on class II object." + Towards the class 0-I objects. L1527. IRC-L1041-2. and IRASO4302. we have detected abundant H»O. CO». and CO ice in the envelope.," Towards the class 0-I objects, L1527, IRC-L1041-2, and IRAS04302, we have detected abundant $_2$ O, $_2$, and CO ice in the envelope." + The column density ratio of CO» to HO ice is 2]—28%.. which coincides with the ratio observed bySST towards YSOs with various inclinations.," The column density ratio of $_2$ to $_2$ O ice is $21-28$, which coincides with the ratio observed by towards YSOs with various inclinations." + The weak absorption at ~4.1 jm can be fitted by HDO τος; the HDO/H:O ratio ranges from 2 to 22%., The weak absorption at $\sim 4.1$ $\mu$ m can be fitted by HDO ice; the $_2$ O ratio ranges from 2 to 22. +. The absorption in the vicinity of the CO band (4.76 um) is double-peaked and fitted by combining CO ice. OCN'. and CO gas.," The absorption in the vicinity of the CO band (4.76 $\mu$ m) is double-peaked and fitted by combining CO ice, $^-$, and CO gas." + The large column density of CO ice suggests that the envelope is still very dense and cold. while OCN™ and CO gas features would originate in the region close to the protostar.," The large column density of CO ice suggests that the envelope is still very dense and cold, while $^-$ and CO gas features would originate in the region close to the protostar." + The column density of OCN7 is as high as 2—6 relative to H3O. which is much higher than previous observations.," The column density of $^-$ is as high as $2-6$ relative to $_2$ O, which is much higher than previous observations." + Our lines of sight (high inclinations from the rotation axis) may preferentially trace the regions with high UV irradiation. such as the surface of a forming disk and/or torus envelope.," Our lines of sight (high inclinations from the rotation axis) may preferentially trace the regions with high UV irradiation, such as the surface of a forming disk and/or torus envelope." +" The spectrum of IRASO4302 includes the 3.5 gm absorption band. but the feature does not match either CH:OH or CH,."," The spectrum of IRAS04302 includes the 3.5 $\mu$ m absorption band, but the feature does not match either $_3$ OH or $_4$." + An OCS absorption band is tentatively detected towards IRC-L1041-2., An OCS absorption band is tentatively detected towards IRC-L1041-2. + Towards the edge-on class II objects. ASR41 and 2MASS J1628137-243139. we have detected the Π.Ο 2005). 2007))," Towards the edge-on class II objects, ASR41 and 2MASS J1628137-243139, we have detected the $_2$ \cite{crbr05}) \cite{terada07})" + Towards the edge-on class II objects. ASR41 and 2MASS J1628137-243139. we have detected the Π.Ο 2005). 2007)).," Towards the edge-on class II objects, ASR41 and 2MASS J1628137-243139, we have detected the $_2$ \cite{crbr05}) \cite{terada07})" +is arelation between aad Hu some galaxies.,is a relation between and in some galaxies. + However. the yplane has been increasingly populated with less certain cestimates that have potentially been included. on the asstuption that r; is spatially resolved. ioc. that they ollow a previously estimated rrelatiou defined by higher quality cestimates.," However, the plane has been increasingly populated with less certain estimates that have potentially been included on the assumption that $r_i$ is spatially resolved, i.e., that they follow a previously estimated relation defined by higher quality estimates." + Iu these cases; as the simulations prescuted vere show. anebserved irelation will avise simply as a result of the +; selection effect. even if lis randomly distributed within ealaxics.," In these cases, as the simulations presented here show, an relation will arise simply as a result of the $r_i$ selection effect, even if is randomly distributed within galaxies." + Does a population of underauassive SMDIIs exist. aud have they been detected?," Does a population of under-massive SMBHs exist, and have they been detected?" + If the irelation is au upper lait. then there should be galaxies that host low mass SMDITIs. as suggested by simulations (Vittorinietal.2005:Volouteri2007)..," If the relation is an upper limit, then there should be galaxies that host low mass SMBHs, as suggested by simulations \citep{2005MNRAS.363.1376V,2007ApJ...663L...5V}." + Indeed. both the simulations preseuted here. and the observatious plotted in Fieure 23. show that uuderanassive SMDIIs cau be. and have been. detected.," Indeed, both the simulations presented here, and the observations plotted in Figure \ref{fig:3}, show that under-massive SMBHs can be, and have been, detected." + Some notable cases are that of NGC 1135 (Coccatoetal.2006).. a sample of barred ealaxies (Cvaham2008a) and uarrow-line Sevtert 1 ealaxies (Mathur&Grupe2005)..," Some notable cases are that of NGC 4435 \citep{2006MNRAS.366.1050C}, a sample of barred galaxies \citep{2008ApJ...680..143G} and narrow-line Seyfert 1 galaxies \citep{2005ApJ...633..688M}." + Tn fact. the evidence to suggest the presence of uuderauassive SMDIIS is not matched by auy evidence for a siguificaut population of over-Inassive SMDIIS leftward of theobserved 1rclation.," In fact, the evidence to suggest the presence of under-massive SMBHs is not matched by any evidence for a significant population of over-massive SMBHs leftward of the relation." + It is iuportaut to note some intricacies with the two dominant techniques for measuring ((stellar and eas dynamics)., It is important to note some intricacies with the two dominant techniques for measuring (stellar and gas dynamics). + Iu the case of stellar cdyvuauuies. there are systematics to the models that may allow a large range of in a eiven bulee (Valhuiotal.2001)..," In the case of stellar dynamics, there are systematics to the models that may allow a large range of in a given bulge \citep{2004ApJ...602...66V}." + Tn the case of eas dvnandes. it is unclear what the inclination of the clear eas disk may be (e.g.Marconietal.2003)..," In the case of gas dynamics, it is unclear what the inclination of the nuclear gas disk may be \citep[e.g.,][]{2003ApJ...586..868M}." + Both uethods allow the potential for many of the current cestimates to iu fact be upper lanits. ic. the true yplane mav have a large distribution of uuderauassive SMBUs.," Both methods allow the potential for many of the current estimates to in fact be upper limits, i.e., the true plane may have a large distribution of under-massive SMBHs." + Therefore. underauassive SMDIIS may have j'en observed. their inass over-cstimated. and their inpact over-looked.," Therefore, under-massive SMBHs may have been observed, their mass over-estimated, and their impact over-looked." + All SMDBII models allow stringent upper limits to be ulaced on citep|e.g..|[|2002ÀpJ...567..2:278.2009A pJ...692..856D..," All SMBH models allow stringent upper limits to be placed on \\citep[e.g.,][]{2002ApJ...567..237S,2009ApJ...692..856B}." + Ilowever. these uupper huits will also be depenudeut on the available spatial resolution.," However, these upper limits will also be dependent on the available spatial resolution." + Αν kineniatical data will lave two velocity points spatially separated on a scale of R., Any kinematical data will have two velocity points spatially separated on a scale of $\Re$. + An upper linut to lis estimated by iucluding au increasing dark mass wutil the derived model becomes iucousisteut with the data.ie. a higher upper luit to Mill be estimated using a lower R.," An upper limit to is estimated by including an increasing dark mass until the derived model becomes inconsistent with the data,i.e., a higher upper limit to will be estimated using a lower $\Re$." +" There are still important coustraints that can be added to the pplane from estimates of ""upper lits. however. as upper lits that fall below theobserved rvolation provide the same evideuce for underauassive SMBs as would a tightly coustrained low massM,."," There are still important constraints that can be added to the plane from estimates of upper limits, however, as upper limits that fall below the relation provide the same evidence for under-massive SMBHs as would a tightly constrained low mass." +. At present. most uupper lanits are based on data clerived from gas dynamics.," At present, most upper limits are based on data derived from gas dynamics." +" Couscequently. these limits generally fall above theobserved ielation due to uninown amounts of Lue broadening frou, non-gravitational processes. anc uncertainties iu the inclination if the nuclear gas disk."," Consequently, these limits generally fall above the relation due to unknown amounts of line broadening from non-gravitational processes, and uncertainties in the inclination if the nuclear gas disk." + Ifthe irelation represcuts au upper lit iu the pplaue. then what is this hnüt?," If the relation represents an upper limit in the plane, then what is this limit?" +" In 77. an upper lait of a,=&.7.5,5.0 was found based on the distribution ofobserved SMDITIS leftward of theobserved irclation."," In \ref{mands} an upper limit of $\alpha_u=8.7, \beta_u=5.0$ was found based on the distribution of SMBHs leftward of the relation." + This is likely a good approximation to the Huit as there are uo reports of a steeper relation., This is likely a good approximation to the limit as there are no reports of a steeper relation. + ILowever. this estinate does not include the poteutial overauassivo SMDIIS from the upper lanits calculated w Beiflorictal.(2009)..," However, this estimate does not include the potential over-massive SMBHs from the upper limits calculated by \cite{2009ApJ...692..856B}." +" Tuclucdine these limits to the upper limit sample gives values of a,=8.8.3,3.9 too3,,."," Including these limits to the upper limit sample gives values of $\alpha_u=8.8, \beta_u=3.9$ to." +". As expected. due to the addition of SAIBIT nuits at lowerM,.. this lait is more shallow with a higher zero-point."," As expected, due to the addition of SMBH limits at lower, this limit is more shallow with a higher zero-point." + Iucludiug hese linüts at the lower eond of the ane addresses. i part. a limitation of the sample used here.," Including these limits at the lower end of the plane addresses, in part, a limitation of the sample used here." + As already noted. the ccatalog used here is likely incomplete due το the difficulty of iieasuriug Hin faint galaxies at ereater distances.," As already noted, the catalog used here is likely incomplete due to the difficulty of measuring in faint galaxies at greater distances." + In addition. the cceatalog likely coutains iuhoimogenuuous nmeasurenienuts that may not translate from bulee to bulec.," In addition, the catalog likely contains inhomogenuous measurements that may not translate from bulge to bulge." + What are the consequences to galaxy evolution models if there is oulv a irclation?, What are the consequences to galaxy evolution models if there is only a relation? + First. galaxies will no longer be required to obey the rrelation. and could host a SMDII with anv j)below0.," First, galaxies will no longer be required to obey the relation, and could host a SMBH with any below." +"),4.. Models that 1uclude feedback from he SMDII to the galaxy will then need to be carefully reconsidered.", Models that include feedback from the SMBH to the galaxy will then need to be carefully reconsidered. + While the SAIBIT will undoubtedly have sole iufiueuce on a portion of the host ealaxy. it would iof need to affect large scale properties: evolution of the SMDBII would be a result of host galaxy evolution.," While the SMBH will undoubtedly have some influence on a portion of the host galaxy, it would not need to affect large scale properties; evolution of the SMBH would be a result of host galaxy evolution." + Aw upper limit iu the yplane would also represeu otιο pinnacle of SMDIT evolution as a function of0... iu which galaxies evolve up to the lanit.," An upper limit in the plane would also represent the pinnacle of SMBH evolution as a function of, in which galaxies evolve up to the limit." + A signature of such a scenario could be a cosmic variation in € (the scatter would iucrease with redshift) aud anobserved rrelation that docs uot exceed05)., A signature of such a scenario could be a cosmic variation in $\epsilon$ (the scatter would increase with redshift) and an relation that does not exceed. + Tf the distribution of lis random within bulges. then when compared with the localobserved relation. cestimates from ligher redshift could fall to the left or the right.," If the distribution of is random within bulges, then when compared with the local relation, estimates from higher redshift could fall to the left or the right." +" Treuetal.(2007). find a population of z=0.36 Sevtert l1 ealaxics that lie above the localobserved rrelation by Aloga= Alogl/,=0.5 L butthispopulationstilllicsbelowthe( ]inüt estimated. here."," \cite{2007ApJ...667..117T} find a population of z=0.36 Seyfert 1 galaxies that lie above the local relation by $\Delta\log\sigma=0.13, \Delta\log$ $=0.54$ but this population still lies below the limit estimated here." + Finally. if SMDIISs can reside auvwhere below the," Finally, if SMBHs can reside anywhere below the" + pdB/By + 2x Dy dpedv| ο iP. 2py0D / Dy+ 2x Dy dpdeyo f pin Dy. where (he approximate expressions retain only linear terms in perturbed quantities.," p_0 B / B_0 + 2 B_0 d d f m )^2, 2 p_0 B / B_0 + 2 B_0 d d f m B_0, where the approximate expressions retain only linear terms in perturbed quantities." + Linearizing the expression lor the relaxed Maxwellian in equation (3)) about. fy. the BGI collision operator is given by ) — vof c vfux (ry m ," Linearizing the expression for the relaxed Maxwellian in equation \ref{eq:BGK}) ) about $f_0$ , the drift-kinetic BGK collision operator is given by ) = f + f_0 u ) - -." +The cdiift-kinetic equation including the DGIx operator can be linearized to obtain the following equation lor 9f Of= αμ.” X. η.A −⊲mw)(ΜΗ, The drift-kinetic equation including the BGK operator can be linearized to obtain the following equation for $\delta f$ f= - )f_0 (-i B )+ +sinall portion of the spectrum of JOLL757 A is not shown because it falls between the atmospheric windows where the telhivic correction was very poor.,small portion of the spectrum of J041757 A is not shown because it falls between the atmospheric windows where the telluric correction was very poor. + The detection of JO11757 D was too weak for a useful spectrum., The detection of J041757 B was too weak for a useful spectrum. + Using IR spectroscopy. Quauzetal.(2010). classified CATIA Tau 1 as à vouug brown dwiuf with a spectral type of L2+0.5.," Using IR spectroscopy, \citet{qua10} classified CAHA Tau 1 as a young brown dwarf with a spectral type of $\pm$ 0.5." + If Ποτοῦ A is a ember of Taurus. Barradoetal.(20090)/ estimated that it should have a temperature of 15501750 I& based on a comparison of its photometry to the fluxes predicted by theoretical evolutionary models. which also correspouds to an L spectral type (Dahnetal.2002).," If J041757 A is a member of Taurus, \citet{bar09} estimated that it should have a temperature of 1550–1750 K based on a comparison of its photometry to the fluxes predicted by theoretical evolutionary models, which also corresponds to an L spectral type \citep{dah02}." +. To assess the accuracy of these classifications. we beein bv comparing our spectra of CATIA Tau 1 and JOll?57 A iu Fieure 1 o a spectrum of oue of the coolest known menibers of Taurus. KPNO 1(M9.5:Driceiioetal.2002).. which has been reddened to roughly match the spectral slopes of the two candidates.," To assess the accuracy of these classifications, we begin by comparing our spectra of CAHA Tau 1 and J041757 A in Figure \ref{fig:spec} to a spectrum of one of the coolest known members of Taurus, KPNO 4 \citep[M9.5;][]{bri02}, which has been reddened to roughly match the spectral slopes of the two candidates." + The most prominent features m the spectrum of IKRPNO. Ll are the deep H2O absorption bands., The most prominent features in the spectrum of KPNO 4 are the deep $_2$ O absorption bands. + CAILA Tau 1 and JO11757 A should show Πο bands that are as strong or stronecr as those in KPNO 1 if they are L-type members of Taurus. but this is nof the case.," CAHA Tau 1 and J041757 A should show $_2$ O bands that are as strong or stronger as those in KPNO 4 if they are L-type members of Taurus, but this is not the case." + Instead. CATA Tau 1 exhibits weak Πο absorption that is indicative of a spectral type of M5.M6 if it is a dwarf and M3ML if it is a voungstar’.," Instead, CAHA Tau 1 exhibits weak $_2$ O absorption that is indicative of a spectral type of M5–M6 if it is a dwarf and M3–M4 if it is a young." +. The absence of detectable Ποο absorption in the spectra of JOILT5T A places constraints of (dwaif) and ΠΟ (voune) on the spectral type., The absence of detectable $_2$ O absorption in the spectrum of J041757 A places constraints of $\lesssim$ M2 (dwarf) and $\lesssim$ M0 (young) on the spectral type. + The M2near-IR absorption features frou the plotospleres of vouug stars can be diluted by contiuuuu veiling from circtustellar dust cluission. but this is nof a possible explanation for the absence of strong IIO absorption for CAITÀA Tau I aud JOI1T57 À since they do not exhibit IR excess cission (Section 3.2.0)).," The near-IR absorption features from the photospheres of young stars can be diluted by continuum veiling from circumstellar dust emission, but this is not a possible explanation for the absence of strong $_2$ O absorption for CAHA Tau 1 and J041757 A since they do not exhibit IR excess emission (Section \ref{sec:mid}) )." + Both candidates are iuuch too faint to be members of Taurus with spectral types of M3.ALL aud ZMAO., Both candidates are much too faint to be members of Taurus with spectral types of M3–M4 and $\lesssim$ M0. + Thus. we classify them as background stars rather than substellar ποος of Taurus.," Thus, we classify them as background stars rather than substellar members of Taurus." + Iu addition to spectroscopy. we have used the available photometry to examine whether the candidates from Quizetal.(2010) and Barracoetal.(2009) ave likely to be voung brown dwarfs.," In addition to spectroscopy, we have used the available photometry to examine whether the candidates from \citet{qua10} and \citet{bar09} are likely to be young brown dwarfs." + We have retrieved photometry in the five optical bands of SDSS (ugriz:Fukueitaotal.1906) from the Sixth Data Release of the survey (Acdelmau-MeCarthyetal.2008). for CAITA Tau 15 aud ΙΟτο A (the D compoucut was not detected by SDSS)., We have retrieved photometry in the five optical bands of SDSS \citep[$ugriz$;][]{fuk96} from the Sixth Data Release of the survey \citep{ade08} for CAHA Tau 1--5 and J041757 A (the B component was not detected by SDSS). + The calibration of these images is described by Padiuauabhnauetal.(2008)., The calibration of these images is described by \citet{pad08}. +".. We selected the data measured with an aperture radius of 1.715"". Barra", We selected the data measured with an aperture radius of $1.745\arcsec$. +doctal.(2009) ineasured / and + photometry for JOLLT57 A and D from archival ππασος from the Canacda-Frauce- Telescope (CFUT)., \citet{bar09} measured $i$ and $z$ photometry for J041757 A and B from archival images from the Canada-France-Hawaii Telescope (CFHT). + The data for J011757 A from Darradoetal.(2009) agree withthose frou SDSS. which sugeests that the two photometric svstems are simular.," The data for J041757 A from \citet{bar09} agree with those from SDSS, which suggests that the two photometric systems are similar." + Therefore. we have adopted the / aud : photometry frou Barradoctal.(2009) for οτον D. We have compiled photometry at J. 77. and A frou the Two-Microu All-Skyv Survey (2\TASS:Shautskie2006) for CAITA Tau 1. 2. Lo and 5. from (2010) for CATIA Tau 3. aud from Dirradoetal. for JOLL757 A and D. All of these sources also appear within archival images at 3.6. 15. 5.8. 8.0. aud 2124 that were obtained by theTelescope.," Therefore, we have adopted the $i$ and $z$ photometry from \citet{bar09} + for J041757 B. We have compiled photometry at $J$, $H$, and $K_s$ from the Two-Micron All-Sky Survey \citep[2MASS;][]{skr06} for CAHA Tau 1, 2, 4, and 5, from \citet{qua10} for CAHA Tau 3, and from \citet{bar09} for J041757 A and B. All of these sources also appear within archival images at 3.6, 4.5, 5.8, 8.0, and 24 that were obtained by the." +" We have ieasured photometry frou, these images using the methods that were applied to the kuowu nembers of Taurus by Lulunanetal.(2010)..", We have measured photometry from these images using the methods that were applied to the known members of Taurus by \citet{luh10tau}. + JO11757 D )econies increasingly dominant over JOLL757 À at ongcr wavelengths. as shown in Fieure 2..," J041757 B becomes increasingly dominant over J041757 A at longer wavelengths, as shown in Figure \ref{fig:sub}." + To measure photometry at 3.6 and 1.5 for 0Ητοῦ A. we subtracted a scaled point spread ποοι (PSF) at the location of J011757 D. (Mareugoetal.2006).," To measure photometry at 3.6 and 4.5 for J041757 A, we subtracted a scaled point spread function (PSF) at the location of J041757 B \citep{mar06}." +. The A component was uot detected in the PSF-subtracted images at 5.8 and &.0sau., The A component was not detected in the PSF-subtracted images at 5.8 and 8.0. + The spatial resolution of the image at 21 jis too low for resolving J011757 A aud D. We have assigned all of the 21 Hux to JOST D since it is much redder than ΠΕτο A at 3.68.0pau., The spatial resolution of the image at 24 is too low for resolving J041757 A and B. We have assigned all of the 24 flux to J041757 B since it is much redder than J041757 A at 3.6–8.0. + We have adopted the average ueasuremieut ina eiven band if an object was observed at inultiple epochs bwSpifzer., We have adopted the average measurement in a given band if an object was observed at multiple epochs by. + We did not measure photometry from the second| epoch of images at 3.6 Tor J01l757 A aud D because they were coutaminatedX cosnuie rays., We did not measure photometry from the second epoch of images at 3.6 for J041757 A and B because they were contaminatedby cosmic rays. + Our measurements ofSpitzer plotometry or CATIA Tau 1:5 aud JOII757 A aud Bare preseuted iu Table 1.., Our measurements of photometry for CAHA Tau 1–5 and J041757 A and B are presented in Table \ref{tab:phot}. + To determine if the candidates from Quangctal.(2010) and Darradoetal.(2009) have the appropriate optical colors aud magnitudes for low-mass members of Taurus we plot them on a diagram of / versus /i in Figure 3. with all known late-tvpe members of Taurus (>AI6) that were detected by SDSS.," To determine if the candidates from \citet{qua10} and \citet{bar09} have the appropriate optical colors and magnitudes for low-mass members of Taurus, we plot them on a diagram of $i$ versus $i-z$ in Figure \ref{fig:col} + with all known late-type members of Taurus $>$ M6) that were detected by SDSS." + We include im that diagram all other point sources in the SDSS nuages that have photometric uncertainties less than 0.1 mae., We include in that diagram all other point sources in the SDSS images that have photometric uncertainties less than 0.1 mag. + All of the candidates appear below the sequence of known members. which iudicates that they are probably field stars or galaxies.," All of the candidates appear below the sequence of known members, which indicates that they are probably field stars or galaxies." + CATIA Tau 2. Land 5 were detected in the survey by Luliuau(2001) and were identified as probable uommembers for the same reason.," CAHA Tau 2, 4, and 5 were detected in the survey by \citet{luh04tau} and were identified as probable nonmembers for the same reason." + One additional candidate. CATIA Tau 3. was euconipassed by the images from Liulunan(2001)... but its sigual-to-noise ratio was too low for useful photometry.," One additional candidate, CAHA Tau 3, was encompassed by the images from \citet{luh04tau}, but its signal-to-noise ratio was too low for useful photometry." + Young stars that are occulted by eireuustellar disks and seeu in scattered liebt can appear anomalously faint for their colors. but JOUTS7 D is the only caudidate that shows possible evidence of a disk im its mid-IR photometry (Section 3.2.0)).," Young stars that are occulted by circumstellar disks and seen in scattered light can appear anomalously faint for their colors, but J041757 B is the only candidate that shows possible evidence of a disk in its mid-IR photometry (Section \ref{sec:mid}) )." + A diagrain of /A. versus JZT is useful for distinguishing between late-tvpe objects aud reddened stars at earlier types (Luliman 2000). , A diagram of $i-K_s$ versus $J-H$ is useful for distinguishing between late-type objects and reddened stars at earlier types \citep{luh00tau}. . +We show a diagram of this Kind in Fieure 3 for the candidates from: Quangctal.(2010) and Biuradoetal. (2009). the known late-type members of Taurus.aud all other sources with errors less than 0.1 mae in the SDSS aud 2MASS images of Taurus.," We show a diagram of this kind in Figure \ref{fig:col} for the candidates from \citet{qua10} and \citet{bar09}, , the known late-type members of Taurus,and all other sources with errors less than 0.1 mag in the SDSS and 2MASS images of Taurus." + (ΑΠΔ Tau 2 is the only candidate that has the appropriate colors for a voung object later than M6., CAHA Tau 2 is the only candidate that has the appropriate colors for a young object later than M6. + Given its position iu / versus / :," Given its position in $i$ versus $i-z$ ," + Given its position iu / versus / :.," Given its position in $i$ versus $i-z$ ," +recent study of the ovewall (infra-red to X-ray) spectrum of hhas been published byv Janiuk ((2001). considering in particular the strong emission in the UV and soft X-ray bands and proposing its origin in à warm optically thick ‘skin’ on the aceretion disc.,"recent study of the overall (infra-red to X-ray) spectrum of has been published by Janiuk (2001), considering in particular the strong emission in the UV and soft X-ray bands and proposing its origin in a warm optically thick `skin' on the accretion disc." + This work also included an analysis of an extended: oobservation in LOOT. suggesting a cold rellection factor Ro = Ofer. where ο is the solid. angle subtended by the reflecting matter. of order unity.," This work also included an analysis of an extended observation in 1997, suggesting a cold reflection factor R = $\Omega$ $\pi$ , where $\Omega$ is the solid angle subtended by the reflecting matter, of order unity." + wavas observed by oon 2001 June 15 vieling a useful exposure of 60 ksec., was observed by on 2001 June 15 yielding a useful exposure of $\sim$ 60 ksec. + In this paper we use data from the EPIC pn camera (Strüdder al.2001). which has he best sensitivity of any instrument llown to date in the 10 keV spectral band. the combined EPIC MOS cameras (Lurner 22001). and the Reflection Crating SpectrometersROS (den Herder 22001).," In this paper we use data from the EPIC pn camera (Strüdder 2001), which has the best sensitivity of any instrument flown to date in the $\sim$ 6-10 keV spectral band, the combined EPIC MOS cameras (Turner 2001), and the Reflection Grating Spectrometer/RGS (den Herder 2001)." + Reference to the Optical Monitor (Mason 22001) confirmed the strong optical ancl UV. emission was close to the tvpical level in143., Reference to the Optical Monitor (Mason 2001) confirmed the strong optical and UV emission was close to the typical level in. +.. AM X-ray data were first. screened. with the NAIA SAS v5.3 software and events corresponding to patterns 0-4. (single ancl double pixel events) were selected for the pn data and patterns 0-12 for MOSI and ALOS2. the latter then being combined.," All X-ray data were first screened with the XMM SAS v5.3 software and events corresponding to patterns 0-4 (single and double pixel events) were selected for the pn data and patterns 0-12 for MOS1 and MOS2, the latter then being combined." + A low οσον cut of 200 eV. was applied to all X-ray data and known hot or bad. pixels were removed., A low energy cut of 200 eV was applied to all X-ray data and known hot or bad pixels were removed. + We extracted source counts within a circular region of rpadius defined around the centroid position οἱ143.. with the background. being taken from a similar region. olfse from but close to the source.," We extracted source counts within a circular region of radius defined around the centroid position of, with the background being taken from a similar region, offset from but close to the source." + The 2-10 keV. X-ray pn light curve is reproduced as figure 1 and shows ~80 percent Lux changes over 6 ksec. similar o those seen in the delata.," The 0.2-10 keV X-ray pn light curve is reproduced as figure 1 and shows $\sim$ 30 percent flux changes over $\sim$ 6 ksec, similar to those seen in the data." + Individual spectra were binned to a minimum of 20 counts per bin. to facilitate use of the X7 minimalisation echnique in spectral fitting.," Individual spectra were binned to a minimum of 20 counts per bin, to facilitate use of the $\chi^2$ minimalisation technique in spectral fitting." + Itesponse functions for spectral itting to the RGS data were generated from the SAS v5.3., Response functions for spectral fitting to the RGS data were generated from the SAS v5.3. + Spectral fitting was based on the XAspec package (Arnaud 1996) and used a grid of ionised absorber mocels calculated with the NSTAR code (Ixallman 11996)., Spectral fitting was based on the Xspec package (Arnaud 1996) and used a grid of ionised absorber models calculated with the XSTAR code (Kallman 1996). + All spectral fits include absorption due to the line-D. . ∪⇂−⊳∖↓⋏∙≟↓∐≺∣⋜↧↓⋯⇍⇂⊔∼≼∼∪↓⊔⊔↓⊔∪⇂⇀∖⊓∶−≽⋅↖∖⋅↱≻↓∪↾≼∼⊔↓⊳∟↓⋅↓⋅∪↓⋅⊳∖, All spectral fits include absorption due to the line-of-sight Galactic column of $N_{H}=2.85\times10^{20}\rm{cm}^{-2}$. +⋅⇁⋅σι2 ∙⊲ ⋜⊔⋅⋖⋅⊏↥⊔∪↿⋖⋅∠⇂⋜∐↿⇂↥∢⋅≤⋗∪⊲∕⋰∣↙≼⇍∪⊔∐∠⇂∢⋅⊔≼⇍∢⊾↓⋖⊾∖⇁∢⊾↓↿∖∆∖⇉∶⇉⋅−⇂∎∪↓⋅∪⊔⋖⋅ interesting parameter)., Errors are quoted at the confidence level $\Delta \chi^{2}=2.7$ for one interesting parameter). + X-ray spectra of AGN at 210 keV are well fitted. to first order. with a power law of photon index E in the range ~1.6-2 for most radio quiet AGN. with a fraction (eg NLS1) jwing somewhat steeper indices.," X-ray spectra of AGN at 2–10 keV are well fitted, to first order, with a power law of photon index $\Gamma$ in the range $\sim$ 1.6-2 for most radio quiet AGN, with a fraction (eg NLS1) having somewhat steeper indices." +" The widelv held: view is hat this ‘hare’ X-ray continuum in Sevlert galaxies arises w Comptonisation of thermal emission from the accretion disc in a ‘hot’ corona (eg Haardt and. Alaraschi 1991). and »oduces. additional spectral features. by ""reflection [rom dense matter in the disc (eg Pounds 11990. Fabian 22000)."," The widely held view is that this `hard' X-ray continuum in Seyfert galaxies arises by Comptonisation of thermal emission from the accretion disc in a `hot' corona (eg Haardt and Maraschi 1991), and produces additional spectral features by `reflection' from dense matter in the disc (eg Pounds 1990, Fabian 2000)." + We began our analysis of bby confirming there were no obvious spectral changes with source [ux and then proceeded to fit the ppn and MOS data integrated. over the full 00 Ksec observation., We began our analysis of by confirming there were no obvious spectral changes with source flux and then proceeded to fit the pn and MOS data integrated over the full $\sim$ 60 ksec observation. + A simple power law fit over the L10 keV band yielded a photon index of P—1.79 (pn) and E—1.71 (MOS). with a broad. excess in the dataumodel ratio between 37 keV. and evidence of absorption at higher energies in both data sets (figure 2).," A simple power law fit over the 1–10 keV band yielded a photon index of $\Gamma$$\sim$ 1.79 (pn) and $\Gamma$$\sim$ 1.71 (MOS), with a broad excess in the data:model ratio between 3–7 keV, and evidence of absorption at higher energies in both data sets (figure 2)." + The fit was statistically unacceptable with an overall x2 /dof of 1541/1176., The fit was statistically unacceptable with an overall $\chi^{2}$ /dof of 1541/1176. +" When extrapolated to 0.3 keV. the 1.10 keV fits to both pn and MOS data revealed a strong ""soft excess’ (figure 3)."," When extrapolated to 0.3 keV, the 1–10 keV fits to both pn and MOS data revealed a strong `soft excess' (figure 3)." + To improve the 110 keV fit we added further. spectra components to match the most obvious features in the data., To improve the 1–10 keV fit we added further spectral components to match the most obvious features in the data. + The indication of an extreme broad emission line suggestec rellection from the inner accretion disc. conventionally modelled with a λος line in Nspec (Laor. 1991).," The indication of an extreme broad emission line suggested reflection from the inner accretion disc, conventionally modelled with a LAOR line in Xspec (Laor 1991)." +" The addition of a LAO. linc. with inclination initially Lixec at 30° and 42,,,7100/2, ( where Re=CMfC is the eravitational radius for mass AZ). resulted in a significan statistical improvement (C dotof 1304/1172). but with an unrealistically large. EW of ~l4 keV (pn) and ον]. keV (MOS)."," The addition of a LAOR line, with inclination initially fixed at $\deg$ and $R_{out}$ $R_{g}$ ( where $R_{\rm g} = GM/c^2$ is the gravitational radius for mass $M$ ), resulted in a significant statistical improvement $\chi^{2}$ /dof of 1304/1172), but with an unrealistically large EW of $\sim$ 1.4 keV (pn) and $\sim$ 1.1 keV (MOS)." + To better fit the broad. line profile we added. a eaussian line with energv tied to that ofthe LAOR. line. (, To better fit the broad line profile we added a gaussian line with energy tied to that ofthe LAOR line. ( + Physically such a gaussian line could. represent emission from larger radii on the disc)., Physically such a gaussian line could represent emission from larger radii on the disc). + This addition gave a further, This addition gave a further +"where a((T,,v) is the Gaunt factor assumed to be 1, a correct value for the range of astrophysical quantities involved in our calculations; v and r, are the frequency measured in MHz and the free-free optical depth, respectively, and T, is the electron temperature in K of the intervening ionized gas.","where $T_{e}$ $\nu$ ) is the Gaunt factor assumed to be 1, a correct value for the range of astrophysical quantities involved in our calculations; $\nu$ and $\tau_{\nu}$ are the frequency measured in MHz and the free-free optical depth, respectively, and $\mathrm{T_{e}}$ is the electron temperature in K of the intervening ionized gas." +" By measuring the relative strengths of the [FelI] lines observed in the near and mid infrared emission, ? conclude that the emitting region behind the J-type shocks as observed in the J and H bands of the 2MASS has a temperature of 12000 K. Although not quantified, the authors recognize a large uncertainty associated with this magnitude as a consequence of different beam size and possibly different filling factors in their measurements."," By measuring the relative strengths of the [FeII] lines observed in the near and mid infrared emission, \citet{rho01} + conclude that the emitting region behind the J-type shocks as observed in the J and H bands of the 2MASS has a temperature of 12000 K. Although not quantified, the authors recognize a large uncertainty associated with this magnitude as a consequence of different beam size and possibly different filling factors in their measurements." + If we assume that Τε is in a reasonable range between 8000-12000 K (which includes the temperature as estimated from the IR observations) and use the optical depth derived from our radio study particularly for this region where the thermal absorption is stronger (r74— 0.3) we obtain an EM between approximately 2.8x10? and 5.0x10? cm~® pc for the eastern rim., If we assume that $\mathrm{T_{e}}$ is in a reasonable range between 8000-12000 K (which includes the temperature as estimated from the IR observations) and use the optical depth derived from our radio study particularly for this region where the thermal absorption is stronger $\tau_{74}\sim0.3$ ) we obtain an EM between approximately $2.8\times10^{3}$ and $5.0\times10^{3}$ $^{-6}$ pc for the eastern rim. +" By combining this emission measure with the postshock electron density, we can roughly calculate the thickness of the molecular gas layer that has been dissociated and ionized by the SNR shock front."," By combining this emission measure with the postshock electron density, we can roughly calculate the thickness of the molecular gas layer that has been dissociated and ionized by the SNR shock front." +" If we assume an electron density of n,~500 cm? as estimated by ? and ? on the basis of forbidden [FelI] lines, we conclude that the dissociation and ionization processes took place ina thin screen of about 3.4 to 6.0 x10!6 cm (—0.01-0.02 pc)."," If we assume an electron density of $n_{e}\sim\,500$ $^{-3}$ as estimated by \citet{fes80} and \citet{rea00} on the basis of forbidden [FeII] lines, we conclude that the dissociation and ionization processes took place in a thin screen of about 3.4 to 6.0 $\times10^{16}$ cm $\sim$ 0.01-0.02 pc)." +" This is a small path compared with the transverse dimensions over which thermal absorption is observed, but is a lower limit if the ionized gas is clumped."," This is a small path compared with the transverse dimensions over which thermal absorption is observed, but is a lower limit if the ionized gas is clumped." +"potential, where no component of angular momentum is conserved and we see an interesting box orbit as a consequence.","potential, where no component of angular momentum is conserved and we see an interesting box orbit as a consequence." +" Frequent, oblique passages through the disk would presumably increase the potential for encounters with massive structures such as stellar clusters or giant molecular clouds and may be partly or wholly responsible for the high velocity dispersion observed in the Hydra I and the EBS stream."," Frequent, oblique passages through the disk would presumably increase the potential for encounters with massive structures such as stellar clusters or giant molecular clouds and may be partly or wholly responsible for the high velocity dispersion observed in the Hydra I and the EBS stream." +" The EBS stream adds to the growing list of halo streams that can be mapped over a sufficient extent that, with suitable follow-up observations, they could be used as probes of the Galactic potential."," The EBS stream adds to the growing list of halo streams that can be mapped over a sufficient extent that, with suitable follow-up observations, they could be used as probes of the Galactic potential." + A preliminary orbit estimate shows that the EBS is unrelated to either the Anticenter or Monoceros streams., A preliminary orbit estimate shows that the EBS is unrelated to either the Anticenter or Monoceros streams. + The somewhat intermediate breadth of the stream together with its relatively high velocity dispersion suggests the possibility that the progenitor could have been more massive than the globular clusters thought to be responsible for the half dozen very cold streams discovered in the SDSS footprint to date., The somewhat intermediate breadth of the stream together with its relatively high velocity dispersion suggests the possibility that the progenitor could have been more massive than the globular clusters thought to be responsible for the half dozen very cold streams discovered in the SDSS footprint to date. +" However, if the progenitor had been a dark matter dominated dwarf galaxy, it would be difficult to understand how it could have held onto it’s dark matter envelope for any length of time in such a confined and eccentric orbit."," However, if the progenitor had been a dark matter dominated dwarf galaxy, it would be difficult to understand how it could have held onto it's dark matter envelope for any length of time in such a confined and eccentric orbit." +" On the other hand, this very orbit may have subjected both the progenitor and the stream to significant heating through encounters with massive structures in the disk."," On the other hand, this very orbit may have subjected both the progenitor and the stream to significant heating through encounters with massive structures in the disk." +" If Hydra I is indeed the progenitor of the EBS, then it is only the second probably unbound progenitor to be associated with a tidal stream."," If Hydra I is indeed the progenitor of the EBS, then it is only the second probably unbound progenitor to be associated with a tidal stream." + A more detailed examination of the structure and stellar kinematics in this remnant may shed new light on the end stage of tidal disruption., A more detailed examination of the structure and stellar kinematics in this remnant may shed new light on the end stage of tidal disruption. +" Though contamination by field stars is high, Hydra I may be particularly attractive in this respect as it is four times closer to us than Bootes III(Grillmair2009)."," Though contamination by field stars is high, Hydra I may be particularly attractive in this respect as it is four times closer to us than Bootes III\citep{grill2009}." +. Refinement of the orbit will require radial velocity and proper motion measurements of carefully selected stars along the length of the stream., Refinement of the orbit will require radial velocity and proper motion measurements of carefully selected stars along the length of the stream. +" Given the very low surface density of stream stars and very high field star contamination, this will necessarily be an ongoing task."," Given the very low surface density of stream stars and very high field star contamination, this will necessarily be an ongoing task." +" In this respect, the EBS may be particularly well situated for follow-up by the upcoming spectroscopic LAMOST survey."," In this respect, the EBS may be particularly well situated for follow-up by the upcoming spectroscopic LAMOST survey." + Gaia and LSST proper motion measurements may also help us to refine the orbit and perhaps trace the stream over a much longer arc., Gaia and LSST proper motion measurements may also help us to refine the orbit and perhaps trace the stream over a much longer arc. + The author is grateful to an anonymous referee for constructive and insightful comments., The author is grateful to an anonymous referee for constructive and insightful comments. + Thanks also go to Kevin Schlaufman for providing the positions of ECHOS member stars., Thanks also go to Kevin Schlaufman for providing the positions of ECHOS member stars. +" Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. Sloan.."," Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. ." + 4.9. the value found by Helsdon Ponman: however. the intercepts for the two samples differ at greater than the 20 level.,"$\sim$ 4.9, the value found by Helsdon Ponman; however, the intercepts for the two samples differ at greater than the $\sigma$ level." + In addition we used a 1-D KS test to compare the temperature and luminosity distributions of sample | and sample 2., In addition we used a 1-D KS test to compare the temperature and luminosity distributions of sample 1 and sample 2. + We find iat the luminosity distributions of the samples are consistent with their being drawn from a single parent population. whereas 1e temperature distributions are different at the ~ 90 per cent Contidence level.," We find that the luminosity distributions of the samples are consistent with their being drawn from a single parent population, whereas the temperature distributions are different at the $\sim$ 90 per cent confidence level." + Therefore. we find a small. but significant difference in ve |uminosity-temperature relation for groups containing radio galaxies. with higher temperatures found at a given luminosity.," Therefore, we find a small, but significant difference in the luminosity-temperature relation for groups containing radio galaxies, with higher temperatures found at a given luminosity." + There are two possible interpretations for this result: either the radio galaxies are heating their environments. or they are reducing the luminosity of the groups.," There are two possible interpretations for this result: either the radio galaxies are heating their environments, or they are reducing the luminosity of the groups." + To test whether the second scenario could be true. we assumed that all of the material through which the radio source passes while expanding has been affected by the radio source and no longer emits in X-rays (e.g. it may have been lifted out to a distance where the subsequent decrease in density produces this effect).," To test whether the second scenario could be true, we assumed that all of the material through which the radio source passes while expanding has been affected by the radio source and no longer emits in X-rays (e.g. it may have been lifted out to a distance where the subsequent decrease in density produces this effect)." + Modelling the expanding radio source as a cone. and comparing the square of number density integrated over the volume through which the radio source will have passed with the same integral for the entire group. we find that only 7 per cent of the luminosity of 3C 66B could have been removed in this way.," Modelling the expanding radio source as a cone, and comparing the square of number density integrated over the volume through which the radio source will have passed with the same integral for the entire group, we find that only 7 per cent of the luminosity of 3C 66B could have been removed in this way." + Using the original group luminosity. a temperature of 0.92+0.21 keV s predicted from our best-fitting relation for radio-quiet groups.," Using the original group luminosity, a temperature of $\pm0.21$ keV is predicted from our best-fitting relation for radio-quiet groups." + The effect will be at a similar level in other sources., The effect will be at a similar level in other sources. + Therefore we conclude that the radio galaxy could not have produced such a decrease in luminosity., Therefore we conclude that the radio galaxy could not have produced such a decrease in luminosity. + For this reason. we interpret these results as providing indirect evidence for radio-source heating of group environments.," For this reason, we interpret these results as providing indirect evidence for radio-source heating of group environments." + This result shows that groups which contain detected radio sources are systematically hotter than those which do not. providing the first evidence that radio-source heating is occurring at a detectable level in the majority of such groups.," This result shows that groups which contain detected radio sources are systematically hotter than those which do not, providing the first evidence that radio-source heating is occurring at a detectable level in the majority of such groups." + It is apparent from Fig., It is apparent from Fig. + 12. that 3C 66B is much hotter than predicted by either of the two relations slotted., \ref{lt} that 3C 66B is much hotter than predicted by either of the two relations plotted. + There are three contributing factors which could explain 3C 66B's exceptional behaviour., There are three contributing factors which could explain 3C 66B's exceptional behaviour. + Firstly. it is more powerful than any of the sources in the radio-loud sample. and so would be expected ο transfer more energy to the group gas.," Firstly, it is more powerful than any of the sources in the radio-loud sample, and so would be expected to transfer more energy to the group gas." + As described above. it is also likely to be comparatively old for an FR-I radio galaxy. so will dave been putting energy into its environment for a longer period of time.," As described above, it is also likely to be comparatively old for an FR-I radio galaxy, so will have been putting energy into its environment for a longer period of time." + Finally. the environment of this source is less massive han the environments of other large. powerful sources. such as 3C 3|.," Finally, the environment of this source is less massive than the environments of other large, powerful sources, such as 3C 31." + It seems plausible that these three factors have combined to make 3C 66B particularly efficient at raising the temperature of its environment., It seems plausible that these three factors have combined to make 3C 66B particularly efficient at raising the temperature of its environment. + Our images of 3C 66B and 3C 449 provide direct evidence that their environments are having a major effect on the evolution of the radio sources., Our images of 3C 66B and 3C 449 provide direct evidence that their environments are having a major effect on the evolution of the radio sources. + In both cases large rounded lobes are associated with large amounts of X-ray-emitting gas. whereas narrower lobes are seen where there is no surrounding material.," In both cases large rounded lobes are associated with large amounts of X-ray-emitting gas, whereas narrower lobes are seen where there is no surrounding material." + In the case of 3C 66B the presence of a blob of gas at the end of the eastern jet and bounded lobe is suggestive of an obstacle having a strong impact on the lobe morphology., In the case of 3C 66B the presence of a blob of gas at the end of the eastern jet and bounded lobe is suggestive of an obstacle having a strong impact on the lobe morphology. + We have revisited the luminosity-temperature relation for, We have revisited the luminosity-temperature relation for +in frequency shifts at selected. values of ας duringA4N2J andΑΛ.,in frequency shifts at selected values of $|m|/\ell$ during and. + We again notice that the lrequencies at all latitudes during24 are lower than (hose during2., We again notice that the frequencies at all latitudes during are lower than those during. + Further. in view of the above discussion. il is not surprising to obtai a single epoch of mininmnm atALV2J for all three values of |m{6," Further, in view of the above discussion, it is not surprising to obtain a single epoch of minimum at for all three values of $|m|/\ell$." + On the contrary. the frequency. shilts hint towards a double minima at ALIN for η = 0.5. and a single minimum [for |m|/( = 0.7 and 0.9.," On the contrary, the frequency shifts hint towards a double minima at for $|m|/\ell$ = 0.5, and a single minimum for $|m|/\ell$ = 0.7 and 0.9." + Although most of well-known surface activity indices do not indicate the minimum in late 2007. as the downward trend in their values continued until (he end of 2008. there are observations where early signs of the beginning of the new cycle were possibly seen.," Although most of well-known surface activity indices do not indicate the minimum in late 2007, as the downward trend in their values continued until the end of 2008, there are observations where early signs of the beginning of the new cycle were possibly seen." + The appearance of a sunspot in Active Region 10981 [or three consecutive davs during 2008 Jan 4. Gal high latitude (80° N) with new evcle polarities hinted towards the rise of evele 24 in early 2008 (Phillips2008a).., The appearance of a sunspot in Active Region 10981 for three consecutive days during 2008 Jan 4 – 6 at high latitude $^{\mathrm{o}}$ N) with new cycle polarities hinted towards the rise of cycle 24 in early 2008 \citep{newcycle}. + ILowever. more sunspots fulfilling the criteria of evele 24 were not visible for several months alter 2008 January.," However, more sunspots fulfilling the criteria of cycle 24 were not visible for several months after 2008 January." + The emergence of three big sunspots al low latitudes will previous evcle polarities a few months later (magnetic polarity in accordance with evele 23) suggested that the minimum after cycle 23 had not vel been reached (Phillips2003b).., The emergence of three big sunspots at low latitudes with previous cycle polarities a few months later (magnetic polarity in accordance with cycle 23) suggested that the minimum after cycle 23 had not yet been reached \citep{oldcycle}. + These observations are not. unusual as sunspots from both eveles are ταον seen during the minimum before approaching the bottom level of the activity., These observations are not unusual as sunspots from both cycles are randomly seen during the minimum before approaching the bottom level of the activity. + The (wo strong dips seen in oscillation Irequencies al certain latitudes can be understood in terms of localized changes in the activitv level., The two strong dips seen in oscillation frequencies at certain latitudes can be understood in terms of localized changes in the activity level. + The mean variation in sunspot number and radio [ux during the two minima discussed here is shown in Figure 5., The mean variation in sunspot number and radio flux during the two minima discussed here is shown in Figure 5. + A close examination of the right panels of Figure 5 reveals that there was indeed a slight rise in activity al (he beginning of 2008 but the trend did not continue., A close examination of the right panels of Figure 5 reveals that there was indeed a slight rise in activity at the beginning of 2008 but the trend did not continue. + The two dips seen in oscillation frequencies mav be interpreted as the manilestation of the competition between the magnetic fields from both the solar eveles., The two dips seen in oscillation frequencies may be interpreted as the manifestation of the competition between the magnetic fields from both the solar cycles. + We also note periodic variations in ὃν thatare addressed in Section 4., We also note periodic variations in $\delta\nu$ thatare addressed in Section 4. +"This ensures that The observed multiplicity frequency fap,=Notuttipte/M is. after this selection. 35+6% (Poisson errors). wheretotal Notutiple 18 the number of binary or multiple systems (38) and N44; 18 the number of observed systems (108).","This ensures that The observed multiplicity frequency $f_\mathrm{obs}=N_\mathrm{Multiple}/N_\mathrm{Total}$ is, after this selection, $35\pm6\%$ (Poisson errors), where $N_\mathrm{Multiple}$ is the number of binary or multiple systems (38) and $N_\mathrm{Total}$ is the number of observed systems (108)." + Figure 5 shows the observed multiplicity fraction for each primary spectral type., Figure 5 shows the observed multiplicity fraction for each primary spectral type. + The multiple systems included in the following analysis can be found in Table 2. and the number of single:binary:triple:quadruple systems in Table 4.," The multiple systems included in the following analysis can be found in Table 2, and the number of single:binary:triple:quadruple systems in Table 4." + To compute the actual multiplicity frequency. we need to consider two effects: (1) at small separations. we detect more equal brightness binaries than systems with large component brightness differences. and (11) a brightness-limited sample is biased in favour of (previously unresolved) binaries or multiple systems compared to single stars.," To compute the actual multiplicity frequency, we need to consider two effects: (i) at small separations, we detect more equal brightness binaries than systems with large component brightness differences, and (ii) a brightness-limited sample is biased in favour of (previously unresolved) binaries or multiple systems compared to single stars." + Assuming that the flux ratio distribution is independent of the separation in the observed range (which can be transformed into a flat mass ratio distribution). we estimate the number of multiple systems of close separations that we miss using the following method.," Assuming that the flux ratio distribution is independent of the separation in the observed range (which can be transformed into a flat mass ratio distribution), we estimate the number of multiple systems of close separations that we miss using the following method." + We divide the number of binaries in Fig., We divide the number of binaries in Fig. + 3 of observed Az' as a function of angular separation p into four different regions of interest., 3 of observed $\Delta z^\prime$ as a function of angular separation $\rho$ into four different regions of interest. +" Assuming that our sample is complete to Ar’<5.5 between angular separation and complete to Az’€2.5 for closer separations. the ratioof companions in the region p=0.5”-33"". ADS225-535 and p=0.5""—PLZ 37, Az=0-2.5 is the same as the ratio of companions in p.=0.1—0.57, Az=2.5—5.5 and p50.1720.57, Az=0—2.5."," Assuming that our sample is complete to $\Delta z^\prime\la5.5$ between angular separation and complete to $\Delta z^\prime\la2.5$ for closer separations, the ratioof companions in the region $\rho=0.5\arcsec-3\arcsec$, $\Delta z^\prime=2.5-5.5$ and $\rho=0.5\arcsec-3\arcsec$ , $\Delta z^\prime=0-2.5$ is the same as the ratio of companions in $\rho=0.1\arcsec-0.5\arcsec$, $\Delta z^\prime=2.5-5.5$ and $\rho=0.1\arcsec-0.5\arcsec$, $\Delta z^\prime=0-2.5$." + This would result in the survey missing two binary companions in the close separation - high flux ratio region. hence the total multiple fraction should be increased to 37+6%.," This would result in the survey missing two binary companions in the close separation - high flux ratio region, hence the total multiple fraction should be increased to $37\pm6\%$." + We compute the multiplicity fraction for a volume-limited sample. f. following the method and Eq. (," We compute the multiplicity fraction for a volume-limited sample, $f^\prime$ , following the method and Eq. (" ++) of ? whereobs δε.=0.37obs is the fraction of observed binaries after sensitivity correction.,"4) of \citet{Burgasser2003} + where $f^\prime_\mathrm{obs}=0.37$ is the fraction of observed binaries after sensitivity correction." + ? consider a values in the range 2.8.) corresponding to only equal brightness systems. to 1.9. which corresponds to a flat flux ratio distribution.," \citet{Burgasser2003} consider $\alpha$ values in the range 2.8, corresponding to only equal brightness systems, to 1.9, which corresponds to a flat flux ratio distribution." + The distribution of z-band brightness ratios (see Table 2) in our sample is more peaked towards unequal systems (on a linear brightness ratio scale). resulting ina=1.73.," The distribution of $z^\prime$ -band brightness ratios (see Table 2) in our sample is more peaked towards unequal systems (on a linear brightness ratio scale), resulting in $\alpha = 1.73$." + According to Eq. (, According to Eq. ( +4) of ?.. this then yields a multiplicity fraction for a volume-limited sample of f!=2546%.,"4) of \citet{Burgasser2003}, this then yields a multiplicity fraction for a volume-limited sample of $f^\prime=25\pm6\%$." +" However. the ? sample is based on a correlation of M dwarf candidates selected from the 400 million sources in the 2MASS point source catalogue (PSC.angularresolution2"",?) with the 1500000 sources in the ROSAT AII Sky Survey (RASS.angularresolution~30.?).. thus the brightness limit is imposed by the X-ray luminosity of the sources."," However, the \citet{Riaz2006} sample is based on a correlation of M dwarf candidates selected from the 400 million sources in the 2MASS point source catalogue \citep[PSC, angular resolution $\sim2\arcsec$,][]{Cutri2003} with the 000 sources in the ROSAT All Sky Survey \citep[RASS, angular resolution $\sim30\arcsec$, ][]{Voges1999}, thus the brightness limit is imposed by the X-ray luminosity of the sources." + Hence. we need to correct for the excess of multiple systems as two or more stellar components emit more X-rays than the corresponding primary component would do if it were single.," Hence, we need to correct for the excess of multiple systems as two or more stellar components emit more X-rays than the corresponding primary component would do if it were single." + We can do this straightforwardly by simply examining all our a posteriori known multiple systems and determining which ones would not have beer included in the sample if the primary had been single., We can do this straightforwardly by simply examining all our a posteriori known multiple systems and determining which ones would not have been included in the sample if the primary had been single. +" X-ray counts and errors are available from ROSAT (?). for each of the 44 multiple systems (except for one system. J20500010-1154092., which is counted as a non-detection here)."," X-ray counts and errors are available from ROSAT \citep{Voges1999} for each of the 44 multiple systems (except for one system, J20500010-1154092, which is counted as a non-detection here)." + Given that the components in any given system should be coeval. it is assumed that the X-ray brightness depends only on the stellar luminosity.," Given that the components in any given system should be coeval, it is assumed that the X-ray brightness depends only on the stellar luminosity." + According to ον. Lx/Lpo is roughly constant as function of spectral type. hence to a reasonable approximation the X-ray count rate can be assumed to be directly proportional to the brightness fraction in. z-band in linear units.," According to \citet{Riaz2006}, $L_{\rm X} / L_{\rm bol}$ is roughly constant as function of spectral type, hence to a reasonable approximation the X-ray count rate can be assumed to be directly proportional to the brightness fraction in $z^\prime$ -band in linear units." + Thus. we use the known Az’ for each system in combination with the unresolved X-ray count rate to estimate the rate for the primary component alone.," Thus, we use the known $\Delta z^\prime$ for each system in combination with the unresolved X-ray count rate to estimate the rate for the primary component alone." + If the new value results in S/N«3.3. the multiple system in question is counted as having been positively selected for and is excluded for the purpose of calculating the multiple fractio for a volume-limited sample. where S/N=3.3 Is the relevant criterion for detection according to the tables of ?..," If the new value results in $S/N < 3.3$, the multiple system in question is counted as having been positively selected for and is excluded for the purpose of calculating the multiple fraction for a volume-limited sample, where $S/N = 3.3$ is the relevant criterion for detection according to the tables of \citet{Voges1999}." + In total. 7 systems are identified as contaminants in this way.," In total, 7 systems are identified as contaminants in this way." + Hence. applying corrections for the X-ray flux limit as described above. it follows that the multiplicity fraction f is given by f=(38—7)/(1087)1.05332+6%.," Hence, applying corrections for the X-ray flux limit as described above, it follows that the multiplicity fraction $f$ is given by $f = (38-7) / (108-7)*1.053 = 32\pm6$." +. While both multiplicity fractions f and f’ agree within the uncertainties. in the following we assume a multiplicity fraction f=32+6%. as the brightness limit is primarily imposed by the X-ray luminosity.," While both multiplicity fractions $f$ and $f^\prime$ agree within the uncertainties, in the following we assume a multiplicity fraction $f=32\pm6\%$, as the brightness limit is primarily imposed by the X-ray luminosity." + We note that some overabundance of short-period binaries (P<20 days) might be present in the X-ray selected sample. but this cannot be quantified until future radial velocity observations have been performed.," We note that some overabundance of short-period binaries $<$ 20 days) might be present in the X-ray selected sample, but this cannot be quantified until future radial velocity observations have been performed." +" We also note that this fraction might still include a small contamination by non-physical (""optical"") binaries. as second-epoch observations for some of the systems are still pending. although we reiterate that the fraction of binaries that are merely optical must be very small (see Fig."," We also note that this fraction might still include a small contamination by non-physical (“optical”) binaries, as second-epoch observations for some of the systems are still pending, although we reiterate that the fraction of binaries that are merely optical must be very small (see Fig." + 4)., 4). + The individual component photometric spectral types from Sect., The individual component photometric spectral types from Sect. + 3.1 are transformed to approximate masses using the mass estimates of ? for young (~500 MMyr) stars., 3.1 are transformed to approximate masses using the mass estimates of \citet{KrausHillenbrand2007} for young $\sim500$ Myr) stars. + We interpolate linearly to obtain masses for subclasses of 0.5 and calculate the mass-ratios. g= ΜΜ.," We interpolate linearly to obtain masses for subclasses of 0.5 and calculate the mass-ratios, $q=M_2/M_1$ ." + The binarieswhere the secondary star is suspected to be an L dwarf are not included in the mass-ratio distribution because of the high uncertainties in mass., The binarieswhere the secondary star is suspected to be an L dwarf are not included in the mass-ratio distribution because of the high uncertainties in mass. +We also exclude components at separations greater than 6” from the,We also exclude components at separations greater than $6\arcsec$ from the +equilibrium: required. for the N-ray mass. modelling will therefore not apply.,equilibrium required for the X-ray mass modelling will therefore not apply. + The A-ray images for the other clusters included in the present sample do not exhibit any dramatic substructure that would clearly invalidate such assumptions., The X-ray images for the other clusters included in the present sample do not exhibit any dramatic substructure that would clearly invalidate such assumptions. + We note. however. the presence of an A-ray luminous subcluster. approximately 2.6 aremin (850 kpe) to the northwest of Abell 2744 (ACIS). visible in the LIRL data.," We note, however, the presence of an X-ray luminous subcluster, approximately 2.6 arcmin (850 kpc) to the northwest of Abell 2744 (AC118), visible in the HRI data." + This subeluster is also identified in the weak lensing analysis of Small (L997)., This subcluster is also identified in the weak lensing analysis of Smail (1997). + The details of the ASCA and ROSAT observations are summarized in Tables 1 and 2. respectively.," The details of the ASCA and ROSAT observations are summarized in Tables 1 and 2, respectively." + The basic X-ray properties of the target clusters are summarized in Table 3., The basic X-ray properties of the target clusters are summarized in Table 3. + The lensing clusters studied in this paper are crawn from a larger sample of X-ray luminous svstems cliscussecl by Allen CI997. in preparation).," The lensing clusters studied in this paper are drawn from a larger sample of X-ray luminous systems discussed by Allen (1997, in preparation)." + A more detailed description of, A more detailed description of +like to be able to recoucile our scheme with the observed abuudances of simple iiolecular ious aud their deuterated counterparts.,like to be able to reconcile our scheme with the observed abundances of simple molecular ions and their deuterated counterparts. + The observed range of values for 77 nuplies a value of S of 500035000 (equ. [1]., The observed range of values for $\bar{R}$ implies a value of $\cal S$ of 5000–35000 (eqn. \ref{DCOP}] ]). + Tn order to calculate the expected value of S in LISIN from equation (3)). we need to know the electron. fraction. ἐς. aud »kjc(mng). the total removal rate of ToD! throughli reactions with heavy species.," In order to calculate the expected value of $\cal S$ in L134N from equation \ref{calS}) ), we need to know the electron fraction, $x_e$, and $\sum_j k_j x(m_j)$, the total removal rate of $_2$ $^+$ through reactions with heavy species." +" We can calculate these values approximately frou the the steady-state IL} concentration, which is given by and the charge conservation equation: Note that the rate coefficients iun equation (13)) are the samme as in equation (3)). suco wo assune that ID! reacts at the same rate as IL}. except for electron recolubination. where we use the value for oULL) of LAS<10""(P/3001I&)999 measured by Suudstrónmui (1991)."," We can calculate these values approximately from the the steady-state $_3^+$ concentration, which is given by and the charge conservation equation: Note that the rate coefficients in equation \ref{h3plus}) ) are the same as in equation \ref{calS}) ), since we assume that $_2$ $^+$ reacts at the same rate as $_3^+$, except for electron recombination, where we use the value for $\alpha_e({\rm H_3^+})$ of $1.15\times10^{-7}(T/300\,{\rm K})^{-0.65}$ measured by Sundströmm (1994)." + TCO! and NoII! have been observed in L131N. aud their respective abuudauces are 1.2«10 “and τς1019 (Seide 1980: Womack 1992a: Dickeus 2000).," $^+$ and $_2$ $^+$ have been observed in L134N, and their respective abundances are $1.2\times10^{-8}$ and $7\times10^{-10}$ (Swade 1989; Womack 1992a; Dickens 2000)." +" Adopting plivsical paramcters of T= LOW and 1.)=2010! ? (Dickens 2000). we thus have three equations (3|. |13]]. and |L1]]) in three unknowns νο, (IL) ame Εμ]."," Adopting physical parameters of $T=10$ K and $n({\rm H_2}) = +2\times10^4$ $^{-3}$ (Dickens 2000), we thus have three equations \ref{calS}] ], \ref{h3plus}] ], and \ref{econs}] ]) in three unknowns $x_e$, $x({\rm H_3^+})$, and $\sum_j k_j x(m_j)$ )." + Solving these equations for different values of S allows us to coustrain the chemical state ofLISIN., Solving these equations for different values of $\cal S$ allows us to constrain the chemical state of. +. Also. because the rate cocfhicicuts for reactions of I. aud IbD! with heavy molecules are z2.10P ems howe can calculate the total abundance of heavy molecules frou oux derived value of »kein).," Also, because the rate coefficients for reactions of $_3^+$ and $_2$ $^+$ with heavy molecules are $\approx 2\times10^{-9}$ $^3$ $^{-1}$, we can calculate the total abundance of heavy molecules from our derived value of $\sum_j k_j x(m_j)$." +" For S=5000. we caleulate x,=Lbs107. (ILI)=1.6«10. ὃν and EE=16«10 Ll consistent with little depletion of CO. No. and ο from the gas phase."," For ${\cal S} = 5000$, we calculate $x_e = 1.4\times10^{-8}$, $x({\rm +H_3^+}) = 1.6\times10^{-9}$ , and $\sum_j x(m_j) = 1.6\times10^{-4}$ consistent with little depletion of CO, $_2$, and O from the gas phase." + Ou the other haud. for S=35000. we derive L9«10 Sey) =6«107. and SSG)=L9«10 . auplving siguificaut depletion.," On the other hand, for ${\cal S} = 35000$, we derive $x_e = +1.9\times10^{-8}$ , $x({\rm H_3^+}) = 6\times10^{-9}$, and $\sum_j +x(m_j) = 1.9\times10^{-5}$ , implying significant depletion." + For the intermediate value of Sz10! interred from the deuterated auunonia fractionation. we find that partial depletion (~50% )) of heavy molecules is required to account for the observed deuterim enhancements.," For the intermediate value of ${\cal S} \approx 10^4$ inferred from the deuterated ammonia fractionation, we find that partial depletion $\sim 50$ ) of heavy molecules is required to account for the observed deuterium enhancements." + A similar couclision has Όσοι reached by Roberts Milhuw (2000)., A similar conclusion has been reached by Roberts Millar (2000). + Finally. as our calculated ionization levels lead to eood agrecment with the observed aabundanee (63). we conclude that the observed abundances aud fractionations in LIStN are wellanatched by steady-state tou-imolecule Chemistry.," Finally, as our calculated ionization levels lead to good agreement with the observed abundance $\S\ref{secgasamm}$ ), we conclude that the observed abundances and fractionations in L134N are well-matched by steady-state ion-molecule chemistry." + Despite reservations concerning the plivsical conditions in LISIN. some as vet unidentified mechanisu may be retimming mmautle-formed molecules to the gas there.," Despite reservations concerning the physical conditions in L134N, some as yet unidentified mechanism may be returning mantle-formed molecules to the gas there." + The uantle abundances of ammonia isotoponiers computed miuericalh bv Brown Mill (1989) (see their table 2) effectively rule out a erain surface origin when scaled to he amunonia abundanceiuEN., The mantle abundances of ammonia isotopomers computed numerically by Brown Millar (1989) (see their table 2) effectively rule out a grain surface origin when scaled to the ammonia abundancein. +. However. their surface reaction scheme permits lareer D/ITI ratios than those xeseuted. since the fractionation is proportional to the eas pliase atomic D/II ratio. R(D). which may be higher hau the value asstmed iu thei calculations.," However, their surface reaction scheme permits larger D/H ratios than those presented, since the fractionation is proportional to the gas phase atomic D/H ratio, $R$ (D), which may be higher than the value assumed in their calculations." + Following he scheme of Brown Millar. we cau derive values for he surface fractionation ratios: Thus.if RED)z 0.05. then it is also possible to explain the observed abundances by surface formation of annona.," Following the scheme of Brown Millar, we can derive values for the surface fractionation ratios: Thus,if $R({\rm D}) \approx 0.05$ , then it is also possible to explain the observed abundances by surface formation of ammonia." +"cut-off located at the peak of the integrand, relatively more will be removed of mass m at the first dip than of other masses.","cut-off located at the peak of the integrand, relatively more will be removed of mass $m$ at the first dip than of other masses." +" We also note that the additive peaks of the minimum (cut-off) mass, shown with red solid curves, are not correlated with the location of the dip in the distribution (marked with a black solid tick mark), but shift as a function of the collision velocity."," We also note that the additive peaks of the minimum (cut-off) mass, shown with red solid curves, are not correlated with the location of the dip in the distribution (marked with a black solid tick mark), but shift as a function of the collision velocity." +" At higher collisional velocity, where waves are stronger, the largest number of cut-off particles are actually produced by the particles in the dip, which would provide a negative feedback according to the traditional picture, canceling the waves."," At higher collisional velocity, where waves are stronger, the largest number of cut-off particles are actually produced by the particles in the dip, which would provide a negative feedback according to the traditional picture, canceling the waves." + Our analysis shows that positive feedback is not the dominant effect in the production of the waves., Our analysis shows that positive feedback is not the dominant effect in the production of the waves. +" The location of the peak of the removal term (Term I in Paper I) is the determining factor in the formation of the waves, which can be given by solving where m gives the location of the first dip and M is given so that X(u,M)=m."," The location of the peak of the removal term (Term I in Paper I) is the determining factor in the formation of the waves, which can be given by solving where $m$ gives the location of the first dip and $M$ is given so that $X(\mu,M)\equiv m$." +" Unfortunately, an analytic solution can not be given, as the derivative is transcendent, but is easily solvable numerically."," Unfortunately, an analytic solution can not be given, as the derivative is transcendent, but is easily solvable numerically." +" The variable that determines the solution will be X, which depends on the collisional velocity, the tensile strength law, and the value of min."," The variable that determines the solution will be $X$, which depends on the collisional velocity, the tensile strength law, and the value of $m_{\rm min}$ ." +" For a given system, where Mmin and the tensile strength law do not change as a function of radial location, the single variable determining the wavy structure is the collisional velocity."," For a given system, where $m_{\rm min}$ and the tensile strength law do not change as a function of radial location, the single variable determining the wavy structure is the collisional velocity." + The minimum mass will change as a function stellar types and some material/velocity dependence of the tensile strength law can be expected between systems., The minimum mass will change as a function stellar types and some material/velocity dependence of the tensile strength law can be expected between systems. +" In Section 2,, we have shown that variations in the proxies for the collisional velocity (such as disk radius) can in fact result in wavy size distributions, even within a narrow debris ring."," In Section \ref{models}, we have shown that variations in the proxies for the collisional velocity (such as disk radius) can in fact result in wavy size distributions, even within a narrow debris ring." +" Extended debris disks will be even more likely to have higher collisional velocities, as the particles with higher 8 values less than get placed on high eccentricity orbits."," Extended debris disks will be even more likely to have higher collisional velocities, as the particles with higher $\beta$ values (but less than 0.5) get placed on high eccentricity orbits." + (butThese small particles0.5) from the inner rings will collide with the particles in the external rings with increased velocities due to the non-zero collisional angles (Thébault&Augereau2007)., These small particles from the inner rings will collide with the particles in the external rings with increased velocities due to the non-zero collisional angles \citep{thebault07}. +". To study the effects of variations only in the collisional velocity, we set the collisional velocity directly within our code to specific values, not changing any other parameters."," To study the effects of variations only in the collisional velocity, we set the collisional velocity directly within our code to specific values, not changing any other parameters." + We present the results from these runs in Figure 9.., We present the results from these runs in Figure \ref{fig:v}. + The figure shows that waves start to appear at collisional velocityvalues of 3kms! and above., The figure shows that waves start to appear at collisional velocityvalues of $3~{\rm km~s}^{-1}$ and above. +10σ deviation from the mean.,$10\sigma$ deviation from the mean. + The inconsistency with the milder decline found in hydrodynamical simulations ?) was also noted by GO9.," The inconsistency with the milder decline found in hydrodynamical simulations \citep[e.g.,][]{roncarelli} was also noted by G09." +eqreffirehose)) and mirror-mode instabilities eqrefmirrormode)) in a manner similar to that (Equationseen in observations of the solar wind (Hellingeretal.2006;Baleetal.2009) and in earlier low- current sheet simulations (Drakeetal.,") and mirror-mode instabilities ) in a manner similar to that seen in observations of the solar wind \citep{Hellinger06, Bale09} + and in earlier $\beta$ current sheet simulations \citep{Drake10}." +" reffiredist shows 2010)..the data for our system in the space of where a= P,/P,.", \\ref{firedist} shows the data for our system in the space of $\alpha$ $\beta_\parallel$ ) where $\alpha = P_\perp/P_\parallel$ . +" This plot is generated by (α,β)calculating the anisotropy o and the β for each grid point.", This plot is generated by calculating the anisotropy $\alpha$ and the $\beta_{\parallel}$ for each grid point. +" The plot is a two-dimensional histogram of grid points in space, where is calculated based on The ("," The plot is a two-dimensional histogram of grid points in $\alpha$ $\beta_{\parallel}$ ) space, where $\beta_{\parallel}$ is calculated based on $P_\parallel$." +"o,8j)parallel and perpendicularBj pressures are calculatedP4. by taking the diagonal components of the pressure tensor after rotating into the frame of the local magnetic field, such that the two perpendicular components are equal."," The parallel and perpendicular pressures are calculated by taking the diagonal components of the pressure tensor after rotating into the frame of the local magnetic field, such that the two perpendicular components are equal." +" We look at the distribution at t—80,120, and "," We look at the distribution at $t = 80, 120,$ and $160\Omega_{\text{ci}}^{-1}$." +"At early times the anisotropies have not yet fully 1609,,'.developed and the plasma still occupies a small region in (a,6)) space.", At early times the anisotropies have not yet fully developed and the plasma still occupies a small region in $\alpha$ $\beta_{\parallel}$ ) space. +" By t=1200! the anisotropy has reached the two stability boundaries, and continues to be confined between these two boundaries at t= even as the average ( increases."," By $t=120\Omega_{\text{ci}}^{-1}$ the anisotropy has reached the two stability boundaries, and continues to be confined between these two boundaries at $t=160\Omega_{\text{ci}}^{-1}$, even as the average $\beta$ increases." +" The anisotropy 1600;,',reaches the stability boundaries at a time after the short wavelength Weibel modes have dissipated.", The anisotropy reaches the stability boundaries at a time after the short wavelength Weibel modes have dissipated. +" Since at this point there are no longer large regions with essentially zero magnetic fields, the firehose and mirror-mode instabilities are what determine the boundaries of the temperature anisotropies."," Since at this point there are no longer large regions with essentially zero magnetic fields, the firehose and mirror-mode instabilities are what determine the boundaries of the temperature anisotropies." + There are no clear signatures of the classical mirror-mode instability at this time., There are no clear signatures of the classical mirror-mode instability at this time. +" The firehose and mirror-mode instabilities may be hard to distinguish among the turbulent interacting magnetic islands, or the islands may just stop generating anisotropy as they approach the instability boundaries."," The firehose and mirror-mode instabilities may be hard to distinguish among the turbulent interacting magnetic islands, or the islands may just stop generating anisotropy as they approach the instability boundaries." +" The islands maintain an elongated form for the simulation shown in clear until t=1200;', the latest time simulated refjzbeta((c)for 6=4.8."," The islands maintain an elongated form for the simulation shown in \\ref{jzbeta}( (c) clear until $t=120\Omega_{\text{ci}}^{-1}$, the latest time simulated for $\beta=4.8$." + This is shown in refjzanisotropy((a) showing the out-of-plane current for t=, This is shown in \\ref{jzanisotropy}( (a) showing the out-of-plane current for $t=120\Omega_{\text{ci}}^{-1}$. +" Since the edges of the islands are pushing against the 120051.firehose instability, the tension force in the magnetic fields is eliminated."," Since the edges of the islands are pushing against the firehose instability, the tension force in the magnetic fields is eliminated." + This can be seen in refjzanisotropy((b) which shows the regions that are unstable to the firehose instability., This can be seen in \\ref{jzanisotropy}( (b) which shows the regions that are unstable to the firehose instability. +" The magnetic islands that reach a significant amplitude are much more elongated at high 8, than at low 8..", The magnetic islands that reach a significant amplitude are much more elongated at high $\beta_e$ than at low $\beta_e$. + These elongated islands should be found even for moderate values of θε at realistic mass ratios., These elongated islands should be found even for moderate values of $\beta_e$ at realistic mass ratios. +" Island elongation is caused by the suppression of the shorter wavelength tearing modes by pressure anisotropies (Pj> P,) that develop due to the Fermi acceleration of electrons.", Island elongation is caused by the suppression of the shorter wavelength tearing modes by pressure anisotropies $P_\parallel > P_\perp$ ) that develop due to the Fermi acceleration of electrons. + Later in time the plasma develops pressure anisotropies of both ions and electrons that are limited by the firehose and Weibel instabilities., Later in time the plasma develops pressure anisotropies of both ions and electrons that are limited by the firehose and Weibel instabilities. + A Weibel mode develops that kinks the magnetic field lines., A Weibel mode develops that kinks the magnetic field lines. +" In the regime with a real mass ratio we would expect even longer islands to form, where multiple wavelengths of the firehose instability could develop."," In the regime with a real mass ratio we would expect even longer islands to form, where multiple wavelengths of the firehose instability could develop." +" At late time the fraction of points unstable to the firehose instability saturates, and the anisotropy is confined between the mirror-mode and firehose instability boundaries."," At late time the fraction of points unstable to the firehose instability saturates, and the anisotropy is confined between the mirror-mode and firehose instability boundaries." + The long islands persist due to the low requirement of anisotropy to reach the marginal firehose condition at high 8., The long islands persist due to the low requirement of anisotropy to reach the marginal firehose condition at high $\beta$. + For even small anisotropies the tension in the magnetic fields is removed., For even small anisotropies the tension in the magnetic fields is removed. +" When encountering magnetic islands in the heliosheath, we predict the formation of similar extended, sausage-shaped islands rather than the more round islands found in low-3 simulations (Drakeet 2010)."," When encountering magnetic islands in the heliosheath, we predict the formation of similar extended, sausage-shaped islands rather than the more round islands found in $\beta$ simulations \citep{Drake10}." +". The cores of these islands should also be at the marginal firehose condition, so the magnetic tension that drives them to become round vanishes."," The cores of these islands should also be at the marginal firehose condition, so the magnetic tension that drives them to become round vanishes." +" We would thus expect these sausage shapes to persist long after the islands have ceased growing, and thus could be found even in regions where reconnection is no longer occurring."," We would thus expect these sausage shapes to persist long after the islands have ceased growing, and thus could be found even in regions where reconnection is no longer occurring." + These elongated islandsexhibit signatures that can be seen in data., These elongated islandsexhibit signatures that can be seen in data. +" In particular, measures all three components of the magnetic field."," In particular, measures all three components of the magnetic field." +" Of particular interest for the explorations of islands that grow in the ecliptic plane is the angle αλ= tan~!(Br/Br), where Br "," Of particular interest for the explorations of islands that grow in the ecliptic plane is the angle $\lambda =\tan^{-1}\left(B_T/B_R\right)$ , where $B_T$ " +lluxes from. Bergeron.ctal.(2001) ancl several similar measurements (AleCook&Sion1999): variability appears possible in this star.,fluxes from \citet{ber01} and several similar measurements \citep{mcc99}; variability appears possible in this star. +.. NLI'T 43806 is [listed among nearby white dwarts with a photometric distance of 15 pe (Llolberectal.2008:WKawka&Vennes 2006).. based on an estimated = mag (Salim&Could.2003).," NLTT 43806 is listed among nearby white dwarfs with a photometric distance of 15 pc \citep{hol08,kaw06}, based on an estimated $V=15.9$ mag \citep{sal03}." +. Phe SDSS photometry implies a distance greater than 20 pe. with g=17.0 mag (V=16.5 mag) and d—24 pe.," The SDSS photometry implies a distance greater than 20 pc, with $g=17.0$ mag $V\approx16.8$ mag) and $d=24$ pc." +.. Vhis white dwarf is neither a binary nor binary suspect., This white dwarf is neither a binary nor binary suspect. + Cireenetal.(2000). incorrectIy identified this extreme ultraviolet source with anearby. unrelated: infrared source (Farihietal.2006).," \citet{gre00} incorrectly identified this extreme ultraviolet source with anearby, unrelated near-infrared source \citep{far06}." +... GD 231 is a double. degenerate: (Morales-etal.2005). and may displav mild. near-infrared excess due to its unseen companion. but the larger 2ALASS excess is not confirmed.," GD 231 is a double degenerate \citep{mor05} and may display mild near-infrared excess due to its unseen companion, but the larger 2MASS excess is not confirmed." + Figure A4O— demonstrates. that he entire spectral energy. distribution. can be decently matched by a 16500 KODA model. but. this may simply reflect a good approximation of the composite light rather han an accurate effective temperature of either Component.," Figure \ref{fig40} demonstrates that the entire spectral energy distribution can be decently matched by a 16500 K DA model, but this may simply reflect a good approximation of the composite light rather than an accurate effective temperature of either component." + Assuming the 18500 Ix. spectroscopic temperature derived » Bergeronetal.(1992). is correct. Figure ΑΕ predicts an excess at JLAN. that is consistent with a 6000 Ix. DC companion.," Assuming the 18500 K spectroscopic temperature derived by \citet{ber92} is correct, Figure \ref{fig41} predicts an excess at $JHK$ that is consistent with a 6000 K DC companion." +. Vhis star was tentatively classified as a magnetic DA. but has since been correctly reclassified as DB (Jordan2001).," This star was tentatively classified as a magnetic DA, but has since been correctly reclassified as DB \citep{jor01}." +. Phe apparent 2\LASS excess at this white dwarl is not confirmed: the 13000 Ix fit to its spectral energy distribution shown in Figure X43. is the first time an elfective temperature has been assessed for this DB star. ., The apparent 2MASS excess at this white dwarf is not confirmed; the 13000 K fit to its spectral energy distribution shown in Figure \ref{fig43} is the first time an effective temperature has been assessed for this DB star. . +. Vhe C-band. photometry for this DZ star appears influenced by its strong Ca LL ancl Ix. absorption. and. possibly other elements (Sionctal.1990).," The $U$ -band photometry for this DZ star appears influenced by its strong Ca H and K absorption, and possibly other elements \citep{sio90}." +... GD 402 is a suspected DA|DC system based on the fact that its optical colors (and. the shape of its optical spectrum) predict a significantly higher effective temperature than do its. relatively weak Balmer lines (Bergeronctal.1990)., GD 402 is a suspected DA+DC system based on the fact that its optical colors (and the shape of its optical spectrum) predict a significantly higher effective temperature than do its relatively weak Balmer lines \citep{ber90}. +.. Although the expanded: spectral energy distribution | including the WYER near-infrared data can be nearly. reproduced by a single 7000 Ix. component as shown in Figure οι the photometry can also be well-modeled with two white dwarf components of approximately SOOO Ix and 5000 Ix. (Figure X47)).," Although the expanded spectral energy distribution – including the IRTF near-infrared data – can be nearly reproduced by a single 7000 K component as shown in Figure \ref{fig46}, the photometry can also be well-modeled with two white dwarf components of approximately 8000 K and 5000 K (Figure \ref{fig47}) )." + This magnetic white dwarf has a spectral energv distribution that cannot be [fitted by a single temperature (non-magnetic) white dwarf. model., This magnetic white dwarf has a spectral energy distribution that cannot be fitted by a single temperature (non-magnetic) white dwarf model. + Figure Ads fits the combined ultraviolet ancl optical fluxes. to vielel a temperature near 13000. WK. while Figure ALO fits the combined optical ancl near-infrarecl colors with a temperature around 10500. Ix. It is unclear. whether this clilliculty is related to the magnetic nature of the star. or if there is a real Jiffy photometric excess (relative to magnetic nmioclels). implying the possibility of binarity. .," Figure \ref{fig48} fits the combined ultraviolet and optical fluxes to yield a temperature near 13000 K, while Figure \ref{fig49} fits the combined optical and near-infrared colors with a temperature around 10500 K. It is unclear whether this difficulty is related to the magnetic nature of the star, or if there is a real $JHK$ photometric excess (relative to non-magnetic models), implying the possibility of binarity. ." +. Vhis white dwarl may have a small A-band excess. but more data are needed to confirm or rule out this possibility: the V-band Dux appears too bright for the mocel shown in Figure A50.. .," This white dwarf may have a small $K$ -band excess, but more data are needed to confirm or rule out this possibility; the $V$ -band flux appears too bright for the model shown in Figure \ref{fig50}. ." +. 6213-07 is by far the best near-infrared excess candidate based on its 2A\LASS data. which appear reliable at all wavelengths.," G273-97 is by far the best near-infrared excess candidate based on its 2MASS data, which appear reliable at all wavelengths." + Llowever. the ΗΤΙ observations demonstrate that the 2ALASS catalog contains significant errors even al S/N 10..," However, the IRTF observations demonstrate that the 2MASS catalog contains significant errors even at S/N $>10$." +. Both Vossetal.(2007) and. Beauchampetal.(1999) give. dup=24500 Ww assuning no hvdrogen for this DBZ star.," Both \citet{vos07} and \citet{bea99} give $T_{\rm eff}\approx +24500$ K – assuming no hydrogen – for this DBZ star." + However. both authors eive alternative effective. temperatures near 22500 Ix. for a nominal hydrogen abundance.," However, both authors give alternative effective temperatures near 22500 K for a nominal hydrogen abundance." + The higher temperature overpredicts the fluxes for this star. and the fit shown in Figure A52. emplovs a temperature. of 21500 Ix to mateh all the photometry: possibly indicating some hvdrogen is present. ancl closer to the 19000 Ix. value eiven bv Koesteretal.(2005a).," The higher temperature overpredicts the fluxes for this star, and the fit shown in Figure \ref{fig52} employs a temperature of 21500 K to match all the photometry; possibly indicating some hydrogen is present and closer to the 19000 K value given by \citet{koe05a}." +. This white dwarf is another exemplary case of a 2MLASS excess not being corroborated by targeted. JLLdy photometry: the lack of infrared excess is also confirmed byIRAC observations (Farihietal.2009)., This white dwarf is another exemplary case of a 2MASS excess not being corroborated by targeted $JHK$ photometry; the lack of infrared excess is also confirmed by IRAC observations \citep{far09}. +. The white dwarf Uluxes in the 2ALASS catalog. were the xime motivation for obtaining follow up data. hence it is relevant to ask how well they predicted. the LYLE Ην xhotometry. ancl the presence of any near-infrared excess.," The white dwarf fluxes in the 2MASS catalog were the prime motivation for obtaining follow up data, hence it is relevant to ask how well they predicted the IRTF $JHK$ photometry, and the presence of any near-infrared excess." + Figure 1. compares the LTE and 2\LASS fluxes (in Vega magnitudes) for all 39 Table 20 white dwarl targets with thi sets of photomoetry: transformations between filter sets were ignored as they are generally within a few percent (Carpenter2001)., Figure \ref{fig53} compares the IRTF and 2MASS fluxes (in Vega magnitudes) for all 39 Table \ref{tbl2} white dwarf targets with both sets of photometry; transformations between filter sets were ignored as they are generally within a few percent \citep{car01}. +. The 2ALASS 1060 minimum detection limit for the whole skv is (CJ.H.IN.)=(15.8.15.1.14.3) mae (Skrutskieetal. 2006).. and there is decent accord for measurements brighter than this limit.," The 2MASS $10\sigma$ minimum detection limit for the whole sky is $(J,H,K_s)=(15.8,15.1, +14.3)$ mag \citep{skr06}, and there is decent accord for measurements brighter than this limit." + The generally clisagrecable behavior of the PALASS data below the 10σ limit in all three filters is not surprising. vet is most pronounced in the A band.," The generally disagreeable behavior of the 2MASS data below the $\sigma$ limit in all three filters is not surprising, yet is most pronounced in the $K$ band." + Somewhat unexpectedly. the J-band shows several discordant measures at the 12a level at relatively bright stars.," Somewhat unexpectedly, the $J$ -band shows several discordant measures at the $1-2\sigma$ level at relatively bright stars." + ln [act. of the co-observec white dwarfs at each bandpass. the //-band 2ALASS data agree most. frequently. with the LYE observations over all brightnesses.," In fact, of the co-observed white dwarfs at each bandpass, the $H$ -band 2MASS data agree most frequently with the IRTF observations over all brightnesses." + At the La level. the 2NLASS and LIEF photometry agree for 29 of 37 stars )) al 44. 25 of 38 stars (66%)) at J. and 1H of 29 stars (48%)) ab dy.," At the $1\sigma$ level, the 2MASS and IRTF photometry agree for 29 of 37 stars ) at $H$, 25 of 38 stars ) at $J$, and 14 of 29 stars ) at $K$." + An infrared excess is revealed by the relative [ux levels of three or more photometric Duxes. at leasttwo of which should. be consistent with photospheric emission ancl at least one of which is significantly higher than expected. for the photosphere alone.," An infrared excess is revealed by the relative flux levels of three or more photometric fluxes, at leasttwo of which should be consistent with photospheric emission and at least one of which is significantly higher than expected for the photosphere alone." + Usingthis practical definition. the PALASS photometry suggest an ff and/or A .-bandexcess relative to 2ALASS J or ultraviolet/optical photometry [ου 27 white dwarls in Table 2..," Usingthis practical definition, the 2MASS photometry suggest an $H$ and/or $K_s$ -bandexcess – relative to 2MASS $J$ or ultraviolet/optical photometry – for 27 white dwarfs in Table \ref{tbl2}. ." + Phe potential excess I[uxes are, The potential excess fluxes are +for each observed svstem for a pulsar population moclel.,for each observed system for a pulsar population model. +" Then. using Baves’ theorem. we ealeulate P(N4,)) from the likelihood P(1: Ny.) and eventually calculate PCR )) using a change of variables."," Then, using Bayes' theorem, we calculate $P$ ) from the likelihood $P$ ) and eventually calculate $P$ ) using a change of variables." + We repeat the whole procedure for all three observed coalescing binaries. and combine the (three individual PDFs to obtain a total PDF of Galactic coalescence rate of NS-WD binaries. ?(R4.4)).," We repeat the whole procedure for all three observed coalescing binaries, and combine the three individual PDFs to obtain a total PDF of Galactic coalescence rate of NS-WD binaries, $P$ )." + In paper I. we showed that a normalized PDF of the coalescence rate for an individual pulsar binary svstem can be written as follows: where C is a coellicient determined by properties of the 7! pulsar: llere. the beaming correction [actor fi is (he inverse of the traction of da sr covered bv the pulsar radiation beam during each rotation.," In paper I, we showed that a normalized PDF of the coalescence rate for an individual pulsar binary system can be written as follows: where $C_{\rm i}$ is a coefficient determined by properties of the $i^{\rm th}$ pulsar: Here, the beaming correction factor $f_{\rm b}$ is the inverse of the fraction of $4 \pi$ sr covered by the pulsar radiation beam during each rotation." +" In (hie case of the (wo DNS svstenis. PSRs D1913—16 and B15344+12. Ixalogeraοἱal.(2001) adopted. f,~6 based on pulse profile and polarization measurements of two pulsars."," In the case of the two DNS systems, PSRs B1913+16 and B1534+12, \citet{k01} adopted $f_b\sim6$ based on pulse profile and polarization measurements of two pulsars." + The lack of such observations for the current sample of binaries means (hat it is difficult to estimate reliable values of {ων, The lack of such observations for the current sample of binaries means that it is difficult to estimate reliable values of $f_b$. +" Therefore. in this paper. we do not correct for pulsar beaming (i.e. fj,= 1)."," Therefore, in this paper, we do not correct for pulsar beaming (i.e. $f_{\rm b}=$ 1)." + As a result. all our values should be considered as lower limits.," As a result, all our values should be considered as lower limits." + In paper I. we calculated PCRGJ) considering (wo observed DNSs systems (labeled by (hesubscripts 1 and 2).," In paper I, we calculated $P$ ) considering two observed DNSs systems (labeled by thesubscripts 1 and 2)." +" We delined the total rate A...=R4+Re and showed that where C, 88^{\circ}$ , $PA = 23^{\circ}$; Sofue $\&$ Nakai 1993)." +" NGC 891 is one of the major members of the 11023 group and its distance is estimated to be MMpe for fy=T5laus+Ape1 making for a scale of 22""/kpe or ppc/""."," NGC 891 is one of the major members of the 1023 group and its distance is estimated to be Mpc for $H_{\rm 0}=75\kms\,{\rm Mpc}^{-1}$, making for a scale of $''$ /kpc or $''$." + Many authors have remarked on the apparent similarity between NGC 891 and the Milky Way., Many authors have remarked on the apparent similarity between NGC 891 and the Milky Way. + The interstellar medium in the disk of NGC 891 has been the subject of several studies. notably radio continuum (Allen et al.," The interstellar medium in the disk of NGC 891 has been the subject of several studies, notably radio continuum (Allen et al." + 1978; Sukumar & Allen 1991). HI (Saneist & Allen 1979: Rupen 1991). dust (Howk & Savage 1997); CO (Garcfaa-Burillo et al.," 1978; Sukumar $\&$ Allen 1991), HI (Sancisi $\&$ Allen 1979; Rupen 1991), dust (Howk $\&$ Savage 1997); CO (Garcíaa-Burillo et al." + 1992: Scoville et al., 1992; Scoville et al. + 1993: Sofue & Nakai 1993; Garcfaa-Burillo & Guéllin. 1995; Sakamoto et al., 1993; Sofue $\&$ Nakai 1993; Garcíaa-Burillo $\&$ Guéllin 1995; Sakamoto et al. + 1997) and CI (Geri & Phillips 1997)., 1997) and CI (Gerin $\&$ Phillips 1997). + Far-infrared dust emission was observed by IRAS (see Wainscoat et al., Far-infrared dust emission was observed by IRAS (see Wainscoat et al. + 1987: Rice et al., 1987; Rice et al. + 1988). and at A = 1.3 mm by Guéllin et al.," 1988), and at $\lambda$ = 1.3 mm by Guéllin et al." + 1993., 1993. + In this paper we present further observations of the submillimeter dust emission from NGC 891 with. the new SCUBA submillimeter continuum array detector., In this paper we present further observations of the submillimeter dust emission from NGC 891 with the new SCUBA submillimeter continuum array detector. + The SCUBA 850 and 450 jim images were obtained in 1997 September on the JCMT'., The SCUBA 850 and 450 $\mu$ m images were obtained in 1997 September on the JCMT. +. SCUBA employs two arrays of cooled bolometers. each covering a field of about 2/3.," SCUBA employs two arrays of cooled bolometers, each covering a field of about $\mind 2.3$." + The respective filters used have central frequencies of 347 GHz and 677 GHz respectively: both filters have a bandwidth of 30 GHz., The respective filters used have central frequencies of 347 GHz and 677 GHz respectively; both filters have a bandwidth of 30 GHz. +" We observed three fields on NGC 891 (center. northeast and southwest) in socalled ""jiggle-mode' in order to cover the whole galaxy."," We observed three fields on NGC 891 (center, northeast and southwest) in socalled `jiggle-mode' in order to cover the whole galaxy." + Chopping distance was 2’ perpendicular to the major axis of the galaxy., Chopping distance was $'$ perpendicular to the major axis of the galaxy. + Total integration time for the three fields was 128 minutes., Total integration time for the three fields was 128 minutes. + Reduction was performed in the standard manner., Reduction was performed in the standard manner. + Photometric calibration was achieved by skydip analysis and mapping of the standard star CRL 618., Photometric calibration was achieved by skydip analysis and mapping of the standard star CRL 618. + More complete details of the instrument as well as observing and reduction methods can be found in Holland et al. (, More complete details of the instrument as well as observing and reduction methods can be found in Holland et al. ( +1998) and Jenness et al. (,1998) and Jenness et al. ( +1998). as well as on the world-wide web (www.jach.hawarn.edu),"1998), as well as on the world-wide web (www.jach.hawaii.edu)." + The resulting images are shown in Figure |., The resulting images are shown in Figure 1. + Especially in the 450;/m image. residual noise is seen at the SCUBA field edges: this should be discounted.," Especially in the $\mu$ m image, residual noise is seen at the SCUBA field edges; this should be discounted." + Guéllin et al. (, Guéllin et al. ( +1993) have fitted two dust components at 15 K and 30 K to their 1.3 mm flux and the IRAS 100j/m flux.,1993) have fitted two dust components at 15 K and 30 K to their 1.3 mm flux and the IRAS $\mu$ m flux. +" From our maps. we have determined integrated flux-densities 5550 = 4.8+0.6 Jy and 5,5)=3947 Jy."," From our maps, we have determined integrated flux-densities $S_{850}$ = $\pm$ 0.6 Jy and $S_{450} = 39\pm7$ Jy." + The observations through the 850m filter include emission from the ./23-2 ?CO line., The observations through the $\mu$ m filter include emission from the $J$ =3–2 $^{12}$ CO line. + In similar beams. the first three CO transitions occur in intensity ratios of 1:0.75:0.4 (see Garefaa-Burillo et al.," In similar beams, the first three CO transitions occur in intensity ratios of 1:0.75:0.4 (see Garcíaa-Burillo et al." + 1992; Gerin & Phillips 1997)., 1992; Gerin $\&$ Phillips 1997). +" From this. we estimate an integrated ,/23—-2 CO contribution of 0.22 Jy to the measured 850;/m flux-density."," From this, we estimate an integrated $J$ =3–2 CO contribution of 0.22 Jy to the measured $\mu$ m flux-density." + Corrected for this contribution. the dust continuum emission becomes κου% 4.6 Jy.," Corrected for this contribution, the dust continuum emission becomes $S_{850} \approx$ 4.6 Jy." + The 26-5 !?CO transition is at the edge of the 450j/m filter: we expect its contribution to the observed flux to be negligible., The $J$ =6–5 $^{12}$ CO transition is at the edge of the $\mu$ m filter; we expect its contribution to the observed flux to be negligible. + With Sy;yy = 0.73 Jy (Guéllin et al., With $S_{1300}$ = 0.73 Jy (Guéllin et al. + 1993). and the IRAS flux densities 54» = 6 Jy. 55; 2 7 Jy.," 1993), and the IRAS flux densities $S_{12}$ = 6 Jy, $S_{25}$ = 7 Jy," +function to model the number of halos of a given mass and redshift for our lens population.,function to model the number of halos of a given mass and redshift for our lens population. + For the internal mass distribution we assume that the region where the majority of strong lensing occurs can be modeled as an elliptical mass profile with a 3D density profile that falls as 1/r?., For the internal mass distribution we assume that the region where the majority of strong lensing occurs can be modeled as an elliptical mass profile with a 3D density profile that falls as $1/r^2$. +" This is an excellent approximation for galaxies (?),, while it is likely not a sufficiently complex model to capture the lensing properties of massive galaxy clusters (?).."," This is an excellent approximation for galaxies \citep{Koopmans:09}, while it is likely not a sufficiently complex model to capture the lensing properties of massive galaxy clusters \citep{Richard:10}." +" We use ray-tracing simulations to explore the impact of lens ellipticity and finite source sizes, assuming constant values of ellipticity and source sizes and exploring the impact of different assumed values."," We use ray-tracing simulations to explore the impact of lens ellipticity and finite source sizes, assuming constant values of ellipticity and source sizes and exploring the impact of different assumed values." +" In all that follows, we assume as our fiducial cosmology a spatially flat universe with 0,,= 0.222, ΠοT1.0kms-! Mpc-!, Ων=0.0449, n,=0.963 and og= 0.801."," In all that follows, we assume as our fiducial cosmology a spatially flat universe with $\Omega_m=0.222$ , $H_\circ=71.0 \,{\rm km\,s^{-1}\, Mpc^{-1}}$ , $\Omega_b=0.0449$, $n_s=0.963$ and $\sigma_8=0.801$ ." + The assumed unlensed source count models (the Durham and UBC have not been calibrated at mm wavelengths; small models)differences in parameters such as dust emissivity that are not large effects at submm wavelengths could lead to large misestimates at mm wavelengths., The assumed unlensed source count models (the Durham and UBC models) have not been calibrated at mm wavelengths; small differences in parameters such as dust emissivity that are not large effects at submm wavelengths could lead to large misestimates at mm wavelengths. + A simple check is to verify that the models produce a reasonable amount of noise power in mm-wave maps; this has been measured in SPT data in ?.., A simple check is to verify that the models produce a reasonable amount of noise power in mm-wave maps; this has been measured in SPT data in \citet{Hall:09}. +" We computed the angular noise power, defined for randomly distributed point sources as with a flux cutoff of 17 mJy (?) for SPT at 220 GHz "," We computed the angular noise power, defined for randomly distributed point sources as with a flux cutoff of 17 mJy \citep{Hall:09} for SPT at 220 GHz (1.4 mm)." +The majority of the noise power comes from sources(1.4 mm).well below the SPT sensitivity limit for detecting individual sources., The majority of the noise power comes from sources well below the SPT sensitivity limit for detecting individual sources. +" The noise power from the UBC model is in excellent agreement with the measurements, and did not require any corrections."," The noise power from the UBC model is in excellent agreement with the measurements, and did not require any corrections." +" The noise power derived from the Durham model, around 60 Jy? /sr, is more than a factor of 3 too high compared to ?.."," The noise power derived from the Durham model, around 60 $^2$ /sr, is more than a factor of 3 too high compared to \citet{Hall:09}." + We scale the flux of each object by a constant amount to match the measured noise power., We scale the flux of each object by a constant amount to match the measured noise power. + In Figure 1 it can be seen that the Durham and UBC models are forced by this constraint on the total power to have comparable number counts at fluxes of a few mJy., In Figure \ref{f1} it can be seen that the Durham and UBC models are forced by this constraint on the total power to have comparable number counts at fluxes of a few mJy. +" After applying a constant flux correction to match the observed noise power, the Durham model still had a problem with the properties of low-redshift dusty galaxies, in that the predicted number of low-redshift galaxies in the SPT sample was too high."," After applying a constant flux correction to match the observed noise power, the Durham model still had a problem with the properties of low-redshift dusty galaxies, in that the predicted number of low-redshift galaxies in the SPT sample was too high." +" In particular, the brightest of the low-z population in the Durham model by are predicted to be bright enough in the (observedSPT maps IRAS)to be found as sources."," In particular, the brightest of the low-z population in the Durham model (observed by IRAS) are predicted to be bright enough in the SPT maps to be found as sources." +" This is not the case, as evidenced by the small fraction of SPT-discovered galaxies that were also observed to be IRAS sources (?)/index.htmg(g"," This is not the case, as evidenced by the small fraction of SPT-discovered galaxies that were also observed to be IRAS sources \citep{Vieira:10}." +") As gravitational lensing is more efficient for high redshift sources, we elected to simply further suppress the flux of low-z galaxies (z« 0.2) to make the low-z galaxy counts agree with the number of SPT sources found to coincide with IRAS sources."," As gravitational lensing is more efficient for high redshift sources, we elected to simply further suppress the flux of low-z galaxies $z<0.2$ ) to make the low-z galaxy counts agree with the number of SPT sources found to coincide with IRAS sources." +" While not rigorously justifiable, the main point of using the Durham model was to get a plausible redshift distribution of sources at high-redshift."," While not rigorously justifiable, the main point of using the Durham model was to get a plausible redshift distribution of sources at high-redshift." + In order to avoid misinterpretations caused by the low-z population we only apply our lensing model to Durham counts with z>0.2 eliminating the low redshift counts and compare the results to SPT’s IRAS removed counts., In order to avoid misinterpretations caused by the low-z population we only apply our lensing model to Durham counts with $z>0.2$ eliminating the low redshift counts and compare the results to SPT's IRAS removed counts. + For the details of gravitational lensing theory we refer the reader to a review by ?.., For the details of gravitational lensing theory we refer the reader to a review by \citet{Bartelmann:01}. + Here we only briefly state a few lensing quantities that are used in this work., Here we only briefly state a few lensing quantities that are used in this work. +" The lens equation that we solve numerically is written as where Dg, D,, and Dg, are the angular diameter distances of deflector (d) and source (s) and @ is the observed position of a point source at 0 deflected by an angle à."," The lens equation that we solve numerically is written as where $D_d$, $D_s$, and $D_{ds}$ are the angular diameter distances of deflector (d) and source (s) and $\vec{\beta}$ is the observed position of a point source at $\vec{\theta}$ deflected by an angle $\vec{\alpha}$." + The total deflection from an ensemble of point masses at a single lens plane is given by The dimensionless surface mass density & is is the critical surface mass density and X(£) is the 2D projected mass density of thelens., The total deflection from an ensemble of point masses at a single lens plane is given by The dimensionless surface mass density $\kappa$ is is the critical surface mass density and $\Sigma(\vec{\xi})$ is the 2D projected mass density of thelens. +In Fig.,In Fig. + 1 the dotted and dashed lines are least. squares fits to the Ist and 2nd harmonics (with all parameters free. but omitting the 2nd harmonic points with Zo 1.15 d) and appear again in the upper part of the diagram multiplied bv two and three respectively to show the approximate evolution of the implied fundamental.," \ref{dno4fig1} the dotted and dashed lines are least squares fits to the 1st and 2nd harmonics (with all parameters free, but omitting the 2nd harmonic points with $T$ $>$ 1.15 d) and appear again in the upper part of the diagram multiplied by two and three respectively to show the approximate evolution of the implied fundamental." + Fig., Fig. + 3 shows a moclified form of Fig. L..," \ref{dno4fig3} shows a modified form of Fig. \ref{dno4fig1}," + where we have now transformed the Ist and 2nd harmonics into the implied Fundamental to show its evolution more clearly., where we have now transformed the 1st and 2nd harmonics into the implied fundamental to show its evolution more clearly. + This shows that the fundamental systematically increases in period )— 105 s before the DNOs disappear completely., This shows that the fundamental systematically increases in period to $\sim$ 105 s before the DNOs disappear completely. + When passing through Nb. — 95 s any appearance of the fundamental in the ET may be confused with the IpDNOs in VW Livi that occur in that. period range., When passing through 85 – 95 s any appearance of the fundamental in the FT may be confused with the lpDNOs in VW Hyi that occur in that period range. + However. the IpDNOs observed so [ar (in. VW. Livi and other CVs) are all pure sinusoids and often long-lasting. so most of them can be recognised for what they are and have been omitted. from Fig. 1..," However, the lpDNOs observed so far (in VW Hyi and other CVs) are all pure sinusoids and often long-lasting, so most of them can be recognised for what they are and have been omitted from Fig. \ref{dno4fig1}." + In VW Livi it therefore appears that at times the fundamental period of the normal DNOs can brielly reach values greater than the period of the IpDNOs. but the tendeney for the 2nd harmonic to stabilize near 30 s at the end of the whole DNO evolution is a strong indication that Pox=Ppxo at the end of outburst.," In VW Hyi it therefore appears that at times the fundamental period of the normal DNOs can briefly reach values greater than the period of the lpDNOs, but the tendency for the 2nd harmonic to stabilize near $\sim$ 30 s at the end of the whole DNO evolution is a strong indication that $P_{DNO} \rightarrow P_{lpDNO}$ at the end of outburst." + The dashed line in Fig., The dashed line in Fig. + 3. is the least squares fit for T — (0.2 to 1.15 d and has the equation The formal errors on the parameters derived. from the least squares fit are elven in Iq., \ref{dno4fig3} is the least squares fit for $T$ = 0.2 to 1.15 d and has the equation The formal errors on the parameters derived from the least squares fit are given in Eq. + 1 in brackets.," \ref{dno4eq1} + in brackets." + 1n order to obtain greater sensitivity to the fast. changing periods we have calculated amplitude and OC values. fittingm sinusoids by [east squares1 to short and usually overlapping sections of the light) curves.," In order to obtain greater sensitivity to the fast changing periods we have calculated amplitude and O–C values, fitting sinusoids by least squares to short and usually overlapping sections of the light curves." + These are, These are +beam were missing our line of sight (Beloborodov|2002).,beam were missing our line of sight \citep{belo02}. +". We are specifically motivated by the X-ray detection of all known MSPs in the globular cluster 47 Tuc (Heinke[Bogdanov allowing predictions o the X-ray emission of etother al.|2006),,MSPs."," We are specifically motivated by the X-ray detection of all known MSPs in the globular cluster 47 Tuc \citep{Heinke05a,Bogdanov06}, allowing predictions of the X-ray emission of other MSPs." +" We choose this sample of MSPs for comparison in part because the distance to globular clusters such as 47 Tuc (4.5kpc,2003updateof are better known than the distances to most MSPs."," We choose this sample of MSPs for comparison in part because the distance to globular clusters such as 47 Tuc \citep[4.5 kpc, 2003 update of][]{harris96} are better known than the distances to most MSPs." + We discuss the significance of our nondetections and conclude in Section ??.. J0, We discuss the significance of our nondetections and conclude in Section \ref{concl}. +917--4638 was observed with the GBT on 2007 November 30., $+$ 4638 was observed with the GBT on 2007 November 30. + The observing set-up and data reduction were the same as described in (2009).., The observing set-up and data reduction were the same as described in \citet{agueros09}. . +" At 820 MHz, the Berkeley-Caltech Pulsar Machine (Backer provided 48 MHz of bandwidth split into 96 spectral channels; total power samples for each channel were recorded every 72s. The total observing time was 13,300 s "," At 820 MHz, the Berkeley-Caltech Pulsar Machine \citep[][]{backer97} provided 48 MHz of bandwidth split into 96 spectral channels; total power samples for each channel were recorded every $72\,\mu$ s. The total observing time was $13,300$ s $3.7$ hr)." +We used standard pulsar search techniques as (3.7implementedhr). in the PRESTO software −− , We used standard pulsar search techniques as implemented in the PRESTO software package \citep{ransom01}. +We calculated the maximum packagedispersion measure (DM) expected in the direction of J0917--4638 using the model for the distribution of free electrons in the Galaxy., We calculated the maximum dispersion measure (DM) expected in the direction of $+$ 4638 using the \citet{cordes02} model for the distribution of free electrons in the Galaxy. +" To account for uncertainties, we dedispersed the data up to a DM limit twice that obtained from the model, i.e., DM =80 cm? pc."," To account for uncertainties, we dedispersed the data up to a DM limit twice that obtained from the model, i.e., DM $= 80$ $^{-3}$ pc." + No convincing pulsar signal was detected in our data., No convincing pulsar signal was detected in our data. + Below we discuss the limitations of our search., Below we discuss the limitations of our search. + The orbital motion of a putative pulsar companion to J0917--4638 could significantly affect its apparent spin period., The orbital motion of a putative pulsar companion to $+$ 4638 could significantly affect its apparent spin period. +" Based on radial velocity monitoring, found that J0917+4638 is in an orbit with a period 7.6 hr."," Based on radial velocity monitoring, \citet{kilic07b} found that $+$ 4638 is in an orbit with a period $7.6$ hr." +" Assuming that the LMWD companion is a 1.4Mo NS, this implies that the maximum orbital acceleration is on the order of 100 m s-?, which is significantly larger than what is typically seen in these systems (for9096ofknownpul-"," Assuming that the LMWD companion is a $1.4\ M_\odot$ NS, this implies that the maximum orbital acceleration is on the order of $100$ m $^{-2}$, which is significantly larger than what is typically seen in these systems \citep[for $90\%$ of known pulsars the maximum orbital acceleration is $\leq |25|$ m s$^{-2}$ ." + Our integration time represents nearly half of the binary orbital period., Our integration time represents nearly half of the binary orbital period. +" As a result, the assumption of a constant apparent acceleration built into PRESTO breaks down."," As a result, the assumption of a constant apparent acceleration built into PRESTO breaks down." +" We therefore divided our GBT data into 14 separate 900 s integrations representing ~3% of an orbit) and one 700 s integration(each and conducted searches for pulsations separately in each of these partial This extended our search sensitivity to accelerations on the order of several hundred m s~?, but as detailed in the following section, reduced our luminosity sensitivity."," We therefore divided our GBT data into $14$ separate $900$ s integrations (each representing $\sim3\%$ of an orbit) and one $700$ s integration and conducted searches for pulsations separately in each of these partial This extended our search sensitivity to accelerations on the order of several hundred m $^{-2}$, but as detailed in the following section, reduced our luminosity sensitivity." + None of these searches uncovered a convincing pulsar signal., None of these searches uncovered a convincing pulsar signal. + We use the standard modifications to the radiometer equation to calculate the minimum detectable period-averaged flux density for our searches., We use the standard modifications to the radiometer equation to calculate the minimum detectable period-averaged flux density for our searches. + We consider a pulsar duty cycle of 20% (typical of MSPs)., We consider a pulsar duty cycle of $20\%$ (typical of MSPs). +" At 820 MHz, the GBT gain is JJy—! and the system temperature is 25KK. The sky temperature at this frequency and a Galactic latitude of b=+44° only adds a few K to the overall temperature."," At $820$ MHz, the GBT gain is $^{-1}$ and the system temperature is K. The sky temperature at this frequency and a Galactic latitude of $b =+44^{\circ}$ only adds a few K to the overall temperature." + We consider an effective threshold signal-to-noise ratio of 10., We consider an effective threshold signal-to-noise ratio of 10. +" For ting=900 s, the sensitivity limit for a long period pulsar at the beam center is ~0.26 muy."," For $t_{int} = 900$ s, the sensitivity limit for a long period pulsar at the beam center is $\sim$ 0.26 mJy." +" Pulsar luminosities are often measured at MMHz; using a typical spectral index of —1.7, the limiting sensitivity at that frequency is 14007:0.10 mJy when searching the 900 s integrations."," Pulsar luminosities are often measured at MHz; using a typical spectral index of $-1.7$ , the limiting sensitivity at that frequency is $S_{1400} \approx 0.10$ mJy when searching the $900$ s integrations." +" For an MSP period of 3 ms, our sensitivity at 1400 MHz was roughly 0.14 mJy for each integration,and it quickly degraded for shorter periods; it was 10x worse for 1 ms."," For an MSP period of 3 ms, our sensitivity at 1400 MHz was roughly $0.14$ mJy for each integration,and it quickly degraded for shorter periods; it was $10\times$ worse for 1 ms." +" The distance to J09174-4638 is estimated to be 2.3 kpc (Kilicοἱal}/2007b),, implying that our Ly490=Sta00d? limits for 3mms periods are zz0.7 mJy kpc? for each 900 s integration."," The distance to $+$ 4638 is estimated to be $2.3$ kpc \citep{kilic07b}, implying that our $L_{1400} \equiv S_{1400} d^2$ limits for ms periods are $\approx\ 0.7$ mJy $^2$ for each $900$ s integration." +" According to the ATNF’s pulsar ⊓ of the 50MSPs (periods <25 ms) outside of globular clusters and with measured luminosities, 64% have L499>0.7 mJy kpc?."," According to the ATNF's pulsar \citep{atnf}, of the $50$MSPs (periods $<25$ ms) outside of globular clusters and with measured luminosities, $64\%$ have $L_{1400} > 0.7$ mJy $^2$." + We would therefore expect our search to detect roughly two-thirds of the known MSPs were one orbiting J09174-4638 and beaming radio waves toward the Earth., We would therefore expect our search to detect roughly two-thirds of the known MSPs were one orbiting $+$ 4638 and beaming radio waves toward the Earth. +" We note that J0917+4638 falls within the FIRST footprint and is not detected in that 1.4 GHz survey, for which the sensitivity limitis roughly 1 mJy 1995).."," We note that $+$ 4638 falls within the FIRST footprint and is not detected in that $1.4$ GHz survey, for which the sensitivity limitis roughly $1$ mJy \citep{first}. ." +" MSP radio beaming fractions are < 100%, and as a result, some MSPs have not yet been detected in the radio in binary systemswhere there is strong evidence"," MSP radio beaming fractions are $< 100\%$ , and as a result, some MSPs have not yet been detected in the radio in binary systemswhere there is strong evidence" +and weaker Ls lines compared to this study.,and weaker $_2$ lines compared to this study. + In addition the relatively strong Fell] emission found. in this study is not observed in starbursts., In addition the relatively strong [FeII] emission found in this study is not observed in starbursts. + Phe excitation temperature of the Lo lines is similar to that found in this study and. others (see Wilman et al., The excitation temperature of the $_2$ lines is similar to that found in this study and others (see Wilman et al. + 2002) but the properties of these two classes of objects are significantly dillerent in the near-infrared., 2002) but the properties of these two classes of objects are significantly different in the near-infrared. + From the properties of the optical emission lines. central cluster galaxies share many properites with LINIZIs.," From the properties of the optical emission lines, central cluster galaxies share many properites with LINERs." + Infra-red. spectroscopy of LINERs by Larkin et ((1998) finds strong molecular lines and a small fraction. with high ionization lines so the resemblance follows into the infra-red as noted by Jalfe et ((2001)., Infra-red spectroscopy of LINERs by Larkin et (1998) finds strong molecular lines and a small fraction with high ionization lines so the resemblance follows into the infra-red as noted by Jaffe et (2001). + This comparison also applies to the Fell] strength as LINERs (unlike starbursts) strong in. Fell] and OI] and weak in Bre (Larkin et 1998)., This comparison also applies to the [FeII] strength as LINERs (unlike starbursts) strong in [FeII] and [OI] and weak in $\gamma$ (Larkin et 1998). + The question of whether the underlying excitation mechanism behind both classes of object is the same is an open one., The question of whether the underlying excitation mechanism behind both classes of object is the same is an open one. + The bulk of the literature on LINERs favours X-rav heating and/or shocks and this is discussed. further in Wilman et ((2002)., The bulk of the literature on LINERs favours X-ray heating and/or shocks and this is discussed further in Wilman et (2002). + Ina few of the lower redshift galaxies our spectra cover the CO stellar absorption bandhead at 2.3/2. The strength of this feature should. provide constraints of the relative mass distribution on the underlving population of stars in the ealaxies., In a few of the lower redshift galaxies our spectra cover the CO stellar absorption bandhead at $\mu$ m. The strength of this feature should provide constraints of the relative mass distribution on the underlying population of stars in the galaxies. + Using the definition of the CO equivalent width of James Mobasher (2000). we calculate values of 3.57. 3.27. 3.96. 3.21. 3.33. 3.45 and 2.89nm. with errors of 0.30nm. for 1UXJ0338|096. LIvelra-X. ASS. ALT95. A2029.. A2052. anc A2199 respectively (only the latter object is in the James Mobasher (2000) sample at 3.23-0.25nm. so. consisten within the errors).," Using the definition of the CO equivalent width of James Mobasher (2000), we calculate values of 3.57, 3.27, 3.96, 3.21, 3.33, 3.45 and 2.89nm, with errors of 0.30nm, for RXJ0338+096, Hydra-A, A85, A1795, A2029, A2052 and A2199 respectively (only the latter object is in the James Mobasher (2000) sample at $\pm$ 0.25nm so consistent within the errors)." + The mean equivalent width of €O οἱ our seven values is 3.38nm2 compared to 3.35nm for the brightest cluster galaxies in James Alobasher (2000) so our results are consistent despite the redshift distribution anc resolution of our spectra not being well suited to studying these lines., The mean equivalent width of CO of our seven values is 3.38nm compared to 3.35nm for the brightest cluster galaxies in James Mobasher (2000) so our results are consistent despite the redshift distribution and resolution of our spectra not being well suited to studying these lines. + The K-band spectra obtained in ths study are dominated by he old stellar population., The K-band spectra obtained in ths study are dominated by the old stellar population. + Llowever. in the case of NCGC1275 here is à strong contribution from the active nucleus which shows a significant upturn at 2yam. This is in part due to he presence of hot dust. (Ixrabbe et 22000).," However, in the case of NGC1275 there is a strong contribution from the active nucleus which shows a significant upturn at $\mu$ m. This is in part due to the presence of hot dust (Krabbe et 2000)." + To test or any possible non-stellar contribution to our spectra we aave determined the spectral slope from 1.9 to 2.15 (rest) or all spectra (excluding emission lines), To test for any possible non-stellar contribution to our spectra we have determined the spectral slope from 1.9 to $\mu$ m (rest) for all spectra (excluding emission lines). + The results are oesented in Figure S plotted against redshift., The results are presented in Figure 8 plotted against redshift. + Only one, Only one +of the [n]giant brauch to its j»osition on the horizontal brauch. the eiaut. branch evolution essentially rus lu reverse.,"of the giant branch to its position on the horizontal branch, the giant branch evolution essentially runs in reverse." + The surface of the star contracts. aud the core of the star expaucls. flattening he rotational profile and raisiug the surlace rotation rate by a factor of 10.," The surface of the star contracts, and the core of the star expands, flattening the rotational profile and raising the surface rotation rate by a factor of $\sim$ 10." + The horizontal bratich star ds 1οἱ rotating as a solid bods. but it is much closer to solid body rotation than the higM7 clifferelally rotating giat brauch starS.," The horizontal branch star is not rotating as a solid body, but it is much closer to solid body rotation than the highly differentially rotating giant branch stars." + yire 2 shows tle impact of mass loss on the evolution of the moment of inertia of πα a& the elant brat(d., Figure 2 shows the impact of mass loss on the evolution of the moment of inertia of the star along the giant branch. + We have plotted total moment of inertia of the star as a funcli ol logCL/L..) for evolttioary tracks witl different. amounts of total ass loss., We have plotted total moment of inertia of the star as a function of $\log (L/L_{\sun})$ for evolutionary tracks with different amounts of total mass loss. + Low on he [n]gi ranch. he star does 100 lose much 1aass in each timestep. and so the stars are not significal different Grom each «Aler.," Low on the giant branch, the star does not lose much mass in each timestep, and so the stars are not significantly different from each other." + ear tlie tip of he giant brauch. however. tlie amouit of mass [9]st dr he star increases. Wwnri reduces the size o‘the giaut couvection zone. resulting in a sinaller Star ald rence a sinaller 1201jent oO ‘inertia than p‘eclicted from non-mass-losiug; uioes.," Near the tip of the giant branch, however, the amount of mass lost from the star increases, which reduces the size of the giant convection zone, resulting in a smaller star and hence a smaller moment of inertia than predicted from non-mass-losing models." + The evoluliorary racks used iu tliis ptiper represent an imp'ovenient over the work done in (1991 in which mass loss was not included in the evolutionary tracks. aud all 'esults were based o ratrack like the 0.8 AZ. track shown in figure 2.," The evolutionary tracks used in this paper represent an improvement over the work done in \cite{PDD}, in which mass loss was not included in the evolutionary tracks, and all results were based on a track like the 0.8 $M_{\sun}$ track shown in figure 2." + Uuder (he assumption that he moment of inertla was not significantly affected by the mass loss. Pinso1eault.Delivanuis(1991). predicted too little angular momentπα loss. particularly fo: the hotter horizontal yanch stars.," Under the assumption that the moment of inertia was not significantly affected by the mass loss, \cite{PDD} predicted too little angular momentum loss, particularly for the hotter horizontal branch stars." + We now examiue the impact of different assumptions about internal auglar moljenun trausport alter the uain sequence turnolL the rotation profile eulorced tu couvective regio1s. auc internal angular 1omentum transport during the helium flash aud on tle horizoutal branel.," We now examine the impact of different assumptions about internal angular momentum transport after the main sequence turnoff, the rotation profile enforced in convective regions, and internal angular momentum transport during the helium flash and on the horizontal branch." + In Figures ὁκ)+) aid [we illustrate the specific aieular luolnnelntlilunu as allncjon of 1lass al different. epochs for two different choices of the rotatjon law in convect]ive regiOLS., In Figures 3 and 4 we illustrate the specific angular momentum as a function of mass at different epochs for two different choices of the rotation law in convective regions. + In eacl1 ligt the top pauel corresj»ouds to the case of local augular nonmentunmi conservatjon in tle raciatiνο cor while tlie bottom se of panels correspotds to the case of solid body rotatior al 1he angular «'eloci of the base of the sulace convection zc16., In each figure the top panel corresponds to the case of local angular momentum conservation in the radiative core while the bottom set of panels corresponds to the case of solid body rotation at the angular velocity of the base of the surface convection zone. +" We have picked a reference 10‘ivoutal brauch Lass 0.6 δι. COLLESPOleing to an ellective eimperature oun the horizonal b""alich «X c 10 000 Ix a total mass loss on tle giant. brauc ιοί)2 AM.."," We have picked a reference horizontal branch mass of 0.6 $M_{\sun}$, corresponding to an effective temperature on the horizontal branch of $\sim$ 10 000 K and total mass loss on the giant branch of 0.2 $M_{\sun}$." + Wetow cdiscuss results for solid body rotation i COLvective regious (Figure 3) aud iuiforu specific auglar momentum in colvecive regious (Fig l1) in turn., We now discuss results for solid body rotation in convective regions (Figure 3) and uniform specific angular momentum in convective regions (Figure 4) in turn. + The assumption of uniform roatlor In convective 'eelons implies that tle matter at the surface wil have relatively Ligh specific augular momentum: as a result. mass loss will ellectively draiu angular momentuim from the enveope {Figure 3+)).," The assumption of uniform rotation in convective regions implies that the matter at the surface will have relatively high specific angular momentum; as a result, mass loss will effectively drain angular momentum from the envelope (Figure 3)." + Furhermore. the base o ‘the surface convection ZOLe will have very low specific angular momentunm: tlis Iniuimizes the angular momentum coutent of he radiative core even if the core angular momentii is not redistributed to the envelope.," Furthermore, the base of the surface convection zone will have very low specific angular momentum; this minimizes the angular momentum content of the radiative core even if the core angular momentum is not redistributed to the envelope." +(NA10; Giordanoetal.2010;Masters 2011)).,"(NA10; \citealp{giordano10,masters11}) )." + Figure 6 shows the dependence of the bar fraction on redshift., Figure \ref{zdepend} shows the dependence of the bar fraction on redshift. +" We find that the fraction of SB1 galaxies (fgpi) is approximately constant, ~25%, until z=0.045, but that it decreases to ~20% at z>0.045."," We find that the fraction of SB1 galaxies $\bfrsbo$ ) is approximately constant, $\sim25\%$, until $z=0.045$, but that it decreases to $\sim20\%$ at $z>0.045$." +" On the other hand, the fraction of SB2 galaxies (faspg2) vary insignificantly from to as the redshift increases."," On the other hand, the fraction of SB2 galaxies $\bfrsbw$ ) vary insignificantly from to as the redshift increases." +" However, even if we restrict our analysis to 5,888 galaxies at z«0.045, we find that there is no significant change in the following results."," However, even if we restrict our analysis to 5,888 galaxies at $z<0.045$, we find that there is no significant change in the following results." +" We investigate the dependence of fa, on several physical parameters of galaxies: r-band absolute magnitude, u—r color, g—i color gradient, equivalent width of the Ha line i-band inverse concentration index (cin=(W(Ho)) Rso/Ro0), and central velocity dispersion (c)."," We investigate the dependence of $\bfr$ on several physical parameters of galaxies: $r$ -band absolute magnitude, $u-r$ color, $g-i$ color gradient, equivalent width of the $H\alpha$ line $W(H\alpha)$ ), $i$ -band inverse concentration index $c_{in}=R_{50}/R_{90}$ ), and central velocity dispersion $\sigma$ )." + Figure 7 presents how the fractions of SBl and SB2 galaxies change with variation of several physical parameters., Figure \ref{params_bf} presents how the fractions of SB1 and SB2 galaxies change with variation of several physical parameters. +" In the case of spectroscopic parameters such as W(Ha) and the central velocity dispersion, we measure[να using data satisfying signal-to-noise ratio S/N>10."," In the case of spectroscopic parameters such as $W(H\alpha)$ and the central velocity dispersion, we measure$\bfr$ using data satisfying signal-to-noise ratio $S/N\ge10$." + We estimate the errors of να by calculating the standard deviation in times-repetitive sampling method.," We estimate the errors of $\bfr$ by calculating the standard deviation in 1,000-times-repetitive sampling method." + Figure 7aa shows να as a function of u—r color., Figure \ref{params_bf}a a shows $\bfr$ as a function of $u-r$ color. + The fraction of SB1 galaxies increases significantly as wu—r increases., The fraction of SB1 galaxies increases significantly as $u-r$ increases. +" The SB1 fraction is less than 20% for galaxies bluer than u—r2.0, but it reaches 4596 when u—rc2.8."," The SB1 fraction is less than $20\%$ for galaxies bluer than $u-r=2.0$, but it reaches $45\%$ when $u-r\simeq2.8$." + Then it decreases at the red color end., Then it decreases at the red color end. + 'This result implies that passively evolving red late-type galaxies are more likely to have a strong bar than blue spirals with some star formation activity., This result implies that passively evolving red late-type galaxies are more likely to have a strong bar than blue spirals with some star formation activity. +" In contrast, the fraction of SB2 galaxies (fsp2) appears to become the maximum when u—r~ 1.4."," In contrast, the fraction of SB2 galaxies $\bfrsbw$ ) appears to become the maximum when $u-r\simeq1.4$ ." +" Then, fgp2 becomes smaller as u—r increases."," Then, $\bfrsbw$ becomes smaller as $u-r$ increases." +" These trends are consistent with the result of Hoyleetal.(2011) who found that longer bars than 5h57! kpc prefer redder late-type galaxies, and that shorter bars inhabit bluer late-type galaxies."," These trends are consistent with the result of \citet{hoyle+11} who found that longer bars than $h^{-1}$ kpc prefer redder late-type galaxies, and that shorter bars inhabit bluer late-type galaxies." +" In Figure 7bb, we plot fj, as a function of r-band absolute magnitude."," In Figure \ref{params_bf}b b, we plot $\bfr$ as a function of $r$ -band absolute magnitude." +" It is seen that fspgi monotonically increases as M,. decreases until it reaches the maximum at M,~—21.2 mag, and then slightly drops at the brightest magnitudes."," It is seen that $\bfrsbo$ monotonically increases as $M_r$ decreases until it reaches the maximum at $M_{r}\simeq-21.2$ mag, and then slightly drops at the brightest magnitudes." +" On the other hand, the SB2 fraction shows no dependence on M,.."," On the other hand, the SB2 fraction shows no dependence on $M_r$." + Méndez-Abreuet also found a similar result using ~190 galaxies in the (2010)Coma cluster., \citet{mabreu10} also found a similar result using $\sim$ 190 galaxies in the Coma cluster. +" In Figure 7cc, we plot the relation between fj, and Ho equivalent width."," In Figure \ref{params_bf}c c, we plot the relation between $\bfr$ and $H\alpha$ equivalent width." +" is an indicator of star formation rate (SFR) in the W(Ha)central region (R« 1"".5) of galaxies.", $W(H\alpha)$ is an indicator of star formation rate (SFR) in the central region $R<1^{\prime\prime}.5$ ) of galaxies. + There were some reports that bars induce central starburst (Hunt&Malkan1999;Eskridgeet 2007)..," There were some reports that bars induce central starburst \citep{hunt99,esk00,jogee05,maj07}. ." +" Therefore,it is expected that [να is higher for galaxies"," Therefore,it is expected that $\bfr$ is higher for galaxies" +separated by some hundred days.,separated by some hundred days. +" All observations of that were public as of 2009 December 01 were used as the data basis for this work, see Table 1.."," All observations of that were public as of 2009 December 01 were used as the data basis for this work, see Table \ref{tab:blocks}." +" While data from individual observation blocks have previously been analyzed in detail (Kreykenbohmetal.,2008;Schanneetal., 2007),, studying all available data allowed us to carry out a deep statistical analysis of the flaring behavior of Vela X-"," While data from individual observation blocks have previously been analyzed in detail \citep{kreykenbohm08a, schanne07a}, studying all available data allowed us to carry out a deep statistical analysis of the flaring behavior of Vela X-1." +" Our analysis was performed with special regard to the giant flares with peak luminosities of more than six times the average luminosity (Staubertetal.,2004;Kreykenbohm2008),, i.e., more than ccps in ISGRI (~1.3 CCrab) between kkeV. It was not clear if these giant flares can be explained by the same physical process as the other, common, flares in or if they are caused by a different mechanism."," Our analysis was performed with special regard to the giant flares with peak luminosities of more than six times the average luminosity \citep{staubert04a, kreykenbohm08a}, i.e., more than cps in ISGRI $\sim$ Crab) between keV. It was not clear if these giant flares can be explained by the same physical process as the other, common, flares in or if they are caused by a different mechanism." +" For this work, only ISGRI data between kkeV were analyzed, because in this energy range the source is bright and the flux is unaffected by photoabsorption, because typical values of the equivalent hydrogen column Ng outside the eclipse are below 3x10? ccm""? (Kreykenbohmetal.,1999)."," For this work, only ISGRI data between keV were analyzed, because in this energy range the source is bright and the flux is unaffected by photoabsorption, because typical values of the equivalent hydrogen column $N_\text{H}$ outside the eclipse are below $3\times 10^{23}$ $^{-2}$ \citep{kreykenbohm99a}." +". Our analysis extends the energy range of prior observations of the stellar wind, which used high-resolution spectra in soft rays to investigate the features of the wind (Watanabeetal.,2006,andreferences therein).."," Our analysis extends the energy range of prior observations of the stellar wind, which used high-resolution spectra in soft X-rays to investigate the features of the wind \citep[and references therein]{watanabe06a}." +" We extracted lightcurves from ISGRI usinglight, a tool distributed as part of the Offline Scientific Analysis (OSA) 7.0 package."," We extracted lightcurves from ISGRI using, a tool distributed as part of the Offline Scientific Analysis (OSA) 7.0 package." + The time resolution was chosen to be ssec to average each data-point over one pulse period to eliminate these fluctuations., The time resolution was chosen to be sec to average each data-point over one pulse period to eliminate these fluctuations. +" The pulse period is changing in a random-walk behavior on all times scales (Deeteretal.,1989),, but no long-term trend is evident, so that the chosen time resolution is sufficient for the purpose of this investigation."," The pulse period is changing in a random-walk behavior on all times scales \citep{deeter89a}, but no long-term trend is evident, so that the chosen time resolution is sufficient for the purpose of this investigation." + As an example for the analyzed data the upper panel of Fig., As an example for the analyzed data the upper panel of Fig. + 1 shows the lightcurve of the kkeV band for data from revolutions 433—440 (Block 3)., \ref{fig:lc_20-60_4xx.pdf} shows the lightcurve of the keV band for data from revolutions 433–440 (Block 3). +" The lightcurve data were filtered according to orbital phase Qo», leaving only data of phases between 0.19 €down< 0.81, for which the system is out of eclipse."," The lightcurve data were filtered according to orbital phase $\phi_\text{orb}$, leaving only data of phases between 0.19 $\le \phi_\text{orb} \le$ 0.81, for which the system is out of eclipse." + These data were then binned into 256 count rate bins., These data were then binned into 256 count rate bins. + The bins are spaced logarithmically between 1ccps and 1000ccps., The bins are spaced logarithmically between cps and cps. +" This binning leads to a histogram of the orbital phase averaged brightness distribution of the source, shown in black in Fig. 2.."," This binning leads to a histogram of the orbital phase averaged brightness distribution of the source, shown in black in Fig. \ref{fig:lcratehistall.ps}." +" The distribution closely resembles a normal distribution in log-space, i.e., a log-normal distribution (Fiirstetal.,2008).."," The distribution closely resembles a normal distribution in log-space, i.e., a log-normal distribution \citep{fuerst08b}." +" In order to quantify the shape of the distribution we fitted a Gaussian function in log-space, i.e., a log-normal distribution in countrate space."," In order to quantify the shape of the distribution we fitted a Gaussian function in log-space, i.e., a log-normal distribution in countrate space." +" Following the approach by Uttleyetal.(2005) the uncertainties of the values N; of each histogram data bin were estimated to be Nj, i.e., assuming Poisson statistics."," Following the approach by \citet{uttley05a} the uncertainties of the values $N_i$ of each histogram data bin were estimated to be $\sqrt{N_i}$, i.e., assuming Poisson statistics." + We show that this assumption is justified in Sect. 4.., We show that this assumption is justified in Sect. \ref{sec:simu}. +" The best-fit function, with a y?-value of 190 for 81 degrees of freedom, is shown as solid line in Fig. 2,,"," The best-fit function, with a $\chi^2$ -value of $190$ for 81 degrees of freedom, is shown as solid line in Fig. \ref{fig:lcratehistall.ps}," +" where in the fit only bins with more than 20 measurements were taken into account (seeGehrels,1986)..", where in the fit only bins with more than 20 measurements were taken into account \citep[see][]{gehrels86a}. + A Gaussian function in log-space is best characterized by its mean and standard deviation., A Gaussian function in log-space is best characterized by its mean and standard deviation. +"These values represent the median of the log-normal function, which will be denoted as (x), and its multiplicative standard deviation, 6, i.e., of all data points fall in the interval [(x)/6.(x)-6] ","These values represent the median of the log-normal function, which will be denoted as $\left$ , and its multiplicative standard deviation, $\tilde \sigma$ , i.e., of all data points fall in the interval $[\left< x\right> / \tilde \sigma, \left< x\right> \cdot \tilde \sigma]$ " +q lor 3x/22. and found that G(/:q) have the similar morphology for />4. which is Clearly shown in Fie. 3.. (,"$\hat{\mathbf q}$ for $3\le l \le 22$, and found that $G(l;\hat{\mathbf q})$ have the similar morphology for $l \ge 4$, which is clearly shown in Fig. \ref{Glq}. (" +Note that. the morphology. of /=3 map is different. which may relate to the unsolved low quadrupole problem as well as the alignment of equacdrupole ancl octupole (Bennettetal. 2011)..),"Note that, the morphology of $l=3$ map is different, which may relate to the unsolved low quadrupole problem as well as the alignment of quadrupole and octupole \citep{wmap_anomly}. .)" + In Table 1.. we list the preferred directions q. where the paritv parameter G(fiq) for each / is minimized (note Chat different from the problem in (Gubitosietal.2011).. here as the widely discussed. alignment. problem of quadrupole aud octupole in (deOliveira-Costaetal.2006).. the uncertainties of the preferred directions are difficult to be defined). which are very close with each other for />4.," In Table \ref{table1}, we list the preferred directions $\hat{\bf q}$, where the parity parameter $G(l;\hat{\bf q})$ for each $l$ is minimized (note that different from the problem in \citep{aaaaa}, here as the widely discussed aligment problem of quadrupole and octupole in \citep{dt2}, the uncertainties of the preferred directions are difficult to be defined), which are very close with each other for $l\ge 4$." + So. we can choose the special direction q (note that —q is another equivalent preferred. direction). where all the parameters G(/:q) are minimized or maximized.," So, we can choose the special direction $\hat{\bf q}$ (note that $-\hat{\bf q}$ is another equivalent preferred direction), where all the parameters $G(l;\hat{\mathbf q})$ are minimized or maximized." + We picked out these regions. and plotted them in the Galactic coordinate svstem in Fig. 4..," We picked out these regions, and plotted them in the Galactic coordinate system in Fig. \ref{dip7}." + It is interesting to find (hat the preferred directions q. where parity violation is largest. are coincident with the WAIAP? KD direction (Jarosiketal.2011).. while the preferred directions q. where parity asvinmnmetry is suiallest. are nearly perpendicular to (he ND direction.," It is interesting to find that the preferred directions $\hat{\bf q}$, where parity violation is largest, are coincident with the WMAP7 KD direction \citep{wmap7_dipole}, while the preferred directions $\hat{\bf q}$, where parity asymmetry is smallest, are nearly perpendicular to the KD direction." + If we assume the parity asvuunetry in the CMD has (he cosmological origin. it is verv hard to explain (hese coincidences.," If we assume the parity asymmetry in the CMB has the cosmological origin, it is very hard to explain these coincidences." + So. (le coincidence ol the preferred direction q with the WMADPT KD direction implies that the CAIB parity asvinmetry may relate to the possible contamination of residual WMAP KD component.," So, the coincidence of the preferred direction $\hat{\bf q}$ with the WMAP7 KD direction implies that the CMB parity asymmetry may relate to the possible contamination of residual WMAP KD component." + Although we will not detailedly study the physical mechanism in (his paper. we could provide some possible explanations for (his coincidence problem.," Although we will not detailedly study the physical mechanism in this paper, we could provide some possible explanations for this coincidence problem." + It is noticed that there is not a great. deal of residual dipole in the WMAP data. so a possible explanation could be connected to the use of the dipole as a photemetric calibrator in (he WALAP data set.," It is noticed that there is not a great deal of residual dipole in the WMAP data, so a possible explanation could be connected to the use of the dipole as a photemetric calibrator in the WMAP data set." + Another possible explanation is related to the contaminations generated by the collective enussion of Ixuiper Bell Objects (INDOs) and other minor bodies in the solar svstem where the KD direction is localized., Another possible explanation is related to the contaminations generated by the collective emission of Kuiper Belt Objects (KBOs) and other minor bodies in the solar system where the KD direction is localized. + Since the emission of INDOs is nearly independent of the frequency in the WMAP frequeney range. this contamination is very hard to be removed in the WMAP data analvsis.," Since the emission of KBOs is nearly independent of the frequency in the WMAP frequency range, this contamination is very hard to be removed in the WMAP data analysis." + In (Marisetal.2011:Hansen2011b).. it was discussed that (his foreground residual could leave signilicant parity asvnunelry in the CAIB data.," In \citep{maris,kbo2}, it was discussed that this foreground residual could leave significant parity asymmetry in the CMB data." + Besides. the explanation may also relate to the measure deviation ofthe WMAP kinematic dipole. which could be caused by the measure error in dipole direction. antenna pointing direction. sidelobe pickup contamination. and so on.," Besides, the explanation may also relate to the measure deviation of the WMAP kinematic dipole, which could be caused by the measure error in dipole direction, antenna pointing direction, sidelobe pickup contamination, and so on." + In (Liu&Li2011).. it was found that this IND deviation could generate the artificial CAIB anisotropies in the low multipoles.," In \citep{liu2011}, it was found that this KD deviation could generate the artificial CMB anisotropies in the low multipoles." + II (his is (rue. these artificial components may account for (he CMD parity violation.," If this is true, these artificial components may account for the CMB parity violation." + In order to cross-check this result. we consider another rotationally variant estimator.," In order to cross-check this result, we consider another rotationally variant estimator," +Most simulations we discuss in this paper do not include explicit feedback: however. we do carry out a few runs which iuclided prompt. energetic feedback from Type II SNe.,"Most simulations we discuss in this paper do not include explicit feedback; however, we do carry out a few runs which included prompt, energetic feedback from Type II SNe." + We apply a simple prescription for stellar feedback to τής the effects of type ID superuovac., We apply a simple prescription for stellar feedback to mimic the effects of type II supernovae. + Because individual star-formation events produce star particles of Aaa~105?AL... we spread feedback over an exteuded period paramcterized by where αι and fy are the initial star particle mass and star particle creation time (seealso?).. and 7 is the maxiuuu of either the dynamical time of the eas out of which the star particle formed. or 10 Alyy (to prevent unrealistically short dvuamical times).," Because individual star-formation events produce star particles of $M_{\rm star} \sim 10^{4-5}$, we spread feedback over an extended period parameterized by where $M_0$ and $t_0$ are the initial star particle mass and star particle creation time \citep[see also][]{Cen:1992p1071}, and $\tau$ is the maximum of either the dynamical time of the gas out of which the star particle formed, or 10 Myr (to prevent unrealistically short dynamical times)." + We take a fraction of the rest mass energv egp of the forming stars as the available feedback cucrev., We take a fraction of the rest mass energy $\epsilon_{\rm FB}$ of the forming stars as the available feedback energy. + This paramctcr can be computed assunine an initial mass function aud a minima initial stellar mass for producing a Type II SN., This parameter can be computed assuming an initial mass function and a minimum initial stellar mass for producing a Type II SN. + As described im ?.. we simply add tlic| thermal energv to the local exid cell.," As described in \citet{Tasker:2006p1072}, we simply add the thermal energy to the local grid cell." + As the surrouudiig gas heats up. it increases the Jeaus length and theoretically. queuches star formation for a time.," As the surrounding gas heats up, it increases the Jeans length and theoretically quenches star formation for a time." +" As was noted in Section L.. if the resolution of the simulation is insufficient to clearly differentiate between a cold. neutral interstellar media and a hot. ionized interstellar 1iediua. then the pumping of this energy iuto the cold star-forming eas results in au ""realistic war componcut (ie. the mixing of hot aud cold phases)."," As was noted in Section \ref{intro}, if the resolution of the simulation is insufficient to clearly differentiate between a cold, neutral interstellar medium and a hot, ionized interstellar medium, then the pumping of this energy into the cold star-forming gas results in an unrealistic warm component (i.e. the mixing of hot and cold phases)." + This wart component sits near the peak of the cooling curve of the gas. aud thus effectively raciates away its enerev very quickly.," This warm component sits near the peak of the cooling curve of the gas, and thus effectively radiates away its energy very quickly." + For the ruu which does include feedback. we take a value of egy=3<109. which corresponds to one 1077 cre SN for every 180 oof stars formed.," For the run which does include feedback, we take a value of $\epsilon_{\rm FB} = 3 \times 10^{-6}$, which corresponds to one $10^{51}$ erg SN for every 180 of stars formed." +" The value of this parameter Is ""uncertain. as it depends on both the initial mass ποο, as well as assuniptious about how the euergv is radiated away imnmediatelv."," The value of this parameter is uncertain, as it depends on both the initial mass function, as well as assumptions about how the energy is radiated away immediately." + Other values are used iu he literature. for example ? cluploved a feedback oxeseription which corresponds to. in our definition. a value of egg5.6«10.," Other values are used in the literature, for example \citet{Abadi:2003p641} employed a feedback prescription which corresponds to, in our definition, a value of $\epsilon_{\rm FB} = 5.6 \times 10^{-6}$." + They. like others. fouud hat this energv was quickly radiated away. aud so the value had little iupact.," They, like others, found that this energy was quickly radiated away, and so the value had little impact." + As we discuss in more detail low. ? added a cooling suppression model. aud argued for a value of egp=13«10.'. although ? adopted egi=2.6«10 © usine a very similar inodel.," As we discuss in more detail below, \citet{Stinson:2006p1023} added a cooling suppression model, and argued for a value of $\epsilon_{\rm FB} = 4.3 \times 10^{-7}$, although \citet{Governato:2007p1022} adopted $\epsilon_{\rm FB} = 2.6 \times 10^{-6}$ , using a very similar model." + Τις from SPIT to AMIR simulations. we note that ? uscd egp5.6«105 for their standard runs (also with a cooling suppression model).," Turning from SPH to AMR simulations, we note that \citet{Agertz:2010p461} used $\epsilon_{\rm FB} = 5.6 \times 10^{-6}$ for their standard runs (also with a cooling suppression model)." + Neither ? nor 7 appear to have used thermal feedback in their snulatious., Neither \citet{Teyssier:2010ia} nor \citet{Gnedin:2011ds} appear to have used thermal feedback in their simulations. + Finally. 2. used egy=10«10.P (lueer than the usual efficiency. because of a postulated contribution from prompt Type Ia SN): this worl: did manage to drive winds without a cooling suppression model. although they added the enerey to the nearest 27 cells weighted inversely by deusitv (ancl also had somewhat worse lass and spatial resolution than used here).," Finally, \citet{Cen:2011bp} used $\epsilon_{\rm FB} = 10 \times 10^{-6}$ (larger than the usual efficiency because of a postulated contribution from prompt Type Ia SN); this work did manage to drive winds without a cooling suppression model, although they added the energy to the nearest 27 cells weighted inversely by density (and also had somewhat worse mass and spatial resolution than used here)." + Ta summary. we see that our chosen value is within the range used by other researchers.," In summary, we see that our chosen value is within the range used by other researchers." + One wav to prevent this cuerev from being quickly lost is to turn off radiative cooling in the region immediately surrounding newly formed stars., One way to prevent this energy from being quickly lost is to turn off radiative cooling in the region immediately surrounding newly formed stars. + This was first attempted. by ?. but has since been explored by a rauge of siuulations (eo...77322)..," This was first attempted by \citet{Gerritsen:1997p1039}, but has since been explored by a range of simulations \citep[e.g.,][]{Thacker:2000p1040, SommerLarsen:2003p1116, Stinson:2006p1023, Governato:2007p1022, Agertz:2010p461, Colin:2010p1053, Piontek:2011p1041, Guedes:2011p1080}." + The idea is to use the Sedow-Tavlor solution for a blastwave to model the suberid shock plysics that the code canunuot resolve., The idea is to use the Sedov-Taylor solution for a blastwave to model the subgrid shock physics that the code cannot resolve. + It, It +"This work confirms that screening in the hot, dense plasma of stellar cores depends on the relative velocities of the interacting ions.","This work confirms that screening in the hot, dense plasma of stellar cores depends on the relative velocities of the interacting ions." + The Debye-Hiicckel screening energy is only valid for describing average properties of the plasma., The Debye-Hücckel screening energy is only valid for describing average properties of the plasma. +" Since faster ions are more likely to engage in nuclear reactions than thermal ions, the mean-field treatment does not provide an accurate representation of this velocity-skewed phenomenon."," Since faster ions are more likely to engage in nuclear reactions than thermal ions, the mean-field treatment does not provide an accurate representation of this velocity-skewed phenomenon." +" In fact, the fast pairs of ions tend to lose energy to the plasma instead of gaining energy from it which would reduce nuclear reaction rates instead of enhancing them."," In fact, the fast pairs of ions tend to lose energy to the plasma instead of gaining energy from it which would reduce nuclear reaction rates instead of enhancing them." + Solar and stellar models should be adjusted to account for this dynamic screening effect., Solar and stellar models should be adjusted to account for this dynamic screening effect. +" Currently, there is no formalism to compute dynamic screening analytically."," Currently, there is no formalism to compute dynamic screening analytically." + This paper is intended to provide insight into the difference between the Salpeter formalismlj and the numerically determined dynamic screening., This paper is intended to provide insight into the difference between the Salpeter formalism and the numerically determined dynamic screening. + A detailed calculation of the dynamic screening correction to the p-p reaction rate in the solar core based on this numerical work is underway for a future publication., A detailed calculation of the dynamic screening correction to the p-p reaction rate in the solar core based on this numerical work is underway for a future publication. +" However, these numerical calculations will not minimize the need for an analytical formalism for dynamic screening in order to generalize the results to other temperatures, densities, compositions, and reactions."," However, these numerical calculations will not minimize the need for an analytical formalism for dynamic screening in order to generalize the results to other temperatures, densities, compositions, and reactions." + Only then can dynamic screening be encorporated consistently in solar and stellar models., Only then can dynamic screening be encorporated consistently in solar and stellar models. + We thank Dan Mao for helpful discussions about the simulations., We thank Dan Mao for helpful discussions about the simulations. + This work was supported in part by grant AST-0708568 of the National Science Foundation., This work was supported in part by grant AST-0708568 of the National Science Foundation. +would like. in the remaining of this section. to see whether it is possible to obtain some. albeit crude. estimates of the optimal softening as a function of quantities linked to the mass distribution. but more straightforward to calculate than the ALASL.,"would like, in the remaining of this section, to see whether it is possible to obtain some, albeit crude, estimates of the optimal softening as a function of quantities linked to the mass distribution, but more straightforward to calculate than the $MASE$." + As we have already seen. smaller values of the softening are necessary for more compact configurations or for representations with a larger number of particles. ic. in cases where the particles are closer together.," As we have already seen, smaller values of the softening are necessary for more compact configurations or for representations with a larger number of particles, i.e. in cases where the particles are closer together." + This suggests that some measure of inter-particle distances could. be used. for estimating the optimal value of the softening., This suggests that some measure of inter-particle distances could be used for estimating the optimal value of the softening. + Furthermore. since MSE is more dependent on the accuracy of the forces to nearby particles. we will try using the distances of the few nearest neighbours.," Furthermore, since $MASE$ is more dependent on the accuracy of the forces to nearby particles, we will try using the distances of the few nearest neighbours." + Let us thus. for a given configuration. measure. for every particle. the distance to its twelve nearest neighbours.," Let us thus, for a given configuration, measure, for every particle, the distance to its twelve nearest neighbours." + Now we need to average this over all particles in the configuration. in order to obtain. for the whole configuration. the mean distance to the nearest. neighbour. he mean distance to the second. nearest. neighbour τοι.," Now we need to average this over all particles in the configuration, in order to obtain, for the whole configuration, the mean distance to the nearest neighbour, the mean distance to the second nearest neighbour etc.," + up to the mean distance to the 12th nearest. neighbour., up to the mean distance to the 12th nearest neighbour. + A standard: arithmetic average would not. be appropriate or this., A standard arithmetic average would not be appropriate for this. + This can be understood if we mentally add to the configuration a single particle. located so far from it that it can. for arguments sake. be considered at infinity.," This can be understood if we mentally add to the configuration a single particle, located so far from it that it can, for argument's sake, be considered at infinity." + This new xwiicle will not inlluence the value of AZ.ASE. the value of coi; that. should. be used. or the accuracy in the force calculations.," This new particle will not influence the value of $MASE$, the value of $\epsilon_{opt}$ that should be used, or the accuracy in the force calculations." + On the other hand it will influence the mean distance to the Ath neighbour., On the other hand it will influence the mean distance to the $k$ th neighbour. +" I. is thus not reasonable to expect a close relation between eoi, and the arithmetic mean of the Ath closest neighbour of all particles.", It is thus not reasonable to expect a close relation between $\epsilon_{opt}$ and the arithmetic mean of the $k$ th closest neighbour of all particles. + For this reason. instead of the standard arithmetic mean. we will prefer using the harmonic mean or. since the force is inversely proportional to the square ol the distance. where the summations are over all the particles in the configuration. and f&=(1....12).," For this reason, instead of the standard arithmetic mean, we will prefer using the harmonic mean or, since the force is inversely proportional to the square of the distance, where the summations are over all the particles in the configuration and $k = (1,..., 12)$." + In this wav we obtain for à given configuration the mean distance to the nearest neighbour. the mean distance to the second. nearest neighbour etc..," In this way we obtain for a given configuration the mean distance to the nearest neighbour, the mean distance to the second nearest neighbour etc.," + up to the mean distance to the 12th nearest neighbour., up to the mean distance to the 12th nearest neighbour. + It is possible to diminish the noise in such calculations by considering several realisations of the same configuration and then mean ruso) OL eoncan2- DOW Using simple arithmetic means over all realisations.," It is possible to diminish the noise in such calculations by considering several realisations of the same configuration and then mean $r_{k,mean1}$ or $r_{k,mean2}$, now using simple arithmetic means over all realisations." +" We calculated these mean inter-particle distances as a function of IN. for the three density. distributions considered. above. using 3 10""£N. realisations."," We calculated these mean inter-particle distances as a function of $N$ for the three density distributions considered above, using 3 $\times 10^6 / +N$ realisations." + In all cases the dependence is roughly linear on a log-log plane. ancl thus can be represented: by power laws of the type πο IS equal ο μυ ," In all cases the dependence is roughly linear on a log-log plane, and thus can be represented by power laws of the type where $r_{k,mean}$ is equal to $r_{k,mean1}$ or $r_{k,mean2}$." +both cases the value of e is around -0.33 or -0.34 ancl does. not depend much on which of the twelve nearest neighbours is considered., In both cases the value of $a$ is around -0.33 or -0.34 and does not depend much on which of the twelve nearest neighbours is considered. + The latter is not truce if we use standard means., The latter is not true if we use standard means. + The values of A depend. of course on which of the nearest neighbours is chosen., The values of A depend of course on which of the nearest neighbours is chosen. + We have thus obtained so far. for a given density istribution and a given number of particles IN. the average istance to the Ath nearest neighbour. for & between 1 and 12. as well as an optimal softening ei; (cf.," We have thus obtained so far, for a given density distribution and a given number of particles $N$, the average distance to the $k$ th nearest neighbour, for $k$ between 1 and 12, as well as an optimal softening $\epsilon_{opt}$ (cf." + Figure 9))., Figure \ref{3eopt}) ). + From jese two we can eliminate the dependence on IN. and obtain 10 dependence of e; directly on the average distance to 16 Ath nearest neighbour., From these two we can eliminate the dependence on $N$ and obtain the dependence of $\epsilon_{opt}$ directly on the average distance to the $k$ th nearest neighbour. + This is given in Figure 11. for 1 sixth. neighbour. after both distances and the softening ve been weighted appropriately. as discussed. in. section 2..," This is given in Figure \ref{eopt_dist} for the sixth neighbour, after both distances and the softening have been weighted appropriately, as discussed in section \ref{sec:notation}." + We have repeated this exercise for all other values of & tween 1 and 12. but since the results are similar we co not reproduce them here.," We have repeated this exercise for all other values of $k$ between 1 and 12, but since the results are similar we do not reproduce them here." + Again a power law gives a good representation of the dependences The values of A and a. for various neighbours are given in Table 2..," Again a power law gives a good representation of the dependences The values of A and a, for various neighbours are given in Table \ref{tab:neib_hpd}." + Figure ll can give some. albeit rough. estima| of the optimal softening to. be used. once àn Meaney has been caleulated. since the dillerences. between the three models are less. or of the order of. 0.5 dex.," Figure \ref{eopt_dist} can give some, albeit rough, estimate of the optimal softening to be used, once an $r_{k,mean}$ has been calculated, since the differences between the three models are less, or of the order of, 0.5 dex." + Lt is. however. possible to narrow down the prediction. further.," It is, however, possible to narrow down the prediction further." + Iideed the three dependences are. ordered. as a function. of the central concentration of the corresponding naioclels. less concentrated models corresponding to higher softening than more concentrated Ones. in σους agreement with what was previously cliseussecl in this section.," Indeed the three dependences are ordered as a function of the central concentration of the corresponding models, less concentrated models corresponding to higher softening than more concentrated ones, in good agreement with what was previously discussed in this section." + Thus comparing the central concentration of the new model. whose optimal softening we want to estimate. with that of the three considered here. should probably narrow the estimate to 0.2 or 0.3 dex.," Thus comparing the central concentration of the new model, whose optimal softening we want to estimate, with that of the three considered here, should probably narrow the estimate to 0.2 or 0.3 dex." + So [ar we have considered. only spherical objects., So far we have considered only spherical objects. + By using. however. the ALASZ rather than the radial. MS47 accuracy estimator it is possible to consider non-spherical configurations.," By using, however, the $MASE$ rather than the radial $MISE$ accuracy estimator it is possible to consider non-spherical configurations." + As an example we will in this section consider, As an example we will in this section consider +SNe Ta at the current epoch (10 Gyr) is 0.001 vetsa factor of two below that of the Model 1 spiral galaxy.,SNe Ia at the current epoch $10 $ Gyr) is $0.001$ $^{-1}$; a factor of two below that of the Model 1 spiral galaxy. + At first. the DDS rate increases and then remains fairl constant until star formation ceases at 10 Cr.," At first, the DDS rate increases and then remains fairly constant until star formation ceases at 10 Gyr." + At all epochs (particularly diving star formation). the DDS channel rates siguificauth outumuber the SDS and AM οΝα channels. but to a lesser degree than when compared with Model 1.," At all epochs (particularly during star formation), the DDS channel rates significantly outnumber the SDS and AM CVn channels, but to a lesser degree than when compared with Model 1." + The SDS channel also exhibits a relatively constant rate at times later than 1 Cyr.," The SDS channel also exhibits a relatively constant rate at times later than $\sim +1$ Gyr." + At the current epoch the SDS rates ave at the level of 0.0002 vrt: a factor of ~5 below the DDS rates., At the current epoch the SDS rates are at the level of $ 0.0002$ $^{-1}$; a factor of $\sim 5$ below the DDS rates. + The rates arising frou AM. CVu proseutors are fairly uceligible at all epochs (<10{sy 4)., The rates arising from AM CVn progenitors are fairly negligible at all epochs $< 10^{-4}$ $^{-1}$ ). + Even when the rates from all three progenitor channels in the considered galaxy 1odel are combined (00.0012 SNe Ia 1). the SN Ia rate at the current epoch falls below the empiricallv-derived rate from Cappellaroetal.(1999). by roughly a factor of two.," Even when the rates from all three progenitor channels in the considered galaxy model are combined (0.0012 SNe Ia $^{-1}$ ), the SN Ia rate at the current epoch falls below the empirically-derived rate from \citet{CET99} by roughly a factor of two." + It is worth comparing our model ealaxy rates to the rates presented in Maunuuccietal.(2005)., It is worth comparing our model galaxy rates to the rates presented in \citet{Man05}. +. Iu that study. Alaunuceietal.(2005) derive SN rates for galaxies of various morphological types. and present the SN rates In SNuM (as inferred from A-band dhuuunositv measurements).," In that study, \citet{Man05} derive SN rates for galaxies of various morphological types, and present the SN rates in SNuM (as inferred from $K$ -band luminosity measurements)." + We cannot directly compareo our rates in Tables 1. 2 to those of Maunuccietal.(2005). since we do not know the exact ages of the galaxies im their saluple., We cannot directly compare our rates in Tables 1 2 to those of \citet{Man05} since we do not know the exact ages of the galaxies in their sample. + However. we note that our Model 1 elliptical ealaxy SN Ia rate at 10 is 0.005 SNUAL (see Table 1 DDS rates). which is nearly a factor of LO below the SN Ia rate in E/S0 ealaxies presented by Maunucciet (2005): 0.011002. the SNuAL(sce. their Table 2).," However, we note that our Model 1 elliptical galaxy SN Ia rate at 10 is $0.005$ SNuM (see Table 1 DDS rates), which is nearly a factor of 10 below the SN Ia rate in E/S0 galaxies presented by \citet{Man05}: $0.044^{+0.016}_{-0.014}$ SNuM (see their Table 2)." + We uote as well that for same niodel galaxy. at f=500 Αν (shortly after à burst of star formation at f= 0) we obtain a SN Ta rate of ~0.18 SNuM (inostlv via the DDS channel with some contribution from AM. CVn). which is about a factor of 2 lower than the rauge of SN Ta rates found for star forming bregular galaxies in Mauunucciet (0.77 22). The Mz," We note as well that for the same model galaxy, at $t = 500$ Myr (shortly after a burst of star formation at $t=0$ ) we obtain a SN Ia rate of $\sim 0.18$ SNuM (mostly via the DDS channel with some contribution from AM CVn), which is about a factor of $2$ lower than the range of SN Ia rates found for star forming Irregular galaxies in \citet{Man05} $0.77^{+0.42}_{-0.31}$ )." +uinecietal.(2005) SN Ta rate for SOa/Lb spirals is found tobe 0.065ratesi5 DtOSNuM. which matches our Model 1 spiral galaxy E 5 Car (0.065 SNuM: mostly DDS AAT CVn progenitors).," The \citet{Man05} SN Ia rate for S0a/b spirals is found to be $0.065^{+0.027}_{-0.025}$ SNuM, which matches our Model 1 spiral galaxy rates at 5 Gyr $0.065$ SNuM; mostly DDS AM CVn progenitors)." + The SN In rate of Shefd spirals frou Manunuccietal.(2005) is 0.17Uca: and matches our Model 1 spiral SN Ia rate ouly at verv carly times (~0.51 Cyr. see Table 1).," The SN Ia rate of Sbc/d spirals from \citet{Man05} is $0.17 ^{+0.068}_{-0.063}$, and matches our Model 1 spiral SN Ia rate only at very early times $\sim 0.5 - 1$ Gyr, see Table 1)." + For the Model 2 elliptical galaxy. we find a low SN Ta rate of ~0.001 SNuM at LO Cor. which is over an order of magnitude below the rate for E/SO galaxies in Alannuccietal.(2005) (0Η16eo SNuM).," For the Model 2 elliptical galaxy, we find a low SN Ia rate of $\sim 0.001$ SNuM at 10 Gyr, which is over an order of magnitude below the rate for E/S0 galaxies in \citet{Man05} $0.044^{+0.016}_{-0.014}$ SNuM)." + The rate at 500 Myr for the same galaxw is ~0.14 SNuM (sce Table 2): mostly arising frou DDS progenitors with some contribution from AM. CVn and SDS., The rate at 500 Myr for the same galaxy is $\sim 0.14$ SNuM (see Table 2); mostly arising from DDS progenitors with some contribution from AM CVn and SDS. + This rate is a factor of a few below the In rates for hregular galaxies in Mannuccietal.(2005) (0.77 0.1)., This rate is a factor of a few below the Ia rates for Irregular galaxies in \citet{Man05} $0.77^{+0.42}_{-0.31}$ ). + Iu comparing our Model 2 spiral rates. we find that our SN Ia rate at3 Car (~0.07 SNuM: mostly DDS progenitors with ποιο --- ποιι SDS) is within the ranee of rates iu Maumnuccietal.(2005). for σαο spirals presented(0.065i culpoe). while oulv our spiral rates for Model 2 at very¢ times (<1 Cyr: 0.13 SNuM. DDS) are high enough to match those of Mauuucciotal.(2005) for Shefd type galaxies (0.17 os).," In comparing our Model 2 spiral rates, we find that our SN Ia rate at 3 Gyr $\sim 0.07$ SNuM; mostly DDS progenitors with some contribution from SDS) is within the range of rates presented in \citet{Man05} for S0a/b spirals $0.065^{+0.027}_{-0.025}$ ), while only our spiral rates for Model 2 at very early times $< 1$ Gyr; $0.13$ SNuM, DDS) are high enough to match those of \citet{Man05} for Sbc/d type galaxies $0.17^{+0.068}_{-0.063}$ )." + We note that in eeneral our rates (por unit mass) are lower than those of Mounuccietal.(2005).. indicating that perhaps other channels leading to the formation of SNe Ia should be considered in evolutionary studies (0.8.. single stars. sub-Chaudrasekliar mass SNe Ia}.," We note that in general, our rates (per unit mass) are lower than those of \citet{Man05}, indicating that perhaps other channels leading to the formation of SNe Ia should be considered in evolutionary studies (e.g., single stars, sub-Chandrasekhar mass SNe Ia)." + We also note however that our predicted rates as a function of time (at least for the DDS) are consistent with the observed rates presented w Cappellaroetal.(1999)., We also note however that our predicted rates as a function of time (at least for the DDS) are consistent with the observed rates presented by \citet{CET99}. +. We fud that in general the DDS outuuuber SDS and AAT CV progenitors., We find that in general the DDS outnumber SDS and AM CVn progenitors. + This effect is somewhat more xonounced in Model L (a factor of &10) than in Model 2 (a factor of ~5). but the reason why is clear for )oth models.," This effect is somewhat more pronounced in Model 1 (a factor of $\gtrsim +10$ ) than in Model 2 (a factor of $\sim 5$ ), but the reason why is clear for both models." + The occurrence rate of a CO-CO WD nerecr with a total mass =1.1 M. (DDS) is higher han that of buikliug up a CO WD's Dues to LIAL. via stable mass transfer (SDS) in a binary., The occurrence rate of a CO-CO WD merger with a total mass $\geq 1.4$ $_{\odot}$ (DDS) is higher than that of building up a CO WD's mass to $\sim 1.4 \msun$ via stable mass transfer (SDS) in a binary. + Formation efficiencies in both of the above cases are very low. after all SNe Ta are rather rare eveuts," Formation efficiencies in both of the above cases are very low, after all SNe Ia are rather rare events." + However. relatively speaking it is easier to find a pair of two CO WDs (DDS). each with a mass of =0.7AL. which is a typical mass for CO WDs. than it is for a CO WD to double its mass through accretion (SDS/AM. CVu).," However, relatively speaking it is easier to find a pair of two CO WDs (DDS), each with a mass of $\gtrsim 0.7 \msun$ which is a typical mass for CO WDs, than it is for a CO WD to double its mass through accretion (SDS/AM CVn)." + This fincdine Is a consequence of the recent updates ou accumulation plivsics calculations (Hachlisuetal.1999:kato&IIachisu1999:Nomotoetal.2007) that we have adopted iu our evolutionary study.," This finding is a consequence of the recent updates on accumulation physics calculations \citep{HKN99,KH99,Nom07} that we have adopted in our evolutionary study." + Basically. the acctuuulation outo a WD is hampered by a umber of processes that tend to remove matter which is transferred from the colupanion ina close binary (¢.¢.. nova explosions. Ifc-shell flashes. optically thick winds from the surface of au accreting WD). in some cases leaciug to the disruption of an accreting WD before it reaches the limiting Chandrasekhar mass: bes accretion of ~0.1AL. of a IIeaich laver which ignites aud disrupts the uuderlviug WD (sub-ChandrasclharmassSNTa.seeegNatoIachisu 1999).," Basically, the accumulation onto a WD is hampered by a number of processes that tend to remove matter which is transferred from the companion in a close binary (e.g., nova explosions, He-shell flashes, optically thick winds from the surface of an accreting WD), in some cases leading to the disruption of an accreting WD before it reaches the limiting Chandrasekhar mass; i.e., accretion of $\sim 0.1 \msun$ of a He-rich layer which ignites and disrupts the underlying WD \citep[sub-Chandrasekhar mass SN Ia, see e.g.,][]{KH99}." +. Since. there is a rather narrow range of nass transfer rates which iav lead to efficicut acctuntlation and there are not that mauy binary configurations (aud we have cousidercd the entire rauge for our adopted evolutionary model) that can sustain mass transter for a prolonged period of time. the SDS and AM. C'Vu channels ave found to produce SNe In at very low rates.," Since, there is a rather narrow range of mass transfer rates which may lead to efficient accumulation and there are not that many binary configurations (and we have considered the entire range for our adopted evolutionary model) that can sustain mass transfer for a prolonged period of time, the SDS and AM CVn channels are found to produce SNe Ia at very low rates." + The receut theoretical study. of Hachisuetal.(2008) &uds à SDS delay tine distribution which follows a power aw., The recent theoretical study of \citet{HKN08} finds a SDS delay time distribution which follows a power law. + Iu their studs. DHachisuetal.(2008) incorporate a now inass stripping effect (based on EHachisuetal. (1999))). where iu the case of high mass trausfer rates he WD blows an optically thick wind strong enough to ‘strip maternal from a iain sequence or giaut donor.," In their study, \citet{HKN08} incorporate a new mass stripping effect (based on \citet{HKN99}) ), where in the case of high mass transfer rates the WD blows an optically thick wind strong enough to `strip' material from a main sequence or giant donor." + This effect in return stabilizes mass trauster. enabliug he binary to avoid a CE phase even in the case of a relatively massive (~6 M...) donor.," This effect in return stabilizes mass transfer, enabling the binary to avoid a CE phase even in the case of a relatively massive $\sim 6$ $_{\odot}$ ) donor." + The result is that the WD can accrete stably up to the Chandrasekhar mass. with a wider rauge of potcutial progenitor donor ZAMS masses: 0.96 ADL. in Bachisuetal.(2008) vs. O7.2.7 AÍ. iu our current study.," The result is that the WD can accrete stably up to the Chandrasekhar mass, with a wider range of potential progenitor donor ZAMS masses: $0.9 - 6$ $_{\odot}$ in \citet{HKN08} vs. $0.7 - 2.7$ $_{\odot}$ in our current study." + Even though we allow for SN Ta progenitors to form from auv initial mass spanning the iuifial mass function. our SDS SNe Ta ouly derive from binaries involving low-mass donors since we do nof take iuto account this stripping effect.," Even though we allow for SN Ia progenitors to form from any initial mass spanning the initial mass function, our SDS SNe Ia only derive from binaries involving low-mass donors since we do not take into account this stripping effect." + We note that the Iachisuctal.(2008) imiodoel predicts the preseuce of a, We note that the \citet{HKN08} model predicts the presence of a +mechanism effect compared to previous estimates that ignore the plunge region (e.g.Ghosh&Abramowicz1997).,mechanism effect compared to previous estimates that ignore the plunge region \citep*[e.g.][]{ga97}. +. Radio galaxies are usually classified as FR I or FR II sources depending on their radio. morphology., Radio galaxies are usually classified as FR I or FR II sources depending on their radio morphology. + 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: Fanaroff Riley 1974)., 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; Fanaroff Riley 1974). + What causes the morphological difference between FR I and FR II radio galaxies is still. unclear., What causes the morphological difference between FR I and FR II radio galaxies is still unclear. + The theoretical models fall into two different groups: (1) the morphological differences arise of the different physical conditions in their ambient medium (seeGopal-Krishna&Wi-ita2000.fora summary): (2) their intrinsic difference of their central engines. 1.e.. different accretion modes and/or Jet formation processes (e.g..Bicknell1995:Reynoldsetal.Woldetal.2007:Hardeastle 2007).," The theoretical models fall into two different groups: (1) the morphological differences arise of the different physical conditions in their ambient medium \citep*[see][for a +summary]{gk00}; (2) their intrinsic difference of their central engines, i.e., different accretion modes and/or jet formation processes \citep*[e.g.,][]{bi95,re96a,me99,gc01,ma04,wo07,ha07}." + Most previous works on the jet power extracted from ADAFs were based on self-similar solutions of ADAFs (e.g..Armitagemenetal. 2007).," Most previous works on the jet power extracted from ADAFs were based on self-similar solutions of ADAFs \citep*[e.g.,][]{an99,me01,cr04,ne07}." +. The self-similar solution can. reproduce the global solution quite well at large radii. while it deviates significantly near the black hole (e.g..Narayanetal.1997). where the relativistic jets are supposed to be formed.," The self-similar solution can reproduce the global solution quite well at large radii, while it deviates significantly near the black hole \citep*[e.g.,][]{na97}, where the relativistic jets are supposed to be formed." + In this paper we calculate the jet power. incorporating some recent MHD simulation results. based on the global ADAF solutions surrounding. Kerr black holes.," In this paper we calculate the jet power, incorporating some recent MHD simulation results, based on the global ADAF solutions surrounding Kerr black holes." + We then compare our model calculations with the accretion power and jet power of XRBs and the radio galaxy FR UFR II dichotomy., We then compare our model calculations with the accretion power and jet power of XRBs and the radio galaxy FR I/FR II dichotomy. + We calculate the global structure of an accretion flow using an approach similar to that of Narayan et al. (, We calculate the global structure of an accretion flow using an approach similar to that of Narayan et al. ( +1997).,1997). + However. the pseudo-Kerr potential for a rotating BH given by Mukhopadhyay(2002) is adopted in solving the equations of the accretion flow. which allows us to calculate the structure of an accretion flow surrounding either a spinning or a nonspinning black hole.," However, the pseudo-Kerr potential for a rotating BH given by \citet{mu02} is adopted in solving the equations of the accretion flow, which allows us to calculate the structure of an accretion flow surrounding either a spinning or a nonspinning black hole." +" The simple a-viscosity (r4,=ap. and p is total pressure. Le. gas pressure plus magnetic pressure) is adopted and all radiative processes (synchrotron.le] bremsstrahlung and Compton scattering) are included consistently in our calculations for ADAP structure."," The simple $\alpha$ -viscosity $\tau_{r\phi}=\alpha p$, and $p$ is total pressure, i.e., gas pressure plus magnetic pressure) is adopted and all radiative processes (synchrotron, bremsstrahlung and Compton scattering) are included consistently in our calculations for ADAF structure." + The advection by tons and electrons has been considered in the energy equation. and a more realistic state of accreting gas (instead of a polytropic index 2.) is employed in the calculations. which ts similar to that used by Manmoto (2000).," The advection by ions and electrons has been considered in the energy equation, and a more realistic state of accreting gas (instead of a polytropic index $\gamma_{\rm g}$ ) is employed in the calculations, which is similar to that used by Manmoto (2000)." +" We solve a set of hydrodynamical equations (1.e.. the radial momentum. angular momentum. and energy equations) for an ADAF. and tune the parameter /,,. the specific angular momentum of the gas swallowed by the black hole. to let the solution passing smoothly through the sonic point near the black hole (seeNarayanetal.1997.fordetails)..."," We solve a set of hydrodynamical equations (i.e., the radial momentum, angular momentum, and energy equations) for an ADAF, and tune the parameter $l_{\rm in}$, the specific angular momentum of the gas swallowed by the black hole, to let the solution passing smoothly through the sonic point near the black hole \citep*[see][for details]{na97}." + We find that the derived global solutions for Kerr black holes can reproduce all the essential properties of the solutions derived in full general relativistic frame by Manmoto(2000) with error less than10%., We find that the derived global solutions for Kerr black holes can reproduce all the essential properties of the solutions derived in full general relativistic frame by \citet{man00} with error less than. +. The global structure of an ADAF surrounding a BH spinning atrate / with mass My can be calculated with proper outer boundaries (e.g.. Manmoto 2000). if the parameters 77. o. ον and à are specified.," The global structure of an ADAF surrounding a BH spinning atrate $j$ with mass $M_{\rm BH}$ can be calculated with proper outer boundaries (e.g., Manmoto 2000), if the parameters $\dot{m}$, $\alpha$ , $\beta$, and $\delta$ are specified." + The parameter j=JGMa| Is the dimensionless angular momentum. where J is the angular momentum of the BH. 7)2M/Mig is dimensionless accretion rate. and Mj is the Eddington accretion rate defined as Mya143«I0Mgy/M..gs.," The parameter $j=J/GM_{\rm BH}^{2}c^{-1}$ is the dimensionless angular momentum, where $J$ is the angular momentum of the BH, $\dot{m}=\dot{M}/\dot{M}_{\rm + Edd}$ is dimensionless accretion rate, and $\dot{M}_{\rm + Edd}$ is the Eddington accretion rate defined as $\dot{M}_{\rm Edd}=1.4\times10^{18}M_{\rm BH}/\msun \rm\ g + \ s^{-1}$ ." + The value of « adopted in ADAF modeling is supposed to be within a very narrow range. Le.. a=0.1—0.3 (e.g..Narayan&McClintock2008.andreferencestherein)... which is supported by MHD numerical simulations of accretion flows ~0.05—0.2 (Hawley&Balbus2002) and the observationally-determmed values ~0.1—0.4 based primarily on studiesof outbursts in dwarf novae and X-ray transients(King2007)..," The value of $\alpha$ adopted in ADAF modeling is supposed to be within a very narrow range, i.e., $\alpha=0.1-0.3$ \citep*[e.g.,][and references + therein]{nm08}, which is supported by MHD numerical simulations of accretion flows $\sim0.05-0.2$ \citep{hb02} and the observationally-determined values $\sim0.1-0.4$ based primarily on studiesof outbursts in dwarf novae and X-ray \citep{ki07}. ." +" The magnetic parameter ./ (defined as ratio of gas to magnetic pressure in the accretion flow, P./Py) is not an independent parameter and can be related to à as 7(0.55—n)f/o. as suggested by MHD simulations (e.g..Hawleyetal.1995).. where P,=B5sane8x and Baynamo 18 the magnetic field strength of the ADAF in the local reference frame."," The magnetic parameter $\beta$ (defined as ratio of gas to magnetic pressure in the accretion flow, $\beta=P_{\rm g}/P_{\rm m}$ ) is not an independent parameter and can be related to $\alpha$ as $\beta\simeq(0.55-\alpha)/\alpha$, as suggested by MHD simulations \citep*[e.g.,][]{ha95}, where $P_{\rm m}=B_{\rm dynamo}^{2}/8\pi$ and $B_{\rm dynamo}$ is the magnetic field strength of the ADAF in the local reference frame." + The parameter ο)~1—5for the typical value of a~0.120.3., The parameter $\beta\simeq1-5$for the typical value of $\alpha\simeq0.1-0.3$ . + Another poorly constrained parameter is 9. which describes the fraction of the turbulent dissipation that directly heats the electrons in the flow.," Another poorly constrained parameter is $\delta$, which describes the fraction of the turbulent dissipation that directly heats the electrons in the flow." + Recent ADAF models typically assume à20.3—0.5 (e.g..Yuanetal.2003:Wu2007.seealsoSharmaetal.2007forslightlylowerà value)..," Recent ADAF models typically assume $\delta\simeq0.3-0.5$ \citep*[e.g.,][see also Sharma et al. 2007 for +slightly lower $\delta$ value]{yu03,wu07}." + Two important modifications of the above global ADAF model are also included (seeMeier2001:Nemmenetal.2007.formoredetails).," Two important modifications of the above global ADAF model are also included \citep*[see][for more details]{me01,ne07}." +". First. as viewed from an outside observer at infinity in the Boyer-Lindquist reference frame. the disk angular velocity O is a sum of its angular velocity relative to the local metric © plus the angular velocity of the metric itself in the Boyer-Lindquist frame w=—g,,/2g,,,. Le. Q= "," First, as viewed from an outside observer at infinity in the Boyer-Lindquist reference frame, the disk angular velocity $\Omega$ is a sum of its angular velocity relative to the local metric $\Omega^{'}$ plus the angular velocity of the metric itself in the Boyer-Lindquist frame $\omega\equiv-g_{\phi t}/g_{\phi\phi }$, i.e., $\Omega=\Omega^{'}+\omega$ ." +Second. we also take into account the field-enhancing shear caused by frame dragging in the Kerr metric. as first suggested by Meier(1999).. which seems to be supported by MHD simulations (e.g..Hawley&Krolik2006).," Second, we also take into account the field-enhancing shear caused by frame dragging in the Kerr metric, as first suggested by \citet{me99}, which seems to be supported by MHD simulations \citep*[e.g.,][]{hk06}." +. Following the work of Meier(2001).. the amplified magnetic field related to the magnetic field produced by the dynamo process in the ADAF can be expressed as B=gBaynano. where g2OQ' is the field-enhaneing factor.," Following the work of \citet{me01}, the amplified magnetic field related to the magnetic field produced by the dynamo process in the ADAF can be expressed as $B=gB_{\rm dynamo}$, where $g=\Omega/\Omega^{'}$ is the field-enhancing factor." + For a black hole of mass Mgy and dimensionless angular momentum /. with magnetic fields δι normal to the horizon at Ry2[EL—P) (Re=GMw/c is the gravitational radius). the power extracted7|R« with the BZ mechanism is given by (e.g..Ghosh&Abramowiez1997;MacDonaldThorne1982) TAE where wp=OOH-OE)On is determined by the angular velocity of field lines. ©) relative to that of the hole Oy ," For a black hole of mass $M_{\rm BH}$ and dimensionless angular momentum $j$, with magnetic fields $B_{\perp}$ normal to the horizon at $R_{\rm H}=[1+(1-j^{2})^{1/2}]R_{\rm g}$ $R_{\rm g }=GM_{\rm +bh}/c^{2}$ is the gravitational radius), the power extracted with the BZ mechanism is given by \citep*[e.g.,][]{ga97,mt82} + c, where $\omega_{\rm F}^{2}\equiv\Omega_{\rm F}(\Omega_{\rm +H}-\Omega_{\rm F})/\Omega_{\rm H}^{2}$ is determined by the angular velocity of field lines $\Omega_{\rm F}$ relative to that of the hole $\Omega_{\rm H}$." +In order to estimate the maximal power extracted from a spinning BH. wy is always required to be 1/2 (e.g..Livioet 2002). ," In order to estimate the maximal power extracted from a spinning BH, $\omega_{\rm F}$ is always required to be 1/2 \citep*[e.g.,][]{li99,ca02}. ." +"The magnetic field B, 1s assumedto approximate tothe poloidal component δρ. andLivioetal. proposed that By~B,y;44, due to the hot thick disk of ADAF (H~ R)."," The magnetic field $B_{\perp}$ is assumedto approximate tothe poloidal component $B_{\rm p}$ , and\citet{li99} proposed that $B_{\rm p}\simeq B_{\rm dynamo}$ due to the hot thick disk of ADAF $H\sim R$ )." + Therefore. we use B4~By—eBay) In our calculations considering the field enhancing effect.," Therefore, we use $B_{\perp}\simeq B_{\rm p}\simeq gB_{\rm dynamo}$ in our calculations considering the field enhancing effect." + This ts consistent with recent MHD simulations in which the poloidal fields are dominant in the z-direction near the BH (Katoetal. 2004).., This is consistent with recent MHD simulations in which the poloidal fields are dominant in the $z$ -direction near the BH \citep{ka04}. . + Following the work of Nemmenetal. (2007). all the physical quantities are evaluated at R= Ri.," Following the work of \citet{ne07}, , all the physical quantities are evaluated at $R=R_{\rm ms}$ ." +Practically. oue needs to truucate the tifinite series by a consistent. manuuer.,"Practically, one needs to truncate the infinite series by a consistent manner." + Fortunately. the i0n-Craussianity generated by gravitationally nonlinear evolution is known to approximately follow he hierarchical model of the higher-order correlations. in which the polyspectra P have the order. PO!~ΟΡ1] (eis.Bernardeau..Colombi.Gaztanaga&Scoccimarro 2002). ," Fortunately, the non-Gaussianity generated by gravitationally nonlinear evolution is known to approximately follow the hierarchical model of the higher-order correlations, in which the polyspectra $P^{(N)}$ have the order, $P^{(N)} \sim {\cal O}[P(k)^{N-1}]$ \citep[e.g.,][]{ber02}. ." +This ueans p) Q(eA;E7). where e~vPO)/V.rv," This means $p^{(N)} \sim {\cal O}(\epsilon^{N-2})$, where $\epsilon \sim \sqrt{P(k)/V}$." + Therefore. one can evaluate the phase correlations. in. yerturbative manuer as long as the expausiou parameter e is s1uall., Therefore one can evaluate the phase correlations in perturbative manner as long as the expansion parameter $\epsilon$ is small. + It ds stralghitforward to perform he above procedure to express P({jag].0&1) in terms of normalized polyspectra to arbitrary order in e.," It is straightforward to perform the above procedure to express ${\cal P}(\{|\alpha_{\sbfm{k}}|, +\theta_{\sbfm{k}}\})$ in terms of normalized polyspectra to arbitrary order in $\epsilon$." + Inthe lowest order approximation. ouly the normalized bispectrum p) eives the term of order O(el).," In the lowest order approximation, only the normalized bispectrum $p^{(3)}$ gives the term of order ${\cal +O}(\epsilon^1)$." + The result is Higher-order terms can be similarly calculated. although they are somehow tedious.," The result is Higher-order terms can be similarly calculated, although they are somehow tedious." + For example. . ↥∐↕∐≺↵⊳∖↩∢∙∩∐≺⇂−∩↕⋅≺⇂≺↵↕⋅⋜↕↥↽≻↥↽∐⋅∩⊸∖∐∐⋜↕↕↥∩∐⋅↺⋖↜←−⋝⋅⋃≺↵↕⋅≺↵⋜↕↥↽≻↥↽≻≺↲⋜⋃∷∖↕⇂≺↵⊳∖≺↽↓⋃⋜⋃⋅≺↲∩↥↕∐≺↲↥∐⋅⊳∖↕−∩↕⋅≺⇂≺↵↓⋅↕≺↵⊔∐⋅⋜↕∐≺⇂. yp . ⋅ terms like with appropriate normalized trispectrum or the product of normalized. bispectra multiplied. and other terms which do not depend ou phases.," For example, in the second-order approximation, ${\cal O}(\epsilon^2)$, there appears the square of the first-order term, and terms like with appropriate normalized trispectrum or the product of normalized bispectra multiplied, and other terms which do not depend on phases." + The pliases always contribute to the probability distribution by the combination of the form. σορόd-dOg). with closed wavevectors: &t4ky=0.," The phases always contribute to the probability distribution by the combination of the form, $\cos(\theta_{\sbfm{k}_1} + \cdots + +\theta_{\sbfm{k}_N})$, with closed wavevectors: $\bfm{k}_1 + \cdots + +\bfm{k}_N = 0$." + This is generally true because the phase depeudence in equation (3)) is the exponential of the sum of phases. aud the probability is the real number so that taking real parts gives the cosine BPuuction.," This is generally true because the phase dependence in equation \ref{eq03}) ) is the exponential of the sum of phases, and the probability is the real number so that taking real parts gives the cosine function." + The reason that phase correlations exlst mly among modes with closed wavevectors comes [rom the trauslational iuvariance., The reason that phase correlations exist only among modes with closed wavevectors comes from the translational invariance. + Iu equations (7)) aud (8)). waveuumbers are restricted to the ulis so that the modes in the lower half spliere are relabeled by 0j=—0κ.," In equations \ref{eq07}) ) and \ref{eq08}) ), wavenumbers are restricted to the uhs so that the modes in the lower half sphere are relabeled by $\theta_{\sbfm{k}} = - \theta_{-\sbfm{k}}$." + The moduli lagk s are easily integratedS in the first-order. approximation of equatiou1 (7)). resulting in The practically useful relatious between phasecorrelations aud the bispectrum are obtained by further inteeratinge[eJ some pliases in equation{ (0)).," The moduli $|\alpha_{\sbfm{k}}|$ 's are easily integrated in the first-order approximation of equation \ref{eq07}) ), resulting in The practically useful relations between phasecorrelations and the bispectrum are obtained by further integrating some phases in equation \ref{eq09}) )." + One obtalius, One obtains +are the same conditious satisfied bv the halo during the accretion process.,are the same conditions satisfied by the halo during the accretion process. + It is therefore well understood why the ensitv profile einereiug from a major merger coincides with the profile that would lave developed by moeaus of pure accretion., It is therefore well understood why the density profile emerging from a major merger coincides with the profile that would have developed by means of pure accretion. + Eurthermore. since the concept of orbital ecav is nieaningless dunnug violent relaxation. there is uo difference between intermediate and simall captures eemine that process apart from the distinct mass mnerease hey produce.," Furthermore, since the concept of orbital decay is meaningless during violent relaxation, there is no difference between intermediate and small captures during that process apart from the distinct mass increase they produce." + Therefore. intermediate captures must be icluded in the formal accretion rate leading to the same eusitv profile as violent relaxation.," Therefore, intermediate captures must be included in the formal accretion rate leading to the same density profile as violent relaxation." + It night be argued that the previous reasoning presunmies that violent relaxation proceeds to completion. while the relaxation time. equal to. sax. three crossing times of the svsteni at the time of the major merecr. with V;(GALAΠ113 the circular velocity of the halo. is lone enough for that condition not always to be satistied.," It might be argued that the previous reasoning presumes that violent relaxation proceeds to completion, while the relaxation time, equal to, say, three crossing times of the system at the time of the major merger, with $\Vc=(GM/R)^{1/2}$ the circular velocity of the halo, is long enough for that condition not always to be satisfied." + ILowever. the NEW expression describes the average profile of halos having a relaxed appearance.," However, the NFW expression describes the average profile of halos having a relaxed appearance." + That is. those halos observed before violent relaxation has gone to completion are not taken into account.," That is, those halos observed before violent relaxation has gone to completion are not taken into account." + Provided that halos do not undergo iutermediate captures after the completion of violent relaxation at their last major merger. their density profile will develop out as explained in reface..," Provided that halos do not undergo intermediate captures after the completion of violent relaxation at their last major merger, their density profile will develop inside-out as explained in \\ref{acc}." +" Strictly, there is no need. iu the preseut case. for aceretion to include intermediate captures. although eivon that the outermost profile is quite iuscusitive to the value of A, used (see Fig. 3))"," Strictly, there is no need, in the present case, for accretion to include intermediate captures, although given that the outermost profile is quite insensitive to the value of $\delm$ used (see Fig. \ref{3}) )" + this makes almost uo difference in practice., this makes almost no difference in practice. + Tf on the contrary. halos uuderego intermediate captures in such a late pliase. as one would naively expect from the larger frequency. of intermediate captures as compared to major ones. the density profile will become cuspicr aud deviate from the NEW profile.," If, on the contrary, halos undergo intermediate captures in such a late phase, as one would naively expect from the larger frequency of intermediate captures as compared to major ones, the density profile will become cuspier and deviate from the NFW profile." + What is therefore crucial to understand the shape of the NEW profile of relaxed halos is to soe that such intermediate captures are quite improbable., What is therefore crucial to understand the shape of the NFW profile of relaxed halos is to see that such intermediate captures are quite improbable. + This is rot in contradistinction with the relatively high average nuubers quoted by Tormen(1997)., This is not in contradistinction with the relatively high average numbers quoted by \citet{T97}. + ere we must oulv consider those intermediate captures restricted to occur after the completion of the violent relaxation of the iio accompanying its last merger carly chough for he captured clump to have completed its orbital decay o» the time the halo is observed., Here we must only consider those intermediate captures restricted to occur after the completion of the violent relaxation of the halo accompanying its last merger early enough for the captured clump to have completed its orbital decay by the time the halo is observed. + This latter bouud Is necessary. indeed. to guarantee that the halo docs rot show any substructure aud cau then be ileuti&ed as a relaxed system.," This latter bound is necessary, indeed, to guarantee that the halo does not show any substructure and can then be identified as a relaxed system." + The low frequency of intermediate captures subject to these two constraints secius to be supported by the results of simulations (Ascasibaretal. 2002)., The low frequency of intermediate captures subject to these two constraints seems to be supported by the results of simulations \citep{aygm}. +. Let us therefore calculate the probability P(M) that the last intermediate merger of a halo with A at f occured after the completion of violent relaxation aud the time at which the eveutual intermediate merger should take place for the mereed chuup to reach the halo ceuter by the time the halo is observed., Let us therefore calculate the probability ${\cal P}(M)$ that the last intermediate merger of a halo with $M$ at $t$ occurred after the completion of violent relaxation and the time at which the eventual intermediate merger should take place for the merged clump to reach the halo center by the time the halo is observed. + This probability is eiven by with the iutegrand giving the joiut probability that the halo was formed from fg to fg|dfg and underwent the last iuteriuediate mereer with a clump of iustautaucous fractional mass between A aud A|dA in the infinitesimal interval of time around f., This probability is given by with the integrand giving the joint probability that the halo was formed from $\ti$ to $\ti+\der\ti$ and underwent the last intermediate merger with a clump of instantaneous fractional mass between $\Delc$ and $\Delc+\der\Delc$ in the infinitesimal interval of time around $t$. + This joiut probability is simply equal to the product ofthe probabilities (fgA)dfe. given in equation (9)). and Α.ΑdAdf. calculated in the Appendix.," This joint probability is simply equal to the product of the probabilities $\Phi(\ti|M)\,\der\ti$, given in equation \ref{dft}) ), and $P\last(\Delc,t|M)\,\der\Delc\,\der t$, calculated in the Appendix." + Iu equation Εμ is the time of completion of violent relaxation after the last major merecr at te. with the relaxation time £44 elven by equation (13)). aud tytegeyf{A) is the time at which the intermediate merecr should occur for the chump to have reached tle halo ceuter at fy.," In equation \ref{paff}) ), $\ti+t\rel(\ti)$ is the time of completion of violent relaxation after the last major merger at $\ti$, with the relaxation time $t\rel$ given by equation \ref{trel}) ), and $t_0-t_{\rm +fric0}(\Delc)$ is the time at which the intermediate merger should occur for the clump to have reached the halo center at $t_0$." +" The characteristic time of orbital decay of a clump of mass AZ, iu a circular orbit of radius + is (e.g. DiuneyTremaine1987)) with Wor}2GM)rt? the circular velocity ator."," The characteristic time of orbital decay of a clump of mass $M\sat$ in a circular orbit of radius $r$ is (e.g., \citealt{binn87}) ) with $\Vc(r)=[GM(r)/r]^{1/2}$ the circular velocity at $r$." + ence. by integrating τας under the approximation of a singular isothermal halo (for which the iutegral is analytical) from 0 to the radius of the initial orbit. R(t). at the time £ of the merger. we have au estimate of the time spent by the chump to spiral to the halo center owing to dviuiuuical friction where AL(t) aud Wet) are. respectively. the total mass and circular velocity of the halo at the time of the merger.," Hence, by integrating $t\orb\,\der r$ under the approximation of a singular isothermal halo (for which the integral is analytical) from 0 to the radius of the initial orbit, $R(t)$, at the time $t$ of the merger, we have an estimate of the time spent by the clump to spiral to the halo center owing to dynamical friction where $M(t)$ and $\Vc(t)$ are, respectively, the total mass and circular velocity of the halo at the time of the merger." + Thus. feyey(A) is given by the previous expression for a ime f solution of the implicit equation t|fg(f.A)=fy.," Thus, $t_{\rm +fric0}(\Delc)$ is given by the previous expression for a time $t$ solution of the implicit equation $t+t\fric(t,\Delc)=t_0$." +" Note that iu deriving equation (16)). we have assmued. or siluplicity. that AZ, remains constant and equal to about of the mass of the inerged chuup after the initial tidal stripping. Le.. M,=0.3V(t)A. with A the yactional mass increase of the halo relative to the final object."," Note that in deriving equation \ref{df}) ), we have assumed, for simplicity, that $M\sat$ remains constant and equal to about of the mass of the merged clump after the initial tidal stripping, i.e., $M\sat=0.3\,M(t\merg)\,\Delc$, with $\Delc$ the fractional mass increase of the halo relative to the final object." +" According to Dekel.Devor.&IHoetzroui(2003).. the nass of cIunips is reduced to of its initial value after rein orbits have already acliieved a substantial decay,"," According to \citet{DDH03}, the mass of clumps is reduced to of its initial value after their orbits have already achieved a substantial decay." + But jese authors assune chuups with a density profile steeper iu the NEW one at large radii. which means that their chuups are more difhcult to tidallv strip than real ones.," But these authors assume clumps with a density profile steeper than the NFW one at large radii, which means that their clumps are more difficult to tidally strip than real ones." + Ou the other haud. we expect the real orbits of merecd chuups to be typically elliptical rather than circular as we are assuiune here. which should diminish their time of orbital decay.," On the other hand, we expect the real orbits of merged clumps to be typically elliptical rather than circular as we are assuming here, which should diminish their time of orbital decay." + However. this effect should be balanced by the fact that. in the elliptical case. chips fall deeper iu the halo at their poericeuter aud. hence. are more severely truucated by tides since the beeiuniug.," However, this effect should be balanced by the fact that, in the elliptical case, clumps fall deeper in the halo at their pericenter and, hence, are more severely truncated by tides since the beginning." +" The upper bound for A iu cquation (11)) is not Ay,= but Au/(11Au)0.33 since the fractional mass", The upper bound for $\Delc$ in equation \ref{paff}) ) is not $\delm=0.5$ but $\delm/(1+\delm)=0.33$ since the fractional mass +(Isockemocretal.2002).,\citep{koek02}. +. Detailed. exaiinatious of the distortion corrected data were performed. given that only subsets of the data were read out.," Detailed examinations of the distortion corrected data were performed, given that only subsets of the data were read out." + The ellipticity of isophotes fitted to 10 stellar objects across cach field never exceeded 0.08 and averaged 0.03., The ellipticity of isophotes fitted to 10 stellar objects across each field never exceeded 0.08 and averaged 0.03. + As a final step. the on and off hine nuages were scaled and subtracted from cach other in order to produce emission line maps.," As a final step, the on and off line images were scaled and subtracted from each other in order to produce emission line maps." + The radio15 source PINS 1519-79 has had imchl previous work cured ¢»t on it., The radio source PKS 1549-79 has had much previous work carried out on it. + A bref outline of its main properties wil be given before the new data are presented., A brief outline of its main properties will be given before the new data are presented. + Readers interested iu more thorough reviews are directed toward Morgantietal.(2001)... Tadluuteretal.(2001). aud Itü6.," Readers interested in more thorough reviews are directed toward \citet{morg01}, \citet{tad01} and H06." + VLBI observations reveal PEK1519-790 to have a relatively small (—150 amas) one-siceddistorted jet structure beuding through G07(INiug1990., VLBI observations reveal PK1549-79 to have a relatively small $\sim$ 150 mas) one-sideddistorted jet structure bending through \citep{king94}. + Lower resolution laps show uo evidence of larec scale structures., Lower resolution maps show no evidence of large scale structures. + The radio jet has a steep radio spectrum and originatesran from an unresolved flat spectrum radio core., The radio jet has a steep radio spectrum and originates from an unresolved flat spectrum radio core. + Altrough the jet is likely close to our line of sieht this souxὉ shows no evideuce for broad omissionu-line features. οςmitrary to the standard oricutation based Uunificatioji schemes.," Although the jet is likely close to our line of sight this source shows no evidence for broad emission-line features, contrary to the standard orientation based unification schemes." + A model describing a cocooned (obseured) quasar. Where a voung radio jet is carving a path through a dense unclear region. has been proposed by Tadlunuteretal.(2001).," A model describing a cocooned (obscured) quasar, where a young radio jet is carving a path through a dense nuclear region, has been proposed by \citet{tad01}." +. A narrow III absorption line eives a redshif of 0.152 (Morgantietal.2001.II06)., A narrow HI absorption line gives a redshift of 0.152 \citep[][H06]{morg01}. +. Prestage&Peacock(1983) were first to associate the radio source with what has been classified as ai 18.5(V) magnitude Sevort 2., \citet{pandp83} were first to associate the radio source with what has been classified as an 18.5(V) magnitude Seyfert 2. + Sieus of star formation are provided by abnormally strong fu-iufrared (FIR) emissions. i.c.. ULIRG. aud an optical contiuttun that is dominated by a population of carly type sas (Rov&Norris1997:Dickson1997 )..," Signs of star formation are provided by abnormally strong far-infrared (FIR) emissions, i.e., ULIRG, and an optical continuum that is dominated by a population of early type stars \citep{randn97,dick97}. ." + Based ou specral svuthesis modcling. 06 have shown," Based on spectral synthesis modeling, H06 have shown" + (Li|s100na101! L.) ja ," $L_{\rm ir} +[8-1000\micron] \sim 10^{14}$ $_\odot$ $\mu$ " +"yagiuentation to proceed. au additional coolne chamuel (such as IDs) is required to allow cooling below ~10119. For simplicity. we shall follow MMNV. auc assume that he eas settles to an isothermal. expoucutial disk. with eas cluperature ως, enibedded ina halo of virial temperature Jus with a Navarro. Frenk White (1997. hereafter NEW) dark matter deusitv profile.","fragmentation to proceed, an additional cooling channel (such as $H_{2}$ ) is required to allow cooling below $\sim 10^{4}$ K. For simplicity, we shall follow MMW, and assume that the gas settles to an isothermal, exponential disk, with gas temperature $T_{\rm +gas}$, embedded in a halo of virial temperature $T_{\rm vir}$ with a Navarro, Frenk White (1997, hereafter NFW) dark matter density profile." + We beein bw listing he characteristic properties of such disks. which will ο relevant for our later studies of IT) formation aud selfshielding.," We begin by listing the characteristic properties of such disks, which will be relevant for our later studies of ${\rm H_2}$ formation and self–shielding." +" Let us assuine that barvons make up the universal mass fraction Q),/Q,, of the halo. of which some yaction f4 have collapsed iuto the disk. ie. Maii=MaMa=fatQu/Odu daw."," Let us assume that baryons make up the universal mass fraction $\Omega_{\rm b}/\Omega_{\rm m}$ of the halo, of which some fraction $f_{\rm d}$ have collapsed into the disk, i.e. $M_{\rm disk}=m_{\rm d} M_{\rm halo}= f_{\rm d} (\Omega_{\rm +b}/\Omega_{\rm m}) M_{\rm halo}$ ." + We also asstune that the angular momentum of the disk J4 is some fraction jq of he halo aneular momentum J. ie. 4—j447.," We also assume that the angular momentum of the disk $J_{\rm d}$ is some fraction $j_{\rm d}$ of the halo angular momentum $J$, i.e. $J_{\rm d}=j_{\rm d} J$." + Henuceforth. we shall assune that the specific angular momentum of he disk is similar to that of the halo. and thus iii=fa.," Henceforth, we shall assume that the specific angular momentum of the disk is similar to that of the halo, and thus $m_{\rm d}=j_{\rm d}$." + The assuniption that the barvous preserve their specific aneular moment duiug collapse results iu a good fit to the observed size distribution of galactic disks (Mo. Mao White 1998).," The assumption that the baryons preserve their specific angular momentum during collapse results in a good fit to the observed size distribution of galactic disks (Mo, Mao White 1998)." + Detailed ποΊσα] simmlatious (Navarro White 1993. Navarro Steiunetz 2000) have iof supported this simple model but instead produced sienificautly simaller disks. due to the transfer of augular uonmniecntunià frou the sas to the dark matter duriug he hiehlv inhomogeneous collapse.," Detailed numerical simulations (Navarro White 1993, Navarro Steinmetz 2000) have not supported this simple model, but instead produced significantly smaller disks, due to the transfer of angular momentum from the gas to the dark matter during the highly inhomogeneous collapse." + These simulations. jowever. also fail to produce the sizes aud properties of observed galaxysized disks.," These simulations, however, also fail to produce the sizes and properties of observed galaxy–sized disks." + Iuclusion of suppression of cooling util late times (Weil. Eke Efstathiou 1998) or supernovae feedback (Thacker Couchiman 2001) reduces the discrepancy. but the issue has vet to be couclusively resolved.," Inclusion of suppression of cooling until late times (Weil, Eke Efstathiou 1998) or supernovae feedback (Thacker Couchman 2001) reduces the discrepancy, but the issue has yet to be conclusively resolved." + We therefore suuplv extend successful seiui-analytic models of disk formation at low redshift to higher redshift. aud note that future nuuerical simulations of such halos may not iu fact produce such disks.," We therefore simply extend successful semi-analytic models of disk formation at low redshift to higher redshift, and note that future numerical simulations of such halos may not in fact produce such disks." +" The hydrogen uuuber density at radius r aud at vertical height + in an isothermal exponeutial disk of radial scale leneth Rg is given by (Spitzer 1912) where s, is the ceutral deusitv. το is the vertical scale height of the disk at radius r. ος 18 the sound speed of the gas. aud sf=0.6 is the mean molecular weight (we choose the definition of the disk scale leneth Ry to couform to the customary assuniption that the surface density has au exponential profile: Xexp(rZ Ra))."," The hydrogen number density at radius $r$ and at vertical height $z$ in an isothermal exponential disk of radial scale length $R_{\rm d}$ is given by (Spitzer 1942) where $n_o$ is the central density, $z_o$ is the vertical scale height of the disk at radius $r$, $c_{\rm s}$ is the sound speed of the gas, and $\mu=0.6$ is the mean molecular weight (we choose the definition of the disk scale length $R_{d}$ to conform to the customary assumption that the surface density has an exponential profile: $\Sigma \propto {\rm +exp}(-r/R_{d})$ )." + For a disk i a halo with spin parameter A. if we asstune the barvous conserve their specific augular momentum when they collapse. then the disk scale leneth is eiven bv Ry=2ου” where regoGEray ds the radius that encloses a mean interior mass deusity of 200i. fle) and Εμίλ.σα.ja) are dimensionless fuuctious of order unity (MN). aud € is the dimensionless concentration parameter (NEW).," For a disk in a halo with spin parameter $\lambda +$, if we assume the baryons conserve their specific angular momentum when they collapse, then the disk scale length is given by $R_{\rm d} = 2^{-1/2} (j_{\rm d}/m_{\rm d}) +\lambda r_{200} f_{c}^{-1/2} f_{R} \approx \frac{\lambda}{\surd 2} r_{200}$, where $r_{200} \approx r_{\rm vir}$ is the radius that encloses a mean interior mass density of $\rho_{\rm crit}$, $f_{c}(c)$ and $f_{R}(\lambda,c,m_{\rm d},j_{\rm d})$ are dimensionless functions of order unity (MMW), and $c$ is the dimensionless concentration parameter (NFW)." +" The central number density of the gas is obtained by setting fd:[2zrdrjimynir.:)=Maga Which vields: When considering the characteristic densities for IL. formation. it will be sufficient to consider densities at nost an order of magnitude below the ceutral deusitv: eas with nomOla, comprises >50% of the mass of the disk."," The central number density of the gas is obtained by setting $\int dz +\int 2 \pi r \, dr \mu m_{\rm p} n(r,z) = M_{\rm disk}$, which yields: When considering the characteristic densities for ${\rm H_2}$ formation, it will be sufficient to consider densities at most an order of magnitude below the central density: gas with $n > 0.1 n_{o}$ comprises $> 50 \%$ of the mass of the disk." + Because hydrostatic support is only relevant iu the vertical direction. as the eas cools the disk becomes thinner. with ⋜↧↥⋅↸∖≼↧⋯⊳↸∖≼↧↴∖↴↸⊳⋜↧↕↸∖∐↸∖↕∶↴∙⊾∐↑−∙↙↗∖⊺∟⋟∓↕∖∙↽∕∏∐∖↖↽↸∖↥⋅," Because hydrostatic support is only relevant in the vertical direction, as the gas cools the disk becomes thinner, with a reduced scale height $z_{o} \propto +T_{\rm gas}^{1/2}$." +⊓↸⊳⋜↧↕⋯↕∏⋯∐ ⋅ ⊥⊐↴ ⋅ deusitv of gas as à function of radius 1x: which is sufficient for selfshiclding of the gas against both ionizing UV radiation aud II2 dissociating raciatiou oei the LW bands to become important., The vertical column density of gas as a function of radius is: which is sufficient for self-shielding of the gas against both ionizing UV radiation and H2 dissociating radiation in the LW bands to become important. + Note that the column density is iudepeudent of the eas temperature., Note that the column density is independent of the gas temperature. + We now cousider the conditions for gravitational instability of the disk., We now consider the conditions for gravitational instability of the disk. +" In computing the rotation curve —“Gry. we use the formalism, of MMW. which takes iuto account the contraction induced in the inner regions of the halo by the cooling aud formation of the disk."," In computing the rotation curve $V(r)$, we use the formalism of MMW, which takes into account the contraction induced in the inner regions of the halo by the cooling and formation of the disk." + This is cone by assuming the disk is assembled slowly aud the aneular momentum of dark matter particles is an adiabatic invariant (Blumenthal et al., This is done by assuming the disk is assembled slowly and the angular momentum of dark matter particles is an adiabatic invariant (Blumenthal et al. + 1986. Flores et al.," 1986, Flores et al." + 1993)., 1993). + For the disk to be locally eravitationally unstable despite the stabilizing effects of tidal shears and pressure forces. we require the Toomre paramcter ο<1. where (oe. Binney Tremaine 1987) Mods the disk surface iaass densitvo and KO=Layyl}dluV/dlur)?? is the epievclic frequency.," For the disk to be locally gravitationally unstable despite the stabilizing effects of tidal shears and pressure forces, we require the Toomre parameter $Q <1$, where (e.g. Binney Tremaine 1987) $\Sigma$ is the disk surface mass density, and $\kappa = 1.41 (V/r) (1 ++ {\rm d \, ln} V/{\rm d \, ln} r)^{1/2}$ is the epicyclic frequency." + Reeious iu local disk ealaxies where Q>1 are observationallv associated with very little star formation. inclicating that the Toomme criterion is obeved remarkably well (I&eunicutt 1989).," Regions in local disk galaxies where $Q>1$ are observationally associated with very little star formation, indicating that the Toomre criterion is obeyed remarkably well (Kennicutt 1989)." + For our purposes. if (Q>1 everywhere throughout the disk. it is exavitationallv stable and we asse no star formation takes place.," For our purposes, if $Q > 1$ everywhere throughout the disk, it is gravitationally stable and we assume no star formation takes place." + Disks with Heh spin paraicters have low surface deusities aud satisfv lus criterion., Disks with high spin parameters have low surface densities and satisfy this criterion. + For any eiven diskhalo svsteni we cau calculate a critical spinparameter Air for Which the disk is uarginallv stable.," For any given disk–halo system, we can calculate a critical spinparameter $\lambda_{\rm crit}$ for which the disk is marginally stable." + Were we define Agar by the requirement hat Q attains a müiniuuii value of Q=1 at least at one position in the disk., Here we define $\lambda_{\rm crit}$ by the requirement that $Q$ attains a minimum value of $Q=1$ at least at one position in the disk. +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴∖"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴∖↴"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴∖↴∙"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴∖↴∙∖"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +" The fraction of halos that remain dark can therefore be found by integrating over the spin mramcter distribution. fq=""m pHA. where ptA) Is Given by with A=0.05 aud σν=0.5 (e.o ↕≧⋜∐⋅∐↸∖↴∖↴∙∖↽"," The fraction of halos that remain dark can therefore be found by integrating over the spin parameter distribution, $f_{\rm dark} = \int_{\lambda_{\rm crit}}^{\infty} +p(\lambda) d\lambda$ , where $p(\lambda)$ is given by with $\bar{\lambda}= 0.05$ and $\sigma_{\lambda} =0.5$ (e.g. Barnes" +orbital ephemeris is an ellective technique as the entire data set ds used together (Paulctal.2002).,orbital ephemeris is an effective technique as the entire data set is used together \citep{Paul02}. +.. Eclipse timing is another technique. used to determine the orbital evolution of LAINBs (Nollctal.2009:Jainet2010).," Eclipse timing is another technique, used to determine the orbital evolution of LMXBs \citep{Wolff09, Jain10}." +. In. some sources. in the absence of pulsations or eclipses. orbital period derivative has been measured with some stable orbital intensity modulation features (4U. 1820-30: (2001): €vg X-3: Singhetal. (2002))).," In some sources, in the absence of pulsations or eclipses, orbital period derivative has been measured with some stable orbital intensity modulation features (4U 1820-30: \citet{Chou01}; Cyg X-3: \citet{Singh02}) )." + ATE JITIO281 is one of the very lew LAINBs. where full. sharp N-rav. eclipses have been observed.," XTE J1710–281 is one of the very few LMXBs, where full, sharp X-ray eclipses have been observed." + The other svstems being. EXO 0748676. (Αλλοetal. 2002).. CRS J1I747312 (intZandοἱal.2003)... AINB 1658298 (Cominsky&Wood1980). and ΑΝ J1745.62901 etal. 1996).," The other systems being, EXO 0748–676 \citep{Wolff02}, , GRS J1747–312 \citep{Zand03}, MXB 1658–298 \citep{Cominsky89} and AX J1745.6–2901 \citep{Maeda96}." +.. Among these sources. EXO 0748.676 (Parmaretal.1986:Woll£2002). is the only system in which a large number of eclipses have been timed with high accuracy Λοctal.2009).," Among these sources, EXO 0748–676 \citep{Parmar86, Wolff02} is the only system in which a large number of eclipses have been timed with high accuracy \citep{Wolff09}." +. In case of GRS J1747312 and AX J1745.62901. the eclipse duration is too long (~43 minutes and ~23 minutes. respectively. intZandetal.(2003):Por-quetetal. (2007))) to carry out monitoring measurements.," In case of GRS J1747–312 and AX J1745.6–2901, the eclipse duration is too long $\sim$ 43 minutes and $\sim$ 23 minutes, respectively, \citet{Zand03, Porquet07}) ) to carry out monitoring measurements." + AX J1745.62901 is located near a bright source. which puts strong constraints on the timing analysis.," AX J1745.6–2901 is located near a bright source, which puts strong constraints on the timing analysis." + The fifth source. MXD 1658298 has been mostlv in inactive state. since its discovery (Wachterctal.2000:Oosterbroeketal. 2001).," The fifth source, MXB 1658–298 has been mostly in inactive state, since its discovery \citep{Wachter00, Oosterbroek01}." +. Phe LAINB NTE J1710281 has a fairly small eclipse duration and has been persistently active since its discovery., The LMXB XTE J1710–281 has a fairly small eclipse duration and has been persistently active since its discovery. + This makes it an ideal source to investigate the orbital evolution., This makes it an ideal source to investigate the orbital evolution. + In this work. we have used the sharp eclipses of NTE J17T10281 as timing markers to determine the orbital period of the system.," In this work, we have used the sharp eclipses of XTE J1710–281 as timing markers to determine the orbital period of the system." + By measuring the mid-eclipse times. over a lone time baseline of ~ 11 vears. we can determine the change in the orbital period and hence estimate the orbital evolution of the system.," By measuring the mid-eclipse times, over a long time baseline of $\sim$ 11 years, we can determine the change in the orbital period and hence estimate the orbital evolution of the system." + Data [or the present analysis were obtained from. observations made with the Proportional Counter —Array (PCA) on board the Rossi X-ray Viming Explorer (/NTE)) satellite (Bracltetal.1993)., Data for the present analysis were obtained from observations made with the Proportional Counter Array (PCA) on board the Rossi X-ray Timing Explorer ) satellite \citep{Bradt93}. +.. The TE--DPC€CNX consists of an array of five collimated xenonmethane multianocde proportional counter units (DCUJ. with a total photon collection. area of 6500 cm. (Jahoclaet.al.1996.2006).. depending on the number of PCUs ON.," The -PCA consists of an array of five collimated xenon/methane multianode proportional counter units (PCU) with a total photon collection area of 6500 $^{2}$ \citep{Jahoda96, Jahoda06}, depending on the number of PCUs ON." + The entire analysis was done using [rom the astronomy software owekase UEASOFLT-ver 6.10., The entire analysis was done using from the astronomy software package -ver 6.10. + The PCA data collected. in he event mode and the Good Xenon. mode. were used o generate the light curves. using theSEENTHCT.," The PCA data collected in the event mode and the Good Xenon mode, were used to generate the light curves, using the." +. ‘The analysis was done in the energy band 20 keV. The rvackerouncl was estimated. using thePCABACKEST., The analysis was done in the energy band $-$ 20 keV. The background was estimated using the. +. Faint source model was taken from the website news.html)., Faint source model was taken from the website $_{-}$ news.html). +" ""hereafter. the background. subtracted ligh CULVOS Were xwvcenter corrected using theFANBARY."," Thereafter, the background subtracted light curves were barycenter corrected using the." +. We have analvzed all the TE-PCAX archival data. which covered a full X-ray eclipse.," We have analyzed all the -PCA archival data, which covered a full X-ray eclipse." + However. since this is a bursting source. we ignored few eclipses. where bursts occurred close to the ingress and egress phase of the eclipse.," However, since this is a bursting source, we ignored few eclipses, where bursts occurred close to the ingress and egress phase of the eclipse." + From data spread over 711 vears (1999-2010). we have found 57 complete eclipses.," From data spread over $\sim$ 11 years (1999-2010), we have found 57 complete eclipses." + Table 1 gives the observation IDs of all the 57 eclipses. observed. with TE--CA.," Table 1 gives the observation IDs of all the 57 eclipses, observed with -PCA." + In Figure l. we have shown a sample background: subtractecl light curve of NTE J1710-281 (Obs LD: 91045-01-01-13). binned with 2 seconds and including an eclipse lasting for 420 s. excluding the ingress and. egress phase.," In Figure 1, we have shown a sample background subtracted light curve of XTE J1710-281 (Obs ID: 91045-01-01-13), binned with 2 seconds and including an eclipse lasting for $\sim$ 420 s, excluding the ingress and egress phase." + For most of the observations. apart from a Lew tvpe- X-ray bursts. the out-of-eclipse count rate did not seem to have anv significant variability.," For most of the observations, apart from a few type-I X-ray bursts, the out-of-eclipse count rate did not seem to have any significant variability." + DPherefore. the variable components of the light curve around the eclipse phase are. the pre-ingress (Csingress). eclipse (νο). ancl post-egress (Chast egress) count rates: the ingress (Δεν). and egress GM.) duration. the eclipse duration (Af...) and the micd-eclipse time.," Therefore, the variable components of the light curve around the eclipse phase are, the pre-ingress $_{pre-ingress}$ ), eclipse $_{eclipse}$ ), and post-egress $_{post-egress}$ ) count rates; the ingress $\Delta t_{in} $ ) and egress $ \Delta t_{eg} $ ) duration, the eclipse duration $ \Delta t_{ecl} $ ) and the mid-eclipse time." + C'onsidering all the components to be freely variable. we first fitted a seven-parameter ramp ancl step model to the eclipse phase (similar to Wollletal. (2009))).," Considering all the components to be freely variable, we first fitted a seven-parameter ramp and step model to the eclipse phase (similar to \citet{Wolff09}) )." + ]t was founc that the value of Cipoqress AMA Castgross were similar and the eclipse ingress and egress duration were also similar withing errors., It was found that the value of $_{pre-ingress}$ and $_{post-egress}$ were similar and the eclipse ingress and egress duration were also similar withing errors. +" The parameter space. thus got reduced. to five: the pre-ingress and the post-egress count rate (Cyniopest og). the eclipse count rate. the ingress ancl egress duration (GM;,, 4). the eclipse duration (λέων). and the mid-eclipse time."," The parameter space, thus got reduced to five: the pre-ingress and the post-egress count rate $_{pre-in-post-eg}$ ), the eclipse count rate, the ingress and egress duration $ \Delta t_{in_eg} $ ), the eclipse duration $ \Delta t_{ecl} $ ), and the mid-eclipse time." + The average eclipse duration in NATE J1710-281 is — 420 s: and for the model fiting. 150 seconds of data was taken before and. after the eclipse phase.," The average eclipse duration in XTE J1710-281 is $\sim$ 420 s; and for the model fitting, $\sim$ 150 seconds of data was taken before and after the eclipse phase." + Lt was also seen that the error in the mid eclispe time measurcment was smaller. when the five-parameter mocel was used., It was also seen that the error in the mid eclispe time measurement was smaller when the five-parameter model was used. + The best fit mocel for the eclipse profile in Figure l. is shown with a solid linc.," The best fit model for the eclipse profile in Figure 1, is shown with a solid line." +" ""Phe best fit had à 47 of 381 for 375 degrees of freedom.", The best fit had a $ \chi^{2} $ of 381 for 375 degrees of freedom. + All the 57 X-ray eclipse light. curves were fitted with a five-parameter ramp function. as described: above.," All the 57 X-ray eclipse light curves were fitted with a five-parameter ramp function, as described above." + The müd-eclipse times and the corresponding le errors were determined., The mid-eclipse times and the corresponding $\sigma$ errors were determined. + The results are given in Table 1., The results are given in Table 1. + Phe errors in the mid-eclipse times vary between 0.000006 d to 0.000064 d. Le. 0.5 s to 5.5 s. This is mainly due to cilference in the relative count rates of the source and the background. ancl partly dueto the number of detectors ON.," The errors in the mid-eclipse times vary between 0.000006 d to 0.000064 d, i.e. 0.5 s to 5.5 s. This is mainly due to difference in the relative count rates of the source and the background, and partly dueto the number of detectors ON." + The orbit numbers are with respect to the first. eclipse detected. in, The orbit numbers are with respect to the first eclipse detected in + The second image is the positive image located outside (he critical curve where it is squeezecl in.,-1 - The second image is the positive image located outside the critical curve where it is squeezed in. + The third image is a negative image located inside (he critical curve where it is bulged out., The third image is a negative image located inside the critical curve where it is bulged out. + The first image is the dim image located near the lens position 2=0., The first image is the dim image located near the lens position $z=0$. + The first and third images become degenerate when w—x., The first and third images become degenerate when $\omega\rightarrow \infty$. + One can see that the third image moves fast and (he first image slowly., One can see that the third image moves fast and the first image slowly. + Because the caustic is a point caustic. all the Uivee image trajectories are continuous.," Because the caustic is a point caustic, all the three image trajectories are continuous." + Physicallv. the first image is lgnorable. and practically there are (wo images as is (he case in the single lens.," Physically, the first image is ignorable, and practically there are two images as is the case in the single lens." + The full lens equation Gn the second order) in eq.(3.2.1)) has a triple pole at the main lens position 2=0 which are image positions of w=o., The full lens equation (in the second order) in \ref{eqApproxLeqTwo}) ) has a triple pole at the main lens position $z=0$ which are image positions of $\omega = \infty$. + Two of them are ignorable because they are confined to the very close proximity to the lens position and have negligible [Iuxes for finite w., Two of them are ignorable because they are confined to the very close proximity to the lens position and have negligible fluxes for finite $\omega$. + In order to find the approximate positions of the ienorable images for small w (in (he near proximity of the caustic). set 2=ο in the lens equation (3.2.1)) and find 4 for which the large terms (depending on the positive power of £) add to zero.," In order to find the approximate positions of the ignorable images for small $\omega$ (in the near proximity of the caustic), set $z = A/\ell$ in the lens equation \ref{eqApproxLeqTwo}) ) and find $A$ for which the large terms (depending on the positive power of $\ell$ ) add to zero." + There are two solutions., There are two solutions. +" ""EE arava The position of the corresponding source is Q(1/6).", z = ) The position of the corresponding source is ${\cal O}(1/\ell)$. + Therefore there are (wo or four images 1n. practice., Therefore there are two or four images in practice. + The lens equation (3.2.1)) can be converted into an analvlic equation and it is a fifth order equation when truncated in the second order in 1/6., The lens equation \ref{eqApproxLeqTwo}) ) can be converted into an analytic equation and it is a fifth order equation when truncated in the second order in $1/\ell$ . + One of (he solutions is an ignorable image., One of the solutions is an ignorable image. + The caustic is centered αἱ οcd(e()!., The caustic is centered at $\omega \sim d(\epsilon\ell)^{-1}$. + A galaxy lens (of finite mass and finite extension) can be expressed in terms of its projected mass density Mo where A/ is the total mass and σ is thenormalized projected mass density., A galaxy lens (of finite mass and finite extension) can be expressed in terms of its projected mass density $M\sigma$ where $M$ is the total mass and $\sigma$ is thenormalized projected mass density. +, = 1 +"black and red crosses represent, respectively, the unweighted mean metallicities, ([Fe/H]), of sub-DLAs in 5 redshift bins containing 16 or 12 systems each and of DLAs in 6 redshift bins containing 27 or 26 systems each.","black and red crosses represent, respectively, the unweighted mean metallicities, $\langle +[{\rm Fe/H}]\rangle$, of sub-DLAs in 5 redshift bins containing 16 or 12 systems each and of DLAs in 6 redshift bins containing 27 or 26 systems each." +The loc errors of the means include sampling and measurement uncertainties and are calculated using the statistical techniques described in(2002).,"The $1\,\sigma$ errors of the means include sampling and measurement uncertainties and are calculated using the statistical techniques described in." +". The unweighted mean metallicities as well as the column density weighted mean metallicities, [(Fe/H)]=log[>;;10!e/#l NCA);/Y; ;], derived in each redshift bin are tabulated in online material."," The unweighted mean metallicities as well as the column density weighted mean metallicities, $[\langle {\rm Fe/H}\rangle] = +\log [\sum_i 10^{{\rm [Fe/H]}_i} N$ $_i / \sum_i N$ $_i]$, derived in each redshift bin are tabulated in online material." +" The unweighted means refer to the average metallicity of galaxies, while theH1 weighted means refer to the mass-weighted metallicity of neutral gas."," The unweighted means refer to the average metallicity of galaxies, while the weighted means refer to the mass-weighted metallicity of neutral gas." +" Since the bulk of N(H 1) is not in sub-DLAs, the unweighted means are the appropriate quantity to trace and compare the metallicity evolution of the two samples of absorbers."," Since the bulk of $N$ ) is not in sub-DLAs, the unweighted means are the appropriate quantity to trace and compare the metallicity evolution of the two samples of absorbers." +" Clearly, the mean [Fe/H] metallicity of sub-DLAs at low redshifts is substantially larger than that of DLAs by about 0.7 dex (4c significance level)."," Clearly, the mean [Fe/H] metallicity of sub-DLAs at low redshifts is substantially larger than that of DLAs by about 0.7 dex $4\,\sigma$ significance level)." + At z>1.7 the mean metallicity of sub-DLAs drops by 0.5 dex and reaches a value similar to the mean metallicity of DLAs., At $z>1.7$ the mean metallicity of sub-DLAs drops by 0.5 dex and reaches a value similar to the mean metallicity of DLAs. + This reveals two major differences in the metallicity evolution of sub-DLAs and DLAs., This reveals two major differences in the metallicity evolution of sub-DLAs and DLAs. +" First,DLAs."," First,." +". Performing a least-square fit to the unweighted mean [Fe/H] metallicity versus redshift data points, we find for sub-DLAs a slope m=—0.41+0.07 dex/Az and a zero point b dex, while the slope for DLAs is less than half this value, m=—0.19+0.04 dex/Az, with a zero point b=—1.03+0.09 dex."," Performing a least-square fit to the unweighted mean [Fe/H] metallicity versus redshift data points, we find for sub-DLAs a slope $m = -0.41\pm 0.07$ $\Delta z$ and a zero point $b = -0.26\pm 0.13$ dex, while the slope for DLAs is less than half this value, $m = -0.19\pm 0.04$ $\Delta z$, with a zero point $b = -1.03\pm 0.09$ dex." +" The best-fit slopes to the1 weighted mean metallicities are identical within statistical uncertainties, both for sub-DLAs and DLAs, to the slopes derived for the unweighted mean metallicities."," The best-fit slopes to the weighted mean metallicities are identical within statistical uncertainties, both for sub-DLAs and DLAs, to the slopes derived for the unweighted mean metallicities." +" They underline the robustness of this result, previously claimed by and(2007)."," They underline the robustness of this result, previously claimed by and." +". Second,redshifts."," Second,." +. This is well illustrated in the two panels on the right-hand side of Fig., This is well illustrated in the two panels on the right-hand side of Fig. +" 1 that show the [Fe/H] distributions of sub-DLAs and DLAs in two separate redshift intervals, z«1.7 and z>1.7."," \ref{metallicity-evolution} that show the [Fe/H] distributions of sub-DLAs and DLAs in two separate redshift intervals, $z<1.7$ and $z>1.7$." +" At z«1.7 sub-DLAs are clearly more metal-rich than DLAs, while at z>1.7 this large metallicity difference is no longer apparent (see the respective cumulative functions)."," At $z<1.7$ sub-DLAs are clearly more metal-rich than DLAs, while at $z>1.7$ this large metallicity difference is no longer apparent (see the respective cumulative functions)." + The KS test shows that the null-hypothesis that the metallicities of the low-redshift sub-DLAs and DLAs are drawn from the same population can be rejected at c.]., The KS test shows that the null-hypothesis that the metallicities of the low-redshift sub-DLAs and DLAs are drawn from the same population can be rejected at c.l. + This c.l., This c.l. + falls to an inconclusive for the metallicity distributions of the high-redshift sub-DLAs and DLAs., falls to an inconclusive for the metallicity distributions of the high-redshift sub-DLAs and DLAs. +" This suggests that at high redshifts sub-DLAs have similar metallicitiesto DLAs, which diverge from that of DLAs at low redshifts only."," This suggests that at high redshifts sub-DLAs have similar metallicitiesto DLAs, which diverge from that of DLAs at low redshifts only." +" With a mean metallicity [Zn/H] =40.18+0.11 in sub-DLAs and [Zn/H] =—0.66+0.09 in DLAs at z<1.7, the [Zn/H] measurements further confirm the bias of low-redshift sub-DLAs towards higher metallicities relatively to DLAs."," With a mean metallicity [Zn/H] $=+0.18\pm 0.11$ in sub-DLAs and [Zn/H] $= -0.66\pm 0.09$ in DLAs at $z< 1.7$, the [Zn/H] measurements further confirm the bias of low-redshift sub-DLAs towards higher metallicities relatively to DLAs." +" found a correlation between the velocity widths, Av, of lines of low-ionization species and metallicities, over two orders of magnitude in metallicity, for DLAs and a few high-N(H1) sub-DLAs at 1.7. oscillations.", The analysis of the LSND data through 1995 \cite{LSND96} strengthens the earlier LSND signal for $\bar\nu_\mu \rightarrow \bar\nu_e$ oscillations. +" Comparison with exchision plots from other experiments inplies a lower nuit Am?fee=Jey,2ome|z-0.2E eV7. Toiuplving. in. firn a lower linüt m,20.15 eV. or O,2050.513,"," Comparison with exclusion plots from other experiments implies a lower limit $\Delta m^2_{\mu e} \equiv +|m(\nu_\mu)^2-m(\nu_e)^2| \gsim 0.2$ $^2$ , implying in turn a lower limit $m_\nu \gsim 0.45$ eV, or $\Omega_\nu \gsim 0.02 (0.5/h)^2$." + This nuplies that the contribution of hot dark luatter o the cosmological deusitv is larger than tiat of all the visible stars (ο.zx0001 [15]))., This implies that the contribution of hot dark matter to the cosmological density is larger than that of all the visible stars $\Omega_\ast \approx 0.004$ \cite{Peebles}) ). +" More data aud analysis are nec.led frou: LSNDs MyM, channel before the initial hint |16] that Any,mG eV? can be confirmed.", More data and analysis are needed from LSND's $\nu_\mu \rightarrow \nu_e$ channel before the initial hint \cite{Cald95} that $\Delta m^2_{\mu e} \approx 6$ $^2$ can be confirmed. + Fortunaely the KARAIEN experiment has just addvd shiclding to decrease its backegrotuid so that it can probe the same region of Aur2t and mining anele. with seusitivitv as ercat as LSND's wihin about two vears.," Fortunately the KARMEN experiment has just added shielding to decrease its background so that it can probe the same region of $\Delta m^2_{\mu e}$ and mixing angle, with sensitivity as great as LSND's within about two years." + TheINamiokanucde data |17 showing that the deficit of Eo>1.3 GeV atmospleric ΠΙΟ neutrinos, TheKamiokande data \cite{KamAtmNu} showing that the deficit of $E > 1.3$ GeV atmospheric muon neutrinos +"For all ofthem we have at haud the photometric fluxes in the R-band from the USNO-D1 catalog. the J. £7. aud Ks bauds. from the Two Micron All Sky Survey (ΛΙΓΑΟΦ), as well as fluxes at 3.6. 1.5. 5.8. 8.0. 21. and 70 jun. obtained byBits.","For all of them we have at hand the photometric fluxes in the $R$ -band from the USNO-B1 catalog, the $J$, $H$, and $K_S$ bands, from the Two Micron All Sky Survey (2MASS), as well as fluxes at 3.6, 4.5, 5.8, 8.0, 24, and 70 $\mu$ m obtained by." +.. We have assumed errors that are within observational standards (S$ 20%)., We have assumed errors that are within observational standards $\la 20 \%$ ). + Tn addition. all of our. targets have heen observed with the Subiuillimeter Array (SALA).," In addition, all of our targets have been observed with the Submillimeter Array (SMA)." + For two of them (Tran 21 aud 32) we obtaimed the fluxes at 1.3 uuu. while for Tran 11 aud 31 we could only derive upper limits (see Paper I for details on SED data).," For two of them (Tran 21 and 32) we obtained the fluxes at 1.3 mm, while for Tran 11 and 31 we could only derive upper limits (see Paper I for details on SED data)." + For Tran 32. νου Was given in Nutter ct al. (," For Tran 32, $F_{850\,\mu{\rm m}}$ was given in Nutter et al. (" +2006).,2006). + One of our targets. Tran Ll. was receutly included in the study ofcold disks (disks with laree ιοί. dust holes) performed by Moerfun et al (2010). 1.6. their source 31.," One of our targets, Tran 11, was recently included in the study of disks (disks with large inner dust holes) performed by n et al (2010), i.e. their source 24." +" Tt was modeled including a positive detection front MIPS photometry: Prou,=537Er uy. aud the measured Spitver—IRS spectrum. but no constraints of the flux at (9ub)-nuüllimeter wavelengths."," It was modeled including a positive detection from MIPS photometry: $F_{70 \mu{\rm m}}= 537\pm78.8$ mJy, and the measured -IRS spectrum, but no constraints of the flux at (sub)-millimeter wavelengths." + For the sake of completeness. we have redone the fit to the SED of this source with our set of free parameters and inchiding the 1.3 nimi fux upper Tait.," For the sake of completeness, we have redone the fit to the SED of this source with our set of free parameters and including the 1.3 mm flux upper limit." + The modeling that we have applied is similar to those performed by Andrews et al. (, The modeling that we have applied is similar to those performed by Andrews et al. ( +2009. 20108) and Brown et al. (,"2009, 2010a) and Brown et al. (" +2009). who have confirmed their plysical estimates from SED iodcling through direct imaging.,"2009), who have confirmed their physical estimates from SED modeling through direct imaging." + A 2D structive model for flared disks is combined with tbe Monte Carlo coutinmun radiative transfer package RADMC v3.1 (Dullemond Dominik. 2001). modified to include a density reduction as an inner cavity.," A 2–D structure model for flared disks is combined with the Monte Carlo continuum radiative transfer package RADMC v3.1 (Dullemond Dominik, 2004), modified to include a density reduction as an inner cavity." + The code computes a temperature structure consistent with the eiven density profile. and in equilibiun with the radiation by the ceutral star.," The code computes a temperature structure consistent with the given density profile, and in equilibrium with the irradiation by the central star." +" The disk is presumed to be passive. au assumption that is supported by the low disk to stellar luminosity ratios of our sample. X0.005 according to estimates in Paper We consider a surface density profile characterized by a power-law, XxHR. with au exponential taper at lavecr radii (xe.1/0) where RB. is the characteristic radius."," The disk is presumed to be passive, an assumption that is supported by the low disk to stellar luminosity ratios of our sample, $\la$ 0.005 according to estimates in Paper We consider a surface density profile characterized by a power-law, $\Sigma \propto R^{-\gamma}$ , with an exponential taper at larger radii $\propto e^{-(R/R_c)^{2-\gamma}}$ ), where $R_c$ is the characteristic radius." + This is phiesicallv motivated by the success of siuularity solutions of viscous disks to reproduce the observed gradual density decay at large radii (Iughes et al., This is physically motivated by the success of similarity solutions of viscous disks to reproduce the observed gradual density decay at large radii (Hughes et al. + 2008)., 2008). + X dis normalized to obtain the total mass of the disk. αμ when integrated.," $\Sigma$ is normalized to obtain the total mass of the disk, $M_d$, when integrated." + The radial iudex was fixed to be 3=L which is a typical value within the range 5=0.11.1 established ly Audrews et al (20102)., The radial index was fixed to be $\gamma=1$ which is a typical value within the range $\gamma = 0.4 - 1.1$ established by Andrews et al (2010a). + Our option could be questioned in the light of results by Isclla ct al. (, Our option could be questioned in the light of results by Isella et al. ( +2009) who have independently inferred slopes from steep to quite shallow in their sample of spatially resolved Resolved images are therefore mandatory to obtain qnore accurate estimates of 5 for our particular tarects.,2009) who have independently inferred slopes from steep to quite shallow in their sample of spatially resolved Resolved images are therefore mandatory to obtain more accurate estimates of $\gamma$ for our particular targets. + Woe set the characteristic radius R..=100 AU., We set the characteristic radius $R_c= 100$ AU. + However. there is no spatially resolved information iu the SEDs alone. and the data can be reproduced equally well with a wide range of outer disk values (Andrews ct al.," However, there is no spatially resolved information in the SEDs alone, and the data can be reproduced equally well with a wide range of outer disk values (Andrews et al." + 20105)., 2010b). + The value we choose is represcutative of the disks with resolved interferometric visibilitics. that are HR.=it198 AU (Andrews et al.," The value we choose is representative of the disks with resolved interferometric visibilities, that are $R_c=14-198$ AU (Andrews et al." +" 2009. 2010b) iu Ophiuchus. aud BR,z30230 AU for Toiurus-Auriga (sella. Carpenter Sargeut. 2009)."," 2009, 2010b) in Ophiuchus, and $R_c\simeq 30-230$ AU for Taurus-Auriga (Isella, Carpenter Sargent, 2009)." + Note that lareer outer radi (100 - 1100 AU) have been obtained with different fitting techniques (sharply truncated power law fits to CO observations) iud are not directly comparable with the A. values (sec Williams Cieza 2011)., Note that larger outer radii (100 - 1100 AU) have been obtained with different fitting techniques (sharply truncated power law fits to CO observations) and are not directly comparable with the $R_c$ values (see Williams Cieza 2011). + Aside from the extreme case of a ucarly edge-on viewing angle. the disk inchnation cannot be determined from waresolved obscrvations.," Aside from the extreme case of a nearly edge-on viewing angle, the disk inclination cannot be determined from unresolved observations." + Scattered light nuages have proben useful in this seuse (Pinte ct al., Scattered light images have proben useful in this sense (Pinte et al. + 2007)., 2007). + We have set an intermediate representative inclination ;=30° in Our There 1 strong observational evidence that circumstellar dust can preseut some degree of settling to the midplane (Furlan et al., We have set an intermediate representative inclination $i=30^\circ$ in our There is strong observational evidence that circumstellar dust can present some degree of settling to the midplane (Furlan et al. + 2006. λοςπιο et al.," 2006, McClure et al." + 2010). which is in aereciment with theoretical predictions(6.9..," 2010), which is in agreement with theoretical predictions(e.g.," + Dominik et al., Dominik et al. + 2007 aud references)., 2007 and references). + Indeed. the erowth," Indeed, the growth" +Turushek. D. et al.,"Turnshek, D. et al." + 1991. ApJ. 128. 93 Uhich. ALT. 2000. ASAApBev. in press Wang. T.. Brinkmann. W.. Beregcron. J. 1996. Ap. 309. 81 Yuan. W.. Siebert. J. Brinkmann. W. 1995. Ap. 331. 198," 1994, ApJ, 428, 93 Ulrich, M-H. 2000, ApRev, in press Wang, T., Brinkmann, W., Bergeron, J. 1996, Ap, 309, 81 Yuan, W., Siebert, J. Brinkmann, W. 1998, Ap, 334, 498" +method similar to the ΑΛΑ ΝΕΟΙ case. aud applied a correction lor the degradation in the low-enerey equantum elficienev of the ACIS chips since the launch.,"method similar to the ${\it XMM}$ ${\it Newton}$ case, and applied a correction for the degradation in the low-energy quantum efficiency of the ACIS chips since the launch." + The backeround spectrum was collected [rom the corresponding 33 region of blank-skv observation datasets (exposure 450 ks) provided bv the ACIS calibration team., The background spectrum was collected from the corresponding S3 region of blank-sky observation datasets (exposure 450 ks) provided by the ACIS calibration team. + It is known that the background rates of blank-sky. data vary. from observation to observation. by up lo ~ in energies below 5 keV. However. this uncertainty can be ignored in the 0.452 keV enerev band. because the observed count rate is an order of magnitude higher than the background level.," It is known that the background rates of blank-sky data vary, from observation to observation, by up to $\sim$ in energies below 5 keV. However, this uncertainty can be ignored in the 0.45–2 keV energy band, because the observed count rate is an order of magnitude higher than the background level." + In the 25 keV band where the background svstematics are more severe. we confirmed that the blank-skv subüracted spectrum obtained in this wav agrees. within6%... with that obtained by subtracting another background. acquired with the other illuminated $1 chip (excluding two point sources) during the same on-source observation: ihe SL and $3 backgrounds are generally known (o coincide with each other within in (his energy band (http://hea-www.harvard.edu/mmaxim/axal/acisbg/data/README).," In the 2–5 keV band where the background systematics are more severe, we confirmed that the blank-sky subtracted spectrum obtained in this way agrees, within, with that obtained by subtracting another background acquired with the other back-illuminated S1 chip (excluding two point sources) during the same on-source observation; the S1 and S3 backgrounds are generally known to coincide with each other within in this energy band maxim/axaf/acisbg/data/README)." + Therefore. we here analvze (he energy spectrum over (he total 0.457 keV energy band. after subtracting the blank-skv background spectrum.," Therefore, we here analyze the energy spectrum over the total 0.45–7 keV energy band, after subtracting the blank-sky background spectrum." + Because of relatively poor statistics. we do nol utilize (hese Chandra data to quantify the spectra of the detected sources.," Because of relatively poor statistics, we do not utilize these ${\it Chandra}$ data to quantify the spectra of the detected sources." + Although SIZAOI already analvzed the summed spectra of the detected point sources with a single power-law (PL for short) model. we here repeat the analvsis for consistency. by consiclering a small amount of flux contribution from the dilhise X-ray emission.," Although SEA01 already analyzed the summed spectra of the detected point sources with a single power-law (PL for short) model, we here repeat the analysis for consistency, by considering a small amount of flux contribution from the diffuse X-ray emission." + Figure shows (he MOS and PN spectra obtained in § 2.1. as à sum over the 92 sources detected within 6.," Figure \ref{fig:lmxb} shows the MOS and PN spectra obtained in $\S$ 2.1, as a sum over the 92 sources detected within $6\arcmin$." + As described in § 2. they have been collected [rom regions of which the summed area amounts to e of the total 6’ radius region.," As described in $\S$ 2, they have been collected from regions of which the summed area amounts to $\sim$ of the total $6\arcmin$ radius region." + According to the spectral decomposition in Paper 1. more than half the overall flux of M 31 below ~2 keV comes from (he diffuse component.," According to the spectral decomposition in Paper 1, more than half the overall flux of M 31 below $\sim 2$ keV comes from the diffuse component." + From these (wo facts. the dilfuse X-ray emission is expected to contribute >5% to the spectra of Figure B3..," From these two facts, the diffuse X-ray emission is expected to contribute $> 5$ to the spectra of Figure \ref{fig:lmxb}." + With this in mind. we analvzed the 0.4.10 keV portion of the MOS and PN spectra of Figure 3..," With this in mind, we analyzed the 0.4–10 keV portion of the MOS and PN spectra of Figure \ref{fig:lmxb}." + We first fitted the spectra above 2 keV. where point sources are dominant (Paper 1). emploving the physical LAINB model (a DBB plus a BB) to represent huminous point sources in M 31 (8EO 1).," We first fitted the spectra above 2 keV, where point sources are dominant (Paper 1), employing the physical LMXB model (a DBB plus a BB) to represent luminous point sources in M 31 $\S$ 1)." +" We allowed the photoelectric absorption column density (ο vary. freely. but constrained it to be higher than Ni=6.7xLO?""20) em7 which corresponds to the Galactic value along the line of sight toward AI 31 (from Einline ancl W3nlI)."," We allowed the photoelectric absorption column density to vary freely, but constrained it to be higher than $N_{\rm H} = 6.7 \times10^{20}$ $^{-2}$ which corresponds to the Galactic value along the line of sight toward M 31 (from Einline and W3nH)." + This LAINB model has, This LMXB model has +At its high metallicity. |Fe/II]|—+0.36£0.07. standard isochrones predict it should be no hotter than Zar225450A if one assumes an old age.,"At its high metallicity, $=+0.36\pm0.07$, standard isochrones \citep{2008ApJS..178...89D} predict it should be no hotter than $T_{\rm{eff}} \approx 5450 K$ if one assumes an old age." + However. its measured temperature is Zur=5944dE68A. viekling the spectroscopic ft /jue2.9 Gyr.," However, its measured temperature is $T_{\rm{eff}} = 5944 \pm 68 K$ , yielding the spectroscopic fit $t_{\rm{Inferred}} \sim 2.9$ Gyr." + Next. consider the metal-rich ([Fe/I1I]—--0.37£0.05) SGD star MOA-2009-DLG-2598.. At ils trenmperature. Zur=4953+93. the expected surface eravily is loggy23.90.," Next, consider the metal-rich $=+0.37\pm0.05$ ) SGB star MOA-2009-BLG-259S. At its temperature, $T_{\rm{eff}} = 4953 \pm 93 K$, the expected surface gravity is $\log{g} \approx 3.90$." + However. (he measured value is logy=3.4040.24. for Mareenv2.0 Gyr.," However, the measured value is $\log{g}=3.40 \pm 0.24$, for $t_{\rm{Inferred}} \sim 3.0$ Gyr." + Therelore. the paradigm-challenging data reported by Bensbyοἱal.(2010.2011). pass the consistency test of being independently demonstrated in two distinct phases of stellar evolution.," Therefore, the paradigm-challenging data reported by \citet{2010A&A...512A..41B,2011A&A...533A.134B} pass the consistency test of being independently demonstrated in two distinct phases of stellar evolution." + In summary: In this we suggest that the origin of this discrepancy is with the isochrones that are used to interpret the data ancl estimate the ages., In summary: In this we suggest that the origin of this discrepancy is with the isochrones that are used to interpret the data and estimate the ages. + Specifically. we question whether the adopted. assumption of scaled-solar helium: abundance is a valid approximation [or the chemical evolution of the Galactic bulge at the metal-rich end.," Specifically, we question whether the adopted assumption of scaled-solar helium abundance is a valid approximation for the chemical evolution of the Galactic bulge at the metal-rich end." + We demonstrate that the resulting impact on stellar parameter determination leads to a bias in inferred ages: photonmetricdeterminations are too old and spectroscopic determinations are (oo voung., We demonstrate that the resulting impact on stellar parameter determination leads to a bias in inferred ages: photometricdeterminations are too old and spectroscopic determinations are too young. +The kinetic equations describing the rate of the magnetized induced scattering are derived in Appendix (see eqs.[A7]).,The kinetic equations describing the rate of the magnetized induced scattering are derived in Appendix (see eqs.[A7]). + To proceed further we introduce several simplifications., To proceed further we introduce several simplifications. + First of all. as we concentrate on the (transverse scattering. only the last terms in the kinetic equations (AT) should be retained.," First of all, as we concentrate on the transverse scattering, only the last terms in the kinetic equations (A7) should be retained." +" Then these equations can be presented as where i,j stand for the polarization states of the incident ancl scattered photons. and One can see that the kinetic equations differ [rom each other only by the factor g. which is generally of order unity. ("," Then these equations can be presented as where $i,j$ stand for the polarization states of the incident and scattered photons, and One can see that the kinetic equations differ from each other only by the factor $g^{ij}$, which is generally of order unity. (" +Note the svmmetry of ϱ with respect to the initial and. final photon states).,Note the symmetry of $g^{ij}$ with respect to the initial and final photon states). + Further. a detailed form of the particle distribution function does not play a crucial role.," Further, a detailed form of the particle distribution function does not play a crucial role." + Therefore we consider a monoenergetie distribution with some characteristic Lorentz-Iactor of (he particles., Therefore we consider a monoenergetic distribution with some characteristic Lorentz-factor of the particles. +" It is convenient (o replace the photon occupation numberswith the intensiües. /,=ο and ij,=2hvini(Ky)/c. and making use of their delta-functional angular distributions to integrate the kinetic equation over (he solid angle."," It is convenient to replace the photon occupation numberswith the intensities, $i_\nu=2h\nu^3n({\bf k})/c^2$ and $i_{\nu_1}=2h\nu_1^3n_1({\bf k_1})/c^2$, and making use of their delta-functional angular distributions to integrate the kinetic equation over the solid angle." +" Then we come to the following system of equations for the spectral intensities of the beams νι=[f /,,,04: where"," Then we come to the following system of equations for the spectral intensities of the beams $I_{\nu,\nu_1}\equiv\int i_{\nu,\nu_1}{\rm d}\Omega_1$ : where" +Iu Fi,In Fig. +e.10. we show the liue profiles resulting from a relmnant expanding into an ambicut medium with 3 sets of parameters.,\ref{fig:lp1} we show the line profiles resulting from a remnant expanding into an ambient medium with 3 sets of parameters. + The normalization of the line xofiles las been set so that the peal emission is approxinately 1.0. and we have preserved the relative scaling between the models(6.9... the central intensity of the à=LO%cm7 e=Ohms bot=20wr profile is approximately [« ereater than the central iuteusitv of the v7=LO’ci e—(0st f£=20vr profile).," The normalization of the line profiles has been set so that the peak emission is approximately 1.0, and we have preserved the relative scaling between the models, the central intensity of the $n = 10^{6} \pcm3$ , $v = 0 \kmps$, $t = 20\yr$ profile is approximately $4\times$ greater than the central intensity of the $n = 10^{5} \pcm3$, $v = 0 \kmps$ , $t = 20\yr$ profile)." + The bottom row in Fie., The bottom row in Fig. + shows the line profile which results if we iuteerate over the age of the reumant., \ref{fig:lp1} shows the line profile which results if we integrate over the age of the remnant. + Iu effect we stun each of the profiles iu the rows above with a Lapxopriate weight which reflects the time between cach “snapshot”., In effect we sum each of the profiles in the rows above with an appropriate weight which reflects the time between each “snapshot”. + Tn general. the nornalization first rises and then falls with time as cool clouds are created aud then destroved.," In general, the normalization first rises and then falls with time as cool clouds are created and then destroyed." + T1o Wieth of the line decreases with time as the expansio- specd of the xenuiuut sows., The width of the line decreases with time as the expansion speed of the remnant slows. + Tl« Ji1ο is clearly flat-topped. Wwch is characteristic of emisS10u from a ecoucetrically aud opically thin spwerical shell (sec. Fie.," The line is clearly flat-topped, which is characteristic of emission from a geometrically and optically thin spherical shell (see, Fig." + EMi 1l Capriotti 1950 with Cyuinfnuax~*~ 1.0)., 3 in Capriotti \cite{CFB1980} with $v_{\rm min}/v_{\rm max} \approx 1.0$ ). +" As th( effective. thiczuess of he ""sher in Fie.", As the effective thickness of the “shell” in Fig. + 2 Is mnininual. a taugeitial ine of seht does uo intercept au mereased 111iiber of ckmds. and “horus” are not seen in the profile.," \ref{fig:nw6_rho} is minimal, a tangential line of sight does not intercept an increased number of clouds, and “horns” are not seen in the profile." + Tn coutrast. he line profiles «is]uaved in tje richtinost columm of Fig.," In contrast, the line profiles displayed in the rightmost column of Fig." +lo 10 show au oeicreasinelv rounded or trianenlar pi'ofile as the remlalit ages.OOS aud the tine-averaged profile displavs a disnnlv ronudec top.," \ref{fig:lp1} show an increasingly rounded or triangular profile as the remnant ages, and the time-averaged profile displays a distintly rounded top." + Iu this model the rCle iod expanding iuto an ACN wind aud the line profies (which are seusitive te| the spatial distribution of cool gas) reflect the iucreasiie distortion of the remnant as it expacds., In this model the remnant is expanding into an AGN wind and the line profiles (which are sensitive to the spatial distribution of cool gas) reflect the increasing distortion of the remnant as it expands. + It is clear from Fie., It is clear from Fig. + 2. that the enussion is from a relatively smuall iiiber of clouds {particularly at later times.e.g... t= 20vr).," \ref{fig:nw6_rho} that the emission is from a relatively small number of clouds (particularly at later times, $t = 20 \yr$ )." + This intrκαπου a great deal of small scale styοure into the luie profiles which we have s:1000thed ou in Fie., This introduces a great deal of small scale structure into the line profiles which we have smoothed out in Fig. + 10 bv iweraging over nui different lines of sigoit;, \ref{fig:lp1} by averaging over many different lines of sight. + Observed line profiles from ACN are in reality very sinooth. aud it has been concluded that the umber of cluitting clouds nust be ~10° (Aray 1997)).. altlx1eh this would be reduced dif there Was sienifieant iudCLCturbulence.," Observed line profiles from AGN are in reality very smooth, and it has been concluded that the number of emitting clouds must be $\sim 10^{5}$ (Arav \cite{A1997}) ), although this would be reduced if there was significant microturbulence." + Aternativelv. electron scattering could hep ο explain siioli broad line profiles. especially in t1e line wines (Enuuieriueota .1992)).," Alternatively, electron scattering could help to explain smooth broad line profiles, especially in the line wings (Emmering \cite{EBS1992}) )." + In oux 2D lhivdrodviuiuaica1 nodels the clotds are in fact rings. alc nunerical viseIS]vy and the fiuitc! nuiniber of exid cells limit the nuuer of«istinet clouds which oru.," In our 2D hydrodynamical models the clouds are in fact rings, and numerical viscosity and the finite number of grid cells limit the number of distinct clouds which form." + Increased ποπΊσα. resolutkn aud 3D siniatious will produce many more istinct clouds. but preseut linitations mean tha the line xofiles from our models are not as s10ootlt as seen dm observations. and their fine structure should be ienored.," Increased numerical resolution and 3D simulations will produce many more distinct clouds, but present limitations mean that the line profiles from our models are not as smooth as seen in observations, and their fine structure should be ignored." + Iu Fig., In Fig. +" 11 we show the line profiles resulting from a remnant expaucdiig into a1 AGN wind of speed 3000. 5000. and. 70ioas 1, for a viewing angle. 135""."," \ref{fig:lp2} we show the line profiles resulting from a remnant expanding into an AGN wind of speed $v = 3000, 5000$ , and $7000 \kmps$ , for a viewing angle, $\theta = 135^{\circ}$ ." + Here 0 is clefec asx the anele between the AGN wind vector and tjo vector from the observer to the relunaut 0=y. COYTOSDouds to the «ybserver facing the side of the reLilaή ονudis iuto the oncoming AGN wind)., Here $\theta$ is defined as the angle between the AGN wind vector and the vector from the observer to the remnant $\theta = 0^{\circ}$ corresponds to the observer facing the side of the remnant expanding into the oncoming AGN wind). + Inspeclon of Figo., Inspection of Fig. +" 5 levicals that coolclouds exist over a wide 2:mec of aneles. naely ron0=0"" to 0zz 135°. For a liie of siely with 0—135 where the flow has a velocity component towards theobserver}."," \ref{fig:nw6_w3e8} reveals that coolclouds exist over a wide range of angles, namely from$\theta = +0^{\circ}$ to $\theta \approx 135^{\circ}$ For a line of sight with $\theta = 135^{\circ}$ where the flow has a velocity component towards theobserver)," +either.,either. +" However, in the simulation with comparatively weak feedback, some satellites have also baryon fractions above the cosmic mean value."," However, in the simulation with comparatively weak feedback, some satellites have also baryon fractions above the cosmic mean value." + These are satellites which lost a lot of dark matter through tidal stripping whereas they could hold on to most of their stars., These are satellites which lost a lot of dark matter through tidal stripping whereas they could hold on to most of their stars. + Both the and models are leading to considerably reduced baryon fractions in low mass satellites., Both the and models are leading to considerably reduced baryon fractions in low mass satellites. +" In the former case, this is readily expected as a signature of the winds."," In the former case, this is readily expected as a signature of the winds." +" In the latter, it is because more baryons stay in a diffuse gaseous phase, allowing them to be more easily ram-pressure stripped."," In the latter, it is because more baryons stay in a diffuse gaseous phase, allowing them to be more easily ram-pressure stripped." + An analysis of the evolution of the baryon fraction between z=6 and z=0 is given in Figure 12.., An analysis of the evolution of the baryon fraction between $z=6$ and $z=0$ is given in Figure \ref{fig:BarFrac_evo}. +" We here only show results for the simulation, as the qualitative behavior of the other simulations is similar."," We here only show results for the simulation, as the qualitative behavior of the other simulations is similar." +" We use two symbols for each satellite, one showing the data point at z=6 (red stars), while the corresponding values at z=0 are given by green stars."," We use two symbols for each satellite, one showing the data point at $z=6$ (red stars), while the corresponding values at $z=0$ are given by green stars." + Every pair of points belonging to the same satellite is connected by a dotted line., Every pair of points belonging to the same satellite is connected by a dotted line. +" Most satellites with circular velocities below ~20kms~! at z=6 lower their baryon fraction substantially until the present epoch, and they also do not tend to grow much."," Most satellites with circular velocities below $\sim 20\,{\rm km\, + s^{-1}}$ at $z=6$ lower their baryon fraction substantially until the present epoch, and they also do not tend to grow much." +" In contrast, most larger satellites tend to keep their baryon fraction or increase it slightly, often accompanied by a significant increase in Umax."," In contrast, most larger satellites tend to keep their baryon fraction or increase it slightly, often accompanied by a significant increase in $v_{\max}$." +" In the most massive satellites, part of this increase stems from modifications of the inner rotation curve due to the formation of a quite concentrated stellar component, ie. these satellites are not really bona-fide dark matter dominated systems as often assumed."," In the most massive satellites, part of this increase stems from modifications of the inner rotation curve due to the formation of a quite concentrated stellar component, i.e. these satellites are not really bona-fide dark matter dominated systems as often assumed." + We note however that the threshold at ~20kms! is not sharp; there are still many examples of satellites with an initially high vmax that end up as low mass satellites with a stripped baryonic component.," We note however that the threshold at $\sim 20\,{\rm km\, s^{-1}}$ is not sharp; there are still many examples of satellites with an initially high $v_{\max}$ that end up as low mass satellites with a stripped baryonic component." +" The last quantity we analyze in this section are the cumulative star formation histories of our satellites, as shown in Figure 13.."," The last quantity we analyze in this section are the cumulative star formation histories of our satellites, as shown in Figure \ref{fig:Satellites_SFRHistory}." +" The solid black line shows the total cumulative star formation history of all satellites in the final virial radius, normalized by their total final stellar mass."," The solid black line shows the total cumulative star formation history of all satellites in the final virial radius, normalized by their total final stellar mass." + The gray shaded area gives the lo scatter around this mean for the ensemble of all satellite star formation histories., The gray shaded area gives the $1\sigma$ scatter around this mean for the ensemble of all satellite star formation histories. +" The vertical dotted, dashed and dot dashed lines mark the times when 1096, 5096 and 9096 of the stars present at z=0 were formed."," The vertical dotted, dashed and dot dashed lines mark the times when $10\,\%$, $50\,\%$ and $90\,\%$ of the stars present at $z=0$ were formed." +" Finally, the dashed blue line repeats the result of the simulation in all the panels corresponding to the other simulations, in order to ease a comparison"," Finally, the dashed blue line repeats the result of the simulation in all the panels corresponding to the other simulations, in order to ease a comparison" +Globular clusters are the oldest. bound. stellar svstenis in our Galaxy.,Globular clusters are the oldest bound stellar systems in our Galaxy. + Their study provides. therefore: valuable information about early Galactic evolution., Their study provides therefore valuable information about early Galactic evolution. + In this respect. a major problem is that we do not know whether what we presently observe ijs still. representative of the initial conditions and. thus. a fossil imprint of the formation process. or whether the initial conditions have been wiped out by a 13GGvr long evolution within the tidal fields of what is now the Milkv Way.," In this respect, a major problem is that we do not know whether what we presently observe is still representative of the initial conditions and, thus, a fossil imprint of the formation process, or whether the initial conditions have been wiped out by a Gyr long evolution within the tidal fields of what is now the Milky Way." + Modelling the dynamical evolution of the Galactic elobular cluster system. is thus of great interest as it helps us to eo back in time to the earliest stages of the cluster svstem and to disentangle the formation and evolutionary fingerprints (sec. eg... Okazaki osa 1995. Baumgardt 1998. Vesperini 1998. Fall Zhang 2001).," Modelling the dynamical evolution of the Galactic globular cluster system is thus of great interest as it helps us to go back in time to the earliest stages of the cluster system and to disentangle the formation and evolutionary fingerprints (see, e.g., Okazaki Tosa 1995, Baumgardt 1998, Vesperini 1998, Fall Zhang 2001)." + In spite of numerous cllorts however. the shape of the initial distribution in mass of the halo globular clusters has remained ill-determined so far.," In spite of numerous efforts however, the shape of the initial distribution in mass of the halo globular clusters has remained ill-determined so far." + In our Galaxy. the mass function of the halo clusters," In our Galaxy, the mass function of the halo clusters" +the solution of the system of equations (13)) is given by: It is AaX;straightforward to see that eqs (16)) and (12)) produce the same results if the infall is negligible (1e. if A- 0).,the solution of the system of equations \ref{eq:newsystem}) ) is given by: It is straightforward to see that eqs \ref{eq:complsol}) ) and \ref{eq:diffwsol}) ) produce the same results if the infall is negligible (i.e. if $\Lambda \sim 0$ ). + Considering again a model with A=1 and Ξξ3 show the trend of Z/yvz in the case a=5 for different values of k., Considering again a model with $\Lambda = 1$ and $\lambda = 3$we can show the trend of $Z/y_Z$ in the case $\alpha = 5$ for different values of $k$. + We notice first that. unlike µ. 7] does not monotonically decrease. but its variation depends on the values of2 and A.," We notice first that, unlike $\mu$, $\eta$ does not monotonically decrease, but its variation depends on the values of $\lambda$ and $\Lambda$." +" It the present example. since >A. M, decreases with time anc Mca; increases with time. therefore 7 ranges between K at the beginning and 0 during the late stages of the evolution of the galaxy."," In the present example, since $\lambda > \Lambda$, $M_g$ decreases with time and $M_{ICM}$ increases with time, therefore $\eta$ ranges between $k$ at the beginning and 0 during the late stages of the evolution of the galaxy." + We show in Fig., We show in Fig. + 5. the evolution of Z/vz as a functioi of g/k for k=O.1 (solid line) and &=0.5 (dotted line)., \ref{fig:compl} the evolution of $Z/y_Z$ as a function of $\eta/k$ for $k=0.1$ (solid line) and $k=0.5$ (dotted line). + The model with &=0.5 predicts larger abundances because the metals carried out by the galactic wind are diluted with less ICM. therefore Ζ is larger.," The model with $k=0.5$ predicts larger abundances because the metals carried out by the galactic wind are diluted with less ICM, therefore $Z_A$ is larger." + These results cannot be compared to the results of eq. (12)), These results cannot be compared to the results of eq. \ref{eq:diffwsol}) ) + in a straightforward manner since they are expressed in different units., in a straightforward manner since they are expressed in different units. + However. we show in this plot the asymptotic trend of the simple differential wind model (short-dashed line in Fig. 3))," However, we show in this plot the asymptotic trend of the simple differential wind model (short-dashed line in Fig. \ref{diff31}) )" + for j£ approaching 0 (dashed line)., for $\mu$ approaching 0 (dashed line). + This comparison allows us to notice that. as expected. eq. (16))," This comparison allows us to notice that, as expected, eq. \ref{eq:complsol}) )" +" predicts larger metallicities compared to models with Z,=0.", predicts larger metallicities compared to models with $Z_A = 0$. + However. as long as &<|. the differences are small (of the order of ~ 0.1 - 0.2 dex).," However, as long as $k < 1$, the differences are small (of the order of $\sim$ 0.1 - 0.2 dex)." +" We can also calculate the ICM metallicity. which turns out to be: In the case we are considering more closely. οἱΞ3 and A=l|. therefore the ICM mass increases with time and X,, ranges between | and (1+κ)."," We can also calculate the ICM metallicity, which turns out to be: In the case we are considering more closely, $\lambda = 3$ and $\Lambda = 1$, therefore the ICM mass increases with time and $X_a$ ranges between 1 and $(1 + k)$." + The larger value Z4/Z can attain is therefore: which is smaller than 1 provided that: Once the ratio Z4/Z between ICM and ISM metallicity is observationally known. eq. (18))," The larger value $Z_A/Z$ can attain is therefore: which is smaller than 1 provided that: Once the ratio $Z_A/Z$ between ICM and ISM metallicity is observationally known, eq. \ref{eq:zratio}) )" + can be used to constrain a., can be used to constrain $\alpha$. + A special case of variable infall metallicity is represented by the situation in which Z4=ZaZ. namely the metallicity of the infalling gas Is set to be always equal to the one of the galactic wind.," A special case of variable infall metallicity is represented by the situation in which $Z_A = Z^o = \alpha Z$, namely the metallicity of the infalling gas is set to be always equal to the one of the galactic wind." + This condition implies therefore that the very same gas that has been driven out of the galaxy by energetic events can subsequently rain back to the galaxy. due to the pull of its gravitational potential.," This condition implies therefore that the very same gas that has been driven out of the galaxy by energetic events can subsequently rain back to the galaxy, due to the pull of its gravitational potential." + This kind of duty cycle is well known in astrophysics and it has been named (Shapiro Field 1976:: Bregman 1980))., This kind of duty cycle is well known in astrophysics and it has been named (Shapiro Field \cite{sf76}; ; Bregman \cite{breg80}) ). + In order to solve the chemical evolution in this special case. we have necessarily to assume that vt>A since the reservoir for the infall gas is given by the gas expelled out of the galaxy through galactic," In order to solve the chemical evolution in this special case, we have necessarily to assume that $\lambda \geq \Lambda$ since the reservoir for the infall gas is given by the gas expelled out of the galaxy through galactic" +where is the radius of the sun and R* is the radius of the earth's orbit Raround sun.,where $R$ is the radius of the sun and $R^{\star}$ is the radius of the earth's orbit around sun. +" Therefore, in view of the given HD line element (1), the metric tensors involved in the above equation (33) reduce to By substituting the expressions of equation (34) in equation (33) for the assumption R*>>p, we get Thus, in the Sun-Earth system we observe that for the usual 4-dimensional case (D= 2), gravitational redshift becomes z~242x1079=5 (say)."," Therefore, in view of the given HD line element (1), the metric tensors involved in the above equation (33) reduce to By substituting the expressions of equation (34) in equation (33) for the assumption $R^{\star}>>\mu$, we get Thus, in the Sun-Earth system we observe that for the usual 4-dimensional case $D=2$ ), gravitational redshift becomes $z \sim +2.12 \times 10^{-6}=z_2$ (say)." +" Therefore, for D>2, z«2 which indicates that as dimension increases the redshift gradually decreases (see Fig."," Therefore, for $D>2$, $zE_{\mu})$, must be compared with the square root of noise, $\sqrt{N}$, given by where $I_{\mu}$ is the background muon intensity due to cosmic ray induced atmospheric showers." + The value of ον necessary to consider an excess signal as a positive detection is above 4., The value of $S/\sqrt{N}$ necessary to consider an excess signal as a positive detection is above 4. + The sensitivity of TUPI muons for gamma bursts is shown in Fig.10. for several spectral indices im the power law energv spectrum and lor a duration of the the burst of AT(=100s).," The sensitivity of TUPI muons for gamma bursts is shown in Fig.10, for several spectral indices in the power law energy spectrum and for a duration of the the burst of $\Delta T (=100\;s)$." + The observation of the bursts at ground level is strongly limited by the spectral index value. as well as by the highest photon energies of the spectrum.," The observation of the bursts at ground level is strongly limited by the spectral index value, as well as by the highest photon energies of the spectrum." + Bursts of very long duration. very high photon maximum energy and with energy spectra not quite so steep can be observed at ground level bv using telescopes wilh a small muon enerev threshold (~0.3 GeV).," Bursts of very long duration, very high photon maximum energy and with energy spectra not quite so steep can be observed at ground level by using telescopes with a small muon energy threshold $\sim 0.3 GeV$ )." + The analvsis on the basis of FLUINA's results shows a real possibility of observing under certain conditions a GRB by the TUPI telescope., The analysis on the basis of FLUKA's results shows a real possibility of observing under certain conditions a GRB by the TUPI telescope. + A more accurate analysis. given the complexity of the processes. requires a full Monte Carlo study including the detector response and geomagnetic effects on the charged muons.," A more accurate analysis, given the complexity of the processes, requires a full Monte Carlo study including the detector response and geomagnetic effects on the charged muons." + We have reported a description and an analysis of (wo GLEs observed on December 2003 during a search for enhancements from the galactic center with mous al sea level and detected by using the TUPI telescope. after an upgrade of the data acquisition svstem.," We have reported a description and an analysis of two GLEs observed on December 2003 during a search for enhancements from the galactic center with muons at sea level and detected by using the TUPI telescope, after an upgrade of the data acquisition system." + The main conclusions are sunmmniarized as follows: (a)Lhe TUPI telescope can detect niuons at sea level with energies greater than the 0.3 GeV required to penetrate the two flagstones or walls surrounding the telescope., The main conclusions are summarized as follows: (a)The TUPI telescope can detect muons at sea level with energies greater than the $0.3$ GeV required to penetrate the two flagstones or walls surrounding the telescope. + The concrete reduces (he noise due to other non-muon particles. for example. il is opaque to electrons.," The concrete reduces the noise due to other non-muon particles, for example, it is opaque to electrons." + The telescope is sensitive to the conventional atmospheric muon flux., The telescope is sensitive to the conventional atmospheric muon flux. + The muon fhix obtained during a raster scan of 12 hours presents the well known West-East effect., The muon flux obtained during a raster scan of 12 hours presents the well known West-East effect. + In addition the geomagnetic effects distort the zenith angle distribution of sub-GeV to GeV ΠΙΟ] (b)Lhe (wo GLEs analyzed here have been mnambiguously detected. because in both cases (he GLE signals survives even when a large amplitude cliscrimination is used. while," In addition the geomagnetic effects distort the zenith angle distribution of sub-GeV to GeV muons (b)The two GLEs analyzed here have been unambiguously detected, because in both cases the GLE signals survives even when a large amplitude discrimination is used, while" +The observation of iillisecond microwave spikes frou the Sun has been iuterpreted as evrosvuchrotron laser enission since the late 1970's (11:: 5)).,The observation of millisecond microwave spikes from the Sun has been interpreted as gyrosynchrotron maser emission since the late 1970`s \cite{Wu}; ; \cite{Holman}) ). + The spikes exhibit a duration of a few milliseconds to some tens of milliseconds. and a uarrow bandwidth (=14 or a few MITIZ).," The spikes exhibit a duration of a few milliseconds to some tens of milliseconds, and a narrow bandwidth $\approx 1\%$ or a few MHz)." + The characteristic high brightness temperature (lot?015 A) deduced for the spikes strongly suggests maser action (3))., The characteristic high brightness temperature $10^{15-18}\ K$ ) deduced for the spikes strongly suggests maser action \cite{BenzBook}) ). + Several papers lave explored the possibilities of evelotron maser., Several papers have explored the possibilities of cyclotron maser. + Melrose aud Dulk (8)) lave used the seni-relativistic approximation to derive approximate ornulae for the frequency of the larecst erowth rate. and for the erowtl rate.," Melrose and Dulk \cite{D+M}) ) have used the semi-relativistic approximation to derive approximate formulae for the frequency of the largest growth rate, and for the growth rate." + Winelee (10)) approximated he effects of the ambicut plasina temperature on the naser enisson., Winglee \cite{Winglee}) ) approximated the effects of the ambient plasma temperature on the maser emission. + Asclavancden (1.. 2)) followed the diffusion of electrons iuto the loss-cone as a result of the maser cnussion. aud therefore the selfcousistent closure of the s-cone as the energy is transferred to the maser.," Aschwanden \cite{Aschwanden a}, \cite{Aschwanden b}) ) followed the diffusion of electrons into the loss-cone as a result of the maser emission, and therefore the self-consistent closure of the loss-cone as the energy is transferred to the maser." +" Iuncic id Robinson (6)) performed rav-traciug in a mode loop with dipole Ποια, aud concluded that maser emissiou cau escape from lower levels."," Kuncic and Robinson \cite{Kuncic}) ) performed ray-tracing in a model loop with dipole field, and concluded that maser emission can escape from lower levels." + Some papers have also attempted to derive analytical approximate foriuulae for the frequencies of the maser., Some papers have also attempted to derive analytical approximate formulae for the frequencies of the maser. + Dulk aud Melrose (8)) concluded that the maser enission is af. or near. 90 deerces to the magnetic field aud the miaxinuun erowth rate they find is very near the exclotrou frequency. or at the second harmonic.," Dulk and Melrose \cite{D+M}) ) concluded that the maser emission is at, or near, $90$ degrees to the magnetic field and the maximum growth rate they find is very near the cyclotron frequency, or at the second harmonic." + Hewitt and Melrose (13) derive the couditions for solutions aud the umber of solutions. but not the frequencies themselves.," Hewitt and Melrose \cite{Hewitt}) ) derive the conditions for solutions and the number of solutions, but not the frequencies themselves." + All previous studies concentrated on therate of the amplified wave. we use thecocfficient approach of Ramaty (9)) and calculate the frequencies at which the absorption coefficient 1s negative.," All previous studies concentrated on the of the amplified wave, we use the approach of Ramaty \cite{R}) ) and calculate the frequencies at which the absorption coefficient is negative." + Iu studies of microwave clussion from solar flares it is customary to assume that the plasina is composed a nonu-thermal population of energetic clecrons. aud a thermal population of much ligher density. but auch lower energv.," In studies of microwave emission from solar flares it is customary to assume that the plasma is composed a non-thermal population of energetic elecrons, and a thermal population of much higher density, but much lower energy." + The thermal plasina is assumed to determines the propagation of waves., The thermal plasma is assumed to determines the propagation of waves. + This plasima can be described by the cold plasina approximation usine magnuetoionic theory , This plasma can be described by the cold plasma approximation using magnetoionic theory \cite{MelroseB}) ). +ECM studies usually compute the erowth rate. while we follow Ramaty (9)) iu using the absorption cocticicut. but for the purposes of this article the detailed equations are uninportaut.," ECM studies usually compute the growth rate, while we follow Ramaty \cite{R}) ) in using the absorption coefficient, but for the purposes of this article the detailed equations are unimportant." + Both the erowth rate equation aud the absorption coefficient equation iuclude a d-function from which the important resonance condition is derived where wp=οληο 1 the evelotrou. frequency. w is the enission frequency. s is an integer. Ó is the anele of the wave vector k to the magnetic field. o is the augle vetween the velocity of the electron aud the magnetic field he pitch angle). aud 5 : ado are the usual Loreutz factor ud velocity.," Both the growth rate equation and the absorption coefficient equation include a $\delta$ -function from which the important resonance condition is derived where $\omega_B=eB/m_e/c$ is the cyclotron frequency, $\omega$ is the emission frequency, $s$ is an integer, $\theta$ is the angle of the wave vector ${\bf k}$ to the magnetic field, $\phi$ is the angle between the velocity of the electron and the magnetic field (the pitch angle), and $\gamma$ and $\beta$ are the usual Lorentz factor and velocity." + The important factor in equation Lois ης je refraction iudex for the Ordinary Mode (OM or 17). aud the eNtra-Ordinary mode (NO or 7-7).," The important factor in equation \ref{resonance} is $n_\pm$ the refraction index for the Ordinary Mode (OM or '+'), and the eXtra-Ordinary mode (XO or '-')." + The refraction oeidex Is a fuuction of he ambicut density. of the evclotrouyequency. of the emission. frequency. aud of the angle to 1e maenctic field," The refraction index is a function of the ambient density, of the cyclotronfrequency, of the emission frequency, and of the angle to the magnetic field" +We present contour maps and profiles of the SB. CI. e. and PA of the Sy galaxies in refcontprof..,"We present contour maps and profiles of the SB, CI, $\epsilon$, and PA of the Sy galaxies in \\ref{contprof}." + We discuss the Sy sample galaxies concerning features found or particularized by this study and cases where the influence of some features on the profiles is essential for the true structural parameters estimation via decomposition., We discuss the Sy sample galaxies concerning features found or particularized by this study and cases where the influence of some features on the profiles is essential for the true structural parameters estimation via decomposition. + When discussing the contour maps and profiles. reference to refcontprof goes without saying.," When discussing the contour maps and profiles, reference to \\ref{contprof} goes without saying." + Furthermore. CI and residual images/maps. contour maps. structure maps. and 2D model-subtracted residual images refMrk3351d--29)) visualize. the individual comments: dashed/dotted lines correspond to ellipticity maxima/minima.," Furthermore, CI and residual images/maps, contour maps, structure maps, and 2D model-subtracted residual images \\ref{Mrk335id}- \ref{38_Mrk541vi}) ) visualize the individual comments; dashed/dotted lines correspond to ellipticity maxima/minima." + In all figures north (N) ts up. east (E) to the left: the rest of the directions have been abbreviated as follows: south (S). west (W). northeast (NE). northwest (NW). southeast (SE). southwest (SW).," In all figures north (N) is up, east (E) to the left; the rest of the directions have been abbreviated as follows: south (S), west (W), northeast (NE), northwest (NW), southeast (SE), southwest (SW)." + In the CI images black is blue. white ts red.," In the CI images black is blue, white is red." + Concerning the residual images. white are the excess structures.," Concerning the residual images, white are the excess structures." + There is an extended feature at PA —4374 with /c peak SB 7/4 away from the nucleus," There is an extended feature at $PA\,$ $\,-43\degr$ with $I_{\rm \scriptstyle C}$ peak SB $7\farcs4$ away from the nucleus" +PACS: 04.80.-y.04.80.Ce Anderson&Nieto(2009) recently examinecl some still unexplained. anomalies connected with astrometric data in the solar svstem.,"PACS: 04.80.-y,04.80.Cc \citet{And09} recently examined some still unexplained anomalies connected with astrometric data in the solar system." + Thev are the fIvby anomaly (Anclerson 2008).. the Pioneer anomaly (Anclersonetal.1998).. the secular change of the Astronomical Unit (Ixrasinskv&Brumberg2004) and the increase in the eccentricity of the Moons orbit (Williams&Boges2008).," They are the flyby anomaly \citep{And08}, the Pioneer anomaly \citep{And98}, the secular change of the Astronomical Unit \citep{Kra04} and the increase in the eccentricity of the Moon's orbit \citep{Wil08}." +. In fact. there is the possibility that also a fifth anomaly does actually exist: (he anomalous perihelion precession of Saturn (Iorio2009a).," In fact, there is the possibility that also a fifth anomaly does actually exist: the anomalous perihelion precession of Saturn \citep{Ior09}." +. The corrections Ac to the standard Newtonian/Einsteinian secular precession of the longitude of the z of Saturn. estimated with the latest versions of the EPM," The corrections $\Delta\dot\varpi$ to the standard Newtonian/Einsteinian secular precession of the longitude of the $\varpi$ of Saturn, estimated with the latest versions of the EPM" +to test the very low rotational velocity preclieted by our interpretation of the pulsation spectrum in later sections. a high-clispersion spectrum. is needed.,"to test the very low rotational velocity predicted by our interpretation of the pulsation spectrum in later sections, a high-dispersion spectrum is needed." + 90700322. was observed. on. 2010. August 12 with the Llieh Resolution Spectrograph (9) on the Llobby-Eberky Telescope at AleDonald. Observatory., 9700322 was observed on 2010 August 12 with the High Resolution Spectrograph \citep{Tull1998} on the Hobby-Eberly Telescope at McDonald Observatory. + “Phe spectrum was taken at f?~30000 using the 316g cross-disperser setting. spanning a wavelength region from 4120τοῦ...," The spectrum was taken at $R \sim 30\,000$ using the 316g cross-disperser setting, spanning a wavelength region from $4120-7850$." + The exposure time was 1500 secs., The exposure time was 1800 secs. + A signal/noise ratio of 194 was found at 593.6 nm., A signal/noise ratio of 194 was found at 593.6 nm. + We reduced the data using standard. techniques with package (?).., We reduced the data using standard techniques with package \citep{vanDokkum2001}. + The effective temperature. Zir. and surface gravity. ogg. can be obtained by minimizing the dillerence between he observed. ancl synthetic spectra.," The effective temperature, $T_{\rm eff}$, and surface gravity, $\log g$, can be obtained by minimizing the difference between the observed and synthetic spectra." + We used a fit to the 113 line to obtain an estimate of the elfective temperature., We used a fit to the $\beta$ line to obtain an estimate of the effective temperature. + For stars. with Zip<7000 IX the Balmer lines are no onger sensitive to gravity. so we used the triplet at 5167.321. 5172.684. anc 5183.604A for this purpose.," For stars with $T_{\rm eff} < 7000$ K the Balmer lines are no longer sensitive to gravity, so we used the triplet at 5167.321, 5172.684, and 5183.604 for this purpose." +" The &oodness-o[-[it. parsumeter. A7. is defined \201=¥(baoday) where N is the total number of pointsand Zi, and Ji are the intensities of the observed. and. computed: profiles. respectively."," The goodness-of-fit parameter, $\chi^2$, is defined $\displaystyle \chi^2 = \frac{1}{N} \sum \bigg(\frac{I_{\rm obs} - I_{\rm th}}{\delta I_{\rm obs}}\bigg)^2$, where $N$ is the total number of pointsand $I_{\rm obs}$ and $I_{\rm th}$ are the intensities of the observed and computed profiles, respectively." +" 344, is the photon noise.", $\delta I_{\rm obs}$ is the photon noise. + The error in a parameter was estimated by the variation required to change v oby unity., The error in a parameter was estimated by the variation required to change $\chi^2$ by unity. + Phe projected. rotational velocity. and. the microturbulence were determined. by matching the metal lines in the range 5160 5200 A., The projected rotational velocity and the microturbulence were determined by matching the metal lines in the range 5160 – 5200 . + From this procedure we obtained Ty 66700 x: 100 Kk. logg——333.7 £ 0l. rsini—-—119 + 1 5. £ 2—22.0 0 0.5 5.," From this procedure we obtained $_{\rm eff}$ 6700 $\pm$ 100 K, $\log g$ 3.7 $\pm$ 0.1, $v \sin i$ 19 $\pm$ 1 $^{-1}$, $\xi$ 2.0 $\pm$ 0.5 $^{-1}$." + In Fig. 2..," In Fig. \ref{spectrum}," + we show the match to observed spectrum., we show the match to observed spectrum. + The theoretical profiles were computed with (?) using atmospheric models. (?).., The theoretical profiles were computed with \citep{Kurucz1981} using atmospheric models \citep{Kurucz1993b}. + The solar opacity. distribution function. was used in these calculations., The solar opacity distribution function was used in these calculations. + “Phe elfective temperature calculated. from the spectrum is somewhat lower than that obtained by matching the spectral energy. distribution discussed in the previous section., The effective temperature calculated from the spectrum is somewhat lower than that obtained by matching the spectral energy distribution discussed in the previous section. + The dillerence is within the statistical uncertainties., The difference is within the statistical uncertainties. + We note that the spectrum. showed no evidence for the presence of a conipanion., We note that the spectrum showed no evidence for the presence of a companion. + Because of problems of line blending. we decided. to use direct matching. of rotationallv-broadened: svnthetic spectra to the observations in order to determine the oxojected rotational velocity.," Because of problems of line blending, we decided to use direct matching of rotationally-broadened synthetic spectra to the observations in order to determine the projected rotational velocity." + For this purpose. we divided he spectrum into several LOOA segments.," For this purpose, we divided the spectrum into several 100 segments." + We derived. the abundances in each segment using X7 minimization., We derived the abundances in each segment using $\chi^2$ minimization. + We used he line lists and atomic parameters in 2? as updated by 2.., We used the line lists and atomic parameters in \citet{Kurucz1995} as updated by \citet{Castelli2004}. + Table 1 shows the abundances expressed. in the usual ogarithmic form relative to the total number of atoms Ni., Table \ref{abund} shows the abundances expressed in the usual logarithmic form relative to the total number of atoms $N_{\rm tot}$. + ‘To more easily compare the chemical abundance pattern in 99700322. Fig.," To more easily compare the chemical abundance pattern in 9700322, Fig." + 3 shows the stellar abuncdances relative o the solar values (?) asa function of atomic number., \ref{abundpat} shows the stellar abundances relative to the solar values \citep{Grevesse2010} as a function of atomic number. + The error in abundance for a particular clement which is shown in Table 1 is the standard error of the mean abundance computed from all the wavelength segments., The error in abundance for a particular element which is shown in Table \ref{abund} is the standard error of the mean abundance computed from all the wavelength segments. + This analysis shows that the chemical abundance in 99700322 is the normal solar abundance., This analysis shows that the chemical abundance in 9700322 is the normal solar abundance. + The effective temperature determined for 99700322 makes the star one of the cooler 0 SSct stars., The effective temperature determined for 9700322 makes the star one of the cooler $\delta$ Sct stars. + Other pulsators with similar temperatures are known. e.g.. IxIx. for p Pup (?) and Why for 44 Tau (7).. TheAepler," Other pulsators with similar temperatures are known, e.g., K for $\rho$ Pup \citep{netopil2008} and K for 44 Tau \citep{LenzPamyatnykhZdravkovBreger2010}." + data of 99700322 were analyzed with the statistical package (?).., The data of 9700322 were analyzed with the statistical package \citep{LenzBreger2005}. + This package carries out multifrequeney analyses with Fourier as well as least-squares algorithms ancl does not rely on the assumption of white noise., This package carries out multifrequency analyses with Fourier as well as least-squares algorithms and does not rely on the assumption of white noise. + Previous comparisons of multifrequencey analyses of satellite cata with other techniques such as (?) have shown that is more conservative in assigning statistical significances. leads to fewer (?).. and hopefully also fewer erroneous. pulsation frequencies. but may consequently alsomiss some valid frequencies.," Previous comparisons of multifrequency analyses of satellite data with other techniques such as \citep{Reegen2007} have shown that is more conservative in assigning statistical significances, leads to fewer \citep{Poretti2009}, , and hopefully also fewer erroneous, pulsation frequencies, but may consequently alsomiss some valid frequencies." + We did not concern ourselves with small instrumental zero-point changes in the data since we have no method, We did not concern ourselves with small instrumental zero-point changes in the data since we have no method +of kineniatical ambiguity the near and far distances are shown). multiplicity aud linear dimension.,"of kinematical ambiguity the near and far distances are shown), multiplicity and linear dimension." + The new infrared clusters. stellay groups aud candidates from the optical nebula survey amount to 123. and from the radio nebula survey 56.," The new infrared clusters, stellar groups and candidates from the optical nebula survey amount to 123, and from the radio nebula survey 56." +" The rates of detection relative to the input nebula catalogues are 12 aud for optical aud radio nebulae. respectively,"," The rates of detection relative to the input nebula catalogues are 12 and for optical and radio nebulae, respectively." + The higher detection rate among radio nebulae is surely because they represent more often overall complexes. while optical nebulae deal more often with structural details of closer complexes.," The higher detection rate among radio nebulae is surely because they represent more often overall complexes, while optical nebulae deal more often with structural details of closer complexes." + Object classes are infrared cluster (IRC). stellar eroup (IIC1). cluster candidate απού. aud open cluster (Πιο).," Object classes are infrared cluster (IRC), stellar group (IRGr), cluster candidate (IRCC), and open cluster (IROC)." + IRCs are iu general populous aud at least, IRCs are in general populous and at least +Iun Fig.,In Fig. + 5. the projected distribution of the detected cluster candidates is shown., \ref{fig:proj_clusters} the projected distribution of the detected cluster candidates is shown. + There is a clear paucitv of clusters in the region a2311 aud 6>10.27.," There is a clear paucity of clusters in the region $\alpha \gsim 341^\circ$ and $\delta \gsim +-40.2^\circ$." + This is probably due to variations iu the completeness of the sinele-frame catalogs at the adopted limiting magnitude of this analysis., This is probably due to variations in the completeness of the single-frame catalogs at the adopted limiting magnitude of this analysis. + In fact. besides the region of clear incompleteness. already removed. there is a significant area of patch A (~25%) which is incomplete at the maguitue adopted here (Paper D.," In fact, besides the region of clear incompleteness, already removed, there is a significant area of patch A $\sim 25\%$ ) which is incomplete at the magnitude adopted here (Paper I)." + Therefore. the definition of a more homogeneous region for the cluster analvsis would require a further trimming of the effective area of the analysis.," Therefore, the definition of a more homogeneous region for the cluster analysis would require a further trimming of the effective area of the analysis." + A ore detailed discussion of this point will be carried out by Scodegeio (1998)., A more detailed discussion of this point will be carried out by Scodeggio (1998). + For cach cluster. cutouts from the coadded inage are created ceutered at the nouinal position of the ideutified cluster covering a region of 7<7! area. which roughly corresponds to the FORS field of view.," For each cluster, cutouts from the coadded image are created centered at the nominal position of the identified cluster covering a region of $7' \times 7'$ area, which roughly corresponds to the FORS field of view." +" These cutouts are available at ""http://wwow.eso.org/eis/datarcl.tui”.", These cutouts are available at “http://www.eso.org/eis/datarel.html”. + Fie., Fig. +" 6 shows the fraction of cluster candidates found in one catalog having a counterpart in the other as function of significauce,", \ref{fig:sigma_dist} shows the fraction of cluster candidates found in one catalog having a counterpart in the other as function of significance. + As expected. highly significant detections are found iu both catalogs. within a search radius of P. arcmin.," As expected, highly significant detections are found in both catalogs, within a search radius of 1 arcmin." + Of the 15 lo detectious found in the even catalog 13 )) have a counterpart in the odd. while for the 15 detections inthe odd catalog 16 )) have a counterpar in the even.," Of the 15 $4\sigma$ detections found in the even catalog 13 ) have a counterpart in the odd, while for the 18 detections in the odd catalog 16 ) have a counterpart in the even." + Note that in this comparison the counterparts nav have a significance lower than to. but still ligher than 20 (because of the choice of the extraction threshold).," Note that in this comparison the counterparts may have a significance lower than $4\sigma$, but still higher than $2\sigma$ (because of the choice of the extraction threshold)." + I cau also be seen that for 26 detections the probability of iiving a counterpart inthe other catalog is stillreasonably ügh for detectious iu the even catalog aud for he odd., It can also be seen that for $3\sigma$ detections the probability of having a counterpart in the other catalog is still reasonably high – for detections in the even catalog and for the odd. +— The estimated redshifts for the detected clusters range vetween 2=0.5 and 2=1.3.," The estimated redshifts for the detected clusters range between $z = +0.3$ and $z = 1.3$." + In Fie., In Fig. + 7 re redshift distribution of the total candidate sample is shown atca compared to the distribution for the candidates reportea in the PDCS., \ref{fig:redshift_dist} the redshift distribution of the total candidate sample is shown and compared to the distribution for the candidates reported in the PDCS. + The shaded area represeuts the redshif distribution of the [o cididates., The shaded area represents the redshift distribution of the $4\sigma$ candidates. + The distribution of the ta candidates is seen to cover the redshift range from 0.3 to 0.9 with a median redshift of ;=0.1 while the total sample exteuds to +=1.2 with a Iiedian of +=0.6.," The distribution of the $4\sigma$ candidates is seen to cover the redshift range from 0.3 to 0.9 with a median redshift of $z = 0.4$, while the total sample extends to $z = 1.3$ with a median of $z = +0.6$." + For couparison. the median redshift of the PDCS is 2=0.1.," For comparison, the median redshift of the PDCS is $z = +0.4$." + The EIS aud PDCS redshift distributions are quite similar. but a small relative shift in redshift may be present.," The EIS and PDCS redshift distributions are quite similar, but a small relative shift in redshift may be present." + This effect might be either due to a sinall bias of the current implementation of the inatched filter algoritluu or to the fact that the EIS data are somewhat deeper than those of the PDCS., This effect might be either due to a small bias of the current implementation of the matched filter algorithm or to the fact that the EIS data are somewhat deeper than those of the PDCS. + Applving the passive evolution R-corrections in the creation of the Likehhood maps. im most cases. does not affect the detection of a caudidate.," Applying the passive evolution K-corrections in the creation of the Likelihood maps, in most cases, does not affect the detection of a candidate." + However. there are a few cases where the candidates detected with the uo- I&-correctious fail to be detected with the passive evolution I&-corrections.," However, there are a few cases where the candidates detected with the no-evolution K-corrections fail to be detected with the passive evolution K-corrections." + The distributions of cstimated cluster richness are shown in Fie. &.., The distributions of estimated cluster richness are shown in Fig. \ref{fig:richness_dist}. + Again the distributions for the total cluster sample is shown. aud the shaded area indicates the distribution for the lo candidates.," Again the distributions for the total cluster sample is shown, and the shaded area indicates the distribution for the $4\sigma$ candidates." + It is ποσα that the A; ricliness spans a wide range extending up to ~110 witha media of ~70., It is seen that the $\Lambda_{cl}$ richness spans a wide range extending up to $\sim 140$ with a median of $\sim 70$. + The Abell richness estimate. Wye. is found to vary between 9 and 102 with a median of 31.," The Abell richness estimate, $N_R$, is found to vary between 9 and 102 with a median of 34." + Note that iu the case of richness an appropriate comparison with the results of P96 cannot be made because of our imposed richness criterion iu the detection and differences between the estimates of the mean background counts in the calculation of the Alcll richuess in this paper aud PDCS., Note that in the case of richness an appropriate comparison with the results of P96 cannot be made because of our imposed richness criterion in the detection and differences between the estimates of the mean background counts in the calculation of the Abell richness in this paper and PDCS. + A comparison between the estimates of the candidates xoperties. discussed in the previous section. is used to obtain a rough estimate of their accuracy.," A comparison between the estimates of the candidates properties, discussed in the previous section, is used to obtain a rough estimate of their accuracy." + Fig., Fig. + 9 shows a comparison of the estimated redshifts for all paired detections. as determined in the odd/even catalogs.," \ref{fig:redshift_comp} shows a comparison of the estimated redshifts for all paired detections, as determined in the odd/even catalogs." + Some of the points represent more than one cluster candidate due to the discreteness of the redshift bius., Some of the points represent more than one cluster candidate due to the discreteness of the redshift bins. + The scatter around the diagonal is found to be ~ 0.06. consistent with the possible accuracy given by the adopted redshift exid.," The scatter around the diagonal is found to be $\sim$ 0.06, consistent with the possible accuracy given by the adopted redshift grid." + In Fig., In Fig. + 10 the richness estimates are compared in the same way and it is found that for X; the scatter is aud for Vy the scatter is., \ref{fig:richness_comp} the richness estimates are compared in the same way and it is found that for $\Lambda_{cl}$ the scatter is and for $N_R$ the scatter is. +". The recently released EIS Ebaud data for Patch A (a~22|n"" and ὃ~——407: see Paper D) have been used to search for clusters of ealaxies over an area of 2.5 square deerces. iu the redshift range 0.2τς1.3."," The recently released EIS I-band data for Patch A $\alpha \sim 22^h +45^m$ and $\delta \sim -40^\circ$; see Paper I) have been used to search for clusters of galaxies over an area of 2.5 square degrees, in the redshift range $0.2 \leq z \leq 1.3$." + The matched filter algorithii has been applied to the even and odd sinele-framie catalogs to assess the performance of the detection technique. to establish the detection threshold for robust detections and to evaluate the quality of the EIS data for this kind of analysis. one of the main goals of the survey.," The matched filter algorithm has been applied to the even and odd single-frame catalogs to assess the performance of the detection technique, to establish the detection threshold for robust detections and to evaluate the quality of the EIS data for this kind of analysis, one of the main goals of the survey." + The candidate cluster sample based on of Lo detections consists of 21 objects. viclding a surface deusitv of 8.1 candidates per square deeree. with a median redshift of :=0.1.," The candidate cluster sample based on of $4\sigma$ detections consists of 21 objects, yielding a surface density of 8.4 candidates per square degree, with a median redshift of $z=0.4$." + When all 36 detectious are cousidered 39 candidates are found. leading to a surface density of 16 per square degree and a mecian redshift of 2= 0.6.," When all $3\sigma$ detections are considered 39 candidates are found, leading to a surface density of 16 per square degree and a median redshift of $z=0.6$ ." + Cutouts for the cluster candidates are available at vlttp://www.eso.ore/eis/ddatarelμα, Cutouts for the cluster candidates are available at datarel.html”. +", These results should be considered preliminary as siguificautlv better data are available for the other EIS patches.", These results should be considered preliminary as significantly better data are available for the other EIS patches. + More importantly. the use of catalogs extracted from the coadded tages will allow a deeper cluster search to be carried out. thereby extending the redshift range for the cluster sample.," More importantly, the use of catalogs extracted from the coadded images will allow a deeper cluster search to be carried out, thereby extending the redshift range for the cluster sample." + Clearly. the EIS data more than fulfills the science requirements of the survey. as originally stated.," Clearly, the EIS data more than fulfills the science requirements of the survey, as originally stated." + Iu this first release of the EIS cluster catalog the effort lias been conceutrated ou the Lbaud data., In this first release of the EIS cluster catalog the effort has been concentrated on the I-band data. + However. a lnmited nuuber of frames iu. V-band have been obtained and will be used to further investigate the candidate clusters over the surveyed region (Olsen 1998).," However, a limited number of frames in V-band have been obtained and will be used to further investigate the candidate clusters over the surveyed region (Olsen 1998)." +"convention for spectral index. o. is such that the variation of flux with [requeney. S,x7mm","convention for spectral index, $\alpha$, is such that the variation of flux with frequency, $S_\nu \propto +\nu^{-\alpha}$." + We have made imaging observations of all quasars in Laing et ((1983) with 0.65<2«1.20 which are accessible from. UINIICE. (Le. 6< 607). except for 3€ 454.3. which only appears in Laing et πο to its strongly Doppler-boosted radio core.," We have made imaging observations of all quasars in Laing et (1983) with $0.65 \leq z < 1.20$ which are accessible from UKIRT (i.e., $\delta < 60^\circ$ ), except for 3C 454.3, which only appears in Laing et due to its strongly Doppler-boosted radio core." + We have also imaged the 1.228 quasar 3C 68.1., We have also imaged the $z=1.228$ quasar 3C 68.1. + With the exception of the images of 3C 68.1 and 3€ 175. which were observed on UT 1998 Sep 17. and 3C 380. which was observed on UT 1999 Mar 12. all the data in this paper were obtained during the nights of UTE 1999 Mar 68 with the HUCAM infrared array imager on UIT.," With the exception of the images of 3C 68.1 and 3C 175, which were observed on UT 1998 Sep 17, and 3C 380, which was observed on UT 1999 Mar 12, all the data in this paper were obtained during the nights of UT 1999 Mar 6–8 with the IRCAM3 infrared array imager on UKIRT." + The. images were cach constructed from a single five-point mosaic with an exposure of GOss ss for 3€ 68.1 and 3€ 175) per position (split into multiple coadcds to avoid saturation while still achieving backgrouncd-limited operation). while the L’ images were produced from a number of five-point mosaics with an exposure of 24ss (30ss for 3€ 68.1 and 3€ 175) per position (again split into multiple coacdcds).," The images were each constructed from a single five-point mosaic with an exposure of s s for 3C 68.1 and 3C 175) per position (split into multiple coadds to avoid saturation while still achieving background-limited operation), while the $L'$ images were produced from a number of five-point mosaics with an exposure of s s for 3C 68.1 and 3C 175) per position (again split into multiple coadds)." + The seeing. as measured. from stars and the quasars themselves at A- band. varied between 0.6 ancl 1.0 aresec.," The seeing, as measured from stars and the quasars themselves at $K$ -band, varied between 0.6 and 1.0 arcsec." + Each five-point mosaic was reduced. separately. in the manner of SRL. and the individual L' mosaics were then averaged together to form the final image.," Each five-point mosaic was reduced separately, in the manner of SRL, and the individual $L'$ mosaics were then averaged together to form the final image." + Many. of our thermal-infrared observations were beset by problems with the array control software. as reported. by SRL.," Many of our thermal-infrared observations were beset by problems with the array control software, as reported by SRL." + Care was taken to ensure that only those mosaics which increased the sensitivity of the observation were added to the final image., Care was taken to ensure that only those mosaics which increased the sensitivity of the observation were added to the final image. + Flax calibration solutions for the nights of Mar 65 were determined [rom observations of UINIIUE stanclard stars throughout the night., Flux calibration solutions for the nights of Mar 6–8 were determined from observations of UKIRT standard stars throughout the night. + The images of 3€ 68.1. 3C 175. and 3€ 380 were photometrically calibrated: using observations of standard stars taken at similar alrmasses to the quasars either immediately. before or after the target. observations.," The images of 3C 68.1, 3C 175, and 3C 380 were photometrically calibrated using observations of standard stars taken at similar airmasses to the quasars either immediately before or after the target observations." + Since the observations through dillerent. filters were taken within minutes of each other (except for 3C 336. whose and L' images were taken on consecutive nights). variability is not à concern.," Since the observations through different filters were taken within minutes of each other (except for 3C 336, whose and $L'$ images were taken on consecutive nights), variability is not a concern." + Photometry for cach of the quasars was measured in a 3-aresec aperture., Photometry for each of the quasars was measured in a 3-arcsec aperture. + For the images. the sky level was determined from an annulus around. the quasar in the usual manner. and any errors introcluced by an incorrect determination were confirmed to be negligible.," For the images, the sky level was determined from an annulus around the quasar in the usual manner, and any errors introduced by an incorrect determination were confirmed to be negligible." + At L'. imperfect matching of the sky level between frames caused bv. residual structure from problems with the control software required a more complex technique to avoid. large systematic errors.," At $L'$, imperfect matching of the sky level between frames caused by residual structure from problems with the control software required a more complex technique to avoid large systematic errors." + polvnomial of first. order in. column and row number was [it to the central region of the mosaic. excluding a 3-aresec aperture around the quasar. ancl this was used as a sky frame [for calculating the photometry.," A polynomial of first order in column and row number was fit to the central region of the mosaic, excluding a 3-arcsec aperture around the quasar, and this was used as a sky frame for calculating the photometry." + Lt was confirmed. that changing the surface fitting parameters varied the measured photometry by less than the photometric uncertainty., It was confirmed that changing the surface fitting parameters varied the measured photometry by less than the photometric uncertainty. + The results are listed in ‘Table 1.., The results are listed in Table \ref{tab:phot}. + We correct the photometry for Galactic extinction using the extinction law of Pei (1992). and then for contamination," We correct the photometry for Galactic extinction using the extinction law of Pei (1992), and then for contamination" +Qual.T) 2 7 AUT).,"(n,T) = n^2 (T)." + where A(7) is the so-called cooling Iunction., where $\Lambda(T)$ is the so-called cooling function. + Correspondingly.equation (2.4)) can be wrillen as AUT) =a.," Correspondingly,equation \ref{eq-entropy-A}) ) can be written as (T) =." + In general. for a plasma with a non-trivial (e.g.. solar) chemical composition. the fictional form of A(T) is quite complicated.," In general, for a plasma with a non-trivial (e.g., solar) chemical composition, the functional form of $\Lambda(T)$ is quite complicated." + In particular.at temperatures between LO? IK and 10* Ix cooling is dominated by various strong atomic lines (mostly. those of carbon. oxvgen. ancl iron) and looks like a mountainous landscape: in solar phivsies il is usuallyrepresented by a broken power-law (e.g..7)..," In particular,at temperatures between $10^5$ K and $10^7$ K cooling is dominated by various strong atomic lines (mostly those of carbon, oxygen, and iron) and looks like a mountainous landscape; in solar physics it is usuallyrepresented by a broken power-law \citep[e.g.,][]{Raymond_Smith-1977}." + Llowever. above about 10* IX atomic lines mostly disappear and the cooling functionis dominated by. a simple bremsstrahlung cooling (?):: — —hZONE," However, above about $10^7$ K atomic lines mostly disappear and the cooling functionis dominated by a simple bremsstrahlung cooling \citep{RL-1979}: (T) = = =." +"S where for a nonrelativistic Maxwellian hydrogen plasma. ος Ko42,2 1 σα, .cmids;—-1IIs:—1/2."," where for a nonrelativistic Maxwellian hydrogen plasma, = c^2 1.4 ." +11 inRusselletal.(2010).,1 in\citet{russell10}. +". An analogue to it is visible in reff:gamma((a), where some points belonging to the soft and intermediate states have Fy; below that of the hard state."," An analogue to it is visible in \\ref{f:gamma}( (a), where some points belonging to the soft and intermediate states have $F_{\rm bol}$ below that of the hard state." +" However, the tracks covered are much more chaotic than those in transients, where the hard-to-soft transitions occur at clearly higher flux than the reverse ones."," However, the tracks covered are much more chaotic than those in transients, where the hard-to-soft transitions occur at clearly higher flux than the reverse ones." +" Although we consider this result for Cyg X-1 secure, it relies on the absorption-corrected ASM observations, and should be tested by pointed X-ray observations."," Although we consider this result for Cyg X-1 secure, it relies on the absorption-corrected ASM observations, and should be tested by pointed X-ray observations." + Our findings concerning radio/X-ray correlations fall into two categories., Our findings concerning radio/X-ray correlations fall into two categories. + Those regarding the dependence of the correlation on the X-ray energy range and on the spectral state are basically independent of the uncertain strength of free-free absorption in the stellar wind from the donor of Cyg X-1., Those regarding the dependence of the correlation on the X-ray energy range and on the spectral state are basically independent of the uncertain strength of free-free absorption in the stellar wind from the donor of Cyg X-1. +" On the other hand, the values of the radio/X-ray correlation indices are, most likely, affected by the free-free absorption."," On the other hand, the values of the radio/X-ray correlation indices are, most likely, affected by the free-free absorption." +" One of our findings is the strong dependence of the radio/X-ray correlation index on theenergy band, see Table 1 and reff:index.."," One of our findings is the strong dependence of the radio/X-ray correlation index on theenergy band, see Table \ref{t:fit} and \\ref{f:index}." + The energy-dependence of the correlation can be understood as an intrinsic dependence of the radio flux on the bolometric flux combined with flux-dependent changes in the spectral shape (seen in our data at E«12 keV and E>150 keV)., The energy-dependence of the correlation can be understood as an intrinsic dependence of the radio flux on the bolometric flux combined with flux-dependent changes in the spectral shape (seen in our data at $E<12$ keV and $E>150$ keV). + The latter leads to the bolometric flux and that in a narrow X-ray band being non-proportional in a way that depends on the band chosen., The latter leads to the bolometric flux and that in a narrow X-ray band being non-proportional in a way that depends on the band chosen. +" This dependence has mostly been ignored in previous work on this subject except, e.g., Merlonietal.(2003),, where its effect has been accounted for by a theoretical model, Zdziarskietal. (2004),, whofitted the radio correlation with Fy in GX 339-4, obtaining Poo=0.79+0.07, and Corbeletal.(2003),, who calculated the correlation in GX 339-4 in four separate bands between 3 keV and 200 keV, but have not calculated the correlation index for F4."," This dependence has mostly been ignored in previous work on this subject except, e.g., \citet{mhd03}, where its effect has been accounted for by a theoretical model, \citet{z04}, , whofitted the radio correlation with $F_{\rm bol}$ in GX 339–4, obtaining $p_{\rm bol}=0.79\pm 0.07$, and \citet{corbel03}, who calculated the correlation in GX 339–4 in four separate bands between 3 keV and 200 keV, but have not calculated the correlation index for $F_{\rm bol}$." +" In models in which the accretion rate, M, governs both the accretion flow and the jet (see Section 4.2.3 below), it is crucial to determine its value."," In models in which the accretion rate, $\dot M$, governs both the accretion flow and the jet (see Section \ref{accretion} below), it is crucial to determine its value." +" It is related to the bolometric flux by, where d is the source distance, and e is the accretion radiative efficiency."," It is related to the bolometric flux by, where $d$ is the source distance, and $\epsilon$ is the accretion radiative efficiency." +" Our results indicate that the jet power is a fraction not of the total accretion power but only of that responsible for emission associated with hot electrons in the accretion flow, Fi."," Our results indicate that the jet power is a fraction not of the total accretion power but only of that responsible for emission associated with hot electrons in the accretion flow, $F_{\rm hot}$." +" In this case we should write Mο:Fjj/ey, where ey is the efficiency of producing the hot flow."," In this case we should write $\dot M \propto F_{\rm hot}/\epsilon_{\rm hot}$, where $\epsilon_{\rm hot}$ is the efficiency of producing the hot flow." +" However, the difference between the two is negligible in the hard state, where the hot flow is energetically dominant."," However, the difference between the two is negligible in the hard state, where the hot flow is energetically dominant." +" The common procedure has been to use a narrow band, such as 3-9 keV (e.g., Coriatetal. 2011)), as the proxy for Fig, i.e. implicitly assuming FxοςFy, and then proceed to theoretical interpretation."," The common procedure has been to use a narrow band, such as 3–9 keV (e.g., \citealt{coriat11}) ), as the proxy for $F_{\rm bol}$, i.e. implicitly assuming $F_{\rm X}\propto F_{\rm bol}$, and then proceed to theoretical interpretation." +" However, we have shown that this can lead to significant errors."," However, we have shown that this can lead to significant errors." +" For our data, we can estimate the 3-9 keV flux as the sum of the 3-5 keV flux and fraction of the 5-12 keV flux, the fraction being determined assuminga a power law with the 3-12keV photon index."," For our data, we can estimate the 3–9 keV flux as the sum of the 3–5 keV flux and a fraction of the 5–12 keV flux, the fraction being determined assuming a power law with the 3–12keV photon index." +" Using the flux in this band, we find p= 1.17+0.10, 1.23+0.10, and 1.48x0.05, for 2.25 GHz, 8.3 GHz,"," Using the flux in this band, we find $p=1.17\pm 0.10$ , $1.23\pm 0.10$ , and $1.48\pm 0.05$ , for 2.25 GHz, 8.3 GHz," +For a slow rotator (Vsinvy=2 kim-slj at low R region (R=5.000-20.000). Qpripix[Uo ,"For a slow rotator $V \sin i$ =2 $\rm{km\cdot s}^{-1}$ ) at low $R$ region (R=5,000-20,000), $Q_{\rm{DFDI}}\propto R^{0.63}$ ." +Since Doppler seusitivity Oty. is inversely proportional to two factors: Q aud WV. according to Equation (6)) aud (13)). the Doppler seusitivity becomes nearly independent of spectral resolution or the DEDI method (x£20-15) if the detection size (or total number of pixels) is fixed.," Since Doppler sensitivity $\delta v_{rms}$ is inversely proportional to two factors: $Q$ and $\sqrt{N_{e^-}}$ according to Equation \ref{eq:overall_Doppler}) ) and \ref{eq:overall_Doppler_2d}) ), the Doppler sensitivity becomes nearly independent of spectral resolution for the DFDI method $\propto R^{-0.13}$ ) if the detection size (or total number of pixels) is fixed." + This indicates that we can use quite moderate resolution spectrograph to disperse the stellar [ringes ooduced by the interferometer in a DEDI instrument while maintaining biel Doppler seusitivity., This indicates that we can use quite moderate resolution spectrograph to disperse the stellar fringes produced by the interferometer in a DFDI instrument while maintaining high Doppler sensitivity. + This opens a major door for multi-object Doppler measurements usiug the DEDI method as »xoposed by Ge (2002)., This opens a major door for multi-object Doppler measurements using the DFDI method as proposed by Ge (2002). +" La comparison. the Doppler seusitivity for the DE method still strongly lepends on spectral resolution [or a fixed number of detector pixels (x29""). indicating that ugher spectral resolution will offer better Doppler seusivity."," In comparison, the Doppler sensitivity for the DE method still strongly depends on spectral resolution for a fixed number of detector pixels $\propto R^{-0.57}$ ), indicating that higher spectral resolution will offer better Doppler sensivity." + As discussed in 8.2. and. )L3.. the DEDI instrument can be designed to have a moderate 'esolutiou spectrograph coupled with a Michelsou type interlerometer.," As discussed in \ref{sec:DFDI_Res} and \ref{sec:CCD_size}, the DFDI instrument can be designed to have a moderate resolution spectrograph coupled with a Michelson type interferometer." + Moderate spectral resolution allows a single order spectrum or a few order spectra to cover a broad waveleugth region iu the NIB 'eelon while keepiug the Doppler seusivitiv similar to a high resolution. DE design which requires a large detector array to cover spectra from a sinele target., Moderate spectral resolution allows a single order spectrum or a few order spectra to cover a broad wavelength region in the NIR region while keeping the Doppler sensivitiy similar to a high resolution DE design which requires a large detector array to cover spectra from a single target. + This indicates that the DEDI method las much greater potential for accomumnocdatiug multiple targets ou the same detector as proposed yw Ge (2002) than the DE instrument., This indicates that the DFDI method has much greater potential for accommodating multiple targets on the same detector as proposed by Ge (2002) than the DE instrument. + Iu order to evaluate the potential impact. of inulti-object DFDI instruments. we redefine the merit function (Q as: where Nop; is the number of objects that cau be wouitored simultaneously. and à is the index of iniportauce for multi-object observations.," In order to evaluate the potential impact of multi-object DFDI instruments, we redefine the merit function $Q^{\prime\prime}$ as: where $N_{obj}$ is the number of objects that can be monitored simultaneously, and $\alpha$ is the index of importance for multi-object observations." + From the perspective of photon count aud S/N. observations are equivalent to an increase of iN... and thus a is 0.5.," From the perspective of photon count and S/N, multi-object observations are equivalent to an increase of $N_{e^-}$, and thus $\alpha$ is 0.5." +" However. [rom au observational efficiency point of view. Q"" should be proportional to IN; because the more objects are observed simultaneously. the quicker the survey is accomplished. aud à is tlierefore equal to 1."," However, from an observational efficiency point of view, $Q^{\prime\prime}$ should be proportional to $N_{obj}$ because the more objects are observed simultaneously, the quicker the survey is accomplished, and $\alpha$ is therefore equal to 1." + We assume a detector that covers from 800 1un to 1350 nim at /22—100.000 so that we can use the (Q [actors obtained in 8L2..," We assume a detector that covers from 800 nm to 1350 nm at $R$ =100,000 so that we can use the $Q$ factors obtained in \ref{sec:DFDI_Res}." + Αρ6 is a coustant since we assume ideutical AA., $N_{e^-}$ is a constant since we assume identical $\Delta\lambda$. + AN); is inversely proportional to the number of pixels per object which is proportional to spectral resolution [ἱ (Equation (18)))., $N_{obj}$ is inversely proportional to the number of pixels per object which is proportional to spectral resolution $R$ (Equation \ref{eq:bandpass}) )). + Note that we do not require Vg; to be an integer because we can. in principle. fit a fraction of spectrum on a detector to make full use of the detector.," Note that we do not require $N_{obj}$ to be an integer because we can, in principle, fit a fraction of spectrum on a detector to make full use of the detector." + Qs for both DEDI aud DE are calculated., $Q^{\prime\prime}$ s for both DFDI and DE are calculated. + Fig., Fig. +" 10. shows the ratio of Q"" and Q,44,444 Lor DEDI under two dillerent assumptions ol a."," \ref{fig:Q_IRET_Q_DE_Res_multi} shows the ratio of $Q^{\prime\prime}$ and $Q^{\prime\prime}_{R=100,000}$ for DFDI under two different assumptions of $\alpha$." + For a=0.5. Le.. increase of Nog; is equivalent to photon gain. only a slight improvement is achieved if the detector is used for multi-object observations at lower resolution than 100.000.," For $\alpha$ =0.5, i.e., increase of $N_{obj}$ is equivalent to photon gain, only a slight improvement is achieved if the detector is used for multi-object observations at lower resolution than 100,000." +" Lu comparison. οι a survey efficiency poiut of view (Le.. a—1). we see a factor of [-6 times boost of Q"" in multi-object observations."," In comparison, from a survey efficiency point of view (i.e., $\alpha$ =1), we see a factor of$\sim$ 4-6 times boost of $Q^{\prime\prime}$ in multi-object observations." + The truncation at R=5.000 is due to a practical reason. that a lower resolution than 2.000 is rarely used in plauet survey using RV techuiques.," The truncation at $R$ =5,000 is due to a practical reason that a lower resolution than 5,000 is rarely used in planet survey using RV techniques." +Ou the other,On the other +from ? to derive a radial velocity distribution.,from \citet{2005MNRAS.360..974H} to derive a radial velocity distribution. + Since the Guitar pulsar already shows a transverse velocity of zz1600km/s this approach would not cover a wide spectrum of radial We take the pulsar position from ? and its proper motion from ?..?, Since the Guitar pulsar already shows a transverse velocity of $\approx\unit[1600]{km/s}$ this approach would not cover a wide spectrum of radial We take the pulsar position from \citet{2004MNRAS.353.1311H} and its proper motion from \citet{1993MNRAS.261..113H}. + give a current distance of 2kpc., \citet{1993ApJ...411..674T} give a current distance of $\unit[2]{kpc}$. +" We use this distance and the distance implied from dispersion measure (galactic electron density distribution model from ?)) of 1.86kpc (?) to determine the parallax (derived from the mean distance of 1.92kpc) and its error (note that this error does not include measurement uncertainties but only corresponds to the maximum deviation between the parallax derived from the mean distance and its lower and upper limit derived from the two cited distance values, those values are given without errors in the literature cited; see also 2)): where 4; is the proper motion in right ascension corrected for declination."," We use this distance and the distance implied from dispersion measure (galactic electron density distribution model from \citealt{2002astro.ph..7156C}) ) of $\unit[1{.}86]{kpc}$ \citep{2005AJ....129.1993M} to determine the parallax (derived from the mean distance of $\unit[1{.}92]{kpc}$ ) and its error (note that this error does not include measurement uncertainties but only corresponds to the maximum deviation between the parallax derived from the mean distance and its lower and upper limit derived from the two cited distance values, those values are given without errors in the literature cited; see also ): where $\mu_{\alpha}^*$ is the proper motion in right ascension corrected for declination." +" Heliocentric coordinates and velocity components of our sample of 140 OB associations and clusters (we use the term ""association"" in the following for both) are available in the Appendix of Paper I. We initially perform 100,000 Monte-Carlo runs for each of the 140 associations in our sample to find those for which close encounters with the Guitar pulsar are possible in the past 5Myr."," Heliocentric coordinates and velocity components of our sample of 140 OB associations and clusters (we use the term “association” in the following for both) are available in the Appendix of Paper I. We initially perform $100{,}000$ Monte-Carlo runs for each of the 140 associations in our sample to find those for which close encounters with the Guitar pulsar are possible in the past $\unit[5]{Myr}$." + Then we select those associations for which the smallest separation between its centre and the pulsar found at some time in the past was within three times the association radius or smaller than 100pc and repeated the procedure carrying out another one million runs., Then we select those associations for which the smallest separation between its centre and the pulsar found at some time in the past was within three times the association radius or smaller than $\unit[100]{pc}$ and repeated the procedure carrying out another one million runs. + lists those 9 associations for which the smallest separation found (column 2) during the second cycle of investigation was consistent with the association Four of them show a peak in their T-dmin contour plot that is consistent with the respective association radius., lists those 9 associations for which the smallest separation found (column 2) during the second cycle of investigation was consistent with the association Four of them show a peak in their $\tau$ $d_{min}$ contour plot that is consistent with the respective association radius. + gives this region of higher probability (columns 2 and 3; boundaries were chosen such that they reflect an approximate decline from the peak — see as an example) along with present-day parameters which the Guitar pulsar would have if it was born in the particular association., gives this region of higher probability (columns 2 and 3; boundaries were chosen such that they reflect an approximate decline from the peak – see as an example) along with present-day parameters which the Guitar pulsar would have if it was born in the particular association. + The last three columns indicate the distance of the potential supernova to the Sun as well as its equatorial coordinates., The last three columns indicate the distance of the potential supernova to the Sun as well as its equatorial coordinates. +" For the definition of selecting the values given in columns 4 to 10 we refer to our Paper I, section 5."," For the definition of selecting the values given in columns 4 to 10 we refer to our Paper I, section 5." +" It should just be mentioned that the errors given specify intervals, however are not 1c Owing to its large transverse velocity we expect the radial velocity to be very small."," It should just be mentioned that the errors given specify intervals, however are not $1\sigma$ Owing to its large transverse velocity we expect the radial velocity to be very small." +" Infact, ? investigated the bow shock which the pulsar is generating due to its high speed motion and measured an inclination of 90+20degrees."," Infact, \citet{2004ApJ...600L..51C} investigated the bow shock which the pulsar is generating due to its high speed motion and measured an inclination of $\unit[90\pm20]{degrees}$." + This means that the radial velocity is nearly zero., This means that the radial velocity is nearly zero. +" Indeed, we find an association in our which is consistent with that: If the pulsar originated from the Cygnus OB3 association (Cyg OB3) its radial velocity would be —27*5.km/s."," Indeed, we find an association in our which is consistent with that: If the pulsar originated from the Cygnus OB3 association (Cyg OB3) its radial velocity would be $\unit[-27^{+81}_{-70}]{km/s}$." + For Vulpecula OB1 (Vul OBI) the current radial velocity would also be relatively small (zz190 km/s); however this is significantly larger than for Cyg OB3., For Vulpecula OB1 (Vul OB1) the current radial velocity would also be relatively small $\approx \unit[190]{km/s}$ ); however this is significantly larger than for Cyg OB3. + The other two parent association candidates infer a substantial larger radial velocity., The other two parent association candidates infer a substantial larger radial velocity. +" We thus conclude that Cyg OB3 is most probably the birth association of the Guitar Note that NGC 6871, Byurakan 1, Byurakan 2 and NGC 6883 are clusters associated with Cyg OB3."," We thus conclude that Cyg OB3 is most probably the birth association of the Guitar Note that NGC 6871, Byurakan 1, Byurakan 2 and NGC 6883 are clusters associated with Cyg OB3." + They also appear in1., They also appear in. +. Their distances to the Sun are still under discussion so that they cover a range between 1.6 to 2.3kpc., Their distances to the Sun are still under discussion so that they cover a range between $1{.}6$ to $\unit[2{.}3]{kpc}$. +" Since NGC 6871 is generally presumed to be the nucleus of Cyg OB3, we repeat our calculations for NGC 6871 adopting the distance of Cyg OB3 (2.5 kpc)."," Since NGC 6871 is generally presumed to be the nucleus of Cyg OB3, we repeat our calculations for NGC 6871 adopting the distance of Cyg OB3 $\unit[2{.}5]{kpc}$ )." + The results do not change significantly: The smallest separation dmin found is 0.4pc and the peak in the T-dmin contour plot is located outside the boundary of NGC 6871 (radius of 7 pc)., The results do not change significantly: The smallest separation $d_{min}$ found is $\unit[0{.}4]{pc}$ and the peak in the $\tau$ $d_{min}$ contour plot is located outside the boundary of NGC 6871 (radius of $\unit[7]{pc}$ ). + We conclude that the Guitar pulsar was born somewhere, We conclude that the Guitar pulsar was born somewhere +"The results described in the previous section show that mm flux densities and gamma ray fluxes in sources with optical polarized emission (HPQ, LPQ) are correlated on monthly timescales, whereas for both quasars without polarization data (QSO) and BL Lacs (BLO) such a correlation is not found.","The results described in the previous section show that mm flux densities and gamma ray fluxes in sources with optical polarized emission (HPQ, LPQ) are correlated on monthly timescales, whereas for both quasars without polarization data (QSO) and BL Lacs (BLO) such a correlation is not found." +" Although the inverse Compton mechanism predicts a correlation between radio and y-ray emission strengths, studies from the era did not find evidence of this correlation (e.g.?).."," Although the inverse Compton mechanism predicts a correlation between radio and $\gamma$ -ray emission strengths, studies from the era did not find evidence of this correlation \citep[e.g.][]{anne_2003}." + This apparent contradiction with results obtained here is probably due to the limited sensitivity and to the sparse y-ray data available at that time., This apparent contradiction with results obtained here is probably due to the limited sensitivity and to the sparse $\gamma$ -ray data available at that time. + The absence of y- and radio flux density correlation for BL Lacs might be an indication that y-ray emission mechanisms or locations are different for quasars and BL Lacs (?).., The absence of $\gamma$ -ray and radio flux density correlation for BL Lacs might be an indication that $\gamma$ -ray emission mechanisms or locations are different for quasars and BL Lacs \citep{anne_2003}. +" In quasars, y-rays and radio flux densities are most significantly correlated for sources with the relativistic jet pointing close to our line of sight (HPQs), which implies that there is a strong coupling between the radio and the y-ray emission mechanisms."," In quasars, $\gamma$ -rays and radio flux densities are most significantly correlated for sources with the relativistic jet pointing close to our line of sight (HPQs), which implies that there is a strong coupling between the radio and the $\gamma$ -ray emission mechanisms." + That the correlation is seen in monthly flux density levels may imply that they have a co-spatial origin., That the correlation is seen in monthly flux density levels may imply that they have a co-spatial origin. + Perhaps the most remarkable result is that the most significant y-ray flux peaks occur when a mm-flare is either rising or peaking., Perhaps the most remarkable result is that the most significant $\gamma$ -ray flux peaks occur when a mm-flare is either rising or peaking. +" As discussed in Section 1, this indicates a sequence of events where first a disturbance (i.e. shock) emerges from the radio core, becoming visible as an increase in the mm radio flux (and as a new VLBI component, (?))), and that the y-ray flux rises and peaks."," As discussed in Section 1, this indicates a sequence of events where first a disturbance (i.e. shock) emerges from the radio core, becoming visible as an increase in the mm radio flux (and as a new VLBI component, \citep{jorstad_2001}) ), and that the $\gamma$ -ray flux rises and peaks." +" In other words, strong y-ray flares are produced in the same disturbances that"," In other words, strong $\gamma$ -ray flares are produced in the same disturbances that" +such as those at 1740 ((Schiminovich et al.,such as those at 1740 (Schiminovich et al. + 2001) are based on large portions of the skv., 2001) are based on large portions of the sky. + We have seen that uniform spherical models of reflection nebulae underestimate e if applied to optically thick situations (large7).. but not necessarily large τοι).," We have seen that uniform spherical models of reflection nebulae underestimate $a$ if applied to optically thick situations (large, but not necessarily large )." + Our reflection nebula models are not ideally suited for interpreting the DGL in the UV. since distributions of the dust and exciting stars are skewed towards (he Galactic plane.," Our reflection nebula models are not ideally suited for interpreting the DGL in the UV, since distributions of the dust and exciting stars are skewed towards the Galactic plane." + However. our results have strong implications as regards (he interpretation of the DGL as discussed in Henry (2002. hereafter H02).," However, our results have strong implications as regards the interpretation of the DGL as discussed in Henry (2002, hereafter H02)." + In that paper it was concluded that ας0.1 for A<1400A., In that paper it was concluded that $a\lesssim0.1$ for $\lambda\lesssim1400$. +. 1102 performed radiative transport calculations by using D.V magnitudes and the spectral (wpe for each star in the entire Hipparchos Input Catalog.," H02 performed radiative transport calculations by using $B, V$ magnitudes and the spectral type for each star in the entire Hipparchos Input Catalog." + The reddening. (0—V). follows from (5—V) aud spectral (wpe.," The reddening, $E(B-V)$, follows from $(B-V)$ and spectral type." + The extinction ly is E(D—V)fàCl/A). where Dads usually taken to be 3.1.," The extinction $A_\lambda$ is $E(B-V)\,R_V\,(A_\lambda/A_V)$, where $R_V$ is usually taken to be 3.1." + The for each star is 4(A)/ 1.036., The for each star is $A(\lambda)$ /1.086. + The heart of H02 is its unique treatment of radiative transfer., The heart of H02 is its unique treatment of radiative transfer. + The scattering is taken io be independently of the geometry., The scattering is taken to be independently of the geometry. + However. Witt. Friedinann. sasseen (1997) have shown from clumpy models that the distribution of cloud optical depths can strongly affect the interstellar radiation field.," However, Witt, Friedmann, Sasseen (1997) have shown from clumpy models that the distribution of cloud optical depths can strongly affect the interstellar radiation field." + At low optical depths. the II02 relation between and approaches the uniform model.," At low optical depths, the H02 relation between and approaches the uniform model." + The dotted lines in Figure 1. show the HO2 mocels for three albedos identified by (he nearby solid lines.," The dotted lines in Figure \ref{fig1} + show the H02 models for three albedos identified by the nearby solid lines." + We see that the HO2 assumption predicts more scattered light than uniform models. and a lower albedo follows.," We see that the H02 assumption predicts more scattered light than uniform models, and a lower albedo follows." + For AS2000A.. the Που most hot stars is probably rather large. since the mean (VW) is 1 magnitude | and the ultraviolet is 3. 5 times larger.," For $\lesssim2000$, the for most hot stars is probably rather large, since the mean $A(V)$ is $\sim$ 1 magnitude $^{-1}$ and the ultraviolet is 3 – 5 times larger." +" For = 1.5. Figure 1 shows that the H02 model with « = 0.6 predicts a scattered [Iux ~exp(0.4) = 1.5 times larger (han (hat ""tvpical hierarchical model. chosen for illustrative purposes only."," For = 1.5, Figure \ref{fig1} shows that the H02 model with $a$ = 0.6 predicts a scattered flux $\sim$ exp(0.4) = 1.5 times larger than that “typical” hierarchical model, chosen for illustrative purposes only." + The conclusion of HO02 that (1000 AA)) is verv low (“perhaps 0.17) is perhaps understandable., The conclusion of H02 that $a$ ) is very low (“perhaps 0.1”) is perhaps understandable. + Such a low albedo is not compatible with the observations of by Witt et al. (, Such a low albedo is not compatible with the observations of by Witt et al. ( +1993). in which the scattered (hix was larger than the stella. or Murthy et al. (,"1993), in which the scattered flux was larger than the stellar, or Murthy et al. (" +1993). which suggests that a(1150. )) e0.55 if of the extinction is within the nebula. as assumed by Witt οἱ al. (,"1993), which suggests that $a$ (1150 ) $\sim 0.55$ if of the extinction is within the nebula, as assumed by Witt et al. (" +1992).,1992). + Aburthy. Henry. Holberg (1991) used.2 to search for DGL for A<1300AA.," Murthy, Henry, Holberg (1991) used to search for DGL for $\lambda\le1300$." +. They failed to detect any. possibly confirming the LI02 picture.," They failed to detect any, possibly confirming the H02 picture." + Our reflection nebula results commonly show that the same nebula can show verv low and rather high scattered, Our reflection nebula results commonly show that the same nebula can show very low and rather high scattered +"ib ds automatically implemented bv (he integral equation for 2, (Eq. 20)):",it is automatically implemented by the integral equation for $P_u$ (Eq. \ref{eq:Pu}) ): + it thus provides iimpor(ant numerical check of our computing accuracy., it thus provides an important numerical check of our computing accuracy. + The slow convergence of the iterative solution of Eq., The slow convergence of the iterative solution of Eq. + 16. has been partly corrected bv use ol the following ruse.," \ref{mainPd} + has been partly corrected by use of the following ruse." +" The conditional probability 2, obevs the following integral constraint.", The conditional probability $P_d$ obeys the following integral constraint. + From Eq., From Eq. + 2. we see that at downstream infinitv the condition Of/O2=0 implies that the only physically acceptable solution for the distribution function is /(j)=constantIN., \ref{eq:transport} we see that at downstream infinity the condition $\partial f/\partial z=0$ implies that the only physically acceptable solution for the distribution function is $f(\mu)=constant=K$. + At downstream infinity we can therefore write which implies the condition for anv value of the exit angle µ., At downstream infinity we can therefore write which implies the condition for any value of the exit angle $\mu$. +" Since the funcüon 2, depends only on the scattering properties of the fluid. which are assumed to be independent of z. Eq."," Since the function $P_d$ depends only on the scattering properties of the fluid, which are assumed to be independent of $z$, Eq." + 50. also applies to any 2. including z—0 which is the position of the shock.," \ref{eq:property} also applies to any $z$, including $z=0$ which is the position of the shock." + This condition is of course automatically satisfied at the end of our ealeulations. while 1 is only approximately satisfied during the iterative solution of Eq. 16..," This condition is of course automatically satisfied at the end of our calculations, while it is only approximately satisfied during the iterative solution of Eq. \ref{mainPd}." + We have found oul that. bv correcting5 the normalization of P;[ so that it obevs the previous integral5 constraint to hold at the end ofevery iteration. the convergence is much accelerated. above all in the SPAS regime.," We have found out that, by correcting the normalization of $P_d$ so that it obeys the previous integral constraint to hold at the end of iteration, the convergence is much accelerated, above all in the SPAS regime." + We are unable to explain why (his occurs. but shortening of the computation times by a [actor of 3. or more. have thus been achieved.," We are unable to explain why this occurs, but shortening of the computation times by a factor of 3, or more, have thus been achieved." +" Once the two functions 2, and 2, have been calculated. (he slope of the spectrum. as discussed in paper 1. is given by the solution of the equation: where Here t,—4T is (he relative velocity between the upstream and downstream [fIuids and gy) is the angular part of the distribution function of the accelerated particles. which contains all the information about the anisotropy."," Once the two functions $P_u$ and $P_d$ have been calculated, the slope of the spectrum, as discussed in paper I, is given by the solution of the equation: where Here $u_{rel}=\frac{u-u_d}{1-u u_d}$ is the relative velocity between the upstream and downstream fluids and $g(\mu)$ is the angular part of the distribution function of the accelerated particles, which contains all the information about the anisotropy." +" A point to notice concerns the function 77, ", A point to notice concerns the function $P_u$ + ~5O% (Furlancttoetal.2006:MeQu," $\sim 50\%$ \citep{fur06,mcq07}." +inu2007).. =50% 2=6.5 (Rhoads&Malhotra2001:Malhotra&Rhoads2006).," $\lesssim 50\%$ $z=6.5$ \citep{rho01,mr04, ste05, kas06,mr06}." +. lower μι2IX zz6.3. τν=10.521.2 (I&oiunatsuetal.2010)., $lower$ $x_{HI} \gtrsim 1 \%$ $z\approx 6.3$ $z_{re}=10.5 \pm 1.2$ \citep{kom10}. +". true for low mass galaxies. às. eenütters are observed to have stellar masses M,X (Cawiseretal.2006:Pirzkal2007:Finkelsteinetal.2007:Peutericci 2009).. appreciably below the stellar masses of Lyman break selected ealaxies (LBC) (Steideletal.1996) at similar redshitts (οιο,Pa-povichetal.2001:ShapleyStark 2009)."," true for low mass galaxies, as emitters are observed to have stellar masses $\rm M_{\star} \lesssim 10^{9} \; +M_{\odot}$ \citep{gaw06, pir07, fin07,pen09}, , appreciably below the stellar masses of Lyman break selected galaxies (LBG) \citep{ste96} at similar redshifts \citep[e.g.][]{pap01, sha01, sta09}." +. Narrow-band imagine is a well established. technique for finding high redshift galaxies (0.9.al. 20093..," Narrow-band imaging is a well established technique for finding high redshift galaxies \citep[e.g.][]{rho00a, rho04, rho03, mr02, + mr04, cow98, hu99, hu02, hu04, kud00, fyn01, pen00, ouc01, ouc03, + ouc08, sti01, shi06, kod03, aji04, tan05, ven04, kas06, iye06, + nil07, fin09}." + The method works because ecluission redshifted iuto a narrow band filter will male je Cluitting ealaxics appear brighter iu images through wat filter than in broadhbauds of similar waveleneth., The method works because emission redshifted into a narrow band filter will make the emitting galaxies appear brighter in images through that filter than in broadbands of similar wavelength. +" A supplemental requireiieut that the selected cussion line ealaxies be faint or undetected in filters bhieward of 1ο narrowband filter effectively weeds out lower redshift cluission line objects οιο,Malhotra&Rhoads2002).", A supplemental requirement that the selected emission line galaxies be faint or undetected in filters blueward of the narrowband filter effectively weeds out lower redshift emission line objects \citep[e.g.][]{mr02}. +. This has proven to be verv efficicut for sclecting star-ormlng ealaxies up to τς7. and remains effective even when those galaxies are too faint iu their coutimmun enissiou to be detected iu typical broadband surveys.," This has proven to be very efficient for selecting star-forming galaxies up to $z\lesssim 7$, and remains effective even when those galaxies are too faint in their continuum emission to be detected in typical broadband surveys." + While large samples of. eenütters have been detected at 2 <6. both survey voles and sample sizes are mach smaller at 2 76.," While large samples of emitters have been detected at $z<$ 6, both survey volumes and sample sizes are much smaller at $z>$ 6." + Since the pphotous are resonantly scattered in neutral IGML a decline in the observed Iuuinositv function (LE) of celuitters would sugeest a chauge in the IGAL phase. assumndue the nmuuber density ofnewly formed ealaxics," Since the photons are resonantly scattered in neutral IGM, a decline in the observed luminosity function (LF) of emitters would suggest a change in the IGM phase, assuming the number density ofnewly formed galaxies" +peak that provide contributions to the x? smaller than the computer precision.,peak that provide contributions to the $\chi^2$ smaller than the computer precision. +" However, this background has no physical significance."," However, this background has no physical significance." + Figure 4 shows how the results change when the random errors are added to the calculated intensities., Figure \ref{1-isot-2} shows how the results change when the random errors are added to the calculated intensities. +" The plasma EM obtained with W=0.05 is displayed as a full black line, and the EM curves obtained by the five different ""noisy"" datasets are shown in red."," The plasma EM obtained with $W=0.05$ is displayed as a full black line, and the EM curves obtained by the five different “noisy” datasets are shown in red." +" The main peak remains unaltered, and since the wings, always present, are much lower than the peak value the technique is still able to recover the single peak."," The main peak remains unaltered, and since the wings, always present, are much lower than the peak value the technique is still able to recover the single peak." +" However, the background noise now is more significant and secondary peaks, much lower than the one at logT=6.0, are sometimes present such as those at logT=6.30 and 6.45."," However, the background noise now is more significant and secondary peaks, much lower than the one at $\log T=6.0$, are sometimes present such as those at $\log T=6.30$ and 6.45." +" Figure 5 displays the results of the reconstruction of the EM(T) curve in the case of two isothermal components with the same EM (1log43.0 each, in cm?, shown as red diamonds) and with variable temperature separation AlogT and no random errors."," Figure \ref{2-isot-1} displays the results of the reconstruction of the $EM(T)$ curve in the case of two isothermal components with the same EM $\log 43.0$ each, in $^{-3}$, shown as red diamonds) and with variable temperature separation $\Delta \log T$ and no random errors." +" The bin width W in each reconstruction was kept smaller than the temperature separation, and when possible we experimented with increasing W to determine how it affected the ability of the technique to separate the two components."," The bin width $W$ in each reconstruction was kept smaller than the temperature separation, and when possible we experimented with increasing $W$ to determine how it affected the ability of the technique to separate the two components." +" We also studied AlogT=0.20 and 0.25, but the results are no different than those obtained with AlogT=0.15."," We also studied $\Delta \log +T=0.20$ and 0.25, but the results are no different than those obtained with $\Delta \log +T=0.15$." +" Table 2 reports the results for the reconstructions obtained with AlogT=0.10,0.15 and 0.20."," Table \ref{metrics} reports the results for the reconstructions obtained with $\Delta \log T=0.10, 0.15$ and 0.20." +" Figure 5 shows that even when random errors are absent the MCMC technique is unable to resolve the two components when their temperature separation is smaller than AlogT 0.10, and provides a single, very noisy peak with a larger width."," Figure \ref{2-isot-1} shows that even when random errors are absent the MCMC technique is unable to resolve the two components when their temperature separation is smaller than $\Delta \log T=0.10$ , and provides a single, very noisy peak with a larger width." +" Also, small values of the bin width W worsen the noise in the solution, further preventing the separation of the two components."," Also, small values of the bin width $W$ worsen the noise in the solution, further preventing the separation of the two components." +" At AlogT=0.10 a complete separation of the two peaks can be achieved only when W=0.05, because noise confuses the result for smaller bin widths and the presence of high EM values in the only bin between the peaks prevents definitive conclusions about separate components."," At $\Delta \log T=0.10$ a complete separation of the two peaks can be achieved only when $W=0.05$, because noise confuses the result for smaller bin widths and the presence of high EM values in the only bin between the peaks prevents definitive conclusions about separate components." + Only when AlogT=0.15 the two components can be fully and convincingly resolved with W=0.05., Only when $\Delta \log T=0.15$ the two components can be fully and convincingly resolved with $W=0.05$. +" Noise in the definition of each peak is significant at any value of ALogT for W«0.05; the bin W=0.1 allows us to resolve the two peaks only when AlogT=0.20 and causes trouble if the separation is AlogT= 0.15, providing a large y?."," Noise in the definition of each peak is significant at any value of $\Delta Log T$ for $W<0.05$; the bin $W=0.1$ allows us to resolve the two peaks only when $\Delta \log T=0.20$ and causes trouble if the separation is $\Delta \log T=0.15$ , providing a large $\chi^2$." + The presence of random errors further limits the ability of the MCMC technique to resolve the two components., The presence of random errors further limits the ability of the MCMC technique to resolve the two components. +" Figure 6 shows the MCMC reconstruction for AlogT=0.15 and 0.20 in the presence of random errors, with W=0.05."," Figure \ref{2-isot-2} shows the MCMC reconstruction for $\Delta \log T=0.15$ and 0.20 in the presence of random errors, with $W=0.05$." +" The two peaks are still visible but badly resolved when AlogT 0.15, while they are well separated in the other case."," The two peaks are still visible but badly resolved when $\Delta \log T=0.15$ , while they are well separated in the other case." +" Thus, we conclude that AlogT=0.20 is a reliable estimate of the smallest temperature distance between two isothermal components which the MCMC techniquecan realistically resolve when the lines listed in Table 1 areavailable."," Thus, we conclude that $\Delta \log T=0.20$ is a reliable estimate of the smallest temperature distance between two isothermal components which the MCMC techniquecan realistically resolve when the lines listed in Table \ref{lines} areavailable." +" If only a smaller number of lines can be used, the minimum AlogT "," If only a smaller number of lines can be used, the minimum $\Delta \log T$ " +We observed Arp220 on 9 November 2003 with a global array of the largest telescopes on Earth.,We observed Arp220 on 9 November 2003 with a global array of the largest telescopes on Earth. + The observations were part of a long-term monitoring campaign aimed at stucdvine the changes in the compact structure of both the eonünuum anc OIL maser emission in Arp220., The observations were part of a long-term monitoring campaign aimed at studying the changes in the compact structure of both the continuum and OH maser emission in Arp220. + 15 antennas took part in the experiment. including the 10 VLBA antennas. the phased VLA. the GBT. Arecibo. and 5 antennas of the European VLBI Network (EVN).," 18 antennas took part in the experiment, including the 10 VLBA antennas, the phased VLA, the GBT, Arecibo, and 5 antennas of the European VLBI Network (EVN)." + All antennas performed nominally except for the phased-VLÀ; it was unable to maintain phase coherence due to the presence of interference. and so all phased-VLA data were flagged.," All antennas performed nominally except for the phased-VLA; it was unable to maintain phase coherence due to the presence of interference, and so all phased-VLA data were flagged." + The observations covered a period of almost 21 hours., The observations covered a period of almost 21 hours. + Data were taken in 2 polarizations and in four & Mz wide intermediate Lrequeney (IE). bands covering (he Irequency. range 1633.99 1665.99 MIIz and with an aggregate data rate of 256 Mbit per second., Data were taken in 2 polarizations and in four 8 MHz wide intermediate frequency (IF) bands covering the frequency range 1633.99 – 1665.99 MHz and with an aggregate data rate of 256 Mbit per second. + The strong recshifted OIL maser emission from Arp220 was centered in the lowest of the four IF bands., The strong redshifted OH maser emission from Arp220 was centered in the lowest of the four IF bands. + The line-[ree portions of the lowest IF. plus the upper 24 MlIIz from the other 3 [Fs were dedicated to observing the continuum emission from. Arp220.," The line-free portions of the lowest IF, plus the upper 24 MHz from the other 3 IFs were dedicated to observing the continuum emission from Arp220." + The data were recorded on tape and shipped to the VLBA correlator in Socorro. New Mexico.," The data were recorded on tape and shipped to the VLBA correlator in Socorro, New Mexico." + There. (wo correlator passes were done. one with 512 channels across the 8 MlIIz of the bottom IF containing the OIL line. and à second with 32 channels per IF across all flour IEs.," There, two correlator passes were done, one with 512 channels across the 8 MHz of the bottom IF containing the OH line, and a second with 32 channels per IF across all four IFs." + The correlated data were sent to Jodrell Bank Observatory [or subsequent data processing., The correlated data were sent to Jodrell Bank Observatory for subsequent data processing. + Due to the presence of strong and compact OIL maser emission (Lonsdale οἱ al., Due to the presence of strong and compact OH maser emission (Lonsdale et al. + 1994. 1993). which could be used for phase calibration. there was no need lor phase relerencing.," 1994, 1998), which could be used for phase calibration, there was no need for phase referencing." +" ""Therefore. the observing scheme was straightforward. with most time being spent on Arp220 and wilh occasional 15 minutes scans on calibrator sources for the purposes of tracking the residual delays aud determining the bandpass corrections."," Therefore, the observing scheme was straightforward, with most time being spent on Arp220 and with occasional 15 minutes scans on calibrator sources for the purposes of tracking the residual delays and determining the bandpass corrections." + The calibrator sources used [or {his purpose were 16134-3241 and 1516-193., The calibrator sources used for this purpose were 1613+341 and 1516+193. + The J2000 coordinates assumed for Arp220 were: RA = 15h 34m 57.225s. Dec = 23° 30° 11.564.," The J2000 coordinates assumed for Arp220 were: RA = 15h 34m 57.225s, Dec = $^{\circ}$ 30' 11.564""." + ΑΙ data calibration was earried out using the ISDECOL release of NILAOs AIPS software and used standard: procedures (Diamond. 1995). to summarize: the data were corrected for digital effects and those of parallactic angle rotation: the eross residual delay. errors in each observing band ancl (he phase offsets between the bands were determined: amplitude calibration was performed using (he svstem temperature and gain curves supplied by the observatories: anv lime variable residual delay offsets and (he full complex bandpass solution for each antenna were determined: the residual antenna-based phase delay rates ol the strongest channel in the OIL spectrum was measured al 5 minute time intervals.," All data calibration was carried out using the 15DEC01 release of NRAO's AIPS software and used standard procedures (Diamond, 1995), to summarize: the data were corrected for digital effects and those of parallactic angle rotation; the gross residual delay errors in each observing band and the phase offsets between the bands were determined; amplitude calibration was performed using the system temperature and gain curves supplied by the observatories; any time variable residual delay offsets and the full complex bandpass solution for each antenna were determined; the residual antenna-based phase delay rates of the strongest channel in the OH spectrum was measured at 5 minute time intervals." + All of these calibration factors were (hen applied to the data to produce a data sel coherent in lime and frequency., All of these calibration factors were then applied to the data to produce a data set coherent in time and frequency. + Then. an iterative loop of sell-calibration and imaging was performed on the OIL reference channel to determine and correct for the short timescale ( 1 minute)," Then, an iterative loop of self-calibration and imaging was performed on the OH reference channel to determine and correct for the short timescale ( 1 minute)" +"Lagil dd"" Now we have ο] -- -29 τμ Lu Ον]. atm(11) where the affine parameter AU) meH! Jas0 Note that D4(à=0|z) A(z).","_0^z dz” Now we have |z) = - _m _0^z |z') + =0|z), where the affine parameter (z) = _0^z Note that $D_A(\tilde{\alpha}=0|z)=\lambda(z)$ ." + Wang(2000b) has shown that 2). where r(z) is the comoving distance in a smooth universe.," \cite{SNfluxavg} has shown that $D_A(\tilde{\alpha}=1|z)=r(z)/(1+z)$ , where $r(z)$ is the comoving distance in a smooth universe." + No approximations have been made in obtaining Eq.(2)) fom the Dver-Roeder equation I«q.(6))., No approximations have been made in obtaining \ref{eq:DR2}) ) from the Dyer-Roeder equation \ref{eq:DR}) ). + We can solve the Eq.(2)) perturbatively by replacing D4(à2’) in the integral on the rght hand side with its previous order approximation., We can solve the \ref{eq:DR2}) ) perturbatively by replacing $D_A(\tilde{\alpha}|z')$ in the integral on the right hand side with its previous order approximation. + In the first order perturbation. we replace D4(6|z) in the integral by its Oth order approximation. D(àz)Da(a-lz:).," In the first order perturbation, we replace $D_A(\tilde{\alpha}|z')$ in the integral by its 0th order approximation, $D_A^{(0)}(\tilde{\alpha}|z')=D_A(\tilde{\alpha}=1|z')$." + This gives us 2)= —0l:) =Hz) —1 where we have defined =| PRSE --- [suse LOUG) (Nang2000b)..," This gives us |z)= + =0|z) =1|z) -1 ) ], where we have defined 1 - = - _0^z (1+z) =1|z)=r(z)/(1+z)$ \citep{SNfluxavg}." + Note thatA(5)<0., Note that $\tilde{\kappa}_{min}(z) <0 $. + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|C," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|C4," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|C4(," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|C4(:," After four more iterations, we find the fifth order perturbative solution =1 - 1-" + After [our more iterations. we lind the fifth order perturbative solution Dy(A=lj=l 1—à|C4(:)," After four more iterations, we find the fifth order perturbative solution =1 - 1-" +hole mass rather than fixing it at &4x109M. appropriate lor Ser A*: second. we focus on capture events that produce non-sell-gravitating disks. which therefore settle into very. thin Keplerian disks rather (han form stars.,"hole mass rather than fixing it at $\approx4\ee 6 \msun $ appropriate for Sgr A*; second, we focus on capture events that produce non-self-gravitating disks, which therefore settle into very thin Keplerian disks rather than form stars." + Consider. then. an extended cloud with hydrogen column density Nj=No;xLO 7 that enters the black hole's sphere of influence with speed e=225(o;kms.|.," Consider, then, an extended cloud with hydrogen column density $N_\mathrm{H}=N_{23}\ee 23 $ $^{-2}$ that enters the black hole's sphere of influence with speed $v = 225\,v_{225}\kms$." + On its passage through the region. the cloud temporarily engulls the black hole and shocks induced by eravitational locussing of the material passing close to the hole lead to capture of some malerial. which circularizes.," On its passage through the region, the cloud temporarily engulfs the black hole and shocks induced by gravitational focussing of the material passing close to the hole lead to capture of some material, which circularizes." + Rapid cooling of the gas leads to formation of a thin. Ixeplerian disk.," Rapid cooling of the gas leads to formation of a thin, Keplerian disk." + In (his scenario (he disk mass is determined by (he mass of gas will impact parameter inside the IIloxleLyttleton radius (Hovle Lyttelton 1939) and ils size is determined by the imperfect cancellation of the angular momentum of the captured material that approached on opposing sides of the black hole due to the cloud morphology and structure., In this scenario the disk mass is determined by the mass of gas with impact parameter inside the Hoyle–Lyttleton radius (Hoyle Lyttelton 1939) and its size is determined by the imperfect cancellation of the angular momentum of the captured material that approached on opposing sides of the black hole due to the cloud morphology and structure. + As in Paper I we characterize the uncertain capture dvnamies using (wo parameters., As in Paper I we characterize the uncertain capture dynamics using two parameters. +" The first of these. &. is (he ratio of the captured mass to the estimate. allowing us to write the disk mass as where mg, is the mass of a hydrogen atom ancl we have assiuned the stancarcd helium to hydrogen of 0.4 by mass."," The first of these, $\kappa$, is the ratio of the captured mass to the Hoyle--Lyttleton estimate, allowing us to write the disk mass as where $m_\mathrm{H}$ is the mass of a hydrogen atom and we have assumed the standard helium to hydrogen of 0.4 by mass." + We expect that &~1 because bulk kinetic energy is efficiently lost through shocking and rapid radiative cooling during the collision Edgar 2004. and references therein).," We expect that $\kappa \sim 1$ because bulk kinetic energy is efficiently lost through shocking and rapid radiative cooling during the collision Edgar 2004, and references therein)." + The second parameter. A. is the average ratio of the specific angular momentum of a [hid element in the resulting Keplerian disk. (GUr)7. to its initial angular momentum be.," The second parameter, $\lambda$, is the average ratio of the specific angular momentum of a fluid element in the resulting Keplerian disk, $(GMr)^{1/2}$, to its initial angular momentum $bv$." + This ignores a number of complications: the cancellation of angular monmentun is sensitive to the distribution of inhomogeneities in the cloud and will be highly variable: it will also depend on the orientation of a fluid elements trajectory relative to the final disk plane., This ignores a number of complications: the cancellation of angular momentum is sensitive to the distribution of inhomogeneities in the cloud and will be highly variable; it will also depend on the orientation of a fluid elements trajectory relative to the final disk plane. + While [Iuid elements with the same impact parameter 5 but with dillerent specific angular momentum vectors will end up on different disk radii. (he convolution in mapping the initial impact parameter 6 (o eventual location in the disk r also has the effect of averaging oul these differences.," While fluid elements with the same impact parameter $b$ but with different specific angular momentum vectors will end up on different disk radii, the convolution in mapping the initial impact parameter $b$ to eventual location in the disk $r$ also has the effect of averaging out these differences." + This assumption allows us to relate the impact parameter of a fluid element in (he approaching cloud. to its typical location in the disk once it has been captured. anc," This assumption allows us to relate the impact parameter of a fluid element in the approaching cloud, to its typical location in the disk once it has been captured and" + This assumption allows us to relate the impact parameter of a fluid element in (he approaching cloud. to its typical location in the disk once it has been captured. ancl," This assumption allows us to relate the impact parameter of a fluid element in the approaching cloud, to its typical location in the disk once it has been captured and" +accretion contribution to the luuinosity vields outliers with lieh accretion rates (see Figure 8)).,accretion contribution to the luminosity yields outliers with high accretion rates (see Figure \ref{fig:acc}) ). +" Iu Figure 7 we show the strong correlation between Που and L,|Lage. consistent with dust sublimation temperatures of T|—L5002000 for cTTs disks and sheltly lower temperatures. ~10001500 for ILAeDe disks."," In Figure \ref{fig:lrin_acc} we show the strong correlation between $R_\mathrm{CO}$ and $L_\star+L_\mathrm{acc}$, consistent with dust sublimation temperatures of $T\sim1500-2000$ for cTTs disks and slightly lower temperatures, $\sim1000-1500$ for HAeBe disks." + Transitional disk radii are larger. but also follow a similar trend (discussed further in Section 6.1)).," Transitional disk radii are larger, but also follow a similar trend (discussed further in Section \ref{sec:trans}) )." +" Tucluding all disks. the p-value associated with linear reeression of log(Rog) against log(L,|Lace) is 610.7. and excluding transitional disks. it is 101"". so the Increase m radius with hnunuinositv is highly statistically sieuificaut."," Including all disks, the $p$ -value associated with linear regression of $\log(R_\mathrm{CO})$ against $\log(L_\star+L_\mathrm{acc})$ is $6\times 10^{-5}$, and excluding transitional disks, it is $10^{-10}$, so the increase in radius with luminosity is highly statistically significant." + Therefore. we believe this is stroug evidence or a dependence of Rew on dust sublimation.," Therefore, we believe this is strong evidence for a dependence of $R_\mathrm{CO}$ on dust sublimation." + The slope of the trend (excluding transitional dis io£01. is somewhat stecper than the slope of 0.5 expected from dust sublimation alone. at the 2 σ level.," The slope of the trend (excluding transitional disks), $0.7 \pm 0.1$, is somewhat steeper than the slope of $0.5$ expected from dust sublimation alone, at the 2 $\sigma$ level." + Tn more detail Ree may be set by ai balance of Niotodissociation (which cats outward in the disk) aud accretion (which replenishes iuwrd).," In more detail, $R_\mathrm{CO}$ may be set by a balance of photodissociation (which eats outward in the disk) and accretion (which replenishes inward)." + This could result in relatively smaller Rew for T Town disks and larecr or TAcBe disks due to the respectively lower aud lieber Niotodissociating UV. fluxes., This could result in relatively smaller $R_\mathrm{CO}$ for T Tauri disks and larger for HAeBe disks due to the respectively lower and higher photodissociating UV fluxes. + Tf dust acts as a shield against photoclissociation. then we should also expect a correlation between Που aud measured dust inuer radii from near-IR iuterferometers.," If dust acts as a shield against photodissociation, then we should also expect a correlation between $R_\mathrm{CO}$ and measured dust inner radii from near-IR interferometers." + We show these radi in Figure 9.. along with a line marking a l:l correlation. aud another marking the best liueoar fit.," We show these radii in Figure \ref{fig:gas_dust}, along with a line marking a 1:1 correlation, and another marking the best linear fit." + Although there is a statistically siguificaut correlation between the two variables. there is also sienificant scatter. of order 0.5 dex.," Although there is a statistically significant correlation between the two variables, there is also significant scatter, of order 0.5 dex." + Uncertaiuties in Rew can be a factor of a few (see Section L2.0)).. which may account for some of the scatter.," Uncertainties in $R_\mathrm{CO}$ can be a factor of a few (see Section \ref{sec:errors}) ), which may account for some of the scatter." + Also. dust iuner radii are measured usime differeut models iu different studies Guclidingthin rings. disks. both iuclined aud not inclined). and so it is possible that some of the scatter arises from the choice of dust model. or from dust mocels uot accounting for the disk iuclination.," Also, dust inner radii are measured using different models in different studies (including thin rings, disks, both inclined and not inclined), and so it is possible that some of the scatter arises from the choice of dust model, or from dust models not accounting for the disk inclination." +" Using a colerent suuple of iuterferometric visibilities. analyzed in the sale wav. and accounting for disk iuclinatiou. could test this hypothesis,"," Using a coherent sample of interferometric visibilities, analyzed in the same way, and accounting for disk inclination, could test this hypothesis." + Another possibility is that there is real scatter due to different rates of photodissociation and/or replenislincut via accretion., Another possibility is that there is real scatter due to different rates of photodissociation and/or replenishment via accretion. + This is consistent with the observation that several WAeBe disks have anomalously large Row. while several cTTs’s have anomalously siuall Που.," This is consistent with the observation that several HAeBe disks have anomalously large $R_\mathrm{CO}$ , while several cTTs's have anomalously small $R_\mathrm{CO}$." + As discussed iu several prior studies (Najitaetal.2003:Rettigetal.200[:Salvk 2009).. Rew lies well iuside Renu. for many transitional disks (where τρως is the radius at which the disk becomes optically thick).," As discussed in several prior studies \citep{Najita03,Rettig04,Salyk09}, $R_\mathrm{CO}$ lies well inside $R_\mathrm{trans}$ for many transitional disks (where $R_\mathrm{trans}$ is the radius at which the disk becomes optically thick)." + This result has also been coufiiied via spectroastrometry for IID 135311 Band TW Uva (Poutoppicdanetal.2008)., This result has also been confirmed via spectroastrometry for HD 135344 B and TW Hya \citep{Pontoppidan08}. +. Iu Figure 10.. we closely reproduce the results of Salvketal. (2009).. who showed that the best-fit Pe lie somewhere between Ray and Ru for all disks.," In Figure \ref{fig:radplot_trans}, we closely reproduce the results of \citet{Salyk09}, who showed that the best-fit $R_\mathrm{CO}$ lie somewhere between $R_\mathrm{sub}$ and $R_\mathrm{trans}$ for all disks." + Au iuportaut caveat is that. as discussed in Section L.2.0.. Rew. the upper limit to Reg is not well constrained for IID 135311 D. ΠΟ 111569 A aud SR 21.," An important caveat is that, as discussed in Section \ref{sec:modelresults}, $R_\mathrm{CO}$ , the upper limit to $R_\mathrm{CO}$ is not well constrained for HD 135344 B, HD 141569 A and SR 21." + The cinission iu SR 21 aud ΠΟ 111560 A could originate at or near Repay. as is observed for the transitional disk IID 1005162009):: however. for ΠΟ 135311 D. some of the CO eas must originate from within the iuner. optically thin reeion (Poutoppidauctal.2008).," The emission in SR 21 and HD 141569 A could originate at or near $R_\mathrm{trans}$, as is observed for the transitional disk HD 100546; however, for HD 135344 B, some of the CO gas must originate from within the inner, optically thin region \citep{Pontoppidan08}." +. Using the procedures described in Section 12.0... we derive sieuificantly larger radii for transitional disks tla for classical cüsks at the same Inninuositv.," Using the procedures described in Section \ref{sec:procedure}, we derive significantly larger radii for transitional disks than for classical disks at the same luminosity." + This ciffereuce is typically an order of magnitude much larger than our uncertainties and holds true for almost three orders of maguitude im Iuninositv., This difference is typically an order of magnitude — much larger than our uncertainties — and holds true for almost three orders of magnitude in luminosity. + Tere we discuss some possible explanations for this difference., Here we discuss some possible explanations for this difference. + Oue possible explanation is that this is the result of ποιο systematic bias du our analysis., One possible explanation is that this is the result of some systematic bias in our analysis. + Two possibilitics cole to mind., Two possibilities come to mind. +" First. since Rew depends on the line wines. one nüeht derive smaller immer radi by “fitting ie noise"" at large velocities."," First, since $R_\mathrm{CO}$ depends on the line wings, one might derive smaller inner radii by “fitting the noise” at large velocities." + Towever. since trausitional disks tend to have lower contiuununu S/N levels. aud ower line/contiuuui ratios than classical disks (Salvk this would teud to bias oue toward ceriving arecr characteristic velocities and lencesmaller ΠΙΟ: radi for transitional disks.," However, since transitional disks tend to have lower continuum S/N levels, and lower line/continuum ratios than classical disks \citep{Salyk11}, this would tend to bias one toward deriving larger characteristic velocities and hence inner radii for transitional disks." + Iu addition. we fud that 1ο transitional disk radius discrepancy holds even if we exunme Z4 derived from the TWOAL," In addition, we find that the transitional disk radius discrepancy holds even if we examine $R_\mathrm{in}$ derived from the HWHM." + A secoud oossibilitv is that Που appears to depend somewhat on ie choice of lines included iu the line composite. with ower-excitation ues vielding larger radi by up to a actor of ~2.," A second possibility is that $R_\mathrm{CO}$ appears to depend somewhat on the choice of lines included in the line composite, with lower-excitation lines yielding larger radii by up to a factor of $\sim2$." +" Since 6/8 trausitional disks iu our survey (GAL Aur. ΠΟ 135311 D. ΠΟ 111569 A. SR 21. TW Ίνα, and UN Tau A) did not have high-excitation cimission ines. their low-excitation lines were analyzed iustead. and this could result i spuriously large radii for these disks."," Since 6/8 transitional disks in our survey (GM Aur, HD 135344 B, HD 141569 A, SR 21, TW Hya, and UX Tau A) did not have high-excitation emission lines, their low-excitation lines were analyzed instead, and this could result in spuriously large radii for these disks." + ILowever. this bias produces ouly up to a factor of two difference in eo. while the observed. difference is an order of magnitude.," However, this bias produces only up to a factor of two difference in $R_\mathrm{CO}$, while the observed difference is an order of magnitude." + Furthermore. the discrepancy is also seeu for DoAr Ll and Lhe 330. which do have lieh-excitatiou emission lines.," Furthermore, the discrepancy is also seen for DoAr 44 and $\alpha$ 330, which do have high-excitation emission lines." + Instead. we suggest that the lack of higl-excitation enmission is a reflection of the relatively larger euüttius radii (and thus lower cutting temperatures) for these disks.," Instead, we suggest that the lack of high-excitation emission is a reflection of the relatively larger emitting radii (and thus lower emitting temperatures) for these disks." + A) phlvsically-imotivated explanation for the radius discrepancy is that Ree iu transitional disks is nof set by sublimation. but rather bv dvuamical truucatiou bv au embedded protoplanet.," A physically-motivated explanation for the radius discrepancy is that $R_\mathrm{CO}$ in transitional disks is not set by sublimation, but rather by dynamical truncation by an embedded protoplanet." + Although ciubedded protoplaucts have been posited as a possible explanation for the inner clearings in transitional disks. they are expected to orbit at radi capable of producing the sharp transition between the optically thin imucr aud optically thick outer disk reeious.," Although embedded protoplanets have been posited as a possible explanation for the inner clearings in transitional disks, they are expected to orbit at radii capable of producing the sharp transition between the optically thin inner and optically thick outer disk regions." + Since the Rew we fiud here is usually significantly smaller than Rian. we would actually need to mvoke dynamical truncation byadditional planets. i smaller orbits.," Since the $R_\mathrm{CO}$ we find here is usually significantly smaller than $R_\mathrm{trans}$ we would actually need to invoke dynamical truncation by planets, in smaller orbits." + In particular. we would predict planets located at radii near ~0.5«Row (Artvinowicz&Lubow1991).," In particular, we would predict planets located at radii near $\sim0.5\times R_\mathrm{CO}$ \citep{Artymowicz94}." +. Tlis is iu line with recent work sugecsting that multiplanet svstenis nay be required to explain transitional disks (Zhuctal.2011:Dodsou-Robinsou&Salvk 2011).," This is in line with recent work suggesting that multiplanet systems may be required to explain transitional disks \citep{Zhu11, Dodson-Robinson11}." +". However. his explanation would ueed to be consistent with he appareut increase in Rew with πιουν]τν,"," However, this explanation would need to be consistent with the apparent increase in $R_\mathrm{CO}$ with luminosity." + This is not unreasonable. as the mass aud location of the xotoplanets may depeud ou the disk audor stellarmass.," This is not unreasonable, as the mass and location of the protoplanets may depend on the disk and/or stellarmass." + Another possible explanation for the radius discrepancy.— auc one qualitatively cousistent— with he inerease of Ree with huunmositv. could be that he dust disk is still truncated at Roy. but that Nàhotodissociatiug UW photous penetrate a finite," Another possible explanation for the radius discrepancy, and one qualitatively consistent with the increase of $R_\mathrm{CO}$ with luminosity, could be that the dust disk is still truncated at $R_\mathrm{sub}$ , but that photodissociating UV photons penetrate a finite" +with different values of the spin aud for (bottom panel) a quasi-I&err black hole with a spin of «=O.LAL and different values of the quadrupolar parameter e.,with different values of the spin and for (bottom panel) a quasi-Kerr black hole with a spin of $a=0.4M$ and different values of the quadrupolar parameter $\epsilon$. + In the case of a Ierr black hole. we find that the asviuuuetiy is neeheible for slowly spiuuiug black holes.," In the case of a Kerr black hole, we find that the asymmetry is negligible for slowly spinning black holes." + In the case of a quasi-Iverr black hole. the asvuuuctry scales lincarly with the quadrupolar parameter e.," In the case of a quasi-Kerr black hole, the asymmetry scales linearly with the quadrupolar parameter $\epsilon$." + We fit the asviunetryv with the expression which is valid for a err with arbitrary spin and for a quasi-lverr black hole with values of the spin 0.0.< aud of the parameter 0.0zxe0.5. respectively.," We fit the asymmetry with the expression which is valid for a Kerr with arbitrary spin and for a quasi-Kerr black hole with values of the spin $0.0\leq a/M \leq0.4$ and of the parameter $0.0\leq\epsilon\leq0.5$, respectively." + The asvuunetry depends strongly on the quadriupolar correction parameter e and the disk inclination / whereas it depends only weakly on the spin as long as the black. hole does not spin rapidly., The asymmetry depends strongly on the quadrupolar correction parameter $\epsilon$ and the disk inclination $i$ whereas it depends only weakly on the spin as long as the black hole does not spin rapidly. + Expression (10)) iuplies that the maxiuuni asviunetrv of a [err black hole is ~ 0.36A7., Expression \ref{asymmetryformula}) ) implies that the maximum asymmetry of a Kerr black hole is $\simeq0.36M$ . + In the event that au asviuuetry larger than this value is detected. which is already. possible for 1iodoerate deviations of the quadrupole moment from the value of a Ikerr quacdaupole moment (c.f.," In the event that an asymmetry larger than this value is detected, which is already possible for moderate deviations of the quadrupole moment from the value of a Kerr quadrupole moment (c.f.," + Figure 8.. bottom panel). then the compact object camnot be a herr black hole.," Figure \ref{asymmetryfits}, bottom panel), then the compact object cannot be a Kerr black hole." + Tn order to test the no-hair theorem. if Is unecessary to measure (at least) three multipole moments of the spacetime (Ryan 1995).," In order to test the no-hair theorem, it is necessary to measure (at least) three multipole moments of the spacetime (Ryan 1995)." + In the following we argue that a measurement of the diameter of the ring deteriuimnes the mass of the ceutral object aud that for a given disk inclination the displacement aud the asvuuuctry of the ving directly measure the spin e& aud the parameter e. respectively.," In the following we argue that a measurement of the diameter of the ring determines the mass of the central object and that for a given disk inclination the displacement and the asymmetry of the ring directly measure the spin $a$ and the parameter $\epsilon$, respectively." + Therefore. the degree of asviuuetry is a direct mmeasure of the violation of the no-hair theorem.," Therefore, the degree of asymmetry is a direct measure of the violation of the no-hair theorem." + Iu Figure 9.. we plot (loft panel) the diameter of the photon ring as a function of the spin e for values of the inclination 177<7x86°.," In Figure \ref{ringstatistics}, we plot (left panel) the diameter of the photon ring as a function of the spin $a$ for values of the inclination $17^\circ\leq i \leq86^\circ$." + Dashed lines correspond to a quasi-I&err black hole with a value of the parameter e—0.5., Dashed lines correspond to a quasi-Kerr black hole with a value of the parameter $\epsilon=0.5$. + The solid line corresponds to a I&err black hole., The solid line corresponds to a Kerr black hole. + Iu all cases. the diameter is practically indepeudent of the spin e.," In all cases, the diameter is practically independent of the spin $a$." + For a I&err black hole. the diameter also depends ouly very weakly ou the disk inclination aud is almost constant for spins iu the ranee 0xαλ<0.1 witha value of Even for large values of the spin. the diameter depends oulv weakly ou the inclination angle causing a svsteniatic nucertaity of only =2% for a value of the spin as large axa=O.OAL.," For a Kerr black hole, the diameter also depends only very weakly on the disk inclination and is almost constant for spins in the range $0\leq a/M \leq0.4$ with a value of Even for large values of the spin, the diameter depends only weakly on the inclination angle causing a systematic uncertainty of only $\simeq2\%$ for a value of the spin as large as $a=0.9M$." + For a quasidI&err black hole. the diameter is affected by the parameter e aud the inclination which leads to a systematic uncertaüntv of c5% if e=0.5.," For a quasi-Kerr black hole, the diameter is affected by the parameter $\epsilon$ and the inclination which leads to a systematic uncertainty of $\simeq5\%$ if $\epsilon=0.5$." + Therefore. a measurement of the rine diameter directly micasures the mass of the central object.," Therefore, a measurement of the ring diameter directly measures the mass of the central object." +" Iu Figure 9.. we also plot (center panel) the ring displacement of a quasiI&orr black hole versus esiu/ for values of the inclination augle 177<7<867"" aud of the piraueter Oxe<0.5."," In Figure \ref{ringstatistics}, we also plot (center panel) the ring displacement of a quasi-Kerr black hole versus $a\sin i$ for values of the inclination angle $17^\circ\leq i \leq86^\circ$ and of the parameter $0\leq\epsilon\leq0.5$." + The solid lines in this figure represent the extreme lanits of the ft function eiven by expression (9)) corresponding to €=0 and e—0.5. siu;= l. respectively.," The solid lines in this figure represent the extreme limits of the fit function given by expression \ref{displacementformula}) ) corresponding to $\epsilon=0$ and $\epsilon=0.5$, $\sin i=1$ , respectively." + Finally. in Figure 9. (aight panel). we show the nue asvuunetrv of quasi-err black holes versus esinD; for values of the spin 0xαλ0.1 aud of the inclination angle 177<7x867.," Finally, in Figure \ref{ringstatistics} (right panel), we show the ring asymmetry of quasi-Kerr black holes versus $\epsilon\sin^{3/2} i$ for values of the spin $0\leq a/M \leq0.4$ and of the inclination angle $17^\circ\leq i \leq86^\circ$." + The solid line is expression (10)) evaluated at a=0.0AL., The solid line is expression \ref{asymmetryformula}) ) evaluated at $a=0.0M$. +" ""Thus. for a given disk inclination. au observation of the displacement aud the asvaiunietzy directly leads to a mcasurement of the spin & and quadrupolar correction paraueter € via the expressious m equations (9)) aud (10)). respectively."," Thus, for a given disk inclination, an observation of the displacement and the asymmetry directly leads to a measurement of the spin $a$ and quadrupolar correction parameter $\epsilon$ via the expressions in equations \ref{displacementformula}) ) and \ref{asymmetryformula}) ), respectively." + A micasurement of the asvuunetry allows for a test of the no-hair theorem., A measurement of the asymmetry allows for a test of the no-hair theorem. + Iu order to determine the displacement it is necessary to nieasure the ceuter of mass of the accretion flow around the black hole., In order to determine the displacement it is necessary to measure the center of mass of the accretion flow around the black hole. + An estimate of the location of the center of mass ay perhaps be obtained from a (time-averaged) πμασο of the outer edges of the accretion flow., An estimate of the location of the center of mass may perhaps be obtained from a (time-averaged) image of the outer edges of the accretion flow. + Iu Paper L we proposed a new framework for testing the no-hair theorem with observatious of black holes iu the electromagnetic spectrum.," In Paper I, we proposed a new framework for testing the no-hair theorem with observations of black holes in the electromagnetic spectrum." + We formulated our tests based on a quasi-lverr metre (Cdampedakis Babak 2006). which deviates smoothly from the err metric in the quadrupole moment.," We formulated our tests based on a quasi-Kerr metric (Glampedakis Babak 2006), which deviates smoothly from the Kerr metric in the quadrupole moment." + Since the no-hair theorem adits exactly two independent iultipole momeuts for a black hole. a 1neasureiment of these three moments will allow us to test the mo-hair theorem.," Since the no-hair theorem admits exactly two independent multipole moments for a black hole, a measurement of these three moments will allow us to test the no-hair theorem." + Tn this paper. we calculated umuuericallv the mapping vetween different locations in the accretion flow around a quasi-I&orr black hole aud positions in the image plane of a distant observer.," In this paper, we calculated numerically the mapping between different locations in the accretion flow around a quasi-Kerr black hole and positions in the image plane of a distant observer." + Our calculations allowed us to study he potential of using Waging observations ofblack holes. hat will become available iu the near future. in order to est the no-hiür theorem.," Our calculations allowed us to study the potential of using imaging observations of black holes, that will become available in the near future, in order to test the no-hair theorem." + We areucd that the expected image of an accretion How will be characterized by a bright cmission ring eenerated by elt ravs that circle iiultiple times around he event horizon before ciucreie towards the observer., We argued that the expected image of an accretion flow will be characterized by a bright emission ring generated by light rays that circle multiple times around the event horizon before emerging towards the observer. + We identified the ring diameter as a direct 1ieasure of the nass of the black hole. aud we quantified the dependence of the displacement aud the asvuuuetry of the rine on the spin and the quadrupolar parameter as well as ou the disk inclination.," We identified the ring diameter as a direct measure of the mass of the black hole, and we quantified the dependence of the displacement and the asymmetry of the ring on the spin and the quadrupolar parameter as well as on the disk inclination." +" For a eiven inclination angle. a measurement of the displacement aud the asvuunetry directly 1ieasures the spin aud the quadrupolar parameter of the system. respectively,"," For a given inclination angle, a measurement of the displacement and the asymmetry directly measures the spin and the quadrupolar parameter of the system, respectively." + The asvunuetry itself provides a direct measure ofthe violation of the no-hair theorem., The asymmetry itself provides a direct measure of the violation of the no-hair theorem. + It is iurportant to emphasize here that only the relative displacement aud asvnunetiy of the vine G.c.. 1mieasured in units of the ving diuneter) aud not then absolute values are necessary m inferring the spin aud quadrupole moment of the black-hole spacetime.," It is important to emphasize here that only the relative displacement and asymmetry of the ring (i.e., measured in units of the ring diameter) and not their absolute values are necessary in inferring the spin and quadrupole moment of the black-hole spacetime." +" As a result. the outcome of such an observation docs not depeud on the distance to the black hole. which might uot be kuown accurately,"," As a result, the outcome of such an observation does not depend on the distance to the black hole, which might not be known accurately." + On the other hand. the aneular diameter of the photou ring is proportional to the mass of the black hole.," On the other hand, the angular diameter of the photon ring is proportional to the mass of the black hole." + This can lead to au accurate measurement of the mass of the black hole if the distance is known. or else of the distance to the black hole. ifits mass is known fron. e.g.. dvnamical observations.," This can lead to an accurate measurement of the mass of the black hole if the distance is known, or else of the distance to the black hole, if its mass is known from, e.g., dynamical observations." + Ser A*. the black hole iu the center of the Milky Way. is the ideal candidate for a test of the no-hair theorem due toits high briglhtuess. large augular size. aud relatively uuinupeded observational accessibility.," Sgr A*, the black hole in the center of the Milky Way, is the ideal candidate for a test of the no-hair theorem due to its high brightness, large angular size, and relatively unimpeded observational accessibility." + Receut VLBI observations (Docleman et al., Recent VLBI observations (Doeleman et al. + 2008) resolved. Ser Α ou horizon scales., 2008) resolved Sgr A* on horizon scales. + Incorporating additional baselines to the VEDI network willlead to the first nuages of, Incorporating additional baselines to the VLBI network willlead to the first images of +(McAlary et al.,(McAlary et al. + 1983) than in 2009-2011 allows us to achieve a linear least-squares fit to the data points with sufficient accuracy., 1983) than in 2009-2011 allows us to achieve a linear least-squares fit to the data points with sufficient accuracy. + The fit yields an AGN slope Lj=1.20€0.04., The fit yields an AGN slope $\Gamma_{BV} = 1.20 \pm 0.04$. + This slope is at the steep end. but within the range determined by Winkler et al. (," This slope is at the steep end, but within the range determined by Winkler et al. (" +1992) for other type-1 AGN.,1992) for other type-1 AGN. + To check the slope of PGOOO03--199 further. we searched in the NASA Extragalactic Database (NED) and found two more B and V data points taken with comparable apertures. one with 1577 by McAlary et al. (," To check the slope of PG0003+199 further, we searched in the NASA Extragalactic Database (NED) and found two more $B$ and $V$ data points taken with comparable apertures, one with $\farcs$ 7 by McAlary et al. (" +1983) and one with 140 by Elvis et al. (,1983) and one with $\farcs$ 0 by Elvis et al. ( +1994).,1994). + These (extinction corrected) data points are displaced to the right of the 778 slope in Fig. 6..," These (extinction corrected) data points are displaced to the right of the $\farcs$ 8 slope in Fig. \ref{fig_fvg_pg0003}," + because they contain additional host emission from the extended disk which is redder than the AGN emission., because they contain additional host emission from the extended disk which is redder than the AGN emission. + These two data points yield a slope Prey=1.25+0.04., These two data points yield a slope $\Gamma_{BV} = 1.25 \pm 0.04$. +" The consistency of Γρ derived from two independent data sets corrobborates the steep AGN slope of PG0003-4199,", The consistency of $\Gamma_{BV}$ derived from two independent data sets corrobborates the steep AGN slope of PG0003+199. + The range of host slopes (0.4=12 and earlier very strong filaments surround the central halo and. appear less ragmentecl into individual smaller objects than at. [ater ines., At $z=12$ and earlier very strong filaments surround the central halo and appear less fragmented into individual smaller objects than at later times. + The right-hand column of Fig., The right-hand column of Fig. + 3. illustrates the larger scale environment of these main haloes., \ref{fig:fig3} illustrates the larger scale environment of these main haloes. +" The images here rave side ους, are projections of a region of depth ου. and again all represent the projected. overdensity using an identical colour scale."," The images here have side $10r_{200}$, are projections of a region of depth $2r_{200}$, and again all represent the projected overdensity using an identical colour scale." + Phe qualitative change in morphology with redshift is more dramatic in these plots., The qualitative change in morphology with redshift is more dramatic in these plots. + As we go back in time the main halo becomes less dominant and less regular. the “filamentary” structures become heavier anc smoother. ancl substructure both within and outside the main object becomes less evident.," As we go back in time the main halo becomes less dominant and less regular, the “filamentary” structures become heavier and smoother, and substructure both within and outside the main object becomes less evident." + At the earliest times. the filaments extend over the full region plotted and are the mos striking structure within it.," At the earliest times, the filaments extend over the full region plotted and are the most striking structure within it." + Note that with the exception of the latest pair (2=5 and z= 0) all the matter shown in each carly image is contained within the central object of the next later one., Note that with the exception of the latest pair $z=5$ and $z=0$ ) all the matter shown in each early image is contained within the central object of the next later one. + The rapid growth in mass of our principal halo was already shown in the top panel of Fig. 1., The rapid growth in mass of our principal halo was already shown in the top panel of Fig. \ref{fig:fig1}. +" At each time we can estimate a representative temperature for the eas associated with this halo from its maximumum circular velocity. Vi: where yam, is the mean molecular weight of the gasand Aw is Boltzmann's constant."," At each time we can estimate a representative temperature for the gas associated with this halo from its maximumum circular velocity, $V_c$: where $\mu m_p$ is the mean molecular weight of the gasand $k_B$ is Boltzmann's constant." + This “virial” temperature is plotted as a function of redshift in Fig. 4.., This “virial” temperature is plotted as a function of redshift in Fig. \ref{fig:fig4}. +" We assume [i1.22 from the earliest. redshift’ until the temperature has risen to 10000Ix. At this time we assume that the gas is collisionally ionised. and has jp=0.59 thereafter: during the transition the halo temperature. is taken to be constant at Ty,=100001. Like the mass. the halo temperature increases very rapidly at early. times."," We assume $\mu=1.22$ from the earliest redshift until the temperature has risen to $10000$ K. At this time we assume that the gas is collisionally ionised and has $\mu=0.59$ thereafter; during the transition the halo temperature is taken to be constant at ${\rm T}_{\rm +Vir}=10000$ K. Like the mass, the halo temperature increases very rapidly at early times." + Already by z=49. it has reached —2000K when. for the hieh densities predicted. at these redshifts. molecular hydrogen should form sullicientIy rapidly to provide efficient radiative cooling.," Already by $z=49$, it has reached $\sim 2000$ K when, for the high densities predicted at these redshifts, molecular hydrogen should form sufficiently rapidly to provide efficient radiative cooling." + As many recent simulations have shown. this is likely to lead to the formation of a massive star at the centre of our object (Abel et al 1998. 2002: Bromm e al.," As many recent simulations have shown, this is likely to lead to the formation of a massive star at the centre of our object (Abel et al 1998, 2002; Bromm et al." + 2002: Yoshida et al., 2002; Yoshida et al. + 2003)., 2003). + The critical temperature for collisional ionisation Z—1011 is reached by z=39 and the greatly increased efficiency. of cooling at this poin could lead to a starburst and the formation of a minigalaxy., The critical temperature for collisional ionisation $T\sim 10^4$ K is reached by $z=39$ and the greatly increased efficiency of cooling at this point could lead to a starburst and the formation of a minigalaxy. + Whether this actually happens depends on the evolution of the first star., Whether this actually happens depends on the evolution of the first star. + Hs ionising radiation and the strong shock i produces if it goes supernova could: already have alfectec surrounding eas over a large region by 2=39 (c.g. Yoshida. Bromm &lIIHLIerneuist 2004).," Its ionising radiation and the strong shock it produces if it goes supernova could already have affected surrounding gas over a large region by $z=39$ (e.g. Yoshida, Bromm Hernquist 2004)." + We studs these issues as wel as others associated with the formation of the first stars in a companion paper (Iteed et al., We study these issues as well as others associated with the formation of the first stars in a companion paper (Reed et al. + 2005c), 2005c). +" Simulations of structure formation in hierarchical clustering cosmogonies (including the concordance ACDAL model) produce dark matter haloes with density. profiles which can be fit bv the simple formula proposed by Navarro. Frenk White (1996. 1907: NEW). Llere ry is à characteristic radius where the profile declines locally as r£ and p,/4 is the density at ry."," Simulations of structure formation in hierarchical clustering cosmogonies (including the concordance $\Lambda$ CDM model) produce dark matter haloes with density profiles which can be fit by the simple formula proposed by Navarro, Frenk White (1996, 1997; NFW), Here $r_s$ is a characteristic radius where the profile declines locally as $r^{-2}$, and $\rho_s/4$ is the density at $r_s$." + At the relatively low redshifts where this formula. has been tested. there is a strong correlation between ες and ps which depends on redshift. on the global cosmological paranicters. and on the initial power spectrum.," At the relatively low redshifts where this formula has been tested, there is a strong correlation between $r_s$ and $\rho_s$ which depends on redshift, on the global cosmological parameters, and on the initial power spectrum." + The larger the mass of a halo. the lower its characteristic density. reflecting the lower density of the universe at the (later) time when more nissive systems were typically assembled.," The larger the mass of a halo, the lower its characteristic density, reflecting the lower density of the universe at the (later) time when more massive systems were typically assembled." + The applicability of this formula in the innermost regions of haloes has been controversial. but recent work suggests that it is a reasonable fit to most haloes over the racial range 0.011900κκο ro2oot," The applicability of this formula in the innermost regions of haloes has been controversial, but recent work suggests that it is a reasonable fit to most haloes over the radial range $0.01r_{200}s80°.," Figure \ref{dmbdt}) ) shows that the subsolar, thermal equilibrium sublimation flux for $A$ = 0 is about $^{-5}$ kg $^{-2}$ $^{-1}$, but this falls by $\sim$ 4.5 orders of magnitude as the albedo rises to $A$ = 0.8 and also drops precipitously at $\theta >$ $\degr$." +" The stability of the ice evidently depends very stronglv on whether it is “clean” (hieh albedo) or ""dirty (low albedo) ancl on the position of the ice relative to (he subsolar point (Figure 3)).", The stability of the ice evidently depends very strongly on whether it is “clean” (high albedo) or “dirty” (low albedo) and on the position of the ice relative to the subsolar point (Figure \ref{dmbdt}) ). + The difference between dirty ancl clean ice can. in (his context. be the result of the admixture of very small mass fractions of absorbing material.," The difference between dirty and clean ice can, in this context, be the result of the admixture of very small mass fractions of absorbing material." + For example. < contamination by mass can. depending on the grain size and composition of the contaminant. depress the optical albedo of ice to $ 8 Gyr) with disc-like kinematics and solar-like abundance exist in significant proportion, then they may represent the signature that the disc, or part of the (old) disc has formed through an evolutionary path that differs from the SCB model evolution." + ‘This stucly was largely inspired by papers by M. Grenon. and his repeated cautionary advice that most local surveys are biased against metal-rich stars.," This study was largely inspired by papers by M. Grenon, and his repeated cautionary advice that most local surveys are biased against metal-rich stars." + | thank the referee Bernard Pagel for his many helpful comments that. improved. this paper., I thank the referee Bernard Pagel for his many helpful comments that improved this paper. + This work has mace extensive use of the Simbad database operated. at CDS. Strasbourg. France.," This work has made extensive use of the Simbad database operated at CDS, Strasbourg, France." + is based on data from ESA astrometric satellite 4/ppercos.,It is based on data from ESA astrometric satellite $Hipparcos$ . +The purpose of this research note 15 to show that the mean polarization angle in the near-infrared was almost constant from 2003-2006. and to infer the position angle of Ser Α from this preferred direction in terms of the orbiting spot model.,"The purpose of this research note is to show that the mean polarization angle in the near-infrared was almost constant from 2003-2006, and to infer the position angle of Sgr A* from this preferred direction in terms of the orbiting spot model." + For a high inclination angle the 3c range for the position angle is derived to be ~60°—1087. thereby supporting the findings of Markoff et al. (2007)).," For a high inclination angle the $3\sigma$ range for the position angle is derived to be $\sim 60\degr - 108\degr$, thereby supporting the findings of Markoff et al. \cite{sera}) )," + and -- less likely — Muzie et al. (2007)), and – less likely – Muzic et al. \cite{kora}) ) + who arrived at position angles of ~105° (see also Muno et al. 2007)), who arrived at position angles of $\sim105\degr$ (see also Muno et al. \cite{muno}) ) + and ~60°. respectively. with complementary methods.," and $\sim60\degr$, respectively, with complementary methods." + However. there are severe caveats.," However, there are severe caveats." + The unknown magnetic field structure. the uncertain foreground polarization. and the lack of a clear understanding of the detailed hydrodynamics of the blob challenge our conclusions.," The unknown magnetic field structure, the uncertain foreground polarization, and the lack of a clear understanding of the detailed hydrodynamics of the blob challenge our conclusions." +for the dark matter.,for the dark matter. + Galaxies live in DM halos with circular velocities of 160km/s.," Galaxies live in DM halos with circular velocities of $160 \, {\rm km/s}$." +" All these experiments are set to have the same elliptical orbits, according to Toomre& (1972),, with pericentre distance of e20 kpc and circularity parameter e~0.20."," All these experiments are set to have the same elliptical orbits, according to \citet{TT72}, with pericentre distance of $\approx 20 $ kpc and circularity parameter $\epsilon \sim 0.20$." + These orbital values are consistent with those obtained from the analysis of cosmological simulations (e.g.Khochfar&Burkert2006)., These orbital values are consistent with those obtained from the analysis of cosmological simulations \citep[e.g.][]{kb06}. +. We have kept them fixed so that changes in the chemical and dynamical properties can only be adscribed to the different angular momentum orientation., We have kept them fixed so that changes in the chemical and dynamical properties can only be adscribed to the different angular momentum orientation. + A larger orbital distribution might be of interest to analyse possible dependences in the future., A larger orbital distribution might be of interest to analyse possible dependences in the future. +" The gas component is initially assumed to have been pre-enriched by the existing stars, following a metallicity gradient statistically consistent with observations."," The gas component is initially assumed to have been pre-enriched by the existing stars, following a metallicity gradient statistically consistent with observations." +" We assigned an initial amount of chemical elements to the gas particles in order to reproduce the mean observational gradient reported by Dutil&Roy(1999),, which for the less gas-rich simulations is —0.08 dex kpc! with a central oxygen abundance of 9.2."," We assigned an initial amount of chemical elements to the gas particles in order to reproduce the mean observational gradient reported by \citet{dutil99}, which for the less gas-rich simulations is $-0.08$ dex $^{-1}$ with a central oxygen abundance of 9.2." +" Thus, the initially simulated galaxies are located onto the observed mass-metallicity relation (Tremontietal.2004)."," Thus, the initially simulated galaxies are located onto the observed mass-metallicity relation \citep{tremonti04}." +". For the more gas-rich simulation (SimVI), the initial metallicity gradient is somewhat steeper (—0.1 dex kpc!) and the central metallicity lower (by 0.7 dex) than the less gas-rich galaxies."," For the more gas-rich simulation (SimVI), the initial metallicity gradient is somewhat steeper $-0.1$ dex $^{-1}$ ) and the central metallicity lower (by 0.7 dex) than the less gas-rich galaxies." +" By doing that, these more gas-rich simulated galaxies are then initially onto the mass-metallicity relation observed at high redshift by Maiolinoetal.(2008)."," By doing that, these more gas-rich simulated galaxies are then initially onto the mass-metallicity relation observed at high redshift by \citet{maiolino08}." +. Note that the exact adopted values are not critical since we are interested in the relative evolution of the metallicity., Note that the exact adopted values are not critical since we are interested in the relative evolution of the metallicity. +" In this section, we analyse the effects of galaxy interactions on the central metallicity in wet local encounters."," In this section, we analyse the effects of galaxy interactions on the central metallicity in wet local encounters." +" Taking SimI as an example, we show in Fig."," Taking SimI as an example, we show in Fig." + 1 the evolution of the gas-phase O/H abundance of one member of the galaxy as the interaction proceeds., \ref{leo1} the evolution of the gas-phase O/H abundance of one member of the galaxy as the interaction proceeds. +" The evolution of the interaction can be followed by the relative distances between the center of masses of the interacting systems, plotted in the same figure."," The evolution of the interaction can be followed by the relative distances between the center of masses of the interacting systems, plotted in the same figure." +" As it can be seen, there is a clear evidence of a decrease in the central gas-phase metallicity which can be correlated with the first and the second pericentre."," As it can be seen, there is a clear evidence of a decrease in the central gas-phase metallicity which can be correlated with the first and the second pericentre." +" Several previous works have shown that during a close encounter, strong tidal torques can develop, re-distributing angular momentum and producing gas infall (e.g.Barnes&Hernquist1996;MihosTissera2000;DiMatteoetal."," Several previous works have shown that during a close encounter, strong tidal torques can develop, re-distributing angular momentum and producing gas infall \citep[e.g.][]{BH96,MH96,tissera00,dimatteo08}." +" 2008)..In the case that a metallicity gradient exists, the infalling gas will tend to have lower metallicity than that of baryons in the central region producing a dilution as it is clearly shown by our simulation (seealsoRupkeetal.2010a;Montuori2010)."," .In the case that a metallicity gradient exists, the infalling gas will tend to have lower metallicity than that of baryons in the central region producing a dilution as it is clearly shown by our simulation \citep[see also][]{rupke10a,montuori10}." +". In order to quantify the chemical changes as a function of time, we compute the central abundances inside concentric spheres of different radius in the range from 1 kpc to 10 kpc."," In order to quantify the chemical changes as a function of time, we compute the central abundances inside concentric spheres of different radius in the range from 1 kpc to 10 kpc." +" As an example, Fig."," As an example, Fig." +" 1 shows the gas-phase oxygen abundance measured inside a sphere of 2 kpc, 4 kpc, and 6 kpc of radius."," \ref{leo1} shows the gas-phase oxygen abundance measured inside a sphere of 2 kpc, 4 kpc, and 6 kpc of radius." +" As it can be appreciated, as one moves inwards, the signals get noisier since the number of particle decreases."," As it can be appreciated, as one moves inwards, the signals get noisier since the number of particle decreases." +" If one moves to larger radius, the signal decreases as chemical abundances start to be dominated by the external regions."," If one moves to larger radius, the signal decreases as chemical abundances start to be dominated by the external regions." +" As a compromised between these two effects, we select Reen=4 kpc (blue line) to define the central region where we will perform the analysis of central metallicities."," As a compromised between these two effects, we select $R_{\rm cen}=4$ kpc (blue line) to define the central region where we will perform the analysis of central metallicities." +" In order to understand the decrease in the central abundances shown by Fig. 1,,"," In order to understand the decrease in the central abundances shown by Fig. \ref{leo1}," + we analyse the evolution of the central gas inflows and their abundances., we analyse the evolution of the central gas inflows and their abundances. + We compute the central gas inflow at a certain time as the ratio between the new accreted gas mass within [ίσοι and the total gas content within the same radius., We compute the central gas inflow at a certain time as the ratio between the new accreted gas mass within $R_{\rm cen}$ and the total gas content within the same radius. + Fig., Fig. + 2aa shows the evolution of this central gas inflow as the interacting galaxies approach each other., \ref{general}a a shows the evolution of this central gas inflow as the interacting galaxies approach each other. + We find that the gas inflow remains constant in around 296 until the first close passage when it rises up to ~8%., We find that the gas inflow remains constant in around $2\%$ until the first close passage when it rises up to $\sim 8\%$. +" From the second pericentre, the gas inflow steadily increases up to around 3096."," From the second pericentre, the gas inflow steadily increases up to around $30\%$." + This gas inflow is dominated by low-metallicity material as it can be appreciated from Fig., This gas inflow is dominated by low-metallicity material as it can be appreciated from Fig. + 2bb (solid line) where we plot its O/H abundance., \ref{general}b b (solid line) where we plot its O/H abundance. + The accretion of this low-metallicity gas dilutes the central oxygen abundance as shown by the continuous decrease detected after the first pericentre (Fig., The accretion of this low-metallicity gas dilutes the central oxygen abundance as shown by the continuous decrease detected after the first pericentre (Fig. +" 2bb, dotted line)."," \ref{general}b b, dotted line)." + 'The main drops in central abundances correlate with the occurrence of the pericentres driven by important gas inflows, The main drops in central abundances correlate with the occurrence of the pericentres driven by important gas inflows + ~ M The evidence for this pattern is based not only on RM observations (Winn et al., $\gsim$ $M_\odot$ The evidence for this pattern is based not only on RM observations (Winn et al. + 2010a) but also on the line-of-sight stellar rotation. velocities of transit hosts (Schlaufman 2010)., 2010a) but also on the line-of-sight stellar rotation velocities of transit hosts (Schlaufman 2010). + This trend may indicate that. planet formation and migration are fundamentally different for low-mass stars than for high-mass stars. for which there is already evidence in the distributions of planet mass and period. (Bowler et al.," This trend may indicate that planet formation and migration are fundamentally different for low-mass stars than for high-mass stars, for which there is already evidence in the distributions of planet mass and period (Bowler et al." + 2010)., 2010). + Another possibility is. that the formation and migration processes are similar. but that the subsequent tidal evolution is different (Winn et al.," Another possibility is that the formation and migration processes are similar, but that the subsequent tidal evolution is different (Winn et al." + 20102)., 2010a). + In this hypothesis. cool stars are observed to have low obliquities because tidal evolution drove them into alignment. while hot stars retain their “primordial” obliquities because of their thinner (or absent) outer convection zones and consequently slower rates of tidal dissipation.," In this hypothesis, cool stars are observed to have low obliquities because tidal evolution drove them into alignment, while hot stars retain their “primordial” obliquities because of their thinner (or absent) outer convection zones and consequently slower rates of tidal dissipation." + Although the trend seems clear. it is difficult to assess its true significance because many possible variables were examined before alighting on stellar temperature and mass.," Although the trend seems clear, it is difficult to assess its true significance because many possible variables were examined before alighting on stellar temperature and mass." + The only way to be sure is to gather more data., The only way to be sure is to gather more data. + This paper presents results for the next two systems we observed after the trend had been identified., This paper presents results for the next two systems we observed after the trend had been identified. +" Both systems have short-period giant planets. but HAT-P-4 is ""cool"" (7yy=5860+80 K: Kovaces et al."," Both systems have short-period giant planets, but HAT-P-4 is “cool” $T_{\rm eff} += 5860\pm 80$ K; Kováccs et al." + 2007) while HAT-P-14 15 “hot” (6600+90 K: Torres et al., 2007) while HAT-P-14 is “hot” $6600\pm 90$ K; Torres et al. + 2010)., 2010). + We present the observations of these systems in 2. the analysis and results in 3. and a discussion in 4.," We present the observations of these systems in 2, the analysis and results in 3, and a discussion in 4." + Our spectroscopic observations employed the High Resolution Spectrograph (HIRES: Vogt et al., Our spectroscopic observations employed the High Resolution Spectrograph (HIRES; Vogt et al. + 1994) of the Keck I 10m telescope. on Mauna Kea. Hawarn.," 1994) of the Keck I 10m telescope, on Mauna Kea, Hawaii." + We gathered 35 spectra of HAT-P-4 on the night of 2010 March 29/30. and 44 spectra of HAT-P-14 on the night of 2010 April 27/28. in both cases spanning a predicted transit of the planet.," We gathered 35 spectra of HAT-P-4 on the night of 2010 March 29/30, and 44 spectra of HAT-P-14 on the night of 2010 April 27/28, in both cases spanning a predicted transit of the planet." + An additional 14. spectra of HAT-P-4 were gathered on other nights. at essentially random orbital phases.," An additional 14 spectra of HAT-P-4 were gathered on other nights, at essentially random orbital phases." + We used the standardinstrument settings and observing, We used the standardinstrument settings and observing +barely degrades detection sensitivity. for smearing widths up to 0.05 periods.,"barely degrades detection sensitivity, for smearing widths up to $0.05$ periods." + Once detected. gamma-photon statistics drive the need for higher precision: the timing residuals should be smaller than the phase histogram bin width. which in turn should be wide enough to have at least several gamma-photons per bin.," Once detected, gamma-photon statistics drive the need for higher precision: the timing residuals should be smaller than the phase histogram bin width, which in turn should be wide enough to have at least several gamma-photons per bin." + A consequence of these relatively modest timing requirements is that a given radio observation need only last the minimum time for detection., A consequence of these relatively modest timing requirements is that a given radio observation need only last the minimum time for detection. + More crucial is the number of observations per year. which depends on the timing noise. correlated with v and ¥ and thus E (??)..," More crucial is the number of observations per year, which depends on the timing noise, correlated with $\nu$ and $\dot\nu$ and thus $\dot{E}$ \citep{CordesHelfand,ZA94}." + Gamma-ray candidates tend to be the noisiest pulsars., Gamma-ray candidates tend to be the noisiest pulsars. + Ilustrations of timing noise in young. high £pulsars can be found. e.g.. in ?..," Illustrations of timing noise in young, high $\dot E$pulsars can be found, e.g., in \citet{GEM06}." + Glitches have been observed for roughlya quarter of the pulsars being monitored (2).., Glitches have been observed for roughlya quarter of the pulsars being monitored \citep{MPW08}. + The bulk of the pulsars in this campaign are observed monthly. and a smaller number are observed weekly or bi-weekly.," The bulk of the pulsars in this campaign are observed monthly, and a smaller number are observed weekly or bi-weekly." + The low radio fluxes of some gamma-ray pulsar candidates require long exposures on the biggest radio telescopes., The low radio fluxes of some gamma-ray pulsar candidates require long exposures on the biggest radio telescopes. + We must devote time to these as they could be bright gamma-ray sources., We must devote time to these as they could be bright gamma-ray sources. + Radio-faint. particularly noisy pulsars could dominate the observation schedules.," Radio-faint, particularly noisy pulsars could dominate the observation schedules." + Radio signals are dispersed by the interstellar medium. with a frequency dependent delay causing signals at high radio frequencies to arrive before those at low radio frequencies.," Radio signals are dispersed by the interstellar medium, with a frequency dependent delay causing signals at high radio frequencies to arrive before those at low radio frequencies." + The pulsar Dispersion Measure (DM). or the integrated column density of free electrons along the line of sight from a pulsar to Earth. usually measured in em? pe. allows extrapolation of the photon arrival times from radio to infinite frequency. as Is required for gamma-ray studies.," The pulsar Dispersion Measure (DM), or the integrated column density of free electrons along the line of sight from a pulsar to Earth, usually measured in $^{-3}$ pc, allows extrapolation of the photon arrival times from radio to infinite frequency, as is required for gamma-ray studies." + The DM. however. can change over timescales of weeks to years (?)..," The DM, however, can change over timescales of weeks to years \citep{You2007}. ." + If the DM is Inaccurate. then the reference phase @ from the radio ephemeris (described below) will change. causing an apparent drift in the gamma-ray absolute phase and a smearing of the resulting gamma-ray pulse profiles.," If the DM is inaccurate, then the reference phase $\Phi_0$ from the radio ephemeris (described below) will change, causing an apparent drift in the gamma-ray absolute phase and a smearing of the resulting gamma-ray pulse profiles." + Such smearing would compromise the multi-wavelength phase comparisons upon which beam geometry studies are based., Such smearing would compromise the multi-wavelength phase comparisons upon which beam geometry studies are based. + Therefore the timing campaign must include occasional monitoring at multiple radio frequencies., Therefore the timing campaign must include occasional monitoring at multiple radio frequencies. + Figure 2. 1s one illustration of the magnitude of the dispersion for different radio frequencies., Figure \ref{residuals} is one illustration of the magnitude of the dispersion for different radio frequencies. + Another illustration of the potential effect of DM changes on a gamma-ray light curve is obtained using the DM values from the Jodrell Bank monthly ephemerides., Another illustration of the potential effect of DM changes on a gamma-ray light curve is obtained using the DM values from the Jodrell Bank monthly ephemerides. +" Over the years of the Compton GRO mission (1991-1999), the maximum excursio in the photon time extrapolated from the radio frequency of 1400. MHz to infinite frequency is 0.3 ms (1% of a rotation of the neutron star)."," Over the years of the Compton GRO mission (1991-1999), the maximum excursion in the photon time extrapolated from the radio frequency of 1400 MHz to infinite frequency is $0.3$ ms $1$ of a rotation of the neutron star)." + For reference. the total DM correction from radio to gamma ray is ~120 ms. which is 4+ rotations.," For reference, the total DM correction from radio to gamma ray is $\sim 120$ ms, which is 4 rotations." + For pulsars faster than theCrab.. the effect could be larger.," For pulsars faster than the, the effect could be larger." + For most pulsars. the effect is minor.," For most pulsars, the effect is minor." + Turbulence in the 1terstellar medium also induces frequency-dependent scattering and refraction of the pulsar signal. due to path-length differences.," Turbulence in the interstellar medium also induces frequency-dependent scattering and refraction of the pulsar signal, due to path-length differences." + Simulations show that those effects are in the order of hundreds of nanoseconds for observations at 1.4 GHz (?).. and hence negligible for gamma-ray astronomy.," Simulations show that those effects are in the order of hundreds of nanoseconds for observations at 1.4 GHz \citep{FC90}, and hence negligible for gamma-ray astronomy." + The radio pulsar monitoring must be sustained throughout the duration of the mission (1.8. for 5 to 10 years). a strain for any observatory. so other contributions are welcome.," The radio pulsar monitoring must be sustained throughout the duration of the mission (i.e. for 5 to 10 years), a strain for any observatory, so other contributions are welcome." + In particular. very frequent monitoring of high £. large S400 pulsars could allow significant contributions to LAT science by smaller radiotelescopes.," In particular, very frequent monitoring of high $\dot{E}$, large $S_{1400}$ pulsars could allow significant contributions to LAT science by smaller radiotelescopes." + The archetypical radio-quiet gamma-ray pulsar is Geminga.. J0633+1746..," The archetypical radio-quiet gamma-ray pulsar is , ." + Biannual XMM satellite measurements, Biannual XMM satellite measurements +"to Pmz2 (as the zero net flux transport vanishes for lower Prandtl number values), and should be somewhat dependent on Pm.","to $Pm \gtrsim 2$ (as the zero net flux transport vanishes for lower Prandtl number values), and should be somewhat dependent on $Pm$." +" As discussed later on, for our weakest field (8= 10*) and smallest Prandtl (Pm= 4), the flow also tends to become bidimensional for low enough Reynolds number."," As discussed later on, for our weakest field $\beta = 10^4$ ) and smallest Prandtl $Pm=4$ ), the flow also tends to become bidimensional for low enough Reynolds number." + Figure 3 displays the behavior of the turbulent transport for the runs we have performed at 8=1000., Figure \ref{ttransp} displays the behavior of the turbulent transport for the runs we have performed at $\beta=1000$. +" The simulations are typically run for 500 shear times, and the transport average is based on the last 400."," The simulations are typically run for $500$ shear times, and the transport average is based on the last $400$." +" In our horizontally extended simulation boxes, the transport fluctuations are substantially reduced with respect to narrower boxes."," In our horizontally extended simulation boxes, the transport fluctuations are substantially reduced with respect to narrower boxes." +" Typically, fluctuations of a factor of 2 are observed, whereas in r:z=1 boxes, fluctuations of an order of magnitude or more are common."," Typically, fluctuations of a factor of $\sim2$ are observed, whereas in $r:z = +1:1$ boxes, fluctuations of an order of magnitude or more are common." +" This reduction provides more precise averages for the transport, even though the averaging time issomewhat smaller than what is used for boxes with aR:Z=1 aspect ratio."," This reduction provides more precise averages for the transport, even though the averaging time issomewhat smaller than what is used for boxes with a $R:Z=1:1$ aspect ratio." +" This aspect ratio dependence is related to the prominence of channel modes in the large transport bursts observed in the narrow boxes (?),, a feature related to the fact that width of the narrow box is comparable to the correlation length of turbulent fluctuations in the horizontal direction (?).."," This aspect ratio dependence is related to the prominence of channel modes in the large transport bursts observed in the narrow boxes \citep{BMCRF08}, a feature related to the fact that width of the narrow box is comparable to the correlation length of turbulent fluctuations in the horizontal direction \citep{GGSJ09}." + The transport time history can be used to quantify the error in the determination of a in the following way., The transport time history can be used to quantify the error in the determination of $\alpha$ in the following way. +" From the raw time history, one can define a series of binned time histories, with a binning time τ ranging from 1 to 100 sheartimes."," From the raw time history, one can define a series of binned time histories, with a binning time $\tau$ ranging from $1$ to $100$ shear." +". From these binned time histories, onecan define a transport standard deviation in the usual way, i.e., o"",=[Y(o;a)?/N]!? where Ν is the number of bins, a; itself being the average value of the transport in bin i."," From these binned time histories, onecan define a transport standard deviation in the usual way, i.e., $\sigma_\alpha=[\sum(\alpha_i-\overline\alpha)^2/N]^{1/2}$ where $N$ is the number of bins, $\alpha_i$ itself being the average value of the transport in bin $i$." + The resulting dependence σα(τ) is shown on Fig., The resulting dependence $\sigma_\alpha(\tau)$ is shown on Fig. + 4 for oneof our runs., \ref{alphadev} for oneof our runs. + Two regimes can be distinguished., Two regimes can be distinguished. +" For t<10 — 20 S, the deviation decreases sharply with r; this is seen directly on Fig. 3,,"," For $\tau \lesssim 10$ – $20\ S^{-1}$ , the deviation decreases sharply with $\tau$ ; this is seen directly on Fig. \ref{ttransp}, ," + where the, where the +operation. where two trees 7j aud 75 with roots are linked resulting in tree 7’. Assume that τι and 7 have hy.ho>0 white nodes. respectively.,"operation, where two trees $\tau_1$ and $\tau_2$ with roots are linked resulting in tree $\tau'$ Assume that $\tau_1$ and $\tau_2$ have $k_1, k_2 > 0$ white nodes, respectively." + The required poteutial ou this link is at most log(A4.|ko)., The required potential on this link is at most $\log{(k_1+k_2)}$. + By induction. P»xloge(hy|ke)vvjsv[Pe log.," By induction, $P_{\tau'} \leq \log{(k_1+k_2)} + \sum_{i=1}^{k_1} \log{i} + \sum_{i=1}^{k_2} \log{i} \leq \sum_{i=1}^{k_1+k_2} \log{i}$ ." + This follows from the fact that Ay!|hel<(G4heA)! for any integers Ay.koLl.," This follows from the fact that $k_1! + k_2! < (k_1 + k_2 - 1)!$, for any integers $k_1, k_2 > 1$." + We maintain the following credits in addition to the potential function: Next. we analyze the time bounds for our operations.," $\Box$ We maintain the following credits in addition to the potential function: Next, we analyze the time bounds for our operations." + Each operation must maintain the potential fuuction. the credits. aud pay for the work it performs.," Each operation must maintain the potential function, the credits, and pay for the work it performs." + No potential or credit changes are required., No potential or credit changes are required. + The actual work of is OCT)., The actual work of is $O(1)$. + It follows that the worst-case cost of is O(L1)., It follows that the worst-case cost of is $O(1)$. + If the inserted node is white. extra potential units may be needed.," If the inserted node is white, extra potential units may be needed." + But. as Leuuua 1 ilustrates. these units are borrowed from the logarithnuc cost perdelete-nin. aud the operation need not pax for that.," But, as Lemma \ref{l1} illustrates, these units are borrowed from the logarithmic cost per, and the operation need not pay for that." + Asstune that as a result of theénsert operation node c is linked to node y., Assume that as a result of the operation node $x$ is linked to node $y$. + Tf y isactiec. theectiec-parent credits need to be increased by O(1).," If $y$ is, the credits need to be increased by $O(1)$." + Hour isaetive. aud the previous leftimost child of yo wasaeetree. the credits need tobe increased by O(1).," If $x$ is, and the previous leftmost child of $y$ was, the credits need tobe increased by $O(1)$." + Since the size of the heap increased by one. theheap credits ueed to be increased by O(1).," Since the size of the heap increased by one, the credits need to be increased by $O(1)$." + The credits uced to be increased by Ollogloeta|1).loglog9) per decreased node. which still sums up to O(1) as indicated by the following proposition.," The credits need to be increased by $O(\log \log (n+1) - \log \log{n})$ per decreased node, which still sums up to $O(1)$ as indicated by the following proposition." + For i72. Dut (1|ly«oc. where e is thebase of the natural ," For $n>2$, But $(1+\frac{1}{n})^{n} < e$, where $e$ is thebase of the natural logarithm." +"οσα], The actual work to link an iusertednode with the main tree is OCT).",$\Box$ The actual work to link an insertednode with the main tree is $O(1)$ . + It follows thatthe amortized cost of 4isert is O(1)., It follows thatthe amortized cost of is $O(1)$ . +assuming that the PAH to dust ratio is homogeneous throughout the disk we find that the total dust mass adds up to 2-10 MMe.,assuming that the PAH to dust ratio is homogeneous throughout the disk we find that the total dust mass adds up to $\cdot$ $^{-6}$ $_{\odot}$. + This means that most of the dust mass resides in large grains (200:1)., This means that most of the dust mass resides in large grains (200:1). + Using the canonical gas to dust ratio of 100 the total disk mass becomes 2-10-+ MMo., Using the canonical gas to dust ratio of 100 the total disk mass becomes $\cdot$ $^{-4}$ $_{\odot}$. +" Note, that a larger gas mass could be required to reproduce the observations."," Note, that a larger gas mass could be required to reproduce the observations." +" The results presented above naturally lead to the picture of a disk in which the dust grains in the outer disk are coagulated and settled towards the midplane, while the gas is still available to keep the PAH molecules in the higher atmosphere of the disk."," The results presented above naturally lead to the picture of a disk in which the dust grains in the outer disk are coagulated and settled towards the midplane, while the gas is still available to keep the PAH molecules in the higher atmosphere of the disk." +" As already noted in the theoretical study by ?,, growth and settling of the dust grains leads to a natural increase of the relative strength of the PAH signature, as is observed in this disk and confirmed by our modeling effort."," As already noted in the theoretical study by \cite{2007A&A...473..457D}, growth and settling of the dust grains leads to a natural increase of the relative strength of the PAH signature, as is observed in this disk and confirmed by our modeling effort." +" In general the study of ? showed that the natural outcome of a group I source after grain-growth and sedimentation of the dust is a group II source, that maintains the flaring structure for the gas."," In general the study of \cite{2007A&A...473..457D} showed that the natural outcome of a group I source after grain-growth and sedimentation of the dust is a group II source, that maintains the flaring structure for the gas." +" However, observational studies have shown that most group II sources lack a flaring gas distribution. ?,,"," However, observational studies have shown that most group II sources lack a flaring gas distribution. \cite{2001A&A...365..476M}," + ? showed that group I sources display significantly more PAH emission.," \cite{2004A&A...422..621A} + showed that group I sources display significantly more PAH emission." + ? showed that group I sources have in general stronger II] emission., \cite{2005A&A...436..209A} showed that group I sources have in general stronger I] emission. + Apparently the gas of most group II sources has either dramatically decreased its scale height because of the lack of heating or the gas has been dispersed from the disk., Apparently the gas of most group II sources has either dramatically decreased its scale height because of the lack of heating or the gas has been dispersed from the disk. +" How disks can lose their gas is currently being debated (see ?)), but photoevaporation seems to be the most likely mechanism."," How disks can lose their gas is currently being debated (see \citealt{2008PhST..130a4024H}) ), but photoevaporation seems to be the most likely mechanism." + On the other hand there is a fair number of group II sources that do show indications of a flaring gas distribution., On the other hand there is a fair number of group II sources that do show indications of a flaring gas distribution. +" Some group II sources display PAH emission in their um spectra,"," Some group II sources display PAH emission in their $\mu$ m spectra," +to convert from counts to electron deusitv.,to convert from counts to electron density. + Frou this. we find an electron density at the relie of 1.«10. 7.," From this, we find an electron density at the relic of $1\times10^{-3}$ $^{-3}$." + Given this aud the average cluster teniperature (6.0 keV) deteruuued im 77.. we fud that the thermal pressure of the siuroundiug hot eas at the location of the radio rele is L.6«4100 dyne 7.," Given this and the average cluster temperature (6.0 keV) determined in \ref{sec:full}, we find that the thermal pressure of the surrounding hot gas at the location of the radio relic is $1.6 \times +10^{-11}$ dyne $^{-2}$." + This value is larger than the estimate for the pressure of the radio plasma. as has been seen in other N-rav cluster radio bubbles2005).," This value is larger than the estimate for the pressure of the radio plasma, as has been seen in other X-ray cluster radio bubbles." +. We note that if the relie were at a larger radius than the projected radius. the electron deusitv and thermal pressure would be lower. bringing it closer to the radio plasumia pressure estinate.," We note that if the relic were at a larger radius than the projected radius, the electron density and thermal pressure would be lower, bringing it closer to the radio plasma pressure estimate." + Iu addition. iu the case of Abell 13. i£ the values assiuned by Sleeetal.(2001) for (1) the ratio of masses m heavy particles and clectrous: (2) the plasma filline factor: (3) the low frequency spectral iudex. and (1) the augle between the," In addition, in the case of Abell 13, if the values assumed by \citet{srm+01} for (1) the ratio of masses in heavy particles and electrons; (2) the plasma filling factor; (3) the low frequency spectral index, and (4) the angle between the" +value.,value. + The central density. ΠΟ(Πσ). and interstellar radiation field (ISRF) incident on the core were then varied until the output model's SED matched the observed data.," The central density, $_0$ $_2$ ), and interstellar radiation field (ISRF) incident on the core were then varied until the output model's SED matched the observed data." + The model SED is shown as a dotted line in Fig. 4.., The model SED is shown as a dotted line in Fig. \ref{seda}. + The final values for no(H2) and the incident ISRF are noted in Table 1.., The final values for $_0$ $_2$ ) and the incident ISRF are noted in Table \ref{coreprop}. + It should be noted that the internal structure of cores within infrared dark clouds have not yet been observed in detail., It should be noted that the internal structure of cores within infrared dark clouds have not yet been observed in detail. + Hence our assumption of an elliptical geometry may be an over-simplification., Hence our assumption of an elliptical geometry may be an over-simplification. + If there is structure on smaller scales than we can resolve this might affect our results., If there is structure on smaller scales than we can resolve this might affect our results. + For example. small scale fragmentation might allow the ISRF to penetrate further into the cores. meaning that our calculated values are probably upper values.," For example, small scale fragmentation might allow the ISRF to penetrate further into the cores, meaning that our calculated values are probably upper values." + Figs. 1-, Figs. \ref{model3095}- +-3 show the output of the model at wavelengths corresponding to the wavelengths of the observed data., \ref{model3176} show the output of the model at wavelengths corresponding to the wavelengths of the observed data. + The modelled images have pixels of 0.02x0.02 ppe in size. which corresponds to ppixels for G031.03+00.26 and pixels for GO30.50+00.95 and GO031.03+00.76.," The modelled images have pixels of $\times$ pc in size, which corresponds to pixels for G031.03+00.26 and pixels for G030.50+00.95 and G031.03+00.76." + All the images have been convolved with the telescope beam., All the images have been convolved with the telescope beam. + For wavelengths where no emission is visible in the model. an image showing background radiation is shown.," For wavelengths where no emission is visible in the model, an image showing background radiation is shown." + As no attempt was made to correctly model the surrounding area. these Images are not an accurate," As no attempt was made to correctly model the surrounding area, these images are not an accurate" +Galactic globular clusters (GGCs) are the most ancient objects known for which reliable ages can be determined. and as the Universe cannot be vounger than the oldest objects it contains. GGCs provide one of the most robust constraints that we have on cosmological models.,"Galactic globular clusters (GGCs) are the most ancient objects known for which reliable ages can be determined, and as the Universe cannot be younger than the oldest objects it contains, GGCs provide one of the most robust constraints that we have on cosmological models." + However. absolute GGCS ages have to be estimated in order (to apply (Bis constraint.," However, absolute GGCs ages have to be estimated in order to apply this constraint." + Although significant improvements in the absolute GGCs age estimates have been obtained in the last decade. they are still allected by both observational and theoretical uncertainties (Vandenberegetal.1996:Cassisi1998:2009) at the zz204 level.," Although significant improvements in the absolute GGCs age estimates have been obtained in the last decade, they are still affected by both observational and theoretical uncertainties \citep{VSB96, + Cass98, Cass09} at the $\approx20$ level." + Nevertheless. il is possible to determine relative GGC ages wilh (he accuracy. required (ο address some outstanding problems. such as those related to the Milkv. Ways formation process.," Nevertheless, it is possible to determine relative GGC ages with the accuracy required to address some outstanding problems, such as those related to the Milky Way's formation process." + The pivotal importance of these problems ancl (he need to improve age estimates as lar as possible are the basis of the huge effort devoted in recent decades to the gathering of the relative ages ol QGCs., The pivotal importance of these problems and the need to improve age estimates as far as possible are the basis of the huge effort devoted in recent decades to the gathering of the relative ages of GGCs. + As a consequence. there exists quite a rich literature dedicated to this fundamental topic (Stetsonetal.1996:SarajediniLOOT:Rosenberg1999:Salaris therein)..," As a consequence, there exists quite a rich literature dedicated to this fundamental topic \citep[][and references +therein]{Stet96, Sar97, R99, SW02, DeA05, MF09}." + The relative age—dating techniques for GGCs - almost universally adopted - can be separated into two basic classes: those methods that are based. on brightness dillerence measurements - (he method -- aud those that are based on color difference measurements - thehorizontal method - in the colormagnitude diagram. (CAID).," The relative $-$ dating techniques for GGCs - almost universally adopted - can be separated into two basic classes: those methods that are based on brightness difference measurements - the method -, and those that are based on color difference measurements - the method - in the color–magnitude diagram (CMD)." + The most commonly adopted of thevertical methods is the magnitude difference between the main sequence turn-off (AISTO) and the zero age horizontal branch (ZAIIB). usually estimated starting from the level of the RI. Lyrae instability strip.," The most commonly adopted of the methods is the magnitude difference between the main sequence turn-off (MSTO) and the zero age horizontal branch (ZAHB), usually estimated starting from the level of the RR Lyrae instability strip." + Thehorizontal method is based on the measurenient of (he color difference between the MSTO and a point in (he lower part of the red giant branch. (RGB)., The method is based on the measurement of the color difference between the MSTO and a point in the lower part of the red giant branch (RGB). +1999).,1999). + Filling a boxy elliptical does not require that a bar be present (see discussion of NGC 2366. IIunter 22001). but it does suggest that one could be present.," Fitting a boxy elliptical does not require that a bar be present (see discussion of NGC 2366, Hunter 2001), but it does suggest that one could be present." + A second observational clue to the presence of a bar can be the twisting of isophotes as one goes [rom the bar to the disk., A second observational clue to the presence of a bar can be the twisting of isophotes as one goes from the bar to the disk. + This is clearly seen in NGC 4449. for example (IIunter 11999). but is not a necessary condition.," This is clearly seen in NGC 4449, for example (Hunter 1999), but is not a necessary condition." + In DDO 43 there is a small shift in the position angle of the major axis as one goes [rom the inner isophotes fit with the boxy curve (με=244 mag 7) to outer isophotes., In DDO 43 there is a small shift in the position angle of the major axis as one goes from the inner isophotes fit with the boxy curve $\mu_{V_0}=24.4$ mag $^{-2}$ ) to outer isophotes. + The inner rectangle is best fit with a boxv curve wilh a position angle of wwhile the outer isophotes have a position augle of6., The inner rectangle is best fit with a boxy curve with a position angle of while the outer isophotes have a position angle of. +57.. This is a shift of aand it is clearly visible in Figure 2.., This is a shift of and it is clearly visible in Figure \ref{fig:v}. + A third characteristic of a bar is a potential mis-alignment between the kinematic line of nodes and the major axis of the bar., A third characteristic of a bar is a potential mis-alignment between the kinematic line of nodes and the major axis of the bar. + In a barred galaxy. the isophotes are not circular in the principal plane. and the stellar orbils may precess. so the line of nodes and morphological axis need not be aligned.," In a barred galaxy, the isophotes are not circular in the principal plane, and the stellar orbits may precess, so the line of nodes and morphological axis need not be aligned." + This is the case in. for example. NGC 1156 (IIunter 22002).," This is the case in, for example, NGC 1156 (Hunter 2002)." + We will discuss the gas kinematics in DDO 43 below. but it is appropriate to mention here that the position angle of the line of nodes of the rotating gas svstem is294.," We will discuss the gas kinematics in DDO 43 below, but it is appropriate to mention here that the position angle of the line of nodes of the rotating gas system is." +.. Thus. the kinematic line of nodes aud (he major axis of the rectangular part of (he galaxy. are nearly perpendicular. being mis-aliened by867.," Thus, the kinematic line of nodes and the major axis of the rectangular part of the galaxy are nearly perpendicular, being mis-aligned by." +.. This is strong evidence (that the boxy appearance of the optical galaxy is in [act due to a stellar bar potential., This is strong evidence that the boxy appearance of the optical galaxy is in fact due to a stellar bar potential. + llowever. if DDO 43 is barred. the bar occupies a large fraction of the optical galaxy.," However, if DDO 43 is barred, the bar occupies a large fraction of the optical galaxy." + This is also the case in NGC 1156. NGC! 2366. and NGC 4449.," This is also the case in NGC 1156, NGC 2366, and NGC 4449." +" In DDO 43 e, λος is 0.5: where Ba; is the radius αἱ a fry, of 25 magnitudes per 7.", In DDO 43 $a_{bar}$ $_{25}$ is 0.8; where $_{25}$ is the radius at a $\mu_{B_0}$ of 25 magnitudes per $^{-2}$. + In NGC! 2366 this ratio is 1.6 and in NGC! 4449 it is 0.6., In NGC 2366 this ratio is 1.6 and in NGC 4449 it is 0.6. +" By contrast. in SdSim spiral galaxies (5, /1155 is 0.050.9 (Ehneereen Elmegreen 1955. Martin 1995. Elueereen 11996)."," By contrast, in Sd–Sm spiral galaxies $a_{bar}$ $_{25}$ is 0.05–0.3 (Elmegreen Elmegreen 1985, Martin 1995, Elmegreen 1996)." + Thus. DDO 43 and other barred Im galaxies are unusual in being dominated optically bv the bar if it is present.," Thus, DDO 43 and other barred Im galaxies are unusual in being dominated optically by the bar if it is present." + In Figuree 3. we show surface brightnesse euts in V alonge the major and minor axis of the bar., In Figure \ref{fig:barcuts} we show surface brightness cuts in V along the major and minor axis of the bar. + The euts are wwide., The cuts are wide. + The cuts from DDO 43 are more rounded at the top Chan cuts across the bar of NGC 4449. but resemble similar cuts made in NGC 2366.," The cuts from DDO 43 are more rounded at the top than cuts across the bar of NGC 4449, but resemble similar cuts made in NGC 2366." + The surface brightness of bars in late-tvpe spirals ave more often fit with an exponential whereas (hose in earlv-tvpe spirals are [latter (Combes Elmeereen 1993. Elineereen," The surface brightness of bars in late-type spirals are more often fit with an exponential whereas those in early-type spirals are flatter (Combes Elmegreen 1993, Elmegreen" +Faint 15 um ssurveys have revealed a large population. ofο strongly evolving. luminous.h infraredn (IR) starburstsNnStS οat lLzc—| FAZAUSNAAussel 1999: Elbaz 1999: Elbaz 2002).,Faint 15 $\mu$ m surveys have revealed a large population of strongly evolving luminous infrared (IR) starbursts at $z\approx1$ Aussel 1999; Elbaz 1999; Elbaz 2002). + The space density. of> these sources is. an order of> magnitude. higher. than found for luminous. IR starbursts in. the local Universe. mat Ο.Ε formingealanspopul, The space density of these sources is an order of magnitude higher than that found for luminous IR starbursts in the local Universe $z\simlt0.1$ ). +ation Althoughthesesourcesontyaccount fForasmall fractionof the thevarelikelytocontributeasigni ficvn FormationhistorvoFtheU niverse(e.g.. Chary&Elbaz2001: Elbazer," Although these sources only account for a small fraction of the star-forming galaxy population, they are likely to contribute a significant fraction of the star-formation history of the Universe (e.g., Chary Elbaz 2001; Elbaz 2002)." +" In this Letter we use the | Ms Paper and↼∙⊲ISOCA ia""DEN-DE-N: BrandtBe apo2001b. hereaftercaofhe: ν HHDF-N (Aussel 1999, 2002) surveys to show that these strongly evolving luminous IR starbursts are identified with the population of apparently normal galaxies detected in deep X-ray surveys GGiaccont 2001: Hornschemeier 2001. hereafter Paper Il: Brandt etal.2001a. hereafter Paper IV)."," In this Letter we use the 1 Ms Deep Field-North (CDF-N; Brandt 2001b, hereafter Paper V) and HDF-N (Aussel 1999, 2002) surveys to show that these strongly evolving luminous IR starbursts are identified with the population of apparently normal galaxies detected in deep X-ray surveys Giacconi 2001; Hornschemeier 2001, hereafter Paper II; Brandt 2001a, hereafter Paper IV)." + Previous cross-identification studies have shown that z of the ssources are AGN-dominated: however. few constraints have been placed on the larger fraction of ssources that do not have obvious AGN activity PPaper II: Paper IV; Fadda 2002).," Previous cross-identification studies have shown that $\approx$ of the sources are AGN-dominated; however, few constraints have been placed on the larger fraction of sources that do not have obvious AGN activity Paper II; Paper IV; Fadda 2002)." + The primary focus of this paper is to place constraints on the X-ray properties of these sources., The primary focus of this paper is to place constraints on the X-ray properties of these sources. +" The Galactic column density toward the CDF-N is (1.60.4)«10°"" emo? (Stark 1992). and Hy=65 km s! Mpe7! Qy= land Q4=2 are adopted throughout this Letter."," The Galactic column density toward the CDF-N is $(1.6\pm 0.4)\times 10^{20}$ $^{-2}$ (Stark 1992), and $H_{0}=65$ km $^{-1}$ $^{-1}$, $\Omega_{\rm M}=\onethird$, and $\Omega_{\Lambda}=\twothirds$ are adopted throughout this Letter." + The | Ms CDF-N observations were centered on the HDF-N (Williams 1996) and cover z450 aremin? (Paper V)., The 1 Ms CDF-N observations were centered on the HDF-N (Williams 1996) and cover $\approx450$ $^2$ (Paper V). + The xy 21.5 aremin? region of the CDF-N survey coincident with the 15 um oobservations is close to the CDF-N aim point., The $\approx$ 21.5 $^2$ region of the CDF-N survey coincident with the 15 $\mu$ m observations is close to the CDF-N aim point. +" Forty-nine high-significance X-ray sources false-positive probabilityM"" threshold ofμμ 107"") are detected in. this. area down (P ONEavi Tς5 keV. soraCON and; 28>Qo Lr"" gaiDC U1X ES EOl στ""ς (Papevc .. dme erg 7S . JeseCOVely, cr m."," Forty-nine high-significance X-ray sources false-positive probability threshold of $^{-7}$ ) are detected in this area down to on-axis 0.5–2.0 keV (soft-band) and 2–8 keV (hard-band) flux limits of $\approx3\times10^{-17}$ erg $^{-2}$ $^{-1}$ and $\approx2\times10^{-16}$ erg $^{-2}$ $^{-1}$, respectively (Paper V)." + The TREOxarservationsr Tawere[ |performedformed at 6.7 aim um ancand 15 a um. Becausecaus tionssigjonsignificantlyIesaü smasmallerler. tield-field-of-view)V c); v 67y um aloο and the small number of sources detected. only the 15 um observations are considered here.," The observations were performed at 6.7 $\mu$ m and 15 $\mu$ m. Because of the significantly smaller field-of-view of the 6.7 $\mu$ m observations and the small number of sources detected, only the 15 $\mu$ m observations are considered here." + Discrete sources are detected down to fisam7 20 wy. although the oobservations are far from complete at this depth.," Discrete sources are detected down to $f_{\rm 15\mu m}\approx$ 20 $\mu$ Jy, although the observations are far from complete at this depth." + In this study we focus on the 41 sources with fisam2 100 gy in the complete 15 um selected sample of Aussel (2002)., In this study we focus on the 41 sources with $f_{\rm 15\mu m}\ge$ 100 $\mu$ Jy in the complete 15 $\mu$ m selected sample of Aussel (2002). + Optical (/-band) sources taken from the photometric catalog produced in Alexander (2001b. hereafter Paper VI) were matched to X-ray and 15 um sources using matching radii of aand3.. respectively (see Paper Vl; Aussel The 15 um sources have J= 18-23. and all but one source have redshifts in the catalogs of Cohen (2000). Cohen (2001) or Dawson (2001).," Optical $I$ -band) sources taken from the photometric catalog produced in Alexander (2001b, hereafter Paper VI) were matched to X-ray and 15 $\mu$ m sources using matching radii of and, respectively (see Paper VI; Aussel The 15 $\mu$ m sources have $I=$ 18–23, and all but one source have redshifts in the catalogs of Cohen (2000), Cohen (2001) or Dawson (2001)." + By contrast. z of the X-ray sources in the 15 gm rregion have /=23.ISOCA although all but two of the J«<23 sources have redshifts.," By contrast, $\approx$ of the X-ray sources in the 15 $\mu$ m region have $I\ge23$, although all but two of the $I<23$ sources have redshifts." + In total 14 of the 41 15 um sources have X-ray counterparts in the high-significance X-ray source However. given," In total 14 of the 41 15 $\mu$ m sources have X-ray counterparts in the high-significance X-ray source However, given" +"���£,/£y, is Ok/T). which is vanishingly small for the preferred mode of /+x.","$\calE_x/\calE_y$ is $O(k/l)$, which is vanishingly small for the preferred mode of $l \to \infty$." + Although now all diffusivities are restored. the plwsical mechanism underlving the instability is unchanged. so we may still expect the ratio of the emls to be Ο/Η). aud hence small.," Although now all diffusivities are restored, the physical mechanism underlying the instability is unchanged, so we may still expect the ratio of the emfs to be $O(k/l)$, and hence small." + It can also be seen that the emf is localised towards the top of the laver., It can also be seen that the emf is localised towards the top of the layer. + Consideration of the simplest instability criterion given by Gilman(1970).. showing that a laver is uustable if the magnetic field strength decreases wilh height. vields for the current problem (he condition C/(L+¢2)>0.," Consideration of the simplest instability criterion given by \citet{Gilman70}, showing that a layer is unstable if the magnetic field strength decreases with height, yields for the current problem the condition $\zeta/(1+\zeta z)>0$." + The left-hand side of this inequality is clearly largest at the top of the laver. leading to the eigenfunctions (and hence the emls) being peaked in this region.," The left-hand side of this inequality is clearly largest at the top of the layer, leading to the eigenfunctions (and hence the emfs) being peaked in this region." + Figure 2 shows the .c- ancl y-components of V.x€: these are the dynamically sienilicant elements in the evolution of the mean field. as explained in relsec:lormulation..," Figure \ref{fig:curle} shows the $x$ - and $y$ -components of $\bfnabla\times\bfcalE$; these are the dynamically significant elements in the evolution of the mean field, as explained in \\ref{sec:formulation}." + It is also of interest (ο examine how the emf varies with the parameters of the problem., It is also of interest to examine how the emf varies with the parameters of the problem. + In order to do this aud to be able to displaythe results in a meaningful way. we must first introduce an overall (z-independent) measure of the eml: with this in mincl. we deline € where κ.€;=Ch£u)>.," In order to do this and to be able to displaythe results in a meaningful way, we must first introduce an overall$z$ -independent) measure of the emf; with this in mind, we define $\bfbarcalE$ where $\barcalE_i=\left(\int_0^1\bfcalE_i^2 dz\right )^{1/2}$." + Figures. 3. 6 show the variation. of (0a£ with. 0. 8. ¢- and By respectively.," Figures \ref{fig:emfsomega} – \ref{fig:emfsb0} show the variation of $\bfbarcalE$ with $\Omega$ , $\theta$ , $\zeta$ and $B_0$ respectively." + Throughout all simulations. the dimensionless magnetic diffusivity.," Throughout all simulations, the dimensionless magnetic diffusivity," + + + + Έτ Ü., + + + = 0. + sellingn4-1 and rearranging gives — which iterates to Atn=1. = Atn=2. equation (À)) is p2= — + ed) d£ Saa," Setting$n+1\rightarrow n$ and rearranging gives = -, which iterates to = + At $n=1$, = = At$n=2$, equation \ref{eq:42}) ) is = - + ]/ ," +ab larec scales (Dellavet. Dvali Gabadacdze 2002: Carloni et al.,"at large scales (Deffayet, Dvali Gabadadze 2002; Carloni et al." + 2005: Damour. Kogan Papazoglou 2002).," 2005; Damour, Kogan Papazoglou 2002)." + The current. cosmological data is only good enough to measure the local value of the dark energy density., The current cosmological data is only good enough to measure the local value of the dark energy density. + Pherefore. the present data prefers a cosmological constant. mainly because it is the simplest. model. whilst. timeevolving models. of dark energy. (which possess more free. parameters) remain relatively unconstrained by present. datasets (see Corasaniti et al.," Therefore, the present data prefers a cosmological constant, mainly because it is the simplest model, whilst time–evolving models of dark energy (which possess more free parameters) remain relatively unconstrained by present datasets (see Corasaniti et al." + 2004: Bassett et al., 2004; Bassett et al. + 2004: Liddle ct al., 2004; Liddle et al. + 2006)., 2006). + Over the next decade. numerous experiments are proposed. over a wide range of redshifts. to explore. the timeevolution of dark energy and determine its density as a function of cosmic time.," Over the next decade, numerous experiments are proposed, over a wide range of redshifts, to explore the time–evolution of dark energy and determine its density as a function of cosmic time." + These experiments (both groundbased and spacebased) represent an investment of billions of dollars. ancl essentially use two general techniques. to probe the dark energy: geometrical tests of the expansion history of the Universe using “stancarel tracers” (see below). and/or observations of the rate of growth of largescale structures (clusters superclusters) in the Universe.," These experiments (both ground--based and space–based) represent an investment of billions of dollars and essentially use two general techniques to probe the dark energy: geometrical tests of the expansion history of the Universe using “standard tracers” (see below), and/or observations of the rate of growth of large–scale structures (clusters superclusters) in the Universe." + The relative merits of these two techniques. aud the proposed experiments. have been explored in detail by the U.S. Dark Ilnergv Task Force (DEVE: Albrecht et al.," The relative merits of these two techniques, and the proposed experiments, have been explored in detail by the U.S. Dark Energy Task Force (DETF; Albrecht et al." + 2006) ancl a similar endeavour in the τν. In this paper. we focus on measurements of the timeevolution of dark energv using observations of the baryon acoustic oscillations (D.XOs) from large galaxy surveys.," 2006) and a similar endeavour in the U.K. In this paper, we focus on measurements of the time--evolution of dark energy using observations of the baryon acoustic oscillations (BAOs) from large galaxy surveys." + The BAOs are generated by acoustic waves in the barvonphoton plasma in the carly Universe. which become frozen into the CAIB radiation. and the distribution of matter. soon after the Universe cools and recombines at zc1100.," The BAOs are generated by acoustic waves in the baryon–photon plasma in the early Universe, which become frozen into the CMB radiation, and the distribution of matter, soon after the Universe cools and re–combines at $z \simeq 1100$." + Over the last 5 vears. the scale of these BAOs has been accurately measured in the CALB by a number of experiments (WALAP. Archeops BOOAERanG). as well as discovered. in. the clistribution of matter (ealaxies and clusters) at low recshift (sce Miller. Nichol Batuski 2001: Eisenstein et al.," Over the last 5 years, the scale of these BAOs has been accurately measured in the CMB by a number of experiments (WMAP, Archeops BOOMERanG), as well as discovered in the distribution of matter (galaxies and clusters) at low redshift (see Miller, Nichol Batuski 2001; Eisenstein et al." + 2005: Cole ct al., 2005; Cole et al. + 2005: Padmanabhan et al., 2005; Padmanabhan et al. + 2006)., 2006). + For example. Eisenstein et al. (," For example, Eisenstein et al. (" +2005) used these observations to constrain he flatness of the Universe to under the assumption hat dark energy is à cosmological constant.,"2005) used these observations to constrain the flatness of the Universe to, under the assumption that dark energy is a cosmological constant." + As the scale of the BAOs can be predicted: to subyereent accuracy (see Eisenstein. Seo White 2006: Eisenstein White 2004). they provide an excellent “standard ruler” which can be used to map the geometry/ of 1ο Universe through the angular-ciameter distance and the —ubble parameter relation (sec Blake Glazebrook 2003: ulCO Eisenstein 2008: Iu IHaiman 2003).," As the scale of the BAOs can be predicted to sub--percent accuracy (see Eisenstein, Seo White 2006; Eisenstein White 2004), they provide an excellent “standard ruler” which can be used to map the geometry of the Universe through the angular-diameter distance and the Hubble parameter relation (see Blake Glazebrook 2003; Seo Eisenstein 2003; Hu Haiman 2003)." + Several spectroscopic experiments have been proposed. to observe nd measure this standard ruler at high redshift and. thus constrain the timeevolution of dark energy. e.g... WEMOS Bassett. Nichol Eisenstein 2005: Glazebrook et al.," Several spectroscopic experiments have been proposed to observe and measure this standard ruler at high redshift and thus constrain the time–evolution of dark energy, e.g., WFMOS (Bassett, Nichol Eisenstein 2005; Glazebrook et al." + 2005a). Darvon Oscillation Probe (Glazebrook et al.," 2005a), Baryon Oscillation Probe (Glazebrook et al." + 2005b). VIRUS (Mill et al.," 2005b), VIRUS (Hill et al." + 2005) in the optical ancl NUR. and the EIubble Sphere Hydrogen Survey (Peterson et al.," 2005) in the optical and NIR, and the Hubble Sphere Hydrogen Survey (Peterson et al." + 2006) and Square WWilometre Array (Blake ct al., 2006) and Square Kilometre Array (Blake et al. + 2004) in the radio., 2004) in the radio. + These new experiments share the desire to measure millions of ealaxy redshifts (at high redshift) over large volumes of the Universe to control the errors from both cosmic variance and Poisson noise (Blake Glazebrook 2003: Glazebrook Blake 2005)., These new experiments share the desire to measure millions of galaxy redshifts (at high redshift) over large volumes of the Universe to control the errors from both cosmic variance and Poisson noise (Blake Glazebrook 2003; Glazebrook Blake 2005). + Given the large investment in time and money for these next generation. BAO experiments. we study here the optimal survey strategy for. galaxy recdshift surveys like those proposed for WEALOS (Nide-Field. Multi-Objeet Spectrograph: see Glazebrook. et al.," Given the large investment in time and money for these next generation BAO experiments, we study here the optimal survey strategy for galaxy redshift surveys like those proposed for WFMOS (Wide-Field Multi-Object Spectrograph; see Glazebrook et al." + 20052))., 2005a). + This work builds upon our previous development. of the Integrated Parameter Survey Optimization (IPSO) framework (Bassett 2005: Bassett. Parkinson Nichol 2005) and addresses kev observational issues such as: Although our answers for these questions are derived for WEMOSlike galaxy surveys. we believe that our results are general to. BAO experiments in. optical wavebancds.," This work builds upon our previous development of the Integrated Parameter Survey Optimization (IPSO) framework (Bassett 2005; Bassett, Parkinson Nichol 2005) and addresses key observational issues such as: Although our answers for these questions are derived for WFMOS–like galaxy surveys, we believe that our results are general to BAO experiments in optical wavebands." + The analysis presented. in. this paper optimises A cosmologies (νο. à cosmological constant) for the underlying cosmological model against which we are optimizing.," The analysis presented in this paper optimises $\Lambda$ --cosmologies (i.e., a cosmological constant) for the underlying cosmological model against which we are optimizing." + Ideally we would like to optimize the observations for à variety of dark energy. models. anc we will present. such. an analysis in a future paper.," Ideally we would like to optimize the observations for a variety of dark energy models, and we will present such an analysis in a future paper." + As a consequence therefore. our conclusions ancl results naturally favour lower-redshilt BAO observations where the elfect of A on the expansion ustory of the Universe is greatest.," As a consequence therefore, our conclusions and results naturally favour lower-redshift BAO observations where the effect of $\Lambda$ on the expansion history of the Universe is greatest." + We also neglect. the improvements to the distance measurements that will arise w reconstructing the acoustic peak in the non-linear regine (sce Eisenstein ct al 200Gb). assuming instead. a non-linear cut-olf for the power spectrum that evolves with recshift.," We also neglect the improvements to the distance measurements that will arise by reconstructing the acoustic peak in the non-linear regime (see Eisenstein et al 2006b), assuming instead a non-linear cut-off for the power spectrum that evolves with redshift." + In Section 2.. we outline our methodology. ancl celine he Figure-of-Merit. (ΟΔΙΕ used. to. judge the optimality of clilferent survey designs.," In Section \ref{optimizingmethod}, we outline our methodology and define the Figure-of-Merit (FoM) used to judge the optimality of different survey designs." + In. Section 3... we lav out. the oocedures. used. to conducting the optimization. while in Section 4.. we describe the model we used to determine the density of target galaxies.," In Section \ref{procedure}, we lay out the procedures used to conducting the optimization, while in Section \ref{LFs}, we describe the model we used to determine the density of target galaxies." + We present and cliscuss our results in Sections 5. and G respectively., We present and discuss our results in Sections \ref{results} and \ref{discussion} respectively. + We conclude in Section 7.., We conclude in Section \ref{conclusions}. + We perform our optimization using the IPSO framework (Bassett 2005)., We perform our optimization using the IPSO framework (Bassett 2005). + Consider a set of allowed survey geometries. indexed bv s.," Consider a set of allowed survey geometries, indexed by $s$." +" For each survey geometry. 5. we compute an appropriate PFigure-of-Merit. (Fo0M) - also. known. as the utility. in Davesian evidence design. risk or [fitness - and optimization then simply proceeds. by selecting the survey geometry which. extremises (müinimising or nmaximising where appropriate) the FoM. Performing such an optimization therefore requires three elements: a survey configuration parameter space S to search through. a target parameter space O of the parameters that we wish to optimally constrain (labelled 0,,,,,). and a numerical Figure-of-Merit. (FoM)."," For each survey geometry, $s$, we compute an appropriate Figure-of-Merit (FoM) - also known as the utility in Bayesian evidence design, risk or fitness - and optimization then simply proceeds by selecting the survey geometry which extremises (minimising or maximising where appropriate) the FoM. Performing such an optimization therefore requires three elements: a survey configuration parameter space $S$ to search through, a target parameter space $\Theta$ of the parameters that we wish to optimally constrain (labelled $\theta_{\mu,\nu...}$ ), and a numerical Figure-of-Merit (FoM)." +. We consider cach of these three elements in turn., We consider each of these three elements in turn. + The survey parameters are those parameters which completely describe the survey we wish to test., The survey parameters are those parameters which completely describe the survey we wish to test. + We start by splitting the survey into two redshift regimes. one at low redshift (2~ 1) and one at high redshift (2~ 3). ancl considering these regimes completely. separately.," We start by splitting the survey into two redshift regimes, one at low redshift $z\simeq1$ ) and one at high redshift $z\simeq3$ ), and considering these regimes completely separately." + The properties of the survey in each regime are described by à set, The properties of the survey in each regime are described by a set +at 8.5 GHz. althoug1 there is an order-o-magnitude discrepaucy in the case of 11291.,"at 8.5 GHz, although there is an order-of-magnitude discrepancy in the case of 4291." + Iu constrast. lor I510 the Bondi rates exceed the ADAF rates for all four galaxies.," In constrast, for $P_6 = 10$ the Bondi rates exceed the ADAF rates for all four galaxies." + Meanineful comiparions between Alp aud Aly must cearly awalt a reliable pressure profile for ellipticals iu eeneral or. better stll. for the galaxies iu this stucv.," Meaningful comparions between $\dot{M}_{\rm B}$ and $\dot{M}_{\rm A}$ must clearly await a reliable pressure profile for ellipticals in general or, better still, for the galaxies in this study." + Futhe‘wore. integration over any adopted pressure profile shoud uot violate Lass listed in the table.," Futhermore, integration over any adopted pressure profile should not violate $L_{\rm RASS}$ listed in the table." + Ftrther radio observations to quantify any jet cojtauninatjon in 11621 and L1660 are also leneeded. for two 'easons.," Further radio observations to quantify any jet contamination in 4621 and 4660 are also needed, for two reasons." + First. renova of jet contamination leads to reduced ADAF emission. which in turn implies even lower ADAF accretion raes.," First, removal of jet contamination leads to reduced ADAF emission, which in turn implies even lower ADAF accretion rates." + Secold. the presence of jet emission could signif the importance of physical processes iguorec ln canoical ADAF models aud. explored by Diateoetal.(2000).," Second, the presence of jet emission could signify the importance of physical processes ignored in canonical ADAF models and explored by \citet{dim00}." +. Follow-up VLA imagine at other frequencies. 7. but at matched augular 'esollois would help assess jet contamination: the [Iux density from an ADAF-dominated source is expected to rise as vsl52 (Mahadevan41997:_Yine&Boughu.1998).. whereas the flux density. frou.. a jet-¢inated source is expected to exhibit a flatter spectral slope.," Follow-up VLA imaging at other frequencies, $\nu$, but at matched angular resolutions would help assess jet contamination: the flux density from an ADAF-dominated source is expected to rise as $\nu^{{1\over3} - {2\over5}}$ \citep{mah97,yi98}, whereas the flux density from a jet-dominated source is expected to exhibit a flatter spectral slope." + Although the 500-4/Jy. ipper liiits at 5 GHz from Wrobe&Heeschen(1991) do |ave similar resolutions. they are too luseusitive to impose useful coustraiits on spectral slopes between 5 and 8.5 GHz.," Although the $\mu$ Jy $^{-1}$ upper limits at 5 GHz from \citet{wro91b} do have similar resolutions, they are too insensitive to impose useful constraints on spectral slopes between 5 and 8.5 GHz." + Moreover. imagine wdh he NRAO VLBA could. provide morphological evideuce for nonthermal svuchrotron jets and spatially sepa‘ate jet from ADAF emissiOl. as was successfully doue for the spira eaανν [1258 (Herrnsteiuetal.1998).," Moreover, higher-resolution imaging with the NRAO VLBA could provide morphological evidence for nonthermal synchrotron jets and spatially separate jet from ADAF emission, as was successfully done for the spiral galaxy 4258 \citep{her98}." +. Finally. the tinage rims values. oxσον. achievect in this study recuirec 2]ours of ineeration time.," Finally, the image rms values, $\sigma_{\rm 8.5GHz}$, achieved in this study required 2 hours of integration time." +" Similar valtes for oxscup, WILL be achieved after a 2-ninute integratio lwith the expanded. VLA (NRAO2000 ).."," Similar values for $\sigma_{\rm +8.5GHz}$ will be achieved after a 2-minute integration with the expanded VLA \citep{nra00}. ." + This will make it feasible. for the first time. to target a| radio-quiesceut. ellipticals studied by agorrianetal.(1098).. whieh should sienificautly advauce our understaucdiug of advectiou-dominaed accretion flows as a new class of extragalactic radio οἱlitters.," This will make it feasible, for the first time, to target all radio-quiescent ellipticals studied by \citet{mag98}, which should significantly advance our understanding of advection-dominated accretion flows as a new class of extragalactic radio emitters." + The authors thaus MMahacevan for discussions., The authors thank Mahadevan for discussions. + This research. has mace use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. Caltech. under coutact with the National Aeronauties and 5yace Administration.," 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 Digitized Sky Surveys were prcxlucec a the Space Telescope Science Iusitute under U.S. Government eraut NAC W-2166., The Digitized Sky Surveys were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. + The ationalT Geographie Society - Palomar OselVaory Sky Atlas (PDOSS-I) was mace by he Califorila Iusitite of Technology with grauts [ro1 the National Geographic Society., The National Geographic Society - Palomar Observatory Sky Atlas (POSS-I) was made by the California Institute of Technology with grants from the National Geographic Society. + The Secoid Palomar Ooe‘vatory Sky Survey (POSS-II) was iade by the California Lustitute of 'Techuology with LucIs fonm tje National Science Foundation. te National Geographic Society. the Sloan Fouudation. tle Samue Oschin Foundation. and the Eastnan Ixodak Corporation.," The Second Palomar Observatory Sky Survey (POSS-II) was made by the California Institute of Technology with funds from the National Science Foundation, the National Geographic Society, the Sloan Foundation, the Samuel Oschin Foundation, and the Eastman Kodak Corporation." + NRAO is a [acility of the Naticla Scieice Foundation operated uuder cooperative agreement by AssociatedUniversities. Lue.," NRAO is a facility of the National Science Foundation operated under cooperative agreement by AssociatedUniversities, Inc." +Weak gravitational lensing is a rapidlv developing subject. with great. progress being mace in many related observational areas.,"Weak gravitational lensing is a rapidly developing subject, with great progress being made in many related observational areas." + The mass ancl density profiles of galaxies have been carefully. explored using galaxy-galaxy shear studies (c.g. Hoekstra ct al 2004). while large-scale structure can be traced. using cosmic shear (see ee. van Waerbeke Alellicr 2003. Relregier 2003. for reviews).," The mass and density profiles of galaxies have been carefully explored using galaxy-galaxy shear studies (e.g. Hoekstra et al 2004), while large-scale structure can be traced using cosmic shear (see e.g. van Waerbeke Mellier 2003, Refregier 2003 for reviews)." +" ""Phis has led το. significant constraints on cosmological parameters. such as the Uuetuation of the matter distribution. the density of matter. and the erowth rate of matter [Luctuations in the Universe."," This has led to significant constraints on cosmological parameters, such as the fluctuation of the matter distribution, the density of matter, and the growth rate of matter fluctuations in the Universe." + CGravitational lensing has received so much interest partially because it allows us. to. measure the mass of structures. with very few physical assumptions., Gravitational lensing has received so much interest partially because it allows us to measure the mass of structures with very few physical assumptions. + The distortion. of background. galaxies depends: only on. the ecometry of the lens system. the mass. and the use of the weak-lield limit of General Relativity.," The distortion of background galaxies depends only on the geometry of the lens system, the mass, and the use of the weak-field limit of General Relativity." + As such. lensing presents us with a method for measuring mass which is free of dynamical uncertainties associated. with questions as to whether the svstem is relaxed.," As such, lensing presents us with a method for measuring mass which is free of dynamical uncertainties associated with questions as to whether the system is relaxed." + I. is a direct measure of the mass present. whether in visible or dark form.," It is a direct measure of the mass present, whether in visible or dark form." + Weak gravitational lensing is typically studied: by examining the cllipticities of source galaxies. seeking a coherent alignment of these. cllipticities (or other combinations of weighted second-order moments of galaxy light) induced by mass along the line of sight (e.g. Ixaiscer. Squires Broadhurst. 1995. [xaiser 2000. Bernstein Jarvis 2002. Refreeicr Bacon 2003. Hirata Seljak 2003).," Weak gravitational lensing is typically studied by examining the ellipticities of source galaxies, seeking a coherent alignment of these ellipticities (or other combinations of weighted second-order moments of galaxy light) induced by mass along the line of sight (e.g. Kaiser, Squires Broadhurst, 1995, Kaiser 2000, Bernstein Jarvis 2002, Refregier Bacon 2003, Hirata Seljak 2003)." + Llowever. Goldberg Natarajan (2002) have shown that significant further. information is available from the skeweclness and arciness of the light cüstribution for source galaxies: we have further developed. this approach in Goldberg Bacon (2005) where we have labelled this third. order cllect as the “Hexion” of these images.," However, Goldberg Natarajan (2002) have shown that significant further information is available from the skewedness and arciness of the light distribution for source galaxies; we have further developed this approach in Goldberg Bacon (2005) where we have labelled this third order effect as the “flexion” of these images." + A related approach using ‘sextupole lensing” has recently been explored by Lewin Shmakova (2005), A related approach using `sextupole lensing' has recently been explored by Irwin Shmakova (2005). + In our previous paper (Goldberg Bacon 2005). we described. the theory of Uexion. anc demonstrated how this effect. can be measured using the Shapelet formalism (Bernstein Jarvis 2002. Relreeicr 2003. Befregier Bacon 2003)," In our previous paper (Goldberg Bacon 2005), we described the theory of flexion, and demonstrated how this effect can be measured using the Shapelet formalism (Bernstein Jarvis 2002, Refregier 2003, Refregier Bacon 2003)." + We also demonstrated: that the Uexion signal is present in Deep Lens Survey data (Wittman et al 2002)., We also demonstrated that the flexion signal is present in Deep Lens Survey data (Wittman et al 2002). + In this paper. we explore and describe what flexion is able to teach us in. the context of several," In this paper, we explore and describe what flexion is able to teach us in the context of several" +products are intermediate to the above cases.,products are intermediate to the above cases. +" According to Macketal.(2007) and Ricotti(2007)... a UCMH (with or without a PBH) at zz, with dark matter mass M, has a truncation radius UCMHs with PBHs at matter-radiation equality have Mppu."," According to \citet{Mack07} and \citet{Ricotti07}, a UCMH (with or without a PBH) at $z \approx z_{\rm eq}$ with dark matter mass $M_{\rm eq}$ has a truncation radius UCMHs with PBHs at matter-radiation equality have $M_{\rm eq} = M_{\rm PBH}$ ." +" The dark matter mass density within Ας) isthe M,sameis.. with a number density n5: at where mio=pe""/(000GeV)."," The dark matter mass density within $R_{\rm tr} (z)$ isthe same, with a number density $n_{\rm tr} \approx 1.8 \times 10^5\ \cm^{-3}\ m_{100}^{-1}$ at $z_{\rm eq}$, where $m_{100} = m_{\rm DM} c^2 / (100\ \GeV)$." +" Within Roy=iil,ισα). the cog.UCMH density profile goes as p r2, because the PBH dominates the mass: at larger radii. a 17 density profile holds (Bertschinger1985)."," Within $R_{\rm eq} = R_{\rm tr} (z_{\rm eq})$, the UCMH density profile goes as $\rho \propto r^{-3/2}$ , because the PBH dominates the mass; at larger radii, a $r^{-9/4}$ density profile holds \citep{Bertschinger85}." +. This density profile is shallower than in Ricotti&Gould(2009).. which assumed px7*7 throughout. and that adiabatic contraction increased the dark matter density further.," This density profile is shallower than in \citet{Ricotti09}, which assumed $\rho \propto r^{-9/4}$ throughout, and that adiabatic contraction increased the dark matter density further." + WIMPs have nearly radial orbits around the PBH. sampling a wide range of densities in each orbit.," WIMPs have nearly radial orbits around the PBH, sampling a wide range of densities in each orbit." + Thus. the dark matter annihilating in theinner regions of the UCMH is often stored much farther out.," Thus, the dark matter annihilating in theinner regions of the UCMH is often stored much farther out." + To answer whether WIMPs with apocenters of survive to the present day. we calculate the number of annihilationsRey a WIMP is expected to experience during the Universe's history.," To answer whether WIMPs with apocenters of $R_{\rm eq}$ survive to the present day, we calculate the number of annihilations a WIMP is expected to experience during the Universe's history." + This is the number of annihilations per orbit times the number of orbital periods 1n a Hubble time: we have defined 0 20as apocenter and P is the orbital period., This is the number of annihilations per orbit times the number of orbital periods in a Hubble time: where we have defined $\theta = 0$ as apocenter and $P$ is the orbital period. + A WIMP can only annihilate once. so the density profile is valid at present only if (Nan) <1.," A WIMP can only annihilate once, so the density profile is valid at present only if $\mean{N_{\rm ann}} < 1$." +"Assuming the orbit is Keplerian and nearly radial. where r, 1s the apocenter radius and v, is the WIMP's tangential velocity at Using Eq."," Assuming the orbit is Keplerian and nearly radial, where $r_a$ is the apocenter radius and $v_a$ is the WIMP's tangential velocity at Using Eq." +" 1. to find Roy and from the bulk velocity dispersion in Ricottt (2009).. we find Therefore. WIMPs with apocenters of R., mostly survive to the present. despite venturing much further in."," \ref{eqn:Rtrz1000} to find $R_{\rm eq}$ and from the bulk velocity dispersion in \citet{Ricotti09}, we find Therefore, WIMPs with apocenters of $R_{\rm eq}$ mostly survive to the present, despite venturing much further in." + The accretion and resultant WIMP orbits at Λο must be nearly radial: σρνΝv/2GMyanRoy., The accretion and resultant WIMP orbits at $R_{\rm eq}$ must be nearly radial: $\sigma_{\rm DM} \ll \sqrt{2 G M_{\rm PBH} / R_{\rm eq}}$. + Using vy (Eq. 4» , Using $v_a$ (Eq. \ref{eqn:vTanDM}) ) +for the initial velocity dispersion apy. we find accretion is radial for all reasonable PBH masses.," for the initial velocity dispersion $\sigma_{\rm DM}$, we find accretion is radial for all reasonable PBH masses." +" WIMPs also have residual thermal dispersion. vinenn71.3Επ. conservatively. we find radial accretion at Ry, holds for Mppu2107muSoy)7M..."," WIMPs also have residual thermal dispersion, $v_{\rm therm} \approx 1.3\ \cm\ \sec^{-1} m_{100}^{-1} (z / z_{\rm eq})^{-1}$; conservatively, we find radial accretion at $R_{\rm eq}$ holds for $M_{\rm PBH} \ga 2 \times 10^{-15} m_{100}^{-3} (z / z_{\rm eq})^{-3} \Msun$." + Accretion at later times than z3000. at smaller radius than A. or from the low end of the WIMP velocity distribution might still form a UCMH: so our limits may roughly hold for smaller PBHs.," Accretion at later times than $z \approx 3000$, at smaller radius than $R_{\rm eq}$ , or from the low end of the WIMP velocity distribution might still form a UCMH; so our limits may roughly hold for smaller PBHs." + Finally. WIMPs must miss the central PBH.," Finally, WIMPs must miss the central PBH." +" WIMPs with an apocenter r, have a pericenter rj~VLL(2GMppuy): if the apocenter is Rey. then Thus. WIMPs with apocenters of A, remain outside the Schwarzschild radius ifMppy=1740M ..."," WIMPs with an apocenter $r_a$ have a pericenter $r_p \approx v_a^2 r_a^2 / (2 G M_{\rm PBH})$; if the apocenter is $R_{\rm eq}$, then Thus, WIMPs with apocenters of $R_{\rm eq}$ remain outside the Schwarzschild radius if$M_{\rm PBH} \la 1740\ \Msun$ ." +" The UCMH annihilation luminosity is) pes3zgrn(r).--toavinmpwicodr. or where fyi, Is à MmMimum-radius cutoff."," The UCMH annihilation luminosity is $L_{\rm ann} = \int_{\rm r_{\rm min}}^{R_{\rm eq}} 2 \pi r^2 n(r)^2 \mean{\sigma_A v} m_{\rm DM} c^2 dr$ , or where $r_{\rm min}$ is a minimum-radius cutoff." + Even if WIMPs with small apocenters annihilate away by the present. WIMPs on radial infall orbits with larger apocenters generate a time-averaged density profile of ρε between apocenter and Therefore. we use Mey U8 Dag.," Even if WIMPs with small apocenters annihilate away by the present, WIMPs on radial infall orbits with larger apocenters generate a time-averaged density profile of $\rho \propto r^{-3/2}$ between apocenter and Therefore, we use $r_{\rm p,eq}$ as $r_{\rm min}$." + The UCMH luminosity is ignoring a very small logarithmic term in Mppgy., The UCMH luminosity is ignoring a very small logarithmic term in $M_{\rm PBH}$. +" The linear sealing with Mpgy arises because the UCMH mass scales linearly with PBH mass. and the WIMP density (and annihilation lifetime) at Ro, is always the same."," The linear scaling with $M_{\rm PBH}$ arises because the UCMH mass scales linearly with PBH mass, and the WIMP density (and annihilation lifetime) at $R_{\rm eq}$ is always the same." + When px r2. each decade in minihalo radius produces the same annihilation luminosity: the total luminosity therefore is not strongly dependent on the WIMP orbits’ eccentricities.," When $\rho \propto r^{-3/2}$ , each decade in minihalo radius produces the same annihilation luminosity; the total luminosity therefore is not strongly dependent on the WIMP orbits' eccentricities." + The eccentricities themselves depend weakly on Mppy. since Facog!l'peqM PRHÁ ," The eccentricities themselves depend weakly on $M_{\rm PBH}$, since $r_{\rm a, eq} / r_{\rm p, eq} \propto M_{\rm PBH}^{0.1}$ ." +Thus the total lummosity is understood as the mass of the UCMH halo (proportional to. Mppy) annihilating over a annihilation timescale (set. by ny) Multiplied by some slowly varying logarithmic factor accounting for the UCMH's inner regions., Thus the total luminosity is understood as the mass of the UCMH halo (proportional to $M_{\rm PBH}$ ) annihilating over a annihilation timescale (set by $n_{\rm tr}$ ) multiplied by some slowly varying logarithmic factor accounting for the UCMH's inner regions. +" A steeper density profile inside A4, would increase the lummosity but shorten the WIMP survival time.", A steeper density profile inside $R_{\rm eq}$ would increase the luminosity but shorten the WIMP survival time. + Eq., Eq. +" 2 implies. for WIMPs with απο density profile and r,=Ra. (Nagy)ΗΕ Ρο&70104,Myr."," \ref{eqn:Nann} implies, for WIMPs with a $r^{-9/4}$ density profile and $r_a = R_{\rm eq}$, $\mean{N_{\rm ann}} = 1500 m_{100}^{-1} (M_{\rm PBH} / \Msun)^{0.08}$, so $t_{\rm ann} \approx 7 m_{100}^{-1}\ \Myr$." + For ar profile. (Nan?=22hMpΜ.Ο (zuzx10 when Hoo= D.," For a $r^{-2}$ profile, $\mean{N_{\rm ann}} = 22 m_{100}^{-1} (M_{\rm PBH} / \Msun)^{0.05}$ $z_{\rm ann} \approx 10$ when $m_{100} = 1$ )." + Efficient annihilation introduces ~Ὁρμῃ worth of high energy radiation into the Universe., Efficient annihilation introduces $\sim \Omega_{\rm PBH}$ worth of high energy radiation into the Universe. + At high redshift. the CMB energy spectrum and reionization history severely constrain the gamma-ray and cosmic ray injection rate (e.g.. 2006)..," At high redshift, the CMB energy spectrum and reionization history severely constrain the gamma-ray and cosmic ray injection rate \citep[e.g.,][]{Fixsen96,Padmanabhan05,Mapelli06}. ." + At low enough redshift. the gamma-ray and neutrino backgrounds still constrain the annihilation products directly.," At low enough redshift, the gamma-ray and neutrino backgrounds still constrain the annihilation products directly." + Thus. it wouldbe difficult to have Oppy7 Qwimp- unlessperhaps the annihilation were entirely into neutrinos.," Thus, it wouldbe difficult to have $\Omega_{\rm PBH} \approx \Omega_{\rm WIMP}$ , unlessperhaps the annihilation were entirely into neutrinos." + Furthermore. suchefficient annihilation would mean that present UCMHs around PBHs have depleted inner halos. with consequencesfor microlensing searches (Ricottt&Gould 2009). ," Furthermore, suchefficient annihilation would mean that present UCMHs around PBHs have depleted inner halos, with consequencesfor microlensing searches \citep{Ricotti09}. ." +More work on the effects of annihilation m steep density profiles is needed., More work on the effects of annihilation in steep density profiles is needed. +onto an absolute scale by requiring the sum of the 303.78 aand 303.32 qquiet-Sun intensities to match the values reported by Mangoetal.(1978) based on earlier observations from satellites and sounding rockets.,onto an absolute scale by requiring the sum of the 303.78 and 303.32 quiet-Sun intensities to match the values reported by \citet{man78} based on earlier observations from satellites and sounding rockets. +" In Mangoetal.(1978) the summed intensity is 7115 ergs cm"" ?s- !sr-! at disk center and brightens by about toward the limb.", In \citet{man78} the summed intensity is 7115 ergs $^{-2}$ $^{-1}$ $^{-1}$ at disk center and brightens by about toward the limb. + The absolute radiometric calibration of SERTS derived this way was estimated to be accurate within a factor better than 2., The absolute radiometric calibration of SERTS derived this way was estimated to be accurate within a factor better than 2. +" The version of SERTS flown in 1997 (SERTS-97;Bro-siusetal.2000a,b) incorporated the same multilayer- toroidal diffraction grating that was flown in 1991 and 1993, but recorded spectrographic data on an intensified CCD-detector."," The version of SERTS flown in 1997 \citep[SERTS-97;][]{bro00a,bro00b} incorporated the same multilayer-coated toroidal diffraction grating that was flown in 1991 and 1993, but recorded spectrographic data on an intensified CCD-detector." +" Its spectral bandpass was 299 — 353À,, with an instrumental resolution (FWHM) of 115mÁ."," Its spectral bandpass was 299 – 353, with an instrumental resolution (FWHM) of 115." +". An end-to-end radiometric calibration of SERTS-97 was performed after the flight at Rutherford-Appleton Laboratory (RAL), UK, in the same facility used to calibrate the Coronal Diagnostic Spectrometer (CDS;Harrisonetal.1995) on theObservatory spacecraft and the Extreme-ultraviolet Imaging (SOHO)Spectrometer (EIS;Cul-haneetal.2007) aboard theHinode satellite (Kosugietal.2007), and using the same EUV light source re-calibrated by Physikalisch-Technische Bundesanstalt (PTB) against the synchrotron radiation of BESSY I (a primary radiation standard)."," An end-to-end radiometric calibration of SERTS-97 was performed after the flight at Rutherford-Appleton Laboratory (RAL), UK, in the same facility used to calibrate the Coronal Diagnostic Spectrometer \citep[CDS;][]{har95} on the ) spacecraft and the Extreme-ultraviolet Imaging Spectrometer \citep[EIS;][]{cul07} aboard the satellite \citep{kos07}, and using the same EUV light source re-calibrated by Physikalisch-Technische Bundesanstalt (PTB) against the synchrotron radiation of BESSY I (a primary radiation standard)." + This was the first time that the SERTS instrument underwent an absolute radiometric calibration while fully assembled., This was the first time that the SERTS instrument underwent an absolute radiometric calibration while fully assembled. + 'The uncertainty on the SERTS-97 absolute radiometric calibration was15%., The uncertainty on the SERTS-97 absolute radiometric calibration was. +". Coordinated, cospatial, time-invariant SERTS-97 and CDS spectra were used to carry out an intensity cross-calibration that yielded an improved responsivity curve for the CDS Normal Incidence Spectrometer’s NIS 1 (308 — 380 À)) waveband (Thomasetal.1999;Thomas2002).."," Coordinated, cospatial, time-invariant SERTS-97 and CDS spectra were used to carry out an intensity cross-calibration that yielded an improved responsivity curve for the CDS Normal Incidence Spectrometer's NIS 1 (308 – 380 ) waveband \citep{tho99, tho02}." +" The importance of a reliable radiometric calibration for SERTS, CDS, EIS, or any EUV instrument cannot be over emphasized."," The importance of a reliable radiometric calibration for SERTS, CDS, EIS, or any EUV instrument cannot be over emphasized." +" Without it, quantitative analyses of properties like temperature, density, emission measure, and element abundances, along with the physical interpretations that emerge from studies of those properties, are impossible."," Without it, quantitative analyses of properties like temperature, density, emission measure, and element abundances, along with the physical interpretations that emerge from studies of those properties, are impossible." + The Extreme Ultraviolet Normal Incidence Spectrograph| (EUNIS) is the successor to SERTS., The Extreme Ultraviolet Normal Incidence Spectrograph (EUNIS) is the successor to SERTS. +" EUNIS contains two independent but co-pointing spectrographs (each with a normal incidence one of which covers a short wavelength (SW;telescope), 170-205 À)) channel while the other covers a long wavelength 300-370 À)) channel"," EUNIS contains two independent but co-pointing spectrographs (each with a normal incidence telescope), one of which covers a short wavelength (SW; 170–205 ) channel while the other covers a long wavelength (LW; 300–370 ) channel." + The instrument and its end-to-end(LW; absolute radiometric calibration is described below., The instrument and its end-to-end absolute radiometric calibration is described below. +" Spectra obtained from the 2006 flight (EUNIS-06) were used to investigate a bright point (Brosius,Rabin,&Thomas2007;Brosiusetal. 2008),, a cool transient brightening (Brosius,Rabin,&Thomas2008), and transition region velocity oscillations (Jessetal.2008),, as well as to derive a calibration update for CDS NIS 1 (Wangetal.2010)."," Spectra obtained from the 2006 flight (EUNIS-06) were used to investigate a bright point \citep{bro07, bro08b}, a cool transient brightening \citep{bro08a}, and transition region velocity oscillations \citep{jes08}, as well as to derive a calibration update for CDS NIS 1 \citep{wan10}." +". EUNIS was last flown on 2007 November 6, when the solar disk contained no active regions."," EUNIS was last flown on 2007 November 6, when the solar disk contained no active regions." + EUNIS-07 obtained coordinated observations of quiet-Sun areas near disk center with both CDS and EIS., EUNIS-07 obtained coordinated observations of quiet-Sun areas near disk center with both CDS and EIS. + In what follows we derive updates to the CDS and EIS absolute radiometric calibrations based on these coordinated observations., In what follows we derive updates to the CDS and EIS absolute radiometric calibrations based on these coordinated observations. + Section 2 describes the observations and data reduction; 83 discusses the absolute radiometric calibration of the EUNIS-06 and EUNIS-07 long wavelength channel; in 84 we derive the calibration update for CDS NIS with EUNIS-07; 85 presents the absolute radiometric calibration of the EUNIS-07 SW channel; in 86 we derive the calibration update for both EIS channels with EUNIS-07; and in 87 we discuss and summarize our conclusions., Section 2 describes the observations and data reduction; 3 discusses the absolute radiometric calibration of the EUNIS-06 and EUNIS-07 long wavelength channel; in 4 we derive the calibration update for CDS NIS with EUNIS-07; 5 presents the absolute radiometric calibration of the EUNIS-07 SW channel; in 6 we derive the calibration update for both EIS channels with EUNIS-07; and in 7 we discuss and summarize our conclusions. +philosophy. particularly with regard to the amount of flux included in identified clumps. as discussed at length by ?..,"philosophy, particularly with regard to the amount of flux included in identified clumps, as discussed at length by \citet{Rosolowsky06}." + Before even identifying individual GMCs. many studies of extragalactic clouds suffer from the more fundamental problem of missing zero spacing information.," Before even identifying individual GMCs, many studies of extragalactic clouds suffer from the more fundamental problem of missing zero spacing information." +" Observing the gas distribution of a galaxy with only an interferometer neglects extended flux on the largest scales. where It is ""resolved out” by the beam of an interferometric array."," Observing the gas distribution of a galaxy with only an interferometer neglects extended flux on the largest scales, where it is “resolved out"" by the beam of an interferometric array." + This effect can be quite significant (?).., This effect can be quite significant \citep{Koda11}. + To investigate GMC evolution in galactic disks and resolve the physics which controls the star formation rate within GMCs. we will present resolved observations of GMCs across the disks of nearby galaxies in a CO Survey of Nearby Galaxies based on observations using the CARMA interferometer and the Nobeyama Radio Observatory 45-meter (NRO45) single dish telescopes.," To investigate GMC evolution in galactic disks and resolve the physics which controls the star formation rate within GMCs, we will present resolved observations of GMCs across the disks of nearby galaxies in a CO Survey of Nearby Galaxies based on observations using the CARMA interferometer and the Nobeyama Radio Observatory 45-meter (NRO45) single dish telescopes." + To date. 13 galactic disks have been observed with both CARMA and NRO45 as part of the survey. and 4 additional galaxies have been observed with the single dish.," To date, 13 galactic disks have been observed with both CARMA and NRO45 as part of the survey, and 4 additional galaxies have been observed with the single dish." + The design of the survey will be discussed in Koda et al. (, The design of the survey will be discussed in Koda et al. ( +in prep).,in prep). + In our survey. we will resolve individual GMCs in a significant sample of nearby spiral galaxies with a variety of morphologies in order to study the evolution of molecular clouds and star formation.," In our survey, we will resolve individual GMCs in a significant sample of nearby spiral galaxies with a variety of morphologies in order to study the evolution of molecular clouds and star formation." + In this paper. we combine observations from CARMA and NRO45 in order to achieve high resolution. as well as extremely high image fidelity. and resolve individual extragalactic GMCs in NGC 6946 to highlight the results made possible by our survey.," In this paper, we combine observations from CARMA and NRO45 in order to achieve high resolution, as well as extremely high image fidelity, and resolve individual extragalactic GMCs in NGC 6946 to highlight the results made possible by our survey." + In Section II we describe the observations. and we present the results in Section HI as well as à discussion of the boundaries of individual GMCs.," In Section II we describe the observations, and we present the results in Section III as well as a discussion of the boundaries of individual GMCs." + We utilize the well known CLUMPFIND algorithm (2). to derive sizes and velocity dispersions for individual GMCs., We utilize the well known $\sc{CLUMPFIND}$ algorithm \citep{Williams94} to derive sizes and velocity dispersions for individual GMCs. + In Section IV. we discuss the properties of our detected GMC sample and determine the conversion factor Xco within individual clouds.," In Section IV, we discuss the properties of our detected GMC sample and determine the conversion factor $X_{\text{CO}}$ within individual clouds." +" We summarize our conclusions in Section V. We observe NGC 6946 in the ""CO (1-0) transition with the Nobeyama 45-meter single dishtelescope"". using the Beam Array Receiver System (BEARS) instrument."," We summarize our conclusions in Section V. We observe NGC 6946 in the $^{12}$ CO (1-0) transition with the Nobeyama 45-meter single dish, using the Beam Array Receiver System (BEARS) instrument." + Observations were performed during the early months of 2008. 2009. and 2010 (throughout our three year observing program at NRO) as part of our CO survey of nearby galaxies.," Observations were performed during the early months of 2008, 2009, and 2010 (throughout our three year observing program at NRO) as part of our CO survey of nearby galaxies." + BEARS is a multi-beam receiver with 25 beams. which are aligned in a 3x5 orientation.," BEARS is a multi-beam receiver with 25 beams, which are aligned in a $\times$ 5 orientation." + The FWHM of the 45-m dish is 15” at 115 GHz (19.7” after regridding). and we observe with channel increments of 500 kHz and Hanning smooth for a velocity resolution of 2.54s!.," The FWHM of the 45-m dish is $\as$ at 115 GHz $\as$ after regridding), and we observe with channel increments of 500 kHz and Hanning smooth for a velocity resolution of 2.54." +. After dropping the edge channels. we use a bandwidth of 265 MHz (690 ," After dropping the edge channels, we use a bandwidth of 265 MHz (690 )." +For most of the scans. the system temperature in the double side band (DSB) of BEARS ranges from 300-400 K: scans with T4; much higher than 400 K (due to being observed during daylight hours) are heavily down-weighted and as a result do not contribute much to the final images.," For most of the scans, the system temperature in the double side band (DSB) of BEARS ranges from 300-400 K; scans with $_{sys}$ much higher than 400 K (due to being observed during daylight hours) are heavily down-weighted and as a result do not contribute much to the final images." + The ratio of the upper-to-lower side band (known as the scale factor) was confirmed each year to be within a few percent by observing Galactic CO sources (e.g.. AGB stars) with BEARS and the single side band (SSB) S100 receiver and making subsequent corrections.," The ratio of the upper-to-lower side band (known as the scale factor) was confirmed each year to be within a few percent by observing Galactic CO sources (e.g., AGB stars) with BEARS and the single side band (SSB) S100 receiver and making subsequent corrections." + We use the S40 receiver and Galactic masers for pointing., We use the S40 receiver and Galactic masers for pointing. +" The pointing was checked roughly every two hours during the observations and was accurate to 2-3"".", The pointing was checked roughly every two hours during the observations and was accurate to $\as$. +" We convert Ty. to T,,6 assuming that the main beam efficiency of the telescope is 0.4 (Le.. Tj; 2 Ty: / 0.4)."," We convert $_{A*}$ to $_{mb}$ assuming that the main beam efficiency of the telescope is 0.4 (i.e., $_{mb}$ = $_{A^*}$ / 0.4)." + Using on-the-fly (OTF) mapping. NGC 6946 was scanned in the RA and Dec directions. and positions external to the galaxy (OFF positions) were observed between scans.," Using on-the-fly (OTF) mapping, NGC 6946 was scanned in the RA and Dec directions, and positions external to the galaxy (OFF positions) were observed between scans." + Extrapolating between OFF scans on opposite sides of the galaxy greatly reduced non-linearities. in the spectral baselines., Extrapolating between OFF scans on opposite sides of the galaxy greatly reduced non-linearities in the spectral baselines. + The duration of each scan (ON + OFF) was -ἰ minute. and the entire galaxy was mapped in ~40 minutes: a total of 36 usable maps were taken for a total of 24 hours of observation time (including ON. OFF. and slew time).," The duration of each scan (ON + OFF) was $\sim$ 1 minute, and the entire galaxy was mapped in $\sim$ 40 minutes; a total of 36 usable maps were taken for a total of 24 hours of observation time (including ON, OFF, and slew time)." +" The scans were separated by 5"". resulting in oversampling by a factor of 3 compared to the 157 FWHM of the beam. which is necessary in order to achieve Nyquist sampling (5.96’) of Aco/D = 11.92” (where co is the observed wavelength. 2.6 mm. and D is the antenna diameter. 45 m)."," The scans were separated by $\as$, resulting in oversampling by a factor of 3 compared to the 15"" FWHM of the beam, which is necessary in order to achieve Nyquist sampling $\as$ ) of $\lambda_{CO}$ /D = $\as$ (where $\lambda_{CO}$ is the observed wavelength, 2.6 mm, and D is the antenna diameter, 45 m)." + The data reduction and sky subtraction. performed by interpolating between the OFF secans for each OTF (ON) scan. were completed using the NOSTAR package developed at the Nobeyama Radio Observatory.," The data reduction and sky subtraction, performed by interpolating between the OFF scans for each OTF (ON) scan, were completed using the NOSTAR package developed at the Nobeyama Radio Observatory." + Spatial baseline-subtracted maps were made separately from the scans in the RA and Dee directions in order to minimize systematic errors in the scan directions. and these were subsequently co-added.," Spatial baseline-subtracted maps were made separately from the scans in the RA and Dec directions in order to minimize systematic errors in the scan directions, and these were subsequently co-added." +" The rms noise of these single dish observations is 0.13 K (0.57 Jy beam""! ).", The rms noise of these single dish observations is 0.13 K (0.57 Jy $^{-1}$ ). + To complement the single dish observations. NGC 6946 was also observed in the 'CO (1-0) transition in April 2009 with the C and D configurations of CARMA.," To complement the single dish observations, NGC 6946 was also observed in the $^{12}$ CO (1-0) transition in April 2009 with the C and D configurations of CARMA." + CARMA is a I5-element interferometer which combines six. DO-meter antennae (originally the Owens Valley Radio Observatory. or OVRO) with nine 6-meter antennae (formerly the Berkeley-Hlinois-Maryland Association. or BIMA) to achieve superior ity-coverage. compared to either of its predecessors.," CARMA is a 15-element interferometer which combines six 10-meter antennae (originally the Owens Valley Radio Observatory, or OVRO) with nine 6-meter antennae (formerly the Berkeley-Illinois-Maryland Association, or BIMA) to achieve superior $\it{uv}$ -coverage compared to either of its predecessors." + The observations of NGC 6946 were performed using three dual side bands. each with 63 channels. for a total bandwidth of ~ 100 MHz (after removing six edge channels per sideband) and a channel width of 2.54.," The observations of NGC 6946 were performed using three dual side bands, each with 63 channels, for a total bandwidth of $\sim$ 100 MHz (after removing six edge channels per sideband) and a channel width of 2.54." +. After a total of ~21 hours on source (including calibrators). we achieve an rms of 0.73 K. In order to combine the data from the single dish and interferometer. the NRO45 image is converted to visibilities. combined with the CARMA visibilities in. the «y-plane. and the new uv dataset is imaged together.," After a total of $\sim$ 21 hours on source (including calibrators), we achieve an rms of 0.73 K. In order to combine the data from the single dish and interferometer, the NRO45 image is converted to visibilities, combined with the CARMA visibilities in the $\it{uv}$ -plane, and the new $\it{uv}$ dataset is imaged together." + We follow the procedure thoroughly described in ?. for the imaging of M51. and we refer the reader to that paper for the details of the NROXHS deconvolution. combination in iv-space. and imaging process.," We follow the procedure thoroughly described in \citet{Koda11} for the imaging of M51, and we refer the reader to that paper for the details of the NRO45 deconvolution, combination in $\it{uv}$ -space, and imaging process." + In that paper. the relative wv-coverage of the single dish and interferometer observations enabled the single dish visibilities to be flagged beyond 4 kt. as CARMA visibilities existed down to this value.," In that paper, the relative $\it{uv}$ -coverage of the single dish and interferometer observations enabled the single dish visibilities to be flagged beyond 4 $\lambda$, as CARMA visibilities existed down to this value." + We keep the ΝΙΚΟΛΑΟΣ visibilities out to 10 kt in order to ensure sufficient overlap between the two sets of uv-coverage., We keep the NRO45 visibilities out to 10 $\lambda$ in order to ensure sufficient overlap between the two sets of $\it{uv}$ -coverage. + The rms of the combined cube is 0.11 Jy beam (1.9 K) using the combined synthesized beam. discussed below.," The rms of the combined cube is 0.11 Jy $^{-1}$ (1.9 K) using the combined synthesized beam, discussed below." +! We maintain the instrumental velocity resolution of 2.54 to optimize our ability to resolve. not only to detect. GMCs.," We maintain the instrumental velocity resolution of 2.54 to optimize our ability to resolve, not only to detect, GMCs." + The observing parameters for both, The observing parameters for both + , +" """, +" ""w", +" ""wa", +" ""wav", +" ""wave", +" ""wave ", +" ""wave i", +" ""wave it", +" ""wave ith", +" ""wave ithe", +" ""wave ithe ", +" ""wave ithe t", +" ""wave ithe ti", +" ""wave ithe tis", +" ""wave ithe tis ", +" ""wave ithe tis p", +" ""wave ithe tis pr", +" ""wave ithe tis pro", +" ""wave ithe tis prog", +" ""wave ithe tis progr", +" ""wave ithe tis progre", +" ""wave ithe tis progres", +" ""wave ithe tis progress", +" ""wave ithe tis progress ", +" ""wave ithe tis progress I", +" ""wave ithe tis progress If", +" ""wave ithe tis progress If ", +" ""wave ithe tis progress If L", +" ""wave ithe tis progress If La", +" ""wave ithe tis progress If Lan", +" ""wave ithe tis progress If Land", +" ""wave ithe tis progress If Land ", +programmes: the SCUILA-2. Cosmology Legacy Survey. S2CLS: and the SCUBA-2 “ALL Sky Survey. SASS (?)..,"programmes: the SCUBA-2 Cosmology Legacy Survey, S2CLS; and the SCUBA-2 `All Sky' Survey, SASSy \citep{Thompson07}." + We have also carried out estimates for surveys at shorter wavelengths. of the sort which might be performed with the 450 yam array of SCUDA-2. or the SPUR instrument on theHoerseheb satellite (operating at 250. 350 and 500jn).," We have also carried out estimates for surveys at shorter wavelengths, of the sort which might be performed with the $450\,\mu$ m array of SCUBA-2, or the SPIRE instrument on the satellite (operating at $250$, $350$ and $500\,\mu$ m)." + While it is clear that there may be many examples of strong lenses in such wide survevs. the fraction of bright sources which are lensed is significantly lower than at longer wavelengths.," While it is clear that there may be many examples of strong lenses in such wide surveys, the fraction of bright sources which are lensed is significantly lower than at longer wavelengths." +" Η one wants to find such ""monsters. either as probes of line-of-sight structure or for their own intrinsic value. then one should turn to ground-based surveys in the 850jim or ~] mm windows."," If one wants to find such `monsters', either as probes of line-of-sight structure or for their own intrinsic value, then one should turn to ground-based surveys in the $850\,\mu$ m or ${\sim}\,1\,$ mm windows." + The S2CLS plans to map approximately 20deg? to an RAIS of OF ms.," The S2CLS plans to map approximately $20\,{\rm deg}^2$ to an RMS of $0.7\,$ mJy." + From our number counts model we estimate that there will be about 96. 44 ancl 27 sources detected with ο720. 25 and 30 mv. respectively.," From our number counts model we estimate that there will be about 96, 44 and 27 sources detected with $S\,{>}\,20$, 25 and $30\,$ mJy, respectively." + The total number of sources above these IHux limits which have sr>>2 will be about 15 (with most of them at the high Εικ end) using our best estimate from the Schechter Function number counts.," The total number of sources above these flux limits which have $\mu\,{>}\,2$ will be about 15 (with most of them at the high flux end) using our best estimate from the Schechter function number counts." + The numbers change only slightlv over this. [lux density range. due to the probability of lensing increasing faster than the decrease in source counts.," The numbers change only slightly over this flux density range, due to the probability of lensing increasing faster than the decrease in source counts." + Above 35 mv the number of expected lensed sources declines steadily.," Above $35\,$ mJy the number of expected lensed sources declines steadily." +" The Iensed fractions are considerably smaller using our ""maximal counts model. as can be seen by referring to Fig. 6."," The lensed fractions are considerably smaller using our `maximal' counts model, as can be seen by referring to Fig. \ref{fig:mmu-combo}." + Phere will of course always be bright sources which have negligible lensing = hence one would like to know how bright to go before the probability of strong lensing is significant., There will of course always be bright sources which have negligible lensing – hence one would like to know how bright to go before the probability of strong lensing is significant. + This can be determined. using Fig. 6.., This can be determined using Fig. \ref{fig:mmu-combo}. + If one is prepared to accept a lin 3 chance ofa source being strongly lensed. then one should. select. sources observed with 5=25 αν.," If one is prepared to accept a 1 in 3 chance of a source being strongly lensed, then one should select sources observed with $S\,{\ga}\,25\,$ mJy." + If. one would like the chance to be 1 in 2. then that [ux rises to about 30 mv.," If one would like the chance to be 1 in 2, then that flux rises to about $30\,$ mJy." + What would one do with such sources in. practice?, What would one do with such sources in practice? + Since strong lensing is likely to come from either a galaxy cluster. or a massive galaxy. then the existence of strong lensing implies strongly. clustered structure along that line of sight.," Since strong lensing is likely to come from either a galaxy cluster, or a massive galaxy, then the existence of strong lensing implies strongly clustered structure along that line of sight." + Hence for cach sulliciently bright source one would use follow-up observations at other wavelengths to try to establish whether strong lensing was likely. and then to see if one could. find the structure which was responsible for the lensing.," Hence for each sufficiently bright source one would use follow-up observations at other wavelengths to try to establish whether strong lensing was likely, and then to see if one could find the structure which was responsible for the lensing." + Multiple images or distorted: morphologies of optical counterparts would be ways of determining that lensing was taking place — these would naturally show up as part of the procedure for trving to determine counterparts in deep data at other wavelengths., Multiple images or distorted morphologies of optical counterparts would be ways of determining that lensing was taking place – these would naturally show up as part of the procedure for trying to determine counterparts in deep data at other wavelengths. +" For SALGs where lensing was stronely suspected. one would target the area to search for the presence of structure along the line of sight. either with X-ray. Sunvaev-Zel'dovich. or ""red. cluster sequence’ observations."," For SMGs where lensing was strongly suspected, one would target the area to search for the presence of structure along the line of sight, either with X-ray, Sunyaev-Zel'dovich or `red cluster sequence' observations." + Phe advantage of this approach is that the high redshift of the SM sources means that it should be feasible to find. cluster (or proto-cluster) lenses at. higher redshifts than are easy to achieve with most other techniques., The advantage of this approach is that the high redshift of the SMG sources means that it should be feasible to find cluster (or proto-cluster) lenses at higher redshifts than are easy to achieve with most other techniques. + Of course. the selection elfects for clusters found in this wav may be complicated to quantify.," Of course, the selection effects for clusters found in this way may be complicated to quantify." + Nevertheless. building up samples of 22 clusters is sullicienthy important. for understanding structure1 formation (as well as constraining dark energy. cte.).," Nevertheless, building up samples of $z\,{>}\,1$ clusters is sufficiently important for understanding structure formation (as well as constraining dark energy, etc.)," + that it is worth using every available method., that it is worth using every available method. + οσον is designed to make a shallow 850 jim map over approximately 4.000deg?. or one tenth of the sky. with the possibility. of extension to a larger area later.," SASSy is designed to make a shallow $850\,\mu$ m map over approximately $4{,}000\,{\rm deg}^2$, or one tenth of the sky, with the possibility of extension to a larger area later." + The RAIS of the maps is planned to be around 30 mv. so that a robust 5e catalogue will have a limit around 150 m.)y.," The RMS of the maps is planned to be around $30\,$ mJy, so that a robust $5\sigma$ catalogue will have a limit around $150\,$ mJy." + H may also be possible to reduce this to nearer to 100 mv. using targetted repeat observations for peaks in the maps.," It may also be possible to reduce this to nearer to $100\,$ mJy using targetted repeat observations for peaks in the maps." + The procedure for identifving strongly lensed: sources in SASSY mav be a little cülferent than for S2C'LS. and because of the brighter Dux. densities. the level of uncertainty in predictions of the number of lensed sources will be considerably higher.," The procedure for identifying strongly lensed sources in SASSy may be a little different than for S2CLS, and because of the brighter flux densities, the level of uncertainty in predictions of the number of lensed sources will be considerably higher." + Phese predictions. depend strongly on the counts model. the normalization. of the ugh amplification tail of the lensing PDE. as well as the amplification. cut-oll imposed. by the finite source size for SAIGs.," These predictions depend strongly on the counts model, the normalization of the high amplification tail of the lensing PDF, as well as the amplification cut-off imposed by the finite source size for SMGs." + Since there are still huge uncertainties in all of these actors. the expectations for SASSV cover a wide range of »ossibilitics.," Since there are still huge uncertainties in all of these factors, the expectations for SASSy cover a wide range of possibilities." + Using the optimistic limit of 100 αν for SASSY. the unlensecl model counts give approximately 1200. sources or the SASSY catalogue.," Using the optimistic limit of $100\,$ mJy for SASSy, the unlensed model counts give approximately 1200 sources for the SASSy catalogue." +" Many. of these will be in the ""Euclidean counts’ regime. and hence should be relatively easy to eliminate as lensing candidates."," Many of these will be in the `Euclidean counts' regime, and hence should be relatively easy to eliminate as lensing candidates." + These will often be alreacly well-known galaxies. with others typically being in theHIS. or radio catalogues.," These will often be already well-known galaxies, with others typically being in the, or radio catalogues." + Colours (e.g.. 850jim. to radio) can be used to distinguish objects which are likely to be at higher redshift and hence have a higher likelihood of being lensed: such methods will also be necessary to eliminate Galactic clouds.," Colours (e.g., $850\,\mu$ m to radio) can be used to distinguish objects which are likely to be at higher redshift and hence have a higher likelihood of being lensed; such methods will also be necessary to eliminate Galactic clouds." + Our best lensing estimate vields a total of 1000 sources in SASSY with lensing amplification po2. with most of them being much more strongly lensed than this limit.," Our best lensing estimate yields a total of 1000 sources in SASSy with lensing amplification $\mu\,{>}\,2$, with most of them being much more strongly lensed than this limit." + We expect that these extremely. Iensed sources will be fairlv. easy to distinguish from relatively nearby. intrinsically bright galaxies. and hence the chances of there being structure along the line of sight to such candidates will be very high.," We expect that these extremely lensed sources will be fairly easy to distinguish from relatively nearby, intrinsically bright galaxies, and hence the chances of there being structure along the line of sight to such candidates will be very high." + Follow-up of these SMCs at other wavelengths. will also be easy. since they should. be ab least an order of magnitude brighter than the typical SCUBA sources which have been followed up in the past.," Follow-up of these SMGs at other wavelengths will also be easy, since they should be at least an order of magnitude brighter than the typical SCUBA sources which have been followed up in the past." + We would like to thank Alexandra Pope for many helpful discussions curing the course of this work., We would like to thank Alexandra Pope for many helpful discussions during the course of this work. + Neal Dalal and Mattia Negrello provided useful feedback on an early draft of this paper., Neal Dalal and Mattia Negrello provided useful feedback on an early draft of this paper. + We also thank Ranga-Ram Chary for providing his submillimetre galaxy. models to us., We also thank Ranga-Ram Chary for providing his submillimetre galaxy models to us. + The models of ? are [ουν availableonline!., The models of \citeauthor{Lagache04} are freely available. +. This work was supported by the Natural Sciences and. Engineering Research Council of Canacla., This work was supported by the Natural Sciences and Engineering Research Council of Canada. + Phe evolutionary model used to determine ccounts as a function. of redshift combines the local measurement of the joint Far-H1t.— colour-Iuminosity clistribution. GL.C) (Lis the far-IR. luminosity. and Cis the logarithm of the," The evolutionary model used to determine counts as a function of redshift combines the local measurement of the joint far-IR colour-luminosity distribution, $\Phi(L,C)$ $L$ is the far-IR luminosity, and $C$ is the logarithm of the" +2001. and 0.5 1.5 aresee in November 2001 (as measured by the ESO Dillerential Image Motion. Monitor).,"2001, and $0.5$ $1.5$ arcsec in November 2001 (as measured by the ESO Differential Image Motion Monitor)." + All the spectra were bias subtracted. Hatfield corrected. cleaned of cosmic ravs. corrected for bad pixels and columns. and wavelength. calibrated. using stancarel routines as in Paper |. We checked that the wavelength rebinning was done properly by measuring the dillerence between the measured ancl predicted wavelengths (Osterbrock et al.," All the spectra were bias subtracted, flatfield corrected, cleaned of cosmic rays, corrected for bad pixels and columns, and wavelength calibrated using standard routines as in Paper I. We checked that the wavelength rebinning was done properly by measuring the difference between the measured and predicted wavelengths (Osterbrock et al." + 1996) for the brightest night-sky emission lines in the observed spectral ranges., 1996) for the brightest night-sky emission lines in the observed spectral ranges. + The resulting accuracy in the wavelength calibration is better than 2, The resulting accuracy in the wavelength calibration is better than 2. + The spectra taken along the same axis for the same ealaxy were co-added. using the center of. the. stellar continuum as reference., The spectra taken along the same axis for the same galaxy were co-added using the center of the stellar continuum as reference. + The contribution of the sky was determined from the outermost ~30 aresee at the two edges of the resulting spectra. where the galaxy light was negligible. and. then subtracted. giving a sky subtraction better than 1 per cent.," The contribution of the sky was determined from the outermost $\sim30$ arcsec at the two edges of the resulting spectra, where the galaxy light was negligible, and then subtracted, giving a sky subtraction better than 1 per cent." + A one-dimensional skv-subtracted spectrum was obtained for cach kinematical template star., A one-dimensional sky-subtracted spectrum was obtained for each kinematical template star. + We measured. the stellar kinematics from the galaxy, We measured the stellar kinematics from the galaxy +data is spurious.,data is spurious. +" Rest-frame galaxy colors aid stellar masses, Which depeud ou the adopted stellar population models. plotometiv. and only coarsely on redshift. are cutirev independent of the cuviroument neasPOs, which depchd upon aueular position aud high-precision redshift iuforinatiou (both for the galaxy in quesloni ixd neleiborius galaxies)."," Rest-frame galaxy colors and stellar masses, which depend on the adopted stellar population models, photometry, and only coarsely on redshift, are entirely independent of the environment measures, which depend upon angular position and high-precision redshift information (both for the galaxy in question and neighboring galaxies)." + There is no reasonable mechanisms that would produce a false correlation between hese independent quantities., There is no reasonable mechanisms that would produce a false correlation between these independent quantities. +" For instance. 6je inieh sugecst that contamination of photomoetrv is an lse in dense regions: however. even iu the most overdeuse environments (c.e.. the top 5% of the euvironmeut distribution). the typical distance to the πο... neighbor corresponds to ~35"" ou the sky, which is much larger than the aperture sizes used in photometry."," For instance, one might suggest that contamination of photometry is an issue in dense regions; however, even in the most overdense environments (e.g., the top $5\%$ of the environment distribution), the typical distance to the $3^{\rm rd}$ -nearest neighbor corresponds to $\sim \! +35^{\prime\prime}$ on the sky, which is much larger than the aperture sizes used in photometry." + Instead. as highlelted iu Section l.. it is far more liselv that the color-deusity relation apparent iu DEEP? has been sincared out in studies using otler data sets due to the smaller sample sizes curyploved aud the sieuificautN larger errors in the enviroment measures derived from. trose data sets.," Instead, as highlighted in Section \ref{sec_intro}, it is far more likely that the color-density relation apparent in DEEP2 has been smeared out in studies using other data sets due to the smaller sample sizes employed and the significantly larger errors in the environment measures derived from those data sets." + Many of the Iuuitations regarding the VVDS data set. as they relate to detectingc» correlations between Ooealaxy xoperties and enuvironmoeut ⋜↧↑−∙∿↕∙ aro cliscussed at length bv là a Colparisou o the work bvG).," Many of the limitations regarding the VVDS data set, as they relate to detecting correlations between galaxy properties and environment at $z +\sim 1$, are discussed at length by in a comparison to the work by." +. Usine DEEP2 data. find a significaut. correlation )etween rest-frame ealaxy color and environment at Mp20.55.logo(1) and 0.9ουν 1one is evident iu the analysis of VVDS data preseuted o»(2006).," Using DEEP2 data, find a significant correlation between rest-frame galaxy color and environment at $M_B < -20.5 - 5 \cdot \log_{10}(h)$ and $0.9 < z < +1.2$, where none is evident in the analysis of VVDS data presented by." +.. While boh of these studies eniployed Iuuinosity-seected galaxy saiuples. the lessons earned from the comparison are no less applicable to analyses of samples selected: accordiie fo stellar nass.," While both of these studies employed luminosity-selected galaxy samples, the lessons learned from the comparison are no less applicable to analyses of samples selected according to stellar mass." + lere. we review the arguments presened by as they relate to the more recent analvses usiug VVDS. zCOSAIOS. and other data sts.," Here, we review the arguments presented by as they relate to the more recent analyses using VVDS, zCOSMOS, and other data sets." + At ;z0.75. the sample size o hoth. VVDS and ZCOSAIOS ave significantly smaller han that collected bv DEEP2.," At $z \gtrsim 0.75$, the sample size of both VVDS and zCOSMOS are significantly smaller than that collected by DEEP2." + At these redshifts. DEEP2 las more thau," At these redshifts, DEEP2 has more than" +when massive stars with short Kelvin times begin to dominate the bolometric output of the stellar population.,when massive stars with short Kelvin times begin to dominate the bolometric output of the stellar population. +" In our simulations, even though we do produce 20 Mo stars with significant internal luminosities toward the end of the simulations, accretion luminosity is the dominant energy source over most of the simulation time."," In our simulations, even though we do produce $\sim 20$ $\msun$ stars with significant internal luminosities toward the end of the simulations, accretion luminosity is the dominant energy source over most of the simulation time." +" 'This morphology of small regions of warm gas strung out along filaments continues to hold to some extent even at time t/t¢=0.6, when the stellar mass has increased to a few percent of the gas mass."," This morphology of small regions of warm gas strung out along filaments continues to hold to some extent even at time $t/t_{\rm ff} = 0.6$, when the stellar mass has increased to a few percent of the gas mass." +" We can still identify distinct heated regions associated with individual stars or small stellar groups, and the bulk of the mass remains near 10 K. In the last two time slices, however, as a larger and larger fraction of the cloud mass is converted into stars, this ceases to be true."," We can still identify distinct heated regions associated with individual stars or small stellar groups, and the bulk of the mass remains near 10 K. In the last two time slices, however, as a larger and larger fraction of the cloud mass is converted into stars, this ceases to be true." +" Even the coldest gas anywhere in the cloud is now at temperatures noticeably larger than the original background temperature, and the regions of very warm gas, ΤΑ’100 K, are beginning to overlap and merge."," Even the coldest gas anywhere in the cloud is now at temperatures noticeably larger than the original background temperature, and the regions of very warm gas, $T\gtsim 100$ K, are beginning to overlap and merge." +" In the last time slice, the coldest gas anywhere in the computational domain is at ~30 K, and much of the mass is concentrated in a few compact regions where the temperature is significantly higher."," In the last time slice, the coldest gas anywhere in the computational domain is at $\sim 30$ K, and much of the mass is concentrated in a few compact regions where the temperature is significantly higher." +" Rather than a few warm, dense regions around individual stars (cf. the bulk of the gas is now concentrated into a smaller ?)number of more massive regions that are heated by the collective effects of large numbers of stars."," Rather than a few warm, dense regions around individual stars \citep[cf.][]{offner09a} the bulk of the gas is now concentrated into a smaller number of more massive regions that are heated by the collective effects of large numbers of stars." + Figure 4 shows the total mass of all stars as a function of time in the runs., Figure \ref{starhist1} shows the total mass of all stars as a function of time in the runs. +" Examining the figure shows that the total mass in stars is nearly identical in the two radiative runs, indicating that this aspect of the simulations is very well converged."," Examining the figure shows that the total mass in stars is nearly identical in the two radiative runs, indicating that this aspect of the simulations is very well converged." +" Run ISO begins to form stars somewhat earlier, and the mass in stars present at equal times is somewhat higher."," Run ISO begins to form stars somewhat earlier, and the mass in stars present at equal times is somewhat higher." +" However, this difference mostly appears to be a time offset."," However, this difference mostly appears to be a time offset." +" T'he overall shape in Figure 4 is the same, indicating a generally similar star formation history."," The overall shape in Figure \ref{starhist1} is the same, indicating a generally similar star formation history." +" The time offset is likely a result of the faster collapse that occurs in the isothermal run, where cooling is assumed to be infinitely rapid and efficient, compared to the radiative run."," The time offset is likely a result of the faster collapse that occurs in the isothermal run, where cooling is assumed to be infinitely rapid and efficient, compared to the radiative run." + Figure 5 shows the number of stars as a function of the total stellar mass in each simulation., Figure \ref{starhist2} shows the number of stars as a function of the total stellar mass in each simulation. +" The total number of stars is somewhat larger in run HR than in run LR,"," The total number of stars is somewhat larger in run HR than in run LR," +"We find that the subhalo profiles are ""anti-biased' relative to the dark matter in the inner regions of the haloes.",We find that the subhalo profiles are `anti-biased' relative to the dark matter in the inner regions of the haloes. + This agrees with the results of 2)., This agrees with the results of \citet{ghigna2}. +. Surprisingly we also find that the radial number density profiles are steeper in low mass haloes than in high mass haloes. a finding that reserves further investigation.," Surprisingly we also find that the radial number density profiles are steeper in low mass haloes than in high mass haloes, a finding that reserves further investigation." + We now use our highest resolution cluster simulation to investigate whether subhalos of dillerent mass have cilferent radial profiles., We now use our highest resolution cluster simulation to investigate whether subhalos of different mass have different radial profiles. + In Fig. 7..," In Fig. \ref{fig:fig5}," + we show the cumulative fraction of substructures as a function of 2/94 lor subbhalos with Aus>0.01Mono (solid line) and Ma0.01Moss (clashecl line).," we show the cumulative fraction of substructures as a function of $R/R_{200}$ for subhalos with $M_{\rm sub}>0.01\,M_{200}$ (solid line) and $M_{\rm sub}\le 0.01\,M_{200}$ (dashed line)." + As Fig., As Fig. + 1 already made clear. there are many more substructures with Ad0.01\,M_{200}$ $9749$ versus $96$ )." + Fig., Fig. + 7. now shows that more massive substructures are preferentially located in the external regions of their parent haloes., \ref{fig:fig5} now shows that more massive substructures are preferentially located in the external regions of their parent haloes. + This can be understood. as a consequence of tidal truncation and stripping elfects that quickly decrease the mass of subhalos as they fall into the cluster and reach the dense inner cores ofthe parent haloes (see Section 7.2. [or a more quantitative analysis of mass-Ioss due to stripping).," This can be understood as a consequence of tidal truncation and stripping effects that quickly decrease the mass of subhalos as they fall into the cluster and reach the dense inner cores of the parent haloes (see Section \ref{sec:mah} + for a more quantitative analysis of mass-loss due to stripping)." + Also note that this finding can be naturally explained as à consequence of the orbital decay experienced. by substructures., Also note that this finding can be naturally explained as a consequence of the orbital decay experienced by substructures. + As shown in 7). the orbital cecay ds consistent with expectations based on the combined effects of dvnamical [friction anel mass-loss.," As shown in \citet*{tordia}, the orbital decay is consistent with expectations based on the combined effects of dynamical friction and mass-loss." + As a result. massive substructures are driven to the centre more rapidly than," As a result, massive substructures are driven to the centre more rapidly than" +given by the Lagrangian displacement ὁ grows if the associated potential energy change is negative.,given by the Lagrangian displacement ${\vec \xi}$ grows if the associated potential energy change is negative. + The potential energy change is given by (Bernsteinetal.1958) W=- |dx where Q=V«(£«B) is the magnetic field perturbation. and the volume integral extends over all space in the radial direction and an averaging over <-direction is assumed.," The potential energy change is given by \citep{Bernstein:58} W = ] where ${\bf Q}\equiv \nabla\times({\vec \xi}\times{\bf B})$ is the magnetic field perturbation, and the volume integral extends over all space in the radial direction and an averaging over $z$ -direction is assumed." +" Following Newcomb(1960).. we consider a perturbation with displacement &2(£,..£5.£-) of the form £- ∖∖⇁∣↴⊖∣⋪⊖≦∣⋅⋖∣⋅⋟⋅≦↙∣↗⋖∣⋅⋟⋅⋅≏↧⋯↿≦⋮⋖∣⋅⋟⋯⋪⊜∣⋪⊜∐∣↑↴∐∏∁⊓∪∏⋋⋂↑↴∣⋪∐↳∐∐⋋↜ ↕∖∫⊜∖∖⇁∁"," Following \citet{Newcomb:60}, we consider a perturbation with displacement ${\vec \xi}=(\xi_r,\xi_\theta,\xi_z)$ of the form = _z) where $\xi_r(r)$, $\xi_\theta(r)$, and $\xi_z(r)$ are real functions of radius." +⋂⋯⇂↴⋋∣⋯∖∖⇁⋋⋔∐↾∏∪∏−∐��↸≣⋋⋝⇁⋯⋯⊜↾⊓∁∣⊃⊜∣↑⋯⋪⇂↴∐⊓⋂∏⋋⋖∣⊔⊐⋯ ↾∣⊺≏∐∏∏∏≣⋯≣∑⊖∏⊤⋅≏∐⋪⊜⋯∁∪⋯∣⊃∣⋪⊜⋋⋋⋔∣⊜⋅∇∙≦∶↭⋅⋅≏↧⋯⊓⊺≏↧∖⇁⊖↑↴⋯⋪⋯ ," Newcomb shows that non-axisymmetric perturbations $(m\neq 0)$ that minimize $W$ are incompressible, $\nabla\cdot{\vec \xi}=0$ , and have form _r = = ) - ] _z = )." +"Note that the function £. identical to the radial displacement £,. has been left unspecified."," Note that the function $\xi$, identical to the radial displacement $\xi_r$, has been left unspecified." +" For perturbations of the form in equation (5)). the potential energy perturbation in equation (3)) can be expressed in terms of the radial displacement only. ""e kr where + pill22 yo B, ).(7) as shown in equations (16-18) of Neweomb(1960) (note that in our case B.= 0)."," For perturbations of the form in equation \ref{eq:displacement_incompressible}) ), the potential energy perturbation in equation \ref{eq:delta_w}) ) can be expressed in terms of the radial displacement only, W = )^2 + ] dr, where + m^2 ^2 - ) as shown in equations (16-18) of \citet{Newcomb:60} (note that in our case $B_z=0$ )." + The perturbations with m=| are special in that they need not vanish at r2O and do not incur substantial cost in. magnetic energy as the field lines are bent only minimally., The perturbations with $m=1$ are special in that they need not vanish at $r=0$ and do not incur substantial cost in magnetic energy as the field lines are bent only minimally. + Therefore one is allowed to assume that the radial displacement is independent of radius. d¢/dr=0. overthe region with the non-vanishing magnetic field. which greatly simplifies the analysis.," Therefore one is allowed to assume that the radial displacement is independent of radius, $d\xi/dr = 0$, overthe region with the non-vanishing magnetic field, which greatly simplifies the analysis." + The true fastest growing mode may not have constant ¢. but one expects the constant & approximation to come close.," The true fastest growing mode may not have constant $\xi$, but one expects the constant $\xi$ approximation to come close." +" For these perturbations. the radial displacements are ο 2o ¢,y", For these perturbations the radial displacements are _r = = = i ) _z = . +",With this. the energy per unit length along the direction ayybecomes W=-7 Aj (9)"," With this, the energy per unit length along the $z$ -direction becomes W = k^2 ^2 dr." + Equation (3)) tells us that all wave numbers &>O are unstable., Equation \ref{eq:energy_solid_shift}) ) tells us that all wave numbers $k>0$ are unstable. + In the long and the short dp limits. the energy can approximately be written Ww Wx (lO)where By is. the RMS magnetic field. and Bo is the magnetic RMS field averaged per unit log-radius.," In the long and the short wavelength limits, the energy can approximately be written as W where $\bar B_\theta$ is the RMS magnetic field, and $\tilde B_\theta$ is the magnetic RMS field averaged per unit log-radius." + Equation (3)) tells us that all ;21 modes with finite wavelengths (4>0) are unstable., Equation \ref{eq:energy_solid_shift}) ) tells us that all $m=1$ modes with finite wavelengths $k>0$ ) are unstable. + An estimate of the linear growth rate is given bywe Dwhere ," An estimate of the linear growth rate is given by, where _z^2) dr." +The value of K for perturbations in equations (8)) equals ΕΙ=ΤΠΕ Where pis the average= mass density., The value of $K$ for perturbations in equations \ref{eq:xi_solid}) ) equals (k^2R^2+1)] where $\bar\rho$ is the average mass density. + In the long and the short wavelength limit. therefore. Note that the growth rates are proportional to the “nonrelativistic” Alfvénn velocity v4=By/4p multiplied by kand ΑΙ. respectively. in the long and the short wavelength limit.," In the long and the short wavelength limit, therefore, Note that the growth rates are proportional to the “nonrelativistic” Alfvénn velocity $\bar v_{\rm A}\equiv +B_\theta / \sqrt{4\pi \bar \rho}$ multiplied by $k$and $R^{-1}$, respectively, in the long and the short wavelength limit." + This 1s expected since the instability is driven by a pressure imbalance., This is expected since the instability is driven by a pressure imbalance. + The growth rate in the long wavelength limit can also be expressed as D—(2)7 HER. where as before wp is the plasma frequency.," The growth rate in the long wavelength limit can also be expressed as $\Gamma\sim (\pi/2)^{1/2}\beta_\parallel kR \omega_{\rm p}$ , where as before $\omega_{\rm p}$ is the plasma frequency." + Note that the magnetic field acquires a component parallel to the axis of the filament., Note that the magnetic field acquires a component parallel to the axis of the filament. +" The RMS strength of the parallel field is B.=Q.~1«zB, in the long wavelength limit.", The RMS strength of the parallel field is $\bar B_z=\bar Q_z\sim \onehalf k \xi B_\theta$ in the long wavelength limit. + The above analysis is valid only as long as the perturbation is nonrelativistic (δις« c/D) and the wavelength of the perturbation is larger than the Larmor radius.," The above analysis is valid only as long as the perturbation is nonrelativistic (i.e., $\xi0.," Consider a particle with unperturbed orbit confined to the region $r0$." + The momentum of the particle p is then a periodic function of z. as are its radial excursion and azimuth.," The momentum of the particle ${\bf p}$ is then a periodic function of $z$, as are its radial excursion and azimuth." + The magnetostaticequilibrium is electrically neutral and the energy of the particle. €=5 nc7.where ~=(1+p fnire7y!?. subject to the magnetic field of the equilibrium only. is a constant of motion.," The magnetostaticequilibrium is electrically neutral and the energy of the particle, ${\cal E}=\gamma mc^2$ ,where $\gamma\equiv(1+p^2/m^2c^2)^{1/2}$ , subject to the magnetic field of the equilibrium only, is a constant of motion." + The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law E- B," The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law = ," + The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law E- Ba," The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law = ," + The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law E- Baa," The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law = ," + The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law E- Baas," The time variation of themagnetic field of the unstable perturbation induces an electric field given by Ohm's law = ," +reasons given earlier.,reasons given earlier. + In. the event (hat thev are formed deep within the star. (μον could be decelerated by overlving lavers. but this is unlikely (ο affect. the timing over the large distances and long time intervals accessible to the observations.," In the event that they are formed deep within the star, they could be decelerated by overlying layers, but this is unlikely to affect the timing over the large distances and long time intervals accessible to the observations." + We therelore confie our attention to the jets., We therefore confine our attention to the jets. + These move through the cireiumstellar environment at high velocity. and from a theoretical point of view their kinematic behavior is uncertain. so we consider various possibilities in turn.," These move through the circumstellar environment at high velocity, and from a theoretical point of view their kinematic behavior is uncertain, so we consider various possibilities in turn." + If the ejection of a jet is short-Iived. the jet head may act as a bullet. in which case it would travel at a constant velocity. or decelerate to the extent that it interacts with ambient cireumstellar gas.," If the ejection of a jet is short-lived, the jet head may act as a bullet, in which case it would travel at a constant velocity, or decelerate to the extent that it interacts with ambient circumstellar gas." + The deceleration of Ile 3-1475 has been considered by and is estimated (o be relatively small. alihough the theory of bullet interactions with ambient eas is not hilly developed.," The deceleration of He 3-1475 has been considered by \citet{rie03} and is estimated to be relatively small, although the theory of bullet interactions with ambient gas is not fully developed." + More generally. if the head decelerates. it can be seen by writing the actual travel time as /=fdr/v(r) that the value of /; given by the current (observed) value of r;/e; (or the equivalent. proper motion o/o) overestimates the travel time.," More generally, if the head decelerates, it can be seen by writing the actual travel time as $t = \int dr/v(r)$ that the value of $t_j$ given by the current (observed) value of $r_j/v_j$ (or the equivalent proper motion $\phi/\dot{\phi}$ ) overestimates the travel time." + Therefore. if the ejections of the jets and tori were simultaneous. the apparent ages of the jets would be older (han the tori. an effect opposite to that seen in Fig.," Therefore, if the ejections of the jets and tori were simultaneous, the apparent ages of the jets would be older than the tori, an effect opposite to that seen in Fig." + 1., 1. + We conclude that deceleration does not dominate the observations. and if it were to play any quantitative role. the true jet-lag would be larger than that tabulated in Table 1.," We conclude that deceleration does not dominate the observations, and if it were to play any quantitative role, the true jet-lag would be larger than that tabulated in Table 1." + A different situation occurs if the jet is continuous over an extended interval of time., A different situation occurs if the jet is continuous over an extended interval of time. + To illustrate this case. we consider a jet driven through the circiumstellar gas by a racial. moment(unm-conserving. fast outflow. wilh a narrow. fixed. opening angle.," To illustrate this case, we consider a jet driven through the circumstellar gas by a radial, momentum-conserving, fast outflow, with a narrow, fixed, opening angle." + For an envelope with a density that varies as r..7. which corresponds to a constant mass loss rate. the velocity ol the jet head is a constant to first order (Leeetal.2001).," For an envelope with a density that varies as $r^{-2}$, which corresponds to a constant mass loss rate, the velocity of the jet head is a constant to first order \citep{lee01}." +. IIowever. according to the simulations of Lee&Sahai(2003).. the effects of the motion of material along the sides of the jeis may lead to à small acceleration.," However, according to the simulations of \cite{lee03}, the effects of the motion of material along the sides of the jets may lead to a small acceleration." + The velocity of the jet head may then be written as e=uj(l+al). where uw; is the initial velocity and a is a parameter with a typical value ~10% ! for the simulations thev present.," The velocity of the jet head may then be written as $v = u_j(1+\alpha t)$, where $u_j$ is the initial velocity and $\alpha$ is a parameter with a typical value $\sim +10^{-3}$ $^{-1}$ for the simulations they present." +" Lf the jets and tori were ejected simultaneously. we can identify / with /, and would expect the apparent ages of (he jets to be smaller (han the ages of the tori bv A’ given by: This expression has values of al/2~ 0. 0.25. and 0.5. lor al«1. al=1. and al>1. respectively,"," If the jets and tori were ejected simultaneously, we can identify $t$ with $t_t$ and would expect the apparent ages of the jets to be smaller than the ages of the tori by $\Delta t$ given by: This expression has values of $\alpha t/2 \sim 0$ , 0.25, and 0.5, for $\alpha t \ll 1$, $\alpha t = 1$, and $\alpha t \gg 1$, respectively." + For an ensemble of objects. we then expect M// to range from near zero for voung svstenms. up lo 0.5 for very old systems.," For an ensemble of objects, we then expect ${\Delta t}/{t}$ to range from near zero for young systems, up to 0.5 for very old systems." + The observed values are listed in column (10) of Table 1., The observed values are listed in column (10) of Table 1. + They show the opposite behavior. with the largest values lor the voungest svstenis.," They show the opposite behavior, with the largest values for the youngest systems." + It is nol possible to rule out acceleration in any individual case. but there is no strong evidence for it from the ensemble of data.," It is not possible to rule out acceleration in any individual case, but there is no strong evidence for it from the ensemble of data." + Ii view of this. and theextensive empirical evidence [or," In view of this, and theextensive empirical evidence for" +results become close to those of the scan-based analysis (Fig. 4)).,results become close to those of the scan-based analysis (Fig. \ref{fig:mapsdvlam}) ). + The same effects are observed in the errors for the flux measurements., The same effects are observed in the errors for the flux measurements. +" However, while there is no bias in the positional measurements observed, there is a significant flux loss (see Table 2))."," However, while there is no bias in the positional measurements observed, there is a significant flux loss (see Table \ref{tab:madresid}) )." +" A simple comparison between the data in Tables 1 and 2 as well as between Figures 2 and 3 clearly demonstrates the performance advantages of the scan-based analysis, providing reliable and generally smaller standard errors and a bias-free flux estimate, with no noticeable dependence on the intercept angle (strictly speaking, an intercept angle can only be defined when two sequences of related scans are present)."," A simple comparison between the data in Tables \ref{tab:ringresid} and \ref{tab:madresid} as well as between Figures \ref{fig:ringsdvcomp} and \ref{fig:mapsdvcomp} clearly demonstrates the performance advantages of the scan-based analysis, providing reliable and generally smaller standard errors and a bias-free flux estimate, with no noticeable dependence on the intercept angle (strictly speaking, an intercept angle can only be defined when two sequences of related scans are present)." +" In order to investigate the origin of the rather large standard deviations observed in the map-based solutions, a set of solutions at latitude 0 degrees was made such that the image was always positioned at the centre of a pixel on the map."," In order to investigate the origin of the rather large standard deviations observed in the map-based solutions, a set of solutions at latitude 0 degrees was made such that the image was always positioned at the centre of a pixel on the map." +" The standard deviations of the positional fits are now closer to what is found for the scan data (Fig. 5)),"," The standard deviations of the positional fits are now closer to what is found for the scan data (Fig. \ref{fig:mapfxlam}) )," + though the noise level on the map data is on average higher and repeatability of those data is a lot poorer., though the noise level on the map data is on average higher and repeatability of those data is a lot poorer. + The relation between those standard deviations and the standard errors is in addition still very poor too., The relation between those standard deviations and the standard errors is in addition still very poor too. + The mean standard error estimates following from the least squares solutions of the map data are systematically overestimated and generally very noisy (Fig. 6))., The mean standard error estimates following from the least squares solutions of the map data are systematically overestimated and generally very noisy (Fig. \ref{fig:mapfxsdv}) ). +" Thus, even in these, what may appear to be relatively favourable, but still unrealistic, conditions, the pixel-map derived positional information is still poor, in the prediction of the errors, while the errors themselves are poor estimates of the actual noise on the data."," Thus, even in these, what may appear to be relatively favourable, but still unrealistic, conditions, the pixel-map derived positional information is still poor, in the prediction of the errors, while the errors themselves are poor estimates of the actual noise on the data." +" This can be observed in Fig. 7,,"," This can be observed in Fig. \ref{fig:mapfimeans}," + where the offsets for the recovered positions in longitude and latitude are shown as a function of longitude of the source., where the offsets for the recovered positions in longitude and latitude are shown as a function of longitude of the source. +" For comparison, the results for the scan-based analysis are also shown."," For comparison, the results for the scan-based analysis are also shown." + An example of how the distribution of pixels affects the data accumulated from the scans is shown in Fig. 8.., An example of how the distribution of pixels affects the data accumulated from the scans is shown in Fig. \ref{fig:pixelgrid}. + A summary of these results as well as a comparison with the scan-based analysis is presented in Table 3.., A summary of these results as well as a comparison with the scan-based analysis is presented in Table \ref{tab:mapfresid}. + Here both the actual standard deviation and the unit-weight standard deviation are shown for the same data analysed from pixel maps and directly from scan data., Here both the actual standard deviation and the unit-weight standard deviation are shown for the same data analysed from pixel maps and directly from scan data. +" The conclusion can be drawn that the scan-based analysis is performing more than an order of magnitude better than the pixel-map based analysis, in the reconstruction of all image parameters as well as in the provision of standard errors."," The conclusion can be drawn that the scan-based analysis is performing more than an order of magnitude better than the pixel-map based analysis, in the reconstruction of all image parameters as well as in the provision of standard errors." +" Also in flux reconstruction the results for the pixel-map derived solutions remain poor, as is shown in Fig. 7.."," Also in flux reconstruction the results for the pixel-map derived solutions remain poor, as is shown in Fig. \ref{fig:mapfimeans}." +" There is still a systematic flux loss of on average about 7 per cent, and a noise that is in particular high for small intercept angles."," There is still a systematic flux loss of on average about 7 per cent, and a noise that is in particular high for small intercept angles." +" None of these features shows in the analysis directly made from the scan data, the results of which are also shown in Fig. 7.."," None of these features shows in the analysis directly made from the scan data, the results of which are also shown in Fig. \ref{fig:mapfimeans}." +" Conclusions we can draw from these experiments are that the position of the source centre with respect to the pixels on the pixel map adds a significant source of noise, in particular for small intercept angles."," Conclusions we can draw from these experiments are that the position of the source centre with respect to the pixels on the pixel map adds a significant source of noise, in particular for small intercept angles." +" But even when this “positional freedom” is removed, and the sources are artificially placed at pixel centres, the results remain significantly inferior to those obtained directly from the scan data, in reproduction of the image parameters (with much higher noise and a flux loss for the map-based analysis) and in accuracy of the standard error estimates."," But even when this “positional freedom"" is removed, and the sources are artificially placed at pixel centres, the results remain significantly inferior to those obtained directly from the scan data, in reproduction of the image parameters (with much higher noise and a flux loss for the map-based analysis) and in accuracy of the standard error estimates." +" Considering that we are still dealing here with a circular-symmetric beam, it"," Considering that we are still dealing here with a circular-symmetric beam, it" +planet formation.],planet formation.] +" In general. the orbit-averaged flix experienced by a star/disk system in a given cluster environment increases with the orbital energy ε and decreases with the orbital angular momentum q/q,,,."," In general, the orbit-averaged flux experienced by a star/disk system in a given cluster environment increases with the orbital energy $\epsilon$ and decreases with the orbital angular momentum $q/q_{max}$." + For similar orbital parameters. stars experience higher fIux values when in more populated clusters. clusters wilh steeper density. profiles. aud in clusters with grealer star lormation efficiencies.," For similar orbital parameters, stars experience higher flux values when in more populated clusters, clusters with steeper density profiles, and in clusters with greater star formation efficiencies." + Given the steep nature of the IME. and the relatively modest. membership sizes V considered herein. (he efficiency with which cireumstellar disks are evaporated by background FUV fields varies greatly from cluster to cluster.," Given the steep nature of the IMF, and the relatively modest membership sizes $N$ considered herein, the efficiency with which circumstellar disks are evaporated by background FUV fields varies greatly from cluster to cluster." + Descriptions of FUV radiation fields mist thus be represented in a statistical manner., Descriptions of FUV radiation fields must thus be represented in a statistical manner. + Here we calculate the probability distribution of the orbit-averaged flux lor five cluster profiles. both with and without dust attenuation.," Here we calculate the probability distribution of the orbit-averaged flux for five cluster profiles, both with and without dust attenuation." + These results indicate that. in the absence of dust. attenuation. between 1/:3 and 2/32 of stars in intermecdiate-sized clusters experience FUV (Iux values capable of inhibiting gian planet formation. where (his fraction increases as (he cluster size increases from ΑΝ = 100 to 1000.," These results indicate that, in the absence of dust attenuation, between 1/3 and 2/3 of stars in intermediate-sized clusters experience FUV flux values capable of inhibiting giant planet formation, where this fraction increases as the cluster size increases from $N$ = 100 to 1000." + Interestingly. although larger clusters are better able to shield stellar svstems via. dust attenuation of the central FUV radiation. roughly (he same percentage of stars enter the central void associated with the most massive star in the svstem.," Interestingly, although larger clusters are better able to shield stellar systems via dust attenuation of the central FUV radiation, roughly the same percentage of stars enter the central void associated with the most massive star in the system." + As a result. roughly 1/3 of stars will experience FUV flux values greater than our benchmark value regardless of Cluster size.," As a result, roughly 1/3 of stars will experience FUV flux values greater than our benchmark value regardless of cluster size." + To summarize. (his study shows (hat a substantial fraction (1/3 to 1/2. depending on dust. attenuation) of the solar svtems born within voung embedded. clusters are exposed to FUV radiation fields more intense than the benchmark value Gy; = 3000 C = 4.8 erg ! 7.," To summarize, this study shows that a substantial fraction (1/3 to 1/2, depending on dust attenuation) of the solar sytems born within young embedded clusters are exposed to FUV radiation fields more intense than the benchmark value $G_B$ = 3000 $G_0$ = 4.8 erg $^{-1}$ $^{-2}$." + By driving photoevaporation. radiation at this level will truncate cireumstellar disks surrounding solar (ype stars to radii r~36 AU over 10 Myr: this timescale is comparable to embedded cluster lifetimes. disk lifetimes. and the Gime required Lor giant planets to form.," By driving photoevaporation, radiation at this level will truncate circumstellar disks surrounding solar type stars to radii $r \sim 36$ AU over 10 Myr; this timescale is comparable to embedded cluster lifetimes, disk lifetimes, and the time required for giant planets to form." + The disks around smaller stars. which are more common. are more easily evaporated.," The disks around smaller stars, which are more common, are more easily evaporated." + As a result. a large fraction of solar svstems are affected by the background radiation in their birth clusters.," As a result, a large fraction of solar systems are affected by the background radiation in their birth clusters." + However. the effects are relatively modest: Although the outer parts of the disks can be truncated. regions near r~5 AU. where planets are most easily formed. are generally left. unseathed.," However, the effects are relatively modest: Although the outer parts of the disks can be truncated, regions near $r \sim 5$ AU, where planets are most easily formed, are generally left unscathed." + As a result. these radiative effects are neither negligible nor dominant.," As a result, these radiative effects are neither negligible nor dominant." + In addition. the FUV radiation fields produced within voune clusters is characterized by a wide distribution. so that assessments of radiative effects must be made statistically.," In addition, the FUV radiation fields produced within young clusters is characterized by a wide distribution, so that assessments of radiative effects must be made statistically." + An important challenge lor the future is to compare the theoretical predictions of this work with observations., An important challenge for the future is to compare the theoretical predictions of this work with observations. + Since photoevaporation acts to disperse circiumnstellar material. disk lifetimes should vary with cluster environment.," Since photoevaporation acts to disperse circumstellar material, disk lifetimes should vary with cluster environment." + As outlined above. however. one expects a wide range of radiative lluxes— and hence evaporation rates and disk lifetimes — within the same cluster.," As outlined above, however, one expects a wide range of radiative fluxes — and hence evaporation rates and disk lifetimes — within the same cluster." + One must (hus compare the full distribution of disk lifetimes [ον a given cluster, One must thus compare the full distribution of disk lifetimes for a given cluster +wwe analyzed have also been studied im detail bx Strolunaveretal.(1996.1997.1998).,"we analyzed have also been studied in detail by \cite{stroh96,stroh97b,stroh98}." +. The Al-Sky Mouitor (ASAT) onboard consists of three Scanniug Shadow Cameras (SSCs) sensitive to photons iu the energv range 210 keV. uounted on a rotating platform (Levineetal.1996).," The All-Sky Monitor (ASM) onboard consists of three Scanning Shadow Cameras (SSCs) sensitive to photons in the energy range 2–10 keV, mounted on a rotating platform \cite[]{asm96}." +. ASMI observatious are performed as sequeuces of dwells asting 90 s. after which the platform: holding the SSCs is rotated by 6 ceerees.," ASM observations are performed as sequences of dwells lasting 90 s, after which the platform holding the SSCs is rotated by 6 degrees." + Most of the sky is viewed once every ew hours., Most of the sky is viewed once every few hours. + The data from each SSC from cach dwell are averaged to obtain the daily intensities of known sources in the field of view., The data from each SSC from each dwell are averaged to obtain the daily intensities of known sources in the field of view. + The long-term 210 keV. RXTE//ASM Hux history of the source is shown in Fig. 1.., The long-term 2–10 keV /ASM flux history of the source is shown in Fig. \ref{asmdat}. + The observed photon flux exhibited variations of ~5counts on time scales of ~10 d. superimposed on a long-term trend of decreasing mica intensity.," The observed photon flux exhibited variations of $\sim5\ \cts$ on time scales of $\sim10$ d, superimposed on a long-term trend of decreasing mean intensity." + We computed the Lomb-normalised periodogram (Press&νοκ19890) of the full dataset. shown also iu Fie. l..," We computed the Lomb-normalised periodogram \cite[]{pr89} of the full dataset, shown also in Fig. \ref{asmdat}." + We found significant evidence for a periodicity with Pysy=38.6 d. with a Lomb-normalized power of 31.9 as well as secondary peaks at 63.7 and 67.5 d. The latter values are close to those found from carlicr measurements. as well as a 2-vr subset of the BRXTE//ASM data by Itongetal.(1998).," We found significant evidence for a periodicity with $\pasm=38.6$ d, with a Lomb-normalized power of 31.9 as well as secondary peaks at 63.7 and 67.5 d. The latter values are close to those found from earlier measurements, as well as a 2-yr subset of the /ASM data by \cite{kong98}." +.. The rims amplitude of the signal at [νι was =894., The rms amplitude of the signal at $\pasm$ was $\simeq 8-9$. +. The dynamical Loub-normalized periodogram shows that this periodicity varied in streneth over the cutire ASM history (Fig. 1))., The dynamical Lomb-normalized periodogram shows that this periodicity varied in strength over the entire ASM history (Fig. \ref{asmdat}) ). + The z39 d periodicity appeared strougly ouly iu the first half of the measurement interval. when the source reached its peak long-ter intensity.," The $\approx39$ d periodicity appeared strongly only in the first half of the measurement interval, when the source reached its peak long-term intensity." + After \LJD 51100. the dominant periodicity appeared to shift to around 30 cd. A broad peak arouud 65 d was also present at most times.," After MJD 51400, the dominant periodicity appeared to shift to around 30 d. A broad peak around 65 d was also present at most times." + The Proportional Counter Array (PCA:Jahodactal. consists of five identical gas-filled proportional counter units (PCUs) with a total effective area of x6000 en. aud sensitivity to X-ray pliotous in the 2GO keV range.," The Proportional Counter Array \cite[PCA;][]{xte96} consists of five identical gas-filled proportional counter units (PCUs) with a total effective area of $\approx6000\ {\rm cm}^2$ , and sensitivity to X-ray photons in the 2–60 keV range." +" We initially scauned 1 s binned lighteurves over the full PCA cnerev range (created from ""Staudiid-1 mode data) in order to locate N-rav bursts from3L.", We initially scanned 1 s binned lightcurves over the full PCA energy range (created from “Standard-1” mode data) in order to locate X-ray bursts from. + In all the PCA observations. data were also collected in user-defined modes offering higher time resolution. compared to the standard data modes.," In all the PCA observations, data were also collected in user-defined modes offering higher time resolution, compared to the standard data modes." + The PCA records the arrival time (Lys resolution) aud energy. (256 channel resolution) of every unrejected photon: data were generally binned on somewhat lower time aud spectral resolution iu order to meet telemetry limits., The PCA records the arrival time (1 $\mu$ s resolution) and energy (256 channel resolution) of every unrejected photon; data were generally binned on somewhat lower time and spectral resolution in order to meet telemetry limits. + Ouce iu N-rav burst was detected. woe extracted. 260 Κον PCA spectra within intervals of 0.25 Ls covering the burst.," Once an X-ray burst was detected, we extracted 2–60 keV PCA spectra within intervals of 0.25–1 s covering the burst." + A response matrix was eenerated separately for cach burst using version 8.0 (part of release 5.2. 2002 June 25) in order to take iuto account the known eain variations over the lite of the iustrmuent.," A response matrix was generated separately for each burst using version 8.0 (part of release 5.2, 2002 June 25) in order to take into account the known gain variations over the life of the instrument." + The eain was manually re-set by the instruct team on 3 occasions (1996 March 21. 1996 April 15 and 1999 March 22): an additional gain epoch (5) began on 2000 May. 13 with the loss of the propane laver in PCU #00," The gain was manually re-set by the instrument team on 3 occasions (1996 March 21, 1996 April 15 and 1999 March 22); an additional gain epoch (5) began on 2000 May 13 with the loss of the propane layer in PCU 0." +", In addition to these abrupt chanecs. a more eradual variation in the iustrmuental response is known to occur. due to a umber of factors."," In addition to these abrupt changes, a more gradual variation in the instrumental response is known to occur, due to a number of factors." + We then analyzed these data using au approach that is often used in X-ray. burst analysis (6.9. Ixuulkersetal.2002b.. although see also vanParadijs&Lewin1986)).," We then analyzed these data using an approach that is often used in X-ray burst analysis (e.g. \citealt{kuul01}, although see also \citealt{vpl86}) )." + For the background to the burst spectra we used a spectrin extracted. from a (typically) 16 s interval prior to the burst., For the background to the burst spectra we used a spectrum extracted from a (typically) 16 s interval prior to the burst. + Each time-resolved backeromnd-subtractec spectrum during the bursts was fit with a blackbody model anultiplied bv a low-cnerey cutoff represcuting interstellar absorption with fixed. abuudances., Each time-resolved background-subtracted spectrum during the bursts was fit with a blackbody model multiplied by a low-energy cutoff representing interstellar absorption with fixed abundances. + The initia spectral fitting was performed with the absorption colin density my free to varv: the resulting fit values typically exhibited very large scatter. particularly towards the eu of the burst when the flux was low.," The initial spectral fitting was performed with the absorption column density $n_{\rm H}$ free to vary; the resulting fit values typically exhibited very large scatter, particularly towards the end of the burst when the flux was low." + Thus. for the fina analvsis we redit cach spectra with vy fixed at the weighted mean value measured over the entire burst.," Thus, for the final analysis we re-fit each spectra with $n_{\rm H}$ fixed at the weighted mean value measured over the entire burst." +" The unabsorbed bolometric fiux £,,,4,; at each timestep f; was then estimated according to where Πω is the neutron star radius. d is the distance to the source. Ty, is the blackbody (color) temperature in mits of keV. aud yy, is the normalisation of the blackbody componcut as returned by the fitting program version 11)."," The unabsorbed bolometric flux $F_{{\rm bol},i}$ at each timestep $t_i$ was then estimated according to where $R_{\rm NS}$ is the neutron star radius, $d$ is the distance to the source, $T_{\rm bb}$ is the blackbody (color) temperature in units of keV, and $K_{\rm bb}$ is the normalisation of the blackbody component as returned by the fitting program version 11)." +" We also estimated the burst fluence £j, by iunteeratiug the measured £1,4,; over the burst duration."," We also estimated the burst fluence $E_{\rm b}$ by integrating the measured $F_{{\rm bol},i}$ over the burst duration." + We discuss the possible cousequences of the bolometric correction implicit iu equation (2)) iu section rofsvsteni. , We discuss the possible consequences of the bolometric correction implicit in equation \ref{flux}) ) in section \\ref{system}. . +Ruulkersetal.(2002b) and other authors havenoted a zz20% svstematic flux offset in, \cite{kuul01} and other authors havenoted a $\approx20$ systematic flux offset in +correspouding color excess was determined using a fixed ratio or Wy1.9.,corresponding color excess was determined using a fixed ratio or $R_R = 1.3$. +" We presettr esults for case (il) wih (A,B)= 020. aud oypp=09025.0035. aud 0705 in the Upper Riel and Lower yanels of Fieure 7.."," We present results for case (ii) with $\langle\Delta_a R\rangle = 0\fm025$ , and $\sigma_{a,R}=0\fm025, 0\fm035,$ and $0\fm05$ in the Upper Right and Lower panels of Figure \ref{dist}." + As before. the corresj»ouding color shift for each galaxy was obtained assuimiug constant Ay ratio.," As before, the corresponding color shift for each galaxy was obtained assuming constant $R_R$ ratio." + From these cases. it is clear that the wider the distribution of shitS. the more difficult it is to recover the correct. valtes of the extinction all| vedclenine shifts.," From these cases, it is clear that the wider the distribution of shifts, the more difficult it is to recover the correct values of the extinction and reddening shifts." + This is due to he [act that ie signal eventually becomes lost in the noise when unplementiug wider and wider distributious., This is due to the fact that the signal eventually becomes lost in the noise when implementing wider and wider distributions. + Heuce. it is necessary to ine a deeper catalog withoa larger coverage a‘ea du order o determine the distribution of extinctious and redcdeniugs inside a set of clusters wrere dust is distributed non-uniformly.," Hence, it is necessary to use a deeper catalog with a larger coverage area in order to determine the distribution of extinctions and reddenings inside a set of clusters where dust is distributed non-uniformly." + To impleuent the third type of dust. which coresponds to dust. distributed patchily. Le. iu ‘pockets’ we raidoinly. picked a specified fraction of galaxies from the Cluster Group aid assigned them exactly he same “cust” induced shifts in maguitude and color.," To implement the third type of dust, which corresponds to dust distributed patchily, i.e. in `pockets' we randomly picked a specified fraction of galaxies from the Cluster Group and assigned them exactly the same `dust' induced shifts in magnitude and color." + Other Cluster Group galaxies were assiimed to lie belii cluster areas that contaited no dust., Other Cluster Group galaxies were assumed to lie behind cluster areas that contained no dust. + The fraction of galaxies that were allected by dust is the clus covering factor of the chSer., The fraction of galaxies that were affected by dust is the dust covering factor of the cluster. + Figure δ slows examples of type (iil) «S., Figure \ref{bimodal} shows examples of type (iii) dust. + The upper right and lower eft show results for bimocal dust: a fraction of the sky in Cluster Gre)ηJs (3:Y and 20%. respectivev) is asstumed to be devoid of dust aud ‘esults in a local ΙΙΙ 11 i £j. plane at (0.0).," The upper right and lower left show results for bimodal dust: a fraction of the sky in Cluster Groups $33\%$ and $20\%$, respectively) is assumed to be devoid of dust and results in a local minimum in the $\xi_{k,l}$ plane at $(0,0)$." +" The res of the sky area is obscurec. and result slut je second local iiuimu1i t non-zero values of 4? anc LA,(By-—2."," The rest of the sky area is obscured, and results in the second local minimum at non-zero values of $\Delta_a R$ and $\Delta_a (B_J-R)$." + The lower isht panel s|ows the expected effect result from a cluster that coutais üunobscurec lines of sit as well as 1wo cdiΠοιοί populatious of d both having Ry=0.5 (ins4.tead of 1.3 used in all ot: tests). but [1ilerent degrees of concentra10," The lower right panel shows the expected effect resulting from a cluster that contains unobscured lines of sight as well as two different populations of dust, both having $R_{R} = 0.5$ (instead of $1.3$ used in all other tests), but different degrees of concentration." +" ' 8 also tests te ability of οιr techuiqi o detect. dust istributed in “pockets""ol small coveri actors.", Figure \ref{bimodal} also tests the ability of our technique to detect dust distributed in `pockets' of small covering factors. + It appea rstat dust with coverii ‘actors less tli: layout 20% would be cliffτιν to detect. at least wit ide pesent sample of cluser.," It appears that dust with covering factors less than about $20\%$ would be difficult to detect, at least with the present sample of clusters." +" Figwes T and & demoistrate that our tech[unique cau detect cluster dust. aud clistiuguisli valloIs vpes of its sλαial «istribution: smooth. 1on-uniforn. cy elumpy. as long as t1e typical values of σι,μ are not toο large. Tz9025—0205. €ald the coveriig factor is not too sinall. 220%."," Figures \ref{dist} and \ref{bimodal} demonstrate that our technique can detect cluster dust, and distinguish various types of its spatial distribution: smooth, non-uniform, or clumpy, as long as the typical values of $\sigma_{a,R}$ are not too large, $\,\lesssim 0\fm025-0\fm05$, and the covering factor is not too small, $\simgt\, 20\%$." + It is important to note that lese limitations as wel as the CN-pane resolutiou are not liuitatious of the imethod. but are ]uste:1d due to the limited uumber galaxies. aud the relatively bieht flux limit of the APM galaxies.," It is important to note that these limitations as well as the CM-plane resolution are not limitations of the method, but are instead due to the limited number galaxies, and the relatively bright flux limit of the APM galaxies." +" Or technique works we Las long as tle dust induced shifts iu color aud iagnitude are z5196 of he total CM-plaue exteut of the nj4í4 atid Neen, distributions.", Our technique works well as long as the dust induced shifts in color and magnitude are $\lesssim 1\%$ of the total CM-plane extent of the $n_{clus}$ and $n_{cont}$ distributions. + Tle extent ol these in magnitude are Liited on the faint. side by the {lus laiit of tlie galaxy survey. while on the bright side the distributions are effectively limited by the limited sky coverage.," The extent of these in magnitude are limited on the faint side by the flux limit of the galaxy survey, while on the bright side the distributions are effectively limited by the limited sky coverage." + Iu this paper we lave developed a technique to detect reddening and obscuration due to dust in galaxy clusters., In this paper we have developed a technique to detect reddening and obscuration due to widely-distributed dust in galaxy clusters. + Welave applied the method to a sample of low redshilt. z<0.08 high Galactic Latitude APM clusters.," Wehave applied the method to a sample of low redshift, $z\leq 0.08$ high Galactic Latitude APM clusters." + Our analysis indicates that the reddeniug and extinction due, Our analysis indicates that the reddening and extinction due +In each figure. it can be seen that the difference in N[CO(¢)] between the À and B parameter sets decreases as the extinction or H» column increases.,"In each figure, it can be seen that the difference in $N$ [CO(g)] between the A and B parameter sets decreases as the extinction or $_{2}$ column increases." + At low H» column. the amount of CO(g) produced in the B models exceeds that in the A models by up to 2 orders of magnitude.," At low $_{2}$ column, the amount of CO(g) produced in the B models exceeds that in the A models by up to 2 orders of magnitude." +" This effect occurs primarily because the warmer grains produce less H»(g) at a given Ay with our simple ""smooth"" grain model: there is much less of a temperature dependence in the ""rough"" grain model of Changetal.(2007).. which utilizes à microscopic Monte Carlo approach."," This effect occurs primarily because the warmer grains produce less $_2$ (g) at a given $A_V$ with our simple “smooth” grain model; there is much less of a temperature dependence in the “rough” grain model of \citet{cch07}, which utilizes a microscopic Monte Carlo approach." + If the abscissa in Figure 3. were the column of total hydrogen rather than of just H». the B and A model results for CO(g) would appear to be more similar.," If the abscissa in Figure \ref{fig-3CO} were the column of total hydrogen rather than of just $_{2}$, the B and A model results for CO(g) would appear to be more similar." + Note also that the two plots with N[CO(g)| emphasize its growth with time in different ranges., Note also that the two plots with $N$ [CO(g)] emphasize its growth with time in different ranges. + In particular. the range of labeled Ay values in Figure 2. of 0.5-3.0 corresponds to only a small portion of the H» column density range in Figure 3..," In particular, the range of labeled $A_{\rm V}$ values in Figure \ref{fig-2NH} of 0.5-3.0 corresponds to only a small portion of the $_{2}$ column density range in Figure \ref{fig-3CO}." + Despite these complications. a salient feature of both figures is that CO(g) becomes the most abundant of the three major forms of carbon as N[Hs].~4107 em (corresponding to an edge-to-center Ay~ 2).," Despite these complications, a salient feature of both figures is that CO(g) becomes the most abundant of the three major forms of carbon as $N[\rm{H_{2}}]\sim 4\times 10^{21}$ $^{-2}$ (corresponding to an edge-to-center $A_V \sim 2$ )." + Thus. long before a cold core with Ay~10 forms. the carbon inventory in the gas phase contains large amounts of CO(g).," Thus, long before a cold core with $A_{\rm V} \sim 10$ forms, the carbon inventory in the gas phase contains large amounts of CO(g)." + Our results. obtained with a gas-grain model. are not identical with those of BO4.," Our results, obtained with a gas-grain model, are not identical with those of B04." + Nevertheless. they confirm how artificial the initial carbon inventory is for pseudo-time-dependent models. in which carbon is assumed to be totally in its atomic tonic form while the density and visual extinction are already at their final values.," Nevertheless, they confirm how artificial the initial carbon inventory is for pseudo-time-dependent models, in which carbon is assumed to be totally in its atomic ionic form while the density and visual extinction are already at their final values." + Moreover. at a visual extinction of 2. the [C + C[/[CO] abundance ratio ts the same as we obtain at steady-state with a gas-phase model at roughly the same density.," Moreover, at a visual extinction of 2, the $^+$ $+$ C]/[CO] abundance ratio is the same as we obtain at steady-state with a gas-phase model at roughly the same density." + Thus. the system has already evolved into a translucent cloud.," Thus, the system has already evolved into a translucent cloud." + In addition to CO(g). calculated column densities for the major ice species H»O(s). CO(s). and CO:(s). along with the somewhat less abundant ices 4(s) and CHiOH(s). are shown in Figure 2..," In addition to CO(g), calculated column densities for the major ice species $_2$ O(s), CO(s), and $_2$ (s), along with the somewhat less abundant ices $_{4}$ (s) and $_3$ OH(s), are shown in Figure \ref{fig-2NH}." + As illustrated in panel (b). the amount of H»O(s) grows relatively quickly to column densities in excess of 1016 cm? between Ay=1.3 and 2.1 for the various models.," As illustrated in panel (b), the amount of $_2$ O(s) grows relatively quickly to column densities in excess of $^{16}$ $^{-2}$ between $A_V = 1.3$ and 2.1 for the various models." + As with CO(g). the results illustrate that this occurs most efficiently for shock models with greater total density.," As with CO(g), the results illustrate that this occurs most efficiently for shock models with greater total density." + The evolution of CO(s). as shown in panel (c). tends to lag behind CO(g). and the plot of its column vs Ay shows a sharp threshold at intermediate visual extinction. which is correlated with a decrease in the rate of synthesis of methanol (see below).," The evolution of CO(s), as shown in panel (c), tends to lag behind CO(g), and the plot of its column vs $A_{\rm V}$ shows a sharp threshold at intermediate visual extinction, which is correlated with a decrease in the rate of synthesis of methanol (see below)." + The molecule forms more efficiently at lower extinction and earlier times for models with greater density. similar to the trend seen for H»O(s).," The molecule forms more efficiently at lower extinction and earlier times for models with greater density, similar to the trend seen for $_2$ O(s)." + There is a large distinction between models with the A and B parameters at low and intermediate. visual extinction., There is a large distinction between models with the A and B parameters at low and intermediate visual extinction. + For some molecular ices. specifically CO(s) and H»O(s). às the visual extinction. grows. the calculated column densities converge towards one another. while others show order of magnitude agreement (CH. CH:OH).," For some molecular ices, specifically CO(s) and $_2$ O(s), as the visual extinction grows, the calculated column densities converge towards one another, while others show order of magnitude agreement $_4$, $_3$ OH)." + No convergence occurs for COs(s). as shown in panel (d). which requires the adoption of the B parameter set to achieve a significant column density. mainly because the warmer grams allow faster. diffusive reaction processes for heavy species.," No convergence occurs for $_2$ (s), as shown in panel (d), which requires the adoption of the B parameter set to achieve a significant column density, mainly because the warmer grains allow faster diffusive reaction processes for heavy species." + Although the effects of individually varying the CO2(s) formation barrier and turning off the Eley-Rideal mechanism on the abundance of CO2(s) were considered. the grain temperature is the most influential parameter for COs(s) formation.," Although the effects of individually varying the $_2$ (s) formation barrier and turning off the Eley-Rideal mechanism on the abundance of $_2$ (s) were considered, the grain temperature is the most influential parameter for $_2$ (s) formation." + The formation of CO(g) and CO(s) also affects. the evolution of methane ice (CH4(s)) and methanol ice (CH;OH(s). shown in panels (e) and (f) of Figure 2.. respectively.," The formation of CO(g) and CO(s) also affects the evolution of methane ice $_4$ (s)) and methanol ice $_3$ OH(s)), shown in panels (e) and (f) of Figure \ref{fig-2NH}, respectively." + Methane on gram surfaces is formed via sequential hydrogenation of C atoms. which accrete from the gas.," Methane on grain surfaces is formed via sequential hydrogenation of C atoms, which accrete from the gas." + The less carbon and the more CO in the gas. the less efficient the surface formation of methane and the more CO(s) aceretes onto grams.," The less carbon and the more CO in the gas, the less efficient the surface formation of methane and the more CO(s) accretes onto grains." + Even at an extinction of 3.0. the CH4(s) column densities for individual models tend to lie below those of CO(s). especially for B parameter (warm grain) models.," Even at an extinction of 3.0, the $_{4}$ (s) column densities for individual models tend to lie below those of CO(s), especially for B parameter (warm grain) models." + As regards methanol. its only formation is via. sequential hydrogenation of CO(s) by reactions involving H atoms on grain surfaces: so that the abundances of CO(s) and CH;OH(s) should be intimately tied together.," As regards methanol, its only formation is via sequential hydrogenation of CO(s) by reactions involving H atoms on grain surfaces: so that the abundances of CO(s) and $_3$ OH(s) should be intimately tied together." + Nevertheless. the extinction dependence of the column densities of methane and methanol is far more complex than that of CO(s). especially with the B parameters. where extra peaks are seen.," Nevertheless, the extinction dependence of the column densities of methane and methanol is far more complex than that of CO(s), especially with the B parameters, where extra peaks are seen." + For both CH4(s) and CH:OH(s) formation. warmer temperatures lead to more desorption of H atoms from the grains and would appear to slow down the hydrogenation of C(s) and CO(s).," For both $_4$ (s) and $_3$ OH(s) formation, warmer temperatures lead to more desorption of H atoms from the grains and would appear to slow down the hydrogenation of C(s) and CO(s)." + This factor clearly affects the formation of methane at most times. but runs counter to the observation that higher surface temperatures aid methanol formation and CO(s) depletion.," This factor clearly affects the formation of methane at most times, but runs counter to the observation that higher surface temperatures aid methanol formation and CO(s) depletion." + Presumably this latter effect occurs because two of the reactions in the hydrogenation of methanol (H + CO and H + H:CO) require overcoming a chemical reaction barrier., Presumably this latter effect occurs because two of the reactions in the hydrogenation of methanol (H + CO and H + $_2$ CO) require overcoming a chemical reaction barrier. + In addition to the results shown in Fig., In addition to the results shown in Fig. + 2. and Fig. 3..," \ref{fig-2NH} and Fig. \ref{fig-3CO}," + we report additional results for major species at Ay = 3., we report additional results for major species at $A_{\rm V}$ = 3. + Specifically. Table 2 contains the fractional abundances of the major gaseous and solid species for the A and B variants of Models 3 and 4.," Specifically, Table \ref{tbl-majform} contains the fractional abundances of the major gaseous and solid species for the A and B variants of Models 3 and 4." + These models span the range of physical conditions considered. and the table illustrates the differences in composition that arise as a result.," These models span the range of physical conditions considered, and the table illustrates the differences in composition that arise as a result." + This table provides a range of input abundances for subsequent dense core models., This table provides a range of input abundances for subsequent dense core models. + Our results are obtained with a model that ignores grain growth., Our results are obtained with a model that ignores grain growth. + But. as grains accumulate icy mantles. the cross sections grow likewise. potentially by substantial factors.," But, as grains accumulate icy mantles, the cross sections grow likewise, potentially by substantial factors." + This effect is often omitted from gas-grain chemistry models., This effect is often omitted from gas-grain chemistry models. + In the single grain-size model considered in this paper. the cross section. gradually doubles following Ay21.5.," In the single grain-size model considered in this paper, the cross section gradually doubles following $A_V \gtrsim 1.5$." + In some preliminary simulations with mantle growth. we found that the onset of CO(s) and CO2(s) growth occurs at slightly lower Av. and that these ices reach abundances slightly different from those of Fig. 2..," In some preliminary simulations with mantle growth, we found that the onset of CO(s) and $_2$ (s) growth occurs at slightly lower $A_V$, and that these ices reach abundances slightly different from those of Fig. \ref{fig-2NH}," + while CO(g) and water ice decrease., while CO(g) and water ice decrease. + Our overall results and conclusions are not changed significantly by consideration of this effect., Our overall results and conclusions are not changed significantly by consideration of this effect. + The inclusion of mantle growth in ἃ chemical model with a distribution of grain sizes is currently being undertaken by K. Acharyya and E. Herbst., The inclusion of mantle growth in a chemical model with a distribution of grain sizes is currently being undertaken by K. Acharyya and E. Herbst. + In addition to mantle growth. the effect of grain-grain collisions ean lead to size distributions with. considerably larger grains at time scales such as those considered here. according to detailed," In addition to mantle growth, the effect of grain-grain collisions can lead to size distributions with considerably larger grains at time scales such as those considered here, according to detailed" +from the nucleus of the host galaxy: where the jet flow collides with the intergalactic mecditm (GAD. inflating the radio lobes.,"from the nucleus of the host galaxy, where the jet flow collides with the intergalactic medium (IGM), inflating the radio lobes." + AGN jets are chormously powerful. with a total bolometric power output that is often comparable to or greater than that from the host galaxy. aud a kinetic enerev flux that can be comparable to the Αν5 bolometric huninosity (Rawlnes Saunders 1991).," AGN jets are enormously powerful, with a total bolometric power output that is often comparable to or greater than that from the host galaxy, and a kinetic energy flux that can be comparable to the AGN's bolometric luminosity (Rawlings Saunders 1991)." + Uutil the last decade. almost all progress towarcls πιακαπο the plysics of jets had come either from nunierical nkxleliug or multi-frequency radio ma)piueg.," Until the last decade, almost all progress towards understanding the physics of jets had come either from numerical modeling or multi-frequency radio mapping." + ILowever. the pace of discovery has accelerated exeatlv in the past decade withHST and observatious.," However, the pace of discovery has accelerated greatly in the past decade with and observations." + One of the first observatious byChendra is a perfect iUlustration: he target. PINS 0637752. was a ight radio-loud quasa.," One of the first observations by is a perfect illustration: the target, PKS 0637–752, was a bright radio-loud quasar." + It was believed that this source would be uuresolvec in the N-ravs. aud therefore ideally suited to focus the telescope.," It was believed that this source would be unresolved in the X-rays, and therefore ideally suited to focus the telescope." + Dustead. we saw a beautiul X-rav jet well over 10 arcsecouds long (Chartas et al.," Instead, we saw a beautiful X-ray jet well over 10 arcseconds long (Chartas et al." + 2000. Schlsvar zct al.," 2000, Schwartz et al." +" 2100), with iiorphologv similar to that seen iu 16 radio."," 2000), with morphology similar to that seen in the radio." + Deeper multi-baud observatious of this jet iive since been done to constrain the nature of its ¢nulisson aid also its matterfenergy couteut (Cr'oorganopotios et al., Deeper multi-band observations of this jet have since been done to constrain the nature of its emission and also its matter/energy content (Georganopoulos et al. + 2005. Uchivama et al.," 2005, Uchiyama et al." + 2005. Mehta et al.," 2005, Mehta et al." + 20093., 2009). + Tudeed. since the launch of the two Creat Observ.atories. the iuuber of extended. arcsecond-scale jets known to enüt in the optical and X-ravs has increased about teold. from less than 5 to more than 50. inchding members of everv huninosity aud morphologicaL class of raclio-loud ACN.," Indeed, since the launch of the two Great Observatories, the number of extended, arcsecond-scale jets known to emit in the optical and X-rays has increased about ten-fold, from less than 5 to more than 50, including members of every luminosity and morphological class of radio-loud AGN." + For jets. uniltiwavebaxl observations eive plivsical," For jets, multiwaveband observations give physical" +oredicted: ancl observed. merger rates.,predicted and observed merger rates. + The merging time-scale is constant in our model. but we cannot exclude an inconsisteney with the time-scale assumed by the CAS and xür methods.," The merging time-scale is constant in our model, but we cannot exclude an inconsistency with the time-scale assumed by the CAS and pair methods." + However. this reason alone is unlikely to explain the ciscrepaney between the merger fractions and rates in the lowest mass bins. given the agreement at the ighest masses.," However, this reason alone is unlikely to explain the discrepancy between the merger fractions and rates in the lowest mass bins, given the agreement at the highest masses." + The one quantity we have not examined vet is. the xedieted number density of galaxies within the Millennium simulation., The one quantity we have not examined yet is the predicted number density of galaxies within the Millennium simulation. + Since we are selecting galaxies through a stellar mass cul. the predicted. number density. of these systems allects the final. values of the merger rates.," Since we are selecting galaxies through a stellar mass cut, the predicted number density of these systems affects the final values of the merger rates." + We examine this in Fig. S.," We examine this in Fig. \ref{massf}," +" which shows a comparison between the Millennium predictions and observations for the number density of galaxies with AZ,21017 ((Lillecl triangles) ancl AZ,71013 ((empty squares).Data points are taken from Conscliceοἱal. (2007).", which shows a comparison between the Millennium predictions and observations for the number density of galaxies with $M_{\star}>10^{10}$ (filled triangles) and $M_{\star}>10^{11}$ (empty squares).Data points are taken from \citet{cons2007}. +. The solid line shows results for the model of Bertoneetal.(2007).. the dashed line for that of DeLucia&Dlaizot (2007).," The solid line shows results for the model of \citet{bertone2007}, the dashed line for that of \citet{delucia2007}." +". The model of Bertoneetal.(2007). roughly. predicts he number density of galaxies with AZ,>1043AL... while he moclel of DeLucia&Blaizot(2007) unclerprecicts it at all redshifts."," The model of \citet{bertone2007} roughly predicts the number density of galaxies with $M_{\star}>10^{11}$, while the model of \citet{delucia2007} underpredicts it at all redshifts." +" Both models overestimate the observed values or galaxies with AJ,>101 bv almost an order of magnitude.", Both models overestimate the observed values for galaxies with $M_{\star}>10^{10}$ by almost an order of magnitude. +"This results in the hieh value derived for the merger rate of galaxies with AZ,«1014 since. as we know from Lig. 5..","This results in the high value derived for the merger rate of galaxies with $M_{\star}<10^{11}$ , since, as we know from Fig. \ref{figcomp}, ," + the predicted merger, the predicted merger +acecretion ancl virialization.,accretion and virialization. + In this paper we investigate in how far we can get insight into this history on the basis of the internal properties of the clusters., In this paper we investigate in how far we can get insight into this history on the basis of the internal properties of the clusters. + This. involves characteristics like their mass ancl mass distribution. their size and their kinetic and gravitational potential energy.," This involves characteristics like their mass and mass distribution, their size and their kinetic and gravitational potential energy." + In particular. we are keen to learn whether these do show any possible trace of a cosmological constant.," In particular, we are keen to learn whether these do show any possible trace of a cosmological constant." + One particular profound manifestation of the virial state of cosmic objects is via scaling relations that connect various structural properties., One particular profound manifestation of the virial state of cosmic objects is via scaling relations that connect various structural properties. + Scaling relations of collapsed and virializecl objects relate two or. three. fundamental characteristics., Scaling relations of collapsed and virialized objects relate two or three fundamental characteristics. + The first involves a quantity measuring the amount of mass AL. often expressed in terms of the amount of light. { emitted by the object.," The first involves a quantity measuring the amount of mass $M$, often expressed in terms of the amount of light $L$ emitted by the object." + The second. quantity involves the size of the object. while the third one quantifies its dynamical state.," The second quantity involves the size of the object, while the third one quantifies its dynamical state." +" For a virialized halo with mass AM. size R and velocity dispersion o,=(0757. the implied scaling relation is where cy is a constant that rellects. the interna dynamics of the svstem."," For a virialized halo with mass $M$, size $R$ and velocity dispersion $\sigma_{v}=\langle v^{2}\rangle^{1/2}$, the implied scaling relation is where $\epsilon_{M}$ is a constant that reflects the internal dynamics of the system." +" (ea, is determined by issues such as he isotropy if the cluster velocity dispersion. its shape anc any substructure)."," $\epsilon_{M}$ is determined by issues such as the isotropy if the cluster velocity dispersion, its shape and any substructure)." +" Systems having similar values of this constant woul of expected. to form a two-parameter family of objects: observationally this manifests itself as the ""Fundamenta lane”.", Systems having similar values of this constant would be expected to form a two-parameter family of objects: observationally this manifests itself as the “Fundamental Plane”. + Objects Iving on the same plane might be expectec o have similar formation histories and. conversely. the nature of the Fundamental Plane is a clue to the underlying ormation mechanism.," Objects lying on the same plane might be expected to have similar formation histories and, conversely, the nature of the Fundamental Plane is a clue to the underlying formation mechanism." + The scaling relations are of great importance for a variety of reasons., The scaling relations are of great importance for a variety of reasons. + First of all. they inform us about the dynamical state of the objects ancl must be a profound reflection of the galaxy formation process (Robertsonοἱal. 2006).," First of all, they inform us about the dynamical state of the objects and must be a profound reflection of the galaxy formation process \citep{robertson06}." +. Also. they have turned out to be of substantial practical importance.," Also, they have turned out to be of substantial practical importance." +" Because they relate. an intrinsic distance independent quantity like velocity. dispersion. to a distance dependent one Like L,. they can be used. as cosmological distance indicators."," Because they relate an intrinsic distance independent quantity like velocity dispersion to a distance dependent one like $L_{e}$, they can be used as cosmological distance indicators." + Since the mid. TOs. we know that the observed properties of elliptical galaxies follow scaling relations.," Since the mid 70s, we know that the observed properties of elliptical galaxies follow scaling relations." + The Faber-Jackson relation (Faber&Jackson1976). relates the luminosity £ and the velocity dispersion σ of an elliptical galaxy., The Faber-Jackson relation \citep{fj76} relates the luminosity $L$ and the velocity dispersion $\sigma$ of an elliptical galaxy. + The ‘Tully-Fisher relation (Tully&Fisher1977) is the equivalent for spiral galaxies., The Tully-Fisher relation \citep{tully77} is the equivalent for spiral galaxies. + A different. though related. scaling is that between the effective radius p and the luminosity £ of the galaxy.," A different, though related, scaling is that between the effective radius $r_{e}$ and the luminosity $L$ of the galaxy." + This is known as the WKormency relation (Ixormendsy1977)., This is known as the Kormendy relation \citep{kormendy77}. +.. These two relations turned out to be manifestations of a deeper scaling relation between three fundamental characteristics. which became known as the Fundamental Plane (Djorgovski&Davis1987:Dresslerοἱal. 1987).," These two relations turned out to be manifestations of a deeper scaling relation between three fundamental characteristics, which became known as the Fundamental Plane \citep{djorgovski87, dressler87}." +. The Fundamental Plane is generally expressecl as a relationship between three parameters. though there. is no consensus as to which three should best. be used. nor precisely how to celine them.," The Fundamental Plane is generally expressed as a relationship between three parameters, though there is no consensus as to which three should best be used, nor precisely how to define them." + This makes detailed comparisons somewhat dillicult., This makes detailed comparisons somewhat difficult. + Some authors use the set (logH.loge.log£). £ being the luminosity in some spectral band within some radius 2. while others use the set (logHR.loga.p). p being the mean surface brightness within that radius.," Some authors use the set $\log R, \log \sigma, \log I$ ), $I$ being the luminosity in some spectral band within some radius $R$, while others use the set $\log R, \log \sigma, \mu$ ), $\mu$ being the mean surface brightness within that radius." + Comparisons are further complicated: by the [act that there appear to be manifest residual. luminosity dependences in the fits. as reported in a recent study of the SDSS by Nigoche-Netroetal.," Comparisons are further complicated by the fact that there appear to be manifest residual luminosity dependences in the fits, as reported in a recent study of the SDSS by \cite{Nigoche-Netro}." +(2009). Care is needed. when interpreting these observed relationships., Care is needed when interpreting these observed relationships. + Observed data generally refers to luminosity. rather than mass. and the radius that is used generally refers to some Liclucial radius such as the halt-light radius or some radius based on profile fitting.," Observed data generally refers to luminosity rather than mass, and the radius that is used generally refers to some fiducial radius such as the half-light radius or some radius based on profile fitting." + Often. the hall-lieht radius. I. as determined. [rom a fit to à de Vaucouleurs profile is usec.," Often, the half-light radius, $R_e$ as determined from a fit to a de Vaucouleurs profile is used." + ‘This situation has been improved somewhat by the >eravitational lensingὃν study of Boltonetal.(2007)., This situation has been improved somewhat by the gravitational lensing study of \cite{bolton07}. +.. These authors presented a new formulation of the FPH using lensing data to replace surface brightness with surface mass density., These authors presented a new formulation of the FP using lensing data to replace surface brightness with surface mass density. +" They also present an interesting alternative. which they refer to as the ""Mass. Plane” (AIP). in which they find the cdepencence of ος} on logfa.2) and surface. mass density X,» within a radius £2,/2."," They also present an interesting alternative, which they refer to as the “Mass Plane” (MP), in which they find the dependence of $\log(R_e)$ on $\log(\sigma_{e2})$ and surface mass density $\Sigma_{e2}$ within a radius $R_e/2$." + Using surface mass density Mee within a radius A42 in place of surface. brightness Il. removes one of the assumptions about the relationship between mass ane light., Using surface mass density $\Sigma_{e2}$ within a radius $R_e/2$ in place of surface brightness $I_e$ removes one of the assumptions about the relationship between mass and light. + Aluch recent galaxy cluster work on the Fundamental Plane has focussed. on the dillerences between the Fundamental Planes of the clusters as defined by their member galaxies (see for example D'Onofrioetal.(2008) and. references therein)., Much recent galaxy cluster work on the Fundamental Plane has focussed on the differences between the Fundamental Planes of the clusters as defined by their member galaxies (see for example \cite{onofrio08} and references therein). + Galaxy cluster scaling relations were discovered by Sehaclleretal.(1993) who studied. a sample of 16 galaxy clusters. concluding that these svstems also populate a Funcdamental Plane.," Galaxy cluster scaling relations were discovered by \cite{schaeffer93} who studied a sample of 16 galaxy clusters, concluding that these systems also populate a Fundamental Plane." + Adamictal.(1998S) used the ESO Nearby Abell Cluster Survey (ENACS) to study the existence of a Fundamental Plane for rich galaxy. clusters. finding that it is significantly different from that for elliptical ealaxies.," \cite{adami98} used the ESO Nearby Abell Cluster Survey (ENACS) to study the existence of a Fundamental Plane for rich galaxy clusters, finding that it is significantly different from that for elliptical galaxies." + \larmoetal.(2004) using data from the WINGS cluster survey found that the dilference is largely a simple shift in the relative positions of the planes., \cite{marmo04} using data from the WINGS cluster survey found that the difference is largely a simple shift in the relative positions of the planes. + The larecly unknown relationship between mass and light frustrates a direct. comparison with the results of N-Jody investigations., The largely unknown relationship between mass and light frustrates a direct comparison with the results of N-Body investigations. + Later. Lanzonictal.(2004) aclelressecl the question using N-Bocly simulations for high mass halos. which are thought to host clusters of galaxies.," Later, \cite{lanzoni04} addressed the question using N-Body simulations for high mass halos, which are thought to host clusters of galaxies." + On the basis of 13 simulated massive dark matter halos ina ACDAL cosmology they found that the dark matter halos follow the EJ. IXormendy. and IEP-like relations.," On the basis of 13 simulated massive dark matter halos in a $\Lambda$ CDM cosmology they found that the dark matter halos follow the FJ, Kormendy and FP-like relations." + In hierarchical scenarios of structure formation. halos build up by subsequent merging of smaller halos into larger and larger halos., In hierarchical scenarios of structure formation halos build up by subsequent merging of smaller halos into larger and larger halos. + Some of these mergers involves sizeable clumps. most involves a more quiescent accretion of malter and small clumps from the surroundings.," Some of these mergers involves sizeable clumps, most involves a more quiescent accretion of matter and small clumps from the surroundings." + This process leaves, This process leaves +deviation over time. we can derive the evolution of the total distribution using a scaling law just for the mean value according to where Ly is the mean log hnuuinositv of the Pleiades (1077720.35 ere/s).,"deviation over time, we can derive the evolution of the total distribution using a scaling law just for the mean value according to where $L_{0}$ is the mean log luminosity of the Pleiades $10^{29.35}$ erg/s)." + The derived sealing law is a first approximation limited by the chisters available!5., The derived scaling law is a first approximation limited by the clusters available. +. Di Fig., In Fig. + 2. we compare the scaling law from Eq., \ref{lumi} we compare the scaling law from Eq. + 2 with those derived by Ribas et al. (, \ref{scaling} with those derived by Ribas et al. ( +2005).,2005). + The solid lue represents the mean value of the log.Oonormal distribution. while the shaded area is bounded by curves giving the lo equivalent spread.," The solid line represents the mean value of the log–normal distribution, while the shaded area is bounded by curves giving the $\sigma$ equivalent spread." + The dashed line gives the 1100 A interval from Ribas et al., The dashed line gives the 1–100 $\mbox{\AA}$ interval from Ribas et al. + To verity the validity of the derived scaling law we construct the distribution fiction for the nearby feld stars by assume a constant stellar birth rate aud a maxima age of 12 Gar., To verify the validity of the derived scaling law we construct the distribution function for the nearby field stars by assuming a constant stellar birth rate and a maximum age of 12 Gyr. + Doing so. we can very well reconstruct the cumulative distribution function for CC stars given in Sclinitt (1997) (Fig. 1)).," Doing so, we can very well reconstruct the cumulative distribution function for G stars given in Schmitt (1997) (Fig. \ref{comp}) )." + To calculate the atinospherie loss from cxoplaucts we use a modified energyvluted approach. which was discussed im details bv Erkaev et al. (," To calculate the atmospheric loss from exoplanets we use a modified energy–limited approach, which was discussed in details by Erkaev et al. (" +2007). s0 we shall be brief here.,"2007), so we shall be brief here." +Galaxy clusters are ideal laboratories in which to study dark matter. being the most massive bound structures in the universe and dominated by their dark matter component (~ 90%).,"Galaxy clusters are ideal laboratories in which to study dark matter, being the most massive bound structures in the universe and dominated by their dark matter component $\sim90\%$ )." + Constraining the clustering properties of dark matter is crucial for refining structure formation models that predict both the shapes of dark matter halos and their mass function(e.g. Navarroetal.(1997): Bahealletal.2003):: Dahle (20062)., Constraining the clustering properties of dark matter is crucial for refining structure formation models that predict both the shapes of dark matter halos and their mass function(e.g. \cite{navarro};; \cite{bahcall}; \cite{dahle}) ). + Several methods are used to measure galaxy cluster dark matter profile shapes and halo masses on a range of scales. including X-ray and Sunyaev-Zeldovich (SZ) studies. dynamical analyses. and gravitational lensing.," Several methods are used to measure galaxy cluster dark matter profile shapes and halo masses on a range of scales, including X-ray and Sunyaev-Zeldovich (SZ) studies, dynamical analyses, and gravitational lensing." + However. all of these methods require simplifying assumptions to be made regarding the shape and/or dynamical state of the cluster in order to derive meaningful constraints from available data.," However, all of these methods require simplifying assumptions to be made regarding the shape and/or dynamical state of the cluster in order to derive meaningful constraints from available data." + Crucially. while most parametric methods typically assume spherical symmetry of the halo. observed galaxy clusters often exhibit significant projected ellipticity and halos in CDM structure formation simulations (e.g. Bettetal.(2007). (using he Millennium simulation): Shawetal. (2006))) show significant riaxiality in cluster-scale halos. with axis ratios between minor and major axes as small as O.4.," Crucially, while most parametric methods typically assume spherical symmetry of the halo, observed galaxy clusters often exhibit significant projected ellipticity and halos in CDM structure formation simulations (e.g. \cite{bett} (using the Millennium simulation); \cite{shaw}) ) show significant triaxiality in cluster-scale halos, with axis ratios between minor and major axes as small as 0.4." +. Understanding and accurately incorporating the impact of this physical reality on cluster mass and xrameter estimates is crucial for accurate comparisons between measured cluster properties and model predictions., Understanding and accurately incorporating the impact of this physical reality on cluster mass and parameter estimates is crucial for accurate comparisons between measured cluster properties and model predictions. + In addition to the determination. of masses. most cluster oofile fits are carried out in the hope of either supporting or refuting the universality of the NFW profile and thus testing the CDM paradigm.," In addition to the determination of masses, most cluster profile fits are carried out in the hope of either supporting or refuting the universality of the NFW profile and thus testing the CDM paradigm." +" The NFW profile is typically parameterised by an approximate virial mass Adsoy and a concentration parameter. 6, and simulations prediet a strong correlation between the two."," The NFW profile is typically parameterised by an approximate virial mass $M_{200}$ and a concentration parameter, $C$, and simulations predict a strong correlation between the two." + For a cluster of AZ=1047 M..C€'— 4.However. several authors (e.g.," For a cluster of $M=10^{15}$ $_{\odot}$ , $C \sim 4$ .However, several authors (e.g." +"It is well known that blazars are variable sources in gamma-ray (seee.g.Abdoetal.2009a,2010d)..","It is well known that blazars are variable sources in gamma-ray \citep[see e.g.][]{abd09_3c454.3,abd10_var}." +" If gamma-ray loud radio galaxies are the misaligned population of blazars, they will also be variable sources."," If gamma-ray loud radio galaxies are the misaligned population of blazars, they will also be variable sources." + Kataokaetal.(2010) has recently that NGC 1275 showed a factor of ~2 variation ofreported gamma-ray flux., \citet{kat10} has recently reported that NGC 1275 showed a factor of $\sim2$ variation of gamma-ray flux. +" For other gamma-ray loud radio galaxies, such a significant variation has not observed yet (Abdoetal.2010b).."," For other gamma-ray loud radio galaxies, such a significant variation has not observed yet \citep{abd10_core}." +" Therefore, it is not straightforward to model the of radio galaxies currently."," Therefore, it is not straightforward to model the variability of radio galaxies currently." +" In this paper, we usedvariability the time-averaged gamma-ray flux of gamma-ray loud radio galaxies in theFermi catalog, which is the mean of theFermi 1-year observation."," In this paper, we used the time-averaged gamma-ray flux of gamma-ray loud radio galaxies in the catalog, which is the mean of the 1-year observation." + More observational information (e.g. frequency) is required to model the gamma-ray variability of radio galaxies., More observational information (e.g. frequency) is required to model the gamma-ray variability of radio galaxies. +" Further long term Fermiobservation will be useful and future ground based imaging atmospheric Cherenkov Telescope, Cherenkov Telescope Array would be a key to understanding short period variabilities."," Further long term observation will be useful and future ground based imaging atmospheric Cherenkov Telescope, Cherenkov Telescope Array would be a key to understanding short period variabilities." +" In this we find that the contribution of gamma- loud radio study,galaxies to the unresolved EGRB above 100 MeV is ~25%."," In this study, we find that the contribution of gamma-ray loud radio galaxies to the unresolved EGRB above 100 MeV is $\sim25$." +. Abdoetal.(2010f) recently showed that unresolved blazars can explain only ~22 of the unresolved EGRB by analyzing one year of the blazars., \citet{abd10_marco} recently showed that unresolved blazars can explain only $\sim22$ of the unresolved EGRB by analyzing one year catalog of the blazars. +" Therefore, the origin of rest ~53% of catalogEGRB is still missing."," Therefore, the origin of rest $\sim$ of EGRB is still missing." + Various gamma-ray emitting extragalactic sources have also been discussed as the GeV EGRB origin., Various gamma-ray emitting extragalactic sources have also been discussed as the GeV EGRB origin. +" Those are intergalactic shocks produced by the large scale structure formation (Loeb&Waxman2000;TotaniKitayamaMiniati2002;Keshetetal.2003;Gabici&Blasi 2003), galaxies (Pavlidou&Fields2002;Thompsonetal.2007;etal.2010;Stecker&Venters 2010),, high Galactic latitude pulsars (Faucher-Giguére&Loeb2010;Siegal-Gaskinsetal. 2010),, size AGN (Stawarzetal.2006), radio quiet AGNs kilo-parsec(Inoue(kpc)etal.2008;jetsInoue&Totani2009),, and GeV mass scale DM annihilation or decay (seee.g.Jung- 2009).."," Those are intergalactic shocks produced by the large scale structure formation \citep{loe00,tot00,min02,kes03,gab03}, galaxies \citep{pav02,tho07,bha09_gal,mak10,fie10,ste10}, high Galactic latitude pulsars \citep{fau10,sie10}, kilo-parsec (kpc) size AGN jets \citep{sta06}, radio quiet AGNs \citep{ino08,ino09}, and GeV mass scale DM annihilation or decay \citep[see e.g.][]{jun96,ber00,uli02,oda05,and06,hor06,and07,ahn07,and09,kaw09}." + Fig., Fig. + 5 shows the gamma-ray loud radio galaxy EGRB spectra at each redshift bins., \ref{fig:egrb_z} shows the gamma-ray loud radio galaxy EGRB spectra at each redshift bins. +" Because of EBL, the spectrum above >30 GeV shows the absorbed signature."," Because of EBL, the spectrum above $>30$ GeV shows the absorbed signature." +" Here, the cosmological sources have their evolution peaks at z=1~2 such as cosmic star formation and AGN activity (seee.g.Hopkins&Beacom2006;historyUedaetal.2003)."," Here, the cosmological sources have their evolution peaks at $z=1\sim2$ such as cosmic star formation history and AGN activity \citep[see e.g.][]{hop06,ued03}." +. This means that the gamma-rays from extragalactic sources (e.g. galaxies and AGNs) will experience the EBL absorption., This means that the gamma-rays from extragalactic sources (e.g. galaxies and AGNs) will experience the EBL absorption. +" However, as shown in 5,,Fermi"," However, as shown in Fig. \ref{fig:egrb_z}," + EGRB spectrum does not show such an absorbed Fig.signature., EGRB spectrum does not show such an absorbed signature. +" This might suggest that nearby gamma-ray emitting sources or sources with very hard gamma-ray spectrum would be the dominant population of EGRB above 10 GeV. To address this issue, we should await the EGRB information above 100 GeV by future observations such asFermi."," This might suggest that nearby gamma-ray emitting sources or sources with very hard gamma-ray spectrum would be the dominant population of EGRB above 10 GeV. To address this issue, we should await the EGRB information above 100 GeV by future observations such as." + CTA would also be possible to see the EGRB at much higher energy band., CTA would also be possible to see the EGRB at much higher energy band. + We also need to examine the EBL models at high redshift., We also need to examine the EBL models at high redshift. + It is expected that CTA will see blazars up to z~1.2 at very high energy band >30 GeV (Inoueetal.2010).., It is expected that CTA will see blazars up to $z\sim1.2$ at very high energy band $>30$ GeV \citep{ino10a}. +" Therefore,Fermi and CTA will be a key to understanding the origin of EGRB."," Therefore, and CTA will be a key to understanding the origin of EGRB." + AGN unification scenario explains various properties of AGNs in terms of viewing angle (Urry&Padovani1995)., AGN unification scenario explains various properties of AGNs in terms of viewing angle \citep{urr95}. +". In the scheme of AGN jet unification scenario, FRI and FRII galaxies are thought as misaligned populations of BL Lacs and FSRQs, respectively."," In the scheme of AGN jet unification scenario, FRI and FRII galaxies are thought as misaligned populations of BL Lacs and FSRQs, respectively." +" From Table. 1,,"," From Table. \ref{tb:sam}," +" the mean photon index of FR Is and FRIIs are 2.27 and 2.58, respectively."," the mean photon index of FR Is and FRIIs are 2.27 and 2.58, respectively." +" Therefore, the FRI population tends to have harder spectra than the FRII population as shown in Abdoetal.(2010b).."," Therefore, the FRI population tends to have harder spectra than the FRII population as shown in \citet{abd10_core}." + This trend is also same as that between BL Lacs and FSRQs (Abdoetal.2010g).., This trend is also same as that between BL Lacs and FSRQs \citep{abd10_agn}. + This result would support that FRIs and FRIIs are the misaligned population of BL Lacs and FSRQs., This result would support that FRIs and FRIIs are the misaligned population of BL Lacs and FSRQs. + It is also important to compare the cosmological evolutions of blazars and radio galaxies based on the recentFermi data., It is also important to compare the cosmological evolutions of blazars and radio galaxies based on the recent data. +" Although theoretical blazar GLF model (Inoue&Totani2009) is a compared with theFermi EGRB result and the cumulative brieflysource count distribution of theFermi blazar data (Inoueetal.2010;Inoue&Totani2011; 2011),, comparison in redshift space has not yet been done."," Although a theoretical blazar GLF model \citep{ino09} is briefly compared with the EGRB result and the cumulative source count distribution of the blazar data \citep{ino10a,ino11_err,ino11}, comparison in redshift space has not yet been done." + This is because redshifts of about a half of BL Lac samples has not yet determined (Abdoetal.2010g).., This is because redshifts of about a half of BL Lac samples has not yet determined \citep{abd10_agn}. +" In addition, Inoue&Totani(2009) treated blazars as one population using blazar sequence (Fossatietal.1998;Kuboetal. 1998)."," In addition, \citet{ino09} treated blazars as one population using blazar sequence \citep{fos98,kub98}." +. Blazar GLF models which divide FSRQs and BL Lacs is required such as Dermer(2007) to interpret the unification scenario., Blazar GLF models which divide FSRQs and BL Lacs is required such as \citet{der07} to interpret the unification scenario. +" Since our model in this paper also does not treat FRI and FRII because of small samples, GLF models of gamma-ray loud separatelyradio galaxies dividing these two population are also required."," Since our model in this paper also does not treat FRI and FRII separately because of small samples, GLF models of gamma-ray loud radio galaxies dividing these two population are also required." +" Therefore, redshift information of all blazars and more data of gamma-ray loud radio galaxies would be to make a comparison of the cosmological evolutions of requiredblazars and radio galaxies."," Therefore, redshift information of all blazars and more data of gamma-ray loud radio galaxies would be required to make a comparison of the cosmological evolutions of blazars and radio galaxies." +" In this paper, we studied the contribution of gamma-ray loud radio galaxies to the EGRB by constructing their GLF."," In this paper, we studied the contribution of gamma-ray loud radio galaxies to the EGRB by constructing their GLF." +with the existing picture of angular momentum evolution in young stars. Where the stars are braked effectively by their accretion discs until the disc disperses.,"with the existing picture of angular momentum evolution in young stars, where the stars are braked effectively by their accretion discs until the disc disperses." + Instead. we argue that. for a given association of young stars. position within the CMD is not primarily a function of age. but of accretion history.," Instead, we argue that, for a given association of young stars, position within the CMD is not primarily a function of age, but of accretion history." + We have shown that this hypothesis can. in principle. explain the correlation we observe between rotation rate and position within the CMD.," We have shown that this hypothesis can, in principle, explain the correlation we observe between rotation rate and position within the CMD." + Since variations in accretion history can lead to spreads in radii and luminosity matching those observed. evidence for age spreads in star forming regions must currently be viewed sceptically.," Since variations in accretion history can lead to spreads in radii and luminosity matching those observed, evidence for age spreads in star forming regions must currently be viewed sceptically." + Furthermore. this implies that masses and ages. which are inferred by comparison to non-aecreting evolutionary tracks. may be in error.," Furthermore, this implies that masses and ages, which are inferred by comparison to non-accreting evolutionary tracks, may be in error." + If true. the initial mass functions and ages of star forming regions may need significant revision.," If true, the initial mass functions and ages of star forming regions may need significant revision." +Here. CreaOLp aud (OLTdoy are the correspoucing temyerature fluctuatious seen through the AMIBA natural bean. e=fizf/eposs.,"Here, $\left(\frac{\delta T}{T}\right)_{\rm CMB}$ and $\left(\frac{\delta +T}{T}\right)_{\rm SZ}$ are the corresponding temperature fluctuations seen through the AMIBA natural beam. $x=h\nu/k_B T_{\rm CMB}$." + We have chosen t1e normalizatiou of the beam such that the noise power spectrum is white with Try=Tove?expir)/[expGe)-1]2.22 IN is the Raleigl-Jeans equivalent CMB temperature al v=90 Ghz., We have chosen the normalization of the beam such that the noise power spectrum is white with $T_{RJ}=T_{\rm CMB} x^2 \exp(x)/[\exp(x)-1]=2.22$ K is the Raleigh-Jeans equivalent CMB temperature at $\nu=90$ Ghz. + The factor of 2 is due to the wo polarizatios of the AMIBA experiment., The factor of $2$ is due to the two polarizations of the AMIBA experiment. + For a single clis1 of infinite aperture. with a single pixel detector. te window would be identically one: the scaned image is just the CMB distribution on the sky.," For a single dish of infinite aperture, with a single pixel detector, the window would be identically one: the scanned image is just the CMB distribution on the sky." + Since AMIBA has many detectors. Olle Call ¢'ombine them to either lower the noise. or to boost t ial.," Since AMIBA has many detectors, one can combine them to either lower the noise, or to boost the signal." + We have chosen to use the equivaleu uolse of a single pixel sinele cdisli ex»eriment. aud iized the beam accordingly.," We have chosen to use the equivalent noise of a single pixel single dish experiment, and normalized the beam accordingly." + We calculate he natural beam Wad) by equatio1 (17) and. (16 of Penetal.(2002)., We calculate the natural beam $W_N(l)$ by equation (17) and (46) of \citet{Pen02}. +. The natural beam in jultipole space and real space are sown in fig. ο.., The natural beam in multipole space and real space are shown in fig. \ref{fig:beam}. + TI αςανα beam has a FWHAL of / 2., The natural beam has a FWHM of $2^{'}$. +]t peass at d;c2xA/D;., It peaks at $l_i\simeq 2\pi \lambda/D_i$. + Here. Ac 3.31un is the AMIBA «yerating waveleugth and D; is the distance of the 7th baselines.," Here, $\lambda\simeq 3.3$ mm is the AMIBA operating wavelength and $D_i$ is the distance of the $i$ th baselines." +" The first peaς n""wk™397.2273 corresponds to he sortest base line D,=12 u. At this angular scale. the sum of all the baselines in'oves througlipt by a factor of almost thee (e& 2.77)."," The first peak $l^{\rm peak}_1\sim 2273$ corresponds to the shortest base line $D_1=1.2$ m. At this angular scale, the sum of all the baselines improves throughput by a factor of almost three $\epsilon\simeq 2.77$ )." + This is analogous to laving a uine pixel detector ou a siigle dish., This is analogous to having a nine pixel detector on a single dish. + More detailed definitions of the beams aud strategies are described in deall in (Pen>etal.2002).., More detailed definitions of the beams and strategies are described in detail in \citep{Pen02}. . + Iu our simulated pipeline. we add primary. CMB map fIuctuations to our simlated sky maps us]—ο CMBFAST-seuerated. (Seljak&Zaldarriaga1996) power spectra with saje. Cosmological parameters except [or the use of COBE normalization lor ox. which is sliehly different from ou simulation value.," In our simulated pipeline, we add primary CMB map fluctuations to our simulated sky maps using CMBFAST-generated \citep{CMBFAST} power spectra with same cosmological parameters except for the use of COBE normalization for $\sigma_8$, which is slightly different from our simulation value." + As we will show below. the primary CMB is not the main source of noise for the SZ power spectrum measurement at AMIBA angular scales ({ 2000) aud is neelieible iu SZ cluster searches.," As we will show below, the primary CMB is not the main source of noise for the SZ power spectrum measurement at AMIBA angular scales $l>2000$ ) and is negligible in SZ cluster searches." + So the effect of this inconsistent oy is isienificaut in our aualvysis., So the effect of this inconsistent $\sigma_8$ is insignificant in our analysis. + Adding SZ maps aud CMB maps. we obtain simulated sky maps (fig. 8)).," Adding SZ maps and CMB maps, we obtain simulated sky maps (fig. \ref{fig:tc}) )." + In these naps. SZ striclures. especially those caused by ciffuse IGA are superimposed witli lie ximary CMB.," In these maps, SZ structures, especially those caused by diffuse IGM are superimposed with the primary CMB." + We then couvove these maps witli the natural beam., We then convolve these maps with the natural beam. +" The beam function decreases 4uickly to zero towards Iaορ alguar scales where the primary CMB clominates. so it efficiently filte5 ost primary CMB st""uctures arger than this scale."," The beam function decreases quickly to zero towards large angular scales where the primary CMB dominates, so it efficiently filters most primary CMB structures larger than this scale." + We then add the noise given by equ. (7)), We then add the noise given by eqn. \ref{eqn:noise}) ) + to this map., to this map. + The normalization in tlie beaniug aud filtering (described below) process is arbit‘ary., The normalization in the beaming and filtering (described below) process is arbitrary. + We choose the uormalizaloh ln sich a way that the power spectrum of the map at the scale of the peak response does not change after )eaniiue or filteriug., We choose the normalization in such a way that the power spectrum of the map at the scale of the peak response does not change after beaming or filtering. + It corresponds to ποιατις the elobal maximiin of the beam fuuctiou to be unity., It corresponds to normalizing the global maximum of the beam function to be unity. + Under such normalization. when we add instrumental noise o the simulated CMB+SZ map. the uoise Is depressed by a uormalizatiou [actor e2 2.77.," Under such normalization, when we add instrumental noise to the simulated CMB+SZ map, the noise is depressed by a normalization factor $\epsilon\simeq 2.77$ ." + Our 1.14X map has 20187 pixels. so the white noise⋅ dominates⋅ on⋅ small scales.," Our $1.19^{\circ}$ map has $2048^2$ pixels, so the white instrumental noise dominates on small scales." + For a⊲ 20 ∙∣hours/deg? survey.9. the d temperatureinstrumentalspersion ofthe noise [uctuation oy =TusyOrelpicfue)~~0.002Ye {μίκο," For a $20$ $/{\rm deg}^2$ survey, the dispersion of the noise temperature fluctuation $\sigma_N=T_{\rm sys}/\sqrt{(2\Delta \nu t_{\rm pixel}}/(\eta +T_{\rm RJ} \epsilon) \sim 0.002\gg \bar{y}$." + Is Lie observing time for each pixel., $t_{\rm pixel}$ is the observing time for each pixel. + So all sigals are hidden uucer the instrumental noise., So all signals are hidden under the instrumental noise. + One nees further filteriug to obtain au image that is uot overwhelmed by the πια] scale noise., One needs further filtering to obtain an image that is not overwhelmed by the small scale noise. + Ignoriug the CMB fIuctuations. we kuow that point sources have the shape of the beam.," Ignoring the CMB fluctuations, we know that point sources have the shape of the beam." + A point sources optimized search would convolve the natural map with the shape of the beam. aud peaksin this ma» correspond to the," A point sources optimized search would convolve the natural map with the shape of the beam, and peaksin this map correspond to the" +Satellite galaxies have attracted an enormous amount of observational ancl theoretical study over the past decade.,Satellite galaxies have attracted an enormous amount of observational and theoretical study over the past decade. + This is partly a consequence of the uerarchical nature of galaxy formation in the currentIv-po»ular ACDM moclels. within which some or all dwarf satellites may represen Loft-«ver building blocks from an earlier assembly phase of |arex| galaxies.," This is partly a consequence of the hierarchical nature of galaxy formation in the currently-popular $\Lambda$ CDM models, within which some or all dwarf satellites may represent left-over building blocks from an earlier assembly phase of large galaxies." + Fhese models also suller from the so-calle ‘substructure problem! (IxIvpinet.al.19990:Mooreet 1999). in that the predicted: numbers. of low-mass dark mater haloes are [ar larecr than the observed. numbers of low-uminositv agalaxies that mighto naurally be expected to occupy them.," These models also suffer from the so-called `substructure problem' \citep{klyp99,moor99}, in that the predicted numbers of low-mass dark matter haloes are far larger than the observed numbers of low-luminosity galaxies that might naturally be expected to occupy them." + The problem has been addressed in the lates mockels (Bensonetal.2002:Simon&Geha20Yr:IxoposovClal.2000) bs decreasing or completely: suppressing the star-Lormation (SE) elliciency in low mass halos.," The problem has been addressed in the latest models \citep{bens02,simo07,kopo09} by decreasing or completely suppressing the star-formation (SF) efficiency in low mass haloes." + One grea success of these models was that they predicted he existence outra-low mass chwarl galaxies substantially fainter than the canonical dwarf spheroidals., One great success of these models was that they predicted the existence of ultra-low mass dwarf galaxies substantially fainter than the canonical dwarf spheroidals. + These ultra-faint cbwarfs were subsequentIy discovered. primarily twoue1 searches of data from the Sloan Digital Sky Survey (SDSS) (Willmanetal.2005:IxXoposov2007:Walshct2 )07).," These ultra-faint dwarfs were subsequently discovered, primarily through searches of data from the Sloan Digital Sky Survey (SDSS) \citep{will05,kopo07,wals07}." +. Satellites are also potentially of im»ortance in other areas of galaxy physics., Satellites are also potentially of importance in other areas of galaxy physics. + Merging of. |w-mass satellites Wil1 their central hosts (minor mergers) is one route OL the formation of thick disk components. and for »ulding bulges.," Merging of low-mass satellites with their central hosts (minor mergers) is one route for the formation of thick disk components, and for building bulges." + For example. Domínguez-DPamero&Bal-celIs.(2008) cite repeated. minor merger episodes as hic‘iv preferred. mechanism through which buσος CD Brow Wirout destroving the surrounding disk. thus preserving hic| strong. correlation between disk anc bulge colours lin individual galaxies found by these authors.," For example, \citet{domi08} cite repeated minor merger episodes as their preferred mechanism through which bulges can grow without destroying the surrounding disk, thus preserving the strong correlation between disk and bulge colours within individual galaxies found by these authors." + Another portant problem concerns the continued. gas supply to ege clisk galaxies. first noted by Larsonetal.(1980).," Another important problem concerns the continued gas supply to large disk galaxies, first noted by \citet{lars80}." +. These LOLS estiniae that the Alilky Way (MW) will consume current disk gas reservoir in —2 Cr. and that he etüvalent timescale for 36 external galaxies is little longer By 1a mean of 3.99 Gye. and certainty less than a Llubxe in16.," These authors estimate that the Milky Way (MW) will consume the current disk gas reservoir in $\sim$ 2 Gyr, and that the equivalent timescale for 36 external galaxies is little longer, with a mean of 3.9 Gyr, and certainly less than a Hubble time." + This points clearly to the need for continued gas suphe ο support ongoing SL in most or all spiral galaxies., This points clearly to the need for continued gas supply to support ongoing SF in most or all spiral galaxies. +" Stich continued or even increasing supply is also indicated by he Ix-chwarf. metalicity distribution in the MW. investigaed w Casuso&Beckman(2004).. who also conclude that he ""burstiness apvent in the SE history of the MW. may indicate significant gas accretion events."," Such continued or even increasing supply is also indicated by the K-dwarf metallicity distribution in the MW, investigated by \citet{casu04}, who also conclude that the `burstiness' apparent in the SF history of the MW may indicate significant gas accretion events." + However. Grcevich&Putman(2009) estimate the total gas mass in the current populations of satellites around the MW. anc M1. (not including the Alagellanie Clouds). finding a total of only 1-2 107 M. for the MW.," However, \citet{grce09} estimate the total gas mass in the current populations of satellites around the MW and M31 (not including the Magellanic Clouds), finding a total of only 1 - $\times10^8$ $_{\odot}$ for the MW." +" Thev conclude that this is too little to provide a satisfactory explanation of the deficiency of low-metallicity να, stars in the Galactic disk (the", They conclude that this is too little to provide a satisfactory explanation of the deficiency of low-metallicity dwarf stars in the Galactic disk (the +∫⊔↕⋮∣∕⊓⊽∪∣⇁↙∣⋮∖∣∷≼∣⋜↧↕⋜⋯⋅↖↽∶∶↴↜⊾↕∪↴⋝∏↕⋜∐⋅↸⊳↕∏↴∖↴↑↸∖↥⋅↴∖↴∶↕∐≼∐↖↽↕≼↧∏⋜↧⊓⋀∖≼∣≼∩⊔∩⋝∶⊸∖," \citep{Alpar82,Backer82}. \citep[see, e.g.,][and references therein]{Poutanen06}." +≓↥⋅⋜↧⋅↖↽↴∖↴∶↴⋝↕∐⋜∐⋅↕↸∖↴∖↴∙↴∖↴↑⋜∐⋅↴∖↴∶∐↸∖∏⊓⋅∪∐ (Alyaretal.1982:Backer1982).. (SCO.referencestherein).. 100 eal.20091. (Altamiranoetal.2008)..," $\sim100$ \citep{Galloway07a,Gavriil07,Casella08,Altamirano08b,Patruno09}, \citep{Altamirano08b}." + LAINBs show pulsations is still uuder debate., LMXBs show pulsations is still under debate. + One of ie proposed scenarios for the uou-pulsatiug svstenis ds iat the neutrou-star maenetic field could be eniporarilv mnned by the accreted matter (C1nuuiueseal.2001)., One of the proposed scenarios for the non-pulsating systems is that the neutron-star magnetic field could be temporarily buried by the accreted matter \citep{Cumming01}. +. Tt the time-averaged mass accretion rate of f10 aACCICTOYS is relatively high. the accreted mater can burv the field ποιοςrout the life of the N-rav bWav SO he neutron star docs not pulsate.," If the time-averaged mass accretion rate of the accretors is relatively high, the accreted matter can bury the field throughout the life of the X-ray binary so the neutron star does not pulsate." + Wowever. if he average accrelon rate Is low enough that the magneic field can diffuse ποιοi the accreted matter faster hau it is nried. then jese systems will probably exhibi oilsatious.," However, if the average accretion rate is low enough that the magnetic field can diffuse through the accreted matter faster than it is buried, then these systems will probably exhibit pulsations." + Another of the proposed scenarios that explaiis why the majority of neuron stars in LAINBs do not pulsate is that of Lambetal.(2009):: if neutrou stars 1udereo loue periods of accretioi (at lugh rates). then tici nuagnetie poCR are naturalv forced to aligu to their spin axes as they spi1 up.," Another of the proposed scenarios that explains why the majority of neutron stars in LMXBs do not pulsate is that of \citet{Lamb09a}; if neutron stars undergo long periods of accretion (at high rates), then their magnetic poles are naturally forced to align to their spin axes as they spin up." + When the acerction rate decreases. the maguetic poles can nove away from the roation axis (due to maeuctic ¢ipole and other braking orques which cause neutroji stars to spin down) and oscillations powered by accretion slowd become visible.," When the accretion rate decreases, the magnetic poles can move away from the rotation axis (due to magnetic dipole and other braking torques which cause neutron stars to spin down) and oscillations powered by accretion should become visible." + Both Cumunineetal.(2001) and Laubetal.(2009) present clear predictio we should fiu pulsatious in iuauv (if not all) of the svstenis which show low time-averaged mass accretio- rates., Both \citet{Cumming01} and \citet{Lamb09a} present clear predictions: we should find pulsations in many (if not all) of the systems which show low time-averaged mass accretion rates. + Verv recenlv. a Chandra ταν observatiocereserenudipitouxlv discovered a new XN-rav trausicut d the elobuluw custer NGC 6110. (IIeiikeetal.2009c:fiedasNGC6110. X-2)..," Very recently, a Chandra X-ray observationserendipitously discovered a new X-ray transient in the globular cluster NGC 6440 \citep[identified as NGC~6440 X-2]{Heinke09a, Heinke09b, + Heinke09c}." +" RXTE and Swift follow-u observations notonly showed thatthis source has a very low time.average accretion: rate (z-y2++102 LAL, ", RXTE and Swift follow-up observations notonly showed thatthis source has a very low time–average accretion rate \citep[$\lesssim2 \cdot 10^{-12}$ $_{\odot}$ +while the shallower CANDELS aud ERS fields coutribute a total of 16 sources (all with Toop&26 mag).,"while the shallower CANDELS and ERS fields contribute a total of 16 sources (all with $H_{160,AB}\lesssim26$ mag)." + Upon inspection of their images. it turus out that all these brighter sources are very well detected in tle IRAC data. even in the shallow 8.0721. band.," Upon inspection of their images, it turns out that all these brighter sources are very well detected in the IRAC data, even in the shallow $\micron$ band." + Their measured Ii[3.6] colors are in the range 1.6 23. which. for a loLO source. would correspond to a UV continua slope 3zo 0.2. or a dust reddening wih Ap>1.6 mae.," Their measured $H_{160}-[3.6]$ colors are in the range $1.6-4.3$ , which, for a $z\sim10$ source, would correspond to a UV continuum slope $\beta\gtrsim-0.2$ , or a dust reddening with $A_V>1.6$ mag." + Cuven that galaxies at the bright ei of the :~7 population are measured to have very low extinction values aud continua slopes of 9&2.040.2 (scee.g. 2001b).. the extremely red colors of these ealaxies rules out :29 solutions with auv sensible SED.," Given that galaxies at the bright end of the $z\sim7$ population are measured to have very low extinction values and continuum slopes of $\beta\simeq-2.0\pm0.2$ \citep[see e.g.][]{Bouwens10b,Finkelstein10,Dunlop11,Wilkins11b}, the extremely red colors of these galaxies rules out $z\gtrsim9$ solutions with any sensible SED." + All sources with IRAC detections are thus removed from our sample of potential 2~10 ealaxies. reclucine the sample to one single candidate in the IIUDE. previously reported in Bowensctal.(2011a.secFig- and Table 20}.," All sources with IRAC detections are thus removed from our sample of potential $z\sim10$ galaxies, reducing the sample to one single candidate in the HUDF, previously reported in \citet[][see Figure \ref{fig:stampz10} and Table \ref{tab:phot} ." + The 16 removed sources are shown iu the appendix in Figure All and listed in Table AL., The 16 removed sources are shown in the appendix in Figure \ref{fig:stampsContamin} and listed in Table \ref{tab:photContamin}. + huages of the only possible 210 ealaxy canclidate are shown in Fieure L., Images of the only possible $z\sim10$ galaxy candidate are shown in Figure \ref{fig:stampz10}. + The source is detected at 6.3 σ in yoy (ueasured in circular apertures of 0725 radius)., The source is detected at 6.3 $\sigma$ in $H_{160}$ (measured in circular apertures of $0\farcs25$ radius). + This is higher but completely consistent with the Bowwensctal.(20114). significance estimates. which are based on smaller apertures.," This is higher but completely consistent with the \citet{Bouwens11} significance estimates, which are based on smaller apertures." + As can he seen. the source is not significantly detected im amy other baud.," As can be seen, the source is not significantly detected in any other band." +" Its value of opt2.77 is very close. but just below he Bit of 42,,=2.8."," Its value of $\chi^2_{opt}=2.77$ is very close, but just below the limit of $\chi^2_{cut}=2.8$." + This is mainly due to a 1.56 Hx excess dn fees. Which appears to be due to an extended structure in the backerouud of that image.," This is mainly due to a $\sigma$ flux excess in $i_{775}$, which appears to be due to an extended structure in the background of that image." + When adopting smaller apertures. the value is fouud o be reduced. indicating also that this wexcess of flix is rot associated with the source itself.," When adopting smaller apertures, the $\chi^2_{opt}$ value is found to be reduced, indicating also that this excess of flux is not associated with the source itself." + The Spitzer TRAC 3.6; data shows some flux from a nearby source., The Spitzer IRAC $\mu$ m data shows some flux from a nearby source. + Towever. after subtraction of all the 1cieliboriug sources in RAC. the candidate is undetected (0.090). with a 27 upper limit on its ΠΑΟ [3.6] uaenitude of >27.2 imag AB (Gonzalez ct al.," However, after subtraction of all the neighboring sources in IRAC, the candidate is undetected $0.09\sigma$ ), with a $2\sigma$ upper limit on its IRAC [3.6] magnitude of $>27.2$ mag AB (Gonzalez et al.," + in xep.)., in prep.). + This thus corresponds to Lyou[3.6]1.6 at 2c. which is much. bluer (bv >0.1 mag) than the vpical low-vedshift coutaminauts that we culled from our sample (see also next section).," This thus corresponds to $H_{160}-[3.6]<1.6$ at $2\sigma$, which is much bluer (by $>0.4$ mag) than the typical low-redshift contaminants that we culled from our sample (see also next section)." + After adding the newly acquired IRAC data from the IUDFLO program (Spitzer xoposal 70115. PE Labbe) to the GOODS IRAC data and removing neieliboriug sources. the +10 candidate is also undetected at [L5] (S/N=0.3) providing added weight to the likehhood of it beme at high rather than ow reshift.," After adding the newly acquired IRAC data from the IUDF10 program (Spitzer proposal 70145, PI: Labbe) to the GOODS IRAC data and removing neighboring sources, the $z\sim10$ candidate is also undetected at [4.5] $S/N=0.3$ ) providing added weight to the likelihood of it being at high rather than low reshift." + We derive the photometric redshift of the candidate using the code ZEBRA (Feldmannetal.2006:Ocschetal.20100) with svuthetic stellar population models roni Druzual&Charlot(2003) to which we added iebular contiuuun aud line enmüssou following. οo.g..οἱ Schaerer&deBarros(2009).," We derive the photometric redshift of the candidate using the code ZEBRA \citep{Feldmann06,Oesch10c} with synthetic stellar population models from \citet{Bruzual03} to which we added nebular continuum and line emission following, e.g., \citet{Schaerer09}." +". Using the full 11 baud fluxes and fux errors. we derive a photometric redshift or this source of z,5,,=LOLYY. with a likelihood of a ow redshift solution at τσ8 of <6%."," Using the full 11 band fluxes and flux errors, we derive a photometric redshift for this source of $z_{phot} = 10.4^{+0.5}_{-0.4}$, with a likelihood of a low redshift solution at $z<8$ of $<6\%$ ." + The full spectral cherey distribution (SED) of the source aud its redshift ikelihood. function are shown in Figure 5.., The full spectral energy distribution (SED) of the source and its redshift likelihood function are shown in Figure \ref{fig:JdropSED}. + The best-fit SED correspouds to a very voung. dust-free. star-liurst (see also Conzalez et al.," The best-fit SED corresponds to a very young, dust-free star-burst (see also Gonzalez et al." + iu prep.)., in prep.). + The best low redshift solution is found at 4.=2.7. or an evolved. very low mass galaxy SED (M=3<10% M.) with moderate extinction.," The best low redshift solution is found at $z_{lowz}=2.7$, for an evolved, very low mass galaxy SED $M=3\times10^8$ ) with moderate extinction." +" Iuterestiuglv. this SED is expected to be detected only at ~1.50 in J1»5. but it nevertheless has a siguificantly higher 4? value (13.7 compared to Az,= 7.0)."," Interestingly, this SED is expected to be detected only at $\sim1.5\sigma$ in $J_{125}$, but it nevertheless has a significantly higher $\chi^2$ value (13.7 compared to $\chi^2_{best}=7.0$ )." + Deeper IST data shortward of the break would be extremely useful in order to constrain the possible non-detection of the source shortward of l.lgn., Deeper HST data shortward of the break would be extremely useful in order to constrain the possible non-detection of the source shortward of 1.4. + As with other high-redshift catalogs. the added shorter-waveleneth optical/uear-IR data would play a kevrole in helping to further tighten its photometric redshift measurement.," As with other high-redshift catalogs, the added shorter-wavelength optical/near-IR data would play a keyrole in helping to further tighten its photometric redshift measurement." + The properties of the 16 intermediate brightuess sources that didnot passthe IRAC nou-detection criteria are discussed in the appendix., The properties of the 16 intermediate brightness sources that didnot passthe IRAC non-detection criteria are discussed in the appendix. + From fitting their SEDs. these sources are found to be mostly massive galaxies GE5<1029 3) with obscuredbut evolved stellar populations at.—2. L.," From fitting their SEDs, these sources are found to be mostly massive galaxies $M>5\times10^{10}$ ) with obscuredbut evolved stellar populations at $z\sim2-4$ ." + Tuterestinely. all these galaxies areesseutiallv. limited to Πω~2126 imag.," Interestingly, all these galaxies areessentially limited to $H_{160}\sim24-26$ mag." + This is ~l mag brighter, This is $\sim$ 1 mag brighter +hey all are good enough to be used. being one slightly vetter than the other depending ou the noise prescut in the data aud the eveuts that are being searched for.,"they all are good enough to be used, being one slightly better than the other depending on the noise present in the data and the events that are being searched for." + We note rere that a possibility to optimize the sensitivity to real rausits is fo limit the automatic scale-selection to trausit-duratious that may be expected from the combination of a elt curves known stellar paramcters (nass. radius) aud the period that is beime searched (e.g. different selections iav be used for different rauges of period-searches).," We note here that a possibility to optimize the sensitivity to real transits is to limit the automatic scale-selection to transit-durations that may be expected from the combination of a light curve's known stellar parameters (mass, radius) and the period that is being searched (e.g. different scale-selections may be used for different ranges of period-searches)." + Tn Table 2 we compare the detections in the 999 test curves as found by TRUFAS with results from the five cals that participated in BTL (see their Tables 1 aud 2). where team 1 used a correlation with a sliding transit cluplate. team 2 used a search for box-shaped rausits with lowpass filtering and brokeu-line detrending. team 3 used he BLS algorithm (Novaccs et al. 20023) ," In Table 2 we compare the detections in the 999 test curves as found by TRUFAS with results from the five teams that participated in BT1 (see their Tables 1 and 2), where team 1 used a correlation with a sliding transit template, team 2 used a search for box-shaped transits with lowpass filtering and broken-line detrending, team 3 used the BLS algorithm (Kováccs et al. \cite{kovacs}) )" +on ight curves that had been detrended through a fitting of 200 harmonics. team l1 emploved a matched filter with imnage-processing detreuding aud team 5 used the yox-shaped transit finding aleorithim by Aigrain Tnwin (2001)) with an iterative 1D filtering.," on light curves that had been detrended through a fitting of 200 harmonics, team 4 employed a matched filter with image-processing detrending and team 5 used the box-shaped transit finding algorithm by Aigrain Irwin \cite{aigrain_irwin}) ) with an iterative 1–D filtering." + For anv further details on hese methods. see DTI.," For any further details on these methods, see BT1." + For a discussion of the detections. we should remember that BTL included iu these light curves 21 transiting planctary systems. as well as 11 low-ampltuce eclipsing binaries and one eclipsing binary in a triple system: all of thei were considered. trausit-like signals.," For a discussion of the detections, we should remember that BT1 included in these light curves 21 transiting planetary systems, as well as 11 low-amplitude eclipsing binaries and one eclipsing binary in a triple system; all of them were considered transit-like signals." + From the 21 planets. 7 were not detected by any of the five algoritliaus in BTL.DTI. 9 planetsplaucts wherewl recovered by all algorithims.algoritl anel 5 only by some of them.," From the 21 planets, 7 were not detected by any of the five algorithms in BT1, 9 planets where recovered by all algorithms, and 5 only by some of them." + No clear “winning algorithm could be ideutified in BT. though team 3 had the best combination of detection of real events aud avoidance of false detections.," No clear “winning algorithm” could be identified in BT1, though team 3 had the best combination of detection of real events and avoidance of false detections." + The performance of TRUFAS is well along the results of the best of these algorithms. except for the one case of lieht-curve 207. which was due to the limiting of the periodicity search to 50 davs due to TRUFAS' requirement of at least 3 ransit eveuts;," The performance of TRUFAS is well along the results of the best of these algorithms, except for the one case of light-curve 207, which was due to the limiting of the periodicity search to 50 days due to TRUFAS' requirement of at least 3 transit events." + Ou the other hand. TRUFAS ound all planets among those that," On the other hand, TRUFAS found planets among those that" +spectrum fc’) from such a hard electron distribution would follow a power law /ic(v)οv~ with index a=0.25 for y«dyn.,"spectrum $j_{\rm IC}(\nu)$ from such a hard electron distribution would follow a power law $j_{\rm IC}(\nu) \propto + \nu^{-\alpha}$ with index $\alpha =0.25$ for $\nu \ll 4\,\gamma_{\rm + b}^2 \nu_s$." + Integrating Eq. (??)), Integrating Eq. \ref{diff}) ) + over y gives the number of particles along a field line N=Q/Q., over $\gamma$ gives the number of particles along a field line $N\simeq Q/\Omega$. +" Electrons. escaping quasi-monoenergetically with y,~107 from the acceleration mechanism and encountering the Comptonized disk photons (v> vw, can Thomson upseatter them to the TeV regime. producing a power law-like energy distribution above αγρνν with index a.~1.2. consistent with the value 1.22x0.15 derived for the HESS 2005 observations of M87 (Aharonian et al."," Electrons, escaping quasi-monoenergetically with $\gamma_b \sim 10^7$ from the acceleration mechanism and encountering the Comptonized disk photons $\nu >\nu_s$ ), can Thomson upscatter them to the TeV regime, producing a power law-like energy distribution above $4 \gamma_{\rm b}^2 \nu_s$ with index $\alpha_c \sim 1.2$, consistent with the value $1.22 \pm 0.15$ derived for the HESS 2005 observations of M87 (Aharonian et al." + 2006). (, 2006). ( +4) The number of escaping particles per unit time is AeHCOypMyp/ Το.,"4) The number of escaping particles per unit time is $n_e \sim + n(\gamma_b)\gamma_b/\tau_{\rm esc}$ ." + Thus. within some time smaller than the cooling time (At=ρίωρ« 1). we accumulate Ns~n;Ar particles that can IC. upseatter Comptonized disk photons.," Thus, within some time smaller than the cooling time $\Delta t=\rho\,t_{\rm cool}, \rho<1$ ), we accumulate $N_b \sim n_e \Delta t$ particles that can IC upscatter Comptonized disk photons." +" We can roughly estimate the associated TeV luminosity from Lic~NpPic. where Pic=L3opey,Uj i5 the single particle Compton power per unit volume."," We can roughly estimate the associated TeV luminosity from $L_{\rm IC} \sim N_b\,P_{\rm IC}$, where $P_{\rm IC}=1.3 \sigma_T c \gamma_b^2 U_{\rm ph}$ is the single particle Compton power per unit volume." +" This gives ""l", This gives ). +eptonic (12)To achieve a Compton luminosity comparable to the observed TeV luminosity of Lj;y=3PETSx erg/sCAharonian et al., To achieve a Compton luminosity comparable to the observed TeV luminosity of $L_{\rm TeV} \simeq 3 \times 10^{40}$ erg/s (Aharonian et al. + 2006). we thus need NV3x10°°/p particles along field ," 2006), we thus need $N \sim 3 \times 10^{36}/\rho$ particles along field lines." +"Denoting the relevant acceleration volume by AV.~qimlines.-Ar, with characteristic length scale Ar=|y/(dy/dr)|~(volyn and 4«I. the corresponding kinetic energy. density UYys,cC/AV (for py2 107) i still well below the energy density Β΄(Απ) of the magnetic field. suggesting that the presumed (quasi force-free) MHD field structure is still a valid approximation (cf."," Denoting the relevant acceleration volume by $\Delta V \sim \eta\,r_{\rm L}^2\,\Delta r$, with characteristic length scale $\Delta r =|\gamma/(\pad\gamma/\pad r)| \sim (\gamma_0/\gamma_b)\, + r_{\rm L}$ and $\eta < 1$, the corresponding kinetic energy density $n(\gamma_b)\gamma_b^2 m_e\,c^2/\Delta V$ (for $\rho\,\eta \gppr + 10^{-5}$ ) is still well below the energy density $B^2/(8 \pi)$ of the magnetic field, suggesting that the presumed (quasi force-free) MHD field structure is still a valid approximation (cf." + also Osmanov et al., also Osmanov et al. + 2007). (, 2007). ( +5) In principle. Τεν gamma-rays can be strongly attenuated due to photon-photon pair production in. the background disk photon field.,"5) In principle, TeV gamma-rays can be strongly attenuated due to photon-photon pair production in the background disk photon field." + The narrow dependence of the cross-section c4. on the product of photon energies implies that VHE photons of energy E interact most efficiently with infrared background photons of energy ει=(ITeV/E) eV. The optical depth τ for à y-ray photon in a background field of infrared luminosity Lj; and size Λιν thus becomes (cf., The narrow dependence of the cross-section $\sigma_{\gamma\gamma}$ on the product of photon energies implies that VHE photons of energy $E$ interact most efficiently with infrared background photons of energy $\epsilon_{\rm IR}\simeq (1 \mathrm{TeV}/E)$ eV. The optical depth $\tau$ for a $\gamma$ -ray photon in a background field of infrared luminosity $L_{\rm IR}$ and size $R_{\rm IR}$ thus becomes (cf. +" NAO7) Ru)= 9.25 tC IU indicating that due to its low bolometric luminosity M87 could be well transparent to VHE gamma-rays. even if almost all of the observed infrared luminosity Lj~101! ere/s (Whysong Antonucer 2004) is (somewhat unrealistically) taken to be produced on a seale Rig~(rj60r,."," NA07) ) = 0.2 ) ), indicating that due to its low bolometric luminosity M87 could be well transparent to VHE gamma-rays, even if almost all of the observed infrared luminosity $L_{\rm IR} \simeq 10^{41}$ erg/s (Whysong Antonucci 2004) is (somewhat unrealistically) taken to be produced on a scale $R_{\rm IR} \sim r_{\rm l}\sim 60\,r_s$." + Note that even if 7 would become larger than one. y-rays from the last transparent layer are still able to escape. so that the VHE flux would not simply decrease exponentially by exp(—r). but only by a factor of ~τ (NAO7). (," Note that even if $\tau$ would become larger than one, $\gamma$ -rays from the last transparent layer are still able to escape, so that the VHE flux would not simply decrease exponentially by $\exp(-\tau)$, but only by a factor of $\sim \tau$ (NA07). (" +6) The number of electrons escaping quasi-monoenergetically from the acceleration mechanism in the vicinity of the light cylinder is of order ΝΕ.,6) The number of electrons escaping quasi-monoenergetically from the acceleration mechanism in the vicinity of the light cylinder is of order $N_b$. +" Once these energetic. particles. encounter. non-vanishing perpendicular and/or the turbulent plasma magnetic fields. they can produce synchrotron emission arising as L,ο below. and decaying exponentially above the peak frequency Vonc30/5x10/3(8sina/1G) Μεν with a total luminosity of order LZ,PonNp~0.06Lic(Bsiner. where P, is the single particle synchrotron power."," Once these energetic particles encounter non-vanishing perpendicular and/or the turbulent plasma magnetic fields, they can produce synchrotron emission arising as $L_{\nu} \propto \nu^{1/3}$ below, and decaying exponentially above the peak frequency $\nu_{\rm + syn} \sim 50 \,(\gamma_b/5\times10^7)^2 (B\,\sin\alpha/1\, + \mathrm{G})$ MeV with a total luminosity of order $L_{\rm syn} \sim + P_{\rm syn}\,N_b \sim 0.06\,L_{\rm IC}\, (B \sin\alpha)^2$, where $P_{\rm syn}$ is the single particle synchrotron power." + In order to satisfy the restrictions imposed by the existing (yet non-contenporaneous) upper limit on the M87 flux in the EGRET energy band above 100 MeV (e.g.. Reimer et al.," In order to satisfy the restrictions imposed by the existing (yet non-contemporaneous) upper limit on the M87 flux in the EGRET energy band above 100 MeV (e.g., Reimer et al." + 2003). the effectively encountered fields should be smaller than ~| Gauss.," 2003), the effectively encountered fields should be smaller than $\sim 1$ Gauss." + This seems consistent with independent estimates suggesting a strength of the random field component close to the black hole of below one Gauss (NAO7)., This seems consistent with independent estimates suggesting a strength of the random field component close to the black hole of below one Gauss (NA07). + The overall spectral energy distribution in M87 is then likely to consist ik a number of different contributions. involving also. other (Georgaiopoulos et al.," The overall spectral energy distribution in M87 is then likely to consist of a number of different contributions, involving also other leptonic (Georganopoulos et al." + 2003: NÀO7) and perhaps even hadronic (Reimer et al., 2003; NA07) and perhaps even hadronic (Reimer et al. + 2004) processes., 2004) processes. + If so. then no straighforward X-ray-TeV correlation might be expected. (," If so, then no straighforward X-ray–TeV correlation might be expected. (" +7) As shown above. accelerating particles up to the light cylinder typically takes à time 71/c. suggesting a characteristic variability time scale for M87 of t.=2hx1.7 days. well consistent with the observed TeV time scale of Ar~2 days. a fact that may further validate the assumptions of the presented model.,"7) As shown above, accelerating particles up to the light cylinder typically takes a time $r_{\rm L}/c$, suggesting a characteristic variability time scale for M87 of $t_v \simeq \frac{r_{\rm L}}{c} + \sim \frac{5\,r_s}{c} \simeq 1.7$ days, well consistent with the observed TeV time scale of $\Delta t \sim 2$ days, a fact that may further validate the assumptions of the presented model." + VHE radiation from. low-luminous. non-blazar AGN jet sources like M87 could provide an ideal test laboratory for the analysis of particle acceleration processes close to the supermassive black hole event horizon.," VHE radiation from low-luminous, non-blazar AGN jet sources like M87 could provide an ideal test laboratory for the analysis of particle acceleration processes close to the supermassive black hole event horizon." + In blazars with their relativistic jets pointing towards us. most of these traces are likely to be masked by strong relativistic beaming effects. while for luminous quasars internal absorption of gamma-rays becomes dominant.," In blazars with their relativistic jets pointing towards us, most of these traces are likely to be masked by strong relativistic beaming effects, while for luminous quasars internal absorption of gamma-rays becomes dominant." + Based on a simple toy model we have shown that efficient centrifugal acceleration of electrons in the vicinity of the light cylinder could provide a natural explanation for variable (time scale of one day) VHE emission with a hard inverse Compton spectrum as observed in M87., Based on a simple toy model we have shown that efficient centrifugal acceleration of electrons in the vicinity of the light cylinder could provide a natural explanation for variable (time scale of one day) VHE emission with a hard inverse Compton spectrum as observed in M87. + Our models fits well with other evidence for advection-dominated accretion in. M87 and may indeed be regarded as providing meisome further corroboration for the presence of such modes in highly underluminous As always. there are a number of subtleties whose impact on the presented results need to be explored in more details including general relativistic effects. anisotropic scattering modifications. quasi rigid rotation andplasma instabilities.," Our models fits well with other evidence for advection-dominated accretion in M87 and may indeed be regarded as providing some further corroboration for the presence of such modes in highly underluminous As always, there are a number of subtleties whose impact on the presented results need to be explored in more details including general relativistic effects, anisotropic scattering modifications, quasi rigid rotation andplasma instabilities." + Theextent to which our conclusions might be affected may require fully relativistic modelling., Theextent to which our conclusions might be affected may require fully relativistic modelling. + Yet. given the demonstrated potential of centrifugal acceleration and our current understanding of relativistic jet formation. this may represent a program worth pursuing.," Yet, given the demonstrated potential of centrifugal acceleration and our current understanding of relativistic jet formation, this may represent a program worth pursuing." +model for the ME.,model for the MF. + However. DAL haloes also possess angular momentum. and taking into account its ellect on the eravitational collapse of rotating shells has been shown to have a detectable consequence on the ME (77).," However, DM haloes also possess angular momentum, and taking into account its effect on the gravitational collapse of rotating shells has been shown to have a detectable consequence on the MF ." +.. Phus. an investigation of the distribution of angular momentum. and its connection with mass is particularly useful. as high-resolution. simulations of increasing resolution will produce MES with a very small statistical An exact determination. of the shape of the DAL halo spin PDE can also have important consequences for the abundance. ofgalarics.. M the latter form preferentially within high-spin DM haloes (??).," Thus, an investigation of the distribution of angular momentum, and its connection with mass is particularly useful, as high-resolution simulations of increasing resolution will produce MFs with a very small statistical An exact determination of the shape of the DM halo spin PDF can also have important consequences for the abundance of, if the latter form preferentially within high-spin DM haloes ." +. Finally. it is of great. importance also in models of the formation and evolution. of central Black Holes.," Finally, it is of great importance also in models of the formation and evolution of central Black Holes." + 1n collapse models where most of the barvons’ specific angular momentum is a {fixed fraction of that of their host. clark matter haloes ).. the angular momentum and extent of the gaseous central accretion dise are strongly dependent on their's spin.," In collapse models where most of the baryons' specific angular momentum is a fixed fraction of that of their host dark matter haloes , the angular momentum and extent of the gaseous central accretion disc are strongly dependent on their's spin." + find that the central density of the cise varies aspycAt thus the initial rate of accretion of the central Black Hole turns out to be a very sensitive function of the spin A For all these reasons. investigations of the origin of the angular momentum growth and of the spin PDE of DM haloes. ancl of their evolution. can have an impact on many different open issues in large-scale structure formation and evolution.," find that the central density of the disc varies as$\rho_{0} \simeq\lambda^{-4}$, thus the initial rate of accretion of the central Black Hole turns out to be a very sensitive function of the spin $\lambda$ For all these reasons, investigations of the origin of the angular momentum growth and of the spin PDF of DM haloes, and of their evolution, can have an impact on many different open issues in large-scale structure formation and evolution." +" ""Theoretical investigations predict. that. to. zeroth-order the angular momentum: should. have a power-law dependeney on the total virialized. mass. with exponent 5/3."," Theoretical investigations predict that to zeroth-order the angular momentum should have a power-law dependency on the total virialized mass, with exponent $5/3$." +" They also make a prediction concerning thespin. a dimensionless quantity defined as: where J=]|. EZ|bai,|Ease are respectively he moduli of the angular momentum. ancl of the total (kinetic plus potential) energv."," They also make a prediction concerning the, a dimensionless quantity defined as: = where $J\equiv\mid\bmath{J}\mid$, $E\equiv\mid E_{kin} + E_{grav} \mid$ are respectively the moduli of the angular momentum, and of the total (kinetic plus potential) energy." + The PDE of A has been »edieted to have an approximately lognormal cistribution., The PDF of $\lambda$ has been predicted to have an approximately lognormal distribution. + This initial prediction. was subsequently found. to. be rmasically valiclh also when higher order cllects were taken into account(2?)., This initial prediction was subsequently found to be basically valid also when higher order effects were taken into account. +. More recently. the racial profile of A ias been derived. [rom modified Jeans’ equations(?7).. and ound to be in good agreement with results from The properties of the spin distribution of DM halocs iive recently been considered. taking advantage of the availability of high spatial- and mass-resolution simulations.," More recently, the radial profile of $\lambda$ has been derived from modified Jeans' equations, and found to be in good agreement with results from The properties of the spin distribution of DM haloes have recently been considered, taking advantage of the availability of high spatial- and mass-resolution simulations." + have analysed a large sample of haloes drawn from the Millennium. simulation. and found (among other things) that the spin. distribution is poorly described by a lognormal distribution: 103) which(ή has the shape of a lognormal except at very high ancl very low values of the spin.," have analysed a large sample of haloes drawn from the Millennium simulation, and found (among other things) that the spin distribution is poorly described by a lognormal distribution: ) = ) They alternatively suggest an empirical fit: ) ] which has the shape of a lognormal except at very high and very low values of the spin." + They also suggest that the actual shape of the distribution depends on the adopted ol Theoretical ancl numerical studies are aimed at understanding the origin of the (almost) lognormal spin distribution. using a limited set of statistical and dynamical assumptions or by performing numerical experiments.," They also suggest that the actual shape of the distribution depends on the adopted of Theoretical and numerical studies are aimed at understanding the origin of the (almost) lognormal spin distribution, using a limited set of statistical and dynamical assumptions or by performing numerical experiments." + have shown that the distribution of specilic angular momentum (which they celine as j=J/AM7) can be described by a rather complex PDE. which. can be approximated. by a lognormal in the central part. but deviates significantly from it at low- and high values of J.," have shown that the distribution of specific angular momentum (which they define as $j' = J/M^{5/3}$ ) can be described by a rather complex PDF, which can be approximated by a lognormal in the central part, but deviates significantly from it at low- and high values of $j'$." + They also make predictions lor the dependence of the peak of the spin distribution for clillerent values of mass. suggesting that it varies linearly with the average as of the distribution.," They also make predictions for the dependence of the peak of the spin distribution for different values of mass, suggesting that it varies linearly with the average $\sigma_{\delta}$ of the distribution." + find a correlation of spin A with mass. albeit a weak one.," find a correlation of spin $\lambda$ with mass, albeit a weak one." + Note that they use a set of simulations using boxes of varving sizes. while in the present work we adopt a single simulation. thus minimising the spurious tidal effects.," Note that they use a set of simulations using boxes of varying sizes, while in the present work we adopt a single simulation, thus minimising the spurious tidal effects." + Furthermore.(???).. studying the spin distributions at. low redshifts. have found that more massive haloes show larger values of A.," Furthermore, studying the spin distributions at low redshifts, have found that more massive haloes show larger values of $\lambda$." + Extensions of these results to higher recdshift jwe been provided by2.. who have studied the evolution of the spin distributions at z26. and showed that more massive haloes (AL~10 M.) tend to have a median A uigher than that of Al~107M. haloes.," Extensions of these results to higher redshift have been provided by, who have studied the evolution of the spin distributions at $z > 6$, and showed that more massive haloes $M\simeq +10^{7} {\rm M}_{\odot}$ ) tend to have a median $\lambda$ higher than that of $M\simeq +10^{6} {\rm M}_{\odot}$ haloes." + Also. high-A idoes tend to cluster more (by a factor 3-5) than low-spin valoes. a trend which strengthens with time.," Also, $\lambda$ haloes tend to cluster more (by a factor 3-5) than low-spin haloes, a trend which strengthens with time." + However. their simulation is restricted to a small box (Li=2.465.1 Mpc). hus making their result more prone to uncertainties [ron COSMIC The steady improvement of the available hardware ancl software. resources makes today possible simulations where the limits imposed by the finite spatial and. mass resolution limits are challenged.," However, their simulation is restricted to a small box $L_{b} = 2.46 h^{-1} {\rm Mpc}$ ), thus making their result more prone to uncertainties from cosmic The steady improvement of the available hardware and software resources makes today possible simulations where the limits imposed by the finite spatial and mass resolution limits are challenged." + The simulations we have performed. in this work were aimed at. providing a sample of reasonably well-resolvecl DAL haloes., The simulations we have performed in this work were aimed at providing a sample of reasonably well-resolved DM haloes. + We have obtained a catalogue of more than 77600 DAL haloes. each resolved by more than 20 particles. and we have used only those haloes with more than 300 particles. thus resulting in a catalogue containing more than 16400 haloes.," We have obtained a catalogue of more than 77600 DM haloes, each resolved by more than 20 particles, and we have used only those haloes with more than 300 particles, thus resulting in a catalogue containing more than 16400 haloes." + We have chosen a box size of L=70h5 Alpe. smaller than the one used in previous papers 2).. and a large number of particles. to maximize the mass In this work we present the results of a numerical," We have chosen a box size of $L = 70 \, h^{-1}$ Mpc, smaller than the one used in previous papers , and a large number of particles, to maximize the mass In this work we present the results of a numerical" +al.,al. + 2000). the ASC'A position deviates from the NED oue of about 137. which is within the error circle (21) of the GIS.," 2000), the ASCA position deviates from the NED one of about $''$, which is within the error circle $''$ ) of the GIS." + We also noted that the image extracted for the events collected during the flare phase has a position 09 59 0.1 27 16.5. which is oulv 117 away from the position of the source im the total observation for the CUS.," We also noted that the image extracted for the events collected during the flare phase has a position 09 59 40.4 -50 27 16.8, which is only $''$ away from the position of the source in the total observation for the GIS." +" This analysis sets that anv confusing source should be within 20"" of the PISS 0558-501.", This analysis sets that any confusing source should be within $''$ of the PKS 0558-504. + The analysis of the SIS data eives a simular conclusion. but detailed uunber is not elven here since our lain analysis below concentrates ou the CIS curve.," The analysis of the SIS data gives a similar conclusion, but detailed number is not given here since our main analysis below concentrates on the GIS curve." + The SIS curve is simular to the CUS one., The SIS curve is similar to the GIS one. + The flare is also clearly present iu the SIS data., The flare is also clearly present in the SIS data. + Due to different screcuing criteria applied to the SIS and CIS. the SIS data just before the peak of flare are not available.," Due to different screening criteria applied to the SIS and GIS, the SIS data just before the peak of flare are not available." + Thus. the flare profile is less well defined than in the GIS curve.," Thus, the flare profile is less well defined than in the GIS curve." + The flare is shown with increased time resolution (CIS 6Ls bius) in Fie., The flare is shown with increased time resolution (GIS 64-s bins) in Fig. + 2b., 2b. + Clearly. the source slows variability on short time scales of 1075. A \? test on the coustancy of count-rate for the second eroup of the data iu Fie 2b vives a \?/do.f=15.6/21. which is at a probability P.=0.002 by chance.," Clearly, the source shows variability on short time scales of $^2$ s. A $\chi^2$ test on the constancy of count-rate for the second group of the data in Fig 2b gives a $\chi^2/d.o.f$ =45.6/21, which is at a probability $P_r$ =0.002 by chance." + The pre-flare has a mean count rate of O.5340.02 cts/s for an interval of 650 seconds., The pre-flare has a mean count rate of $\pm$ 0.02 cts/s for an interval of 650 seconds. + Unfortunately. there is a 1000 see gap just before the on-set of the flare.," Unfortunately, there is a 1000 sec gap just before the on-set of the flare." + Nevertheless. the flare rose very fast.," Nevertheless, the flare rose very fast." + The couut rate increased from (.82+0.05 cts/s (average over three bins) to the peak count rate 1250.05. cts/s (three bius average) in less than 128 sec., The count rate increased from $\pm$ 0.05 cts/s (average over three bins) to the peak count rate $\pm$ 0.05 cts/s (three bins average) in less than 128 sec. + The statistical significance of this chauge is about 1.5 o., The statistical significance of this change is about 4.5 $\sigma$. + This gives a vate of change in the count-rate ACR/At~(2.5£0.6)10? ets s2.," This gives a rate of change in the count-rate $\Delta CR/\Delta t +\simeq (2.5\pm 0.6) 10^{-3}$ cts $^{-2}$." + A more subjective estimation of the rate by linear fit to the rising part of the curve (first 7 poiuts in the second eroup of data. see Fig 2b).," A more subjective estimation of the rate by linear fit to the rising part of the curve (first 7 points in the second group of data, see Fig 2b)." + This eives a rate ACR/At~(1.1140.28) ets 7. a factor of two lower than the one from direct visual inspection.," This gives a rate $\Delta CR/\Delta t \simeq +(1.14\pm 0.28)$ cts $^{-2}$, a factor of two lower than the one from direct visual inspection." + We will quote the latter more conserved uuuber in later discussion., We will quote the latter more conserved number in later discussion. + The source was then dime with possible flicker till the eud of this orbit., The source was then dimming with possible flicker till the end of this orbit. + The presence of the second peak probably is also real since it appears in the SIS curve as well., The presence of the second peak probably is also real since it appears in the SIS curve as well. + The flux returus to the pre-flare level at the beeiuniug of the next orbit observation., The flux returns to the pre-flare level at the beginning of the next orbit observation. + We define two bands based ou the spectral ranee of the soft N-rav excess (see 83.1)., We define two bands based on the spectral range of the soft X-ray excess (see 3.1). +" The energy ranges are 0.8-2.0, 2.0-10 keV for soft and hard bands for CIS data."," The energy ranges are 0.8-2.0, 2.0-10 keV for soft and hard bands for GIS data." + The soft ancl extends to 0.5 keV for the SIS data since there is a substantial fraction of couuts in the euergv range 0.5-0.8 seV. The light curves for these two bauds were extracted oe1 order to examine possible spectral variations caine the flare., The soft band extends to 0.5 keV for the SIS data since there is a substantial fraction of counts in the energy range 0.5-0.8 keV. The light curves for these two bands were extracted in order to examine possible spectral variations during the flare. + The flare is seen in both bauds., The flare is seen in both bands. + However. it shows arecr amplitude and decays faster in the hard baud thu in the soft one.," However, it shows larger amplitude and decays faster in the hard band than in the soft one." + The ratios of the count rate for the fare oeak to the pre-flare are 1.700.183 aud 2.1340.17 for the μαoft aud. hard. bands. respectively.," The ratios of the count rate for the flare peak to the pre-flare are $\pm$ 0.13 and $\pm$ 0.17 for the soft and hard bands, respectively." +" There is no significant ecrease in the count rate of the soft baud chiving the fare orbit (A2 /d.o.f 3.1/3. D.>0.3 for a coustant count rate oe ithe last 1021 oseconds during the flare orbit). in contrast to the clear rising aud fading in the curves of hard baud (\?fd.o. f=15.0/3. D,c0.002 for the same \? test) (Fie 3)."," There is no significant decrease in the count rate of the soft band during the flare orbit $\chi^2/d.o.f$ = 3.4/3, $P_r>0.3$ for a constant count rate in the last 1024 seconds during the flare orbit), in contrast to the clear rising and fading in the curves of hard band $\chi^2/d.o.f$ =15.0/3, $P_r\simeq 0.002$ for the same $\chi^2$ test) (Fig 3)." + We examine whether there is a correlation between the hardness ratio and the total count-rate and whether the flare follows the same correlation., We examine whether there is a correlation between the hardness ratio and the total count-rate and whether the flare follows the same correlation. +" The harduess ratio is the ratio of count rates in the 2-I0keV. baud aud iu the 0,8-2.0. keV band.", The hardness ratio is the ratio of count rates in the 2-10keV band and in the 0.8-2.0 keV band. + In order to achieve a reasonable S/N ratio for cach bin. we use a 1021 sec binning light curves.," In order to achieve a reasonable S/N ratio for each bin, we use a 1024 sec binning light curves." + The larduess ratio is weakly auticorelated with the total couut-rate for normal variations., The hardness ratio is weakly anti-correlated with the total count-rate for normal variations. + A Spearman rank correlation analysis gives a correlation cocticient of r—-0.371 (n=62) for which the probability for uull liypothesis is P.—0. for the CUS data (we have ignored these data points with uncertainty iu the ratio larger than 0.15)., A Spearman rank correlation analysis gives a correlation coefficient of r=-0.371 (n=62) for which the probability for null hypothesis is $P_r$ for the GIS data (we have ignored these data points with uncertainty in the ratio larger than 0.15). + In figure [| we add the hardness ratios for the flare phase. which is binned with 512 seconds binine.," In figure 4, we add the hardness ratios for the flare phase, which is binned with 512 seconds bining." + The flare is distinguished itself from others for having large harduess ratio and large count-rate (Fie., The flare is distinguished itself from others for having large hardness ratio and large count-rate (Fig. + 1)., 4). + Iu fact. the first three data points of the flare pliase have harducss ratio among he lareest.," In fact, the first three data points of the flare phase have hardness ratio among the largest." + And the last data point of the flare returns to he average harducss ratio., And the last data point of the flare returns to the average hardness ratio. + If tle spectrum curing the flare is a power-law. the power-law indices is about 1.9 for the frst three poiuts of the flare and 2.2 for the last point of the Hare (Fie.," If the spectrum during the flare is a power-law, the power-law indices is about 1.9 for the first three points of the flare and 2.2 for the last point of the flare (Fig." + 1)., 4). + It is not clear that if the decreasing harducss ratio indicates a time delay between the two bands., It is not clear that if the decreasing hardness ratio indicates a time delay between the two bands. + Cross correlation analysis does not found auv definite delays )otwoeen the soft N-ravs aud the hard N-ravs., Cross correlation analysis does not found any definite delays between the soft X-rays and the hard X-rays. + Since the first three data poiuts diving the flare imply a spectral iudex of 1.9 (see Fig., Since the first three data points during the flare imply a spectral index of 1.9 (see Fig. + E). which is significautlv wader than the menn spectral iudex in the 2-10 keV xuid. the change iu the harduess ratio cannot be explained with the change of relative coutributiou of the soft excess alone.," 4), which is significantly harder than the mean spectral index in the 2-10 keV band, the change in the hardness ratio cannot be explained with the change of relative contribution of the soft excess alone." + The spectruui in the 2-10 keV band must also have changed., The spectrum in the 2-10 keV band must also have changed. + However. if the hard X-ray spectimim is as flat as D—1.9. sugeested by the NMM observation (0Drien et al.," However, if the hard X-ray spectrum is as flat as $\Gamma=1.9$, suggested by the XMM observation (O'Brien et al." + 2000). then during the large dare. the N-vav cussion is domiuated by the hard power-law.," 2000), then during the large flare, the X-ray emission is dominated by the hard power-law." +" We observed a rapid flare in the PISS 0558-501. during which the N-rav count rate increased by a factor of nearly two iu 33 nuünutes. and possibly in as short as two τος,"," We observed a rapid flare in the PKS 0558-504, during which the X-ray count rate increased by a factor of nearly two in 33 minutes, and possibly in as short as two minutes." + This result independently confirms the existence of a rapid flare im this object obtained by Cüuga (Remillard et al..," This result independently confirms the existence of a rapid flare in this object obtained by Ginga (Remillard et al.," + 1992) and further sugecsts that the flare is repetitive., 1992) and further suggests that the flare is repetitive. + ROSAT IIRI observation also found that the object is hiehlv variable in the soft N-rav baud. but no such clearly vapid flare eveut was secu (Cliozz et al.," ROSAT HRI observation also found that the object is highly variable in the soft X-ray band, but no such clearly rapid flare event was seen (Gliozzi et al." + 2000)., 2000). + This might be due to the relatively fat spectrum duriug the flare phase (à much lower amplitude in the soft X-rav band) aud/or to the low photon collecting area of the ROSAT IIRI or it does not happen to catch a flare., This might be due to the relatively flat spectrum during the flare phase (a much lower amplitude in the soft X-ray band) and/or to the low photon collecting area of the ROSAT HRI or it does not happen to catch a flare. + Neither the ASCA observation in 1999 nor the NADAL observation in 2000 (Ciliozzi et al 2001) detected a rapid flare., Neither the ASCA observation in 1999 nor the XMM observation in 2000 (Gliozzi et al 2001) detected a rapid flare. + The total exposure time. by πας up those of Ciuga. ASCA and XMM. observatious. is about 150 ks. in which rapid flares were detected twice.," The total exposure time, by summing up those of Ginga, ASCA and XMM observations, is about 150 ks, in which rapid flares were detected twice." + This suggests that rapid flares occur uot infrequently in this object., This suggests that rapid flares occur not in-frequently in this object. + The fastest variation during the rie of flare phase sugeests a rate of change in the count rate ACR/AF=(1.1120.28)10.7 cts 2 in a linear fit.," The fastest variation during the rise of flare phase suggests a rate of change in the count rate $\Delta CR/\Delta t = (1.14\pm +0.28)~10^{-3}$ cts $^{-2}$ in a linear fit." +" Asstuning the N-rav spectzrun can be deseribed by a power-law. the flare has a spectrum with a photon iudex of ~ 1.9 from its harduess ratio. the rate of change in the count-rate vields a AL/At=(1.5+£0.10"" eres 7 in the 05-10 keV band (assuming fy = 75 and qj= 0.5) if the X- Cluits isotropically."," Assuming the X-ray spectrum can be described by a power-law, the flare has a spectrum with a photon index of $\simeq$ 1.9 from its hardness ratio, the rate of change in the count-rate yields a $\Delta L/ +\Delta t = (1.8\pm 0.4) 10^{42}$ ergs $^{-2}$ in the 0.8-10 keV band (assuming $H_0$ = 75 and $q_0 = 0.5$ ) if the X-ray emits isotropically." + Note this value is similar to the one obtained diving the Cduea observation., Note this value is similar to the one obtained during the Ginga observation. + Since the flare spectrum must be uot Linited iu the 0.5 to 10 keV band. the actual AL/Af should be lareer.," Since the flare spectrum must be not limited in the 0.8 to 10 keV band, the actual $\Delta L/ +\Delta t$ should be larger." + If the lard Nara spectrum extends to energv as biel as 100 keV. then the," If the hard X-ray spectrum extends to energy as high as 100 keV, then the" +to 20 clean lines in the 2008 HITRAN database (Rothmanetal.2009).,to 20 clean lines in the 2008 HITRAN database \citep{rothman09}. +. Strong. clear svimnmetric lines were favored and a calibration [actor was then applied to correct the waveniumber scale.," Strong, clear symmetric lines were favored and a calibration factor was then applied to correct the wavenumber scale." + A typical calibration factor (obtained by dividing the HITRAN line center by observation) for region | was 1.000001739 (Irom the 900 °C! line list) resulting in a shift of 0.00209 tat 1200 | ancl a tvpical calibration factor [or region 2 was 1.000000990 (rom the 900 °C! line list) resulting in a shift of 0.00158 + at 1600 1., A typical calibration factor (obtained by dividing the HITRAN line center by observation) for region 1 was 1.000001739 (from the 900 $^{\circ}$ C line list) resulting in a shift of 0.00209 $^{-1}$ at 1200 $^{-1}$ and a typical calibration factor for region 2 was 1.000000990 (from the 900 $^{\circ}$ C line list) resulting in a shift of 0.00158 $^{-1}$ at 1600 $^{-1}$. + The overall accuracy of our wavenumber scale is better than £0.002 em| after calibration., The overall accuracy of our wavenumber scale is better than $\pm$ 0.002 $^{-1}$ after calibration. + Qur analysis method is not unique and has successfully been applied by many workers. most recently in the sub-nillimeter regime to study. astrophysical ‘weeds’ 2010).," Our analysis method is not unique and has successfully been applied by many workers, most recently in the sub-millimeter regime to study astrophysical `weeds' \citep{fortman10}." +". It involves a comparison of the observed intensities to a reliable data set (in our case, the 2008 LITRAN database)."," It involves a comparison of the observed intensities to a reliable data set (in our case, the 2008 HITRAN database)." + In order to make our emission line lists comparable to IIETRAN we converted them into absorption intensities (Nassar&Bernath 2003).., In order to make our emission line lists comparable to HITRAN we converted them into absorption intensities \citep{nassar03}. . +" The relationship between emission and absorption intensities 1s given bv where 5,,,::,, IS the intensity of the emitted line. v is the Ilrequency and L is the (emperature."," The relationship between emission and absorption intensities is given by where $S_{emission}$ is the intensity of the emitted line, $\nu$ is the frequency and $T$ is the temperature." + The HITRAN database provides line intensities at 296 IX so they were converted to the relevant temperatures of the fitted line lists., The HITRAN database provides line intensities at 296 K so they were converted to the relevant temperatures of the fitted line lists. +" The intensity of a line (5) is delined in SI units as where 5,» is thesquare of the transition dipole moment. (Q is (he total internal"," The intensity of a line $S'$ ) is defined \citep{bernath05} in SI units as where $S_{J'J''}$ is thesquare of the transition dipole moment, $Q$ is the total internal" +where ¢ and + are the internal rotation velocity and characteristic size of the disk. respectively. V. and R are the orbit velocity and separation at pericenter. and ο the eccentricity of the orbit (Binney&Tremaine1987).,"where $v$ and $r$ are the internal rotation velocity and characteristic size of the disk, respectively, $V$ and $R$ are the orbit velocity and separation at pericenter, and $e$ the eccentricity of the orbit \citep{BT87}." +. The work of Barnes(1988.1992) and Hernquist(1992.1993) generalized TT72's modeling by treating the internal gravity of colliding galaxies self-consistently.," The work of \citet{B88,B92} and \citet{H92,H93} + generalized TT72's modeling by treating the internal gravity of colliding galaxies self-consistently." + These and related studies verified TT72's claim that tail-making is a kinematic process and examined the sensitivity of tail-making efficiency to the distribution of stars in interacting and merging galaxies.," These and related studies verified TT72's claim that tail-making is a kinematic process \citep[see +Figure 3 of][]{DMH99} and examined the sensitivity of tail-making efficiency to the distribution of stars in interacting and merging galaxies." + Subsequently. Dubinskietal.(1996) and Μοςetal.(1998) demonstrated empirically that the lengths and kinematics oftidal tails also depend critically on the distribution of dark matter surrounding each galaxy.," Subsequently, \citet{DMH96} and \citet{MDH98} demonstrated empirically that the lengths and kinematics of tidal tails also depend critically on the distribution of dark matter surrounding each galaxy." + For the same radial stellar profiles. longer (shorter) and more (less) prominent tidal tails form in shallower (deeper) dark matter potential wells.," For the same radial stellar profiles, longer (shorter) and more (less) prominent tidal tails form in shallower (deeper) dark matter potential wells." + In particular. and Springel&White(1999) showed that the depth of the dark matter potential is as important in determining the lengths of tidal tails as the orbit (seee.g.Figure4inDubinskietal.1999).," In particular, \citet{DMH99} and \citet{SW99} showed that the depth of the dark matter potential is as important in determining the lengths of tidal tails as the orbit \citep[see e.g. Figure 4 in][]{DMH99}." + Similar conclusions have been reached in studies of local dwarf spheroidal galaxies (seee.g.Mayeretal.2002:Read2006).," Similar conclusions have been reached in studies of local dwarf spheroidal galaxies \citep [see e.g.] +[]{M02, Read06}." +. Most recently. it has been shown that resonant stripping of stars in disks can alter the mass to light ratios of dwarf galaxies when they encounter more massive systems by removing luminous material more efficiently than dark matter.," Most recently, it has been shown that resonant stripping of stars in disks can alter the mass to light ratios of dwarf galaxies when they encounter more massive systems by removing luminous material more efficiently than dark matter." + The minimal response of the dark matter is expected if its particles move on random orbits. in which case the net perturbation on the halo mostly averages out (D'Onghiaetal.2009)..," The minimal response of the dark matter is expected if its particles move on random orbits, in which case the net perturbation on the halo mostly averages out \citep{Don09b}." + A similar resonant phenomenon has been suggested to cause the LMC disk to thicken by interactions with the Milky Way (Weinberg2000) or more generally in the context of heating and disruption of satellites (Choietal. 2009).., A similar resonant phenomenon has been suggested to cause the LMC disk to thicken by interactions with the Milky Way \citep{W00} or more generally in the context of heating and disruption of satellites \citep{CWK09}. . + In addition. the origin of the Magellanic Stream might have a tidal origin (Beslaetal.2010) through interactions between dwarfs in groups (D'Onghia&Lake2008).. provided that the Magellanic Clouds are on their first pericentric passage 2007).," In addition, the origin of the Magellanic Stream might have a tidal origin \citep{B10} through interactions between dwarfs in groups \citep{DL08}, provided that the Magellanic Clouds are on their first pericentric passage \citep []{K06, B07}." +. While these works have elucidated the physies of tail-making and proven that observed peculiar galaxies are indeed a consequence of collisions and mergers. they have left open à number of important issues.," While these works have elucidated the physics of tail-making and proven that observed peculiar galaxies are indeed a consequence of collisions and mergers, they have left open a number of important issues." + Most previous studies of tidal interactions between disk galaxies have been empirical in nature., Most previous studies of tidal interactions between disk galaxies have been empirical in nature. + Thus. they have identified broad conditions for tidal features to be produced. but have not provided simple criteria for isolating the efficiency of this process on specific parameters. of an encounter.," Thus, they have identified broad conditions for tidal features to be produced, but have not provided simple criteria for isolating the efficiency of this process on specific parameters of an encounter." + Therefore. the circumstances under which very long tails could result. like those in the Superantennae. remain uncertain.," Therefore, the circumstances under which very long tails could result, like those in the Superantennae, remain uncertain." + Moreover. even nearly 40 years after TT72. there is still considerable confusion in the community regarding galactic bridges and tails. especially in regards to the role played by resonances in their origin.," Moreover, even nearly 40 years after TT72, there is still considerable confusion in the community regarding galactic bridges and tails, especially in regards to the role played by resonances in their origin." + Finally. attempts to reproduce the morphology and kinematies of individual observed systems is greatly complicated bythe sensitivity of tail-making to the," Finally, attempts to reproduce the morphology and kinematics of individual observed systems is greatly complicated bythe sensitivity of tail-making to the" +the differences fi‘om the lower resolution run to be small.,the differences from the lower resolution run to be small. + In the present study we aim at unclerstaucing the basic role of gravity. neglecting the role of orbital motio1 and the acceleration zone of tle primary wind.," In the present study we aim at understanding the basic role of gravity, neglecting the role of orbital motion and the acceleration zone of the primary wind." + Iucludiug these two effects will increase the acc‘eliou rate and makes it start ealier even (see section 2.2))., Including these two effects will increase the accretion rate and makes it start earlier even (see section \ref{sec:assumptions}) ). + We set the distance between the sta ‘sto be constant at each numerical run. aud let the flow reach a. more or less. steady state.," We set the distance between the stars to be constant at each numerical run, and let the flow reach a, more or less, steady state." + It is not a strict steady state. as the winds interaction zone is wieoling. as was already found by Pittare et al. (," It is not a strict steady state, as the winds interaction zone is wiggling, as was already found by Pittard et al. (" +1998). Pittard Corcora1 (2002). Okazaki et al. (,"1998), Pittard Corcoran (2002), Okazaki et al. (" +2008). axl Parkin et al. (,"2008), and Parkin et al. (" +2009).,2009). + We coudict the runs with and without tie secoudary stellar gravity. and lor several orbital separatioins.," We conduct the runs with and without the secondary stellar gravity, and for several orbital separations." + At the location of each star we inject its appropriate wind., At the location of each star we inject its appropriate wind. + The winds were injected [roi a square of size 8x8 cells around eacl star., The winds were injected from a square of size $8\times 8$ cells around each star. + Namely. 1 cells [rom the star aloug tlie axes.," Namely, 4 cells from the star along the axes." + This prevents the accreting primary wiud to reac1 a distance closer than 5 cells from the center o‘the secondary (as at each time step the coucitiou of outflow is imposecl there)., This prevents the accreting primary wind to reach a distance closer than 5 cells from the center of the secondary (as at each time step the condition of outflow is imposed there). + As in most runs tlere are 56 cells between the stars. this closest cista eis ~Q.1r.," As in most runs there are 56 cells between the stars, this closest distance is $\sim 0.1 r$." + As the condition of accretion we tase tlie preseuce oL dense primary wiud gas at a distance of 5 cells (0.17) from the center of the secoπο star., As the condition of accretion we take the presence of dense primary wind gas at a distance of 5 cells $0.1 r$ ) from the center of the secondary star. + The boundary couditious of tle simulation box are outflow at the 6 sides of tle box: the gas that flows out of the box cannot flow back into the box., The boundary conditions of the simulation box are outflow at the $6$ sides of the box; the gas that flows out of the box cannot flow back into the box. + Our calculations have the aim of revealing the role of gravity in the winds collision process when y Car approaches its periastron: near apastrou the role of gravity is small., Our calculations have the aim of revealing the role of gravity in the winds collision process when $\eta$ Car approaches its periastron; near apastron the role of gravity is small. + For that we varied ouly the distauce between the stars (orbital separation). aud let the flow reach a steady state (up to a wigeline motion) with aud without gravity.," For that we varied only the distance between the stars (orbital separation), and let the flow reach a steady state (up to a wiggling motion) with and without gravity." + We now discuss the implications of the main assuiptions and approximations. (, We now discuss the implications of the main assumptions and approximations. ( +1) While we inject the primary wind at its terminal veocity of e4=500kinsJ|. 4he winds from OB super-giants are accelerated over a large distauce o “about several stellar radii.,"1) While we inject the primary wind at its terminal velocity of $v_1=500 \km \s^{-1}$, the winds from OB super-giants are accelerated over a large distance of about several stellar radii." +" À commonly fitting formula at several stellar radii is e2wil[—(R&Rπα. with 8~1-335 (e.δι, Ixraus et al."," A commonly fitting formula at several stellar radii is $v \simeq v_1[1-(R_1/r)]^{\beta}$, with $\beta \simeq 1-3.5$ (e.g., Kraus et al." + 2007. and. refereices herein).," 2007, and references therein)." + Lf. Or example. we tase Ry=0.8AU aud 32. we find that the wiud velocity is ve(2AU)=0.360 (just before »eriastron) aud oflAU)=0.6143.," If, for example, we take $R_1=0.8 \AU$ and $\beta=2$, we find that the wind velocity is $v(2 {\AU})=0.36 v_1$ (just before periastron) and $v(4 {\AU})=0.64 v_1$." + The density is vy/e times as high coupared witl the value according to our assumption of a constant wind speed., The density is $v_1/v$ times as high compared with the value according to our assumption of a constant wind speed. + When the orbital motion is «'onsidered. it is fou that the stagnation point of the collici winds moves toward the secoucary. ald by that streugtheniug the accretion process.," When the orbital motion is considered, it is found that the stagnation point of the colliding winds moves toward the secondary, and by that strengthening the accretion process." + Also. when acceleration zone is ¢‘OLSidered. the accretion rate Is Inereased. by a substantial [raclion (Ixashi Soker 2009c).," Also, when the acceleration zone is considered, the accretion rate is increased by a substantial fraction (Kashi Soker 2009c)." + Neglecing the acceleralon zone is οιr strongest assumption., Neglecting the acceleration zone is our strongest assumption. + lucludiug it will mak our conclusious that accretion is inevitable stronger even., Including it will make our conclusions that accretion is inevitable stronger even. +of about. 1.2 by the reaction (10) at 3. x 104 vr.,of about 1.2 by the reaction (10) at 3 $\times$ $^4$ yr. + Figure 8 shows the density dependence of the CECS/PCCS ratio., Figure 8 shows the density dependence of the $^{13}$ $^{13}$ CCS ratio. + Although the peak isotopomer ratio increases wilh density. ihe maximum ratio is 1.3. which is much lower than observed.," Although the peak isotopomer ratio increases with density, the maximum ratio is 1.8, which is much lower than observed." + Unlike CCII. the exchange reaction (10) is inellicient. because (he sulfur atom is less abundant than the hydrogen atom by three orders of magnitude.," Unlike CCH, the exchange reaction (10) is inefficient, because the sulfur atom is less abundant than the hydrogen atom by three orders of magnitude." + The gas density in TMC-I is 107 - 10° 7? (Pratap et al., The gas density in TMC-1 is $^4$ - $^5$ $^{-3}$ (Pratap et al. + 1997)., 1997). +" In the model with my, — 5 x 10! *. the abundances of carbon-chain molecules. such as HC4N. reach their maxima at 2 x 10! vr."," In the model with $n_{\rm {H_2}}$ = 5 $\times$ $^4$ $^{-3}$, the abundances of carbon-chain molecules, such as $_3$ N, reach their maxima at 2 $\times$ $^4$ yr." + Since carbon-chain molecules are abundant in TAIC-1. we compare the obtained molecular abundaces anc isotope and isolopomer ratios in the fiducial model al 2 x 107 vr to the observed values in TMC-I.," Since carbon-chain molecules are abundant in TMC-1, we compare the obtained molecular abundaces and isotope and isotopomer ratios in the fiducial model at 2 $\times$ $^4$ yr to the observed values in TMC-1." + Pratap et al. (, Pratap et al. ( +1997) made maps along the TAIC-1 ridge. and determine the column densities of CO and CCII.,"1997) made maps along the TMC-1 ridge, and determine the column densities of $^{18}$ O and CCH." + The molecular abundance of CCLL in TMC-I is 5.7 x ! - 4.0 x 7. assuming a Ο/Η» ratio of 1.7 x *i (Frerking et al.," The molecular abundance of CCH in TMC-1 is 5.7 $\times$ $^{-10}$ - 4.0 $\times$ $^{-9}$, assuming a $^{18}$ $_2$ ratio of 1.7 $\times$ $^{-7}$ (Frerking et al." + 1982)., 1982). + In our model. we obtain the abundance of CCII to be 3.0 x )°. which is in reasonable agreement with (he observed. value.," In our model, we obtain the abundance of CCH to be 3.0 $\times$ $^{-9}$, which is in reasonable agreement with the observed value." + In the fiducial model. we obtain the CCIL/PCCIL CCIL/CECIH and CPCIT/PCCIUH ralios (ο be 99. 63 and 1.6. respectively. at 2 x LO! vi.," In the fiducial model, we obtain the $^{13}$ CCH, $^{13}$ CH and $^{13}$ $^{13}$ CCH ratios to be 99, 63 and 1.6, respectively, at 2 $\times$ $^4$ yr." + The observed. CCIL/CECIT and CCHIL/PCCIE ratios are higher than 170 and 250. respectively (Sakai οἱ al.," The observed $^{13}$ CH and $^{13}$ CCH ratios are higher than 170 and 250, respectively (Sakai et al." + 2010)., 2010). + The average Isotope ratio. which is caleulated by," The average isotope ratio, which is calculated by" +" ).. aud = X11) with m=| aud 5,=DBjD'"". D""=ch mee, ","+ ^2 ), and = ) with $\mnt = 1 + 2 \szn \kn \gmn$ and $\gamma_n = B / B_{c}^{n}$, $B_{c}^{n} = e \hbar / m_{n}^{2} c^{3}$ ." +Iu Eqs. (, In Eqs. ( +"5)-(7).24756,5, à=e?fhe is the fine structure constant. s! ids the neutron spi projection in the magnetic field direction. aud νι6), are the anomalous magnetic moments for protons aud neutrons respectively. as eiven below iu Eq. (","5)-(7), $\alpha = e^2 / \hbar c$ is the fine structure constant, $s_{z}^{n}$ is the neutron spin projection in the magnetic field direction, and $\kp, \, \kn$ are the anomalous magnetic moments for protons and neutrons respectively, as given below in Eq. (" +10).,10). + Tere we use the same notation as in Suh Mathews (2001a) for the particle Fermi energv aud magnetic field streuet[u, Here we use the same notation as in Suh Mathews (2001a) for the particle Fermi energy and magnetic field strength. +m For a system of strongly interacting barvous (neutrons and protous). the relativistic mean field (Hartree) theory should be a reasonable approximation for the description of the equation of state for maeuetarmatter at lieh density (Brodericketal.2000:Chakrabarty1997) through the exchange of σ and vector z.p imesous in a strong magnetic field.," For a system of strongly interacting baryons (neutrons and protons), the relativistic mean field (Hartree) theory should be a reasonable approximation for the description of the equation of state for magnetar-matter at high density \citep{broderick,chakra} through the exchange of $\sigma$ and vector $\omega, \rho$ mesons in a strong magnetic field." +" In the barvon Lagrangian for the relativistic ILurtree theory. the anomalous magnetic moments are included through the coupling of the barvons to the clectromaguetic field tensor with 5tyete] and the streneths &, aud 5, given by Lops A(12) where g,=5.58 aud g,=—X82 are the Laude g-factors for protons and neutrons. respectively."," In the baryon Lagrangian for the relativistic Hartree theory, the anomalous magnetic moments are included through the coupling of the baryons to the electromagnetic field tensor with $\sigma_{\mu \nu} = \frac{i}{2} \left[\gamma_{\mu},\gamma_{\nu} \right]$ and the strengths $\kp$ and $\kn$ given by _p = - 1 ), _n =, where $g_p = 5.58$ and $g_n = -3.82$ are the Lande $g$ -factors for protons and neutrons, respectively." + Iu this work. we can ignore the possible scalar 6. the vector w and the iso-vector p meson seltiuteractious.," In this work, we can ignore the possible scalar $\sigma$, the vector $\omega$ and the iso-vector $\rho$ meson self-interactions." + Therefore. although the electromagnetic field is included iu the total Lagrangian. it assumed to be externally ecucrated (and thushas no associated feld equation) aud ouly frozeu-field configurations will be considered.," Therefore, although the electromagnetic field is included in the total Lagrangian, it assumed to be externally generated (and thushas no associated field equation) and only frozen-field configurations will be considered." +" The effective barvon mass ngνι is then given by the coupling to the σ ΠΟΣΟΙ. = 1 ος loa) (03) where gs aud i, are ma)the σ meson coupling coustaut and mass respectively,"," The effective baryon mass $m_{b=p,n}$ is then given by the coupling to the $\sigma$ meson, = m_b - + ) , where $g_{\sigma}$ and $m_{\sigma}$ are the $\sigma$ meson coupling constant and mass respectively." + Iu Eq. (, In Eq. ( +"11). ny is the scalar uuuber density for protons. Sol. { }3 z——h|t ).(11) whereTove nm_=\;ηl|σος255,SHUT,EA aud. poc—Ga,fmP yes. whileB the scalar number density. for. neutrons PU,15So = ( )H - {(UR ) |. with. =.1d|ostiat, and 55—(inefinPa,5 ","11), $n_{p}^{S}$ is the scalar number density for protons, = ( )^3 ), where $\mbps = \sqrt{1 + 2 \nfp \gmp{^*}} - \szp \kp \gmp^{*}$ and $\gmp^{*} = (m_e / \mps)^2 \gme$ , while the scalar number density for neutrons is = ( )^3 - ( ) ], with $\mbns = 1 + \szn \kn \gmn{^*}$ and $\gmn^{*} = (m_e / \mns)^2 \gme$." +"For simplicity. the nucleon rest dass is faken as my,—any ta the munerical calculation (Chakrabartyotal."," For simplicity, the nucleon rest mass is taken as $m = m_n = m_p$ in the numerical calculation \citep{chakra}." +"1997).. Assunniug a mixture of neutrons. protons. aud clectrous in chemical equilibrium. the chemical potentials are related by May, EL while the condition of charge neutrality gives uj(e,..55)) = nete ted)A"," Assuming a mixture of neutrons, protons, and electrons in chemical equilibrium, the chemical potentials are related by _n = _p + _e, while the condition of charge neutrality gives n_p, ) = n_e, )." +T) Given the uncleou-aneson coupling coustaut aud the cocficicuts in the scalar sel£iuteractious. the field equations can be solved selt-cousisteutlv for the chemical potentials. py(j-—n.p.c) and the inesou field streneths in a uniform imaeuetic Ποια B along the : axis corresponding to the choice of the gauge for the vector poteutial AM” (Brodericketal.2000)..," Given the nucleon-meson coupling constant and the coefficients in the scalar self-interactions, the field equations can be solved self-consistently for the chemical potentials, $\mu_{j} \; (j=n,p,e)$ , and the meson field strengths in a uniform magnetic field $B$ along the $z$ axis corresponding to the choice of the gauge for the vector potential $A^{\mu}$ \citep{broderick}." +" In this work. we adopt the following coupling coustauts and mesous asses: q2(nnsfiio)?=357.17, getingnu)=2737S. and wuninpy?=97.0 (Horowitz&SerotI981)."," In this work, we adopt the following coupling constants and mesons masses: $g_{\sigma}^{2} (m_{N}/m{\sigma})^2 = 357.47$, $g_{\omega}^{2} (m_{N}/m{\omega})^2 = 273.78$, and $g_{\rho}^{2} (m_{N}/m{\rho})^2 = 97.0$ \citep{horowitz}." +". Figure d shows the effective barvon mass fi as ἃ πιοΠο of baryon deusitv p for maguetic field strengths of +,=0.01 (solid line) and 10° (dash line) calculated in he model of Horowitz aud Serot (1981).", Figure 1 shows the effective baryon mass $m^{*} / m$ as a function of baryon density $\rho$ for magnetic field strengths of $\gme = 0.01$ (solid line) and $10^5$ (dash line) calculated in the model of Horowitz and Serot (1981). + For a magnetic field stroieth less than ~1015 C. This figure shows that he effective nucleon mass is uot significauth affected * naegnetie feld streugth.," For a magnetic field strength less than $\sim 10^{18}$ G, This figure shows that the effective nucleon mass is not significantly affected by magnetic field strength." + Broderick et al. (, Broderick et al. ( +2000) ancl Chakrabarty et al. (,2000) and Chakrabarty et al. ( +1997) have obtained simular results.,1997) have obtained similar results. + This effective baryon mass modifies the barvou dispersion relation in dense maenetar-matter., This effective baryon mass modifies the baryon dispersion relation in dense magnetar-matter. +" For matter iun thermodvuamical equilibria below the neutron drip density. pai,99L3«&101 ο), we adopt the maeguetic Bavin. Pethick. and Sutherland (BPS) model (Lai&Shapiro1991). aud use the seui-clupirical mass formula (Shapiro&Teukolsky1983)."," For matter in thermodynamical equilibrium below the neutron drip density, $\rho_{drip} \approx 4.3 \times 10^{11}$ $^3$, we adopt the magnetic Baym, Pethick, and Sutherland (BPS) model \citep{lai} + and use the semi-empirical mass formula \citep{shapiro}." +. For simplicity. we only consider 32 Fe unclei in the umucerical calculation.," For simplicity, we only consider $^{56} _{26}$ Fe nuclei in the numerical calculation." +" Then. the magnetization aud the differeutial susceptibility forthe magnetic BPS equation of state are eiven bv — Hy where M,. is themagnetization of the electron gas and gj. is given in Eq. ("," Then, the magnetization and the differential susceptibility forthe magnetic BPS equation of state are given by = _e + _L , = _e + where ${\cal M}_e$ is themagnetization of the electron gas and $\eta_e$ is given in Eq. (" +5).,5). + In Eq. (, In Eq. ( +16). My is the magnetization for the beeCoulomb lattice energy.,"16), ${\cal M}_L$ is the magnetization for the $bcc$Coulomb lattice energy." +" Then we drive here the lattice cdiffercutial susceptibility ij, to be. L lli"," Then we drive here the lattice differential susceptibility $\eta_L$ to be, _L = - 1.444" +"where M and ce, are (he mass accretion rate and radial velocity. respectively.","where $\dot{M}$ and $v_{r}$ are the mass accretion rate and radial velocity, respectively." +" Given the post-shock values for density. radial velocity. entropy and electron [raction (Ppost- Cpost» Spost- Yepont) the shock radius is determined so that the density obtained at the inner boundary should agree with the fixed value. pj,=1011g/em*."," Given the post-shock values for density, radial velocity, entropy and electron fraction $\rho_{\text{post}}$ , $v_{\text{post}}$, $s_{\text{post}}$, ${Y_{\text{e}}}_{\text{post}}$ ), the shock radius is determined so that the density obtained at the inner boundary should agree with the fixed value, $\rho_{\rm in} = 10^{11}~\mbox{g/cm}^{3}$." + The post-shock values are calculated from the corresponding pre-shock values by the Tankine-Hugoniot relations., The post-shock values are calculated from the corresponding pre-shock values by the Rankine-Hugoniot relations. + The pre-shock values. on Cie other hand. are obtained from the nass accretion rate. the outer boundary. conditions for entropy and electron fraction (sou. )cou ) 20 the assumption that matter falls freely with no interaction with neutrinos outside (he shock: where Ais (he shock radius.," The pre-shock values, on the other hand, are obtained from the mass accretion rate, the outer boundary conditions for entropy and electron fraction $s_{\text{out}}$, ${Y_{\text{e}}}_{\text{out}}$ ) and the assumption that matter falls freely with no interaction with neutrinos outside the shock; where $R_{\rm s}$ is the shock radius." + The entropy aud the electron fraction at the outer boundary are assiuned to be s44=3Ap/barvon and οι=0.5. respectively.," The entropy and the electron fraction at the outer boundary are assumed to be $s_{\rm out} = 3 ~k_{\rm B}\mbox{/baryon}$ and ${Y_{\rm e}}_{\rm out} = 0.5$, respectively." + In the munerical simuladions. axisvanmetry but no equatorial svmmelry is assumed.," In the numerical simulations, axisymmetry but no equatorial symmetry is assumed." + At the outer boundary. we adopt the fixed boundary condition consistent with the unperturbed state.," At the outer boundary, we adopt the fixed boundary condition consistent with the unperturbed state." + On the other hand. the free-flow-in boundary condition is used at the inner boundary.," On the other hand, the free-flow-in boundary condition is used at the inner boundary." + A realistic tabulated EOS by is used both for the simulations and for the preparation of the initial conditions., A realistic tabulated EOS by is used both for the simulations and for the preparation of the initial conditions. +" For all the models investigated in (his paper. tlie mass accretion rate and the mass of central object are fixedto be M-1A./s and Mi,=1.4AL.. respectively."," For all the models investigated in this paper, the mass accretion rate and the mass of central object are fixedto be $\dot{M} = 1~M_{\odot}\mbox{/s}$ and $M_{\rm in} = 1.4~M_{\odot}$, respectively." +" The neutrino temperatures are also fixed (o Z7,=4 MeV and Zz.=5 MeV. the twpical values in the phase."," The neutrino temperatures are also fixed to $T_{\nu_{\text{e}}} = 4$ MeV and $T_{\bar{\nu}_{\text{e}}} = 5$ MeV, the typical values in the post-bounce phase." +" On the other hand. we svstematically vary the neutrino luminosity within the range of L,=3.027.0-107?erg/s for different models."," On the other hand, we systematically vary the neutrino luminosity within the range of $L_{\nu} = \mbox{3.0--7.0}\cdot 10^{52}~\mbox{erg/s}$ for different models." + Note that these values are constant in time for each model., Note that these values are constant in time for each model. + This is necessary. (o realize the steady unperturbed states., This is necessary to realize the steady unperturbed states. +" Tostudy the instability. we add angular-dependent perturbations. d7,(7.4). to the racial velocity initially:"," Tostudy the instability, we add angular-dependent perturbations, $\delta v_{r}(r, \theta)$ , to the radial velocity initially;" +quantity seems to represent a lower limit to the binary fraction in OC's.,quantity seems to represent a lower limit to the binary fraction in OCs. + The existing estimates of the binary fraction in other nearby OC's from radial velocity surveys agree with this lower limit. (Mermilliod. Grenon Mayor 2008: Alermilliod. Queloz Mayor 2008: Mermilliod et al.," The existing estimates of the binary fraction in other nearby OCs from radial velocity surveys agree with this lower limit (Mermilliod, Grenon Mayor 2008; Mermilliod, Queloz Mayor 2008; Mermilliod et al." + 2008)., 2008). + The fraction of binaries obtained in this paper range from toTO%., The fraction of binaries obtained in this paper range from to. +. Xecording to the theoretical. N-body simulations of Portegies-Zwart et al. (, According to the theoretical N-body simulations of Portegies-Zwart et al. ( +2004) the binary fraction in a stellar svstem with a mass of AJ=1600Al. and a cistance to the Galactic center of d—12.1Ape (i.c. a typical OC) would. remain actually constant during the entire cluster evolution. rising in the cluster core from an initial up to in 1 Gyr of evolution as a result of mass segregation.,"2004) the binary fraction in a stellar system with a mass of $M=1600~M_{\odot}$ and a distance to the Galactic center of $d=12.1~Kpc$ (i.e. a typical OC) would remain actually constant during the entire cluster evolution, rising in the cluster core from an initial up to in 1 Gyr of evolution as a result of mass segregation." + Phe fraction measured here are marginallv smaller than the final result of these simulations., The fraction measured here are marginally smaller than the final result of these simulations. + Following these considerations. the initial binary fraction in our target clusters could be ~20504.. comparable to that observed in the solar neighborhood (Abt Levy 1976: Duquennoy Mawvor 1991: Reid Cizis 1997).," Following these considerations, the initial binary fraction in our target clusters could be $\sim30-50$, comparable to that observed in the solar neighborhood (Abt Levy 1976; Duquennoy Mayor 1991; Reid Gizis 1997)." + The estimated: binary fractions of the clusters of our sample and those of the sample of GC's presented in Sollima et al. (, The estimated binary fractions of the clusters of our sample and those of the sample of GCs presented in Sollima et al. ( +2007) reveal a significant correlation with the absolute visual magnitude and cannot exclude a dependence on the cluster age.,2007) reveal a significant correlation with the absolute visual magnitude and cannot exclude a dependence on the cluster age. + The possible correlation with age found in Sollima ct al. (, The possible correlation with age found in Sollima et al. ( +2007) is indeed not clearly visible among the sample of OC's.,2007) is indeed not clearly visible among the sample of OCs. + Phe OC's analyvsecl here hold binary fractions which are on average Larger than those measured in GC's., The OCs analysed here hold binary fractions which are on average larger than those measured in GCs. + However. the small number of OC's and the large errors in the binary fraction estimates make cdillicult to assess the existence (or absence) of a smooth correlation between these two parameters.," However, the small number of OCs and the large errors in the binary fraction estimates make difficult to assess the existence (or absence) of a smooth correlation between these two parameters." + A homogeneous analysis of a larger sample of clusters spanning a wide range of ages is therefore required to clarity this issue., A homogeneous analysis of a larger sample of clusters spanning a wide range of ages is therefore required to clarify this issue. + On the other hand. the correlation between binary fraction. ane luminosity. already reported bv Milone et al. (," On the other hand, the correlation between binary fraction and luminosity, already reported by Milone et al. (" +2008) in a sample of 50 GCs. has been confirmed also in the range of mass covered by our sample of OCs 10004. ).,"2008) in a sample of 50 GCs, has been confirmed also in the range of mass covered by our sample of OCs $M_{\odot}$ )." + Vheoretical models of evolution of stellar systems with a population of primordial binaries (Sollima 2008) indicates that 1e elliceney. of the binary ionization process. which represents together with mass segregation the dominant process in determining the binary fraction. has the same dependence as the mass on the cluster density and. velocity dispersion.," Theoretical models of evolution of stellar systems with a population of primordial binaries (Sollima 2008) indicates that the efficency of the binary ionization process, which represents together with mass segregation the dominant process in determining the binary fraction, has the same dependence as the mass on the cluster density and velocity dispersion." + An anticorrelation between these two quantities is therefore expected., An anticorrelation between these two quantities is therefore expected. + La practice. in low-mass clusters. a larger fraction of binary systems can survive to the process of binary clestruction as a consequence of the lower rate of collisions and the smaller mean kinetic energy. οἱ collidingstars.," In practice, in low-mass clusters, a larger fraction of binary systems can survive to the process of binary destruction as a consequence of the lower rate of collisions and the smaller mean kinetic energy of colliding." +. The results presented here conlirm the prediction of theoretical mocels indicating that the process of binary ionization is the dominant process which determines the fraction of binaries in any relaxed stellar svstem., The results presented here confirm the prediction of theoretical models indicating that the process of binary ionization is the dominant process which determines the fraction of binaries in any relaxed stellar system. +example. wilh Dy=8 kpc. a neutron star would have to have e=435 km/s. and would have to be located within roughly a pe of the source star.,"example, with $D_S=8$ kpc, a neutron star would have to have $v=435$ km/s, and would have to be located within roughly a pc of the source star." + This is so unlikely. compared to the probability of a stellaz-speed. planet-mass object within a few hundred pe. that it can be effectively ruled out.," This is so unlikely, compared to the probability of a stellar-speed planet-mass object within a few hundred pc, that it can be effectively ruled out." + In fact. if free-floating planets are common. we will have dozens to hundreds of similar cases. and (he hypothesis (hat the majority of the events (μον generate were actually generated by high-velocity stellar remnants can be eliminated.," In fact, if free-floating planets are common, we will have dozens to hundreds of similar cases, and the hypothesis that the majority of the events they generate were actually generated by high-velocity stellar remnants can be eliminated." + Note that the arguments above apply to the case in which the Einstein angle has been measured and is small., Note that the arguments above apply to the case in which the Einstein angle has been measured and is small. + If (he light curve shape provides only an upper limit on the value of θες.L then it is still possible that the lens was a high-velocity5S object.," If the light curve shape provides only an upper limit on the value of $\theta_E,$ then it is still possible that the lens was a high-velocity object." + High-resolution5S images9 that check for a shift in the sources light centroid can measure 85. if the lens is indeed a nearby stellar remnant.," High-resolution images that check for a shift in the source's light centroid can measure $\theta_E,$ if the lens is indeed a nearby stellar remnant." + Alternatively. they can place an upper limit on the value of θε. thereby at least providing a constraint on stellar-mass mocdels for the lens.," Alternatively, they can place an upper limit on the value of $\theta_E,$ thereby at least providing a constraint on stellar-mass models for the lens." + When an event with a small measured value of τι occurs. we are guaranteed that the lens is of special interest.," When an event with a small measured value of $\tau_E$ occurs, we are guaranteed that the lens is of special interest." + It could be a low-mass object such as a brown cdwarL. a planet bound (o a star. or a free-floating planet.," It could be a low-mass object such as a brown dwarf, a planet bound to a star, or a free-floating planet." + Or it could be a high-velocity object. either a dim member of the Galactic halo or else a high-velocity stellar remnant. a product of an earlier high-enerev event.," Or it could be a high-velocity object, either a dim member of the Galactic halo or else a high-velocity stellar remnant, a product of an earlier high-energy event." + Because short events are being discovered by current monitoring teams (Ucdalski 2003: Bond et 2001). it is possible to initiate programs designed to learn about both and high-speed lenses.," Because short events are being discovered by current monitoring teams (Udalski 2003; Bond et 2001), it is possible to initiate programs designed to learn about both low-mass and high-speed lenses." + The science returis will be high. because each class of lenses eeneraling short events is important. each for ils own reasons.," The science returns will be high, because each class of lenses generating short events is important, each for its own reasons." + Furthermore. the lenses will be nearby (within about a kpc) lor a large fraction of short events.," Furthermore, the lenses will be nearby (within about a kpc) for a large fraction of short events." + This means that (μον are amenable to a variety of complementary studies., This means that they are amenable to a variety of complementary studies. + Short events could. for example. provide a unique wav to discover local [astmoving neutron stars which would otherwise be known only as dim. unexceptional sources.," Short events could, for example, provide a unique way to discover local fast-moving neutron stars which would otherwise be known only as dim, unexceptional sources." + Short events will almost certainly identify planetary systems which can be studied in a variety of wavs., Short events will almost certainly identify planetary systems which can be studied in a variety of ways. +" Some of these will become the ""gold standards” of (heir class. because kev quantities such as the eravitational mass of the star and planet can be measured. sometimes in complementary wavs (DiStelano 2009)."," Some of these will become the “gold standards” of their class, because key quantities such as the gravitational mass of the star and planet can be measured, sometimes in complementary ways Stefano 2009)." + We have shown that even short events generated bv lree-[loating planets can be correctly identified. and planet properties can be measured.," We have shown that even short events generated by free-floating planets can be correctly identified, and planet properties can be measured." + It is significant (hat (his can be accomplished without sophisticated astrometry missions (cf., It is significant that this can be accomplished without sophisticated astrometry missions (cf. + Lana 2006)., Han 2006). + Our result relies on careful study of short event light curves. including multiwavelength monitoring to identify the effects of blending and finite-source size.," Our result relies on careful study of short event light curves, including multiwavelength monitoring to identify the effects of blending and finite-source size." + We also require comprehensive catalogs and/or additional hieh-qualitv observations of the region in which the event occurred to search for counterparts. sometimes during aud sometimes vears alter the event.," We also require comprehensive catalogs and/or additional high-quality observations of the region in which the event occurred to search for counterparts, sometimes during and sometimes years after the event." + Lensing will even provide gravitational mass measurements Lor some free-floating planets., Lensing will even provide gravitational mass measurements for some free-floating planets. + In the companion paper we demonstrate (hat the expected rate of short events is hieh enough that programs designed (ο study them will be productive., In the companion paper we demonstrate that the expected rate of short events is high enough that programs designed to study them will be productive. + We also suggest a set of procedures designed to optimize their efficacy., We also suggest a set of procedures designed to optimize their efficacy. + These procedures will not only. identilv individual interesting lenses. but will also allow population properties to be measured.," These procedures will not only identify individual interesting lenses, but will also allow population properties to be measured." + We, We + 2007).," \citep{you91, omo07}." + |. (Omout2007) (Meier&Turner2005)., $^+$ \citep{omo07} \citep{meie05}. +". roni these dense gas tracers. such as ICN, have been argued to be a much better probe of the star-forming nolecular gas than emission from CO. (Cao&Solomon200 1)."," from these dense gas tracers, such as HCN, have been argued to be a much better probe of the star-forming molecular gas than emission from CO \citep{gao04}." +. The utility of cussion from these high dipole momeut nolecules’ to probe the properties of the nuclear regious of galaxies is well established. and there have Όσοι jiunnerous papers on this subject iu recent vears (see review by Onout (2007))).," The utility of emission from these `high dipole moment molecules' to probe the properties of the nuclear regions of galaxies is well established, and there have been numerous papers on this subject in recent years (see review by \citet{omo07}) )." + Iu addition. several emission ine ratios. such as those for the low-lving trausitions of IICN/CO. IICO /IICN. and IINC/IICN. have been sugeested as good. !diagnostics of the properties of the deuse star-forming eas in galaxies (Aalto2008).," In addition, several emission line ratios, such as those for the low-lying transitions of HCN/CO, $^+$ /HCN, and HNC/HCN, have been suggested as good diagnostics of the properties of the dense star-forming gas in galaxies \citep{aal08}." +.. Nearly all studies of the molecular cussion iu galaxies have oen Dinuited to just a few tarected molecular dues., Nearly all studies of the molecular emission in galaxies have been limited to just a few targeted molecular lines. + However. Martíuetal.(2006) recently preseuted the first spectral scan of a galaxy providing an inveutory of the molecular lines ia NGC 253 within the 2 nuu wavelength baud.," However, \citet{mar06} recently presented the first spectral scan of a galaxy providing an inventory of the molecular lines in NGC 253 within the 2 mm wavelength band." + Such spectral scans can provide a nmcli nore complete description of the chemical complexities of he molecular gas in galaxies., Such spectral scans can provide a much more complete description of the chemical complexities of the molecular gas in galaxies. + The 3 mun wavelength bane is equally well suited. and past spectral surveys have xovided imuportaut iuformiation on the chemical iu physical properties of Ciaut Molecular Clouds (6MCs) iu ιο Milly Way (Johanssonetal.198f£.1985:ϱπαπατιςetal.1986:Turner 1989).," The 3 mm wavelength band is equally well suited, and past spectral surveys have provided important information on the chemical and physical properties of Giant Molecular Clouds (GMCs) in the Milky Way \citep{joh84, joh85, cum86, tur89}." +. One major disadvantage of -ost spectral line surveys is that the data are assemble Eoni Inany observations with varving pointing accuracy and with potentially systematic calibration problems., One major disadvantage of most spectral line surveys is that the data are assembled from many observations with varying pointing accuracy and with potentially systematic calibration problems. + The observations described here have been obtaiuc« Ssimultancousty the full spectral baud for cach galaxy. and hence many oversvsteniatie problems are eliminated.," The observations described here have been obtained over the full spectral band for each galaxy, and hence many systematic problems are eliminated." + Iuthis paper we present the first 3 nuu spectral scans of the, Inthis paper we present the first 3 mm spectral scans of the +"starts when the first resolved halos (which correspond to the first ionizing sources) form, around z~20.","starts when the first resolved halos (which correspond to the first ionizing sources) form, around $z\sim20$." +" Ouradopted f, values yield final overlap of ionized regions at Zoy~9 (8) for the Model 1 (Model 2) case, in rough agreement with the current observational constraints."," Ouradopted $f_\gamma$ values yield final overlap of ionized regions at $z_{\rm ov}\sim9$ (8) for the Model 1 (Model 2) case, in rough agreement with the current observational constraints." +" The corresponding integrated Thomson scattering optical depth seen by the Cosmic Microwave Background photons is Tes=0.094(0.069) is also in agreement with the latest constraints from WMAP satellite combined with the other available datasets, Τος=0.084+0.016 (2).."," The corresponding integrated Thomson scattering optical depth seen by the Cosmic Microwave Background photons is $\tau_{\rm es}=0.094 (0.069)$ is also in agreement with the latest constraints from WMAP satellite combined with the other available datasets, $\tau_{\rm es}=0.084\pm0.016$ \citep{2009ApJS..180..330K}." +" The early reionization (z> 14) is driven primarily by the low-mass sources, which have similar efficiencies in the two cases (the slightly lower source efficiency in the Model 2 case is compensated for by its higher collapsed fraction at the same redshift) and as a consequence the two reionization histories are initially very similar."," The early reionization $z>14$ ) is driven primarily by the low-mass sources, which have similar efficiencies in the two cases (the slightly lower source efficiency in the Model 2 case is compensated for by its higher collapsed fraction at the same redshift) and as a consequence the two reionization histories are initially very similar." +" Later on the larger sources take over, both because of their rapidly rising collapsed fraction (cf."," Later on the larger sources take over, both because of their rapidly rising collapsed fraction (cf." +" Figure 1)) and the strong Jeans suppression of the low-mass sources, and thus reionization proceeds more slowly in the Model 2 case due to the lower efficiency adopted for its high-mass sources."," Figure \ref{fcoll_sourcenum_fig}) ) and the strong Jeans suppression of the low-mass sources, and thus reionization proceeds more slowly in the Model 2 case due to the lower efficiency adopted for its high-mass sources." +" The fact that the mean mass-weighted ionized fraction, x is always larger than the corresponding volume-weighted one, x, (Figure 3,, upper panel) indicates that reionization proceeds in an inside-out fashion (i.e. high-density regions are preferentially ionized first) in both cases, in agreement with previous simulation results based on non-constrained realizations (?).."," The fact that the mean mass-weighted ionized fraction, $x_m$ is always larger than the corresponding volume-weighted one, $x_v$ (Figure \ref{global_reion_hist_fig}, upper panel) indicates that reionization proceeds in an inside-out fashion (i.e. high-density regions are preferentially ionized first) in both cases, in agreement with previous simulation results based on non-constrained realizations \citep{2006MNRAS.369.1625I}." +" For the Model 2 realization this ratio is noticeably higher, up to ~3 at z=20, due to its more advanced structure formation at any given redshift."," For the Model 2 realization this ratio is noticeably higher, up to $\sim3$ at $z=20$, due to its more advanced structure formation at any given redshift." +" This also boosts the clumpiness of the gas, and therefore the recombinations, which in turn extends reionization even further."," This also boosts the clumpiness of the gas, and therefore the recombinations, which in turn extends reionization even further." + The mean reionization history presented above ensures that the currently available global observational constraints - the electron-scattering optical depth and overlap epoch are satisfied., The mean reionization history presented above ensures that the currently available global observational constraints - the electron-scattering optical depth and overlap epoch are satisfied. +" However, to achieve our present goals we need to track separately the reionization history of the progenitors of each object of interest, namely the Local Group, as well as the nearby clusters of galaxies."," However, to achieve our present goals we need to track separately the reionization history of the progenitors of each object of interest, namely the Local Group, as well as the nearby clusters of galaxies." +" To this purpose, we extracted the Lagrangian mass distribution for each object (i.e. the mass which eventually will end up in that object by the present day) and followed the reionization history of all radiative transfer cells containing at least one particle which ends up in that object by z=0."," To this purpose, we extracted the Lagrangian mass distribution for each object (i.e. the mass which eventually will end up in that object by the present day) and followed the reionization history of all radiative transfer cells containing at least one particle which ends up in that object by $z=0$." +" The resulting local reionization histories are shown in Figure 4,, along with the global one for direct comparison."," The resulting local reionization histories are shown in Figure \ref{reion_hist_fig}, along with the global one for direct comparison." +" The Lagranian regions of both the LG and the nearby clusters are significantly overdense in either constrained realization and at all times, reflecting the fact that all of these objects correspond to high peaks of the density field."," The Lagranian regions of both the LG and the nearby clusters are significantly overdense in either constrained realization and at all times, reflecting the fact that all of these objects correspond to high peaks of the density field." +" The proto-Local Group region starts only moderately overdense, by ~7% (~9%) in the Model 1 (Model 2) case, which rises over time as the corresponding object collapses gravitationally, to reach ~16% (~25%) by the global overlap epoch at z~9 (8)."," The proto-Local Group region starts only moderately overdense, by $\sim7\%$ $\sim9\%$ ) in the Model 1 (Model 2) case, which rises over time as the corresponding object collapses gravitationally, to reach $\sim16\%$ $\sim25\%$ ) by the global overlap epoch at $z\sim9$ (8)." + The proto-clusters correspond to still higher peaks of the density field., The proto-clusters correspond to still higher peaks of the density field. +" Initially the proto-Virgo region is overdense by 12% for both simulations, rising over time to 24% (30%) for Model 1 (Model 2)."," Initially the proto-Virgo region is overdense by $12\%$ for both simulations, rising over time to $24\%$ $30\%$ ) for Model 1 (Model 2)." +" The proto-Fornax region (Model 2) starts 10% overdense, rising to 26% by the global overlap."," The proto-Fornax region (Model 2) starts $10\%$ overdense, rising to $26\%$ by the global overlap." + The higher local density yields an (exponentially) larger halo collapsed fraction and thus ionizing photon production., The higher local density yields an (exponentially) larger halo collapsed fraction and thus ionizing photon production. +" Therefore, for both LG and clusters we can expect local reionization to occur earlier than average, which is confirmed by our simulation results (Figure 4))."," Therefore, for both LG and clusters we can expect local reionization to occur earlier than average, which is confirmed by our simulation results (Figure \ref{reion_hist_fig}) )." +" In both simulations the LG reionization starts at about 12.5, at which time its oldest progenitor halos form."," In both simulations the LG reionization starts at about $z\sim12.5$ , at which time its oldest progenitor halos form." +" Before z~ the LG ionized fraction is tiny, below 3x10? (2x107°) for Model 1 (Model 2)."," Before $z\sim12.5$ the LG ionized fraction is tiny, below $3\times10^{-9}$ $2\times10^{-5}$ ) for Model 1 (Model 2)." +" Thereafter the (proto-)LG reionization accelerates, albeit only gradually."," Thereafter the (proto-)LG reionization accelerates, albeit only gradually." +" For Model 1 (Figure 4,, left) the LG ionized fraction reaches by z~12 and by z~10.5."," For Model 1 (Figure \ref{reion_hist_fig}, left) the LG ionized fraction reaches by $z\sim12$ and by $z\sim10.5$." + After that point the evolution becomes very fast and full ionization (fm> 99%) is achieved by z~10., After that point the evolution becomes very fast and full ionization $x_m>99\%$ ) is achieved by $z\sim10$. +" In contrast, the reionization history of the (proto-)Virgo cluster in the same Model 1 simulation is quite different."," In contrast, the reionization history of the (proto-)Virgo cluster in the same Model 1 simulation is quite different." + Proto-Virgo is a higher density peak and the formation of the local nonlinear structures is therefore accelerated., Proto-Virgo is a higher density peak and the formation of the local nonlinear structures is therefore accelerated. +" Hence the local reionization proceeds faster, as well."," Hence the local reionization proceeds faster, as well." +" The evolution remains smooth throughout, with no sudden changes of slope, unlike in the LG case."," The evolution remains smooth throughout, with no sudden changes of slope, unlike in the LG case." +" The mass-weighted ionized fraction reaches by z—12.6, by z— 11.3, and by z— 10.5."," The mass-weighted ionized fraction reaches by $z=12.6$, by $z=11.3$ , and by $z=10.5$ ." +" The reionization histories are similar in the Model 2 case (Figure 4,, right)."," The reionization histories are similar in the Model 2 case (Figure \ref{reion_hist_fig}, right)." +" Once again, the proto-cluster regions, both Virgo and Fornax, reionize earlier than LG and much ealier than an average region - z,,=0.1 is reached by z=12.3 (12.6), Im=0.5 by z=10.25 (10.5) and am=0.9 by z=9.2 (9.3) for Virgo (Fornax)."," Once again, the proto-cluster regions, both Virgo and Fornax, reionize earlier than LG and much ealier than an average region - $x_m=0.1$ is reached by $z=12.3$ (12.6), $x_m=0.5$ by $z=10.25$ (10.5) and $x_m=0.9$ by $z=9.2$ (9.3) for Virgo (Fornax)." + For both proto-clusters overlap (tm= 0.99) is reached at z=8.7 andthe evolution remains smooth throughout., For both proto-clusters overlap $x_m=0.99$ ) is reached at $z=8.7$ andthe evolution remains smooth throughout. +" Interestingly, most reionization stages (but not the local overlap, which is roughly simultaneous) of the Fornax reionization occur earlier than the corresponding ones for Virgo, even though Fornax has lower mass at the present epoch."," Interestingly, most reionization stages (but not the local overlap, which is roughly simultaneous) of the Fornax reionization occur earlier than the corresponding ones for Virgo, even though Fornax has lower mass at the present epoch." +" In comparison, the reionization of the LG occurs later, reaching tm=0.1 by z= 11.6, x4,=0.5 by z= 9.4, tm=0.9 by z=8.6 and local overlap is achieved by z=8.4."," In comparison, the reionization of the LG occurs later, reaching $x_m=0.1$ by $z=11.6$ , $x_m=0.5$ by $z=9.4$ , $x_m=0.9$ by $z=8.6$ and local overlap is achieved by $z=8.4$." +" It lags the global mean in its earliest stages (z> 12.5), but as more progenitor halos form it catches upand then speedsahead after z= 12."," It lags the global mean in its earliest stages $z>12.5$ ), but as more progenitor halos form it catches upand then speedsahead after $z=12$ ." +" The Local Group reionization history isagain much less smooth than the proto-cluster ones, with significant changes of slope around z~ 12, z~11 and 2~ 9.6."," The Local Group reionization history isagain much less smooth than the proto-cluster ones, with significant changes of slope around $z\sim12$ , $z\sim11$ and $z\sim9.6$ ." +Nipoti 2010; Haghi et al.,Nipoti 2010; Haghi et al. +" 2009, 2011)."," 2009, 2011)." +" Newtonian and MOND velocity dispersion profiles are clearly distinct, regardless of the assumed anisotropy profile."," Newtonian and MOND velocity dispersion profiles are clearly distinct, regardless of the assumed anisotropy profile." +" As already reported above, all the models presented in the previous section are based on the assumption that the cluster is isolated."," As already reported above, all the models presented in the previous section are based on the assumption that the cluster is isolated." +" However, it is known that the presence of the Galactic field is important for Pal 14 (see Sect. 1)),"," However, it is known that the presence of the Galactic field is important for Pal 14 (see Sect. \ref{intro_sec}) )," + with different effects depending of the considered gravity law., with different effects depending of the considered gravity law. + In particular: To account properly for the above effects we used a set of N- simulations performed in both Newtonian gravity and MOND., In particular: To account properly for the above effects we used a set of N-body simulations performed in both Newtonian gravity and MOND. + In the Newtonian case we studied the tidal effects by simulating the evolution of the cluster orbiting within the Galactic potential., In the Newtonian case we studied the tidal effects by simulating the evolution of the cluster orbiting within the Galactic potential. +" In the MOND case we used N-body simulations to model the cluster in the presence of a uniform external field, so for we neglected the tidal effects."," In the MOND case we used N-body simulations to model the cluster in the presence of a uniform external field, so for simplicity we neglected the tidal effects." +" We note that, given the very simplicitylong two-body relaxation time of Pal 14 (4,~19.9Gyr; S11), its time evolution can be simulated also with collisionless N-body codes, which is especially convenient in MOND, because of the difficulty in realizing a collisional N-body code in this non-linear theory."," We note that, given the very long two-body relaxation time of Pal 14 $t_{rh}\sim19.9\,Gyr$; S11), its time evolution can be simulated also with collisionless N-body codes, which is especially convenient in MOND, because of the difficulty in realizing a collisional N-body code in this non-linear theory." +" As already discussed in Sect. 3.2,"," As already discussed in Sect. \ref{nbody_sec}," + the tidal interaction between Pal 14 and the Milky Way is expected to heat the outskirts of the cluster., the tidal interaction between Pal 14 and the Milky Way is expected to heat the outskirts of the cluster. + This can be the result of both compressive shocks occurred during the disk crossing and perigalactic passages (Ostriker et al., This can be the result of both compressive shocks occurred during the disk crossing and perigalactic passages (Ostriker et al. + 1972; Gnedin Ostriker 1997) and the sudden change in the underlying potential (like in the case of the Sagittarius galaxy; Taylor Babul 2001)., 1972; Gnedin Ostriker 1997) and the sudden change in the underlying potential (like in the case of the Sagittarius galaxy; Taylor Babul 2001). +" After these episodes, the kinetic energy of a fraction of the cluster stars can exceed the boundaries of the cluster potential well and these stars become ""potential escapers"" (Kupper et al."," After these episodes, the kinetic energy of a fraction of the cluster stars can exceed the boundaries of the cluster potential well and these stars become ""potential escapers"" (Kupper et al." + 2010)., 2010). + These stars have larger than the other bound stars mightand remaina velocitywithin the dispersioncluster in spite of their positive energy for a timescale comparable to the cluster orbital period (Lee Ostriker 1987)., These stars might have a velocity dispersion larger than the other bound stars and remain within the cluster in spite of their positive energy for a timescale comparable to the cluster orbital period (Lee Ostriker 1987). +" Moreover, the"," Moreover, the" + ACDM cosmological models (e... Dekel&Silk(1986):IXIvpiuetal. (1999))) for structure formation in the universe predict niuiv compact darkanatter structures.ménmihalos. in the regions near galaxies and ealaxy groups.,"$\Lambda$ CDM cosmological models (e.g., \citet{dek86, kly99}) ) for structure formation in the universe predict many compact dark-matter structures, in the regions near galaxies and galaxy groups." + Ricotti(2009) predicts that nmünibalos may have a late phase of eas accretion. which will be detectable iu observations.," \citet{ric09} predicts that minihalos may have a late phase of gas accretion, which will be detectable in observations." + It has Όσο proposed that at least part of the population of neutral hydrogen. Tigh Velocity Clouds (IITVCs) surrounding the Milky Wav πας represent these embedded: dark. matter halos (Blitzctal.1999:Braun&Bin-ton1999:Ciovanellietal. 2010).," It has been proposed that at least part of the population of neutral hydrogen High Velocity Clouds (HVCs) surrounding the Milky Way may represent these -embedded dark matter halos \citep{blitz99,braunburt99,gio10}." +. Other possible origins of the IWVCs include galactic fountaius, Other possible origins of the HVCs include galactic fountains +"The results do not change significantly if the heat pulses are deposited at the loop footpoints, instead of uniformly, one of the options explored by ?)..","The results do not change significantly if the heat pulses are deposited at the loop footpoints, instead of uniformly, one of the options explored by \cite{Reale_2008}." + The time dependence is a pulse function: where f is the start time of heat pulse., The time dependence is a pulse function: where $t_0$ is the start time of heat pulse. + The amplitude of the pulse is Hy=0.38ergcms., The amplitude of the pulse is $H_0 = 0.38 ~ erg ~ cm^{-3} ~ s^{-1}$. +" Note that the duration of heat pulse is short enough to have a multi-temperature loop system, and not short enough to have large effects on the overall evolution and emission due to non-equilibrium of ionization (?).."," Note that the duration of heat pulse is short enough to have a multi-temperature loop system, and not short enough to have large effects on the overall evolution and emission due to non-equilibrium of ionization \citep{Reale_2008}." + The equations are solved numerically by means of the Palermo-Harvard loop code (??)..," The equations are solved numerically by means of the Palermo-Harvard loop code \citep{Peres_1982,Betta_1997}." +" This is a well tested and highly stable code, used for both flaring (??) and quiescent loops (?).."," This is a well tested and highly stable code, used for both flaring \citep{Peres_Reale_1987,Betta_al_2001} and quiescent loops \citep{Reale_al_2000_II}." + The Palermo-Harvard code has an adaptive grid (?).. to better describe the steep gradients along the strands and during the evolution.," The Palermo-Harvard code has an adaptive grid \citep{Betta_1997}, to better describe the steep gradients along the strands and during the evolution." +" For the case of post-processing, the code output results were interpolated on a fixed equispaced erid for the post-processing."," For the ease of post-processing, the code output results were interpolated on a fixed equispaced grid for the post-processing." + The grid is made of 1024 cells along the strand., The grid is made of 1024 cells along the strand. +" The code output consists of temperature, density and velocity distributions along the loop strand sampled with a regular cadence during the strand evolution driven by the heat pulse."," The code output consists of temperature, density and velocity distributions along the loop strand sampled with a regular cadence during the strand evolution driven by the heat pulse." +" As mentioned above we model only one strand and replicate it so that one strand is different from the other only for the start time of the heat pulse, 1.6. £j."," As mentioned above we model only one strand and replicate it so that one strand is different from the other only for the start time of the heat pulse, i.e. $t_0$." +" This choice minimizes the number of tree parameters, and a single simulation was enough for our further analysis."," This choice minimizes the number of free parameters, and a single simulation was enough for our further analysis." +" The simulation computes the evolution of the strand for 2000 s after the start of the heat pulse, and the solutions"," The simulation computes the evolution of the strand for $2000$ s after the start of the heat pulse, and the solutions" +We find some of our conclusions of the optimal BAO surveys to be comparable to the optimal configurations of other Dark Energy surveys.,We find some of our conclusions of the optimal BAO surveys to be comparable to the optimal configurations of other Dark Energy surveys. + The maximization of the survey area is the optimal configuration in both. Weak. Lensing (Yamamoto et al..," The maximization of the survey area is the optimal configuration in both Weak Lensing (Yamamoto et al.," + 2007) ancl ISW surveys (Douspis. ct al.," 2007) and ISW surveys (Douspis, et al.," + 2008)., 2008). + However. these types of surveys are not so sensitive to the choice of redshift range as BAO surveys.," However, these types of surveys are not so sensitive to the choice of redshift range as BAO surveys." + Lt would. be. possible το go bevond. the [lat ACDAL model with the dark energy. equation. of state described by something cilferent to the CPL parameterization., It would be possible to go beyond the flat $\Lambda$ CDM model with the dark energy equation of state described by something different to the CPL parameterization. + One example. would. be a form of ws) that. remains constant at early times before undergoing a rapid transition ab some redshift to a negative value. at. late times to drive the acceleration., One example would be a form of $w(z)$ that remains constant at early times before undergoing a rapid transition at some redshift to a negative value at late times to drive the acceleration. + This paramecterization has been studied in a number of publications (Bassett et 22002: Corasaniti Copeland 2003. ete).," This parameterization has been studied in a number of publications (Bassett et 2002; Corasaniti Copeland 2003, etc)." + However. the constraints on the parameters of this parameterization are often non-Gaussian. and so the predicted: constraints using a Fisher matrix approach are often incorrect. (when checked against à more rigorous analysis using ALCALC techniques).," However, the constraints on the parameters of this parameterization are often non-Gaussian, and so the predicted constraints using a Fisher matrix approach are often incorrect (when checked against a more rigorous analysis using MCMC techniques)." + While we investigated. optimal surveys using this dark cnereyv parameterization. the results proved. using present methodology. to be uncomfortably unreliable.," While we investigated optimal surveys using this dark energy parameterization, the results proved, using present methodology, to be uncomfortably unreliable." + We also showed that the Dexibility bounds on the survey parameters expand. as other datasets are added in as prior information., We also showed that the flexibility bounds on the survey parameters expand as other datasets are added in as prior information. + While the Hexibility bounds should not be too narrow. as this could Lead to fine tuning of the survey which mav not be realizable in practice. when the Uexibility bounds become too large it is because the instrument is having too small an impact its contribution to the total science from. all surveys up to that point will be negligible.," While the flexibility bounds should not be too narrow, as this could lead to fine tuning of the survey which may not be realizable in practice, when the flexibility bounds become too large it is because the instrument is having too small an impact — its contribution to the total science from all surveys up to that point will be negligible." +" This leaves a ""window of opportunity for à WEMOS-like survey such that PUN will be of scientific benefit if it is performed alter WigeleZ ixl BOSS. but will become obsolete if it post-dates a 4ky BAO survey performed. by JDIZM. Euelic or SIA."," This leaves a `window of opportunity' for a WFMOS-like survey such that it will be of scientific benefit if it is performed after WiggleZ and BOSS, but will become obsolete if it post-dates a full-sky BAO survey performed by JDEM, Euclid or SKA." + Finally. the conclusion that an optimal WEMOS-like gsurvey should target. exclusively 0.13$." +" However, the distribution of egw in function of the mass of a black hole is unknown."," However, the distribution of $\epsilon_{\rm GW}$ in function of the mass of a black hole is unknown." +" In particular, let us think of what occurs with other compact objects — namely, the neutron stars — to see if we can learn something from them."," In particular, let us think of what occurs with other compact objects $-$ namely, the neutron stars $-$ to see if we can learn something from them." +" A newly born neutron star could lose angular momentum due to gravitational waves associated with non-radial oscillations (Ferrari,Miniutti&Pons2003).", A newly born neutron star could lose angular momentum due to gravitational waves associated with non-radial oscillations \citep{ferr03}. +. This could explain why all known young neutron stars are relatively slow rotators., This could explain why all known young neutron stars are relatively slow rotators. +" 'The black holes could have had a similar history, i.e, they could have been formed rapidly rotating and lost momentum to gravitational radiation via their quasi-normal modes."," The black holes could have had a similar history, i.e, they could have been formed rapidly rotating and lost momentum to gravitational radiation via their quasi-normal modes." +" If this was the case, the value of egw could be near the maximum one, or in the worst case, it could have a value to produce (S/N)>3 for a LIGO III pair."," If this was the case, the value of $\epsilon_{\rm GW}$ could be near the maximum one, or in the worst case, it could have a value to produce ${\rm (S/N)} > 3$ for a LIGO III pair." +" In this work, we have used the hierarchical formation scenario derived from the Press-Schechter formalism to build the cosmic star formation rate - CSFR in a self-consistent way."," In this work, we have used the hierarchical formation scenario derived from the Press-Schechter formalism to build the cosmic star formation rate - CSFR in a self-consistent way." +" Our paper differs from earlier works basically in the form as is obtained the function p, (or CSFR).", Our paper differs from earlier works basically in the form as is obtained the function $\dot\rho_{\star}$ (or CSFR). +" In particular, from the hierarchical scenario we obtain the baryon accretion rate, a»(t), that supplies the gaseous reservoir in the halos."," In particular, from the hierarchical scenario we obtain the baryon accretion rate, $a_{\rm b}(t)$, that supplies the gaseous reservoir in the halos." +" Thus, the term ap(t) is treated as an infall term in our model."," Thus, the term $a_{\rm b}(t)$ is treated as an infall term in our model." +" This scenario is in agreement with the cold dark matter model of cosmological structure formation, where the first sources of light are expected to form in ~10°Mo dark matter potential at z>20."," This scenario is in agreement with the cold dark matter model of cosmological structure formation, where the first sources of light are expected to form in $\sim 10^{6}{\rm M}_{\odot}$ dark matter potential at $z\geq 20$." +" Using p, we calculate the stochastic background of gravitational waves produced by pre-galactic black holes.", Using $\dot\rho_{\star}$ we calculate the stochastic background of gravitational waves produced by pre-galactic black holes. +" We show that a significant amount of GWs is produced related to the history of CSFR studied here, and this SBGWs can in principle be detected by a pair of LIGO III interferometers."," We show that a significant amount of GWs is produced related to the history of CSFR studied here, and this SBGWs can in principle be detected by a pair of LIGO III interferometers." +" Note that signal-to-noise ratios (S/N)90 could be obtained if the efficiency of generation of GWs is close to the maximum value (ecw,,,,=7x 10-2), if the IMF produces a high number of massive remnants (x=0.35), and if 20."," Note that signal-to-noise ratios ${\rm (S/N)}\sim 90$ could be obtained if the efficiency of generation of GWs is close to the maximum value $\epsilon_{\rm GW_{max}}=7\times 10^{-4}$ ), if the IMF produces a high number of massive remnants $x=0.35$ ), and if $z_{\rm ini}\sim 20$ ." +" Considering a Salpeter IMF (x— 1.35), we obtain to-noise ratios (S/N)~10."," Considering a Salpeter IMF $x=1.35$ ), we obtain signal-to-noise ratios ${\rm (S/N)}\sim 10$." + 'The critical parameter to be constrained in the case of a non-detection is egw., The critical parameter to be constrained in the case of a non-detection is $\epsilon_{\rm GW}$. + A non-detection would mean that the efficiency of GWs during the formation of black holes is not high enough., A non-detection would mean that the efficiency of GWs during the formation of black holes is not high enough. +" In reality, eaw,,,, should be divided by a factor >35 in the case of a non-detection."," In reality, $\epsilon_{\rm GW_{max}}$ should be divided by a factor $> 35$ in the case of a non-detection." + It is worth mentioning that an IMF with x—2.35 could also be responsible for a non-detection same with egw= €GWyax- , It is worth mentioning that an IMF with $x=2.35$ could also be responsible for a non-detection same with $\epsilon_{\rm GW}=\epsilon_{\rm GW_{max}}$ . +"However, z=2.35 produces a high number of low mass stars that is not in agreement with recent numerical simulations of the collapse and fragmentation of primordial clouds (see, e.g., Abel,Bryan&Norman 2002))."," However, $x=2.35$ produces a high number of low mass stars that is not in agreement with recent numerical simulations of the collapse and fragmentation of primordial clouds (see, e.g., \citealp{abel}) )." + Another possibility for a non-detection is that the pre-galactic stars are such that the black holes formed had masses >500Mo., Another possibility for a non-detection is that the pre-galactic stars are such that the black holes formed had masses $>500{\rm M}_{\odot}$. +" In this case, the GW frequency band would be out of the LIGO bandwidth."," In this case, the GW frequency band would be out of the LIGO bandwidth." +" However, considering black holes formed from stars with masses 25ΜοSm140Mo, then the sensitivity of the future third generation of detectors could be high enough to increase one order of magnitude in the expected value of (S/N)."," However, considering black holes formed from stars with masses $25\ {\rm M}_{\odot}\lesssim m \lesssim 140\ {\rm M}_{\odot}$, then the sensitivity of the future third generation of detectors could be high enough to increase one order of magnitude in the expected value of (S/N)." + Examples of such detectors are the Large Scale Cryogenic Gravitational Wave Telescope (LCGT) and the European antenna EGO (see Regimbau&deFre-itasPacheco2006 and the references therein for a short discussion on thissubject)., Examples of such detectors are the Large Scale Cryogenic Gravitational Wave Telescope (LCGT) and the European antenna EGO (see \citealp{regimbau} and the references therein for a short discussion on thissubject). +" Specifically, around 650Hzthe planned strain noise for EGO will be a factor of ~4 higher than that provided for advanced LIGO configuration."," Specifically, around $650\ {\rm Hz}$the planned strain noise for EGO will be a factor of $\sim 4$ higher than that provided for advanced LIGO configuration." + This could represent a gain of a factor ~5—20 for the value of (S/N) considering two interferometers located at the same place (see, This could represent a gain of a factor $\sim 5-20$ for the value of (S/N) considering two interferometers located at the same place (see +It is. on the other hand. rather tricky to assess the effect of finite resolution on the coherence leneth of our model IGAIF. because the definition of coherence length for RAL is not completely clear and the estimation of coherence length. for imstance. for the filament IGME alone is not trivial.,"It is, on the other hand, rather tricky to assess the effect of finite resolution on the coherence length of our model IGMF, because the definition of coherence length for RM is not completely clear and the estimation of coherence length, for instance, for the filament IGMF alone is not trivial." + We (ried (to quantify coherence length in the following (three wavs: 1) We directly caleulated the coherence length of 2). that is. the length with the same sign of Dj. along LOSs.," We tried to quantify coherence length in the following three ways: 1) We directly calculated the coherence length of $B_{\parallel}$, that is, the length with the same sign of $B_{\parallel}$, along LOS's." + Figure 6 shows the PDF of the resulting coherence length through the WILLM. which composes mostly filaments.," Figure 6 shows the PDF of the resulting coherence length through the WHIM, which composes mostly filaments." + It peaks at the length of 3 zones corresponding to 586h.! kpe., It peaks at the length of 3 zones corresponding to $586\ h^{-1}$ kpc. + 2) We caleulated 3/4 times the integral scale. which is the coherence length defined for RAL in Che incompressible limit (see Introduction). for the IGME inside the whole computational box of (100'Mpe)*ni volume.," 2) We calculated 3/4 times the integral scale, which is the coherence length defined for RM in the incompressible limit (see Introduction), for the IGMF inside the whole computational box of $(100\ h^{-1}{\rm Mpc})^3$ volume." + Here. 23?(h) is the (hree-climensional power spectrum of magnetic fields (the third panel of Figure 4).," Here, $P_B^{3D}(k)$ is the three-dimensional power spectrum of magnetic fields (the third panel of Figure 4)." + We found the value to be ~800fA! kpe for our model IGMP., We found the value to be $\sim 800\ h^{-1}$ kpc for our model IGMF. + 3) We also calculated. the largest energv containing scale in the whole computational box. which is the peak scale of KkDP(E) (not shown).," 3) We also calculated the largest energy containing scale in the whole computational box, which is the peak scale of $kP_B^{3D}(k)$ (not shown)." + It is ~9005.! kpe for our model IGMP., It is $\sim 900\ h^{-1}$ kpc for our model IGMF. + Note that the latter two values include contributions from the IGME in Lilaments as well as in clusters. sheets. and voids.," Note that the latter two values include contributions from the IGMF in filaments as well as in clusters, sheets, and voids." + All (he three scales are comparable., All the three scales are comparable. + These length scales are ~ 3 (o 5 times larger than the grid resolution of our simulations. 195fh! kpe.," These length scales are $\sim$ 3 to 5 times larger than the grid resolution of our simulations, $195\ h^{-1}$ kpc." + Cho&Ryu(2009). studied characteristic lengtühs in incompressible simulations of MIID turbulence (see Introduction): based on it. they. predicted that the coherence length for RAL would be a few x1005.! kpe in filaments. while a lew x105.+ kpe in clusters.," \citet{cr09} studied characteristic lengths in incompressible simulations of MHD turbulence (see Introduction); based on it, they predicted that the coherence length for RM would be a few $\times\ 100\ h^{-1}$ kpc in filaments, while a few $\times\ 10\ h^{-1}$ kpc in clusters." + With our erid resolution of 195fh! kpc. the coherence length of the IGME in clusters should not be resolved ancl so our estimation of RAI lor clusters should be resolution-allected. as pointed in Section 3.," With our grid resolution of $195\ h^{-1}$ kpc, the coherence length of the IGMF in clusters should not be resolved and so our estimation of RM for clusters should be resolution-affected, as pointed in Section 3." +" On the other hand. while the predicted. coherence length for the IGME in filaments is still larger than the grid resolution. the estimated coherence length of D, for the ΥΛΗΝΤ is a couple of times lareer than the prediction for filaments."," On the other hand, while the predicted coherence length for the IGMF in filaments is still larger than the grid resolution, the estimated coherence length of $B_{\parallel}$ for the WHIM is a couple of times larger than the prediction for filaments." + It could be partly due to the limited resolution in our simulations., It could be partly due to the limited resolution in our simulations. + However. as noted in Section 3. IAM is dominantly contributed by the density peak along LOSs (Figure 2).," However, as noted in Section 3, RM is dominantly contributed by the density peak along LOS's (Figure 2)." + The above statements indicate Chat our estimate of the IAM through filaments is expected {ο have uncertainties. especially due to limited resolution of our simulations: (he error in our estimation could be up to a [factor of several.," The above statements indicate that our estimate of the RM through filaments is expected to have uncertainties, especially due to limited resolution of our simulations; the error in our estimation could be up to a factor of several." +"551741, with zj5o;=5.07 in our catalogue) that is neither uds.iir45.an ERG nor any kind of dwarf star.","51741, with $z_{phot}=5.07$ in our catalogue) that is neither an ERG nor any kind of dwarf star." + The left-hand panels of Fig., The left-hand panels of Fig. +" 12 show the z-band and images of this source, as well as its best-fitting SED."," \ref{fig_z5cand} show the $z$ -band and images of this source, as well as its best-fitting SED." +" Although other objects present in the field are close to this galaxy, they are at a distance d >2 and d> arcsec, in the z-band and images, respectively."," Although other objects present in the field are close to this galaxy, they are at a distance $d>$ 2 and $d\gsim 4$ arcsec, in the $z$ -band and images, respectively." +" So, the4 aperture photometry for uds.iir45.551741 should not significantly be contaminated by the light of neighbour sources."," So, the aperture photometry for 51741 should not significantly be contaminated by the light of neighbour sources." +" Therefore, we believe that uds.iir45.551741 could be a genuine z»5 galaxy in spite of its (R—z)<3 colour."," Therefore, we believe that 51741 could be a genuine $z>5$ galaxy in spite of its $(R-z)< 3$ colour." +" If it were the case, this would be a rare example of an old and massive galaxy present in the early Universe: the best-fitting SED suggests an age of ~ 1 Gyr and stellar mass Mz3.8x10!! Mo."," If it were the case, this would be a rare example of an old and massive galaxy present in the early Universe: the best-fitting SED suggests an age of $\sim$ 1 Gyr and a stellar mass $M \approx 3.8 \times 10^{11} \, M_\odot$ ." +" Finally,a one of our 10 z>5 candidates (id uds.iir45.996606) has (R—z)=3.03+0.56 and (z 0.37."," Finally, one of our 10 $z>5$ candidates (id 96606) has $(R-z)=3.03\pm0.56$ and $(z-J)=0.51\pm0.37$ ." + The estimated redshift of this source is 5.23., The estimated redshift of this source is $z_{phot}=5.23$ . + The fitting of its SED with an M dwarf star template can be rejected with more than 3c confidence (with respect to the fitting with a galaxy template)., The fitting of its SED with an M dwarf star template can be rejected with more than $\sigma$ confidence (with respect to the fitting with a galaxy template). +" This is a well isolated object in the optical bands, and does not show any sign of blending in the 4.5band (see Fig. 12))."," This is a well isolated object in the optical bands, and does not show any sign of blending in the band (see Fig. \ref{fig_z5cand}) )." +" Thus, it is very likely another genuinem z>5 source."," Thus, it is very likely another genuine $z>5$ source." + Our SED fitting indicates that this is a ~0.2 Gyr old galaxy with a stellar mass M£z4.6x101?Mo.," Our SED fitting indicates that this is a $\sim$ 0.2 Gyr old galaxy with a stellar mass $M \approx 4.6 \times 10^{10} \, M_\odot$." +" The bias introduced by doing a flux-limited selection of z>3 galaxies is quite clear: the most massive galaxies will be included in the sample, while there is a poor constraint on the less massive galaxy populations."," The bias introduced by doing a flux-limited selection of $z\geq 3$ galaxies is quite clear: the most massive galaxies will be included in the sample, while there is a poor constraint on the less massive galaxy populations." +" Optical selections that map the rest-frame UV light of z>3 galaxies, instead, favour the selection of sources with high levels of star formation and/or little dust obscuration, but their biases with respect to a stellar mass selection are not obvious."," Optical selections that map the rest-frame UV light of $z\geq 3$ galaxies, instead, favour the selection of sources with high levels of star formation and/or little dust obscuration, but their biases with respect to a stellar mass selection are not obvious." + Our aim in this section is to use our z>3 sample to investigate this problem., Our aim in this section is to use our $z \geq 3$ sample to investigate this problem. +" Current optical surveys can reach faint limits, usually R=27 or z=26 AB magnitudes (measured through 2-arcsec-diameter apertures) over reasonable-size fields, such as the UDS or the COSMOS field."," Current optical surveys can reach faint limits, usually $R=27$ or $z=26$ AB magnitudes (measured through 2-arcsec-diameter apertures) over reasonable-size fields, such as the UDS or the COSMOS field." +" So, within our sample, we searched for galaxies that would be missed even in these deep optical surveys."," So, within our sample, we searched for galaxies that would be missed even in these deep optical surveys." +" For a clear comparison, we limited the analysis to our catalogue with completeness, i.e. [4.5]<22.4, and all redshifts 3€z<4, to be able to explore a more or less wide stellar mass range."," For a clear comparison, we limited the analysis to our catalogue with completeness, i.e. $\leq 22.4$ , and all redshifts $3 \leq z<4$, to be able to explore a more or less wide stellar mass range." + Fig., Fig. +" 13 shows the percentage of our [4.5]22.4 galaxies at 3xz<4 that have R>26.6 and z>25.6 AB total magnitudes (roughly equivalent to R>27 and z>26 2-arcsec-diameter aperture magnitudes), i.e. that will be missed even by deep optical selections."," \ref{fig_missperc} shows the percentage of our $\leq 22.4$ galaxies at $3 \leq z<4$ that have $R>26.6$ and $z>25.6$ AB total magnitudes (roughly equivalent to $R>27$ and $z>26$ 2-arcsec-diameter aperture magnitudes), i.e. that will be missed even by deep optical selections." +" The total number of galaxies analysed in each stellar mass bin are 69, 53, 41, 12 and 5 for logy).€[10.7;10.9),[10.9;11.1),[11.1;11.3),[11.3;11.5) and [11.5;M 11.7), respectively."," The total number of galaxies analysed in each stellar mass bin are 69, 53, 41, 12 and 5 for $\log_{10} M \in [10.7;10.9), [10.9;11.1),[11.1;11.3),[11.3;11.5)$ and $[11.5;11.7)$ , respectively." + Our results show that typical deep optical surveys miss a significant fraction of massive galaxies at z=3—4., Our results show that typical deep optical surveys miss a significant fraction of massive galaxies at $z=3-4$ . +" The percentage of missed galaxies clearly increases with stellar mass: it is (20.3+4.8)% for galaxies with M~6x10?Mo, and as high as (60+22)% for the rare M~3—4x10!!Mo galaxies."," The percentage of missed galaxies clearly increases with stellar mass: it is $(20.3\pm4.8)\%$ for galaxies with $M\sim 6 \times 10^{10} \, M_\odot$, and as high as $(60 \pm 22)\%$ for the rare $M\sim 3-4 \times 10^{11} \, M_\odot$ galaxies." +" Note that, although the exact figures depend on the different magnitude cuts chosen for the R and z bands, our conclusions are still valid even when considering slightly deeper magnitudes."," Note that, although the exact figures depend on the different magnitude cuts chosen for the $R$ and $z$ bands, our conclusions are still valid even when considering slightly deeper magnitudes." +" For example, the fraction of galaxies with [4.5]&22.4 at 3€z«4 that have R>27 and z>26 AB total magnitudes range from (11.6+3.9)% to (29.34-7.0)96 and (60+22)%, for stellar masses M6x1019, 1.6x1011 and M—34x10!!Mo, respectively."," For example, the fraction of galaxies with $\leq 22.4$ at $3\leq z<4$ that have $R>27$ and $z>26$ AB total magnitudes range from $(11.6\pm3.9)\%$ to $(29.3 \pm 7.0)\%$ and $(60 \pm 22)\%$, for stellar masses $M\sim 6 \times 10^{10}$, $1.6 \times 10^{11}$ and $M\sim 3-4 \times 10^{11} \, \rm M_\odot$, respectively." +" 'The fact that the fraction of sources missed by deep optical surveys increases with stellar mass is directly related to an increase in internal extinction: the median extinction of our [4.5]«22.4 galaxies at 3 pJy, Spitzer/indicating that these galaxies indeed host dust-obscured activity."," A total of 4 out of the 5 galaxies with $\log_{10} M \in [11.5;11.7)$ in our $3\leq z<4$ sample are detected in the /MIPS 24 $\rm \mu m$ band with flux densities $S_\nu(24 \, \rm \mu m) > 95 \, \mu Jy$ , indicating that these galaxies indeed host dust-obscured activity." +" For two of these galaxies, in particular, the 24 yum fluxdensities are quite high (6ν(24m)>300 Jy), which at these redshifts could suggest the presence ofan obscured active galactic nucleus (AGN; see e.g.Desai etal."," For two of these galaxies, in particular, the 24 $\rm \mu m$ fluxdensities are quite high $S_\nu(24 \, \rm \mu m) > 300 \, \mu Jy$ ), which at these redshifts could suggest the presence ofan obscured active galactic nucleus (AGN; see e.g.Desai etal." + 2008)., 2008). +" However, we note that, for none of these two galaxies, the UV through near-IR SED shows any obvious sign of an AGN component, i.e. we get goodSED fittings using only stellar templates (Fig. 14))."," However, we note that, for none of these two galaxies, the UV through near-IR SED shows any obvious sign of an AGN component, i.e. we get goodSED fittings using only stellar templates (Fig. \ref{fig_massbright24}) )." +Two leading scenarios to explain the origin of gamma-ray bursts are (he collapsar and supranova (SA) models (see Mészáros(2002):Dermer(2002) [or review).,"Two leading scenarios to explain the origin of gamma-ray bursts are the collapsar and supranova (SA) models (see \citet{mes02,der02} for review)." + The collapsar model (e.g..Wang&Wooslev2002:1993) assumes that GRBs originate from the collapse of a massive star to a black hole.," The collapsar model \citep{ww02,woo93} assumes that GRBs originate from the collapse of a massive star to a black hole." + During the collapse process. a nuclear-density.," During the collapse process, a nuclear-density," +toward the host galaxy of GRB 000118 using the Plateau de Bure Interferometer (PdBlCaoulloteaue,toward the host galaxy of GRB 000418 using the Plateau de Bure Interferometer \citep[PdBI;][]{guil92}. . +tal. 1992).. 77. describes the host galaxy of GRB 0001185., \ref{sec:host} describes the host galaxy of GRB 000418. + ?7? outlines the observations aud data reduction. aud the results are presented iu 77.," \ref{sec:observation} outlines the observations and data reduction, and the results are presented in \ref{sec:result}." + In ?7.. we derive constraints on phlivysical quantities of the host galaxy aud discuss the nature of the galaxy.," In \ref{sec:discussion}, we derive constraints on physical quantities of the host galaxy and discuss the nature of the galaxy." +" A παν is preseuted is οο,", A summary is presented in \ref{sec:summary}. +" Throughout the paper. we adopt a cosmology with fly=70 kas | ο. Oy,=0.3. aud Ὃν=0.7."," Throughout the paper, we adopt a cosmology with $H_0=70$ km $^{-1}$ $^{-1}$, $\Omega_{\rm{M}}=0.3$, and $\Omega_{\Lambda}=0.7$." + UVfoptical/NIR observatious show that the host ealaxy is a blue. compact. subluuninous galaxy (Mj= at 2αι.1σι+0.0001 (Bloometal.2003).," UV/optical/NIR observations show that the host galaxy is a blue, compact, subluminous galaxy \citep[$M_B = -20.6$;][]{goro03} at $z=1.1181 \pm 0.0001$ \citep{bloo03}." +. The extinction-corrected SER derived. from the i|] line luninosity is 15.1 Af. yr+ (Bloometal.2003:Corosabel2003).," The extinction-corrected SFR derived from the ] line luminosity is 15.4 $M_{\odot}$ $^{-1}$ \citep{bloo03, goro03}." +. SED fits to the UV/optical/NIR. data show that the host is a vouus star-forming ealaxy with an SFR of ~10 20 AL. + (Gorosabeletal.2003:Clhiistensenct2001:Savaglioetal.2009:Svensson 2010).," SED fits to the UV/optical/NIR data show that the host is a young star-forming galaxy with an SFR of $\sim$ 10--20 $M_{\odot}$ $^{-1}$ \citep{goro03, chri04, sava09, sven10}." +. Subbu and radio observations with the Subuullimetre Comou- Bolometer Array (SCUBA) and the Χαν Large Airrav (VLA) detected a source at the position of the host galaxy with fluxes of S(S850¢an)=3.2+0.9 mih. S(1.13€IIz)=69+15 py. οσοκ)=16413 μ.]ν. and $(8.16GIIz)=514312 Jv (Bergeretal.2003).," Submm and radio observations with the Submillimetre Common-User Bolometer Array (SCUBA) and the Very Large Array (VLA) detected a source at the position of the host galaxy with fluxes of $S(\rm 850\mu m) = 3.2 \pm 0.9$ mJy, $S(\rm 1.43GHz) = 69\pm 15$ $\mu$ Jy, $S(\rm 4.86GHz) = 46 \pm 13$ $\mu$ Jy, and $S(\rm 8.46GHz) = 51 \pm 12$ $\mu$ Jy \citep{berg03}." +. The SERs derived from the subi aud radio ΟΙΙΟΣ are SFR(subuun)=690+195 AL. | and SERGadio)=330£75 M. E (Berecretal.2003)., The SFRs derived from the submm and radio emissions are ${\rm SFR(submm)} = 690 \pm 195$ $M_{\odot}$ $^{-1}$ and ${\rm SFR(radio)} = 330 \pm 75$ $M_{\odot}$ $^{-1}$ \citep{berg03}. +. The SED fit of Michalowskietal.(2008). raugiug from UV to radio waveleugths shows that the host galaxy is à young star-forming ealaxy with Ly;=Lo«10077 £. and SFR=2as M. vrο., The SED fit of \cite{mich08} ranging from UV to radio wavelengths shows that the host galaxy is a young star-forming galaxy with $L_{\rm IR} = 4.6 \times 10^{12}$ $L_{\odot}$ and ${\rm SFR} = 288$ $M_{\odot}$ $^{-1}$. + The IR huninosity. classifies the host ealaxy as an ultralwmunous infrared galaxy (UCLIRG)., The IR luminosity classifies the host galaxy as an ultraluminous infrared galaxy (ULIRG). + The large discrepancy between the SER based ou UV/optical/NIR observations aud the SER based ou subuun/racdio observations indicates that the bulk of the star formation is obscured by dust., The large discrepancy between the SFR based on UV/optical/NIR observations and the SFR based on submm/radio observations indicates that the bulk of the star formation is obscured by dust. + The PdBI observations were conducted on Aueust G and 7. 2006 using the D coufguration with five anteunas aud on March 13. 2007 usine the D configuration with six antennas.," The PdBI observations were conducted on August 6 and 7, 2006 using the D configuration with five antennas and on March 13, 2007 using the B configuration with six antennas." +" The phase center was positioned at a(J2000) = 12 25"" 19.35 ane à(J2000) = |20* 06 11"".0.", The phase center was positioned at $\alpha$ (J2000) = $12^h$ $25^m$ $19.3^s$ and $\delta$ (J2000) = $+20^{\circ}$ $06'$ $11''.0$. + The redshifted CO (21) ane CO (13) lines were sinultaneouslv observed at. 2-nua and 1.3-mun bauds. respectively.," The redshifted CO (2–1) and CO (4–3) lines were simultaneously observed at 3-mm and 1.3-mm bands, respectively." + Receiver 1 vas tuned to 108.812 Cz (πα baud) for the upper sidehbancl are the receiver 2 was tuned to 217.667 GIIz 1απ baud) for the lower sideband., Receiver 1 was tuned to 108.842 GHz (3-mm band) for the upper sideband and the receiver 2 was tuned to 217.667 GHz (1.3-mm band) for the lower sideband. + The correlator was equippcc with 58O-MITz bandwidth im cach sideband im the 2006 observations and 1-GIIz bandwidth iu cach sidebar in the 2007 observations., The correlator was equipped with 580-MHz bandwidth in each sideband in the 2006 observations and 1-GHz bandwidth in each sideband in the 2007 observations. + The system temperature of receiver d was To.~200 300 Is in 2006 aud To.~100 200 I in 2007 (SSB)., The system temperature of receiver 1 was $T_{\rm sys} \sim 200$ –300 K in 2006 and $T_{\rm sys} \sim 100$ –200 K in 2007 (SSB). + Because the atmospheric conditions in the 1αι baud for the CO (13) line were unfavorable (TL.21000 KR). we use only 22m data for the CO (21) line in what follows.," Because the atmospheric conditions in the 1.3-mm band for the CO (4–3) line were unfavorable $T_{\rm sys} \gsim 1000$ K), we use only 3-mm data for the CO (2–1) line in what follows." + Data reduction and imaging were carried out using the CLIC program in the GILDAS package (Cuilloteau&Lucas2000)., Data reduction and imaging were carried out using the CLIC program in the GILDAS package \citep{guil00}. +. Passhand calibratious were performed using bright QSOs observed during the track., Passband calibrations were performed using bright QSOs observed during the track. + Flux calibrations were performed using standard calibrators., Flux calibrations were performed using standard calibrators. + Αααι seuxitivitv was achieved by adopting natural weiehtius. which gave a final svuthesized beam size of OSL41756 (position augle = 13°).," Maximum sensitivity was achieved by adopting natural weighting, which gave a final synthesized beam size of $2\farcs54 \times 1\farcs56$ (position angle = $13^{\circ}$ )." + Neither CO line emission nor continuuni enussion is detected (Fig. 1)., Neither CO line emission nor continuum emission is detected (Fig. \ref{fig:map}) ). + The ruis noise level is 1.3 1àJy: beam with 100 kan + resolution., The rms noise level is 1.3 mJy $^{-1}$ with 100 km $^{-1}$ resolution. + Suunuing the signals within the bandwidth. we obtained au rus noiselevel of (0.15 uy + for the 3-man continuum (rest frame 1.3 nun)," Summing the signals within the bandwidth, we obtained an rms noiselevel of 0.15 mJy $^{-1}$ for the 3-mm continuum (rest frame 1.3 mm)." + The 26 upper limits on the CO flux and contiuuuu flux at the position of the host ave 2.4 mw (100. n resolutiou) aud 0.62 mJy +. respectively.," The $\sigma$ upper limits on the CO flux and continuum flux at the position of the host are 2.4 mJy $^{-1}$ (100 km $^{-1}$ resolution) and 0.62 mJy $^{-1}$, respectively." + This upper Buton contimunu fiux is consistent with the SED iuodel of Alichalowskietal.(2008)., This upper limiton continuum flux is consistent with the SED model of \cite{mich08}. +". The CO hue huninosity (Lig) is given as follows (Solomonctal.1992):: where Leg is measured in EK kin 1 per. Seo ds the observed CO flux in Jv. Ac is the velocity widtli in lau ο, aud Dy is the huninosity distance in Ape."," The CO line luminosity $L'_{\rm CO}$ ) is given as follows \citep{solo92}: where $L'_{\rm CO}$ is measured in K km $^{-1}$ $^2$, $S_{\rm CO}$ is the observed CO flux in Jy, $\Delta v$ is the velocity width in km $^{-1}$, and $D_L$ is the luminosity distance in Mpc." +" Assnnins a velocity width of 300kins+t.which is the typical value for local ULIRGs (Solomonetal.1997).. the 26 upper limit of CO (21) line luninosity is Γιο» 1) <6.9410) Klaus | pe. Molecular oOeas mass is Oogiven by where aco is the CO-to-II; couversion factor iu AL. (KR kan ft pet) ο,"," Assuming a velocity width of $300\ \rm{km\ s^{-1}}$,which is the typical value for local ULIRGs \citep{solo97}, , the $\sigma$ upper limit of CO (2–1) line luminosity is $L'_{\rm CO}$ (2–1) $< 6.9 \times 10^9$ K km $^{-1}$ $^2$ Molecular gas mass is given by where $\alpha_{\rm CO}$ is the $_2$ conversion factor in $M_{\odot}$ (K km $^{-1}$ $^2$ $^{-1}$." + We assunie that the gas is optically thick and thermalized and has a CO (2 (10) Πο ratio of uuitv., We assume that the gas is optically thick and thermalized and has a CO (2--1)/CO (1–0) luminosity ratio of unity. +" The 26 upper limit ou molecular gas nues is Mg,<55«LO AL... and it is obtained by adopting a conversion factor of oco=O08 AL. US das | pet) Hl which is the standard value for ULIRGs (Downes&Solomon 19958).."," The $\sigma$ upper limit on molecular gas mass is $M_{\rm H_2} < 5.5 \times 10^9$ $M_{\odot}$, and it is obtained by adopting a conversion factor of $\alpha_{\rm CO} = 0.8$ $M_{\odot}$ (K km $^{-1}$ $^2$ $^{-1}$ , which is the standard value for ULIRGs \citep{down98}. ." +" This is lower than the median value of (3.0+1.6)«1079. AL, obtained for a sample of submillimeter galaxies (SACs:2005)..Note that if we adopta Galactic conversion factor of aco=L6 AL. (K dans ! per} E (Solomon&Barrett 1991).. the 26. upper luit would lncrease by about a factor of 6."," This is lower than the median value of $(3.0 \pm 1.6) \times 10^{10}$ $M_{\odot}$ obtained for a sample of submillimeter galaxies \citep[SMGs;][]{grev05}.Note that if we adopta Galactic conversion factor of $\alpha_{\rm CO} = 4.6$ $M_{\odot}$ (K km $^{-1}$ $^2$ $^{-1}$ \citep{solo91}, the $\sigma$ upper limit would increase by about a factor of 6." + The dust mass can be derived fom subi flux.," The dust mass can be derived from submm flux," +"At energies above 14keV, measured by Swift/BAT, marginally significant variation is detected in the monthly lightcurve with a probability for the fit of a constant of p,»~1% over the 58 months of (?)).","At energies above 14keV, measured by /BAT, marginally significant variation is detected in the monthly lightcurve with a probability for the fit of a constant of $p_{\chi^2} \approx 1\% $ over the 58 months of \cite{Baumgartner2010}) )." +" However, since a long integration time is needed, a observatio]behaviour similar to that in the 2-10 keV range would be difficult to detect."," However, since a long integration time is needed, a behaviour similar to that in the 2-10 keV range would be difficult to detect." +" The simultaneous observations in the optical by ATOM, the UV, and X-rays by are considered to study the synchrotron spectrum (see Fig. [f],"," The simultaneous observations in the optical by ATOM, the UV, and X-rays by are considered to study the synchrotron spectrum (see Fig. \ref{SED_Sy}," +" simultaneous data of August, 21 2009)."," simultaneous data of August, 21 2009)." +" The host galaxy of 1ES 02294200 is an elliptical galaxy with a brightness of postr=15.85€0.01mag and a half-light radius of re=3.25+0.07"" (?).."," The host galaxy of 1ES 0229+200 is an elliptical galaxy with a brightness of $m_{\rm{host,R}} = 15.85 \pm 0.01\;\rm{mag}$ and a half-light radius of $r_e = 3.25\pm 0.07''$ \citep{Urry2000}." +" Other observations in the Bessel R-band with the Nordic Optical Telescope (NOT) (?) show results with mpostr=15.76mag and rer=4""."," Other observations in the Bessel R-band with the Nordic Optical Telescope (NOT) \citep{Falomo1999} show results with $m_{\rm{host,R}} = 15.76 \; \rm{mag}$ and $r_{\rm{e,R}} = 4 ''$." +" The host galaxy profile of 1ES 0229-200 was also studied in the Bessel U, B, and V-bands with the Nordic Optical Telescope (NOT) by ?.."," The host galaxy profile of 1ES 0229+200 was also studied in the Bessel U, B, and V-bands with the Nordic Optical Telescope (NOT) by \cite{Hyvoenen2007}." +" The results are rmgo,;g=18.75mag, Mpos,u=18.83mag, and Mpos,v=17.58mag with half light radii rep=5.65"", 2.75"", and Τον=4.90”."," The results are $m_{\rm{host,B}} = 18.75\;\rm{mag}$, $m_{\rm{host,U}} = 18.83\;\rm{mag}$, and $m_{\rm{host,V}} = 17.58\;\rm{mag}$ with half light radii $r_{\rm{e,B}} = 5.65 ''$, $r_{\rm{e,U}} = 2.75''$ , and $r_{\rm{e,V}} = 4.90''$." +" In order to correct for the host galaxy light, a de Vaucouleurs profile of the galaxy was assumed and the measured brightness was transformed to that of an aperture of 4"" using equation (4) of ?.."," In order to correct for the host galaxy light, a de Vaucouleurs profile of the galaxy was assumed and the measured brightness was transformed to that of an aperture of $4''$ using equation (4) of \cite{Young1976}." + ATOM photometry was performed with a 4” aperture., ATOM photometry was performed with a $4''$ aperture. +" The resulting host-galaxy corrected fluxes would result in an unphysical bump in the V band in the SED (in vF,,) since the calculated influence of the host galaxy in the V band is only 30%, while in the R and B band it is 90% and 30%, respectively."," The resulting host-galaxy corrected fluxes would result in an unphysical bump in the V band in the SED (in $\nu F_{\nu}$ ) since the calculated influence of the host galaxy in the V band is only $30\%$, while in the R and B band it is $90\%$ and $30\%$, respectively." +" Hence, the influence of the host galaxy in the R, B, V, and filters were also calculated using a spectral template for a nearbyU elliptical galaxy by ?.."," Hence, the influence of the host galaxy in the R, B, V, and U filters were also calculated using a spectral template for a nearby elliptical galaxy by \cite{Fukugita1995}." +" Here, the R-band magnitudes and the half-light radius detected by ? were used."," Here, the R-band magnitudes and the half-light radius detected by \cite{Falomo1999} were used." +" The influence of the host galaxy is then ~90%, ~80%, and ~57% for R, V, and B, respectively, which are the values that we used to correct the measured fluxes (shown in Fig. B)."," The influence of the host galaxy is then $\sim 90\%$, $\sim 80\%$, and $\sim 57\%$ for R, V, and B, respectively, which are the values that we used to correct the measured fluxes (shown in Fig. \ref{SED_Sy}) )." + The host-galaxy corrected flux in the R- is compatible with the detected nucleus magnitude of and ? as expected from the absence of variability., The host-galaxy corrected flux in the R-band is compatible with the detected nucleus magnitude of \cite{Falomo1999} and \cite{Urry2000} as expected from the absence of variability. + The host galaxy influence in the UVW1 and UVM2 bands is unknown., The host galaxy influence in the UVW1 and UVM2 bands is unknown. +" Figure4] therefore shows two values, connected by a bar."," Figure \ref{SED_Sy} therefore shows two values, connected by a bar." +" The upper ones correspond to the measured values corrected only for extinction, the lower ones also assume a correction for the host galaxy of 30% (the value derived in the adjacent U band)."," The upper ones correspond to the measured values corrected only for extinction, the lower ones also assume a correction for the host galaxy of $30\%$ (the value derived in the adjacent U band)." +" In an independent check, the spectral slope was extracted using the Sloan Digital Sky Survey (SDSS) observations (five bands taken simultaneously)."," In an independent check, the spectral slope was extracted using the Sloan Digital Sky Survey (SDSS) observations (five bands taken simultaneously)." + The resulting slope is identical to the one measured by ATOM., The resulting slope is identical to the one measured by ATOM. +" Since the data do not match the epoch of the ATOM observations, the absolute fluxes were not considered."," Since the data do not match the epoch of the ATOM observations, the absolute fluxes were not considered." +" Since the host galaxy magnitudes were not extinction corrected, the influence of the host galaxy was first subtracted."," Since the host galaxy magnitudes were not extinction corrected, the influence of the host galaxy was first subtracted." + The extinction correction was thereafter applied to the AGN http://heasarc.gsfc.nasa.gov/docs/Swift/results/bs58monjight., The extinction correction was thereafter applied to the AGN light. +literature (see for instance Benn Ellison 1998)).,literature (see for instance Benn Ellison \cite{benn}) ). +" Actually, one interesting fact that emerges from this analysis is that, when the nights are short (filled symbols), i.e. during austral summer, most of the line enhancements tend to take place in the second half of the night."," Actually, one interesting fact that emerges from this analysis is that, when the nights are short (filled symbols), i.e. during austral summer, most of the line enhancements tend to take place in the second half of the night." +" On the other hand, when the nights are long (i.e. lasting more than 8.5 hours), flux changes tend to become smaller and smaller during the last two hours of the night."," On the other hand, when the nights are long (i.e. lasting more than 8.5 hours), flux changes tend to become smaller and smaller during the last two hours of the night." +" No such dichotomy is observed for the [OI]5577 line (see Fig. 12,,"," No such dichotomy is observed for the [OI]5577 line (see Fig. \ref{fig:decay}," + lower panel)., lower panel). +" Finally, as for the broad band data (see Sec. 4.1)),"," Finally, as for the broad band data (see Sec. \ref{sec:timescale}) )," + I have calculated the correlation coefficient between line fluxes and SFD for different values of the time delay τ and time window At., I have calculated the correlation coefficient between line fluxes and SFD for different values of the time delay $\tau$ and time window $\Delta \tau$. +" An instructive example is presented in Fig. 13,,"," An instructive example is presented in Fig. \ref{fig:scoi5577}," +" showing the case of the [05577 line, which displays the strongest dependency on solar activity."," showing the case of the [OI]5577 line, which displays the strongest dependency on solar activity." + The correlation function shows a peak at about t=15 days for Ar-1 day.," The correlation function shows a peak at about $\tau$ =15 days for $\Delta +\tau$ =1 day." +" The maximum correlation increases for a Ar=10 days, peaking at τ ~12.5 days."," The maximum correlation increases for a $\Delta \tau$ =10 days, peaking at $\tau\simeq$ 12.5 days." +" As for the photometric data (see Fig. 7)),"," As for the photometric data (see Fig. \ref{fig:sunanal}) )," + spurious correlation peaks due to the 27 days solar rotation are present., spurious correlation peaks due to the 27 days solar rotation are present. + A similar analysis for the [OI]6300 line shows two similar peaks at t 513 days and 27 days later., A similar analysis for the [OI]6300 line shows two similar peaks at $\tau\simeq$ 13 days and 27 days later. + The Na I D doublet is known to originate in a layer placed at about 92 km and to undergo strong seasonal variations around an average value of 50 R (Roach Gordon 1973))., The Na I D doublet is known to originate in a layer placed at about 92 km and to undergo strong seasonal variations around an average value of 50 R (Roach Gordon \cite{roach}) ). +" For the Paranal site, the doublet varies by a factor larger than 15, reaching a maximum value of 160 R in April."," For the Paranal site, the doublet varies by a factor larger than 15, reaching a maximum value of 160 R in April." + The maximum contribution of the Na I D doublet to the global night sky brightness reaches ~0.1 mag arcsec~? both in V and R passbands (see Table 5))., The maximum contribution of the Na I D doublet to the global night sky brightness reaches $\sim$ 0.1 mag $^{-2}$ both in $V$ and $R$ passbands (see Table \ref{tab:oi}) ). + While the Na I D doublet intensity does not show any significant correlation with solar activity (the correlation coefficient is only 0.07; see Fig., While the Na I D doublet intensity does not show any significant correlation with solar activity (the correlation coefficient is only 0.07; see Fig. +" 14 and Table 6)), it does show a clear SAO, with a peak in April and a secondary peak in October/November (see Fig. 14,,"," \ref{fig:naid} and Table \ref{tab:lines}) ), it does show a clear SAO, with a peak in April and a secondary peak in October/November (see Fig. \ref{fig:naid}," +" upper panel), as it is seen in all atomic lines analyzed in this work."," upper panel), as it is seen in all atomic lines analyzed in this work." +" The SAO displayed by mesospheric sodium column density is a well studied phenomenon (Kirchoff 1986)), also because of the importance of the sodium layer for the laser guide star adaptive optics systems (see for instance Ageorges Hubin 2000))."," The SAO displayed by mesospheric sodium column density is a well studied phenomenon (Kirchoff \cite{kirkhoff}) ), also because of the importance of the sodium layer for the laser guide star adaptive optics systems (see for instance Ageorges Hubin \cite{ageorges}) )." + The fact that different species like Na I and [OI] all show a SAO is interpreted in the light of the role of O atoms in the source photochemical reactions (Slanger 2005))., The fact that different species like Na I and [OI] all show a SAO is interpreted in the light of the role of O atoms in the source photochemical reactions (Slanger \cite{slanger}) ). +" Since the resolution provided by the grism 600R is sufficient to resolve the D; and D2 components (see Fig. 15)),"," Since the resolution provided by the grism 600R is sufficient to resolve the $_1$ and $_2$ components (see Fig. \ref{fig:uves}) )," +" following the work done by Slanger et al. (2005)),"," following the work done by Slanger et al. \cite{slanger}) )," +" I have estimated the intensity ratio D2/D; during dark time, on a total of 147 spectra."," I have estimated the intensity ratio $_2$ $_1$ during dark time, on a total of 147 spectra." +" The region of interest is contaminated by the presence of at two features belonging to the OH(8-2) band, which I will indicate as B and C (see Fig."," The region of interest is contaminated by the presence of at two features belonging to the OH(8-2) band, which I will indicate as B and C (see Fig." + 15 for their identification)., \ref{fig:uves} for their identification). +" Using a high resolution UVES spectrum of the night sky (Hanuschik 2003)), I have estimated the intensity ratio between these two features and another OH(8-2) line at"," Using a high resolution UVES spectrum of the night sky (Hanuschik \cite{hanuschik}) ), I have estimated the intensity ratio between these two features and another OH(8-2) line at" +Oxygen abundance estimates in star-forming regions play a crucial role in the understanding of galaxy evolution.,Oxygen abundance estimates in star-forming regions play a crucial role in the understanding of galaxy evolution. + For example. oxygen radial gradients in spiral galaxies obtained by region observations (e.g. Stanghellinietal.2010: Kewleyetal. 2010:: Bresolinetal. 2009:: Krabbeetal.2008: Dors&Copetti 2005:: Kennicuttetal. 2003)) are essential to test chemical evolution models (see Molla&Díaz 20055) and to investigate the effect of environment on galaxy interactions (Ellisonetal.2010:Dors&Copetti2006:Skillmanetal.1996) as well as the mass-metallicity relation of galaxies (e.g. Pilyuginetal.2004: Pérez-Montero&Contini 20093).," For example, oxygen radial gradients in spiral galaxies obtained by region observations (e.g. \citealt{stanghellini10}; \citealt{kewley10}; \citealt{bresolin09}; ; \citealt{krabbe08}; \citealt{dors05}; \citealt{kennicutt03}) ) are essential to test chemical evolution models (see \citealt{molla05}) ) and to investigate the effect of environment on galaxy interactions \citep{ellison10,dors06,skillman96} as well as the mass-metallicity relation of galaxies (e.g. \citealt{pyliugin04}; \citealt{perez09}) )." + Likewise. oxygen abundance estimates in metal- galaxies are also important to test theories of chemical evolution of galaxies because these are the least chemically evolved objects (Kunth&Sargent1983).," Likewise, oxygen abundance estimates in metal-poor galaxies are also important to test theories of chemical evolution of galaxies because these are the least chemically evolved objects \citep{kunt83}." + Unfortunately. for the most of star-forming regions. only collisionally excited. emission-lines (CELs) in the optical are bright enough to be used for the derivation of elemental abundance.," Unfortunately, for the most of star-forming regions, only collisionally excited emission-lines (CELs) in the optical are bright enough to be used for the derivation of elemental abundance." +" CELs are temperature sensitive. thus. only an accurate determination of the metallicity can be achieved from the previous estimation of the electron temperature (this method will be called) 7Z.-method) using. for instance. the ratio of different CELs 959+,15007)/,14363. which are weak or unobservable in star-forming regions with high metallicity and/or low excitation (Dorsetal.2008:Díazal. 2007)."," CELs are temperature sensitive, thus, only an accurate determination of the metallicity can be achieved from the previous estimation of the electron temperature (this method will be called $T_{\rm e}$ -method) using, for instance, the ratio of different CELs $(\lambda4959+\lambda5007)/\lambda4363$ which are weak or unobservable in star-forming regions with high metallicity and/or low excitation \citep{dors08,diaz07}." +. In these cases. oxygen abundances can be obtained by empirical (ie. using oxygen determinations via 7\.-method) or theoretical (i.e. using photoionizationmodels) calibrations between oxygen abundances and more easily measured line ratios (hereafter strong-line methods).," In these cases, oxygen abundances can be obtained by empirical (i.e. using oxygen determinations via $T_{\rm e}$ -method) or theoretical (i.e. using photoionizationmodels) calibrations between oxygen abundances and more easily measured line ratios (hereafter strong-line methods)." +" The oxygen abundance indicator Rz(QOuL037274[O mL143959,15007 yHB proposed by Pageletal.(1979) has found large acceptance in this context and several authors have calibrated this line ratio with O/H abundance (e.g. Edmunds&Pagel1984: Dopita&Evans 1986: Pilyugin 2001::Dors&Copetti 2005::among others).Additional O/H indicators based on other emission lines such as N»2[N 16584/He 1994)... u06583/[O15007"," The oxygen abundance indicator $R_{23}$ $\lambda$ $\lambda$ $\lambda$ $\beta$ proposed by \citet{pagel79} has found large acceptance in this context and several authors have calibrated this line ratio with O/H abundance (e.g. \citealt{edmunds84}; \citealt{dopita86}; \citealt{pilyugin01}; \citealt{dors05}; among others).Additional O/H indicators based on other emission lines such as $N_{2}$ $\lambda$ $\alpha$ \citep{thaisa94}, , $\lambda$ $\lambda$ 5007" +in the cores by a factor of its initial mass-to-maeuetic field flux ratio.,in the cores by a factor of its initial mass-to-magnetic field flux ratio. + In other words. a stronger magnetic feld with respect to the mass of a core causes more effective segregation.," In other words, a stronger magnetic field with respect to the mass of a core causes more effective segregation." + Although this scerceation occurs while the aubipolar diffusion appears. before dvuamical collapse. the signature footprint could remain iu the envelopes of Class 0. YSOs.," Although this segregation occurs while the ambipolar diffusion appears, before dynamical collapse, the signature footprint could remain in the envelopes of Class 0 YSOs." + On the other hand. although this effect would be minor to the features we have discussed because the segregation is effective. to relatively πια. eraius (¢Xm10 0). it is noteworthy that it would set the initial erain distribution of Class 0 YSO envelopes for more efficient erowth iu the ceutral region.," On the other hand, although this effect would be minor to the features we have discussed because the segregation is effective to relatively small grains $a \lesssim 10^{-4}$ cm), it is noteworthy that it would set the initial grain distribution of Class 0 YSO envelopes for more efficient growth in the central region." + We carried out interferometric observations towards three Class 0 YSOs (Lillis IRS 2. Lills IRS 3. aud Ll157)at aud 2.7 nua coutiuuun using CARMA.," We carried out interferometric observations towards three Class 0 YSOs (L1448 IRS 2, L1448 IRS 3, and L1157) at and 2.7 mm continuum using CARMA." + The coutimmun at these millumeter wavelengths is mainly thermal dust enmission of their envelopes., The continuum at these millimeter wavelengths is mainly thermal dust emission of their envelopes. + Our observations have Όσοι designed. particularly to cover comparablewe ranges at +he two waveleneths. which allowed us to tackle dust eran opacity spectral indexes (7) of Class 0 YSOs. using unprecedented compact coufiguration and high nage fidelity.," Our observations have been designed particularly to cover comparable ranges at the two wavelengths, which allowed us to tackle dust grain opacity spectral indexes $\beta$ ) of Class 0 YSOs, using unprecedented compact configuration and high image fidelity." + Through simultaneous modeling of the wo wavelength visibilitics as well as comparisons of the ο.nages and visibilities for the first time. we found not xlv the 3 of Class 0 YSOs but also its racial dependence.," Through simultaneous modeling of the two wavelength visibilities as well as comparisons of the images and visibilities for the first time, we found not only the $\beta$ of Class 0 YSOs but also its radial dependence." + —1 addition. we addressed the single power-law density oeidex p of Class 0 YSO cuvelopes.," In addition, we addressed the single power-law density index $p$ of Class 0 YSO envelopes." + 1., 1. + We found that the dust opacity spectral index. of +je earliest YSOs. so-called Class 0. is around 1.," We found that the dust opacity spectral index $\beta$ of the earliest YSOs, so-called Class 0, is around 1." + This oeuplies that dust eras have significantly: erown already at the carlicst stage., This implies that dust grains have significantly grown already at the earliest stage. + 2., 2. + We obtained the power-law deusitv iudex p of 1.85. —2.6. aud ~1.7 for ΤάΤΙ IRS 2. L11Is IRS 3B. and LILS7T. respectively.," We obtained the power-law density index $p$ of $\sim 1.8$ , $\sim +2.6$, and $\sim 1.7$ for L1448 IRS 2, L1448 IRS 3B, and L1157, respectively." + Athough we did uot attempt to constrain star formation theories. we pointed out the difference between that of Τά115 IRS 3B aud those of the other two.," Although we did not attempt to constrain star formation theories, we pointed out the difference between that of L1448 IRS 3B and those of the other two." + Based on differcut properties of LILLs IRS 3B from the other two sources. we sugeested that “binary system” YSOs and/or vounger YSOs in terms of kinematic fine scales of their bipolar outflows would have steeper density distributions.," Based on different properties of L1448 IRS 3B from the other two sources, we suggested that “binary system” YSOs and/or younger YSOs in terms of kinematic time scales of their bipolar outflows would have steeper density distributions." + 3., 3. + We found radial dependences of 3., We found radial dependences of $\beta$. + Iu particular. the dependence is distinct in Litls IRS 3B. We verified. it bv models eiiploxiug ο) as a function of radius.," In particular, the dependence is distinct in L1448 IRS 3B. We verified it by models employing $\beta$ as a function of radius." +" Iu addition. we discussed that the erain growth causing the depenudeuce can be achieved iu a time scale of 101 years. corresponding to the kinematic time scale of bipolar outflows of Class 0 YSOs,"," In addition, we discussed that the grain growth causing the dependence can be achieved in a time scale of $10^4$ years, corresponding to the kinematic time scale of bipolar outflows of Class 0 YSOs." + First of all we thauk the CARMA staffs for their dedicated work to commission and operate CARALA.," First of all, we thank the CARMA staffs for their dedicated work to commission and operate CARMA." + W. Ik. thanks M. W. Kunz. T. Ch.," W. K. thanks M. W. Kunz, T. Ch." + Mouschovias. aud C. EF. Camuiue for helpful discussions and coments.," Mouschovias, and C. F. Gammie for helpful discussions and comments." + lu addition. we thank anouviuous referee for valuable conuuents to improve this paper.," In addition, we thank anonymous referee for valuable comments to improve this paper." + W. Kk. aud L. W. L. ackuowledge support from NASA Ovigins Craut No., W. K. and L. W. L. acknowledge support from NASA Origins Grant No. + NNCOGCELCG. L.C. M. acknowledges support frou NASA Origius Caaut No., NNG06GE41G. L. G. M. acknowledges support from NASA Origins Grant No. + NNCGOGQGELG6C. Support for CARMA construction was derived. from the states of Illinois. California. aud Marvlaud. the Cordon aud Betty Moore Foundation. the Eileen aud I&euneth Norris Foundation. the Caltech Associates. and the National Science Foundation.," NNG06GE16G. Support for CARMA construction was derived from the states of Illinois, California, and Maryland, the Gordon and Betty Moore Foundation, the Eileen and Kenneth Norris Foundation, the Caltech Associates, and the National Science Foundation." + Ougoine CARMA development and operations are supported bv the National Science Foundation under cooperative agreement. AST-05LO159. and by the CABRMA partuer universities.," Ongoing CARMA development and operations are supported by the National Science Foundation under cooperative agreement AST-0540459, and by the CARMA partner universities." + Facilities:, Facilities: +Fig.,Fig. + E. gives a clear overview of the dominant large-scale structures in the Great Attractor region., \ref{lssoverview} gives a clear overview of the dominant large-scale structures in the Great Attractor region. + Support for the prominence of the Norma SCL has come from various complementary multiwavelength studies at lower Galactic latitudes such as the detection of several further clusters embedded. in the Norma SCL., Support for the prominence of the Norma SCL has come from various complementary multiwavelength studies at lower Galactic latitudes such as the detection of several further clusters embedded in the Norma SCL. + An X-ray search lor hiehly obscured. clusters in the ZOA (Ebcline. Alullis Tully 2002) revealec the second. most massive cluster in the Norma SCL. namely CIZAJJ1324.7.5736.," An X-ray search for highly obscured clusters in the ZOA (Ebeling, Mullis Tully 2002) revealed the second most massive cluster in the Norma SCL, namely J1324.7–5736." + This cluster is ~50% less massive than the Norma cluster. (Raclburn-Smith et al., This cluster is $\sim$ less massive than the Norma cluster (Radburn-Smith et al. + 2006) and is located a (b.e)~(307.47.15.07.5700 kms 13.," 2006) and is located at $(\ell, b, v) \sim (307.4^{\circ}, +5.0^{\circ}, 5700$ km $^{-1}$ )." + Deep near-infrare observations (Nagavama et al., Deep near-infrared observations (Nagayama et al. + 2004) furthermore uncoverec à low-mass cluster around 11343-601. at. (6.6.0)~(3090.77.|1.77.3900 km 1). also within the Norma SCL.," 2004) furthermore uncovered a low-mass cluster around 1343-601 at $(\ell, b, v) \sim (309.7^{\circ}, +1.7^{\circ}, 3900$ km $^{-1}$ ), also within the Norma SCL." + Apart [rom this significant> collection of clusters. a ὃνgenera overdensity along the Norma SCL is also clearly presen in the Parkes deep ΓΗ multibeam ZOA survey ]xorteweg et al.," Apart from this significant collection of clusters, a general overdensity along the Norma SCL is also clearly present in the Parkes deep I multibeam ZOA survey (Kraan-Korteweg et al." + 2005)., 2005). + In the first of a series of papers investigating the Norma cluster. we present a detailed. dynamical analysis of this cluster. the most massive cluster in the Great Attractor overdensity. centrally located in a cosmic web of filaments ane wall-like structures.," In the first of a series of papers investigating the Norma cluster, we present a detailed dynamical analysis of this cluster, the most massive cluster in the Great Attractor overdensity, centrally located in a cosmic web of filaments and wall-like structures." + Figure 2. shows the distribution of the optically-detected: galaxies (Woudt WKraan-hortewee 2001) in the eencral direction. of the Norma cluster. where the bell radius (24— U7/2) of the Norma cluster is indicated. by the dashed. circle.," Figure \ref{distribution} shows the distribution of the optically-detected galaxies (Woudt Kraan-Korteweg 2001) in the general direction of the Norma cluster, where the Abell radius $R_A \equiv 1\farcm7/z$ ) of the Norma cluster is indicated by the dashed circle." + At the redshift of the Norma cluster (see Sect., At the redshift of the Norma cluster (see Sect. + 3). the Abell radius corresponds to an angular radius of 1.757.," 3), the Abell radius corresponds to an angular radius of $^{\circ}$." + Assuming a Hubble constant of Ly=τὸ km s.+ + and the cosmological concordance model (assumed throughout this paper). the Abell radius corresponds to a physical size of 2.0 Alpe (the cosmology-correctecd angular scale at this distance is 1.16. Alpe per degree).," Assuming a Hubble constant of $H_0 = 73$ km $^{-1}$ $^{-1}$ and the cosmological concordance model (assumed throughout this paper), the Abell radius corresponds to a physical size of 2.0 Mpc (the cosmology-corrected angular scale at this distance is 1.16 Mpc per degree)." + Contours of equal. Galactic foreground. extinction. taken from. the DIRBE/URAS Galactic reddening map (Schlegel. Finkbeiner Davis 1998). are overlaved on the galaxy distribution in lig.," Contours of equal Galactic foreground extinction, taken from the DIRBE/IRAS Galactic reddening map (Schlegel, Finkbeiner Davis 1998), are overlayed on the galaxy distribution in Fig." + 2 (ιν = 1.0. 1.5 and 3.0 mag. respectively: Carcdelli. Clavton Mathis 1989). and show that. the Galactic foreground extinction within the Abell radius of the Norma cluster is mocerate. elyx1.5 mag.," \ref{distribution} $A_{\rm B}$ = 1.0, 1.5 and 3.0 mag, respectively; Cardelli, Clayton Mathis 1989) and show that the Galactic foreground extinction within the Abell radius of the Norma cluster is moderate, $A_{\rm B} \le 1.5$ mag." + Within the Abell radius. there are 603. optically-detected. galaxies with observed diameters in excess of 12” (Woudt Ixraan-Ixorteweg 2001) ancl 219 (near-inlrareel-detected) galaxies in the extended source catalogue (NSC) of the 2 Micron Al-Sky Survey (24LASS. Skrutskie et al.," Within the Abell radius, there are 603 optically-detected galaxies with observed diameters in excess of $''$ (Woudt Kraan-Korteweg 2001) and 219 (near-infrared-detected) galaxies in the extended source catalogue (XSC) of the 2 Micron All-Sky Survey (2MASS, Skrutskie et al." + 2006)., 2006). + The 2ALASS galaxies in the Norma cluster represent a subset of the 603 opticallv-detected. galaxies. although not all the PALASS galaxies have an optical counterpart: 165 of the 219 2PALASS galaxies (75%)) were also found by Woudt Ixraan-Ixorteweg (2001).," The 2MASS galaxies in the Norma cluster represent a subset of the 603 optically-detected galaxies, although not all the 2MASS galaxies have an optical counterpart; 165 of the 219 2MASS galaxies ) were also found by Woudt Kraan-Korteweg (2001)." +" For the brighter 2ALASS galaxies (10""-aperture. dvy-band « 12.5 mag). the overlap between 2NLASS and the optical survey is excellent: of the 2NLASS galaxies have an optical counterpart."," For the brighter 2MASS galaxies $''$ -aperture $K_s$ -band $<$ 12.5 mag), the overlap between 2MASS and the optical survey is excellent: of the 2MASS galaxies have an optical counterpart." + Lt should be noted that at the position of the Norma cluster (f£.b.~3257. T) star-crowding is the primary limiting factor. not the Galactic foreground extinction.," It should be noted that at the position of the Norma cluster $\ell, b \sim 325^{\circ}, -7^{\circ}$ ) star-crowding is the primary limiting factor, not the Galactic foreground extinction." + “Phe star-crowding leaves a Zone of Avoidance imprint on the 2ALASS NSC catalogue near the Galactic Bulee (Ixraan-Ixorteweg Jarrett 2005) and the Norma cluster is located on the edge of this Zone of Avoidance., The star-crowding leaves a Zone of Avoidance imprint on the 2MASS XSC catalogue near the Galactic Bulge (Kraan-Korteweg Jarrett 2005) and the Norma cluster is located on the edge of this Zone of Avoidance. + At moderate extinction (ip<3 mag). but in the presence of severe star-crowcding. optical surveys still retrieve the most complete galaxy distribution in the Zone of Xvoidance (Ixraan-Ixorteweg Jarrett 2005).," At moderate extinction $A_B \le 3$ mag), but in the presence of severe star-crowding, optical surveys still retrieve the most complete galaxy distribution in the Zone of Avoidance (Kraan-Korteweg Jarrett 2005)." + We have obtainted 129 new redshifts of galaxies within the Abell radius. of the Norma cluster using. the 2dl spectrograph at the Anelo-Australian Observatory., We have obtainted 129 new redshifts of galaxies within the Abell radius of the Norma cluster using the 2dF spectrograph at the Anglo-Australian Observatory. + These new observations are presented in Section 2., These new observations are presented in Section 2. + In. Section 3 all the redshifts obtained. to date are combined. and a detailed: dynamical analysis of the cluster. based. on 296 cluster members is presented., In Section 3 all the redshifts obtained to date are combined and a detailed dynamical analysis of the cluster based on 296 cluster members is presented. + La Section 4. we cliscuss a few individual galaxies in the Norma cluster of dynamical interest.," In Section 4, we discuss a few individual galaxies in the Norma cluster of dynamical interest." + Spectra were obtained. with the 2dE facility. (Lewis. et al., Spectra were obtained with the 2dF facility (Lewis et al. + 2002) on the 3.9m Anelo-Australian Telescope., 2002) on the 3.9m Anglo-Australian Telescope. + Full details of the observing 2dE setup used for observations are given in Table 1..., Full details of the observing 2dF setup used for observations are given in Table \ref{2dfsetup}. + As the main objective was to measure the velocity disperisons of the eluster's early-type galaxies the 1200V eratings were used in cach of the 2dE spectrographs., As the main objective was to measure the velocity disperisons of the cluster's early-type galaxies the 1200V gratings were used in each of the 2dF spectrographs. + These gave a EWLAL resolution of ~ 125 km + at, These gave a FWHM resolution of $\sim$ 125 km $^{-1}$ at +The physical meaning of the particle creation process is the lollowing: a wave packet of negative frequencies incident on an electric field oriented along the z—direction will be partly reflected. by the electric Ποιά and partly transmitted to 2—x as a wave with positive frequencies.,The physical meaning of the particle creation process is the following: a wave packet of negative frequencies incident on an electric field oriented along the $z-$ direction will be partly reflected by the electric field and partly transmitted to $z\rightarrow \infty $ as a wave with positive frequencies. + This process is nothing else than tunnelling (Soffeletal.1952:Greinerelal.1985:WangandWong 1933).," This process is nothing else than tunnelling \citep{So82,Gr85,Wa88}." +. Ilowever. in the following we restrict the tunnelling probability to the transmission probability through the potential barrier but exclude any non-zero (ransimission probability above the potential barrier.," However, in the following we restrict the tunnelling probability to the transmission probability through the potential barrier but exclude any non-zero transmission probability above the potential barrier." +" In the WIXB approximation (he (ranusmission probability 7. can be approximately given as""TN |4ΤΙ−—=exp-4(9~lusκ.ας)exp(-2o ). where o=mh.|d2 and b. is the longitudinal momentum of the electron (SoffeletPage 2002)."," In the WKB approximation the transmission probability $\mathcal T$ can be approximately given as $% +\left| \mathcal T\right| ^{2}=\exp \left( -2\int_{barrier}\left| +k_{z}\right| dz\right) =\exp\left(-2\sigma \right ), where $\sigma += \int_{barrier}\left| k_{z}\right| dz$ and $k_z$ is the longitudinal momentum of the electron \citep{So82,Wa88,Kim02}." +. For a fixed frequency w. the momentum ff. of an electron in the electrosphere of the quark stars is given by where 4?=he+ho (Solfeletal.1982;NimandPage2002).," For a fixed frequency $\omega $, the momentum $k_{z}$ of an electron in the electrosphere of the quark stars is given by where $k_{\perp }^{2}=k_{x}^{2}+k_{y}^{2}$ \citep{So82,Kim02}." +. The mean particle production number (NV) is given by summing over all modes (Greineretal.1085).. In order to estimate (he electron-positron pair production rate in (he electrosphere. we have to find first the thickness of the classically forbidden zone. or. equivalently. the limits of integration ο. in the transmission probability.," The mean particle production number $\left\langle N\right\rangle$ is given by summing over all modes \citep{Gr85}, In order to estimate the electron-positron pair production rate in the electrosphere, we have to find first the thickness of the classically forbidden zone, or, equivalently, the limits of integration $z_{\pm }$ in the transmission probability." + They. can be obtained as solutions of the equation A.(2)=0 and are given by Therefore. bv taking into account the form of the electrostatic potential in the electrosphere we obtain lor σ the expression With the help of the transformation η=(uvan+i)ζω—V?zT/sinhEVax/31(225)\J," They can be obtained as solutions of the equation $k_{z}\left( z\right) =0$ and are given by Therefore, by taking into account the form of the electrostatic potential in the electrosphere we obtain for $\sigma $ the expression With the help of the transformation $\eta =\left( +1/\sqrt{m_{e}^{2}+k_{\perp +}^{2}}\right) \left\{ \omega -\sqrt{2}\pi T/\sinh \left[ 2\sqrt{\alpha \pi /3% +}T\left( z+z_{0}\right) \right] \right\}" +aaxd irafios between the low aud high SER subsamples of the iutOracting svstenis.,and ratios between the low and high SFR subsamples of the interacting systems. + Similarly. a factor of 2 reduction in Που the active interacting class 5 systems would also acc“OILut for inost of the crease in aaxd rratios frou class 1 to class 5Γ svsteuas.," Similarly, a factor of 2 reduction in for the active interacting class 5 systems would also account for most of the increase in and ratios from class 1 to class 5 systems." + ILowever. changes in the conversion factor of this order mot account for the whole range of fractional CCCmutent seen within the eutire isolated anc interacting eaaxy salples.," However, changes in the conversion factor of this order cannot account for the whole range of fractional content seen within the entire isolated and interacting galaxy samples." + The aadxd rratios for individual systems cover a range of a factor of 10 !u1 both quantities, The and ratios for individual systems cover a range of a factor of 100 in both quantities. + Therefore a change iu the conversion factor on the order of 2 over the range of | souples cannot account for the stroug correlation seen ween fractional ccoutent aud SFR. as iicasured by id irafios (soc. e.g. Figs.," Therefore a change in the conversion factor on the order of 2 over the range of the samples cannot account for the strong correlation seen between fractional content and SFR, as measured by and ratios (see, e.g., Figs." + 7 and 9)), \ref{fig_irb_h2b} and \ref{fig_hab_h2b}) ). + Therefore even if the aLeut increase in the ccoutent of the interacting axiesfor the sample as a whole relative to the isolated spirals aud as a function of SER and interaction strenethis nof real. there still is a correlation between the ccontent aud the current SFR in individual galaxies. for both the isolated aud interacting svstenis.," Therefore even if the apparent increase in the content of the interacting galaxies—for the sample as a whole relative to the isolated spirals and as a function of SFR and interaction strength---is not real, there still is a correlation between the content and the current SFR in individual galaxies, for both the isolated and interacting systems." + We lave analyzed CO (10) emission-line observatious for a sauple of 37 interacting galaxy svstenis., We have analyzed CO (1–0) emission-line observations for a sample of 37 interacting galaxy systems. + The parcut saluple of interacting svstenis was compiled ou the basis of morphological evidence for current involvement in an interaction. without bias towards kuown optical. infrared. or radio flux levels or indicators of high current SERs or nuclear activity.," The parent sample of interacting systems was compiled on the basis of morphological evidence for current involvement in an interaction, without bias towards known optical, infrared, or radio flux levels or indicators of high current SFR's or nuclear activity." + The svstenis in this study were selected from the parent sample in order to span a laree rauge of interaction strengths aud star formation properties., The systems in this study were selected from the parent sample in order to span a large range of interaction strengths and star formation properties. + If we asunue that the ccouversion factor ds approximately coustant from galaxy to galaxy within our saluples of interacting and isolated galaxies. our results indicate the following.," If we assume that the conversion factor is approximately constant from galaxy to galaxy within our samples of interacting and isolated galaxies, our results indicate the following." + There are dudicatious. however. that the cconversion factor may not be constaut either within individual galaxies or from galaxy to galaxy.," There are indications, however, that the conversion factor may not be constant either within individual galaxies or from galaxy to galaxy." + Furthermore. the level of observed CO emission itself may be depenudeut ou the level of star formation activity within a galaxx. which will teud to increase the “visibility” of molecular clouds.," Furthermore, the level of observed CO emission itself may be dependent on the level of star formation activity within a galaxy, which will tend to increase the “visibility” of molecular clouds." + Reductious in the ccouversion factor that have been sugeested for some ultraluniuous infrared ealaxies would be amore than cuough to nullifv the average increases in derived ccoutent for the various subsamples of interacting galaxies in this study., Reductions in the conversion factor that have been suggested for some ultraluminous infrared galaxies would be more than enough to nullify the average increases in derived content for the various subsamples of interacting galaxies in this study. + Thev cannot. however. account for the total range of cconteut seen within the cutive samples of isolated aud interacting galaxies.," They cannot, however, account for the total range of content seen within the entire samples of isolated and interacting galaxies." + Therefore the observed correlation, Therefore the observed correlation +Because we have observed both ΠΟ 199113 and IID 358623 several times (Dec. 2000. July 2001. Dec. 2001). we can check. whether the companion candidates share the oxoper motion with their respective primarics.,"Because we have observed both HD 199143 and HD 358623 several times (Dec. 2000, July 2001, Dec. 2001), we can check, whether the companion candidates share the proper motion with their respective primaries." + Iu the case of TID 3586223. we have detected the companion candidate at all three epochs. while iu the case of ΠΟ 199113. he companion candidate is only marginallv resolved iu he two Soff images.," In the case of HD 358623, we have detected the companion candidate at all three epochs, while in the case of HD 199143, the companion candidate is only marginally resolved in the two SofI images." + We could compare our separations vetween ΠΟ 199113 A and D with those given in JBOL. but hat does not vield any significant conclusion. because the JBOL observations were obtained only one mouth before our SUARP-T images.," We could compare our separations between HD 199143 A and B with those given in JB01, but that does not yield any significant conclusion, because the JB01 observations were obtained only one month before our SHARP-I images." + The separations are nieasured as differences. between the photo-ceuters of the objects determined by a Gaussian centering withΗΕ?.. with the pixel scales uentioned above.," The separations are measured as differences between the photo-centers of the objects determined by a Gaussian centering with, with the pixel scales mentioned above." + The north-south aliguiieut of SIARP-I during our observing nights was determined to be precise within less than half a degree using nuages of the galactic center and known bound binarics taken iu the same nights., The north-south alignment of SHARP-I during our observing nights was determined to be precise within less than half a degree using images of the galactic center and known bound binaries taken in the same nights. + The separations in a and à as well as the position angles PA eiwen in Table 3H include an additional error of Ε.Τ due to an uncertainty in the north-south aliguimenut of the detectors. while the errors in total separation are independant of that uncertaimty.," The separations in $\alpha$ and $\delta$ as well as the position angles PA given in Table 3 include an additional error of $\pm 0.5^{\circ}$ due to an uncertainty in the north-south alignment of the detectors, while the errors in total separation are independant of that uncertainty." + Fist. we lave measured 18 other comparison stars iu the ΠΟ 358623 feld. iu order to determine the precison of our astrometiy.," First, we have measured 18 other comparison stars in the HD 358623 field, in order to determine the precison of our astrometry." + The separations between cach pair of conarisou stars did not change within the LO to 20 mas errors., The separations between each pair of comparison stars did not change within the 10 to 20 mas errors. + Heuce. they are very likely nounauoviug backeround stars.," Hence, they are very likely non-moving background stars." + IID 358623 moves relative to the comparison stars o the SE., HD 358623 moves relative to the comparison stars to the SE. + From the differeuce in the separations betwee- LS comparison stars and TD 358623. we cau measure itpA xoper notion (103 mas/vr). which is consisteut wit[um its known Tycho proper motion (86 imas/vr) withi- ess than 20 mas/yr.," From the difference in the separations between 18 comparison stars and HD 358623, we can measure its proper motion $\sim 103$ mas/yr), which is consistent with its known Tycho proper motion $\sim 86$ mas/yr) within less than 20 mas/yr." + This shows how precise relative astrometry can be. even after only oue wear of epoc[um difference. due to (4) the simall Soff pixels together wit[um he small FWIIM due to excellent secing and active optics (lence αμα. error im Caussian ceuterins). (i) he relatively large Sofl ficld-ofview even iu the so-called field (houce niu. suticicutly bright coluparison stars). (d) and the stability of the telescope and iustiruueut optics over the oue vear epoch difference.," This shows how precise relative astrometry can be, even after only one year of epoch difference, due to (i) the small SofI pixels together with the small FWHM due to excellent seeing and active optics (hence small error in Gaussian centering), (ii) the relatively large SofI field-of-view even in the so-called (hence many sufficiently bright comparison stars), (iii) and the stability of the telescope and instrument optics over the one year epoch difference." + If the close conmrpanion candidate would be au related backeround object. it should also move relative o IID 358623 as the other very widely separated vackevound objects. namely by almost ~LOO mas/vr according to the PA.," If the close companion candidate would be an unrelated background object, it should also move relative to HD 358623 as the other very widely separated background objects, namely by almost $\sim 100$ mas/yr according to the PA." + The separations measured by us aud x JBOL are given in Table 3., The separations measured by us and by JB01 are given in Table 3. + One can clearly see that the companion candidate TD 358623 D did not move relative ΠΟ 358623 A (see Fig., One can clearly see that the companion candidate HD 358623 B did not move relative HD 358623 A (see Fig. + 3) from one epoch to aux of the hree other epochs. no within the measurement errors and certainly uot with a motion as large as ~100 mas/vr.," 3) from one epoch to any of the three other epochs, not within the measurement errors and certainly not with a motion as large as $\sim 100$ mas/yr." + Iu particular. the comparison between the Dec. 2000 Sofl iuaege lniageandaix the| Dec. Dec.2001 Sofl nuage nuageis very significant.siguificaut οσα they have been taken with the same iustruueut at the same airmass and also with exactly a full vear epoch difference. so that there are no ciffereutial parallax tor refraction effects;," In particular, the comparison between the Dec. 2000 SofI image and the Dec. 2001 SofI image is very significant, because they have been taken with the same instrument at the same airmass and also with exactly a full year epoch difference, so that there are no differential parallax nor refraction effects." + Because the companion candidate is located SE of the primary and because the primary is shown to move SE. the separation should have decreased No100 mas. if the pair would not be bouud.," Because the companion candidate is located SE of the primary and because the primary is known to move SE, the separation should have decreased by $\sim 100$ mas, if the pair would not be bound." + This cau be excluded with a significance of 67., This can be excluded with a significance of $6 \sigma$. + Deuce. the companion candidate is co-moving with the star.," Hence, the companion candidate is co-moving with the star." + Therefore. it is very ikelv a trucly bouud companion of the vouus nearby star IID 358623.," Therefore, it is very likely a truely bound companion of the young nearby star HD 358623." + To coufiru this conclusion. we have taken a ρουσπιν of IID 358623 D.," To confirm this conclusion, we have taken a spectrum of HD 358623 B." +different criteria.,different criteria. + Observational astronomers classify small bodies having transient. wnbotnel abmospheres (usually mace visible by the scattering of sunlight from entrained micron-sized dust particles) as “comets”.," Observational astronomers classify small bodies having transient, unbound atmospheres (usually made visible by the scattering of sunlight from entrained micron-sized dust particles) as “comets”." + Dodies having imstead a constant. geometric cross-section are called “asteroids”., Bodies having instead a constant geometric cross-section are called “asteroids”. + To planetary scientists. comets and asteroids are distinguished by (heir ice content or perhaps by their formation location.," To planetary scientists, comets and asteroids are distinguished by their ice content or perhaps by their formation location." +" Comets are icy (because (μον formed bevond the ""snow-line') while asteroids are not (supposedly because thev formed at higher mean temperatures inside 1).", Comets are icy (because they formed beyond the “snow-line”) while asteroids are not (supposedly because they formed at higher mean temperatures inside it). + Lastlv. to dvnamücisis. comets and asteroids are broadly distinguished by a dynamical parameter. most usually Che Tisserand parameter measured with respect to Jupiter (INresak 1982. INosai 1992).," Lastly, to dynamicists, comets and asteroids are broadly distinguished by a dynamical parameter, most usually the Tisserand parameter measured with respect to Jupiter (Kresak 1982, Kosai 1992)." +" It is defined by where a. e and 7 are the semimajor axis. eccentricity and inclination of the orbit while a, = 5.2 AU is the semimajor axis of the orbit of Jupiter."," It is defined by where $a$ , $e$ and $i$ are the semimajor axis, eccentricity and inclination of the orbit while $a_J$ = 5.2 AU is the semimajor axis of the orbit of Jupiter." + This parameter. which is conserved in the circular. restricted 3-bodsy problem. provides a measure of the close-approach speed to Jupiter.," This parameter, which is conserved in the circular, restricted 3-body problem, provides a measure of the close-approach speed to Jupiter." + Jupiter itself has 77; = 3., Jupiter itself has $T_J$ = 3. +" Main belt asteroids have α 3 while dynamical comets have Ty<3.", Main belt asteroids have $a \le a_J$ and $T_J >$ 3 while dynamical comets have $T_J < 3$. + The three svstems of classification (observational. compositional aud dynamical) are independent but imperlect.," The three systems of classification (observational, compositional and dynamical) are independent but imperfect." + For example. whether a coma or tail is detected on a given object depends strongly on the parameters of the observing svstem used.," For example, whether a coma or tail is detected on a given object depends strongly on the parameters of the observing system used." + A puny telescope nav not reveal a coma (hat is easily rendered visible by a more powerlul one., A puny telescope may not reveal a coma that is easily rendered visible by a more powerful one. + There is either an agreed quantitative ice Traction with which to divide comets [rom asteroids nor. nore importantly. any reliable way to measure the ice fraction in a small body.," There is neither an agreed quantitative ice fraction with which to divide comets from asteroids nor, more importantly, any reliable way to measure the ice fraction in a small body." + The formation ocalions and dvnanmical histories of small bodies are rendered uncertain by. haid-to-mocdel 1on-gravitational forces [rom both electromagnetic radiation (he Yarkovsky effect) and mass oss Itself (the rocket effect of Whipple 1950) and also by (the chaotic nature of solar system denamics., The formation locations and dynamical histories of small bodies are rendered uncertain by hard-to-model non-gravitational forces from both electromagnetic radiation (the Yarkovsky effect) and mass loss itself (the rocket effect of Whipple 1950) and also by the chaotic nature of solar system dynamics. + The Tisserand criterion is an imperfect classifier because it is based on an idealized representation of the Solar svstem (e.g. Jupiters orbit is not a circle. (he eravily of other planets is not entirely negligible. and so on).," The Tisserand criterion is an imperfect classifier because it is based on an idealized representation of the Solar system (e.g. Jupiter's orbit is not a circle, the gravity of other planets is not entirely negligible, and so on)." + Therefore. the utility of Equation (1)) as a dvnamical discriminant is limited lor objects with 7; verv close to 3.," Therefore, the utility of Equation \ref{tisserand}) ) as a dynamical discriminant is limited for objects with $T_J$ very close to 3." + As an example. the quasi-llilda comets at «~ 4.0 AU have Ty 2.9 - 3.04 but are clearly interacting with Jupiter through the 3:2 mean-motion resonance (Toth 2006).," As an example, the quasi-Hilda comets at $a \sim$ 4.0 AU have $T_J \sim$ 2.9 - 3.04 but are clearly interacting with Jupiter through the 3:2 mean-motion resonance (Toth 2006)." +" Some recognized Jupiter [amily comets (e.g. 2P/Encke with 7, = 3.03) also fall in this category. (Fernandez οἱ al.", Some recognized Jupiter family comets (e.g. 2P/Encke with $T_J$ = 3.03) also fall in this category (Fernandez et al. + 2002. Levison et al.," 2002, Levison et al." + 2006)., 2006). + Given these and other imperfections it is remarkable that. for a majority of objects. theobservational. compositional and. dynamical definitions of asteroids and comets lie in," Given these and other imperfections it is remarkable that, for a majority of objects, theobservational, compositional and dynamical definitions of asteroids and comets lie in" +value (0.93) is used for extended sources (objects with C>2.3) in each filter.,value (0.98) is used for extended sources (objects with $C > 2.3$ ) in each filter. +" These mean aperture corrections are determined from the magnitude difference between apertures of 3 and 12.5 pixel (1... 0.5"") for a sample of zz 50 relatively isolated. high S/N stars ancl clusters."," These mean aperture corrections are determined from the magnitude difference between apertures of 3 and 12.5 pixel (i.e., $0.5^{\prime\prime}$ ) for a sample of $\approx$ 50 relatively isolated, high S/N stars and clusters." +" We assume an additional 0.10 mag is needed to correct the photometry from 0.5"" to infinity. based on new measurements lor WFEC3."," We assume an additional 0.10 mag is needed to correct the photometry from $0.5^{\prime\prime}$ to infinity, based on new measurements for WFC3." + This assumption is good (o a lew percent aud is dominated by the aperture correction. in any case.," This assumption is good to a few percent and is dominated by the aperture correction, in any case." + A weakness of this approach is that bv adopting a single value of the aperture correction for all clusters. we overestimate the total Iuminosity of more compact clusters and underestimate the total luminosity of more extended clusters.," A weakness of this approach is that by adopting a single value of the aperture correction for all clusters, we overestimate the total luminosity of more compact clusters and underestimate the total luminosity of more extended clusters." + Our second approach applies aperture corrections (to each object based on its measured size (C)., Our second approach applies aperture corrections to each object based on its measured size $C$ ). + Unresolved objects (those with C<2.3) all receive a sinele. filter-depencdent aperture correction (0.30).," Unresolved objects (those with $C < 2.3$ ) all receive a single, filter-dependent aperture correction (0.30)." + For extended objects. we find the following relationship between aperture correction and C. lor the range 2.3 70/2. then we complete (he integration up to 75 and record the final eccentricity as ej=e(79). talàng into account the incomplete Ixozai cycle.," However, if the Kozai period is comparable to the assumed age of the system, with $P_{\rm KOZ}> \tau_0 /2$ , then we complete the integration up to $\tau_0$ and record the final eccentricity as $e_{\rm f}=e(\tau_0)$, taking into account the incomplete Kozai cycle." + Applving this method for each of the 5000 sample svstems. we Chen derive the cumulative probability distribution of ey.," Applying this method for each of the 5000 sample systems, we then derive the cumulative probability distribution of $e_{\rm f}$." + The results for representative models are presented. together with the observed cumulative distribution in refresult.., The results for representative models are presented together with the observed cumulative distribution in \\ref{result}. + For each model. we have plotted the final eccentricities in histograms with normalized probabiliies as well as cumulative distributions.," For each model, we have plotted the final eccentricities in histograms with normalized probabilities as well as cumulative distributions." + These are compared to the distribution derived [rom (the observed single planets with c4>0.1. from the California Carnegie Planet Search Catalogue.," These are compared to the distribution derived from the observed single planets with $a_1>0.1$, from the California Carnegie Planet Search Catalogue." + In all the models. a signilicant fraction of planets have failed [or various reasons to achieve high eecentricitv.," In all the models, a significant fraction of planets have failed for various reasons to achieve high eccentricity." + The analvsis of the svstems retaining a low final eccenlricily is presented in Table 2.., The analysis of the systems retaining a low final eccentricity is presented in Table \ref{lowecc}. + The [ist four models have iniGial parameter distributions that (8) ave compatible with our current knowledge of stellar and substellar binary companions. and Gi) can produce the closest result to the observed eccentricity distribution of extrasolar planets.," The first four models have initial parameter distributions that (i) are compatible with our current knowledge of stellar and substellar binary companions, and (ii) can produce the closest result to the observed eccentricity distribution of extrasolar planets." + The results are shown in Figure 3.., The results are shown in Figure \ref{BDratio}. + Each of the four different models represents 5000 sample systems wilh a different assumed ratio of brown dwarf companions to FGIX dwarf companions., Each of the four different models represents 5000 sample systems with a different assumed ratio of brown dwarf companions to FGK dwarf companions. + Although the differences between (hese models are rather small. the results show that a higher fraction of brown chvarl companions leads to more planetswith low eccentricities. as expected.," Although the differences between these models are rather small, the results show that a higher fraction of brown dwarf companions leads to more planetswith low eccentricities, as expected." + All the, All the +"have been followed for more than a mouth bySwift. because usually the X-ray afterglow fades below the AN-Rax telescope (XRT:Durrowsetal.2005) detection lunitLE ~10515 eres ται ? roughly a week after the trigger,","have been followed for more than a month by, because usually the X-ray afterglow fades below the X-Ray telescope \citep[XRT; ][]{burrows05} detection limit $\sim 10^{-14}$ ergs $^{-1}$ $^{-2}$ roughly a week after the trigger." + One of these exceptions is the bright ταν iterelow of CRB 060729 which was detected by NART even 125 davs after the burst (Capeetal.2007).., One of these exceptions is the bright X-ray afterglow of GRB 060729 which was detected by XRT even 125 days after the burst \citep[][]{grupe07}. . + The Burst Alert Telescope (BAT:Baxtheliuny2005) trigeered on GRB060720 on 2006 July 29. 19:12:29 UT (Gaupectal.2007) and à).a redshift of >=(51 was measured (Thocneetal.2006," The Burst Alert Telescope \citep[BAT; ][]{barthelmy05} + triggered on GRB060729 on 2006 July 29, 19:12:29 UT \citep{grupe07} and a redshift of $z=0.54$ was measured \citep{thoene06}." + The XRT and the UV/Optical Telescope (UVOT:Rominectal. 2005).. started observing the burst about 2 ninuutes after the trigecr.," The XRT and the UV/Optical Telescope \citep[UVOT; ][]{roming05}, started observing the burst about 2 minutes after the trigger." + The UUVOT was able to followt» afterelow in the UVWT filter up to 31 days after the BAT tieeer., The UVOT was able to follow this afterglow in the UVW1 filter up to 31 days after the BAT trigger. + Ta X-rays Suwiffss NRT was still detectine the N-rav afterglow of CRB 060729 at the eud of 2006 November. 125 cays after the burst (προctal.2007).," In X-rays s XRT was still detecting the X-ray afterglow of GRB 060729 at the end of 2006 November, 125 days after the burst \citep{grupe07}." +. However. by 2006 December the Susft--NRT. detection lait was reached and only 36 upper limit couldbe eiven for the 63.5 ks exposure timea obtained in December 2006.," However, by 2006 December the -XRT detection limit was reached and only a $\sigma$ upper limit could be given for the 63.5 ks exposure time obtained in December 2006." + By that time the X-ray afterglow of GRB 060729 did not show any clear evidence for a jetbreak. giving a lower lait on the jet opening angle of 0=28° (Caupeetal.2007). based ou the asstmption of a coustaut cireuuburst medi.," By that time the X-ray afterglow of GRB 060729 did not show any clear evidence for a jet break, giving a lower limit on the jet opening angle of $\theta=28^{\circ}$ \citep{grupe07} based on the assumption of a constant circumburst medium." + Iu order to extend the elt curve of this exceptional X-ray afterglow. we observed it five times with AACTS in 2007 and 2008.," In order to extend the light curve of this exceptional X-ray afterglow, we observed it five times with ACIS in 2007 and 2008." + We report on the detections of the N-rav. afterglow of CRB 060729 (Carpeetal.2007) with uup to 612 davs after the burst - the latest detection ever of au X-ray afterglow of a CRB at cosmological distance., We report on the detections of the X-ray afterglow of GRB 060729 \citep{grupe07} with up to 642 days after the burst - the latest detection ever of an X-ray afterglow of a GRB at cosmological distance. + Previously the burst with the latest detection of an X-ray afterglow was GRD 030329 (Ticuegoetal.2003.2001). which hac a detection258 davs after the burst byN," Previously the burst with the latest detection of an X-ray afterglow was GRB 030329 \citep{tiengo03, +tiengo04}, which had a detection 258 days after the burst by." +"ewton, Our paper is organized as follows: in the observations and data reduction are explained. iu— rofresults the measurements ofthe X-ray light curve are shown. and iu refdiscuss woe discuss the nmuplieations of this μον curve."," Our paper is organized as follows: in \\ref{observe} the observations and data reduction are explained, in \\ref{results} the measurements of the X-ray light curve are shown, and in \\ref{discuss} we discuss the implications of this light curve." +" Throughout this paper the N-vayv flux dependence on time and frequency is defined as Εκtov 8, Launinosities ave caleulated assuming a ACDAL cosmologv with O4; = 0.30. O4 = 0.70 and a Hubble coustant of Z7, = llus ! | corresponding to a huninosity distance dp = 3061 Mpc for CRB 060729."," Throughout this paper the X-ray flux dependence on time and frequency is defined as $F\propto +t^{-\alpha}\nu^{-\beta}$ Luminosities are calculated assuming a $\Lambda$ CDM cosmology with $\Omega_{\rm M}$ = 0.30, $\Omega_{\Lambda}$ = 0.70 and a Hubble constant of $H_0$ = 71 km $^{-1}$ $^{-1}$ corresponding to a luminosity distance $d_L$ = 3064 Mpc for GRB 060729." + oobserved CRB 060729 three times between 2007 Marchi 16 aud 2007 June 30., observed GRB 060729 three times between 2007 March 16 and 2007 June 30. + Two very late-time observations were performed in 2007 December/2008 Jaunary for 72.7 ks andiu 2008 ÁApzil/Maw for 117.3 ks., Two very late-time observations were performed in 2007 December/2008 January for 72.7 ks and in 2008 April/May for 117.3 ks. + Due to pitch angle coustraiuts some of these had tobe split iuto several visits., Due to pitch angle constraints some of these had to be split into several visits. + All observations. with start and end times and exposure times. are listed in rofxravjog..," All observations, with start and end times and exposure times, are listed in \\ref{xray_log}." + All of these observations were performed with the standardposition 3.28 readout time in Very Faint mode on the On-axis on the back-ilhuninated ACTS-83 CCD., All of these observations were performed with the standard 3.2s readout time in Very Faint mode on the on-axis position on the back-illuminated ACIS-S3 CCD. + Data reduction was performed with the aanalysis software CIAO version L0 aud the calibration database CALDD version 3.1.3., Data reduction was performed with the analysis software CIAO version 4.0 and the calibration database CALDB version 3.4.3. + Tn order to reduce the ACTS particle background all sstage Ll event data were reprocessed usine CIAO with the VF mode cleaning., In order to reduce the ACIS particle background all stage 1 event data were reprocessed using CIAO with the VF mode cleaning. + Ouly ACTS erades 0. 2. 3. L. aud 6 were selected. for further analvsis.," Only ACIS grades 0, 2, 3, 4, and 6 were selected for further analysis." + The background was further reduced by using ouly photons iu the 0.5 - 8.0 keV enerev range., The background was further reduced by using only photons in the 0.5 - 8.0 keV energy range. + Before further analysis. the observations were combined to one eveut file cach using the CTAO taska," Before further analysis, the observations were combined into one event file each using the CIAO task." +"ll, Source photons were selected i a circle with a radius Ξ- 1 and background photons in a close-by source-free region with a radius r = 10.", Source photons were selected in a circle with a radius $r$ = $^{''}$ and background photons in a close-by source-free region with a radius $r$ = $^{''}$. + Count rates were converted into fluxes by using PIMMS veπο” 3.9) using the parameters from the spectral fits to the ddata after the break at 1 Ms atten the burst (see below) with an absorption cohunuu density Ng=1.31<1074 ceni7 and an N-ray spectral index Jg 0.89., Count rates were converted into fluxes by using PIMMS version 3.9b using the parameters from the spectral fits to the data after the break at 1 Ms after the burst (see below) with an absorption column density $N_{\rm H}=1.34 \times 10^{21}$ $^{-2}$ and an X-ray spectral index = 0.89. + A description of the reduction aud analysis of the ddata can be found in Capeetal.(2007)., A description of the reduction and analysis of the data can be found in \citet{grupe07}. +. For display purposes and fitting the late-time elt curve we rebiuned the NART Photon Counting data with 250 counts per bin for the times up to 2 Ms after the burst and 100 counts per bin for the times thereafter., For display purposes and fitting the late-time light curve we rebinned the XRT Photon Counting data with 250 counts per bin for the times up to 2 Ms after the burst and 100 counts per bin for the times thereafter. + Spectral analvses were performed for the times 300 - S00 ks and TZ2 1 Mx after the bist.," Spectral analyses were performed for the times 300 - 800 ks and $T +>$ 1 Ms after the burst." + Source photons were collected iu a circle with &=23/75 aud backeround photous with ro—96 with evade selection 0 - 12., Source photons were collected in a circle with $r = 23\farcs5$ and background photons with $r = 96^{''}$ with grade selection 0 - 12. + The response matrix20010100010. παν used., The response matrix was used. + The spectra were rebinned with 20 counts per bin for the 300-800. ks after the burst spectra aud 15 counts per bin for the spectimm with T> 1 Als. The spectra were analyzed with NSPEC version. 12.0x (Armand1996)., The spectra were rebinned with 20 counts per bin for the 300-800 ks after the burst spectrum and 15 counts per bin for the spectrum with $T >$ 1 Ms. The spectra were analyzed with XSPEC version 12.4.0x \citep{arnaud96}. +. To search for chauges in the N-rav spectruu we applied a hardness ratio segment by seement., To search for changes in the X-ray spectrum we applied a hardness ratio segment by segment. + Because of the low-unuboer statistics iu some of the later segments of the NNART and all oobservations. we applied Bayesian statistics to determine the harduess ratios as described by Park.al. (2006).," Because of the low-number statistics in some of the later segments of the XRT and all observations, we applied Bayesian statistics to determine the hardness ratios as described by \citet{park06}." + The late-time X-ray light curve incliding the five ppoiutiues is shown in retsravye., The late-time X-ray light curve including the five pointings is shown in \\ref{xray_lc}. + Weiquorethe firstdayofthe oobscrvcationbeeauscitisnotrclevant forthestudyo fthelate tinmelighteurce , We ignore the first day of the observation because it is not relevant for the study of the late-time light curve. +Phecearlylighteurceandadetailcddiscussioncanbe fo ,The early light curve and a detailed discussion can be found in \citet{grupe07}. . +The late-time light curve jas Fittedbgsevceralpowerlaweandimultiply brokenpowerlawimodclsaslistedinT its. , The late-time light curve \\ref{xray_lc}) ) was fitted by several power law and multiply-broken power law models as listed in \\ref{lc_fits}. . +Fittingthelighteurcewiththedecauystopeoa? , Fitting the light curvewith the decay slope $\alpha_3$ +write with- 0.5FoxàS«Lj1.,write with $0.5 \la \delta \la 1$. +" This suggests that the lower frequencies found in ULXs imav also be the result of jeher accretion rates, not only of higher masses, in aerecinent with our proposed interpretation of the spectral data."," This suggests that the lower frequencies found in ULXs may also be the result of higher accretion rates, not only of higher masses, in agreement with our proposed interpretation of the spectral data." + The different dependence of spectral and timing parameters ou jj and AL suggests that it is in principle possible to diseutanele the two effects., The different dependence of spectral and timing parameters on $\dot{m}$ and $M$ suggests that it is in principle possible to disentangle the two effects. +" Preliminary back-of-the-envelope calculations sugeest that accretion rates mee20 and masses AM~50. 100, are consistent witli both the spectral and timing data of typical ULXs.", Preliminary back-of-the-envelope calculations suggest that accretion rates $\dot{m} \approx 20$ and masses $M \sim 50$ $100 M_{\odot}$ are consistent with both the spectral and timing data of typical ULXs. + Iu addition to providing an independent coustraiut on the effects of accretion rate and BIT nass. the observed auticorrelation jetween a ddsk parameter (inner radius) and the frequency of oscillation of the power-law ciission component is also a fundamental clue to understand the origin of such oscillations.," In addition to providing an independent constraint on the effects of accretion rate and BH mass, the observed anticorrelation between a disk parameter (inner radius) and the frequency of oscillation of the power-law emission component is also a fundamental clue to understand the origin of such oscillations." + One can speculate that LF-QPOs originate at the interface of outer disk aud inner flow. at the transition radius £2.," One can speculate that LF-QPOs originate at the interface of outer disk and inner flow, at the transition radius $R_{\rm c}$." +" We will present a more detailed analvsis of the timing results, in the framework of our ultraluininous branch model. in Soria and I&uncie (2007)."," We will present a more detailed analysis of the timing results, in the framework of our ultraluminous branch model, in Soria and Kuncic (2007b)." +" It is often noted that ULNs aud stellaramass DIIS share N-rvay spectral aud timing properties typical of the Wuninous, power-law dominated accretion state (very Πο] state)."," It is often noted that ULXs and stellar-mass BHs share X-ray spectral and timing properties typical of the luminous, power-law dominated accretion state (very high state)." +" If we plot the spectral evolution of typical accreting DIIS in the (Zi.Lai) plane, the very high state is located on the right-hand-side (higher disk eniperatures) of the thermal track characteristic of he disk-dominated hieh/soft state."," If we plot the spectral evolution of typical accreting BHs in the $(T_{\rm in}, L_{\rm disk})$ plane, the very high state is located on the right-hand-side (higher disk temperatures) of the thermal track characteristic of the disk-dominated high/soft state." +" This would imply hat ULXs and stellar-1uass DIIs are entirely. separate species, with a two-order-ofanaecnitude eap in DII nasses between them (Gutermecdiate-mass BIT scenario)."," This would imply that ULXs and stellar-mass BHs are entirely separate species, with a two-order-of-magnitude gap in BH masses between them (intermediate-mass BH scenario)." +" However,RNTE spectral fits to the stellar-amass DIT ATE 561 in its 19981999 outburst (Sobezak et al."," However, spectral fits to the stellar-mass BH XTE $-$ 564 in its 1998–1999 outburst (Sobczak et al." +" 2000: I&ubota and Done 2001) clearly show that he very high state can in fact be divided into two sub-states, located to the right (ligher teniperature) and to the left (lower temperature) of the thermal rack, respectively."," 2000; Kubota and Done 2004) clearly show that the very high state can in fact be divided into two sub-states, located to the right (higher temperature) and to the left (lower temperature) of the thermal track, respectively." + In the latter state (ultraliuninous xanch). the fitted inner-disk radius is much larger han the innermost stable circular orbit. and. the peak colour temperature much cooler than in the highκο state.," In the latter state (ultraluminous branch), the fitted inner-disk radius is much larger than the innermost stable circular orbit, and the peak colour temperature much cooler than in the high/soft state." + This is consistent with a standard outer disk runeated or obscured bevond a trausition radius RoeRisco., This is consistent with a standard outer disk truncated or obscured beyond a transition radius $R_{\rm c} \gg R_{\rm ISCO}$. +" The inner region of the inflow contributes mostly to the power-law component, perhaps through upscattering in a imoderately optically-thick corona. or in a magnetized wind (I&uucic and Dicknell, 2001)."," The inner region of the inflow contributes mostly to the power-law component, perhaps through upscattering in a moderately optically-thick corona, or in a magnetized wind (Kuncic and Bicknell, 2004)." + The observed auticorrelation between bolometric disk huninosity and peak colour temperature. aud the observed sequence of spectral state transitions imply that the ultraluninous branch occurs at even higher accretion rates than the classical very high state. hence probably at AL> a few Altaa- ," The observed anticorrelation between bolometric disk luminosity and peak colour temperature, and the observed sequence of spectral state transitions imply that the ultraluminous branch occurs at even higher accretion rates than the classical very high state, hence probably at $\dot{M} >$ a few $M_{\rm Edd}$." +The ultrahuninous brauch providesÜULXs., The ultraluminous branch provides. + It supports the cool-disk spectral model for ULXs. »t not its interpretation in terms of intermediate-nass ΟΙ».," It supports the cool-disk spectral model for ULXs, but not its interpretation in terms of intermediate-mass BHs." +" At the other end of the bridge, we showed that two prototypical ULXs (NGCIIS13 X-l and X-2) also move along a track consistent with the ultrahuunmous brauch. with anticorrelations between disk luminosity and temperature. aud between imucr-disk radius and temperature"," At the other end of the bridge, we showed that two prototypical ULXs 1313 X-1 and X-2) also move along a track consistent with the ultraluminous branch, with anticorrelations between disk luminosity and temperature, and between inner-disk radius and temperature." +" If this interpretation is correct, all the main spectral and timing properties of accreting BUS (including their characteristic LF-QPO frequencies) depend simultancoush on two factors: a imass scaling and an accretion rate SCine. at least when AL> afew A/pag."," If this interpretation is correct, all the main spectral and timing properties of accreting BHs (including their characteristic LF-QPO frequencies) depend simultaneously on two factors: a mass scaling and an accretion rate scaling, at least when $\dot{M} >$ a few $M_{\rm Edd}$ ." + We sugecstee that 16 observed ULX parameters may. be best explained with BIT masses ~50 10037. and ALz20Mpgqq., We suggest that the observed ULX parameters may be best explained with BH masses $\sim 50$ $100 M_{\odot}$ and $\dot{M} \approx 20 M_{\rm Edd}$. + This uso iuplies that the radiative. cfiicicney c| the. )912al medium is &0.02., This also implies that the radiative efficiency of the non-thermal medium is $\la 0.02$. + Iu this scenario. ULXs persistently occupy the accretion-rate ultraluninous branch. whereas stellar-nass DIIS would ouly rarely reach such rates; and onlv near the peak of their outbursts (for example. we showed iu Section l1 that the stellaranass BIT T1713) 322 never does).," In this scenario, ULXs persistently occupy the high-accretion-rate ultraluminous branch, whereas stellar-mass BHs would only rarely reach such rates, and only near the peak of their outbursts (for example, we showed in Section 4 that the stellar-mass BH $-$ 322 never does)." + We are currently investieating whether brief transitions to an ultralunuinous branch nav have occurred in other historical outbursts of Galactic DIIs: and whether other variable ULXs move along similar tracks., We are currently investigating whether brief transitions to an ultraluminous branch may have occurred in other historical outbursts of Galactic BHs; and whether other variable ULXs move along similar tracks. + A significant difference between he two classes ofobjects iiay be caused by the different vpes of Roche-lobe mass trauster., A significant difference between the two classes of objects may be caused by the different types of Roche-lobe mass transfer. +" ULXs may accrete a ew M. over €10° vr from B-type donor stars; while ransicut Calactic DIIS accretefrom older solar-iiass stars, which do not persistently fill their Roche lobes."," ULXs may accrete a few $M_{\odot}$ over $\la 10^6$ yr from B-type donor stars, while transient Galactic BHs accretefrom older solar-mass stars, which do not persistently fill their Roche lobes." +An equivalently good fit was provided by the BMC model.,An equivalently good fit was provided by the BMC model. + This model has the same number of free parameters as the CUTOFFPL+BB (see Table 3)). and would predict similar properties for the temperature and the size in which the soft photons are produced (thenormalizationoftheBMCmodelandreferences therein)..," This model has the same number of free parameters as the CUTOFFPL+BB (see Table \ref{tab:igrbfit}) ), and would predict similar properties for the temperature and the size in which the soft photons are produced \citep[the normalization of the BMC model is defined as the ratio of the source luminosity +to the square of the distance in units of 10~kpc, see e.g.][and references therein]{sidoli09b}." + Similar values of the fit parameters were also obtained from the analysis of the spectrum. of eextracted by using the total available exposure time of the observation (see Sect., Similar values of the fit parameters were also obtained from the analysis of the spectrum of extracted by using the total available exposure time of the observation (see Sect. + 4.2. and Table 4))., \ref{sec:igrresults} and Table \ref{tab:igrbfittotal}) ). + According to the discussion in Sect. 5.1..," According to the discussion in Sect. \ref{sec:xtediscussion}," + a BB emission with these properties seems unlikely in the case of and thus we suggest that the CUTOFFPL+MKL model can provide à more reasonable description of the data., a BB emission with these properties seems unlikely in the case of and thus we suggest that the CUTOFFPL+MKL model can provide a more reasonable description of the data. + We note that in the sspectrum of the supergiant HMXB JJ16320-4751 a similar soft component was found that could be fit with à BB of 0.07 keV but was attributed to a cloud surrounding the NS (Rodriguezetal..2006)., We note that in the spectrum of the supergiant HMXB J16320-4751 a similar soft component was found that could be fit with a BB of 0.07 keV but was attributed to a cloud surrounding the NS \citep{rodriguez06}. +. Following the CUTOFFPL+MKL interpretation. the rate resolved analysis carried out for the observation of wwould indicate that the properties of the MKL component do not change significantly with the souree count rate and the increase in the hardness ratio observed in Figs.," Following the CUTOFFPL+MKL interpretation, the rate resolved analysis carried out for the observation of would indicate that the properties of the MKL component do not change significantly with the source count rate and the increase in the hardness ratio observed in Figs." + 11. and 10. ts most likely due to a change in the CUTOFFPL photon index., \ref{fig:igrbhardness} and \ref{fig:igrb} is most likely due to a change in the CUTOFFPL photon index. + Furthermore. no significant variations in the absorption column density were revealed in the different rate-resolved spectra.," Furthermore, no significant variations in the absorption column density were revealed in the different rate-resolved spectra." + This is similar to the results discussed above for As for aa comparison between the results of the present oobservation and the observations carried out in the same energy band (0.5-10 keV) when this source was in outburst does not indicate a clear correlation between the power law photon index. the absorption column density. and the source X-ray flux (see Sect. 1)).," This is similar to the results discussed above for As for a comparison between the results of the present observation and the observations carried out in the same energy band (0.5-10 keV) when this source was in outburst does not indicate a clear correlation between the power law photon index, the absorption column density, and the source X-ray flux (see Sect. \ref{sec:intro}) )." + It is interesting that the soft component in this source detected by the oobservation appears to have a different origin from. that detected by Sidolietal.(2009b) when wwas in outburst (see Sect. 1))., It is interesting that the soft component in this source detected by the observation appears to have a different origin from that detected by \citet{sidoli09b} when was in outburst (see Sect. \ref{sec:intro}) ). + On that occasion. the soft component appeared to be caused by thermal emission from a hot-spot on the NS surface.," On that occasion, the soft component appeared to be caused by thermal emission from a hot-spot on the NS surface." + We note that. even if the soft component observed by iis interpreted in terms of a BB emission. the emitting region derived from the fit is considerably larger than the NS radius and it is thus unlikely that it is produced on the NS surface.," We note that, even if the soft component observed by is interpreted in terms of a BB emission, the emitting region derived from the fit is considerably larger than the NS radius and it is thus unlikely that it is produced on the NS surface." + Finally. for both aand wwe investigated whether the harder spectral component detected in these sources might be produced by the X-ray emission from the NS supergiant companion.," Finally, for both and we investigated whether the harder spectral component detected in these sources might be produced by the X-ray emission from the NS supergiant companion." + The time-averaged X-ray luminosity that we measured from these sources in quiescence matches quite well the luminosity expected from an isolated OB supergiant or from colliding winds in a binary containing OB supergiant stars (seee.g..Gudel&Nazel.2009.forareview)..," The time-averaged X-ray luminosity that we measured from these sources in quiescence matches quite well the luminosity expected from an isolated OB supergiant or from colliding winds in a binary containing OB supergiant stars \citep[see e.g.,][for a review]{gudel09}." + However. this interpretation appears to be contrived for the following reasons.," However, this interpretation appears to be contrived for the following reasons." + The X-ray spectrum of isolated or colliding wind binaries with OB supergiant stars is usually described well by a model comprising one or more thermal components (MKLinXSPEC.. see 2009).," The X-ray spectrum of isolated or colliding wind binaries with OB supergiant stars is usually described well by a model comprising one or more thermal components \citep[MKL in {\sc xspec}, see ." + The softer MKL component has a typical temperature of ~0.2-0.7 keV. and is thus similar to those we detected in aand This component is thought to be generated by the shocks within the stellar wind.," The softer MKL component has a typical temperature of $\sim$ 0.2-0.7 keV, and is thus similar to those we detected in and This component is thought to be generated by the shocks within the stellar wind." + The hotter MKL component. extending up to several keV. can have a temperature as high as ~1-3 keV and is characterized by à number of very prominent emission lines (seealsoRaassenetal..2008).," The hotter MKL component, extending up to several keV, can have a temperature as high as $\sim$ 1-3 keV and is characterized by a number of very prominent emission lines \citep[see also][]{raassen08}." +. This hard component is usually interpreted in terms of magnetically confined wind shocks. highly compressed wind shocks. or inverse Compton scattering of photospheric UV photons by relativistic particles accelerated within the shocks (AlbaceteColomboetal..2007).," This hard component is usually interpreted in terms of magnetically confined wind shocks, highly compressed wind shocks, or inverse Compton scattering of photospheric UV photons by relativistic particles accelerated within the shocks \citep{colombo07}." + Possible detections of a non-thermal X-ray emission from OB supergiant stars were reported only in only two cases. but they still lack confirmation (Gudel&Nazel.2009).," Possible detections of a non-thermal X-ray emission from OB supergiant stars were reported only in only two cases, but they still lack confirmation \citep{gudel09}." +. The X-ray spectra of aand wwere reproduced well using a CUTOFFPL model. and no prominent emission lines were detected.," The X-ray spectra of and were reproduced well using a CUTOFFPL model, and no prominent emission lines were detected." + The values of the photon index. T. derived from aare also comparable to those obtained previously when the sources were in outburst. thus suggesting that a common mechanism produces their harder X-ray component.," The values of the photon index, $\Gamma$, derived from are also comparable to those obtained previously when the sources were in outburst, thus suggesting that a common mechanism produces their harder X-ray component." + Furthermore. the relatively rapid X-ray variability (of period few thousands seconds) observed in the lightcurves of aand iisan not reminiscent of the typical X-ray variability of the OB stars. which. when present. takes place on longer timescales (tensofks.seee.g.AlbaceteColomboetal..2007).," Furthermore, the relatively rapid X-ray variability (of period few thousands seconds) observed in the lightcurves of and is not reminiscent of the typical X-ray variability of the OB stars, which, when present, takes place on longer timescales \citep[tens of ks, see e.g,][]{colombo07}." +. We conclude that the harder X-ray emission from aand iis most likely due to residual accretion taking place onto the NS at a much lower rate than during outburst., We conclude that the harder X-ray emission from and is most likely due to residual accretion taking place onto the NS at a much lower rate than during outburst. + The three oobservations that we have analyzed in the present work. indicate that the quiescent spectra of the two prototypical SEXT aand aare characterized by two different spectral components. one," The three observations that we have analyzed in the present work, indicate that the quiescent spectra of the two prototypical SFXT and are characterized by two different spectral components, one" +If the image has been sampled at the Nyquist deusity or higher. then we cau apply Parseval's Theorem to obtain the simple form Thus in the limit of faint. Nvquist-sampled. unresolved sources. the S/N for detection/photomery depeuds up the effective area Aga of the PSF. aud we cau easily see how the PRE affects this.,"If the image has been sampled at the Nyquist density or higher, then we can apply Parseval's Theorem to obtain the simple form Thus in the limit of faint, Nyquist-sampled, unresolved sources, the $S/N$ for detection/photometry depends up the effective area $A_{SN}$ of the PSF, and we can easily see how the PRF affects this." + The f-space integral form is particularly convenient since the Airy PSF is bounced to &<2D/X. aud uo convolutious must be executed. Walser.Toury.&Lup," The $k$ -space integral form is particularly convenient since the Airy PSF is bounded to $k<2D/\lambda$, and no convolutions must be executed. \citet{KTL}," +pino(2000).. for example. derive aud inake use of this fori.," for example, derive and make use of this form." + ead noise aud dark current produce white nolse that cau be subsumed iuto » in the simple formula (12))., Read noise and dark current produce white noise that can be subsumed into $n$ in the simple formula \ref{AreaSN}) ). + When the image is uot Nyquist sampled. it is easier to transforin tle ePSE to z-space aud use(10).," When the image is not Nyquist sampled, it is easier to transform the ePSF to $x$ -space and use." +. The Fisher matrix also allows us to evaluate the astrometric accuracy for point sources., The Fisher matrix also allows us to evaluate the astrometric accuracy for point sources. + This is not a primary goal for SNAP. but [will present some results in relastrometry..," This is not a primary goal for , but I will present some results in \\ref{astrometry}." + Iu the background-lunited case. the oue-dimiensioual uncertainty of poiut-source astrometry is simply quantified as This aucl related forms are given by Ixaiser.Tonry.&Luppino(2000) aud used lor analysis ol the proposedPOI project.," In the background-limited case, the one-dimensional uncertainty of point-source astrometry is simply quantified as This and related forms are given by \citet{KTL} and used for analysis of the proposed project." + The incorporation of cosmic-ray (CI) hits iuto the Fisher formalisin is easy., The incorporation of cosmic-ray (CR) hits into the Fisher formalism is easy. + We just remove [rom the sum (10)) the information contributed by pixels that are ruiued., We just remove from the sum \ref{fisher}) ) the information contributed by pixels that are ruined. + Iu theSNAP mission. we expect the CRs to span many pixels. while the ePSF will be <2 pixels across.," In the mission, we expect the CRs to span many pixels, while the ePSF will be $\lesssim 2$ pixels across." + Hence the probability ol losing the entire exposures information is esseutially equal to the probability Pry of the central pixel being contaminated during an exposure., Hence the probability of losing the entire exposure's information is essentially equal to the probability $P_{CR}$ of the central pixel being contaminated during an exposure. + For detectors with uou-destructive readout (suchas, For detectors with non-destructive readout (suchas + The quest to fud ordered structures in the svstem of satellites of the Mille Way galaxy dates back to 30 vears ago when possible aliguiieuts of globular clusters aud/or dwairt galaxies along wide streams were first noted 1982)., The quest to find ordered structures in the system of satellites of the Milky Way galaxy dates back to 30 years ago when possible alignments of globular clusters and/or dwarf galaxies along wide streams were first noted . +. The mounting couseusus for scenarios im which the accretionof satellites has a major role in the formationof the outer halo of the Galaxy proiupted a new. burst of such kind of studies since the mid “90s up to the present dav 2002)., The mounting consensus for scenarios in which the accretion of satellites has a major role in the formation of the outer halo of the Galaxy prompted a new burst of such kind of studies since the mid '90s up to the present day . + Despite the many interesting suggestions. none of the quoted studies was able to provide a couclusive proof of the reality of the aliments. mainly because of the overwhehuing difficulty to assess the statistical significance of structures forme by inherently πια] wmmbers of objects.," Despite the many interesting suggestions, none of the quoted studies was able to provide a conclusive proof of the reality of the alignments, mainly because of the overwhelming difficulty to assess the statistical significance of structures formed by inherently small numbers of objects." +" However, recent theoretical aud observational achievements may help— us to look iuto the problem frou a different aud more fruitful perspective: Thevere example of a satellite accretion/disruptiou is provided by the Sagittarius dwarf Spheroidal galaxy 1997).. which is currently moreing with the Milkv Way. aud is carrving its own elobular cluster svstem (i0. M 51. Ter 8. Arp 2 and Ter 7. previously believed to be normal Galactic elobulars)."," However, recent theoretical and observational achievements may help us to look into the problem from a different and more fruitful perspective: The example of a satellite accretion/disruption is provided by the Sagittarius dwarf Spheroidal galaxy , which is currently merging with the Milky Way, and is carrying its own globular cluster system (i.e., M 54, Ter 8, Arp 2 and Ter 7, previously believed to be normal Galactic globulars)." + There is now clear observational evidence that the Ser dSph is loosing stars uncer the strain of the Milky Way tidal field., There is now clear observational evidence that the Sgr dSph is loosing stars under the strain of the Milky Way tidal field. + These tidallv-reumioved stars are found. along a huge (aud quite colerent) stream extending all over the sky therein). tracing the orbit of the pareut ealaxy.," These tidally-removed stars are found along a huge (and quite coherent) stream extending all over the sky , tracing the orbit of the parent galaxy." +"with the stellar photons (e,:e;5,= 1).",with the stellar photons $\vec{e}_{\star}\cdot\vec{e}_{obs}=1$ ). + This effect is smoothed if the finite size of the massive star is considered., This effect is smoothed if the finite size of the massive star is considered. + The escaping gamma-ray density at inferior conjunction 1s more important than at superior conjunction as TeV photons suffer less from absorption., The escaping gamma-ray density at inferior conjunction is more important than at superior conjunction as TeV photons suffer less from absorption. + The semi-analytical method is ideal to study the first generation of particles in the cascade as it provides quick and accurate solutions., The semi-analytical method is ideal to study the first generation of particles in the cascade as it provides quick and accurate solutions. + In. principle. this method can be extended to ral arbitrary number of generation but the computing time increases tremendously.," In principle, this method can be extended to an arbitrary number of generation but the computing time increases tremendously." + The Monte Carlo approach is well suited to treat complex three dimensional radiative transfer problems., The Monte Carlo approach is well suited to treat complex three dimensional radiative transfer problems. + With this method. the full cascade radiation (neluding all generations) can be computed with a reasonable amount of time but a large number of events is required to have enough statistics for accurate predictions.," With this method, the full cascade radiation (including all generations) can be computed with a reasonable amount of time but a large number of events is required to have enough statistics for accurate predictions." + Figure 4 gives the escaping gamma-ray spectra at both conjunctions in LS 5039., Figure \ref{compare} gives the escaping gamma-ray spectra at both conjunctions in LS 5039. + The Monte Carlo output is compared with the semi-analytical results in the same configuration as in Fig., The Monte Carlo output is compared with the semi-analytical results in the same configuration as in Fig. + 3. for Ξ30° and 1507., \ref{map} for $\psi=30\degr$ and $150\degr$. + Both approaches give similar results for the first generation of gamma rays., Both approaches give similar results for the first generation of gamma rays. + There are slight differences mainly due to statistical and. binning effect in the Monte Carlo result. particularly at u=307 where the absorption is high.," There are slight differences mainly due to statistical and binning effect in the Monte Carlo result, particularly at $\psi=30\degr$ where the absorption is high." + The contribution from additional generations of pairs to the cascade radiation is of major importance as it dominates the overall escaping gamma-ray flux where the primary photons are fully absorbed., The contribution from additional generations of pairs to the cascade radiation is of major importance as it dominates the overall escaping gamma-ray flux where the primary photons are fully absorbed. + The Monte Carlo approach is needed to compute the cascade radiation where absorption is strong at superior conjunction., The Monte Carlo approach is needed to compute the cascade radiation where absorption is strong at superior conjunction. + In practice. the one-generation approximation catches the main features of the full three-dimensional pair cascade calculation elsewhere along the orbit.," In practice, the one-generation approximation catches the main features of the full three-dimensional pair cascade calculation elsewhere along the orbit." + Three-dimensional cascade radiation presents identical spectral features to the one-dimensional limit (2). (Fig. 5))., Three-dimensional cascade radiation presents identical spectral features to the one-dimensional limit \citep{2009A&A...507.1217C} (Fig. \ref{anis}) ). + Below the threshold energy for pair production. ej«mz[2e5(1—cosAa) with ey the stellar photon energy. pairs cool down inverse Compton scattering in the Thomson regime and accumulate at lower energy in a ~—1.5 photon index power-law tail.," Below the threshold energy for pair production, $\epsilon_1 f44).," Figure \ref{mag} shows the effects of an uniform ambient magnetic field on the cascade radiation for $B=0$, 3 and 10 G. The VHE emission is quenched as synchrotron radiation becomes the dominant cooling channel for electrons produced in the cascade $t_{ic}>t_{syn}$ )." + The large contribution of the cascade in the TeV band is preserved if the magnetic field does not exceed a few Gauss (see Fig. 1)., The large contribution of the cascade in the TeV band is preserved if the magnetic field does not exceed a few Gauss (see Fig. \ref{domain}) ). + Synchrotron radiation contributes to the total flux im the X-ray to soft gamma-ray energy band., Synchrotron radiation contributes to the total flux in the X-ray to soft gamma-ray energy band. + These photons do not participate to the cascade as their energy does not exceed 100 MeV. which ts insufficient for pair production with the stellar photons.," These photons do not participate to the cascade as their energy does not exceed 100 MeV, which is insufficient for pair production with the stellar photons." +he racio core resolved.,the radio core resolved. + Our best constraint on the size scale of the compact jet comes from the first. brightest observation on 2010 April 25. when the source was unresolved. down to he beam size of 12.5«5.2 numas? in position angle eeast of north.," Our best constraint on the size scale of the compact jet comes from the first, brightest observation on 2010 April 25, when the source was unresolved down to the beam size of $12.5\times5.2$ $^2$ in position angle east of north." + his. corresponds to a source size ol 12.5(d/kpe) aau. where d is the source distance (ie.«4ttau 2010).," This corresponds to a source size of $<12.5(d/{\rm kpc})$ au, where $d$ is the source distance \citep[i.e.\ $<44$\,au for the distance of 3.5\,kpc claimed by][]{Sha10}." +. Following the detection of the radio core. we re-reduced he VLBA data of 2010 February 18. 23 and 26 (program code BB290) presented by Yangetal.(2000)... using a larger image size to search for receding ejecta to the south-cast of he core.," Following the detection of the radio core, we re-reduced the VLBA data of 2010 February 18, 23 and 26 (program code BB290) presented by \citet{Yan10}, using a larger image size to search for receding ejecta to the south-east of the core." + No new components were detected. to Sa upper imits of 0.62. 0.74. and rrespectively.," No new components were detected, to $5\sigma$ upper limits of 0.62, 0.74, and respectively." + Using high-precision optical astrometry. we have been able to locate the core of the X-ray binary svstem NTE J1752-223.," Using high-precision optical astrometry, we have been able to locate the core of the X-ray binary system XTE J1752-223." + From VLBI observations made during the hard spectral state of the svstem. when the radio emission is dominated by a compact. unresolved. core jet. we were able to further refine the core position.," From VLBI observations made during the hard spectral state of the system, when the radio emission is dominated by a compact, unresolved core jet, we were able to further refine the core position." + In light of this new information. we now reanalvze the VLBI data presented by. Yang (2010).," In light of this new information, we now reanalyze the VLBI data presented by \citet{Yan10}." +. The VLBA image of 2010 February 26 published by Yangetal.(2010). shows two components. labelled A and D. While component A was detected in all four VEDI images from 2010 February. component. D. is detected in only this one image. and was interpreted. by Yangctal.(2010). as à receding component. leading the authors to infer that the true core of the system lay between components Ao and D in this image.," The VLBA image of 2010 February 26 published by \citet{Yan10} shows two components, labelled A and B. While component A was detected in all four VLBI images from 2010 February, component B is detected in only this one image, and was interpreted by \citet{Yan10} as a receding component, leading the authors to infer that the true core of the system lay between components A and B in this image." + Our new determination of the true core position implies that both these components are located to the north-west of the core Clable 13)., Our new determination of the true core position implies that both these components are located to the north-west of the core (Table \ref{tab:components}) ). + From this. we infer that components A and D must arise from separate ejection events.," From this, we infer that components A and B must arise from separate ejection events." + While we are unable to reliably. pinpoint the exact epoch of ejection of each of those events. we can use X-ray spectral and timing information together with constraints from the integrated. radio light curves to obtain a rough estimate for component A. Radio Wares in N-rav binaries have been linked to the transition from a hard intermediate state (LIENS) to a soft intermediate state (SIMS) during a rapid. phase of X-ray spectral softening at the peak of the outburst (Fenderctal.2004).," While we are unable to reliably pinpoint the exact epoch of ejection of each of those events, we can use X-ray spectral and timing information together with constraints from the integrated radio light curves to obtain a rough estimate for component A. Radio flares in X-ray binaries have been linked to the transition from a hard intermediate state (HIMS) to a soft intermediate state (SIMS) during a rapid phase of X-ray spectral softening at the peak of the outburst \citep{Fen04}." +. Also associated with this transition are a sharp crop in the integrated rms variability of the X-ray emission and a reduction in the coherence of the associated. quasi-periocdic oscillations (QPOs). from high-coherence. Type € QPOs associated: with [lat-topped. noise in the power spectrum. to. lower-coherence Type A or Type B QPOs associated with weak red noise (Belloniοἱal.2005).," Also associated with this transition are a sharp drop in the integrated rms variability of the X-ray emission and a reduction in the coherence of the associated quasi-periodic oscillations (QPOs), from high-coherence Type C QPOs associated with flat-topped noise in the power spectrum, to lower-coherence Type A or Type B QPOs associated with weak red noise \citep{Bel05}." +. However. we note that Fenderctal.(2009) found that while these changes in the variability properties were closely associated with radio ection events. the association was not exact. such that one could. precede the other by up to a few clays.," However, we note that \citet{Fen09} found that while these changes in the variability properties were closely associated with radio ejection events, the association was not exact, such that one could precede the other by up to a few days." + On MJD555215.9. NPE 1752-223 was in à ΗΛ. with a 2.21IHIz Type € QDPO (Shaposhnikovetal.2010).," On 55215.9, XTE J1752-223 was in a HIMS with a Hz Type C QPO \citep{Sha10}." +". The integrated ris variability then decreased. [from 25 to 18 per cent as the QPO frequency. rose to 2.311, by NMJD555217.9. and. by 555218.8. the observed QPOs had changed from Type € to Pype A/B. This sugeests that the transition from HIEMS to SIMS occurred around 555218."," The integrated rms variability then decreased from 25 to 18 per cent as the QPO frequency rose to Hz by 55217.9, and by 55218.8, the observed QPOs had changed from Type C to Type A/B. This suggests that the transition from HIMS to SIMS occurred around 55218." + Supporting this inference is the bright mm.Jy). flat-spectrum radio emission observed on 55!5217 between 1.2 and 19€CGllz (Brocksoppctal. 2010).," Supporting this inference is the bright mJy), flat-spectrum radio emission observed on 55217 between 1.2 and GHz \citep{Bro10}." +. This represents an increase of the radio brightness by an order of magnitude as compared to Che initial 22mJv racio detection in the rising hard state (Brocksoppetal.2009).. suggesting the onset of a radio Dare.," This represents an increase of the radio brightness by an order of magnitude as compared to the initial 2-mJy radio detection in the rising hard state \citep{Bro09}, suggesting the onset of a radio flare." + Vhis increase in racio brightness corresponds to an increase of only a factor of ~2 inthe 5OkkeV. τήBAL and kkeV ALAXI/GSC X- count rates (Nakahiraetal.2010) over the same period. which is not consistent with the racio/X-ray luminosity correlation (Cialloetal.2003). found. for the compact. jets of many hard state black hole candidates (althoughnoteanever-inereasingnumberofoutliers:e.g.Gallo 2007).," This increase in radio brightness corresponds to an increase of only a factor of $\sim 2$ in the keV /BAT and keV /GSC X-ray count rates \citep{Nak10} over the same period, which is not consistent with the radio/X-ray luminosity correlation \citep{Gal03} found for the compact jets of many hard state black hole candidates \citep[although note an ever-increasing number of outliers; e.g.][]{Gal07}." +.. While this would tend to support the interpretation of a radio Hare on 555217. the flat radio spectrum over more than a decade in frequency instead argues that this radio emission most probably still arises from a compact jet. rather than discrete. optically-thin transient. ejecta.," While this would tend to support the interpretation of a radio flare on 55217, the flat radio spectrum over more than a decade in frequency instead argues that this radio emission most probably still arises from a compact jet, rather than discrete, optically-thin transient ejecta." + The anomalously bright radio emission could then correspond to the period of jet instability known to occur immediately preceding a Iarge racio [lare (Fenderetal.2004)., The anomalously bright radio emission could then correspond to the period of jet instability known to occur immediately preceding a large radio flare \citep{Fen04}. +. Further racio information is available from the APCA monitoring. which covered the entire outburst. from. the initial rising hard state through to the ἄοσαν back to quiescence.," Further radio information is available from the ATCA monitoring, which covered the entire outburst from the initial rising hard state through to the decay back to quiescence." + Although a full analysis is bevond the scope of this paper ancl will be presented by Brocksopp ct ((in prep.).," Although a full analysis is beyond the scope of this paper and will be presented by Brocksopp et (in prep.)," + we summarize the relevant. information here., we summarize the relevant information here. + The integrated. radio light curve shows at least. two large flares. followed. by a. few. smaller events.," The integrated radio light curve shows at least two large flares, followed by a few smaller events." + The [last [Bat-spectrum radio detection was made on 555217. after which the 9-Gllz racio (ux density dropped to mmy on 555220. before peaking at 9.0mmJyv on 555221.," The last flat-spectrum radio detection was made on 55217, after which the 9-GHz radio flux density dropped to mJy on 55220, before peaking at mJy on 55221." + This suggests an initial ejection date between 555217 and 55220., This suggests an initial ejection date between 55217 and 55220. + Phe second Hare was somewhat broader. with the rise phase beginning after 555226 and the Hare peaking at 10.9mmJyv. at Οι on 555242.," The second flare was somewhat broader, with the rise phase beginning after 55226 and the flare peaking at mJy at GHz on 55242." + Phe. double- light curve supports our conclusion from Section 4.1 that the outburst comprised at least two ejection events., The double-peaked light curve supports our conclusion from Section \ref{sec:double} that the outburst comprised at least two ejection events. + Negoroetal.(2010) reported a sharp increase of the soft (< 4tkkeV) NX-rav [lux and a decline of the hard, \citet{Neg10} reported a sharp increase of the soft $<4$ keV) X-ray flux and a decline of the hard +One such example is the high euergy emitting binary LS7+617303 (Massi.etal.2001)..,"One such example is the high energy emitting binary $LS\,I\,+61^\circ 303$ \citep{Massi}." + Lt was first detected iu the TeV ranee with MAGIC (AlbertJ..etal.2006) anc further observed with VERITAS at [lus levels ranging between aud of the Crab Nebula (Acciarietal.2008).., It was first detected in the TeV range with MAGIC \citep{Magic} and further observed with VERITAS at flux levels ranging between and of the Crab Nebula \citep{Veritas}. + 'Ehis source has been observed throughout most of the electromagnetic spectrum starting with racio frequencies aud extending to VHE eamima rays (Lealy2001)., This source has been observed throughout most of the electromagnetic spectrum starting with radio frequencies and extending to VHE gamma rays \citep{Leahy}. +". This broad spectral study inclicates that the system consists of a main sequence Be star of mass M,=12.542.5AL.(Casares.elal. 2005). 3.uwrounded by a circumstellar cisk (Cirundstrometal.2007:Paredes2007).. and a compact companion separated by teus of solar radii at periastron."," This broad spectral study indicates that the system consists of a main sequence Be star of mass $M_1=12.5\pm2.5\,M_{\odot}$\citep{Casares}, surrounded by a circumstellar disk \citep{Grundstrom, Paredes}, and a compact companion separated by tens of solar radii at periastron." + The compact companion cau be either a neutron star or a black bole (Casares.elal.2005).. aud its exact uature is still subject ol investigation aud debate (Zdziarskietal.2010).," The compact companion can be either a neutron star or a black hole \citep{Casares}, and its exact nature is still subject of investigation and debate \citep{Neronov}." +. The maximum VHE einission occurs close to apastron (Acciarietal.2008:AlbertJ..2006).. suggestingMD that absorption plays an important role in the The outline of the paper is as follows: first we describe tlie moclel of attenuation due to pair production.," The maximum VHE emission occurs close to apastron \citep{Veritas, Magic}, suggesting that absorption plays an important role in the The outline of the paper is as follows: first we describe the model of attenuation due to pair production." + Then we cousider the particular case of LS [+61 303., Then we consider the particular case of LS I +61 303. + We assume that the high energy radiation is emitted [roi the vicinity of the compact object aud that its emission is isotropic aud constant in time., We assume that the high energy radiation is emitted from the vicinity of the compact object and that its emission is isotropic and constant in time. + The modulation due to photou-photon tuteractious is found to be iusullicieut to account for the VERITAS observations., The modulation due to photon-photon interactions is found to be insufficient to account for the VERITAS observations. + For this reason. we include additional interactions of VHE photons with circumstellar material.," For this reason, we include additional interactions of VHE photons with circumstellar material." + This model permits to coustrain the orbital parameters aud the mass of the compact object as well as the density of ejected material from the companion star., This model permits to constrain the orbital parameters and the mass of the compact object as well as the density of ejected material from the companion star. + The radiative transfer equation (Chaudrasekhar1960) [or the intensity Z(s.E). where s is the distauce traveled by a photon of energy E from the emission point is Where n(s.€) is the spectral density of background photous of euergy ε emitted by the main sequence star. e(E.e.£) is the cross for the interaction between photons colliding at angle £. and JCE.s) is a source term.," The radiative transfer equation \citep{chandra} for the intensity $I(s,E)$, where $s$ is the distance traveled by a photon of energy $E$ from the emission point is Where $n(s,\epsilon)$ is the spectral density of background photons of energy $\epsilon$ emitted by the main sequence star, $\sigma(E,\epsilon,\xi)$ is the cross for the interaction between photons colliding at angle $\xi$ , and $j(E,s)$ is a source term." +Open star clusters (OCS) play an important role in studying the formation and evolution ol the Galactic disk and the stellar evolution as well.,Open star clusters (OCs) play an important role in studying the formation and evolution of the Galactic disk and the stellar evolution as well. + Also. OCs are the most suitable objects lor studying the space-age structure of the Galactic disk.," Also, OCs are the most suitable objects for studying the space-age structure of the Galactic disk." + The fundamental physical parameters of OCs. e.g. distance. reddening. age. and metallicity are necessary [or studyiug the Galactic disk.," The fundamental physical parameters of OCs, e.g. distance, reddening, age, and metallicity are necessary for studying the Galactic disk." + The Galactic. racial ancl vertical. abundance gradient can be studied by OCs (ILou et al.," The Galactic, radial and vertical, abundance gradient can be studied by OCs (Hou et al." + 2000: Chen et al., 2000; Chen et al. + 2003: INim and Sung 2003: Tadross 2003: Nim et al., 2003; Kim and Sung 2003; Tadross 2003; Kim et al. + 2005)., 2005). + Thev are excellent probes of the Galactic disk structure (Janes Phelps. 1994: Dica et al.," They are excellent probes of the Galactic disk structure (Janes Phelps, 1994; Bica et al." + 2006)., 2006). + The strong interest of OCs results come from their fundamental properties., The strong interest of OCs results come from their fundamental properties. + Among the 1787 currently. OC's. most than half of them have been poorly studied or even unstudied up to now. Piatti et al. (," Among the 1787 currently OCs, most than half of them have been poorly studied or even unstudied up to now, Piatti et al. (" +2011).,2011). + The current. paper is thus part of our continuation series whose goal is to obtain (he main astroplivsical properties of previously unstudied OCS using modern databases., The current paper is thus part of our continuation series whose goal is to obtain the main astrophysical properties of previously unstudied OCs using modern databases. + The most important thing for using PPMXL database lies in containing the positions. proper motions of and the Near Infrared (NUR) photometry of," The most important thing for using PPMXL database lies in containing the positions, proper motions of and the Near Infrared (NIR) photometry of" +for computing the curve per night and filter is ~75.,for computing the curve per night and filter is $\sim75$. + Notice the wide color coverage for the standard stars., Notice the wide color coverage for the standard stars. + The last step on the calibration is the aperture correction., The last step on the calibration is the aperture correction. + As no available bright and isolated stars exist on the cluster images. we used DAOPHOT to subtract from the image the stars in the neighborhood of the brightest ones. in order to compute the difference between the aperture and the PSF-fitting magnitudes.," As no available bright and isolated stars exist on the cluster images, we used DAOPHOT to subtract from the image the stars in the neighborhood of the brightest ones, in order to compute the difference between the aperture and the PSF-fitting magnitudes." + In view of the stable seeing conditions. we used the same aperture for calculating the aperture photometry of the standard and cluster stars.," In view of the stable seeing conditions, we used the same aperture for calculating the aperture photometry of the standard and cluster stars." + In order to check the photometric homogeneity of the data and of the calibration to the standard photometric system. one cluster (NGC 3201) was observed in both runs.," In order to check the photometric homogeneity of the data and of the calibration to the standard photometric system, one cluster (NGC 3201) was observed in both runs." + Having one common field. it is possible to analyze the individual star photometry. and test if any additional zero point difference and/or color term exist.," Having one common field, it is possible to analyze the individual star photometry, and test if any additional zero point difference and/or color term exist." + The latter check is crucial when measures of the relative position of CMD features are going to be done., The latter check is crucial when measures of the relative position of CMD features are going to be done. + The comparison between the two runs ts presented in Fig. 3..," The comparison between the two runs is presented in Fig. \ref{compare}," + where 456 common stars with internal photometric errors (as given by ALLFRAME) smaller than 0.02 mag are used., where 456 common stars with internal photometric errors (as given by ALLFRAME) smaller than 0.02 mag are used. + Fig., Fig. + 3 shows that there are no systematic differences between the two runs., \ref{compare} shows that there are no systematic differences between the two runs. + The slope of the straight lines best fitting all the points in both the AVTU ) plane and the (ΝΑ plane is < 0.002. and <0.002+0.003 in the (V-LA(U— Do plane.," The slope of the straight lines best fitting all the points in both the $\Delta V^{\rm apr}_{\rm dec}$ ) plane and the $\Delta I^{\rm +apr}_{\rm dec}$ ) plane is $\leq 0.001\pm0.002$ , and $\leq +0.002\pm0.003$ in the $\Delta (V-I)^{\rm apr}_{\rm dec}$ ) plane." + The zero point differences are always <0.01 mag., The zero point differences are always $\le 0.01$ mag. + This ensures the homogeneity of our database. particularly for relative measurements within the CMDs.," This ensures the homogeneity of our database, particularly for relative measurements within the CMDs." + In order to facilitate the readers work. we present in Tables3.. 4 and$ the basic parameters available for our GGCssample!.," In order to facilitate the readers work, we present in Tables, \ref{param2} and \ref{param3} the basic parameters available for our GGCs." +. In Table we give the coordinates. the position. and the metallicity of the clusters: right ascension and. declination (epoch J2000. columns 3 and 4» Galactic longitude and latitude (columns 5 and 6): Heliocentrie. (column 7) and Galactocentric (column 8) distances (assuming 7?; =8.0 kpe): spatial components (X.Y.Z) (columns 9. IO and 11) in the Sun-centered coordinate system (X pointing toward the Galactic center. Y in direction of Galactic rotation. Z toward North Galactic Pole) and. finally. the metallicity given in Rutledge et al. (1997)).," In Table we give the coordinates, the position, and the metallicity of the clusters: right ascension and declination (epoch J2000, columns 3 and 4); Galactic longitude and latitude (columns 5 and 6); Heliocentric (column 7) and Galactocentric (column 8) distances (assuming $R_{\sun}$ =8.0 kpc); spatial components (X,Y,Z) (columns 9, 10 and 11) in the Sun-centered coordinate system (X pointing toward the Galactic center, Y in direction of Galactic rotation, Z toward North Galactic Pole) and, finally, the metallicity given in Rutledge et al. \cite{rutledge97}) )," + on both the Zinn West (1984)) and Carretta Gratton (1997)) scales., on both the Zinn West \cite{zinnwest84}) ) and Carretta Gratton \cite{carretagratton97}) ) scales. + In Table 4.. the photometric parameters are given.," In Table \ref{param2}, the photometric parameters are given." + Column 3 lists the foreground reddening: column 4. the V magnitude level of the horizontal branch: column 5. the apparent visual distance modulus: integrated V magnitudes of the clusters are given i column 6; column 7 gives the absolute visual magnitude.," Column 3 lists the foreground reddening; column 4, the $V$ magnitude level of the horizontal branch; column 5, the apparent visual distance modulus; integrated $V$ magnitudes of the clusters are given in column 6; column 7 gives the absolute visual magnitude." + Columns 8 to 11 give the integrated color indices (uncorrected for reddening)., Columns 8 to 11 give the integrated color indices (uncorrected for reddening). + Column 12 gives the specific frequency of RR Lyrae variables. while column 13 list the horizontal-branch morphological parameter (Lee 1990)).," Column 12 gives the specific frequency of RR Lyrae variables, while column 13 list the horizontal-branch morphological parameter (Lee \cite{lee90}) )." +"coellicient. 0""NNl accounts for the WKlhein-Nishina elfect.",coefficient $0\leq \eta_{_{KN}}\leq 1$ accounts for the Klein-Nishina effect. +" For 5,=mintr..g~105. dus1 (see Appendix D. of Fan Piran 2006)."," For $\bar{\gamma}_e=\min \{ +\gamma_c, \gamma_m\}\sim 10^3$, $\eta_{_{KN}}\sim 1$ (see Appendix B. of Fan Piran 2006)." + To get σος we also need to estimate the parameter Y.," To get $\gamma_c$, we also need to estimate the parameter $Y_{_{EIC}}$." + In the shock front. the magnetic energy. density can be estimated as At P2 the forward shock front χο)reachesLAL the radius A?7107iW.emEySa)eeeFC(0p2)/2] .," In the shock front, the magnetic energy density can be estimated as At $t$, the forward shock front reaches the radius $R\simeq +1.9\times 10^{17}~{\rm +cm}~E_{k,53}^{1/4}n_0^{-1/4}t_3^{1/4}[(1+z)/2]^{-1/4}$ ." +" In the forwardoscar shock region. the energysun, density of FUY photons of the Hare can be estimated as where Lp, is the Luminosity of the Hare."," In the forward shock region, the energy density of FUV photons of the flare can be estimated as where $L_{ph}$ is the luminosity of the flare." + Note that the Hare lasts AY(]!1 and no009DNOQESIOALOl4ο.”UNOdpouxM3)5SUOdgY):|n"," Now $\gamma_m\simeq 970 \epsilon_{e,-1} C_p +E_{k,53}^{1/4}A_*^{-1/4}t_3^{-1/4}[(1+z)/2]^{1/4}$ and $\gamma_c\simeq +280\epsilon_{B,-2}^{-1}E_{k,53}^{1/4}A_*^{-5/4}t_3^{3/4}[(1+z)/2]^{-3/4}(1+Y)^{-1}$." +" Similar to the ISM case. at / the Forward shock front reaches the radius /?22.71075emLye,1HTbqoa]bo. Up.~13.eresem.3(B,obLASTAysE(iBytes Cun=STerescm|Lynanf,2d=,2ARS(L:)77. and the ELC parameter can be estimated by Provided that 77/70.3. we have Y=Vowe|Vy— and 5.~TO."," Similar to the ISM case, at $t$ the forward shock front reaches the radius $R \simeq 2.7\times 10^{16}~{\rm +cm}~E_{k,53}^{1/2}A_*^{-1/2}t_3^{1/2}[(1+z)/2]^{-1/2}$, $U_B +\simeq 13~{\rm ergs +~cm^{-3}}~\epsilon_{B,-2}E_{k,53}^{-1/2}A_*^{3/2}t_3^{-3/2}({1+z +\over 2})^{3/2}$, $U_{ph} = 87~{\rm ergs +~cm^{-3}}~L_{ph,49}E_{k,53}^{-3/2}A_*^{3\over 2}t_3^{-1 \over +2}({1+z \over 2})^{1/2}$, and the EIC parameter can be estimated by Provided that $\Delta T/t\sim 0.3$, we have $Y=Y_{_{SSC}}+Y_{_{EIC}}\simeq 3.4$ and $\gamma_c\sim 70$." + So 5.~10 for fF~10° s. Phe possibility of one Hare photon being scattered. (Le. the optical depth) can be estimated by Now most scattered. photons are. in the sub. MeV band and the count number is ~0.05 7 (where eq. (11))," So $\bar{\gamma}_e\sim 70$ for $t\sim 10^3$ s. The possibility of one flare photon being scattered (i.e., the optical depth) can be estimated by Now most scattered photons are in the sub MeV band and the count number is $\sim 0.05$ $^{-2}$ (where eq. \ref{eq:N_tot}) )" + and eq.(14)) have been taken into account). which is uncetectable for the BAT.," and \ref{eq:f_cor}) ) have been taken into account), which is undetectable for the BAT." + The tens MeV. photons (resulting in the keV. [lare photons-forward shock electrons interaction) may be still detectable for the GLAST., The tens MeV photons (resulting in the keV flare photons-forward shock electrons interaction) may be still detectable for the GLAST. + But he counts rate is not higher than that of the ISAL case., But the counts rate is not higher than that of the ISM case. + On the other hand. there are just a small fraction of bursts were born in the stellar wind (e.g.. Chevalier Li 2000: 'unaltescu Ixumar 2002).," On the other hand, there are just a small fraction of bursts were born in the stellar wind (e.g., Chevalier Li 2000; Panaitescu Kumar 2002)." + So we do not discuss this case 'urther., So we do not discuss this case further. + As pointed out in Beloborocoy (2005). the upscattered photons are de-collimated and their arrival time is allected by the spherical curvature of the blast wave.," As pointed out in Beloborodov (2005), the upscattered photons are de-collimated and their arrival time is affected by the spherical curvature of the blast wave." + The duration of the high energy. emission thus can be estimated as we have 7~4 in the ISM case and Z7.—2/ in the wind case. which could be much longer than AY.," The duration of the high energy emission thus can be estimated as we have $T\sim 4t$ in the ISM case and $T\sim 2t$ in the wind case, which could be much longer than $\Delta T$." + The duration increases when the anisotropic radiation of the up-seattered photons (sce cq. 13]]), The duration increases when the anisotropic radiation of the up-scattered photons (see eq. \ref{eq:j1}] ]) + has been taken into account because now the strongest emission are [rom 8—L/LI.," has been taken into account because now the strongest emission are from $\theta \sim +1/\Gamma$." + As a result. most of the up-scatterecl photons will arrive alter the PUY Uare.," As a result, most of the up-scattered photons will arrive after the FUV flare." + This lagging behavior is a signature of upscattering of internal radiation in the external blast wave. which may be tested by the observations.," This lagging behavior is a signature of upscattering of internal radiation in the external blast wave, which may be tested by the observations." + Assuming that the spectrum of the flare has the Form. Γρκια[ore>my. where τν~12. as reported in most X-ray. [lares (e.g. ODrien et al.," Assuming that the spectrum of the flare has the form $F_{\nu}\propto \nu^{-\beta_{_{XRT}}}~{\rm for~\nu>\nu_{\rm uv}}$, where $\beta_{_{XRT}}\sim 1.2$, as reported in most X-ray flares (e.g. O'Brien et al." + 2006)., 2006). + Phe total number of soft FUY photons reaching us (in unit area ancl without absorption) can be estimated hy ∖∖⋎↓↥∢⊾↓⋅⋖⋅∼⊼⊲↓⊳∖⇂↓↕∢⊾⋖⋅⊔∢⊾↓⋅⋏∙≟∙∖⇁∐⋯⋅⊔≼↛⋖⋅∪⇂⋅↿↓↥⋖⋅∐, The total number of soft FUV photons reaching us (in unit area and without absorption) can be estimated by where ${\cal F}$ is the energy fluence of the flare. +⋜⊔⋅⋖⊾⊳ The interaction between the photon beam and the isotropic relativistic electrons. (i.c.. the anisotropic LC scattering) has been discussed externsively (Brunetti 2001 ancl the references therein).," The interaction between the photon beam and the isotropic relativistic electrons (i.e., the anisotropic IC scattering) has been discussed externsively (Brunetti 2001 and the references therein)." +" Here the scattering is in the Thompson regime anc the electrons. are. ultra-relativistic (their distribution is n(5.)=A5 ΠΠ t$." +" anopumal algorithm. For any 1> 2, the strategy isa,,,-compe"," Hence, at any time a buffer can store the candidate jobs to be migrated." +"llive, where o, isthe solution ofan equation representing load in a"," On the other hand, to the best of our knowledge, the algorithms by Englert et al." +nideal machine profile forasubset of the jobs. Forij!, \cite{EOW} do not translate into strategies with job migration. +" = 2, the competitive ratio15 4/3, Theratios are non-decreasing and converge toM 4(—1/« 2)/AOAM. 167 leES 1.1659 asiitends"," All the algorithms in\cite{EOW} use the given buffer of size $cm$ , for some constant $c$, to store the $cm$ largest jobs of the job sequence." +" toinfinity. Again, V 4isthe lower branch ofthe Lambert VV function."," However in our setting, a migration of the largest jobs does not generate good schedules." +" The algorithm uses Πω. 7) + Ljjob migrations. For i >11,this expression isatmost 777.For smaller m"," The problem are shrinking jobs, jobs that are among the largest jobs at some time $t$ but not at later times." +achine numbers1tis5r» to10). Wenote that the competi,"We cannot afford to migrate all shrinking jobs, unless we invest $\Theta(n)$ migrations." +tiveness of1.4659 isconsiderably below the factorof roughly 1.9 obta," With limited job migration, scheduling decisions are final for almost all of the jobs." +inedby deterministic algorithms in the standard online setüng. Itis alsobe, Hence the corresponding algorithms are more involved than in the setting with a reordering buffer. +low theratioof e/(e—1) attainable if randomization or job preemption areallowed., For the description of the algorithm and the attained competitive ratio we define a function $f_m(\alpha)$. + In Secon??. we giveamatching lower migrations. Our algorithms relyona numberof newideas.," Intuitively, $f_m(\alpha)$ represents accumulated normalized load in a “perfect” machine profile for a subset of the jobs." +All strategies classifyincoming Jobs into small andlarge depending ona careful esumate onthe optimum makespan. Thealgorithms consist ofa jobarrival phase followedbyamigrationphase.Theoptimal," In such a profile the load ratios of the first $\lfloor m/\alpha\rfloor$ machines follow a Harmonic series of the form $(\alpha-1)/(m-1), \ldots, (\alpha-1)/(m-\lfloor m/\alpha\rfloor)$ while the remaining ratios are $\alpha/m$." +" algorithm, inthe arrivalphase, maintains aload", Summing up these ratios we obtain $f_m(\alpha)$. + profile on the machines with respectto jobsthat main challenge," Formally, let for any machine number $m\geq 2$ and real-valued $\alpha>1$." + inthe analyses of thevarious algorithmsis toboundthe numberof jobs that havetobe migrated," Here $H_k = \sum_{i=1}^k 1/i$ denotes the $k$ -th Harmonic number, for any integer $k\geq 1$." + [rom each mach, We set $H_0 = 0$. +ine. We finallyrelate our contributions to some exisüngresults.," For any fixed $m\geq 2$, let $\alpha_m$ be the value satisfying $f_m(\alpha)=1$." + First wepoint out thatthe goal inonline, Lemma \ref{lem:l1} below implies that $\alpha_m$ is well-defined. + minimization is to construct a good schedule, The algorithm we present is exactly $\alpha_m$ -competitive. +"whenjobs arrive oneby one.Once theschedule isconstructed,theproce"," By Lemma \ref{lem:l2}, the values $\alpha_m$ form a non-decreasing sequence." +ssingofjobsmaystart.Itnot sapulatedthat machines start executing jobs while other Jobsof& sull need tobe scheduled., There holds $\alpha_2 = 4/3$ and $\lim_{m\rightarrow \infty} \alpha_m = W_{-1}(-1/e^2)/(1+ W_{-1}(-1/e^2))\approx 1.4659$. + This frameworkis assumedin all the l, This convergence was also stated by Englert et al. +iterature ononline given.fac, \cite{EOW} but no thorough proof was presented. +"tThe optimalcompetitive ratioandof thata,,,in1sthe maintains", The following two technical lemmas are proven in the appendix. +olutionof certainanequationload that profilealso the arises machines.," Let $m\geq 2$ and $M_1, \ldots, M_m$ be the available machines." +in|8].. ThisOuris dueto doesthe thatour," Furthermore, let $\alpha_m$ be as defined above." + algorithm small while[8]. considers on machines.strategyIn so [rameworkw.r.t. gobs that are currentlyharder maintain becausethe strateg," The algorithm, called, operates in two phases, a and a. In the job arrival phase all jobs of $\sigma = J_1, \ldots, J_n$ are assigned one by one to the machines." +yin |8] shrinking jobs. alljjobsjobsthat assignedt, In this phase no job migrations are performed. +o time ourbutsmall latertheprofile is to job phaseof i.e.," Once $\sigma$ is scheduled, the job migration phase starts." + reschedules jobsare removedlargeat fromsome / machines., First the algorithm removes some jobs from the machines. + Thisat operationtimes{ή> /.In correspondsthe the migrat, Then these jobs are reassigned toother machines. +"ion“finalphase""our the algorithmin policies by of in[8] |8].. virtualHowever, schedule.our algorithm directly applies Graham |12,"," In thisphase classifies jobs into small and large and, moreover, maintains a load profile with respect to the small jobs on the machines." + 13] whilethe algorithm computes a, At any time the load of a machine is the sum of +"with the |?CO PdBI contours allow to investigate the nuclear region ofNGC3627,, while the torques derived from the Spitzer--IRAC image in combination withthe ""CO BIMA contours are much better adapted to visualize the whole spiral structure of the galaxy.","with the $^{12}$ CO PdBI contours allow to investigate the nuclear region of, while the torques derived from the -IRAC image in combination withthe $^{12}$ CO BIMA contours are much better adapted to visualize the whole spiral structure of the galaxy." +" We perform the subtraction of foreground stars, deprojection, and resampling, as described in other NUGA papers (e.g.,García-Burilloetal.2005)."," We perform the subtraction of foreground stars, deprojection, and resampling, as described in other NUGA papers \citep[e.g.,][]{santi05}." +". Here, we briefly recall some definitions and assumptions used to evaluate the gravitational torques."," Here, we briefly recall some definitions and assumptions used to evaluate the gravitational torques." +" NIR images are completed in the vertical dimension by assuming an isothermal plane model with a constant scale height, equal to ~1/12th of the radial scale-length of images."," NIR images are completed in the vertical dimension by assuming an isothermal plane model with a constant scale height, equal to $\sim$ 1/12th of the radial scale-length of images." +" With a Fourier transform method we derive the potential and we assume a constant mass-to-light (M/L) ratio able to reproduce the observed ""CO RCs.", With a Fourier transform method we derive the potential and we assume a constant mass-to-light (M/L) ratio able to reproduce the observed $^{12}$ CO RCs. +" Beyond a radius of 20""(or kkpc in diameter), the mass density is set to 0 in the HST--NICMOS F160W image, thus suppressing any spurious m=4 terms."," Beyond a radius of (or kpc in diameter), the mass density is set to 0 in the -NICMOS F160W image, thus suppressing any spurious $m=4$ terms." + This assumption is sufficient to compute the potential over the PdBI '*CO(1-0) primary beam., This assumption is sufficient to compute the potential over the PdBI $^{12}$ CO(1–0) primary beam. +" For the Spitzer--IRAC iimage, this radius truncation is done at 169""(or kkpc in diameter)."," For the -IRAC image, this radius truncation is done at (or kpc in diameter)." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thet"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Theto"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetor"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorq"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorqu"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorque"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquem"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquema"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemap"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapi"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapis"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapiso"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisor"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisori"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorie"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorien"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorient"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisoriente"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisoriented"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorienteda"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedac"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedacc"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedacco"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccor"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccord"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordi"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordin"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccording"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordingt"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordingto"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordingtot"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordingtoth"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +" For the non-axisymmetric part of the potential ®(R,0), we decompose O(R,8) in Fourier components (m-modes), following Combes&Sanders (1981):: 8) =DR) + ®,,(R) t(x,y) 2 x F,—y F,.Thetorquemapisorientedaccordingtothe"," For the non-axisymmetric part of the potential $\Phi(R,\theta)$, we decompose $\Phi(R,\theta)$ in Fourier components $m$ -modes), following \citet{francoise81}: : $$ ) = _0(R) + _m(R) t(x,y) = x F_y -y F_x." +history.,history. + The isochrones allow us to get an estimate of the mass and radius of component Aa using (he spectroscopic gravity. although details of the distribution of the stars chemicals introduce some uncertainty (e.g. Sills et al.," The isochrones allow us to get an estimate of the mass and radius of component Aa using the spectroscopic gravity, although details of the distribution of the star's chemicals introduce some uncertainty (e.g. Sills et al." + 1997)., 1997). + This is possible because of the relation that exists along each isochrone. with vounger isochrones giving the overall largest mass and radius values.," This is possible because of the mass-radius relation that exists along each isochrone, with younger isochrones giving the overall largest mass and radius values." +" We find nearly identical ranges from both the Yietal. and Girardietal.(2000). isochrones: Ma,=1.242 and Ry,=137miH..."," We find nearly identical ranges from both the \citet{yy} and \citet{gir00} + isochrones: $M_{Aa} = 1.24^{+0.19}_{-0.08}$ and $R_{Aa} = +1.37^{+0.24}_{-0.16}$." +" We can also estimate the Z-band magnitude of component Aa using the isochrones and (he spectroscopic measurements. and again we find similar ranges from the two sets of isochrones (Mj,44=3.08EN )"," We can also estimate the $I$ -band magnitude of component Aa using the isochrones and the spectroscopic measurements, and again we find similar ranges from the two sets of isochrones $M_{I,Aa} = 3.08^{+0.20}_{-0.28}$ )." + The primary source of uncertainty in all three estimates is from the measurement of the surface gravitv., The primary source of uncertainty in all three estimates is from the measurement of the surface gravity. + In order to make (he final comparison wilh our observed values [rom (he eclipse analysis. we use the models of Daraffe et al. (," In order to make the final comparison with our observed values from the eclipse analysis, we use the models of Baraffe et al. (" +1998; herealter DCAII) since they eurrently provide the best fit to the lower main sequence of M6T (although {μον begin to become too blue by Vzz17; Sandquist 2003).,"1998; hereafter BCAH) since they currently provide the best fit to the lower main sequence of M67 (although they begin to become too blue by $V +\approx 17$; Sandquist 2003)." + The DCAII models are also the ones that have been most thoroughly compared to observed colors of low-mass stus. and seem (o agree well down to Tigy~3600 IK (DCAID).," The BCAH models are also the ones that have been most thoroughly compared to observed colors of low-mass stars, and seem to agree well down to $T_{eff} \sim 3600$ K (BCAH)." + The isochrones of the Y? and Padova groups employ different tables to convert [rom 7;y; to color. and get widely different results.," The isochrones of the $^2$ and Padova groups employ different tables to convert from $T_{eff}$ to color, and get widely different results." + The BCAIL models most jearly agree with those of the Y? or Padova groups near solar mass. where the agreement is orced by calibration to the Sun.," The BCAH models most nearly agree with those of the $^2$ or Padova groups near solar mass, where the agreement is forced by calibration to the Sun." + We made a correction of 0.077 mag in { to account for differences between the Padova and Y? isochrones and those of DCAL for solar-mass nodels., We made a correction of 0.077 mag in $I$ to account for zero-point differences between the Padova and $^2$ isochrones and those of BCAH for solar-mass models. + Such a correction is justified by current uncertainties in the zero-points of T;yj-color relations lor different sets of isochrones., Such a correction is justified by current uncertainties in the zero-points of $T_{eff}$ -color relations for different sets of isochrones. + Our final comparison is shown in Fie. 8.., Our final comparison is shown in Fig. \ref{eclcomp}. + The ligure indicates that component Ab is smaller and/or brighter in J band than the models predict., The figure indicates that component Ab is smaller and/or brighter in $I$ band than the models predict. + The uncertainties in the spectroscopic properties of component Aa mean. however. that the inconsistency is only al a little over lo.," The uncertainties in the spectroscopic properties of component Aa mean, however, that the inconsistency is only at a little over $1 \sigma$." +" A reasonable estimation of the effective temperature of star Ab is next to impossible. partly because there are not enough well-measured cool dwarf stars to set an empirical T,;- relation or calibrate a theoretical one."," A reasonable estimation of the effective temperature of star Ab is next to impossible, partly because there are not enough well-measured cool dwarf stars to set an empirical $T_{eff}$ -color relation or calibrate a theoretical one." + In addition. there is some evidence that the giant stars (vpically used (ο calibrate relations at low temperatures may follow a systematically different relation than the clwarls (IIoudasheltetal.," In addition, there is some evidence that the giant stars typically used to calibrate relations at low temperatures may follow a systematically different relation than the dwarfs \citep{houda}." +2000).. The depth of secondary eclipse is sensitive to T;yy. but the interpretation requiresthe use of updated stellar atmosphere models because of strongly non-grey nature of the atiospheres of low-mass stars (e.g. Chabrier Baralle 1997).," The depth of secondary eclipse is sensitive to $T_{eff}$, but the interpretation the use of updated stellar atmosphere models because of strongly non-grey nature of the atmospheres of low-mass stars (e.g. Chabrier Baraffe 1997)." +" From the BCAIL models and a distance modulus (00—M);=9.67&0.03 (Sandquist 2003). we roughly estimate that Zip,=31002:200 Ix. but we again note that this is near where their models start to diverge from color-magnitude diagrams of field dwarls."," From the BCAH models and a distance modulus $(m - M)_{I} = 9.67 \pm 0.03$ (Sandquist 2003), we roughly estimate that $T_{eff} = 3700 \pm 200$ K, but we again note that this is near where their models start to diverge from color-magnitude diagrams of field dwarfs." +vanish: at the other end. o; decreases for weaker and weaker shocks corresponding to smaller Ac and/or relatively stronger preheating levels.,"vanish; at the other end, $a_R$ decreases for weaker and weaker shocks corresponding to smaller $\Delta\phi$ and/or relatively stronger preheating levels." + As no other major sources or sinks of entropy occur from the boundaryat a few Mpes down to the central r10° kpe. in the outer range the entropy is deposited and during the stage of slow accretion: thus a powerlaw radial run is set with slope a(r) that stays close to its boundary value CGR.," As no other major sources or sinks of entropy occur from the boundary at a few Mpcs down to the central $r\sim 10^2$ kpc, in the outer range the entropy is deposited and during the stage of slow accretion; thus a powerlaw radial run is set with slope $a(r)$ that stays close to its boundary value $a_R$." + At the center. instead. entropy may be intermittently injected by shocks driven by mergers reaching down there (McCarthy et al.," At the center, instead, entropy may be intermittently injected by shocks driven by mergers reaching down there (McCarthy et al." + 2007: Balogh et al., 2007; Balogh et al. + 2007). and by powerful AGNS residing in the central massive galaxies (see Valageas Silk 1999; Wu et al.," 2007), and by powerful AGNs residing in the central massive galaxies (see Valageas Silk 1999; Wu et al." + 2000; Scannapieco Oh 2004; Lapi et al., 2000; Scannapieco Oh 2004; Lapi et al. + 2005) as observed and reviewed by MeNamara Nulsen (2007) and Markevitch Vikhlinin (2007)., 2005) as observed and reviewed by McNamara Nulsen (2007) and Markevitch Vikhlinin (2007). + Thus the full entropy profile that combines central injections. with outer stratification may be described by the simple parametric expression (see Voit 2005 and references therein) in fact. this approaches a constant value at small radit. and smoothly goes into a powerlaw at large radii.," Thus the full entropy profile that combines central injections with outer stratification may be described by the simple parametric expression (see Voit 2005 and references therein) in fact, this approaches a constant value at small radii, and smoothly goes into a powerlaw at large radii." + Entropy profiles similar to Eg. (, Entropy profiles similar to Eq. ( +5) have been recently reported by Cavagnolo et al. (,5) have been recently reported by Cavagnolo et al. ( +2009) from an analysis of a archival sample comprising 239 clusters.,2009) from an analysis of a archival sample comprising $239$ clusters. + In fact. most of them are well fit by a power law at large radii plus a constant value peakedkozκ. at small radii. a basically bimodal distributioncm. at kyz20 keV em? withand at kyz150 keV ," In fact, most of them are well fit by a power law at large radii plus a constant value $k_0\ga k_c$ at small radii, with a basically bimodal distribution peaked at $k_0\approx 20$ keV $^{2}$ and at $k_0\approx 150$ keV $^{2}$." +On the other hand. some clusters show evidence of a sharper entropy floor (e.g.. Pratt et al.," On the other hand, some clusters show evidence of a sharper entropy floor (e.g., Pratt et al." + 2005 for A2218: see also Fig., 2005 for A2218; see also Fig. + 5 in Cavagnolo et al., 5 in Cavagnolo et al. + 2009) that we represent not only with à level & but also with a definite extension ry. so that the corresponding radial entropy run reads for rry.," 2009) that we represent not only with a level $k_c$ but also with a definite extension $r_{f}$, so that the corresponding radial entropy run reads for $\bar{r}\leq \bar{r}_{f}$, and as for $\bar{r}>\bar{r}_{f}$." + The scale r; may be interpreted as the terminal radius just reached by an outbound blastwave driven by a violent energy input at the center (see Lapi et al., The scale $r_f$ may be interpreted as the terminal radius just reached by an outbound blastwave driven by a violent energy input at the center (see Lapi et al. +" 2005): at r, its decreasing Mach number ντο) has decayed to unity and the blast has stalled and degraded into adiabatie sound waves. as caught in action by Fabian et al. ("," 2005); at $r_f$ its decreasing Mach number $\mathcal{M}(r)$ has decayed to unity and the blast has stalled and degraded into adiabatic sound waves, as caught in action by Fabian et al. (" +2006) in the Perseus Cluster (see also 5 for a discussion).,2006) in the Perseus Cluster (see also 5 for a discussion). + In this picture. Eq. (," In this picture, Eq. (" +5) represents a later stage caused by diffusive smoothing and mixing of such à radial imprint. while the enhanced entropy level is still high before radiative erosion has set in.,"5) represents a later stage caused by diffusive smoothing and mixing of such a radial imprint, while the enhanced entropy level is still high before radiative erosion has set in." + From the gas temperature and density profiles given by (, From the gas temperature and density profiles given by Eq. ( +2) we also derive the distribution of the gravitating mass in Eq.the form (see Sarazin 1988) where the quantity riz67) grows slowly with 7 and saturates to unity.,"2) we also derive the distribution of the gravitating mass in the form (see Sarazin 1988) where the quantity $\bar{r}\,\bar{v}^2_c(r)$ grows slowly with $\bar{r}$ and saturates to unity." +" Here we describe in detail how we use the SM to account for the X-ray brightness and temperature profiles for both classes of CC and NCC clusters (Molendi Pizzolato 2001: Leccardi Molendi 2009): in. particular. we will examine in detail the following six clusters: A2199. A2597, AI689. ΑΙό56. A2256. A644."," Here we describe in detail how we use the SM to account for the X-ray brightness and temperature profiles for both classes of CC and NCC clusters (Molendi Pizzolato 2001; Leccardi Molendi 2009); in particular, we will examine in detail the following six clusters: A2199, A2597, A1689, A1656, A2256, A644." + These have been collected from the literature on the basis of the quality and detail of the X-ray data. keeping a balance between the CC and NCC classes.," These have been collected from the literature on the basis of the quality and detail of the X-ray data, keeping a balance between the CC and NCC classes." + We fit with the SM the profiles of emission weighted temperature and X-ray brightness. as given by Eqs. (," We fit with the SM the profiles of emission weighted temperature and X-ray brightness, as given by Eqs. (" +9) and (10) below.,9) and (10) below. + The parameter values are pinned down on using a standard 4 minimization procedure. and the uncertainties are quoted at the 68% confidence level.," The parameter values are pinned down on using a standard $\chi^2$ minimization procedure, and the uncertainties are quoted at the $68\%$ confidence level." + The free parameters are set as follows., The free parameters are set as follows. + The DM halo distributions depend weakly on the index « that we fix at the value «àz1.27. and are strongly marked by the c that we leave as a free parameter.," The DM halo distributions depend weakly on the index $\alpha$ that we fix at the value $\alpha\approx 1.27$ , and are strongly marked by the $c$ that we leave as a free parameter." +" The ICP profiles are parameterized by the a and me value k, that define the radial entropy run after Eq. (", The ICP profiles are parameterized by the $a$ and the value $k_c$ that define the radial entropy run after Eq. ( +5):s in some clusters an acceptable fit requires to introduce the ry of the central floor after Eqs. (,5); in some clusters an acceptable fit requires to introduce the $r_{f}$ of the central floor after Eqs. ( +6)withand (7).,6) and (7). + In the SM as all models the virial radius Α intervenes to set the data distance scale. as well as the bound to Lo.s.," In the SM as with all models the virial radius $R$ intervenes to set the data distance scale, as well as the bound to l.o.s." +" integrations: we keep it fixed when a robust determination 1s available from the literature. e.g.. from observations of galaxy dynamies. or ""red sequence’ termination. or gravitational lensing."," integrations; we keep it fixed when a robust determination is available from the literature, e.g., from observations of galaxy dynamics, or `red sequence' termination, or gravitational lensing." + Otherwise. we determine it from fitting (in the range where the data are reliable) the projected profile of the emission-weighted temperature in terms of the projected radius w—w/R.," Otherwise, we determine it from fitting (in the range where the data are reliable) the projected profile of the emission-weighted temperature in terms of the projected radius $\bar{w}\equiv w/R$." + In the following we approximate the detailed cooling function with A(T)x as appropriate for the ICP in hot clusters., In the following we approximate the detailed cooling function with $\Lambda(T)\propto T^{1/2}$ as appropriate for the ICP in hot clusters. +" The above relation not only pins down the (horizontal) scale R. but also sets the normalization (vertical) scale 7,5."," The above relation not only pins down the (horizontal) scale $R$, but also sets the normalization (vertical) scale $T_R$." +" The knowledge of R and 7 allows us to derive the ICP density 7/4 from fitting the brightness distribution here $5223.4«107(14z fem s en? aremin™. and the factor ΕΕ.Es.T)=ebar—e0T takes into account specific instrumental bands E»—E, (e.g.. Ettori 2000)."," The knowledge of $R$ and $T_R$ allows us to derive the ICP density $n_R$ from fitting the brightness distribution here $S_0\approx 3.4\times 10^{-13}\, (1+z)^{-4}$ erg $^{-1}$ $^{-2}$ $^{-2}$, and the factor $F(E_1,E_2,T)\simeq +e^{-E_1/k_BT}-e^{-E_2/k_BT}$ takes into account specific instrumental bands $E_2-E_1$ (e.g., Ettori 2000)." +" The SM actually predicts the values of 74. jg and ke="" from extrapolating the profiles into the outer cluster nsregions: observing the latter challenges the sensitivity and defies the resolution of most current instruments. but will constitute a main target for the next-generation X-ray telescopes planned to study low surface brightness plasmas (see 5)."," The SM actually predicts the values of $T_R$, $n_R$ and $k_R=k_B T_R/n_R^{2/3}$ from extrapolating the profiles into the outer cluster regions; observing the latter challenges the sensitivity and defies the resolution of most current instruments, but will constitute a main target for the next-generation X-ray telescopes planned to study low surface brightness plasmas (see 5)." +" With the use of such facilities R will be read out from the as the position of the shock discontinuitiesin(7) and T(r) 2003),profilesabove the values prevailing in the IGM (see Lapi et al.", With the use of such facilities $R$ will be read out from the profiles as the position of the shock discontinuities in$n(r)$ and $T(r)$ above the values prevailing in the IGM (see Lapi et al. + even though such discontinuities may be blended by," 2005), even though such discontinuities may be blended by" +accounted for.,accounted for. + That the tov model represents the σας fraction best for halos beginning al 2=10 adds support to the idea that the Milkv. Way. clwarls were formed at or before this redshift (Luxetal.2010)., That the toy model represents the gas fraction best for halos beginning at $z=10$ adds support to the idea that the Milky Way dwarfs were formed at or before this redshift \citep{Lux2010}. +. This simple tov model reproduces the observed. [raction of gas-rich. cdwarls. however. there are several [actors that were assumed to be negligible which could affect the amount ol gas that survives.," This simple toy model reproduces the observed fraction of gas-rich dwarfs, however, there are several factors that were assumed to be negligible which could affect the amount of gas that survives." + The accretion of gas onto dwarls is unaccounted [or in this mocel., The accretion of gas onto dwarfs is unaccounted for in this model. + The prospect of low redshift accretion (Ricotti2009) in particular would greatly increase the chance of a dwarl surviving with eas to (he present day. and if included for all dwarfs would likely result in an overestimation of the gas fraction even for dwarls beginning al late redshifts.," The prospect of low redshift accretion \citep{Ricotti2009} in particular would greatly increase the chance of a dwarf surviving with gas to the present day, and if included for all dwarfs would likely result in an overestimation of the gas fraction even for dwarfs beginning at late redshifts." + The effects of dust through photoelectric heating aud cooling is also ignored., The effects of dust through photoelectric heating and cooling is also ignored. + Due to the low metallicity environments that dwarf galaxies (vpically have. this effect. will be smaller than in larger galaxies. but may still be an important source of heating or cooling.," Due to the low metallicity environments that dwarf galaxies typically have, this effect will be smaller than in larger galaxies, but may still be an important source of heating or cooling." + We also assume (hat (he early extragalactic UV field is uniform in space. a elunipy radiation lield around the time of reionization. may have a large impact on the amount of eas that remains cold ancl protected [rom ram pressure stripping. (his early gas loss could greatly impact the survival around (he earliest pericentre passages.," We also assume that the early extragalactic UV field is uniform in space, a clumpy radiation field around the time of reionization, may have a large impact on the amount of gas that remains cold and protected from ram pressure stripping, this early gas loss could greatly impact the survival around the earliest pericentre passages." + The use of a smooth medium for the gas.compared to a more realistic fractal mediumminimises the cooling of warm eas and allows it to extend bevond the tidal radius., The use of a smooth medium for the gas—compared to a more realistic fractal medium—minimises the cooling of warm gas and allows it to extend beyond the tidal radius. + Even with a large filling factor. the gas will cool much more quickly via metal line cooling than in the smooth medium mininising (he large gas loss at the beeinning.," Even with a large filling factor, the gas will cool much more quickly via metal line cooling than in the smooth medium minimising the large gas loss at the beginning." + Potentially the biggest limitation is (he assumption that all of the supernova energv goes into heating (he eas (o an extremely. hot state. much of (his energy likely goes into raising cold and warm gas oul of the potential well of the clwarl. allowing it to be much more easily stripped.," Potentially the biggest limitation is the assumption that all of the supernova energy goes into heating the gas to an extremely hot state, much of this energy likely goes into raising cold and warm gas out of the potential well of the dwarf, allowing it to be much more easily stripped." + In particular this will predominantly impact the lowest gas masses. where the smooth eas assumption versus a fractal medium is more likely to have a large impact (Bland-Lawthornetal.2007).," In particular this will predominantly impact the lowest gas masses, where the smooth gas assumption versus a fractal medium is more likely to have a large impact \citep{JBH2007}." +. Due to the large initial gas loss. our star formation rates do not represent (hat of the majority of cwarfs. with most stars forming early in the cdwarls life. as opposed to a roughly continuous star formation rate with some bursts (Weiszetal.2011).," Due to the large initial gas loss, our star formation rates do not represent that of the majority of dwarfs, with most stars forming early in the dwarfs life, as opposed to a roughly continuous star formation rate with some bursts \citep{Weisz2011}." +. This makes our model more suited to explaining dwarls with these early periods of star formation where a majority ol starsare formed. e.g. DINBN. KDG52 (Weiszetal. 2011)..," This makes our model more suited to explaining dwarfs with these early periods of star formation where a majority of starsare formed, e.g. BK5N, KDG52 \citep{Weisz2011}. ." +(Reiners&Basi2006).. but this approach is vet to be implemented for anv star bevoud ALLS.,"\citep{rb06}, but this approach is yet to be implemented for any star beyond M4.5." + The presence aud dissipation of magnetic fields can © alternatively traced through activity indicators such as Πα. N-vay and radio emission.," The presence and dissipation of magnetic fields can be alternatively traced through activity indicators such as $\alpha$, X-ray and radio emission." + The Πα sud N-rav Cluission are secondary iudicators since they arise frou dlasina prestunably heated by the dissipation of maguctic fields. through for cxample maguetie reconnection.," The $\alpha$ and X-ray emission are secondary indicators since they arise from plasma presumably heated by the dissipation of magnetic fields, through for example magnetic reconnection." + Iu he standard scenario the input of onergv drives an outflow of hot plasma iuto the corona through evaporation of the underlvine chromosphere. leading iu urn to bremsstraliluug X-ray cussion and Te ciission (Neupert1968:Tawleyetal.1995:Cmedol1996).," In the standard scenario the input of energy drives an outflow of hot plasma into the corona through evaporation of the underlying chromosphere, leading in turn to bremsstrahlung X-ray emission and $\alpha$ emission \citep{neu68,hfs+95,gbs+96}." +. The radio cussion. ou the other laud. arises frou evroresonance or coherent processes. which trace the preseuce aud properties of maguetic fields directly.," The radio emission, on the other hand, arises from gyroresonance or coherent processes, which trace the presence and properties of magnetic fields directly." + Thus. radio observatious can be used to iufer the field streneth of individual objects directly. whereas Ho aud X-ray euission provide a useful statistical measure and insight iuto the influcuce of magnetic fields on the outer lavers of dwiuf stars.," Thus, radio observations can be used to infer the field strength of individual objects directly, whereas $\alpha$ and X-ray emission provide a useful statistical measure and insight into the influence of magnetic fields on the outer layers of dwarf stars." + Observationalh. the lack of siguificaut change in the ueasured level of Πα aud N-rav activity with the onse of full convection at about spectral type MD. suggests hat at least in the earl-M chwarfs the putative turbuleu or distributed dynamo can operate effiiientlv.," Observationally, the lack of significant change in the measured level of $\alpha$ and X-ray activity with the onset of full convection at about spectral type M3, suggests that at least in the early-M dwarfs the putative turbulent or distributed dynamo can operate efficiently." + It is also yossible that the magnetic field itself acts to reduce the nass at which the transition to fully couvective structure akes place (Atlan&MacDonald2001)., It is also possible that the magnetic field itself acts to reduce the mass at which the transition to fully convective structure takes place \citep{mm01}. +". However. ονομα, spectral type MT there is a precipitous drop imn Te and X-rav persistent activity. aud ouly a few perceut of the objects exhibit flares (e.g... Reidetal.1999:CazisWestetal. 2001))."," However, beyond spectral type M7 there is a precipitous drop in $\alpha$ and X-ray persistent activity, and only a few percent of the objects exhibit flares (e.g., \citealt{rkg+99,gmr+00,rbm+00,lkc+03,whw+04}) )." + Furthermore. unlike iu the eurlv-M dwarts (Rosneretal.1985:Fleming1993:Mohantyetal.2002:Pizzolato 2003).. nianu late-type rapidly rotating dwarts exhibit little or no discernible activity iu these hands (Basti&Marcy1995:MobautyBasri2003).," Furthermore, unlike in the early-M dwarfs \citep{rgv85,fgs+93,mbs+02,pmm+03}, many late-type rapidly rotating dwarfs exhibit little or no discernible activity in these bands \citep{bm95,mb03}." +. These patterus are consistent with either a decrease in the eeneration or dissipation of the magnetic fields. or both.," These patterns are consistent with either a decrease in the generation or dissipation of the magnetic fields, or both." + Ou the other hand. radio cimission has been detected from several late-M. aud L disvarfs (Bergerotal.2001:2005) sugeesting that at least some of these objects are capable of eeneratiug and dissipating magnetic fields.," On the other hand, radio emission has been detected from several late-M and L dwarfs \citep{bbb+01,ber02,brr+05,bp05} suggesting that at least some of these objects are capable of generating and dissipating magnetic fields." + Surprisingly. the ratio of radio to N-vav huninositv in the detected objects exceeds by. several orders of magnitude the value measured for carly M chwarts and a variety of other stars (includiug the Sun: Guecel&Benz1993:&Cuedel 1991)). and there is uo obvious correlation with IIo emission.," Surprisingly, the ratio of radio to X-ray luminosity in the detected objects exceeds by several orders of magnitude the value measured for early M dwarfs and a variety of other stars (including the Sun; \citealt{gb93,bg94}) ), and there is no obvious correlation with $\alpha$ emission." + Thus. radio observations present a powerful. and perhaps unique approach for iuferiiie the magnetic field streugth of late-tvpe stars and brow1 dwarts.," Thus, radio observations present a powerful, and perhaps unique approach for inferring the magnetic field strength of late-type stars and brown dwarfs." + ere we exploit this approach and continue our investigation of radio cussion frou late-M. L aud T dwarts by expanding the observed sample o about a factor of three (to 90 sources).," Here we exploit this approach and continue our investigation of radio emission from late-M, L and T dwarfs by expanding the observed sample by about a factor of three (to $90$ sources)." + With this extended: saunple we find continued evidence for a sharp transition iu the ratio of radio to N-rav. huuinosity at spectral ype MT. as well as an increased level of activity with later spectral type.," With this extended sample we find continued evidence for a sharp transition in the ratio of radio to X-ray luminosity at spectral type M7, as well as an increased level of activity with later spectral type." + We show. however. that as in the case of Πα aud X- observations. the fraction of objects producing radio cluission drops from about 30% in the AI chwarfs to oulv ~5% in the L dwarts.," We show, however, that as in the case of $\alpha$ and X-ray observations, the fraction of objects producing radio emission drops from about $30\%$ in the M dwarfs to only $\sim 5\%$ in the L dwarfs." + Most importantly. we preseut for the first time estimates of the maguetic field. strength of a laree sample of late-M and L chwarts. aud show that for the active sour‘es there is at most a modest drop iu the field streneth f£on earlv-M to early-L dwarts.," Most importantly, we present for the first time estimates of the magnetic field strength of a large sample of late-M and L dwarfs, and show that for the active sources there is at most a modest drop in the field strength from early-M to early-L dwarfs." + We observec a snple of 21 late-M. L and T dwarts with the Verv Large Array Iuc.)) at 8.16 CIIz using the standard continu mode with 2«50 MIIz contiguous bands at each frequency.," We observed a sample of $21$ late-M, L and T dwarfs with the Very Large Array ) at 8.46 GHz using the standard continuum mode with $2\times 50$ MHz contiguous bands at each frequency." + The flux. density scale was determüned usiie the standard extragalactic calibrator sources 3€ [s (JOIST|331). 3C 117 (J0512|198) aud 3€ 286 (J1331|305). while the phase was monitored using calibrators locaed within LO° of the targets sources.," The flux density scale was determined using the standard extragalactic calibrator sources 3C 48 (J0137+331), 3C 147 (J0542+498) and 3C 286 (J1331+305), while the phase was monitored using calibrators located within $10^\circ$ of the targets sources." + The data were reduced and analyzed using the Astronomical huage Processiic System., The data were reduced and analyzed using the Astronomical Image Processing System. + Iu addition to our observations we obtained auk reduced all publicly available observations of late-AL L αιid Todwarfs from the VLA archive. and collected all τιeasurenients published in the literature.," In addition to our observations we obtained and reduced all publicly available observations of late-M, L and T dwarfs from the VLA archive, and collected all measurements published in the literature." + This resulted 11a total 58 objects rangiug from MT. to Tea. as well a sugle ALS dwart and a single M55 chwart.," This resulted in a total 88 objects ranging from M7 to T8, as well a single M5 dwarf and a single M5.5 dwarf." + A sunuuarv of all observations and the relevant source properties are eiveu in Table 1 aud Figure 1.., A summary of all observations and the relevant source properties are given in Table \ref{tab:obs} and Figure \ref{fig:prop}. + For the deteected objects we searched for variability (flares) using the followiug method., For the detected objects we searched for variability (flares) using the following method. +" We removed all the bright feld sources using the AIPS/IMAGR routine to CLEAN the region around each source (vith the exception of the target xmace), and the AIPS/UVSUB routine to subtract the resulting source models from the visibility data."," We removed all the bright field sources using the AIPS/IMAGR routine to CLEAN the region around each source (with the exception of the target source), and the AIPS/UVSUB routine to subtract the resulting source models from the visibility data." + We theu plottec the real part of the complex visibilitics at the position of the science target as a fuuctiou of time using the AIPS/UVPLT routine., We then plotted the real part of the complex visibilities at the position of the science target as a function of time using the AIPS/UVPLT routine. + The subtraction of field. sources Is necessary since their sidelobes and the change in the shape of the svuthesized beam during the observation restut in flux variations over the map. which may contaminate auv real variability or generate false variability.," The subtraction of field sources is necessary since their sidelobes and the change in the shape of the synthesized beam during the observation result in flux variations over the map, which may contaminate any real variability or generate false variability." + Quiescent raclio e1ussiou has been previously detected roni six late-type dwarfs ranging from spectral type M? o L3.5 (Berecretal.2001:Berger2002:Berecr2005:Bureasser&PutmanOstenctal. 2006).," Quiescent radio emission has been previously detected from six late-type dwarfs ranging from spectral type M7 to L3.5 \citep{bbb+01,ber02,brr+05,bp05,ohb+06}." +. Four of these objects also produced. short-lived. liehh-volarized flares with a typical timescale of ~10 iin and a fux increase compared to the quiescent level of at cast a factor of few.," Four of these objects also produced short-lived, highly-polarized flares with a typical timescale of $\sim 10$ min and a flux increase compared to the quiescent level of at least a factor of few." + In addition. the L3.5 dwarf 2MASS JOO3G61617|1821101 was shown to exhibit a periodicity of 3 hr in its quiescent radio cussion. whose origin is uot lly understood. but may arise from a closely-orbiting companion (Bergerctal.2005).," In addition, the L3.5 dwarf 2MASS J00361617+1821104 was shown to exhibit a periodicity of 3 hr in its quiescent radio emission, whose origin is not fully understood, but may arise from a closely-orbiting companion \citep{brr+05}." +. Iu the extended sample we detect radio enission frou lvee additional dwarf stars: LOS 1070 (MD5.5). LSR J1835(3259 (M8.5). and 2MÁAÀSS JO52338221103022 (L2.5) with fluxes of 161415 pJv. 525+15 py aud 231+11 py. respectively.," In the extended sample we detect radio emission from three additional dwarf stars: LHS 1070 (M5.5), LSR J1835+3259 (M8.5), and 2MASS $05233822-1403022$ (L2.5) with fluxes of $161\pm 15$ $\mu$ Jy, $525\pm +15$ $\mu$ Jy and $231\pm 14$ $\mu$ Jy, respectively." + LOS 1070 has heen detected on two separate occasions with fluxes of 153+23 aud, LHS 1070 has been detected on two separate occasions with fluxes of $153\pm 23$ and +"which, given the featureless nature of the SN continuum, is a realisticvalue?.","which, given the featureless nature of the SN pseudo-continuum, is a realistic." +". The results obtained for N=10'? cm""?, and SNR=100 for different values of FWHM, b (1, 3 and 5 km s! corresponding to EW of 58, 113 and 138 respectively ), and ó4 are presented in Table 1.."," The results obtained for $N$ $^{12}$ $^{-2}$, and =100 for different values of $FWHM$ , $b$ (1, 3 and 5 km $^{-1}$ corresponding to $EW$ of 58, 113 and 138 respectively ), and $\delta \lambda$ are presented in Table \ref{tab:err}." +" As the RMS errors are inversely proportional toSNR, these values can be readily scaled to different signal-to-noise ratios."," As the RMS errors are inversely proportional to, these values can be readily scaled to different signal-to-noise ratios." + These results have been checked against Monte-Carlo simulations and were found to be consistent to within a few 0.1mA., These results have been checked against Monte-Carlo simulations and were found to be consistent to within a few 0.1. +". Incidentally, this questions the need for a revision of the Chalabaev Maillard formula discussed by Vollmann Eversberg (2006))."," Incidentally, this questions the need for a revision of the Chalabaev Maillard formula discussed by Vollmann Eversberg \cite{vollmann}) )." + In the following we will consider an equivalent width variation AEW detectable if |JAEW|>5V20gw., In the following we will consider an equivalent width variation $\Delta EW$ detectable if $|\Delta EW|\geq 5\sqrt{2} \sigma_{EW}$. +" For a typical case where FWHM=7 km s!, b=1 km s!, 6A=0.01 pix-!, and SNR=100, this turns into a 5-σ detection limit AEWiim=4.4 (AE Wrim=5.3 for b-5 !)."," For a typical case where $FWHM$ =7 km $^{-1}$, $b$ =1 km $^{-1}$, $\delta +\lambda$ =0.01 $^{-1}$, and =100, this turns into a $\sigma$ detection limit $\Delta EW_{lim}$ =4.4 $\Delta EW_{lim}$ =5.3 for $b$ =5 $^{-1}$ )." +" Although the model can be used for any inter-stellar absorption line, in the following we present the results obtained for D», because it is a strong transition, it falls in a region almost free of telluric absorption features, and in a spectral interval where most optical, high-resolution spectrographs have their maximum sensitivity."," Although the model can be used for any inter-stellar absorption line, in the following we present the results obtained for $_2$, because it is a strong transition, it falls in a region almost free of telluric absorption features, and in a spectral interval where most optical, high-resolution spectrographs have their maximum sensitivity." + Example EW evolutions for two different cloud offsets xc (0 and 64 AU) and a number of cloud radii rc are presented in Fig., Example $EW$ evolutions for two different cloud offsets $x_C$ (0 and 64 AU) and a number of cloud radii $r_C$ are presented in Fig. + 3 up to 2 months after maximum light., \ref{fig:offsphere} up to 2 months after maximum light. +" In general, the maximum variability is expected when the cloud is close to the center of the photodisk."," In general, the maximum variability is expected when the cloud is close to the center of the photodisk." + The maximum variation is achieved for cloud radii between 64 and 128 AU., The maximum variation is achieved for cloud radii between 64 and 128 AU. +" Also, small clouds are better detected during the early phases (when their size is comparable to that of the photosphere), while the detection of large clouds requires a larger time span."," Also, small clouds are better detected during the early phases (when their size is comparable to that of the photosphere), while the detection of large clouds requires a larger time span." +" If the cloud is too large (rc 2512 AU), then the EW variation is not sufficiently ample to be detected."," If the cloud is too large $r_C>$ 512 AU), then the $EW$ variation is not sufficiently ample to be detected." +" With a minimum set of two observations taken 10 days apart in the pre-maximum phases, for typical high-resolution setup, a 5-c- detection limit of 4.4mA,, b=1 km s! and No=5x10!! cm~?, the simulations show that one is able to detect clouds with rc between 16 and 128 AU up to a maximum offset of 64 AU."," With a minimum set of two observations taken 10 days apart in the pre-maximum phases, for typical high-resolution setup, a $\sigma$ detection limit of 4.4, $b$ =1 km $^{-1}$ and $N_0$ $\times$ $^{11}$ $^{-2}$, the simulations show that one is able to detect clouds with $r_C$ between 16 and 128 AU up to a maximum offset of 64 AU." + For larger offsets the cloud starts to intersect the photodisk when its size is too large and the corresponding covering factor is too small., For larger offsets the cloud starts to intersect the photodisk when its size is too large and the corresponding covering factor is too small. +" Besides implying probably unrealistic density contrasts, increasing the column density does not enhance the detectability of a small size cloud, since this rapidly becomes totally opaque."," Besides implying probably unrealistic density contrasts, increasing the column density does not enhance the detectability of a small size cloud, since this rapidly becomes totally opaque." +" We note that, while this causes the saturation of the covering factor, it does not produce a saturated profile in the emerging absorption line."," We note that, while this causes the saturation of the covering factor, it does not produce a saturated profile in the emerging absorption line." +" A distinctive feature of small (rc<64 AU), offset clouds is an EW growth followed by a decrease during the pre-maximum light phase (Fig. 3,,"," A distinctive feature of small $r_C\leq$ 64 AU), offset clouds is an $EW$ growth followed by a decrease during the pre-maximum light phase (Fig. \ref{fig:offsphere}," + lower panel)., lower panel). + This is due to the growth of the covering factor as the photodisk starts intersecting the off-centered knot., This is due to the growth of the covering factor as the photodisk starts intersecting the off-centered knot. +" Once a maximum value is reached, the subsequent increase in the photodisk size causes the covering factor to drop."," Once a maximum value is reached, the subsequent increase in the photodisk size causes the covering factor to drop." +" Although this mechanism can produce an absorption feature which grows in strength and then disappears on timescales of a month, this is expected to happen only during the pre-maximum epochs, when the relative increase in the photodisk surface is very fast."," Although this mechanism can produce an absorption feature which grows in strength and then disappears on timescales of a month, this is expected to happen only during the pre-maximum epochs, when the relative increase in the photodisk surface is very fast." +" Larger clouds placed at larger offsets (xc 2256 AU) produce features which start appearing around maximum light, but keep steadily increasing in strength up to several months after maximum."," Larger clouds placed at larger offsets $x_C\geq$ 256 AU) produce features which start appearing around maximum light, but keep steadily increasing in strength up to several months after maximum." +" Finally, for central column densities (or density contrasts) smaller than ~2x10!° cm’, the EW variations are always below the detection limit."," Finally, for central column densities (or density contrasts) smaller than $\sim$ $\times$ $^{10}$ $^{-2}$, the $EW$ variations are always below the detection limit." +" Because of the stochastic geometry of the power-law clouds, their effects were evaluated using a statistical approach."," Because of the stochastic geometry of the power-law clouds, their effects were evaluated using a statistical approach." + For a given cloud realizationwe computed EW(t) and derived the absolute peak-to-peak variation AEW over the whole time interval —10 to +50 days., For a given cloud realizationwe computed $EW(t)$ and derived the absolute peak-to-peak variation $\Delta EW$ over the whole time interval $-$ 10 to +50 days. +" To mimic a more realistic situation, we did this also on two sub-sets of data-points including only post-maximum observations (0, +10, +20, +30, +40, +50 and"," To mimic a more realistic situation, we did this also on two sub-sets of data-points including only post-maximum observations (0, +10, +20, +30, +40, +50 and" +observational Gamma-ray bursts. (GRBs) are detected by high-energy satellites.,"}\tikzmark{mainBodyEnd1} + + \date{Received ?, ?, ?; accepted ?, ?, ?} + + +% \abstract{}{}{}{}{} +% 5 {} token are mandatory + + \abstract + % context heading (optional) + % {} leave it empty if necessary + {Multi-wavelength observations of gamma-ray burst (GRB) afterglows provide + important information about the activity of their central engines and their + environments. In particular, the short timescale variability, such as bumps + and/or rebrightening features visible in the m\tikzmark{mainBodyStart2}multi-wavelength\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}light\tikzmark{mainBodyEnd3} + \tikzmark{mainBodyStart4}curves,\tikzmark{mainBodyEnd4} \tikzmark{mainBodyStart5}is\tikzmark{mainBodyEnd5} \tikzmark{mainBodyStart6}still\tikzmark{mainBodyEnd6} \tikzmark{mainBodyStart7}poorly\tikzmark{mainBodyEnd7} \tikzmark{mainBodyStart8}understood.}\tikzmark{mainBodyEnd8} + % aims heading (mandatory) + {We analyze the multi-wavelength observations of the GRB\,100219A afterglow at redshift + 4.7. In particular, we attempt to identify the physical origin of the late + achromatic flares/bumps detected in the X-ray and optical bands.}\tikzmark{mainBodyStart9}}\tikzmark{mainBodyEnd9} + % methods heading (mandatory) + {We present ground-based optical photometric data and \textit{Swift} X-ray + observations on GRB\,100219A. We analyzed the temporal behavior of the X-ray + and optical light curves, as well as the X-ray spectra.}\tikzmark{mainBodyStart10}}\tikzmark{mainBodyEnd10} + % results heading (mandatory) + {The early flares in the X-ray and optical light curves peak simultaneously at about 1000 s + after the burst trigger, while late achromatic bumps in the X-ray and optical bands + appear at about $2\times 10^4$~s after the burst trigger. + These are uncommon features in the afterglow phenomenology. Considering the + temporal and spectral properties, we argue that both optical and X-ray + emissions come from the same mechanism. The late flares/bumps may be + produced by late internal shocks from long-lasting activity of the central + engine. An off-axis origin for a structured jet model is also discussed to + interpret the bump shapes. The early optical bump + can be interpreted as the afterglow onset, while the early X-ray flare could + be caused by the internal activity. GRB\,100219A exploded in a dense + environment as revealed by the strong attenuation of X-ray emission and the + optical-to-X-ray spectral energy distribution.}\tikzmark{mainBodyStart11}}\tikzmark{mainBodyEnd11} + % conclusions heading (optional), leave it empty if necessary + {}\tikzmark{mainBodyStart12}} Gamma-ray bursts (GRBs) are detected by high-energy observational satellites." + Ground-based telescopes can be subsequently alerted and perform follow-up observations., Ground-based telescopes can be subsequently alerted and perform follow-up observations. + The accurate positions delivered by the satellite provide an excellent opportunity for multi-wavelength observations., The accurate positions delivered by the satellite provide an excellent opportunity for multi-wavelength observations. + The analysis of GRB light curves can provide plenty of information about the central engine and the surrounding environment., The analysis of GRB light curves can provide plenty of information about the central engine and the surrounding environment. + A canonical shape has been identified for the X-ray light curves (Nousek et al., A canonical shape has been identified for the X-ray light curves (Nousek et al. + 2006: Zhang et al., 2006; Zhang et al. + 2006). as obtained by the X-ray Telescope (XRT).," 2006), as obtained by the X-ray Telescope (XRT)." + In the optical band. following the alert by the Burst Alert Telescope (BAT). ground-based follow-up observations carried out by robotic telescopes increased the number of well-sampled optical light curves significantly (see. e.g.. Melandri et al.," In the optical band, following the alert by the Burst Alert Telescope (BAT), ground-based follow-up observations carried out by robotic telescopes increased the number of well-sampled optical light curves significantly (see, e.g., Melandri et al." + 2008: Klotz et al., 2008; Klotz et al. + 2009; Rykotff et al., 2009; Rykoff et al. + 2009: Cenko et al., 2009; Cenko et al. + 2009)., 2009). + Apparently. these light curves present different temporal behaviors.," Apparently, these light curves present different temporal behaviors." + It is hard to identify a uniform characterization for them., It is hard to identify a uniform characterization for them. + An attempt to classify the temporal properties has been made by Melandri et al. (, An attempt to classify the temporal properties has been made by Melandri et al. ( +2008).,2008). + Rykoff et al. (, Rykoff et al. ( +2009) attempted to find commonalities within a sample of optical light curves. finding that the external forward shock is a possible origin for the overall optical emission.,"2009) attempted to find commonalities within a sample of optical light curves, finding that the external forward shock is a possible origin for the overall optical emission." + Panaitescu et al. (, Panaitescu et al. ( +2006) have studied in detail the decay of light curves both in the X-ray and optical bands from a theoretical point of view.,2006) have studied in detail the decay of light curves both in the X-ray and optical bands from a theoretical point of view. + Chromatic breaks identified by comparing optical and X-ray light curves indicate that most likely the optical and X-ray emissions arise from a different origin., Chromatic breaks identified by comparing optical and X-ray light curves indicate that most likely the optical and X-ray emissions arise from a different origin. + While light curves generically decay in time. in several cases rebrightenings or bumps are observed in the X-ray or optical bands.," While light curves generically decay in time, in several cases rebrightenings or bumps are observed in the X-ray or optical bands." + These features call for a more detailed investigation of the physies of GRB and afterglow., These features call for a more detailed investigation of the physics of GRB and afterglow. + However. the situation is quite complicated.," However, the situation is quite complicated." + Most rebrightenings are only observed in the X-ray band., Most rebrightenings are only observed in the X-ray band. + Some early X-ray bumps. usually called X-ray flares. have no corresponding optical features (see the statistics by Melandri et al.," Some early X-ray bumps, usually called X-ray flares, have no corresponding optical features (see the statistics by Melandri et al." + 2008 and Rykoff et al., 2008 and Rykoff et al. + 2009. as well as Uehara et al.," 2009, as well as Uehara et al." + 2010 for the individual cases of GRB0071112C and 0080506)., 2010 for the individual cases of 071112C and 080506). +" In contrast. but less frequently. a rebrightening feature may be seen in the optical but not in the X-ray band (e.g. 0050721, Antonelli et al."," In contrast, but less frequently, a rebrightening feature may be seen in the optical but not in the X-ray band (e.g. 050721; Antonelli et al." + 2006)., 2006). + Some GRBs with both. X-ray and optical bumps shown before 1000 s in the observers frame have been identified: 0060418 and 0060607A (Molinari et al., Some GRBs with both X-ray and optical bumps shown before 1000 s in the observer's frame have been identified: 060418 and 060607A (Molinari et al. + 2007). 0060904B (Klotz et al.," 2007), 060904B (Klotz et al." + 2008) and 0071031 (Krühhler et al., 2008) and 071031 (Krühhler et al. + 2009)., 2009). + The X-ray flare and optical bump of 0060418 have the same peak time., The X-ray flare and optical bump of 060418 have the same peak time. + However. it is likely that the X-ray flare has an internal origin. while the optical bump is the result of external shock onset (see the optical statistics from Oates et al.," However, it is likely that the X-ray flare has an internal origin while the optical bump is the result of external shock onset (see the optical statistics from Oates et al." + 2008)., 2008). + The optical bump and giant X-ray flare of 0060904B are clearly chromatic., The optical bump and giant X-ray flare of 060904B are clearly chromatic. + The optical rising and X-ray flare of 0071031. are not exactly simultaneous. but their observed correlation suggests à common origin caused by late central engine activity.," The optical rising and X-ray flare of 071031 are not exactly simultaneous, but their observed correlation suggests a common origin caused by late central engine activity." + More importantly. it is worth noting that GRBs with late bumps/rebrightenings shown after 100—10 s in the observers frame are very rare.," More importantly, it is worth noting that GRBs with late bumps/rebrightenings shown after $10^3$ $10^4$ s in the observer's frame are very rare." + We mention that the X-ray light curves of 0050502B. 0050724. 0050904. 0060413. 0060906. and 0070311 have similar late bump/flare/wiggle features.," We mention that the X-ray light curves of 050502B, 050724, 050904, 060413, 060906, and 070311 have similar late bump/flare/wiggle features." + Some of them have been identified as late X-ray flares (Curran et al., Some of them have been identified as late X-ray flares (Curran et al. + 2008; Bernardini et al., 2008; Bernardini et al. + 2011)., 2011). + It is rare. however. to have well-sampled.," It is rare, however, to have well-sampled," +quiescent O-rich stellar outflows.,quiescent O-rich stellar outflows. +" Conductive needles may therefore be a noticeable dust constituent of the general interstellar medium (ISAT),", Conductive needles may therefore be a noticeable dust constituent of the general interstellar medium (ISM). + Needle properlies required to explain the GC extinction curve are derived in 822. and the astrophysical implication. whether (he GC extinction is a special case or more characteristic of the general ISM is discussed in 823.," Needle properties required to explain the GC extinction curve are derived in 2, and the astrophysical implication, whether the GC extinction is a special case or more characteristic of the general ISM is discussed in 3." + The exünction optical depth 7(À) at wavelength A to a source located at distance 5 due to an assembly of dust particles along the line of sight (LOS) is given by: where (he sum is over dust composition j., The extinction optical depth $\tau(\lambda)$ at wavelength $\lambda$ to a source located at distance $S$ due to an assembly of dust particles along the line of sight (LOS) is given by: where the sum is over dust composition $j$. +" The integral in square brackets is over the grain size distribution /;(e) (which is normalized to unitv over the [u4,4,5. (,,4,] size interval). my,j(a.s)ο is the mass of a dust particle of radius a and mass density pj located at 5s. and Ελ.a) is the mass extinction coellicient of the dust."," The integral in square brackets is over the grain size distribution $f_j(a)$ (which is normalized to unity over the $a_{min,j}$, $a_{max,j}$ size interval), $m_{d,j}(a,\ s) = 4\pi \rho_j a^3/3$ is the mass of a dust particle of radius $a$ and mass density $\rho_j$ located at $s$, and $\kappa(\lambda,\ a)$ is the mass extinction coefficient of the dust." + The outer integral is over the LOS. where ny(5) is the II-number density at distance s.," The outer integral is over the LOS, where $n_H(s)$ is the H-number density at distance $s$." + Equation (1) can be written in a simplified form as: where &;(AÀ) represents an average over (he grain size distribution. and M; is (he mass column density of the dust of composition J. pj; is the average dust mass density in the ISAL. and L is the total distance along the LOS.," Equation (1) can be written in a simplified form as: where $\kappa_j(\lambda)$ represents an average over the grain size distribution, and ${\cal M}_{d,j}$ is the mass column density of the dust of composition $j$, $\rho_{d,j}$ is the average dust mass density in the ISM, and $L$ is the total distance along the LOS." + Figure 1 depicts (he average interstellar extinction curve for a mixture of bare graphite and silicate particles and PAIIs that satisfies Che UV-optical extinction curve. the diffuse IR. emission. as well as the solar interstellar abundances constraints (model. in Zubko. Dwek. Arendt. 2004. ZDA).The extinction was normalized to unitv at [x (A=2.2 yom). for which 7(IN)—3.47/1.086—3.20 (Rieke. Rieke. Paul 1989. Lutz et al.," Figure 1 depicts the average interstellar extinction curve for a mixture of bare graphite and silicate particles and PAHs that satisfies the UV-optical extinction curve, the diffuse IR emission, as well as the solar interstellar abundances constraints (model BARE-GR-S in Zubko, Dwek, Arendt 2004, ZDA).The extinction was normalized to unity at K $\lambda= 2.2\ \mu$ m), for which $\tau$ (K)=3.47/1.086=3.20 (Rieke, Rieke, Paul 1989, Lutz et al." + 1996)., 1996). + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟��⋯⋅≼↲≀, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧↴, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧↴↕, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧↴↕⋅, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧↴↕⋅≼, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are + At 10 jan the extinction is dominated by the silicate absorption feature which has ⋅⋅ ↼ ⋅ ⋅ ⋅≓ ≀↧↴∐⋯↪∖⇁⋝∖⊽≀↧↴∣↽≻⊳∖⇁∪↕⋅↕↽≻⊔∪∐≺∢∪≼↲∐↓≺∢↕≼↲∐↥∣↘⇁⊽∖∣⋅∣⋜↜≤∍⋅↙∕∣∐↓∣⋝∶≡↽⊰⋗⋖∐≱≺∢∐↓−↖↳↴↓⋅↼≚↥⊳∖⊽∪⊳∖⊽∐⋯∖↽∐↕∐⊔∐↲∐↖∟↴⋯⋅≼↲≀↧↴↕⋅≼↲, At $\sim10\ \mu$ m the extinction is dominated by the silicate absorption feature which has a mass absorption coefficient $\kappa_{sil}(9.7\ \mu$ m) = $\times10^3$ $^2$ $^{-1}$ Also shown in the figure are +Cravitational waves (GAs). ripples in space-time. perturh our four dimensional space-time background. on which the pulsed racio radiation from. pulsars propagates.,"Gravitational waves (GWs), ripples in space-time, perturb our four dimensional space-time background, on which the pulsed radio radiation from pulsars propagates." + Phrough such space-time background: perturbations. the CAV leaves its fingerprint in the arrival times of pulsar signals by introducing an extra correlated component. (??)..," Through such space-time background perturbations, the GW leaves its fingerprint in the arrival times of pulsar signals by introducing an extra correlated component \citep{Sazhin78, Detweiler79, BCR84, Wahlquist87, +BH86}." + By timing multiple pulsars, By timing multiple pulsars +"ueasurelents of r,; bv using stall scale temperature Huctuations. eiven that the uucertainty for our Model A ics af m4,=OTL.","measurements of $\tau_{ri}$ by using small scale temperature fluctuations, given that the uncertainty for our Model A lies at $\sigma_{A_D}=0.71$." + When DPlaucks power spectra are imnodceled with he “patchy amplitude” parameter. constraints for the standard cosmological paraicters are slightly improved over the analysis that abandoned teniperature data vevoud 7=1000 (Table 21). despite the imtroductiou of an extra parameter.," When Planck's power spectra are modeled with the “patchy amplitude” parameter, constraints for the standard cosmological parameters are slightly improved over the analysis that abandoned temperature data beyond $l=1000$ (Table \ref{bias_reduced}) ), despite the introduction of an extra parameter." +" C'oncretelv. we find iu this analvsis that στ στ, oo.=0.0L1. 0, 0.0018. 0, 70.00021. oo, —0.0058. a). —0.0052."," Concretely, we find in this analysis that $\sigma_{\tau_{ri}}$ =0.037, $\sigma_{\Omega_\Lambda}$ =0.011, $\sigma_{\omega_{dm}}$ =0.0018, $\sigma_{\omega_b}$ =0.00021, $\sigma_{\Omega_{n_s}}$ =0.0058, $\sigma_{A_s}$ =0.0052." + The Doppler power spectrum can bias the result of piuraueter analyses with future CAB iumvevs., The Doppler power spectrum can bias the result of parameter analyses with future CMB surveys. + Plauck may be able to prove constramts on models with au extended epoch of patelhiuess., Planck may be able to improve constraints on models with an extended epoch of patchiness. + We hope to have shown in this section that careful modeling of the epoch of first stars will become crucial for dome precision cosmology with progressively more refined CMD experiments., We hope to have shown in this section that careful modeling of the epoch of first stars will become crucial for doing precision cosmology with progressively more refined CMB experiments. + We lave presented simulations of secondary alisotropies in the cosnüc nücrowave backerownd— calculated using smoothed particle lydrodvuamics siuulatious of large scale structure. focusing on the effect produced by patchy reiouiz:ition.," We have presented simulations of secondary anisotropies in the cosmic microwave background calculated using smoothed particle hydrodynamics simulations of large scale structure, focusing on the effect produced by patchy reionization." + We incorporated an analytic inodel for the uxpholoev of the IIII regions iuto our nunerical treatnent and investigated whether such patchy scenarios car1 be distinguished from homogceucous reionization., We incorporated an analytic model for the morphology of the HII regions into our numerical treatment and investigated whether such patchy scenarios can be distinguished from homogeneous reionization. + An important advantage of our technique over pure analytical predictions of patchy reionization morpholoey is that we follow the complicated clustering of dark matter and barvons iuto the slightly mon-linear reeiue., An important advantage of our technique over pure analytical predictions of patchy reionization morphology is that we follow the complicated clustering of dark matter and baryons into the slightly non-linear regime. + Iu contrast to full radiative trausfor calculations of the reiouization epoch. we combine uucertaiufies iu the plysics of source formation. feedback processes and radiative transfer iuto a sinele parameter and explore the consequences of varving this parameter.," In contrast to full radiative transfer calculations of the reionization epoch, we combine uncertainties in the physics of source formation, feedback processes and radiative transfer into a single parameter and explore the consequences of varying this parameter." + This simplification allows us to make predictions on scales au order of magnitude larger than current radiative transfer schemes can accomplish with a small expense of iiemiory and CPU., This simplification allows us to make predictions on scales an order of magnitude larger than current radiative transfer schemes can accomplish with a small expense of memory and CPU. + We extracted power spectra from sky maps produced bv tracing ravs across our simulation voluues., We extracted power spectra from sky maps produced by tracing rays across our simulation volumes. +" The patchy reionization signal peaks ou iuultipole scales of {σε2000. and it increases the amplitude of the ""cleaned? CAB power spectrum by up to 30% on scales /21000. so that the total level of Doppler related auisotropies is AT~—2.dpA."," The patchy reionization signal peaks on multipole scales of $l\simeq 2000$, and it increases the amplitude of the “cleaned” CMB power spectrum by up to $30\%$ on scales $l\geq4000$, so that the total level of Doppler related anisotropies is $\Delta T \simeq 2.4 \mu +K$." + Wo found that with the next eeneration of eround based CAIB experiments (ACT. SPT) the differeut reionization models we investigated could be distinguished with high significance by using he power spectrum.," We found that with the next generation of ground based CMB experiments (ACT, SPT) the different reionization models we investigated could be distinguished with high significance by using the power spectrum." +" Additional information about the norphological properties of that epoch that could be obtained for instance by iueasurme the four poiut ""uctiou or other deviations from Caussianity will xobablv be difficult to obtain. because the reioulzation signal peaks onu angular scales where the primorcial CXMB anisotropies dominate."," Additional information about the morphological properties of that epoch that could be obtained for instance by measuring the four point function or other deviations from Gaussianity will probably be difficult to obtain, because the reionization signal peaks on angular scales where the primordial CMB anisotropies dominate." + In the future. if may boe yossible to combine measurements of the CAIB with other observations such as 21 cnm fluctuatious frou jeutral hydrogen (Zaldariiagaetal.2001:Cooray2001) or Lxauiuralpha cussion from high redshift galaxies (e.g. Furluettoetal.20014011 to further coustrain the opologv of the reionization process.," In the future, it may be possible to combine measurements of the CMB with other observations such as 21 cm fluctuations from neutral hydrogen \citep{Zaldarriaga:2004cm,coo04} + or Lyman-alpha emission from high redshift galaxies (e.g. \citealt{Furlanetto:2004to}) ) to further constrain the topology of the reionization process." + We investigated the bias iu the determination of cosmological parameters that will be produced by the additional patchy reionization signal. when extracting cosmological parazuuneters frou CAMB anisotropy micasurements.," We investigated the bias in the determination of cosmological parameters that will be produced by the additional patchy reionization signal, when extracting cosmological parameters from CMB anisotropy measurements." +— We find that for Plauck this bias is siguificant., We find that for Planck this bias is significant. + The bias uiav be circumvented by focusing completely on polarizaion information in the umultipole reeiue where patchiness peaks. with only a slight disadvantage in parameter constraints.," The bias may be circumvented by focusing completely on polarization information in the multipole regime where patchiness peaks, with only a slight disadvantage in parameter constraints." + Alternatively. a template for the Doppler spectra could be introduced in the parameter analysis which may lead to a detection of the effects of an extended reiouization pliase.," Alternatively, a template for the Doppler spectrum could be introduced in the parameter analysis which may lead to a detection of the effects of an extended reionization phase." + O. Z. aud M. Z. ave supported by NSF eraut AST 0095606 and bv the David aud Lucille Packud Foundation Fellowship for Scieuce aud. Eugineeriug aud by the Sloan Foundation., O. Z. and M. Z. are supported by NSF grant AST 0098606 and by the David and Lucille Packard Foundation Fellowship for Science and Engineering and by the Sloan Foundation. + This work is also supported by NSF erants ACT 96-19019. AST 00-71019. AST 02-06299. and AST 0383-07690. aud NASA ATP erauts NACS5-12110. NAG5-13292. and NAG5S-13381.," This work is also supported by NSF grants ACI 96-19019, AST 00-71019, AST 02-06299, and AST 03-07690, and NASA ATP grants NAG5-12140, NAG5-13292, and NAG5-13381." + The simulations were performed at the Center for Parallel Astrophysical Computing at the ILhuwird-Suüthsonian Center for Astrophysics., The simulations were performed at the Center for Parallel Astrophysical Computing at the Harvard-Smithsonian Center for Astrophysics. +and by Nelson&Ixeeidl(1993). do not detect. pulsation in this star.,and by \cite{Nelson93} do not detect pulsation in this star. + This can be understood. in. terms of rotational modulation in an oblique pulsator. and/or beating of multiple frequencies.," This can be understood in terms of rotational modulation in an oblique pulsator, and/or beating of multiple frequencies." + The pulsation amplitudes in many roAp stars ave modulated with the rotation period of the star., The pulsation amplitudes in many roAp stars are modulated with the rotation period of the star. + Thus some of the photometric observations may not have been in best aspect., Thus some of the photometric observations may not have been in best aspect. + To judge this the rotational period and ephemeris of 996237 needs to be determined., To judge this the rotational period and ephemeris of 96237 needs to be determined. + From the AI Sky Automated Survey (ASAS) and Llipparcos photometry. Frevhanuneretal.(2008). found |1otometric variations with a period of 20.91dd. which may be the rotational period.," From the All Sky Automated Survey (ASAS) and Hipparcos photometry, \citet{Freyhammer08b} found photometric variations with a period of d, which may be the rotational period." + Using two seasons of data [rom the WASP (Wide Anele Search for Planets) Project. (Pollaccoetal. 2006)) covering the intervals 2007 January 4 to 2007 June 3 and 2008 January 5 to 2008 Alay 28 we obtained a similar period., Using two seasons of data from the WASP (Wide Angle Search for Planets) Project \citealt{Pollacco06}) ) covering the intervals 2007 January 4 to 2007 June 3 and 2008 January 5 to 2008 May 28 we obtained a similar period. + Combining the WASDP (passband from 400 to nnm) and ASAS (V-band) photometry we find the following rotational ephemeris for the photometric maximum: The rotational phases calculated. from this ephemeris are also shown in ref06237:phot.., Combining the WASP (passband from 400 to nm) and ASAS (V-band) photometry we find the following rotational ephemeris for the photometric maximum: The rotational phases calculated from this ephemeris are also shown in \\ref{96237:phot}. + The bottom panel is for data observed. in the same rotation period as the upper panel. but. does not display any pulsations.," The bottom panel is for data observed in the same rotation period as the upper panel, but does not display any pulsations." + Ehe spectral observations presented at reLOG237:ampspec were also obtained at a similar rotation phase using the above ephemeris., The spectral observations presented at \\ref{96237:ampspec} were also obtained at a similar rotation phase using the above ephemeris. + Assuming our case for pulsation in this star to be good. we suggest two possible solutions: 1) the rotation period may be not dd. but double that.," Assuming our case for pulsation in this star to be good, we suggest two possible solutions: 1) the rotation period may be not d, but double that." + This ambiguity happens for some peculiar stars (see for example Wadeetal. (1997)))., This ambiguity happens for some peculiar stars (see for example \cite{Wade97}) ). + Longitudinal magnetic field. measurements over the rotational period. can resolve this: 2) the star may. be multiperiodic., Longitudinal magnetic field measurements over the rotational period can resolve this; 2) the star may be multiperiodic. + There is some hint in ref06237:phot of a peak at QO.79mmllz., There is some hint in \\ref{96237:phot} of a peak at mHz. + Additional observations are required to resolve these questions., Additional observations are required to resolve these questions. + The chemically peculiar magnetic star 996237. is an interesting object which has a very peculiar spectrum with significant spectral variability., The chemically peculiar magnetic star 96237 is an interesting object which has a very peculiar spectrum with significant spectral variability. + Phe star demonstrates large overabundances of rare earth elements., The star demonstrates large overabundances of rare earth elements. + A high resolution spectrum obtained with the ESO 2.2-m telescope. anc FEROS spectrograph resembles the spectrum. of another highly peculiar star. 1101065.," A high resolution spectrum obtained with the ESO 2.2-m telescope and FEROS spectrograph resembles the spectrum of another highly peculiar star, 101065." + Xbundances of and determined. from. this spectrum. are even higher than in 1101065., Abundances of and determined from this spectrum are even higher than in 101065. + Other rare earth elements also show large overabuncdances similar to those found in 1101065., Other rare earth elements also show large overabundances similar to those found in 101065. +The mechanisms of the energy. transfer. from the pulsar to the pulsar wind. nebula still. remain obscure.,The mechanisms of the energy transfer from the pulsar to the pulsar wind nebula still remain obscure. + Lt is widely acceptec that pulsars emit. an electron-positron plasma. which form an ultrarelativistic magnetizecl wind.," It is widely accepted that pulsars emit an electron-positron plasma, which form an ultrarelativistic magnetized wind." + The rotational energy. of the neutron star is carried mostly by electromagnetic fields as Povnting flux (Michel 1982) and should be eventually transferred to the radiating particles., The rotational energy of the neutron star is carried mostly by electromagnetic fields as Poynting flux (Michel 1982) and should be eventually transferred to the radiating particles. +" Vhe wind terminates at oa shock. front located. in the case of the Crab Nebula. some 107"" em from the pulsar."," The wind terminates at a shock front located, in the case of the Crab Nebula, some $10^{17}$ cm from the pulsar." + The postshock Dow parameters may be matched: with the observed. Crab. structure if the enerey ας at the shock front is carried mainly by the particles (Rees Ciunn 1974: kennel Coroniti 1984: Emumering Chevalier 1987). (, The postshock flow parameters may be matched with the observed Crab structure if the energy flux at the shock front is carried mainly by the particles (Rees Gunn 1974; Kennel Coroniti 1984; Emmering Chevalier 1987). ( +lomimatsu 1994: DBogovalov 1997: 2001b: Deskin. ]xuznetsova Ralikoy 1998: Chiuch. Li Beeelman 1998: Lyubarsky Eichler 2001).,"Tomimatsu 1994; Bogovalov 1997; 2001b; Beskin, Kuznetsova Rafikov 1998; Chiueh, Li Begelman 1998; Lyubarsky Eichler 2001)." + Phe necessary conversion may be provided by dissipation mechanisms., The necessary conversion may be provided by dissipation mechanisms. + These mechanisms operate faster at small scales aud first alfect waves generated in the wind bv the rotating. oblique pulsar magnetosphere.," These mechanisms operate faster at small scales and first affect waves generated in the wind by the rotating, oblique pulsar magnetosphere." + There is à variety of electromagnetic waves in rarefied magnetized plasmas (ee. Akhiezer ct 11975) however we assume that only true AUID waves. (those satisfving the condition E=v DB) may be generated by. the rotating magnetosphere.," There is a variety of electromagnetic waves in rarefied magnetized plasmas (e.g., Akhiezer et 1975) however we assume that only true MHD waves (those satisfying the condition $\bf E=v\times B$ ) may be generated by the rotating magnetosphere." + The reason is that according to the convenient view. the plasma density in the pulsar. wind is sullicienthy large such that electromagnetic waves are heavily cdamped (e... Assco et 11978: Melatos Melrose 1996).," The reason is that according to the convenient view, the plasma density in the pulsar wind is sufficiently large such that electromagnetic waves are heavily damped (e.g., Asseo et 1978; Melatos Melrose 1996)." +" ""here are four tvpes of MILD: waves but only the entropy ancl the fast magnetosonic (EMS) waves may propagate [ar bevond the light evlinder where the magnetic field is predominantly. toroidal.", There are four types of MHD waves but only the entropy and the fast magnetosonic (FMS) waves may propagate far beyond the light cylinder where the magnetic field is predominantly toroidal. +2001a). The entropy wave decavs because of the current starvation in current sheets separating strips with the opposite magnetic field (Usoy 1975: Michel 1982. 1994: Coroniti 1990).," The entropy wave decays because of the current starvation in current sheets separating strips with the opposite magnetic field (Usov 1975; Michel 1982, 1994; Coroniti 1990)." + Lyubarsky Wirk (2001) showed that the low significantly accelerates in the course of reconnection and this clilates the timescale over which the wave clecavs., Lyubarsky Kirk (2001) showed that the flow significantly accelerates in the course of reconnection and this dilates the timescale over which the wave decays. + A typical conditions. the dissipation radius exceeds the radius of the termination shock therefore one should conclude tha the Povnting [lux in the striped wind does not dissipate unti the wind enters the termination shock.," At typical conditions, the dissipation radius exceeds the radius of the termination shock therefore one should conclude that the Poynting flux in the striped wind does not dissipate until the wind enters the termination shock." + All the energy shoule release within the shock where the Low decelerates., All the energy should release within the shock where the flow decelerates. + 1n this article. the fate of PMS waves is considered.," In this article, the fate of FMS waves is considered." + | will be shown that these waves may be described: within the ALD framework throughout the pulsar wind up to the termination shock and ever bevond., It will be shown that these waves may be described within the MHD framework throughout the pulsar wind up to the termination shock and ever beyond. + Phe reason is that in a magnetically dominated plasma. the PAIS waves excite small," The reason is that in a magnetically dominated plasma, the FMS waves excite small" +Figure 2..,Figure \ref{fig:2}. + For forces slightly greater than f—lOnm! the DNA undergoes an over-stretching transition |?].. hence the elastic rod model is not relevant.," For forces slightly greater than $\tilde{f}\sim +10\,\mathrm{nm^{-1}}$ the DNA undergoes an over-stretching transition \cite{Smith2}, hence the elastic rod model is not relevant." + We have therefore. picked the force range of 0—f5nm.," It can be seen that $\frac{\langle z^{\,(2)}\rangle}{\langle z^{\,(0)}\rangle}$ does not exceed $10^{-2}$ for $A \geq 5 \mathrm{nm} +$." + As can be seen from Figure 2.. for A= 50nm. where the theoretical curve is best fitted to theexperimental data |?].. one must measureexperiments al least with the accuracy LO1 to detect (2'7').," As can be seen from Figure \ref{fig:2}, for $A=50\,\mathrm{nm}$ where the theoretical curve is best fitted to theexperimental data \cite{Marko}, one must measure $\frac{\langle z\rangle}{L}$ at least with the accuracy $10^{-4}$ to detect $\langle z^{\,(2)}\rangle $." + Since L~10yim in /2].. minimum accuracy of Lum is required in measuring (2).," Since $L\sim +10\,\mathrm{\mu m}$ in experiments \cite{Smith}, minimum accuracy of $1\,\mathrm{nm}$ is required in measuring $\langle z\rangle $." + However. the accuracy of the experiments is by far less than this limit [?].. therefore ἐς7) can not be detected by stretching experiments.," However, the accuracy of the experiments is by far less than this limit \cite{Smith}, therefore $\langle z^{\,(2)}\rangle $ can not be detected by stretching experiments." + We now show that (zU) is also small.," We now show that $\langle z^{\,(3)}\rangle $ is also small." + It is obvious that when W(O.0) is independent of the Euler angle 5. the partition function is invariant. under the transformation A——A.," It is obvious that when $\Psi (\Theta ,0)$ is independent of the Euler angle $\gamma$, the partition function is invariant under the transformation $\lambda\rightarrow -\lambda $." + This means that odd powers of À are not present in the expansion of (=). ie.. (27?1)= 0.," This means that odd powers of $\lambda $ are not present in the expansion of $\langle z\rangle $, i.e., $\langle z^{\,(2p+1)}\rangle =0$ ." + In addition. the effect of the initial conditions on the force extension curve ofDNA is suppressed if DNA is long enough.," In addition, the effect of the initial conditions on the force extension curve ofDNA is suppressed if DNA is long enough." + As a result. one expects — to be small even when Ψ(Ο.0) depends on 5.," As a result, one expects $\frac{\langle z^{\,(2p+1)}\rangle}{L}$ to be small even when $\Psi (\Theta ,0)$ depends on $\gamma $." + In other words. odd powers of A have no significant contribution to the end-to-end DNA extension.," In other words, odd powers of $\lambda$ have no significant contribution to the end-to-end DNA extension." + Therefore. to the third order of A. the response of an anisotropic DNA to the external tension is close to an isotropic DNA with the effective bending constant To justifv our result. we must show that the condition AE?Lx1 which corresponds to the limit of long DNA. is satisfied in experiments as well.," Therefore, to the third order of $\lambda$ , the response of an anisotropic DNA to the external tension is close to an isotropic DNA with the effective bending constant To justify our result, we must show that the condition $\Delta\mathcal{E}^{R}L\gg 1$ which corresponds to the limit of long DNA, is satisfied in experiments as well." + Figure 5.shows AE’ as a [function of fA for A= 50nm.," Figure \ref{fig:4} shows $\Delta\mathcal{E}^{R}A$ as a function of $\tilde{f}A$ for $A=50\,\mathrm{nm}$ ." + As can be seen.AEM4> 1.," As can be seen,$\Delta\mathcal{E}^{R}A\geq +1$ ." + As a result. the condition AE’Lc| is equivalent to the condition £X.1. which is well known in polvimer physics.," As a result, the condition $\Delta\mathcal{E}^{R}L\gg 1$ is equivalent to the condition $L\gg A$, which is well known in polymer physics." + Since A=50nm and £~10 imn. this condition is satisfiedin the streching experiments.," Since $A=50\,\mathrm{nm}$ and $L\sim +10\,\mathrm{\mu m}$ , this condition is satisfiedin the streching experiments." +mechanisms appear not viable.,mechanisms appear not viable. + Relativisiic beaming effects due to changes in Lorentz factor and emission direction sill need to be investigated and are plausibly very important., Relativistic beaming effects due to changes in Lorentz factor and emission direction still need to be investigated and are plausibly very important. + Longer duration. more sensitive observations of giant phenomena over a larger range of [luxes. together wilh associated fitting of multiple vectoriallv-convolved wave distributions should resolve these observational issues.," Longer duration, more sensitive observations of giant phenomena over a larger range of fluxes, together with associated fitting of multiple vectorially-convolved wave distributions should resolve these observational issues." + Outstanding theoretical issues may be resolved by extending current simulations and theories for relativistic beaming ancl wave collapse. together with any other radiation mechanisms required. (o conditions appropriate for pulsar magnetosplhieres.," Outstanding theoretical issues may be resolved by extending current simulations and theories for relativistic beaming and wave collapse, together with any other radiation mechanisms required, to conditions appropriate for pulsar magnetospheres." +The model described here constructs the particular spatial distribution of E. based ou the input surface magnetic map and its potential field.,The model described here constructs the particular spatial distribution of $E_\gamma$ based on the input surface magnetic map and its potential field. + It then solves the set of conservation laws for nass. monmeutun. magnetic iuduction.and enerev (the ideal MITID equations): with P23/2 until a steady-state stellar wind solution is obtained.," It then solves the set of conservation laws for mass, momentum, magnetic induction,and energy (the ideal MHD equations): with $\Gamma=3/2$ until a steady-state stellar wind solution is obtained." + The stellar iuput parameters required for the model are the boundary. value for the density. py. as well as the stellar radius. νι mass. AZ... and rotation frequency. Οι.," The stellar input parameters required for the model are the boundary value for the density, $\rho_0$, as well as the stellar radius, $R_\star$, mass, $M_\star$, and rotation frequency, $\Omega_\star$." + À παν of the stella parameters of AD Dor adopted for the simulation. based on the references cited in Section 1.. is provided im Table 1..," A summary of the stellar parameters of AB Dor adopted for the simulation, based on the references cited in Section \ref{sec:Intro}, is provided in Table \ref{table:t1}." + Iu the simulations preseuted here. we assune that the relation between the flux tube expansion aud the terminal speed obtained from that flux tube is a universal process that occurs on AB Dor in a similar manner to the Sun.," In the simulations presented here, we assume that the relation between the flux tube expansion and the terminal speed obtained from that flux tube is a universal process that occurs on AB Dor in a similar manner to the Sun." + Observatious of the corona of AB Dor indicate dominant plasina temperatures peaking in the range 330 AT (e.g.Sauz-Forcadaetal.2003:Carcia-Alvarezetal.," Observations of the corona of AB Dor indicate dominant plasma temperatures peaking in the range 3--30 MK \citep[e.g.][]{SanzForcada03,GarciaAlvarez05}." + 2005).. Ty is. in principle. the average temperature of the stellar corona.," $T_0$ is, in principle, the average temperature of the stellar corona." + However. we stress that iu our model. Zp is essentially a free parameter for the boundary condition that controls the euergization of the stellar wind.," However, we stress that in our model, $T_0$ is essentially a free parameter for the boundary condition that controls the energization of the stellar wind." + Further description of the adaptation of a solar corona mmocdel to stellaz coronae cau be fouud in (Cohenetal. 2010)., Further description of the adaptation of a solar corona model to stellar coronae can be found in \citep{Cohen10}. +. Figure d. shows the the iuput surface magnetic fiek map adopted for the simulations., Figure \ref{fig:f1} shows the the input surface magnetic field map adopted for the simulations. + This i$ based o-— spectropolariaetriec observations obtained in DoecenibeLt 2007 (Ilusmün ct al., This is based on spectropolarimetric observations obtained in December 2007 (Hussain et al. + submitted to MNRAS). iux analyzed im a similar manner to the maps described iu IIussainetal.(2002.2005. 2007)..," submitted to MNRAS), and analyzed in a similar manner to the maps described in \cite{Hussain02,Hussain05,Hussain07}. ." + The reader is referre to those works for further details., The reader is referred to those works for further details. + Since AB Dor has au inclination of 60deg (I&uersteretal.1991).. the part of the stellar surface near the far pole is alwavs hidden from view.," Since AB Dor has an inclination of $60\deg$ \citep{Kuersteretal94}, the part of the stellar surface near the far pole is always hidden from view." + Consequently. surface maps for AB Dor are intrinsically incomplete.," Consequently, surface maps for AB Dor are intrinsically incomplete." +" The initial. iucouiplete map is shown in the top left panel of Figure Ἐν,"," The initial, incomplete map is shown in the top left panel of Figure \ref{fig:f1}." + To construct a complete map. we cuforced hemispherical reflection svinmnetry on the magnetic Ποια across the equatorial aue.," To construct a complete map, we enforced hemispherical reflection symmetry on the magnetic field across the equatorial plane." + Those parts of the southern hemisphere with uaenetie field magnitude of less than 50C were assigned naenetic field values from the same longitude at the corresponding northern hemisphere latitude. but with he opposite polavity.," Those parts of the southern hemisphere with magnetic field magnitude of less than $50\;G$ were assigned magnetic field values from the same longitude at the corresponding northern hemisphere latitude, but with the opposite polarity." + The complete map used in the sinulation is shown as a loueitude-latitude contour map (top-rielt panel) and as spherical plots of the two ouegitudinal hemuspheres. colored with contours of the surface magnetic feld (lower panels).," The complete map used in the simulation is shown as a longitude-latitude contour map (top-right panel) and as spherical plots of the two longitudinal hemispheres, colored with contours of the surface magnetic field (lower panels)." + Figure 5 slows he three-dinieusional distribution of the potential field calculated based on this input surface map., Figure \ref{fig:f2} shows the three-dimensional distribution of the potential field calculated based on this input surface map. + It cau e seen that the loops extend up to the height of the source surface (located at 7=LOR.) aud that they have uo toroidal compoucnt., It can be seen that the loops extend up to the height of the source surface (located at $r=10R_\star$ ) and that they have no toroidal component. + The field is fully radial above the source surface as required by the analytical solution., The field is fully radial above the source surface as required by the analytical solution. + We caution that there are a umber of sources of systematic uncertainty preseut in the simulations., We caution that there are a number of sources of systematic uncertainty present in the simulations. + First. the ZDI maps are missing part of the stellar surface. ar we have extrapolated the field. assuming antisviunietry. to those non-visible areas.," First, the ZDI maps are missing part of the stellar surface, and we have extrapolated the field, assuming antisymmetry, to those non-visible areas." + Second. the maguetic fiek naps are of limuted resolution (latitudinal resolution of 3 degrees) aud do not resolve the details of the active reeions.," Second, the magnetic field maps are of limited resolution (latitudinal resolution of 3 degrees) and do not resolve the details of the active regions." + Third. large areas in the maps that appear to have strong fields may in fact be dominated by localizes active regions that are smeared out due to the lower resolution: such a scenario would lead to a significautlv different MIID solutiou.," Third, large areas in the maps that appear to have strong fields may in fact be dominated by localized active regions that are smeared out due to the lower resolution; such a scenario would lead to a significantly different MHD solution." + While the ZDI maps do likely muss low level structure. the maps should recover the strongest field reeious. which are likely to dominate over the laree-scale global models such as those considered here.," While the ZDI maps do likely miss low level structure, the maps should recover the strongest field regions, which are likely to dominate over the large-scale global models such as those considered here." + The weaker complex fields would only be interesting much closer to the surface., The weaker complex fields would only be interesting much closer to the surface. + Observations of AB Dor reveal that the coronal base density. iy. ranges between 1019104?e7? (Sauz-Forcadaetal. 2003).," Observations of AB Dor reveal that the coronal base density, $n_0$, ranges between $10^{10}-10^{12}\;cm^{-3}$ \citep{SanzForcada03}." +. These measurements. however. were nade based on neasurenmients of strong clussion lues that are associated with the denser. closed. loops.," These measurements, however, were made based on measurements of strong emission lines that are associated with the denser, closed loops." +" Therefore. the density in the ""quiet star” (analogous to the ""quiet Sun”) where the wind originates should be lower."," Therefore, the density in the “quiet star” (analogous to the “quiet Sun”) where the wind originates should be lower." +" In order to partially cover this density range. we siuulate three test cases with different values for the coronal base density. """," In order to partially cover this density range, we simulate three test cases with different values for the coronal base density. “" +"Case A7. with oyο105eim7. which is the value used for simulations of the solar wincl. “Case D with my=10°enm7. and ""Case CU with ny=101en 7.","Case A”, with $n_0=2\cdot 10^8\;cm^{-3}$, which is the value used for simulations of the solar wind, “Case B” with $n_0=10^{9}\;cm^{-3}$, and “Case C” with $n_0=10^{10}\;cm^{-3}$ ." +" We simulate the wind and corona using a Cartesian box of 30,x30R,x30R,. in the frame of reference rotating with the star (to expedite couverecuce using the local time step algorithm (Cohenetal.2008)))."," We simulate the wind and corona using a Cartesian box of $30R_\star$ $30R_\star$ $30R_\star$, in the frame of reference rotating with the star (to expedite convergence using the local time step algorithm \citep{cohen08b}) )." +" We use a non-muiform erid with a maxima resolution of 2.102R, prescribed new the surface.", We use a non-uniform grid with a maximum resolution of $2\cdot10^{-2}R_\star$ prescribed near the surface. + The erid is dvnanücallv refined durius the simulation so that lieh resolution is applied at location of magnetic field Inversion (current sheets)., The grid is dynamically refined during the simulation so that high resolution is applied at location of magnetic field inversion (current sheets). + We performed the simulation using the PLETADES super computer at the NASA AMES center., We performed the simulation using the PLEIADES super computer at the NASA AMES center. + The steady state MIID solutious for the three test cases are shown in Figure 8.., The steady state MHD solutions for the three test cases are shown in Figure \ref{fig:f3}. + The most notable feature of the solutious for all cases is the taugline of the field iu theazinmuthal direction due to the rapid rotation of the star., The most notable feature of the solutions for all cases is the tangling of the field in the direction due to the rapid rotation of the star. + This feature clearly does not appear in the poteutial field solution. for which stellar rotation is not a relevant parameter. nor m similar ΑΠΟ solutious for the Sun (c.g.Cohenetal. 2008)..," This feature clearly does not appear in the potential field solution, for which stellar rotation is not a relevant parameter, nor in similar MHD solutions for the Sun \citep[e.g.][]{cohen08b}." + As inielit be expected. solutions for he three cases are qualitatively quite similar. though coser mnnspection does reveal significant differences.," As might be expected, solutions for the three cases are qualitatively quite similar, though closer inspection does reveal significant differences." + It cai be seen from the aiddle pancl of Figure 8. that fjo radial wind speed decreases with an increase of the se deusitv., It can be seen from the middle panel of Figure \ref{fig:f3} that the radial wind speed decreases with an increase of the base density. + Iu addition. the corona ix denser in the solition for Case C as compared to Case A. While many ofthe taneled field lines in Cases A aud Bare open due to tje strong racial stretching by the stellar wind. in Case C. most of the taneled field lines are closed.," In addition, the corona is denser in the solution for Case C as compared to Case A. While many of the tangled field lines in Cases A and B are open due to the strong radial stretching by the stellar wind, in Case C, most of the tangled field lines are closed." + The closed oops in the lowcoronaare raclially stretched in Cases A aud D. while the," The closed loops in the lowcoronaare radially stretched in Cases A and B, while the" +to be taken into account.,to be taken into account. + We used in our spectrum synthesis calculation the line list and atomie data of ? for this spectral region. supplemented by a few additional. weak lines identified with the Vienna Atomic Line (VALD:??).," We used in our spectrum synthesis calculation the line list and atomic data of \cite{Johnson/Bolte:2001} for this spectral region, supplemented by a few additional, weak lines identified with the Vienna Atomic Line \cite[VALD;][]{VALD2a,VALD2b}." + Whenever the atomic data retrieved form VALD disagreed with those listed in Johnson Bolte. we adopted the former.," Whenever the atomic data retrieved form VALD disagreed with those listed in Johnson Bolte, we adopted the former." + This was the case only for very few lines., This was the case only for very few lines. + The most important case isthe 4019.057 line. for which we adopt the VALD value of loggf=0.213.," The most important case isthe $4019.057$ line, for which we adopt the VALD value of $\log gf = -0.213$." + This is 0.306 ddex lower than that used by Johnson Bolte. which in turn goes back to ?..," This is $0.306$ dex lower than that used by Johnson Bolte, which in turn goes back to \citet{Snedenetal:1996}." + The latter authors artificially increased the logef value of this line by ddex to improve their fit to the spectrum of CS 22892-052.," The latter authors artificially increased the $\log +gf$ value of this line by dex to improve their fit to the spectrum of CS 22892-052." + However. since ? have shown that at least part of the missing absorption on the blue side of the Th line is due toCH.. we do not use the scaled logef value for the Ce line. but thatlisted in VALD.," However, since \citet{Norrisetal:1997b} have shown that at least part of the missing absorption on the blue side of the Th line is due to, we do not use the scaled $\log gf$ value for the Ce line, but thatlisted in VALD." + For the 4019.129 line we use loggf=0.228 (?).. which ts 0.042 ddex higher than the value adopted by Johnson Bolte. and 0.423 ddex higher than that listed in VALD.," For the $4019.129$ line we use $\log gf = -0.228$ \citep{Nilssonetal:2002}, which is $0.042$ dex higher than the value adopted by Johnson Bolte, and $0.423$ dex higher than that listed in VALD." + For the partition functions of thortum we interpolate between the values listed in Table 3 of ?.. which are based on calculations provided by Holweger.," For the partition functions of thorium we interpolate between the values listed in Table 3 of \citet{Morelletal:1992}, which are based on calculations provided by Holweger." + In the temperature range covered by our stars.the partition functions for are higher by a factor of 3.6-4.4 than those listed in 2.. translating to abundance differences of the order of 0.5 ddex.," In the temperature range covered by our stars,the partition functions for are higher by a factor of $3.6$ $4.4$ than those listed in \citet{Irwin:1981}, translating to abundance differences of the order of $0.5$ dex." + For most of the elements which have lines in the relevant spectral region. abundances were available from clean lines in other wavelength regions.," For most of the elements which have lines in the relevant spectral region, abundances were available from clean lines in other wavelength regions." + We adopted these abundances for the spectrum synthesis., We adopted these abundances for the spectrum synthesis. + The abundances of elements which we could not determine from lines in other wavelength regions (e.g.. Co. Sc) were set to the solar abundance minus 2.64 ddex (1.e.. the abundances were scaled to the [Fe/H] of 004).," The abundances of elements which we could not determine from lines in other wavelength regions (e.g., Co, Sc) were set to the solar abundance minus $2.64$ dex (i.e., the abundances were scaled to the [Fe/H] of )." +" We first used ""C/C=10 in our spectrum synthesis. as constrained from the absence of features in other wavelength regions."," We first used $\element[][12]{C}/\element[][13]{C} = 10$ in our spectrum synthesis, as constrained from the absence of features in other wavelength regions." +" However. it was found that the fit could be considerably improved when it was assumed that carbon in is present in the form of""C. ie. ""C/C=,"," However, it was found that the fit could be considerably improved when it was assumed that carbon in is present in the form of, i.e., $\element[][12]{C}/\element[][13]{C} = \infty$." + Hence we adopted this value for our final fit., Hence we adopted this value for our final fit. + The result is loge(Th)= 0.96ddex., The result is $\log\epsilon (\mbox{Th}) = -0.96$ dex. + We estimate the uncertainty to be of the order of 0.15 ddex. which includes uncertainties introduced by the fit procedure. continuum placement. and the choice of stellar parameters.," We estimate the uncertainty to be of the order of $0.15$ dex, which includes uncertainties introduced by the fit procedure, continuum placement, and the choice of stellar parameters." + Non-local thermodynamical equilibrium (non-LTE) line formation calculations for have been carried out forBall.II. and using the methods described in ?.. 2.. and ?.. respectively.," Non-local thermodynamical equilibrium (non-LTE) line formation calculations for have been carried out for, and using the methods described in \citet{Mashonkinaetal:1999}, , \citet{Mashonkina/Gehren:2000}, and \citet{Mashonkina/Gehren:2001}, respectively." + A modified version of the code NONLTE3. based on the complete linearization method as described by ?.. has been used (?)..," A modified version of the code NONLTE3, based on the complete linearization method as described by \citet{Auer/Heasley:1976}, has been used \citep{Kampetal:2004}." + In cool giants.Bar. and are all. majority species.," In cool giants, and are all majority species." + Therefore. departures from LTE level populations of these species are mainly caused by radiative bound-bound transitions.," Therefore, departures from LTE level populations of these species are mainly caused by radiative bound-bound transitions." + Non-LTE effects strengthen the and lines compared with the LTEcase resulting in. negative non-LTE abundance correction Avy.=log£4log£4. and weaken the lines resulting in Άν of opposite sign.," Non-LTE effects strengthen the and lines compared with the LTEcase resulting in negative non-LTE abundance correction $\Delta_{\mbox{\tiny NLTE}} = \log +\varepsilon_{\mbox{\tiny NLTE}} - \log \varepsilon_{\mbox{\tiny LTE}}$ and weaken the lines resulting in $\Delta_{\mbox{\tiny NLTE}}$ of opposite sign." + According to our calculations for the investigated spectral lines A44.= 0.09ddex 4554A)). 0.23 ddex 4934A)). 1 0.03ddex 4129A)) απά 0.12 ddex 4215 A)).," According to our calculations for the investigated spectral lines $\Delta_{\mbox{\tiny NLTE}} = -0.09$ dex $4554$ ), $-0.23$ dex $4934$ ), $+0.03$ dex $4129$ ) and $-0.12$ dex $4215$ )." + In Figure + we show a comparison of the abundance pattern of with the solar r-process pattern. as listed in Table 5 of ?.. minus ddex.," In Figure \ref{Fig:CS29497_solar_r} we show a comparison of the abundance pattern of with the solar r-process pattern, as listed in Table 5 of \citet{Burrisetal:2000}, minus dex." + This is theaverage, This is theaverage +For each observation in which the CME or shock was visible the front was identified.,For each observation in which the CME or shock was visible the front was identified. + A nunmber of points alone (his front were then manually chosen., A number of points along this front were then manually chosen. + For the observations where the CME or shock was observed from both spacecralt. the front was localised in three dimensions using the tie-point method (Inhester2006:Byrneetal.2010;οἱal. 2009)..," For the observations where the CME or shock was observed from both spacecraft, the front was localised in three dimensions using the tie-point method \citep{Inhester:2006p2249,2010NatCo...1E..74B,Mierla:2010p7463,Temmer:2009p5011}." + As the CME or shock was only observed by one spacecraft in the ILI we used (he additional assumption of pseucdo-radial propagation. based on the direction derived from CORT and COR2 to localise the front (Maloneyetal.2009)..," As the CME or shock was only observed by one spacecraft in the HI field-of-view, we used the additional assumption of pseudo-radial propagation, based on the direction derived from COR1 and COR2 to localise the front \citep{Maloney:2009p6617}." + The resulting dala consisted of a series of points in 3D lor the CME aud shock for each observation time., The resulting data consisted of a series of points in 3D for the CME and shock for each observation time. + Figure 3((a) shows the 3D reconstruction of both the shock front and CAIE front. viewed perpendicular to the direction of propagation (assumed to be a cross-section).," Figure \ref{f3}( (a) shows the 3D reconstruction of both the shock front and CME front, viewed perpendicular to the direction of propagation (assumed to be a cross-section)." + The techniques for deriving the 3D coordinates of features in the CORT/2 ancl especially the ILE field-of-view are not. without error., The techniques for deriving the 3D coordinates of features in the COR1/2 and especially the HI field-of-view are not without error. + In the case of the event studied here. the CATE was close (107) to the plane-ol-skv of STEREO A. As a result. errors in position should be small.," In the case of the event studied here, the CME was close $<$ $^{\circ}$ ) to the plane-of-sky of STEREO A. As a result, errors in position should be small." + In order to compare with relations in Section 1.. the data was transformed. into a coordinate svstem centred on the CME.," In order to compare with relations in Section \ref{s_intro}, the data was transformed into a coordinate system centred on the CME." + To accomplish Chis. each CALE front was fit with an ellipse.," To accomplish this, each CME front was fit with an ellipse." + The centre coordinates of these fits were then used to collapse all the data on to a common coordinate svstem centred on the CME., The centre coordinates of these fits were then used to collapse all the data on to a common coordinate system centred on the CME. + The shock front was fit with Equation (5)).," The shock front was fit with Equation \ref{Eq6}) )," +The recent results reported [rom the WAIAP CMD probe have provided a wealth of new precision data for cosmology and astrophysics.,The recent results reported from the WMAP CMB probe have provided a wealth of new precision data for cosmology and astrophysics. + While some of the results have been anticipated by earlier ground-based. anisotropy. measurements (2000))). several new and surprising claims have been made on the basis of the first νου) WAIAP data (Spergeletal. (2003))).," While some of the results have been anticipated by earlier ground-based anisotropy measurements \cite{boomerang}) ), several new and surprising claims have been made on the basis of the first year WMAP data \cite{spergel}) )." + Particular attention has been paid to the facts that the power spectrum at low 1 (large angle) seems to differ somewhat from that predicted bv a single power-law inflationarv spectrum. and the fact that the epoch of reionization.," Particular attention has been paid to the facts that the power spectrum at low l (large angle) seems to differ somewhat from that predicted by a single power-law inflationary spectrum, and the fact that the epoch of reionization," +al.,al. + 2002) on SUBARU has been reported by (Ouchi et al., 2002) on SUBARU has been reported by (Ouchi et al. + 2003). however that survey is not as deep as the survey presented here and it targets a higher redshift (z24.86).," 2003), however that survey is not as deep as the survey presented here and it targets a higher redshift (z=4.86)." + A narrow-band image targeting z=3 LEGOs to the same depth as in the present survey. but obtained with the Suprime Camera will provide of order 500 candidates per field.," A narrow-band image targeting z=3 LEGOs to the same depth as in the present survey, but obtained with the Suprime Camera will provide of order 500 candidates per field." + Furthermore. it is mandatory that candidates based on narrow-band imaging are subsequently confirmed (or rejected) based on spectroscopy to make sure that conclusions based or surveys of LEGOs can be trusted.," Furthermore, it is mandatory that candidates based on narrow-band imaging are subsequently confirmed (or rejected) based on spectroscopy to make sure that conclusions based on surveys of LEGOs can be trusted." + A major reason for the success of the Lyman-Break surveys ts the spectroscopic confirmation of most of their candidates., A major reason for the success of the Lyman-Break surveys is the spectroscopic confirmation of most of their candidates. + Studies of LEGOs should follow this good example., Studies of LEGOs should follow this good example. + Furthermore. the measurement of several hundred redshifts in one field can be used to map out the underlying filamentary structure (Moller Fynbo 2001) and even provide an independent measurement of the cosmological constant (Weidinger et al.," Furthermore, the measurement of several hundred redshifts in one field can be used to map out the underlying filamentary structure ller Fynbo 2001) and even provide an independent measurement of the cosmological constant (Weidinger et al." + 2002)., 2002). + We thank our referee C. Steidel for comments that help us improve the discussion and the Paranal staff for excellent support during the visitor run in July 2002., We thank our referee C. Steidel for comments that help us improve the discussion and the Paranal staff for excellent support during the visitor run in July 2002. + JPUF and CL gratefully acknowledge the receipt of an ESO research fellowship., JPUF and CL gratefully acknowledge the receipt of an ESO research fellowship. + We acknowledge helpful discussions with Bruno Leibundgut. Nicolas Bouche. Masami Ouchi and Vincenzo Mainiert.," We acknowledge helpful discussions with Bruno Leibundgut, Nicolas Bouche, Masami Ouchi and Vincenzo Mainieri." + JPUF. gratefully acknowledge support from the Carlsberg Foundation., JPUF gratefully acknowledge support from the Carlsberg Foundation. +achieved through two different dispersers (two Amici prisms - see Oliva. 2000)).,achieved through two different dispersers (two Amici prisms - see \citealt{Oli00}) ). + The IFS is composed of several subsystems: The novel lenslet IFU concept upon which this spectrograph ts based (BIGRE. Antichietal... 2009)) allows the entrance slits plane to be made of images of the telescope focal plane and not of images of the telescope pupil. as orderec in the classical TIGER design (Baconetal...1995).," The IFS is composed of several subsystems: The novel lenslet IFU concept upon which this spectrograph is based (BIGRE, \citealt{JA09}) ) allows the entrance slits plane to be made of images of the telescope focal plane and not of images of the telescope pupil, as ordered in the classical TIGER design \citep{Ba95}." +. In this design. each lenslet is an afocal system with two powered surfaces.," In this design, each lenslet is an afocal system with two powered surfaces." + The thickness of the array 1s then given by the sum of the focal lengths of the lenslets of the two arrays., The thickness of the array is then given by the sum of the focal lengths of the lenslets of the two arrays. + The matt advantage of the BIGRE configuration over the TIGER one is that it allows a strong reduction of the cross-talk betweer adjacent lenslets as demostrated by Antichietal.(2009)., The main advantage of the BIGRE configuration over the TIGER one is that it allows a strong reduction of the cross-talk between adjacent lenslets as demostrated by \citet{JA09}. +. The microlens array is composed of 145x hexagonal lenslets with a pitch of 161.5 pan (corresponding to ~0.012 aresec)., The microlens array is composed of $145\times 145$ hexagonal lenslets with a pitch of 161.5 $\mu m$ (corresponding to $\sim$ 0.012 arcsec). + Each lenslet is masked with a circular aperture with a factor of 0.98 to avoid straylight., Each lenslet is masked with a circular aperture with a factor of 0.98 to avoid straylight. + The full field of view (FOV) of the instrument is a square with a side of 1.77 arcsec., The full field of view (FOV) of the instrument is a square with a side of 1.77 arcsec. + The total length of the whole instrument from the first surface of the IFU to the detector plane is 1061.89 mm., The total length of the whole instrument from the first surface of the IFU to the detector plane is 1061.89 mm. + A more detailed description of the whole instrument can be found. e.g. in Claudietal.(2010).," A more detailed description of the whole instrument can be found, e.g., in \citet{Cl10}." +. We exploited two software tools for our simulations: The CAOS system (Carbilletetal.2004) 1s an IDL based software that aims to simulate the behavior of a generic adaptive optic (AO) system from the atmospheric propagation of light to the sensing of the wavefront aberrations and the corrections through a deformable mirror., We exploited two software tools for our simulations: The CAOS system \citep{Car04} is an IDL based software that aims to simulate the behavior of a generic adaptive optic (AO) system from the atmospheric propagation of light to the sensing of the wavefront aberrations and the corrections through a deformable mirror. + This is done with a Fraunhofer approach. so that it cannot be used to properly evaluate the impact of Fresnel propagation (see Section ??)).," This is done with a Fraunhofer approach, so that it cannot be used to properly evaluate the impact of Fresnel propagation (see Section \ref{fresnelimpact}) )." + An end-to-end numerical tool has been developed for the simulation of the whole SPHERE instrument within the CAOS environment., An end-to-end numerical tool has been developed for the simulation of the whole SPHERE instrument within the CAOS environment. + [t contains detailed instrumental modeling of the Extreme adaptive optics systems. of IRDIS and ZIMPOL (Carbilletetal.2008).," It contains detailed instrumental modeling of the Extreme adaptive optics systems, of IRDIS and ZIMPOL \citep{Car08}." +.. A module simulating the SPHERE IFS has been also developed to properly take both the real and the imaginary parts of the image forming on the lenslet plane into account., A module simulating the SPHERE IFS has been also developed to properly take both the real and the imaginary parts of the image forming on the lenslet plane into account. + In principle. this could allow a complete treatment of the cross-talk among the lenslets when studying the impact of light propagation through the BIGRE.," In principle, this could allow a complete treatment of the cross-talk among the lenslets when studying the impact of light propagation through the BIGRE." + However. the execution of this module turned out to be very time consuming so that it was not possible to use it for a large number of detailed simulations.," However, the execution of this module turned out to be very time consuming so that it was not possible to use it for a large number of detailed simulations." + To overcome this difficulty. we used a shorter code that calculates the impact of the cross-talk between adjacent lenslets (coherent) and adjacent spectra (Incoherent) by providing the beam propagation over a sub-sample of 7 hexagonal lenslets.," To overcome this difficulty, we used a shorter code that calculates the impact of the cross-talk between adjacent lenslets (coherent) and adjacent spectra (incoherent) by providing the beam propagation over a sub-sample of 7 hexagonal lenslets." + This code is described in detail in Antichietal.(2009)., This code is described in detail in \citet{JA09}. +. After running this code we concluded that a value of the cross-talk equal to or less than 10 was completely adequate for meeting the objectives of our We then decided to use our (IDL oriented) code called CSP to perform all the simulations of light propagation within the IFS. while we decided to use the SPHERE CAOS package to provide real intensities over the [FU entrance focal plane as input for CSP.," After running this code we concluded that a value of the cross-talk equal to or less than $10^{-2}$ was completely adequate for meeting the objectives of our We then decided to use our (IDL oriented) code called CSP to perform all the simulations of light propagation within the IFS, while we decided to use the SPHERE CAOS package to provide real intensities over the IFU entrance focal plane as input for CSP." + For this. we performed simulations using the CAOS IRDIS module with 100 atmospheric phase screens at 64 different wavelengths ranging between 0.95 and 1.35 jun (or between 0.95 and 1.65 jun in the Y-H-mode case).," For this, we performed simulations using the CAOS IRDIS module with 100 atmospheric phase screens at 64 different wavelengths ranging between 0.95 and 1.35 $\mu m$ (or between 0.95 and 1.65 $\mu m$ in the Y-H-mode case)." + There are enough atmospheric screens is large enough to ensure that static speckles dominate noise. as expected in real cases. and to ensure that the PSF has an overall shape representing a realistic stellar," There are enough atmospheric screens is large enough to ensure that static speckles dominate noise, as expected in real cases, and to ensure that the PSF has an overall shape representing a realistic stellar" +relation )otwoeen colors and AL./£ ratios.,relation between colors and $M_*/L$ ratios. + Recdder ealaxies should have higher Af./£L ( sec. e.g. Bell&deJone2001.Belleta.2003.Porinarictal. 20013).," Redder galaxies should have higher $M_*/L$ ( see, e.g., \cite{bel01, bel03, port04}) )." + The slope of this lncar reation does uot depend onu the exact details of he history of star formation. Le. the assiuned IME.," The slope of this linear relation does not depend on the exact details of the history of star formation, i.e. the assumed IMF." + But depenudi i how nav stars are present at the low-Lass Cl of the stelar IME. the colorM./L curve shifts up aud ¢own.," But depending on how many stars are present at the low-mass end of the stellar IMF, the $M_*/L$ curve shifts up and down." + This is because low-mass stars οςntribute sieuificautlv to he mass of a population. but not as much o its Iuuinositv aux color (Bell&deJong2001 )).," This is because low-mass stars contribute significantly to the mass of a population, but not as much to its luminosity and color \cite{bel01}) )." + Iu the SPS scheme Salpeter’s (1955) IME overestimates the AM./L ratios of many of the ealaxies aud vioates the condiion of less disk mass than the lass of naxinmun disk., In the SPS scheme Salpeter's (1955) IMF overestimates the $M_*/L$ ratios of many of the galaxies and violates the condition of `less disk mass than the mass of maximum disk'. + To remedy this. Bell et al. (," To remedy this, Bell et al. (" +2003) scale down Salο TAIF and come up with a limit for the coloa-M./L reation above which the phivsical viabilitv is not euaraieed.,2003) scale down Salpeter's IMF and come up with a limit for the $M_*/L$ relation above which the physical viability is not guaranteed. + Their suggested relation is There are other INF that lead to shehthy cüffereut relations., Their suggested relation is There are other IMF that lead to slightly different relations. + For example. based on an analysis of the vertical velocity dispersion of stars. Bottcme| (1997) argues for a substanjallv subnaxiual AZ/L ratio for al cisk-donuirated galaxies;," For example, based on an analysis of the vertical velocity dispersion of stars, Bottema (1997) argues for a substantially submaximal $M/L$ ratio for all disk-dominated galaxies." + Alteruatively IKroupa. (2001) nroduces a urnover at the low-mass end of liis IME., Alternatively Kroupa (2001) introduces a turnover at the low-mass end of his IMF. + Tn Eq. (10 , In Eq. \ref{color}) ) +yt1ο sloe LF Lis jof sensitive to Vvzudations iu IMF. but he gyintercept is.," the slope $1.74$ is not sensitive to variations in IMF, but the $y$ -intercept is." + To obtain the equivaleut relation for staard Salpeters. Iroupas. aud Dottema’s IMP one shotId shift Eq. (10))," To obtain the equivalent relation for standard Salpeter's, Kroupa's, and Bottema's IMF one should shift Eq. \ref{color}) )" + and the plots in Fig. (1)), and the plots in Fig. \ref{ML-Color1}) ) + up aud down x about (0.15. -0.15. -0.35) dex. respectively (Belletal.20 By).," up and down by about (0.15, -0.15, -0.35) dex, respectively \cite{bel03}) )." + Iu Fig., In Fig. + 1 we contrast AL./L ratios of the three eravity nodes against the predictious of SPS. where we use t16 D-baud. Iuninosities of Sanders MeCGaugh (2002).," \ref{ML-Color1} we contrast $M_*/L$ ratios of the three gravity models against the predictions of SPS, where we use the B-band luminosities of Sanders McGaugh (2002)." + I reach frame the solid line is the best ft to the data points obtained from the analysis of the rotation curves., In each frame the solid line is the best fit to the data points obtained from the analysis of the rotation curves. + The theoretical SPS predictions of BeJl de Jong (2001). aud Bell et al. (," The theoretical SPS predictions of Bell de Jong (2001), and Bell et al. (" +2003) for differeut IAIFs are also plottec.,2003) for different IMFs are also plotted. + The slope of MONDI. 1.78+0.23. of MIOND?. LI+0.21. aud of MOND3. 1.75+0.26 are reasonabv close to that of Eq. (103).," The slope of MOND1, $1.78 \pm 0.23$, of MOND2, $1.81\pm 0.21$, and of MOND3, $1.75\pm 0.26$ are reasonably close to that of Eq. \ref{color}) )." +" The corresponding j--Mtercepts. —0,88dE0.11. LOGEOAZand 1.1320.15. respectively. are also in harmouy with that of Eq. (10))."," The corresponding -intercepts, $-0.88\pm0.14$, $-1.06\pm0.12$, and $-1.13\pm0.15$, respectively, are also in harmony with that of Eq. \ref{color}) )." + The uncertainties in slopes aud j-utercepts are iu the a7 error range., The uncertainties in slopes and -intercepts are in the $\sigma$ error range. + The errors in y—intercepts are simia lerough to enable one to distiuguish: between different IAIFs., The errors in -intercepts are small enough to enable one to distinguish between different IMFs. + MONDI falls somewhere between standard Salpeters and scaled Salpeter’s JAIF., MOND1 falls somewhere between standard Salpeter's and scaled Salpeter's IMF. + MOND2 agrees with Ίντοιpas aud scaed Salpeter’s IME., MOND2 agrees with Kroupa's and scaled Salpeter's IMF. + MOND? is iu good agreement with Iroupa's IMF., MOND3 is in good agreement with Kroupa's IMF. + The slope for MOC. 1.06+0.21. cannot be reconciled with SPS predictions.," The slope for MOG, $1.06 \pm 0.21$, cannot be reconciled with SPS predictions." + The case of NEW js. also questionable., The case of NFW is also questionable. +" Although the slope. 2.33στ, Is Consistcut witi L7L. considering its huge error bar. the dispersion of the simulated data poiuts is too large to conuclide a meaninueful color-M./£ correlation."," Although the slope, $2.33\pm 0.67$, is consistent with 1.74, considering its large error bar, the dispersion of the simulated data points is too large to conclude a meaningful $M_*/L$ correlation." + Auv alternative eravitv can bave a dark matter equivalent., Any alternative gravity can have a dark matter equivalent. + Deviations from the Newtoman eravity can be attributed to a hypothetical dark cutity and a dark density profile calculated through Poisson's equation. for instance.," Deviations from the Newtonian gravity can be attributed to a hypothetical dark entity and a dark density profile calculated through Poisson's equation, for instance." + Oue feature however distinetishes this interpretatiou fron the couventional CDAL scenarios., One feature however distinguishes this interpretation from the conventional CDM scenarios. + Here. there is a well defiuecL relation between the barvonie matter and its so-iuterpreted dark companion.," Here, there is a well defined relation between the baryonic matter and its so-interpreted dark companion." + While in CDM models. Dbarvouic axd «ark matters miav co-exist iudepenudoeutlv.," While in CDM models, baryonic and dark matters may co-exist independently." + Iu our opiuio1. the reason for the good agreenen of MOND with SPS precictions and non-conipliauce of CDAL wit[um it des in the existence or non-existence of fus relalon between tιο observable ancl nou-observale natters.," In our opinion, the reason for the good agreement of MOND with SPS predictions and non-compliance of CDM with it lies in the existence or non-existence of this relation between the observable and non-observable matters." + Iu MOND. xuvonic matter plavs a pivotal vole.5 :id the cark eutitv owes its existence to it.," In MOND, baryonic matter plays a pivotal role, and the dark entity owes its existence to it." + This is not the case iu CDAL, This is not the case in CDM. + Dark liater is allowed to play a role indepeudeutly othe observale matter., Dark matter is allowed to play a role independently ofthe observable matter. + As for MOG. we are not in à posilon to express ani ΟΡΙΟ.," As for MOG, we are not in a position to express an opinion." + Le Us sSUuDunuuize our conclusion: a) the SPS scheme can distiicuish between different eravity models aud. b) the two tevecther cau choose between different INMEs., Let us summarize our conclusion: a) the SPS scheme can distinguish between different gravity models and b) the two together can choose between different IMFs. + The luere ac that a eravity theory reproduces the observed rotation curves satisfactorily docs not tell the whole story., The mere fact that a gravity theory reproduces the observed rotation curves satisfactorily does not tell the whole story. + At least on galactic scales. cdiwuamiics of spirals cast doubt on the viabilitv of the classical theories of eravitation.," At least on galactic scales, dynamics of spirals cast doubt on the viability of the classical theories of gravitation." + A umber of alternative theories are capable of reproducing the rotation curves of spirals with acceptable detail. a uoutrivial fact that deserves atteutiou.," A number of alternative theories are capable of reproducing the rotation curves of spirals with acceptable detail, a nontrivial fact that deserves attention." + In this paper we used two alernative theories of exavitation. MOND and MOC. and a CDAL model to de«uce the dynamics of a well-sizec saluple of hieh- aud ow-urface brightness galaxy types. aud checked the results agaiust observations.," In this paper we used two alternative theories of gravitation, MOND and MOG, and a CDM model to deduce the dynamics of a well-sized sample of high- and low-surface brightness galaxy types, and checked the results against observations." + The τοσο» are not equivaeut although they all simulate the rotation curves 1h niore or less to the same deeree of accuracy..," The models are not equivalent although they all simulate the rotation curves in more or less to the same degree of accuracy,." + Tn MOND aud MOC. rotation curves are constructed with oily one free adjustable parameter. the stellar mass-to-ligh: ratio.," In MOND and MOG, rotation curves are constructed with only one free adjustable parameter, the stellar mass-to-light ratio." + This is iu contrast to the CDM model. where an aditional parameter is needed to describe the dark conipovent.," This is in contrast to the CDM model, where an additional parameter is needed to describe the dark component." + There are cases of buleed salaxies where fits to observations lead to lower ML ratios for the bulee than for the disk., There are cases of bulged galaxies where fits to observations lead to lower $M_*/L$ ratios for the bulge than for the disk. + This might be owing to the low resolution of the III data and to the inner e(e). aud/or uucertaiu size of the nilee.," This might be owing to the low resolution of the HI data and to the inner $v(r)$, and/or uncertain size of the bulge." + Stelay population svuthesis models impose constraints O1 ALL., Stellar population synthesis models impose constraints on $M_*/L$. +" Redder galaxies should lave higher AL.1, ratios.", Redder galaxies should have higher $M_*/L$ ratios. + MOND with different interpolating functions meets his expectation. albeit with different I\IFs.," MOND with different interpolating functions meets this expectation, albeit with different IMFs." + This is ronrarkable. because there is πο explicituplicit COLLCCjon between the basic tenets of the SPS aud MOND.," This is remarkable, because there is no explicit/implicit connection between the basic tenets of the SPS and MOND." + On the other haud. MOC does uot meet the SPS," On the other hand, MOG does not meet the SPS" +2002).,. +. Theoretical studies of magnetic fields in AGB stars have produced a range of predicted magnetic field magnitudes and morphologies., Theoretical studies of magnetic fields in AGB stars have produced a range of predicted magnetic field magnitudes and morphologies. +Dlackmanεἰaf(2001) have proposed a conventional aw dynamo driven by an inner. differentiallv-rotating stellar core.,"\citet{blackman01} have proposed a conventional $\alpha +\omega$ dynamo driven by an inner, differentially-rotating stellar core." + Photospheric fields of theorder of ~400 G are predicted in (his model. but for the ease of an isolated (non-binary) ACB star. this model would need a dvnamo resupply mechanism {ο sustain a magnetic field for long enough to shape global mass oss (Nordhaus.Blackman.&Frank2007).," Photospheric fields of theorder of $\sim 400$ G are predicted in this model, but for the case of an isolated (non-binary) AGB star, this model would need a dynamo resupply mechanism to sustain a magnetic field for long enough to shape global mass loss \citep{nordhaus07}." +. The role of magnetic fields in shaping the global uass-]oss process is considered by (Garcfa-Seguraetal. 1999)., The role of magnetic fields in shaping the global mass-loss process is considered by \citep{garcia-segura99}. . +. In contrast. a convective a? dvnamo bas been proposed by Soker&Zoabi(2002): in this model. turbulent dvnanmo fields nav produce magnetic cool spots (ο10—100 G) that regulate local dust formation. but have 10 global shaping role in mass-loss (Soker2002)..," In contrast, a convective $\alpha^2 \omega$ dynamo has been proposed by \citet{soker02}; in this model, turbulent dynamo fields may produce magnetic cool spots $\sim 10-100$ G) that regulate local dust formation, but have no global shaping role in mass-loss \citep{soker02a}." + Droader arguments against (he proposition that elobal magnetic fields can shape planetary nebulae are summarized by Soker(2006)., Broader arguments against the proposition that global magnetic fields can shape planetary nebulae are summarized by \citet{soker06}. +. In a nunerical MIID study. Dorch(2004) have examined a convective dynamo in a supergianl abmosphere and report predicted local fields of up to ~500 G. These fields would have significant. dvnanmical influence. albeit localized with low filling factors.," In a numerical MHD study, \citet{dorch04} have examined a convective dynamo in a supergiant atmosphere and report predicted local fields of up to $\sim 500$ G. These fields would have significant dynamical influence, albeit localized with low filling factors." + Recent water maser polarimetrv of late-tvpe. evolved stars has found evidence for >107 mG magnetic fields at the shell radius of (he HI5O0 masers: this extrapolates to a surface field of ~10? G (Vlenunines.vanLangevelde.&Diamond2005).," Recent water maser polarimetry of late-type, evolved stars has found evidence for $> 10^2$ mG magnetic fields at the shell radius of the $_2$ O masers; this extrapolates to a surface field of $\sim 10^2$ G \citep{vlemmings05}." +. Recent optical spectroscopy of unobsceured central stars in planetary nebulae (PNe) has measured photospheric magnetic fields of several Κ in these stars (Jordan.Werner.&O'Toole2005)., Recent optical spectroscopy of unobscured central stars in planetary nebulae (PNe) has measured photospheric magnetic fields of several kG in these stars \citep{jordan05}. +. Magnetic fields of (his magnitude would be dynamically significant ancl of great. astvoplvsical importance {ο related studies of collimation mechanisms for PPN (Aleixnere£al.2004)., Magnetic fields of this magnitude would be dynamically significant and of great astrophysical importance to related studies of collimation mechanisms for PPN \citep{meixner04}. +. A small subset of stellar water maser sources (hat show linear jels also appear lo be magnetically collimated (Vlemmings.Diamond.&Imai2006)., A small subset of stellar water maser sources that show linear jets also appear to be magnetically collimated \citep{vlemmings06}. +. As part of a svnoplic monitoring campaign of the J=1—0 SiO maser emission toward the Mira. variable TX Cam at fine time-ssunpling. (his paper presents linear polarization images for the [ist 43 epochs observed. over a pulsation phase range ὁ=0.68 lo ὁ=1.82.," As part of a synoptic monitoring campaign of the $J=1-0$ SiO maser emission toward the Mira variable TX Cam at fine time-sampling, this paper presents linear polarization images for the first 43 epochs observed, over a pulsation phase range $\phi=0.68$ to $\phi=1.82$." + The reduction aud analvsis of the corresponding total intensity data was published earlier by Diamoud&Ixemball(2003.hereafterPaperI).., The reduction and analysis of the corresponding total intensity data was published earlier by \citet[hereafter Paper I]{diamond03}. + This paper serves (ο present the time-series of linear polarization images. and provides an analvsis of the linear polarization morphology and early results on associated proper motions.," This paper serves to present the time-series of linear polarization images, and provides an analysis of the linear polarization morphology and early results on associated proper motions." + The associated circular polarization data will be presented in a future paper., The associated circular polarization data will be presented in a future paper. + TX Cam is aan M8-MIO Mira variable with a pulsation period of 557.4 davs (IxXholopovetal.1985).. and has a Mira. period-Iuminositw. distance estimate of 390pe (Olivier.Whitelock.&Marang 2001)...," TX Cam is a an M8-M10 Mira variable with a pulsation period of 557.4 days \citep{kholopov85}, and has a Mira period-luminosity distance estimate of $\sim 390$pc \citep{olivier01}. ." + Itwas first observed in total intensity. VLBI observations by Diamondetal. (1994)., Itwas first observed in total intensity VLBI observations by \citet{diamond94}. +. Recent simultaneous total intensity VLBI imaging of the e={1.2}.J1—0 transitions is reported by Yiefal. (2005)..," Recent simultaneous total intensity VLBI imaging of the $v=\{1,2\},\ J=1-0$ transitions is reported by \citet{yi05}. ." +The data used here are the same as (hose used in BS. which provides further details.,"The data used here are the same as those used in BS, which provides further details." + They consist of a 60 day sequence of 1024 [ull disk Doppler velocity images obtained by the MDI instrument (Scherrer.οἱal.1995). on the SOLIO spacecraft. covering the time period 1996 May 24 to 1996 July 22.," They consist of a 60 day sequence of $^2$ full disk Doppler velocity images obtained by the MDI instrument \citep{scherrer} on the SOHO spacecraft, covering the time period 1996 May 24 to 1996 July 22." + These images were derotated and averaged to a GO minute cadence. remapped to a uniform grid in longitude and latitude and binned down to have a resolution of 1200 points covering 360° of longitude aud GOO points in latitude between x90.," These images were derotated and averaged to a 60 minute cadence, remapped to a uniform grid in longitude and latitude and binned down to have a resolution of 1200 points covering $^\circ$ of longitude and 600 points in latitude between $\pm 90^\circ$." + A temporal average was removed as part of this process., A temporal average was removed as part of this process. + SUips in longitude were extracted. multiplied by a weighting function. tracked. αἱ a latitude dependent rate. zero padded to 360° by 12° and passed through a 3 dimensional Fourier transform.," Strips in longitude were extracted, multiplied by a weighting function, tracked at a latitude dependent rate, zero padded to $^\circ$ by $^\circ$ and passed through a 3 dimensional Fourier transform." + The weighting function is a product of a function of 0 and a function of o., The weighting function is a product of a function of $\theta$ and a function of $\phi$. + The function of 0 is 1 for |o00|<4.5°. 0 for |90|>5.5 and falls off as a cosine bell for 4.5°€|o0|x5.5. where 08 is the distance from the target latitude.," The function of $\theta$ is 1 for $|\delta\theta| \le 4.5^\circ$, 0 for $|\delta\theta| \ge 5.5^\circ$ and falls off as a cosine bell for $4.5^\circ \le |\delta\theta|\le 5.5^\circ$, where $\delta\theta$ is the distance from the target latitude." + Two weighting functions in ó were applied., Two weighting functions in $\phi$ were applied. + One is optimized for recovering V5. the other for recovering Vo.," One is optimized for recovering $V_\phi$ , the other for recovering $V_\theta$." +" The first weighting function is given by HW,=Asigno/sino+0.01. where 4A, is an apodization function which is 0 for |o|=O° and for |o|>80°. E for 10«Jol<τη and cosine bell apodized for 0°<|o|LO and τοῦ80^\circ$, 1 for $ 10^\circ < |\phi| <70^\circ$ and cosine bell apodized for $0^\circ < |\phi| < 10^\circ$ and $70^\circ < |\phi| < 80^\circ$." + The second weighting function is given by Wy=docos?o+0.01. where Ay is only apodized between 70 and 807.," The second weighting function is given by $W_\theta = A_\theta / \sqrt{\cos^2 \phi + 0.01}$, where $A_\theta$ is only apodized between $70^\circ$ and $80^\circ$." + The small bias οἱ 0.01 was added to avoid singularities., The small bias of 0.01 was added to avoid singularities. + Figure 1 shows the power as a function of direction of propagation., Figure \ref{pphi} shows the power as a function of direction of propagation. + Note (that there is significantly more power in (he prograde than in the retrograde direction. consistent. with the findings of GDS.," Note that there is significantly more power in the prograde than in the retrograde direction, consistent with the findings of GDS." + Also there is little power in the N-S direction where GDS did observe significant. power., Also there is little power in the N-S direction where GDS did observe significant power. + Since (he method used by GDS is sensitive to both components of the velocity (his indicates (hat (he displacement. at least at the surface. is almost entirely in the direction of propagation.," Since the method used by GDS is sensitive to both components of the velocity this indicates that the displacement, at least at the surface, is almost entirely in the direction of propagation." +" Given the above waves in the © direction should show modulation given by A, when using the 1, weighting and waves in the 8 direction should be modulated by sly when using Wy.", Given the above waves in the $\phi$ direction should show modulation given by $A_\phi$ when using the $W_\phi$ weighting and waves in the $\theta$ direction should be modulated by $A_\theta$ when using $W_\theta$. + For other directions the modulation is more complex and only waves propagating near the longitude and Latitude directions are studied here., For other directions the modulation is more complex and only waves propagating near the longitude and latitude directions are studied here. + The velocity projection also means that waves traveling in latitude are essentially unobservable at the equator., The velocity projection also means that waves traveling in latitude are essentially unobservable at the equator. + At high latitudes the foreshortening becomes a problem., At high latitudes the foreshortening becomes a problem. + Fieure 2 shows examples of the resulting power spectra., Figure \ref{plotx} shows examples of the resulting power spectra. + Note the different power level in prograde waves (lower right) ancl retrograde waves (upper right) in the left panel., Note the different power level in prograde waves (lower right) and retrograde waves (upper right) in the left panel. + In (he rightpanel there is excess power in the southbound waves (upper right) relative to the northbound waves (lower right)., In the rightpanel there is excess power in the southbound waves (upper right) relative to the northbound waves (lower right). +have a decreased turbulent intensity.,have a decreased turbulent intensity. + This is seen in observational scintillation techniques (??) and data (?) in the fast solar wind.," This is seen in observational scintillation techniques \citep{Woo_etal1995,Woo1996} and data \citep{MarschTu1990} in the fast solar wind." + The observed values for An/n are as low as 0.3% at distances <0.1 AU , The observed values for $\Delta n/n$ are as low as $0.3\%$ at distances $<0.1$ AU \citep{Woo_etal1995}. +A much higher magnitude of plasma wave energy density(?).. was achieved close to the Sun with smaller levels of fluctuations., A much higher magnitude of plasma wave energy density was achieved close to the Sun with smaller levels of fluctuations. +" To estimate how density flucutations might radially evolve, we varied the initial conditions of the simulations."," To estimate how density flucutations might radially evolve, we varied the initial conditions of the simulations." + We used a variety of different initial electron beam spectral indices (6 in Equation 3)) and different radial dependence of density fluctuations (4 in Equation 9))., We used a variety of different initial electron beam spectral indices $\delta$ in Equation \ref{init_f}) ) and different radial dependence of density fluctuations $\psi$ in Equation \ref{pertm3}) ). +" Using the resulting fluence spectra near the Earth for each simulation, we compared the spectral indices above and below the break energy (Figure 9))."," Using the resulting fluence spectra near the Earth for each simulation, we compared the spectral indices above and below the break energy (Figure \ref{fig:SI_high_low_vary}) )." + The spectral index becomes smaller below the break energy for larger values of Ψ., The spectral index becomes smaller below the break energy for larger values of $\psi$. + We have also overplotted the correlation of spectral indices above and below the break energy of peak flux measurements taken from a statistical survey (?) of impulsive electron events detected by the three-dimensional Plasma and Energetic Particles experiment on the WIND spacecraft (?).., We have also overplotted the correlation of spectral indices above and below the break energy of peak flux measurements taken from a statistical survey \citep{Krucker_etal09} of impulsive electron events detected by the three-dimensional Plasma and Energetic Particles experiment on the WIND spacecraft \citep{Lin_etal95}. + A level of fluctuations with i» around 0.25 would give a similar correlation to the observational line., A level of fluctuations with $\psi$ around 0.25 would give a similar correlation to the observational line. +" We note, however, that the observational line presented from ? fitted a large scatter of data points."," We note, however, that the observational line presented from \citet{Krucker_etal09} fitted a large scatter of data points." + The ratio of low:high spectral index for all simulated results presented in figure 9 lies between 0.42 and 0.58 which is within the narrow range presented in ?.., The ratio of low:high spectral index for all simulated results presented in figure \ref{fig:SI_high_low_vary} lies between 0.42 and 0.58 which is within the narrow range presented in \citet{Krucker_etal09}. +" A variety of simulation variables can affect the energy of the spectral break at the Earth: the model of radial background density decrease, the density fluctuations, the initial spectral index of the beam, the number density of injected electrons, the spatial distribution of injected electrons, the temporal nature of the injection, and the initial coronal background density where the electrons are injected."," A variety of simulation variables can affect the energy of the spectral break at the Earth: the model of radial background density decrease, the density fluctuations, the initial spectral index of the beam, the number density of injected electrons, the spatial distribution of injected electrons, the temporal nature of the injection, and the initial coronal background density where the electrons are injected." + The spectral index below the break energy of the resultant double power-law in fluence spectra near the Earth is increased when density fluctuations have a larger effect on the level of induced plasma waves., The spectral index below the break energy of the resultant double power-law in fluence spectra near the Earth is increased when density fluctuations have a larger effect on the level of induced plasma waves. +" It is important to note, however, the spectra below the break energy is only approximately a power-law."," It is important to note, however, the spectra below the break energy is only approximately a power-law." + The presence of density fluctuations causes fine structure to be present., The presence of density fluctuations causes fine structure to be present. +" A bump around 10-20 keV was found, caused by acceleration of electrons at the back of the beam through absorbed plasma waves."," A bump around 10-20 keV was found, caused by acceleration of electrons at the back of the beam through absorbed plasma waves." + The onset of this bump appears to be close to the Sun where plasma wave energy density is high., The onset of this bump appears to be close to the Sun where plasma wave energy density is high. + The magnitude of this bump is reflected in the size of the spectral index error bars in Figure 9 with a larger bump corresponding to a larger error., The magnitude of this bump is reflected in the size of the spectral index error bars in Figure \ref{fig:SI_high_low_vary} with a larger bump corresponding to a larger error. +" With the prospect of Solar Orbiter and Solar Probe Plus, it is very attractive to extend these studies further to understand the spectral evolution of the electron beam between the Sun and the Earth."," With the prospect of Solar Orbiter and Solar Probe Plus, it is very attractive to extend these studies further to understand the spectral evolution of the electron beam between the Sun and the Earth." + This work is partially supported by a STFC rolling grant and STFC Advanced Fellowship (EPK)., This work is partially supported by a STFC rolling grant and STFC Advanced Fellowship (EPK). +" Financial support by the Royal Society grant (RG090411), and by the European Commission through the SOLAIRE"," Financial support by the Royal Society grant (RG090411), and by the European Commission through the SOLAIRE" +Qvlich we neglect) could significautly affect the results.,(which we neglect) could significantly affect the results. +" Iu order to bracket the size of this effect we consider two spectral shapes for the radiation field. J~p"" with a=1 ando d."," In order to bracket the size of this effect we consider two spectral shapes for the radiation field, $J \sim \nu^{\alpha}$ with $\alpha = +1$ and $-1$." + The softer spectrum. 6=1. is the spectral slope expected in this cnerey range from a typical OB stir.," The softer spectrum, $\alpha =1$, is the spectral slope expected in this energy range from a typical OB star." + The steeper slope. à=1. is the type of spectral shape one nuelit expect for a backeround radiation field built up frou a number of sources;," The steeper slope, $\alpha = -1$, is the type of spectral shape one might expect for a background radiation field built up from a number of sources." +" The cross section (iu. cm?) for plioto-detaclinent is given by (yang Shapiro 1987. Abel 1997). where Pp is Plauck's coustaut aud Jp,=0.755 eV. The photo-detachineut rate coefficient ρε Can then be calculated with hprj;=13.6 eV. For the two spectral shapes given above we find Aycan,=1.910.Frys | fora= Laud G.[410 Ες τα Lwhere Fry isthe radiation flux in|."," The cross section (in $^2$ ) for $^-$ photo-detachment is given by (Kang Shapiro 1987, Abel 1997), where $h_P$ is Planck's constant and $h_P \nu_{th} = 0.755$ eV. The $^{-}$ photo-detachment rate coefficient $k_{detach}$ can then be calculated with $h_P\nu_{H}=13.6$ eV. For the two spectral shapes given above we find $k_{detach}=1.9 +\times 10^{-12} F_{LW}$ $^{-1}$ for $\alpha = 1$ and $6.4 \times +10^{-11}F_{LW}$ $^{-1}$ for $\alpha = -1$ where $F_{LW}$ is the radiation flux in." + The dominant processes other than photodoetachuneut that determine the abundance of are shown in Equations 1 aud 20 with rate cocffcieuts Ay aud hs. crespeetivelv," The dominant processes other than photodetachment that determine the abundance of $^-$ are shown in Equations \ref{eq:hminusform} and \ref{eq:hminus2h2} with rate coefficients $k_7$ and $k_8$, respectively." + Furthermore the processes proceed sufficieutlv rapidly compared to the chemistry that ag ds well approximated by its equilibrimn value CAbel 1997) such that (αποπιο photo-detachiueut): We see that photo-detachment can be neglected as loug as μικι , Furthermore the processes proceed sufficiently rapidly compared to the chemistry that $n_{H-}$ is well approximated by its equilibrium value (Abel 1997) such that (including photo-detachment): We see that $^-$ photo-detachment can be neglected as long as $k_{detach} \ll k_8 n_H$. +"This leads to a condition ou the hydrogen number density vy29KFkgqaü4q,/[k«s that can be checked against the conditions found withinthe collapsing structures.", This leads to a condition on the hydrogen number density $n_H \gg k_{detach}/k_8$ that can be checked against the conditions found withinthe collapsing structures. + For gas temperatures found iu our simulations. ksz1.13«10? em?s 1 (Abel 1997).," For gas temperatures found in our simulations, $k_8 \approx 1.43 \times +10^{-9}$ $^3$ $^{-1}$ (Abel 1997)." +" Thus photo-detachiment is uniuportant provided that ""ngLOISCErp/1071) for the steeper a=1l spectruni or. ἃ πιο weaker coustraiut. that ο2»1.1«10""(Fg/1071) for the softer a=1 spectu."," Thus photo-detachment is unimportant provided that $n_H \gg 0.045(F_{LW}/10^{-21})$ for the steeper $\alpha = -1$ spectrum or, a much weaker constraint, that $n_H +\gg 1.4 \times 10^{-3}(F_{LW}/10^{-21})$ for the softer $\alpha = 1$ spectrum." + Let us check these constraints against the hvdrogen ΠΡΟ densities found within the first peak to collapse at our lughest fux level (Fru1ο21 ip shown iu Figure 1.., Let us check these constraints against the hydrogen number densities found within the first peak to collapse at our highest flux level $F_{LW} = 10^{-21}$ ) shown in Figure \ref{fig:P0evolve}. + At +=25. before cooling commuciuces. hydrogen uuuber densities iu the core region where will form already exeeed pg;7.," At $z=25$, before cooling commences, hydrogen number densities in the core region where will form already exceed $n_H \sim 7$." + This umber density satisfies the constraint for the steepest spectrum (60= 1) by more than two orders of magnitude so that the neglect of photo-detachiuent does not significautly affect our results., This number density satisfies the constraint for the steepest spectrum $\alpha =-1$ ) by more than two orders of magnitude so that the neglect of $^{-}$ photo-detachment does not significantly affect our results. + Analytical arguineuts often provide qualitative insielt into the key plysical principles at work iu a complex process., Analytical arguments often provide qualitative insight into the key physical principles at work in a complex process. + Since we have seen that the fraction is close to photodissociation equilibrium throughout our objects. we use these equilibrium abundances to develop analytical estimates of the critical fraction and mass thresholds for collapse.," Since we have seen that the fraction is close to photodissociation equilibrium throughout our objects, we use these equilibrium abundances to develop analytical estimates of the critical fraction and mass thresholds for collapse." + We adopt the simplified collapse criterion. simular to that of Tegmark (1997). that au object will collapse if the cooling time f. Given by Equation 10.. is less than the Uhibble time fuabire. defined in Equation 12..," We adopt the simplified collapse criterion, similar to that of Tegmark (1997), that an object will collapse if the cooling time $t_{cool}$, given by Equation \ref{eq:tcool}, is less than the Hubble time $t_{Hubble}$, defined in Equation \ref{eq:thub}. ." + Thus The equality|NS eives the critical πο. density of necessary for collapse., Thus The equality gives the critical number density of necessary for collapse. +" Solving for ""p we obtain Iu this equation. A is. in our case. the Lepp Shull (1983) cooling fuuction."," Solving for $n_{H2}^{crit}$ we obtain In this equation, $\Lambda$ is, in our case, the Lepp Shull (1983) cooling function." + In the low density limit (ayo«Lot )) for temperatures 0017 drops adiabatically with ον~O.OLBS(L|:)2K .," The gas temperature after thermal decoupling from the CMB at $z \sim 200$ but before reionization $z > +7$ drops adiabatically with $T_{IGM} \sim 0.0135(1+z)^2 \kel$ ." + If we asstime that the barvous in the central regions of dark matter halos were never shocked. we fix the masiuuun density that can be reached hy adiabatic evolution IS mo»ΠΡ...," If we assume that the baryons in the central regions of dark matter halos were never shocked, we find the maximum density that can be reached by adiabatic evolution is $n_{max} = +n_B(T_{vir}/T_{IGM})^{1/(\gamma-1)}$." + Tere γι denotes the virial temperature of the dark matter halo aud 5=5/3 is the ratio of specific heats for au ideal eas., Here $T_{vir}$ denotes the virial temperature of the dark matter halo and $\gamma = 5/3$ is the ratio of specific heats for an ideal gas. + This gives a collapse redshift independent maximal density as a function of virial temperature for adiabatic evolution of the cloud as tyrescστοCP/A000y7?," This gives a collapse redshift independent maximal density as a function of virial temperature for adiabatic evolution of the cloud as $n_{max} \sim 187 +\Omega_b h^2 (T/1000)^{3/2}$." + We expect adiabatic evolution to be a eood approximation for low mass clouds. but caution that for larger mass structures niereiues aud shock heating become inportaut.," We expect adiabatic evolution to be a good approximation for low mass clouds, but caution that for larger mass structures merging and shock heating become important." +" Inserting the expressions for »,,,,4 and we. into Equation 25. woe obtain: The structures that are first able to collapse iuour simulations all have virial temperatures between 700Is and30001. where the exponential iu Equation 21. cau be approximated by a power law. e.MPFzz0. Is(P/Lo0g)8"", "," Inserting the expressions for $n_{max}$ and $x_e$ into Equation \ref{eq:fluxcrit} we obtain: The structures that are first able to collapse inour simulations all have virial temperatures between $700 \kel$ and$3000 \kel$ where the exponential in Equation \ref{eq:fluxcrit2} can be approximated by a power law, $e^{-730/T} \approx 0.48(T/1000)^{0.7}$ ." +Tusciting this approximation into Equation 21— aud, Inserting this approximation into Equation \ref{eq:fluxcrit2} and +The backeround-sublractecl spectrum shown in Figure 3. has many emission lines.,"The background-subtracted spectrum shown in Figure \ref{spectrum} + has many emission lines." + We hence determined the peak energies of the 5 brightest lines with a phenomenological model. a power-law continuumn plus Gaussian lines.," We hence determined the peak energies of the 5 brightest lines with a phenomenological model, a power-law continuum plus Gaussian lines." + The most strong line structures are the peak at 0.55 keV and the Lump al 0.67 keV. These energies are nearly equal to the Ίνα and Lyman a lines of He- and II-like oxvgen. hence are attributable to highly ionized oxvgen.," The most strong line structures are the peak at 0.55 keV and the hump at 0.67 keV. These energies are nearly equal to the $\alpha$ and Lyman $\alpha$ lines of He- and H-like oxygen, hence are attributable to highly ionized oxygen." + Likewise. the other clear peaks at 0.87. 1.31. and 1.76 keV are most likely IIe-like Ka of Ne. Mg and Si. respectively.," Likewise, the other clear peaks at 0.87, 1.31, and 1.76 keV are most likely He-like $\alpha$ of Ne, Mg and Si, respectively." + However. in detail all (he observed line energies are svstematically smaller than those of the relevant atomic data.," However, in detail all the observed line energies are systematically smaller than those of the relevant atomic data." + These apparent energy shifts have been usually observed ina voung SNK plasma in non-equilibrium ionization (NEI)., These apparent energy shifts have been usually observed in a young SNR plasma in non-equilibrium ionization (NEI). +" The ""energy shift in (his case is due to the different line ratio of many sub-levels aud/or different ionization states.", The “energy shift” in this case is due to the different line ratio of many sub-levels and/or different ionization states. + The oxvgen Lyman à is isolated from the other lines of different ionization states. hence the NEI effect eives no energy shift.," The oxygen Lyman $\alpha$ is isolated from the other lines of different ionization states, hence the NEI effect gives no energy shift." + Still we see apparent down-shilt of the observed line energy from that of the laboratory data., Still we see apparent down-shift of the observed line energy from that of the laboratory data. + Ie-like Ίνα lines are complex of many fine structures with the split-energv of at most ~ 25 eV (for He-like silicon)., He-like $\alpha$ lines are complex of many fine structures with the split-energy of at most $\sim$ 25 eV (for He-like silicon). + Although the energy shift of IHe-like Ίνα lines due to NET should be smaller (han (his split-energv. the observed energy shilts are systematically larger than the split-energv.," Although the energy shift of He-like $\alpha$ lines due to NEI should be smaller than this split-energy, the observed energy shifts are systematically larger than the split-energy." + We therefore regard (hat the apparent energy shifts ave due mainly to energy calibration errors. hence fine-tuned (he energy gain to reduce by3.8%... the average shift of the 5 brightest lines.," We therefore regard that the apparent energy shifts are due mainly to energy calibration errors, hence fine-tuned the energy gain to reduce by, the average shift of the 5 brightest lines." + We then fitted the spectrum with a thin thermal plasma model in NEI calculated by Borkowskietal.(2001a)., We then fitted the spectrum with a thin thermal plasma model in NEI calculated by \citet{borkowski2001a}. +". The abundances of C. N. O. Ne. Meg. οἱ, 5. and Fe in the plasma were treated to be [ree parameters. whereas those of the other elements were fixed to the solar values (Anders&Grevesse1989)."," The abundances of C, N, O, Ne, Mg, Si, S, and Fe in the plasma were treated to be free parameters, whereas those of the other elements were fixed to the solar values \citep{anders}." +. The absorption column was calculated using the cross sections bv Morrison& with the solar abundances., The absorption column was calculated using the cross sections by \citet{morrison} with the solar abundances. + Since this NEI model exhibited svstematic data excess at high enerev above 2 keV. we added a power-law component and the fit improved. dramatically.," Since this NEI model exhibited systematic data excess at high energy above 2 keV, we added a power-law component and the fit improved dramatically." + Figure 3. and Table 1. show the best-fit. models (dashed. and solid lines for thermal and power-law components) ancl parameters. respectively.," Figure \ref{spectrum} and Table \ref{thermal_para} show the best-fit models (dashed and solid lines for thermal and power-law components) and parameters, respectively." + Instead of the phenomenological power-law model. we applied in the NSPEC package as a more physical model.," Instead of the phenomenological power-law model, we applied in the XSPEC package as a more physical model." + Details of theΡΟΗ model fitting are given in A.1.., Details of the model fitting are given in \ref{srcut}. + The best-li€ ego 1s 9.2 (8.610.3) xLOM Lz. with better reduced 4? of 389.0/215 than that of the power-law model of 447.9/215 (see Table 1)).," The best-fit $\nu_{\rm rolloff}$ is 9.2 (8.6–10.3) $\times 10^{16}$ Hz, with better reduced $\chi^2$ of 389.0/215 than that of the power-law model of 447.9/215 (see Table \ref{thermal_para}) )." + As [or the thermal components. we also tried the fitting with a plane shock model (XSPEC model epshock) plus either a power-law or aud found no essential difference from the case of an NEI model.," As for the thermal components, we also tried the fitting with a plane shock model (XSPEC model $vpshock$ ) plus either a power-law or and found no essential difference from the case of an NEI model." + Although these simple models globally follow the data very. well. all are rejected in the statistical point of view. leaving wavy residuals near the line structure as shown in Figure (lower panel).," Although these simple models globally follow the data very well, all are rejected in the statistical point of view, leaving wavy residuals near the line structure as shown in Figure \ref{spectrum} (lower panel)." + This may be caused by improper response function in enerev scale and/or, This may be caused by improper response function in energy scale and/or +Iuminosity and. displacement from the nearest star cluster would arise if the source lifetime varies inversely with luminosity.,luminosity and displacement from the nearest star cluster would arise if the source lifetime varies inversely with luminosity. + An upper bound on the source lifetime can be obtained from the time required to fully accrete the stellar companion., An upper bound on the source lifetime can be obtained from the time required to fully accrete the stellar companion. + For a companion mass Af and an ellicieney. for the conversion of mass lost by the companion to Luminosity of yp. the source lifetime must be 7gAIC/L where L is the average luminosity ancl e is the speed. of light.," For a companion mass $M$ and an efficiency for the conversion of mass lost by the companion to luminosity of $\eta$, the source lifetime must be $T \le \eta M c^2 / L$ where $L$ is the average luminosity and $c$ is the speed of light." + For sources traveling with a speed ο perpendieular to the line of sight. the displacement from the point of origin will then be dxegALc£L.," For sources traveling with a speed $v$ perpendicular to the line of sight, the displacement from the point of origin will then be $d \le v \eta M c^2 / L$." + Uo the companion mass is independent of the compact object mass. then this would reproduce the required dependence of source lifetime on Luminosity.," If the companion mass is independent of the compact object mass, then this would reproduce the required dependence of source lifetime on luminosity." + Given a tvpical runaway velocity οl0kms1 Wwe must have yll~O.2AL. to match the line plotted in Fig.," Given a typical runaway velocity $v \sim 10 \, \rm km \, s^{-1}$, we must have $\eta M \sim 0.2 M_{\odot}$ to match the line plotted in Fig." + 3 which bounds the region where X-ray. sources are found., \ref{lumd} which bounds the region where X-ray sources are found. + If Roche lobe overllow is occurring. then aecretion. onto the compact object may be ellicient. with little mass loss giving nO.1.," If Roche lobe overflow is occurring, then accretion onto the compact object may be efficient with little mass loss giving $\eta \sim 0.1$." + In this case. the companion mass would be Al~2M...," In this case, the companion mass would be $M \sim 2 M_{\odot}$." + Such intermediate mass companions could be captured via dynamical interactions in the cluster., Such intermediate mass companions could be captured via dynamical interactions in the cluster. + However. the capture must be directly into a Roche-lobe filling orbit or the binary must hardened into a Roche-lobe filling orbit via successive interactions in order to begin accretion promptly. since the evolutionary time scale of the companion is long.," However, the capture must be directly into a Roche-lobe filling orbit or the binary must hardened into a Roche-lobe filling orbit via successive interactions in order to begin accretion promptly, since the evolutionary time scale of the companion is long." + A better understanding of the IME anc the dynamical interactions within the clusters is needed to determine if this scenario is viable., A better understanding of the IMF and the dynamical interactions within the clusters is needed to determine if this scenario is viable. + Even with accretion via Roche lobe overllow. the efficiency. 7 may be less than 0.1 since outLlows ave often observed in X-ray binaries.," Even with accretion via Roche lobe overflow, the efficiency $\eta$ may be less than $0.1$ since outflows are often observed in X-ray binaries." + Super-IEXdcdington mass transfer would also produce 5«0.1., Super-Eddington mass transfer would also produce $\eta < 0.1$. + oy«0.1. then a higher companion mass niav be compatible with the data shown in Lig. 3..," If $\eta < 0.1$, then a higher companion mass may be compatible with the data shown in Fig. \ref{lumd}." + If the N-rav. sources have high-mass companions and maximum speeds near 50kms then the absence of high luminosity source at large clisplacements implies a limit on the X-ray emitting lifetime of the sources.," If the X-ray sources have high-mass companions and maximum speeds near $\sim 50 \rm \, km \, s^{-1}$, then the absence of high luminosity source at large displacements implies a limit on the X-ray emitting lifetime of the sources." + We find. no sources at luminosities above 107ergs.| at displacements larger than 200 pe.," We find no sources at luminosities above $10^{38} \rm \, erg \, s^{-1}$ at displacements larger than 200 pc." + At 50knis+. this would imply that the lifetimes of these Luminous sources must be less than 4 Myr. corresponding to very massive stars.," At $50 \rm \, km \, s^{-1}$, this would imply that the lifetimes of these luminous sources must be less than 4 Myr, corresponding to very massive stars." + An alternative is that the Iuminosity of the sources decreases with age., An alternative is that the luminosity of the sources decreases with age. + Ελπίς would require an evolutionary. path for binaries which produces an accretion rate which decreases with age., This would require an evolutionary path for binaries which produces an accretion rate which decreases with age. + The X-ray binaries may also be beamed (xingetal.2001:IxOÓrdingetal.2002:Ixaaret 2003).," The X-ray binaries may also be beamed \cite{king01,kording02,kaaret03}." +. I£ the X- sources are high-mass svstems with high velocities. then the observed: correlation would imply that beaming only occurs when the binaries are quite voung.," If the X-ray sources are high-mass systems with high velocities, then the observed correlation would imply that beaming only occurs when the binaries are quite young." + Wine et ((2001) suggest that the ULXs in starburst galaxies are high mass X-ray binaries with beamed X-ray emission in a phase of hermal-timescale mass. transfer., King et (2001) suggest that the ULXs in starburst galaxies are high mass X-ray binaries with beamed X-ray emission in a phase of thermal-timescale mass transfer. + Phe delav between the formation of the black hole (and. presumably. the start of he binarv's motion away from it point of origin) and the onset of the thermal-timescale mass transfer phase depends on the stellar evolution of the companion.," The delay between the formation of the black hole (and, presumably, the start of the binary's motion away from it point of origin) and the onset of the thermal-timescale mass transfer phase depends on the stellar evolution of the companion." + The delay. could oe c 20 Myr for a DAZ. companion. which would. imply hat very bright X-rav sources should be visible out. to lkpe.," The delay could be $\sim$ 20 Myr for a $9 M_{\odot}$ companion, which would imply that very bright X-ray sources should be visible out to 1 kpc." + The data appear inconsistent with this. unless highly »xuned X-ray emission occurs in the thermal-timescale mass ransfer phase only for very massive companions.," The data appear inconsistent with this, unless highly beamed X-ray emission occurs in the thermal-timescale mass transfer phase only for very massive companions." + Large displacements of high Dux sources. inconsistent with the data. also appear allowed in the relativistic beamecl model of Ixording ct ((2002).," Large displacements of high flux sources, inconsistent with the data, also appear allowed in the relativistic beamed model of Kording et (2002)." + We thank the Aspen Center for Physics for its hospitality during the workshop where this work was begun., We thank the Aspen Center for Physics for its hospitality during the workshop where this work was begun. + Ply acknowledges partial support from NASA erant. NAC5-7405 and Chandra grant. O2-3102X. JSC thanks the University of Wisconsin-Macdison for support of this research., PK acknowledges partial support from NASA grant NAG5-7405 and Chandra grant GO2-3102X. JSG thanks the University of Wisconsin-Madison for support of this research. +" (1253-055: 2=0.536). 0. }=v/e3οἱ at 130137 and 1309.28À."," In total, the blend consists of three lines decaying to the ground-state fine-structure levels of ] at 1302.17, 1304.86, 1306.03 and two lines of $ +[(3s\,3p^2) \, ^{2} {\rm S_{1/2}} \rightarrow (3s^2\,3p)\, ^{2} {\rm P}_{1/2,3/2}]$ at 1304.37 and 1309.28." +". DallAglio ((2008a) micasured the redshifts of the strongest enission lues to be ta,=2.877£0.003 ενα) and 2.861+0.008. 1V9).", Dall'Aglio (2008a) measured the redshifts of the strongest emission lines to be $z_{\rm em} = 2.877 \pm 0.003$ ) and $2.861 \pm 0.003$ ). +" However. there are well-known shifts iu enüssiou-liue redshifts between the QSO systemic redshift and the ""higher ionization ues” (e.g.Lwa.. A071. ALD19)."," However, there are well-known shifts in emission-line redshifts between the QSO systemic redshift and the “higher ionization lines"" (e.g., $\lambda$ 977, $\lambda$ 1549)." + For mstance. 1000-3000 ooffsets are conmuouly seen in the eniüssion-liue spectra of AGN (Esper 11989: Corbin 1990).," For instance, 1000-3000 offsets are commonly seen in the emission-line spectra of AGN (Espey 1989; Corbin 1990)." +" These issues provide systematic ""nucertainties iu deteriuning 44. either from the weak. blend or from higher ionization lines."," These issues provide systematic uncertainties in determining $z_{\rm em}$, either from the weak blend or from higher ionization lines." + Thus. the precise QSO systemic redshift remains imncertain.," Thus, the precise QSO systemic redshift remains uncertain." + Because of these problems with cuussionu-line redslifts. we propose to use IGM absorption to determine the QSO redshift.," Because of these problems with emission-line redshifts, we propose to use IGM absorption to determine the QSO redshift." + We suggest that the QSO svsteimic redshift is toso=2.901c0.002 (Figure 1) based ou the exteut of the ad aabsorptiou., We suggest that the QSO systemic redshift is $z_{\rm QSO} = 2.904 \pm 0.002$ (Figure 4) based on the extent of the and absorption. + This redshift is chosen as the midpoint between the centroid of the stroug absorption feature seen iu both aand and the red edge of aabsorptiou., This redshift is chosen as the midpoint between the centroid of the strong absorption feature seen in both and and the red edge of absorption. + Figure Ll provides an overview of the COS/CGI30M spectrum from 1100.1190A.. together with a blowup of the region from L173 tto the eedee near 1186A.," Figure 4 provides an overview of the COS/G130M spectrum from 1100–1190, together with a blowup of the region from 1173 to the edge near 1186." +" We see absorption extending well bevond the waveleneths (1180.21180.5. À)) corresponding to the QSO ciuission-line redshifts. 2= and 2.886τε0.003. sugeested by Reimers (1997) and Dall’Aglio ((2008b) respectively,"," We see absorption extending well beyond the wavelengths (1180.2–1180.5 ) corresponding to the QSO emission-line redshifts, $z = 2.885 \pm 0.005$ and $2.886 \pm 0.003$, suggested by Reimers (1997) and Dall'Aglio (2008b) respectively." + The absorption between 1151Lls6 ssugeests that the QSO actually lies at zz2.90., The absorption between 1181--1186 suggests that the QSO actually lies at $z \approx 2.90$. + Tuterestingly. we see no ionization effects frou proximity to the QSO. a topic discussed further in Section 3.6.," Interestingly, we see no ionization effects from proximity to the QSO, a topic discussed further in Section 3.6." + The aabsorption profile shows an absorption feature at 1186.5À.. with optical depth 7z0.7. equivalent width 9511À.. and FWIIM = 0.11A.," The absorption profile shows an absorption feature at 1186.5, with optical depth $\tau \approx 0.7$, equivalent width 95, and FWHM = 0.14." +. This may be proximate aabsorptiou or absorption from an intervene svsteni., This may be proximate absorption or absorption from an intervening system. + The spectruuii near 1186.5 sshows a clear flux muita aud recovery toward shorter waveleugths— typical of weak absorption lines., The spectrum near 1186.5 shows a clear flux minimum and recovery toward shorter wavelengths typical of weak absorption lines. +" A distinct aabsorbeLael would appear as a broader ""shelf? ou the side of the τω>1 trough."," A distinct absorber would appear as a broader “shelf"" on the side of the $\tau_{\rm HeII} > 1$ trough." + We do not make a specific line identification for the feature., We do not make a specific line identification for the feature. + It does not ποσα to he low- nor any of the usual IGAL metalline svstems (Danfortl Shull 2008)., It does not seem to be low-redshift nor any of the usual IGM metal-line systems (Danforth Shull 2008). + οσους ((1997) attribute the aac aabsorption at 22.885 to a nmlticoniponent svsteni of associated absorbers between Lisl1186 (Qvedshifts 2=2.891 to =2.901)., Reimers (1997) attribute the and absorption at $z > 2.885$ to a multicomponent system of associated absorbers between 1181–1186 (redshifts $z = 2.891$ to $z = 2.904$ ). + Many of these systems have stroug metal lines ofIV...δν aud VL. aud a few have anomalous ratios of aand aabsorptiou.," Many of these systems have strong metal lines of, and , and a few have anomalous ratios of and absorption." + These observations sugeest that the eas might be exposed to high fluxes of ionizing radiation and affected by: uueleosvuthetic anomalies mio nuclear outflows., These observations suggest that the gas might be exposed to high fluxes of ionizing radiation and affected by nucleosynthetic anomalies in nuclear outflows. +" Sinctte ((2002) aud Fechuer ((2001) both noted evidence for ""luue-dockiug iu aand ((althoueh the latter eroup fouud no statistical evideuce conrpared to simmlated line Lists)."," Smette (2002) and Fechner (2004) both noted evidence for “line-locking"" in and (although the latter group found no statistical evidence compared to simulated line lists)." + Line-locking occurs when two absorbers are separated by a velocity equal to the separation of the two doublet lines. resulting from radiation pressure on material plivsically uear tle QSO (Wevinaun. Carswell. Smith 1981).," Line-locking occurs when two absorbers are separated by a velocity equal to the separation of the two doublet lines, resulting from radiation pressure on material physically near the QSO (Weymann, Carswell, Smith 1981)." + These radiative forces will drive gas outward from the QSO. whereas the redshifted associated absorbers would have infall velocity ~1500 rrelative to the QSO if toso= 2.585.," These radiative forces will drive gas outward from the QSO, whereas the redshifted associated absorbers would have infall velocity $\sim1500$ relative to the QSO if $z_{\rm QSO} = 2.885$ ." + Thus. our proposed redshift. tos0=2.901. provides a more plausible explanation.," Thus, our proposed redshift, $z_{\rm QSO} = 2.904$, provides a more plausible explanation." + The uncertainty in QSO systemic redshift illustrates the need for mfrared observations., The uncertainty in QSO systemic redshift illustrates the need for infrared observations. + Unfortunately. the redshitted hydrogen Daliuer lines fall in difficult spectral bands (Ila at 2.563 jun. IL) at 1.5908 jnu). as does the TI]A3727 doublet at 1.155 aud 1.156 jan. The forbidden|O lines of [O IH] A5007.1959 occurat more promising wavelengths. 1.955 jan and 1.936 gan. Thus. infrared spectroscopic searches for the |O III] forbidden," Unfortunately, the redshifted hydrogen Balmer lines fall in difficult spectral bands $\alpha$ at 2.563 $\mu$ m, $\beta$ at 1.898 $\mu$ m), as does the [O II]$\lambda 3727$ doublet at 1.455 and 1.456 $\mu$ m. The forbidden lines of [O III] $\lambda 5007, 4959$ occurat more promising wavelengths, 1.955 $\mu$ m and 1.936 $\mu$ m. Thus, infrared spectroscopic searches for the [O III] forbidden" +"The auxiliary field7Tj,. prepotentialsolu","Our ansatz is where and $\beta_U,\beta_f,\beta_1,\beta_2,\beta_3,\beta_0,\beta_T$ are constants satisfying $\mu\abs{\beta}\ll1$ , and $e^{-K}$ also contains $\beta$ 's. Note that besides the explicit $\mu$ -corrections above, some of the fields will also have implicit $\mu$ dependence via $e^{-K}$ and $A(z)$." +tion.introduce linear proves p-correctionstothe fields.," In addition we assume The equation of motion for the $SU(2)$ connection is always satisfied by the vanishing $SU(2)$ connection, for a bosonic background and withour choice of $V$ -gauge (also assuming no" +from the Sedov solution is four times smaller than that analyzed by Pfrommer et al. (,from the Sedov solution is four times smaller than that analyzed by Pfrommer et al. ( +2005) at a frequency of 144 GHz (see Figs.,2005) at a frequency of 144 GHz (see Figs. + of this paper and 1 of Pfrommer et al., \ref{sz217} of this paper and 1 of Pfrommer et al. + 2005)., 2005). + The signal-to- ratio for a detection of an AGN cocoon is proportional the square root of the exposure time (e.g. Pfrommer et al., The signal-to-noise ratio for a detection of an AGN cocoon is proportional the square root of the exposure time (e.g. Pfrommer et al. +" 2005; for a review, see a document P)) and, therefore, to detect the signal at a frequency of 217 GHz the ALMA exposure time of 80 (i.e. 5x4?) hours is required."," 2005; for a review, see a document ) and, therefore, to detect the signal at a frequency of 217 GHz the ALMA exposure time of 80 (i.e. $5\times4^2$ ) hours is required." + The ALMA Design Reference Science Plan (Version 2.2) contains two in SZ studies (1.4.1 and 1.4.2) with integration time for each program of 400 hours and new projects concern the SZ effect on AGN cocoons should be promising., The ALMA Design Reference Science Plan (Version 2.2) contains two in SZ studies (1.4.1 and 1.4.2) with integration time for each program of 400 hours and new projects concern the SZ effect on AGN cocoons should be promising. +" The Chandra X-ray satellite with a high spatial resolution of 0.5"" has studied a non-relativistic electron component in the ICM, while ALMA high-resolution observations of AGN cocoons in galaxy clusters should permit us to reveal a new (mildly relativistic) electron component."," The Chandra X-ray satellite with a high spatial resolution of $''$ has studied a non-relativistic electron component in the ICM, while ALMA high-resolution observations of AGN cocoons in galaxy clusters should permit us to reveal a new (mildly relativistic) electron component." + Recent hydrodynamic simulations do not take into account relativistic corrections to the SZ effect to derive SZ intensity maps (e.g. Sijacki et al., Recent hydrodynamic simulations do not take into account relativistic corrections to the SZ effect to derive SZ intensity maps (e.g. Sijacki et al. + 2008)., 2008). + Such relativistic corrections are necessary for calculating the SZ effect on AGN cocoons produced by electrons with high temperatures., Such relativistic corrections are necessary for calculating the SZ effect on AGN cocoons produced by electrons with high temperatures. +" So far, studies of the SZ effect in the relativistic formalism were done only for an analytic toy cocoon model (see Pfrommer et al."," So far, studies of the SZ effect in the relativistic formalism were done only for an analytic toy cocoon model (see Pfrommer et al." + 2005)., 2005). +" In this section, we derive the SZ intensity map by using hydrodynamical simulations and the Wright relativistic formalism."," In this section, we derive the SZ intensity map by using hydrodynamical simulations and the Wright relativistic formalism." + In previous papers (Antonuccio-Delogu Silk 2008; Tortora et al., In previous papers (Antonuccio-Delogu Silk 2008; Tortora et al. +" 2009), we used an Adaptive Mesh Refinement (AMR) code to follow the evolution of the cocoon produced by the jet propagating in the interstellar (ISM)/intergalactic medium (IGM)."," 2009), we used an Adaptive Mesh Refinement (AMR) code to follow the evolution of the cocoon produced by the jet propagating in the interstellar (ISM)/intergalactic medium (IGM)." +" To perform the simulation, we used FLASH v.2.5 (Fryxell et al."," To perform the simulation, we used FLASH v.2.5 (Fryxell et al." +" 2000), a parallel, AMR code, which implements a second order, shock-capturing, Piecewise-Parabolic-Method (PPM) solver."," 2000), a parallel, AMR code, which implements a second order, shock-capturing, Piecewise-Parabolic-Method (PPM) solver." +" The jet is modelled as a one-component fluid, characterized by a density which is a small fraction of the initial density of the ISM."," The jet is modelled as a one-component fluid, characterized by a density which is a small fraction of the initial density of the ISM." +" In this simulation, the power injected by the jet is 10*° erg/s. We model the environment, where the jet propagates, as a two-phase ISM, comprising a hot, diffuse, low-density component having a temperature and a central density at 107 K and 1 cm’, respectively, and a cold, clumped system of high-density clouds in pressure equilibrium with the diffuse component."," In this simulation, the power injected by the jet is $10^{46}$ erg/s. We model the environment, where the jet propagates, as a two-phase ISM, comprising a hot, diffuse, low-density component having a temperature and a central density at $10^7$ K and $1$ $^3$, respectively, and a cold, clumped system of high-density clouds in pressure equilibrium with the diffuse component." + Such values of temperature and pressure are typical of the ISM in the central parts of an elliptical at high redshift., Such values of temperature and pressure are typical of the ISM in the central parts of an elliptical at high redshift. +" We assume that the diffuse gas is embedded within a dark matter (DM) halo, the latter being described by a White (NFW) density profile."," We assume that the diffuse gas is embedded within a dark matter (DM) halo, the latter being described by a Navarro-Frenk-White (NFW) density profile." +" We choose a simulation box having a size [πο=40h. kpc, where h is defined such that the Hubble constant Ho is 100 h km s! Mpc!."," We choose a simulation box having a size $L_{\mathrm{box}} = 40 +h_{-1}$ kpc, where $h$ is defined such that the Hubble constant $H_0$ is 100 $h$ km $^{-1}$ $^{-1}$." + The spatial resolution attained is a function of the maximum refinement level and of the structure of the code., The spatial resolution attained is a function of the maximum refinement level and of the structure of the code. +" For a block-structured AMR code like FLASH, where each block is composed by n,xny cells, the maximum resolution along each direction is given by Loox/(n,2)), where | is the maximum refinement level."," For a block-structured AMR code like FLASH, where each block is composed by $n_x \times n_y$ cells, the maximum resolution along each direction is given by $L_{\mathrm{box}}/(n_x +2^l)$, where $l$ is the maximum refinement level." +" In this simulation, n, = n, = 8 and 1 = 6, thus the minimum resolved scale is 78.125 pc."," In this simulation, $n_x$ = $n_y$ = 8 and $l$ = 6, thus the minimum resolved scale is 78.125 pc." +" Note that we are performing a 2D simulation, but we do not impose any special symmetry."," Note that we are performing a 2D simulation, but we do not impose any special symmetry." + Numerical simulations by Antonuccio-Delogu Silk (2008) and Tortora et al. (, Numerical simulations by Antonuccio-Delogu Silk (2008) and Tortora et al. ( +"2009) show that the temperatures can reach very high values within the cocoons (T «10?—10!! K), if the ISM is dominated by a population of cold, star forming clouds embedded into and in approximate pressure equilibrium with a hot, diffuse phase.","2009) show that the temperatures can reach very high values within the cocoons (T $\approx 10^9-10^{11}$ K), if the ISM is dominated by a population of cold, star forming clouds embedded into and in approximate pressure equilibrium with a hot, diffuse phase." + The pressure within the cocoon can reach high values because the temperatures are on average very high., The pressure within the cocoon can reach high values because the temperatures are on average very high. + The SZ effect should be significant since it determined by the pressure of the electron population integrated along a line of sight., The SZ effect should be significant since it determined by the pressure of the electron population integrated along a line of sight. + Our simulation temperature and pressure maps are shown in Fig., Our simulation temperature and pressure maps are shown in Fig. +[7] and}., \ref{Tsim} and \ref{Psim}. + Figure[7]| reveals a hot shell with a temperature of koT.>2 MeV around the internal region of the simulation AGN cocoon., Figure \ref{Tsim} reveals a hot shell with a temperature of $k_\mathrm{b} T_\mathrm{e}>2$ MeV around the internal region of the simulation AGN cocoon. + The temperatures of the internal region are smaller than that of, The temperatures of the internal region are smaller than that of +" and the pulsar term becomes à source with the same amplitude compared to the ""signal, term.", and the pulsar term becomes a source with the same amplitude compared to the `signal' term. + The SNR of single pulsar data is then always less than one., The SNR of single pulsar data is then always less than one. + On the other hand. if we can use the pulsar term as a ‘signal’> term. we can sum the signal from multiple pulsar coherently and the SNR. of single pulsar data can be arbitrarily large. depending on the GAV amplitude.," On the other hand, if we can use the pulsar term as a `signal' term, we can sum the signal from multiple pulsar coherently and the SNR of single pulsar data can be arbitrarily large, depending on the GW amplitude." + The second. important benefit of utilizing the pulsar term is that it allows us to get a precise measurement for the GW source location., The second important benefit of utilizing the pulsar term is that it allows us to get a precise measurement for the GW source location. + The interference between the pulsar term and the Earth term introduces spiky lobes in the single-pulsar response pattern (sce Figure 2))., The interference between the pulsar term and the Earth term introduces spiky lobes in the single-pulsar response pattern (see Figure \ref{fig:beampat}) ). + In. Figure 23.. where the small scale ring-tvpe structures come from individual response patterns and the larec scale quacruple structure is introduced bv the coherence of the Earth term.," In Figure \ref{fig:skymap}, where the small scale ring-type structures come from individual response patterns and the large scale quadruple structure is introduced by the coherence of the Earth term." + This large scale euuacdruple structure is. in fact. the response pattern. when one only uses the Earth termi as a signal.," This large scale quadruple structure is, in fact, the response pattern, when one only uses the Earth term as a signal." +is perpendicular to the associated outflow. (his velocity structure can be interpreted. as an expanding motion in the flattened disk (e.g. Nitanmwa el al.,"is perpendicular to the associated outflow, this velocity structure can be interpreted as an expanding motion in the flattened disk (e.g. Kitamura et al." + 1996)., 1996). + The other velocity structure in (he X-shape. wilh (wo components at the velocity range of Viag 9.9 to 10.4lans + (blueshifted: Gii)-l in Figure ??b)andV ju5—10.5 to 11.0 kins ! (yvedshifted: (1)-2 in Figure ??b). hasthesamevelocilysenseaslhatoftheassocialedoul flow.," The other velocity structure in the X-shape, with two components at the velocity range of $V_{\rm{LSR}}$ =9.9 to 10.4km $^{-1}$ (blueshifted; (iii)-1 in Figure \ref{H13_PV}$ $b$ ) and $V_{\rm{LSR}}$ =10.5 to 11.0 km $^{-1}$ (redshifted; (iii)-2 in Figure \ref{H13_PV}$ $b$ ), has the same velocity sense as that of the associated outflow." +E heseticocomponentsarelocatedo," These two components are located outside the central condensed gas, suggesting that these components are not associated with the central protostar." +ul. updensegasbylheassociatedoul flowperpendicutarlolhedish—likeenvelope.," Thus, it is natural to interpret that these gas components trace swept-up dense gas by the associated outflow perpendicular to the disk-like envelope." + Asimilarecamplehasalsobeenr densi, A similar example has also been reported from observations in L1228 \citep{taf94}. . +lytraceraf heC sls (245 - 194) emission. whieh has (he same velocity sense as that of the associated CO bipolar outflow. and (μον interpreted this shift as an interaction between (he dense core and (he high-velocity outflow.," They detected a velocity shift in the high-density tracer of the $_3$ $_2$ $_{12}$ - $_{01}$ ) emission, which has the same velocity sense as that of the associated CO bipolar outflow, and they interpreted this shift as an interaction between the dense core and the high-velocity outflow." +" The eastern component. indicated by (1) in Figure T. ds located. at 25"" east of the continuum peak position and is a part of the [an-shaped structure seen in the combined total intensity map."," The eastern component, indicated by (i) in Figure 7, is located at $''$ east of the continuum peak position and is a part of the fan-shaped structure seen in the combined total intensity map." + We interpret that. this component is remnant dense gas interacting with the associated molecular outflow., We interpret that this component is remnant dense gas interacting with the associated molecular outflow. + In stummary. the P-V diagram can be interpreted as a mixture of the dispersing gas along the disk-lhke envelope and the interacting dense gas with the associated: outflow perpendicularly to the disk-like envelope.," In summary, the P-V diagram can be interpreted as a mixture of the dispersing gas along the disk-like envelope and the interacting dense gas with the associated outflow perpendicularly to the disk-like envelope." + Presumably due to the insufficient spatial resolution. an infall eas motion in the protostellar source of MMS 7 has not been clearly identified in our observations.," Presumably due to the insufficient spatial resolution, an infall gas motion in the protostellar source of MMS 7 has not been clearly identified in our observations." + The velocity structure of the clisk-like envelope shows a dispersing gas motion as discussed in the last section., The velocity structure of the disk-like envelope shows a dispersing gas motion as discussed in the last section. +" The virial mass of the clisk-like envelope is estimated to be 23-30 M. assuming D, = 0.15 pe. Tay = 26 - 50 I. and Ac = LO km J|. which is larger than the LTE mass of the disk-like envelope (5.1 - 9.1 M.)."," The virial mass of the disk-like envelope is estimated to be 23-30 $_{\odot}$ assuming $D_{\rm{env}}$ = 0.15 pc, $T_{\rm{env}}$ = 26 - 50 K, and $\Delta v$ = 1.0 km $^{-1}$, which is larger than the LTE mass of the disk-like envelope (5.1 - 9.1 $_{\odot}$ )." + This result suggests that the clisk- envelope is gravitationally unbound. supporting our interpretation of the dispersing gas motion.," This result suggests that the disk-like envelope is gravitationally unbound, supporting our interpretation of the dispersing gas motion." + We can estimate physical parameters of the dispersing motion. such as an expanding velocity (V5.5). momentum (7745). expanding energy (£5). and mechanical power (Lexy). as Vup(mVua)=1.2kinsL Day=ALVinayAP.kinsb. Boxy=MVZ/2M.kn?s 7. and Lay=MVE/2RL..," We can estimate physical parameters of the dispersing motion, such as an expanding velocity $V_{\rm{exp}}$ ), momentum $P_{\rm{exp}}$ ), expanding energy $E_{\rm{exp}}$ ), and mechanical power $L_{\rm{exp}}$ ), as $V_{\rm{exp}}({\equiv}~V_{\rm{max}})~=~1.2~\rm{km~s^{-1}}$, $P_{\rm{exp}}=MV_{\rm{max}}~\rm{M_{\odot} km~s^{-1}}$, $E_{\rm{exp}}=MV_{\rm{max}}^2/2~\rm{M_{\odot} km ^{2}~s^{-2}}$ , and $L_{\rm{exp}}=MV_{\rm{max}}^3/2R~\rm{L_{\odot}}$." + We adopt the disk inclination angle from the plane of the sky to be iον 80°. which is consistent with the morphology of the CO outflow and the reflection nebula.," We adopt the disk inclination angle from the plane of the sky to be $i~\sim$ $^{\circ}$, which is consistent with the morphology of the CO outflow and the reflection nebula." + Table 4. lists the estimated physical parameters of the dispersing motion., Table \ref{COMP_T} lists the estimated physical parameters of the dispersing motion. + The dispersing envelope has also been observed in the ο emission around a low-mass YSO of DG Tau (xilamuraetal...1996)., The dispersing envelope has also been observed in the $^{13}$ CO(1-0) emission around a low-mass YSO of DG Tau \citep{kit96}. +. We compare physical parameters specified in (the dispersing process of MMS 7 (0 those of DG Tau in Table 4.., We compare physical parameters specified in the dispersing process of MMS 7 to those of DG Tau in Table \ref{COMP_T}. . + These parameters of MAIS 7 are twoorders of magnitude larger than those in DG Tan., These parameters of MMS 7 are twoorders of magnitude larger than those in DG Tau. + The interlerometric observations of DG Lau mav sulfer from the problem of themissing flix. which prevents us from making a direct," The interferometric observations of DG tau may suffer from the problem of themissing flux, which prevents us from making a direct" +the number of iterations. it is merely necessary to test for convergence by calculating the dillerence to the corrected image after each step.,"the number of iterations, it is merely necessary to test for convergence by calculating the difference to the corrected image after each step." + Fvpical science images from. carly 2010 change by only one electron in a handful of pixels after three iterations. and by less than an electron in every. pixel after four.," Typical science images from early 2010 change by only one electron in a handful of pixels after three iterations, and by less than an electron in every pixel after four." + Since cach iteration has a [arge price in run time. we shall henceforth stop at the third iteration.," Since each iteration has a large price in run time, we shall henceforth stop at the third iteration." + We use our updated CCD readout model to correct science imagine throughout the lifetime of ACS. following the same procedure as Bristow(2003a)..," We use our updated CCD readout model to correct science imaging throughout the lifetime of ACS, following the same procedure as \citet{bristow03im}." + Phe points with dotted error bars in figure 2. show the effective density of charge traps after correction. which are a factor of 20 lower than in the paw data and Consistent with image quality in the first. six months of operations.," The points with dotted error bars in figure \ref{fig:postsm4} show the effective density of charge traps after correction, which are a factor of 20 lower than in the raw data and consistent with image quality in the first six months of operations." + For the sake of clarity. equivalent post-correction measurements are not shown for earlier epochs. but these recover about the same factor of 10-15 correction seen in Massey.efaf(2010).," For the sake of clarity, equivalent post-correction measurements are not shown for earlier epochs, but these recover about the same factor of 10-15 correction seen in \citet{m10}." +. Thus. ironically. as the ςΤΙ has got worse. the trailing has become easier to measure and the correction has become more accurate!," Thus, ironically, as the CTI has got worse, the trailing has become easier to measure and the correction has become more accurate!" + Figure 3. shows a region of a typical exposure. which was intentionally»0/ used when measuring parameters of the reacout model.," Figure \ref{fig:beforeafter} shows a region of a typical exposure, which was intentionally used when measuring parameters of the readout model." + The charge trailing that is now reaclily apparent in visual inspection of recent ACS images is successfully removed by our correction scheme., The charge trailing that is now readily apparent in visual inspection of recent ACS images is successfully removed by our correction scheme. + We have developed. a physically-motivated model. of. the readout and Charge ‘Transfer Inellicieney (CFI) in the ACS/WEC detectors throughout their lifetime., We have developed a physically-motivated model of the readout and Charge Transfer Inefficiency (CTI) in the ACS/WFC detectors throughout their lifetime. + We find tha here are approximately 1.3 charge traps per pixel in mid-2010. split between three different: species.," We find that there are approximately 1.3 charge traps per pixel in mid-2010, split between three different species." + The extende rails produced. by these traps can be accurately mocellec as a sum of three decaving exponentials., The extended trails produced by these traps can be accurately modelled as a sum of three decaying exponentials. + We also used. our model to correct images. reducing the amount of trailing by a actor of 20. to a level seen in the first six months of orbita operations.," We also used our model to correct images, reducing the amount of trailing by a factor of $\sim$ 20, to a level seen in the first six months of orbital operations." + As with Chiabereeefa£.(2009).. we still find no evidence for significant serial CIEL (trailing perpendicular to he main trails. created by charge traps in the serial readou register). and therefore ignore this cllect.," As with \citet{chiaberge}, we still find no evidence for significant serial CTI (trailing perpendicular to the main trails, created by charge traps in the serial readout register), and therefore ignore this effect." + When building our mocel. we adopted the best availablemeasurements from science imaging (which we performed) and dark exposures (fromAnderson&Beclin2010).," When building our model, we adopted the best availablemeasurements from science imaging (which we performed) and dark exposures \citep[from][]{a10}." +. The dark exposures were particularly useful. to constrain the extended shape of trails out to 7100 pixels and thus provide better correction of object photometry (Ithodes.efa. 2010)., The dark exposures were particularly useful to constrain the extended shape of trails out to $\sim$ 100 pixels and thus provide better correction of object photometry \citep{rhodes10}. . +. Where measurements disagreed. the data's support," Where measurements disagreed, the data's support" +(factor?,factor? +" ""ποσο issues will all be visited ancl cüscussed in a forthcoming paper.", These issues will all be visited and discussed in a forthcoming paper. + We have performed raciative transfer. calculations to investigate the nature of τςὉ CJ=1-0) emission [rom simulations of molecular clouds (MCS)., We have performed radiative transfer calculations to investigate the nature of $^{12}$ CO (J=1-0) emission from simulations of molecular clouds (MCs). +. The. ος are modeled through hverodynamic simulations of a turbulent. magnetized. non self-eravitating gaseous medium. along with a treatment of chemistry to track the formation of aand CO (Paper E and Paper 1).," The MCs are modeled through hydrodynamic simulations of a turbulent, magnetized, non self-gravitating gaseous medium, along with a treatment of chemistry to track the formation of and CO (Paper I and Paper II)." + As part of the radiative transfer calculations. we use the Sobolev (Ας) method to solve for the CO level populations.," As part of the radiative transfer calculations, we use the Sobolev (LVG) method to solve for the CO level populations." + We analyze the probability distribution functions. and [factor properties. using the velocity integrated. CO intensitiesΕλ. alone with the column density of CO andIle. aadNy... respectively.," We analyze the probability distribution functions, and factor properties, using the velocity integrated CO intensities, along with the column density of CO and, and, respectively." + Our main findings are: 1) In all models. iincreases with increasing total hydrogen column density (for extinction ον).," Our main findings are: 1) In all models, increases with increasing total hydrogen column density (or extinction )." + However. for the Milkv Way and high density mocdels (9300 and n1000). which have the highest CO abundances. we find a threshold in z65 aat high extinction due to saturation of the CO line.," However, for the Milky Way and high density models (n300 and n1000), which have the highest CO abundances, we find a threshold in $\approx$ 65 at high extinction due to saturation of the CO line." + 2) All models have log-normal aand cdeistributions., 2) All models have log-normal and distributions. + In general. however. the aand PPDES are not log-normal.," In general, however, the and PDFs are not log-normal." + Further. since CO is optically hick. the >PDEs do not have similar shapes to the corresponding >PDEs.," Further, since CO is optically thick, the PDFs do not have similar shapes to the corresponding PDFs." + 3) In some mocdels for which the CO line is saturated. he peak in iis olfset towards higher densities than the peak inW.. hough the two PDEs seem to be correlated at low densities.," 3) In some models for which the CO line is saturated, the peak in is offset towards higher densities than the peak in, though the two PDFs seem to be correlated at low densities." + However. such a 7piled-up PPDE does not necessarily arise in clouds with saturated CO emission. especially those with a limited. range in densities (as for the high density model n1000. Figs. 1--2)).," However, such a “piled-up” PDF does not necessarily arise in clouds with saturated CO emission, especially those with a limited range in densities (as for the high density model n1000, Figs. \ref{wvsav}- \ref{pdfpans}) )." +" Independent measurements of oorJVg,.. along withΕν are needed. to. unambiguously identify saturation."," Independent measurements of or, along with, are needed to unambiguously identify saturation." + +) The factor is not constant within individual molecular clouds. and in models with low CO fractions. can vary by up to 4 orders of magnitude.," 4) The factor is not constant within individual molecular clouds, and in models with low CO fractions, can vary by up to 4 orders of magnitude." + The low density regions have the üghest factor. in agreement with previous moceling efforts.," The low density regions have the highest factor, in agreement with previous modeling efforts." +" 5) In most simulations. the averaged Tactor is found to be similar to ((—2.I0 s5)). masking. the variation""M of wwith density within clouds."," 5) In most simulations, the averaged factor is found to be similar to $ \sim2 \times 10^{20}$ ), masking the variation of with density within clouds." + In clouds with low CO abundances relative to the Galaxy (such as the model MC in à cwarl galaxv or the LAIC/SAIC). the densest: regions have ~Nou," In clouds with low CO abundances relative to the Galaxy (such as the model MC in a dwarf galaxy or the LMC/SMC), the densest regions have $\sim$." + 6) Emission weighted averaged flactors from all models provide values = with the exception of the very low metallicity MC Z01)., 6) Emission weighted averaged factors from all models provide values $\approx$ with the exception of the very low metallicity MC (n300-Z01). + Similarly. at extinctions 77. the ITactor for abwavs falls in a narrow range ~102U 107s.," Similarly, at extinctions 7, the factor for always falls in a narrow range $\sim 10^{20}$ $10^{21}$." + As discussed. in Paper LL. observations targeting CO bright regions may be unable to detect real variations in the ITactor.," As discussed in Paper II, observations targeting CO bright regions may be unable to detect real variations in the factor." + 7) In the low density gas within the MC mocels. the CO fraction feo is Found to be ~ three orders of magnitude lower than feo measured by ? in diffuse gas.," 7) In the low density gas within the MC models, the CO fraction $f_{CO}$ is found to be $\sim$ three orders of magnitude lower than $f_{CO}$ measured by \citet{Burghetal07} in diffuse gas." + That aand aare comparable between the model and the observations suggest that the cvnamics and chemical evolution of dilfuse eas in our molecular cloud model are rather dillerent. from that in the [arge scale dilluse eas of the Galaxy., That and are comparable between the model and the observations suggest that the dynamics and chemical evolution of diffuse gas in our molecular cloud model are rather different from that in the large scale diffuse gas of the Galaxy. + S) We do not assume that clouds or cores are in virial equilibrium. and the [factor variations we [ind are in general agreement with observational analyses emploving independent measures of total molecular column densities.," 8) We do not assume that clouds or cores are in virial equilibrium, and the factor variations we find are in general agreement with observational analyses employing independent measures of total molecular column densities." + In ai follow up investigation. we will use the spectral information of our models to assess the [factor when virial equilibrium is assumed: we will then be able to address the [factor discrepaney found. in observational works when assuming virtalized CO clouds or when using independent molecular mass estimates.," In a follow up investigation, we will use the spectral information of our models to assess the factor when virial equilibrium is assumed; we will then be able to address the factor discrepancy found in observational works when assuming virialized CO clouds or when using independent molecular mass estimates." + We are grateful. ο Eve Ostriker. Alyssa Goodman. Jaime Pineda. Morelecai Mac. Low. Frank Bigiel. jXnucdrew Harris. and Christoph Federrath for stimulating cliscussions regarding CO emission and the Ifactor.," We are grateful to Eve Ostriker, Alyssa Goodman, Jaime Pineda, Mordecai Mac Low, Frank Bigiel, Andrew Harris, and Christoph Federrath for stimulating discussions regarding CO emission and the factor." + We also thank an anonvmous referee. for useful comments., We also thank an anonymous referee for useful comments. + ltS.. acknowledges financial support from the Landesstiftung DBaden-Würrrtemberg via their program International Collaboration LL (grant. P-LS-SPIL/IS) and from the German Dundesministerium Ltr Bildung. und Forschung via the ASTIVONIZE project STAR EFOIULVE (erant 05A00VILA)., R.S.K. acknowledges financial support from the Landesstiftung Baden-Würrrtemberg via their program International Collaboration II (grant P-LS-SPII/18) and from the German Bundesministerium fürr Bildung und Forschung via the ASTRONET project STAR FORMAT (grant 05A09VHA). +. It.S.Ix.. furthermore acknowledges subsidies from the DPC under grants. no., R.S.K. furthermore acknowledges subsidies from the DFG under grants no. + INL1358/1. IXL1358/4. IXL1358/5. IXL1358/10. and IxLI358/11. as well as from a Frontier grant of Lleiclelbere University sponsored by the German Excellence Initiative.," KL1358/1, KL1358/4, KL1358/5, KL1358/10, and KL1358/11, as well as from a Frontier grant of Heidelberg University sponsored by the German Excellence Initiative." + This work was supported in part bv the U.S. Department of Enerey contract no., This work was supported in part by the U.S. Department of Energy contract no. + DIZ-AC-02-768E00515., DE-AC-02-76SF00515. + RSJ. also thanks the Ίαν Institute for Particle Astrophysics anc Cosmology at Stanford University and the Department of Astronomy ancl Astrophysics at the University of California at Santa Cruz for their warm hospitality curing a sabbatical stay in spring 2010., R.S.K. also thanks the Kavli Institute for Particle Astrophysics and Cosmology at Stanford University and the Department of Astronomy and Astrophysics at the University of California at Santa Cruz for their warm hospitality during a sabbatical stay in spring 2010. +bv the best estimates of observations.,by the best estimates of observations. + The velocity. with higher uncertainty. is randomly altered for each run assuming independent Cassia uncertainties in the measurements. the Lue of sight velocity aud the proper motions are calculated aud then transformed to Cartesian velocities around the Calaxy The inasses of the LMC aud SMC have. larger uucertainties than that of the position. ancl velocity.," The velocity, with higher uncertainty, is randomly altered for each run assuming independent Gaussian uncertainties in the measurements, the line of sight velocity and the proper motions are calculated and then transformed to Cartesian velocities around the Galaxy The masses of the LMC and SMC have larger uncertainties than that of the position and velocity." + obtained a mass for the LAIC of (8.7=L3)ς10°AL. within 8.9 kpe une an analysis of carbon stars., obtained a mass for the LMC of $(8.7\pm{}4.3)\times10^{9}$$_\odot$ within $8.9$ kpc using an analysis of carbon stars. +" This mass is less than half that estimated by who derive a niass of 2«1010 NL, using radial velocities for several of the oldest star clusters in the LAIC that lie well bevouk 6 kpc of its center.", This mass is less than half that estimated by who derive a mass of $2\times10^{10}$ $_\odot$ using radial velocities for several of the oldest star clusters in the LMC that lie well beyond $6$ kpc of its center. +" We match the circular velocities of with that of an Einasto halo at 8.9 kpe to get a virial mass of 2\Lott AL, for a halo virialized today.", We match the circular velocities of with that of an Einasto halo at $8.9$ kpc to get a virial mass of $2\times10^{11}$ $_\odot$ for a halo virialized today. +" Using a simplified model of tidal radius 77.Ia=nuafMcaaeWe fd that the LMC will have a mass of 5.6«1010 NE, within a tida radius of 30 kpc. oulv sheltly more massive than that of(1992)."," Using a simplified model of tidal radius $r_{\rm tid}^3/R_{\rm peri}^3 = m_{\rm tid}/M_{\rm Galaxy}$we find that the LMC will have a mass of $5.6\times10^{10}$ $_\odot$ within a tidal radius of $30$ kpc, only slightly more massive than that of." +. For the SAIC. obtained a lower mass limit of 1.0.10? ML. from observations of carbon stars. while obtained 0.9«LO? from simular observations of planetary nebulae.," For the SMC, obtained a lower mass limit of $1.0\times10^{9}$ $_\odot$ from observations of carbon stars, while obtained $0.9\times10^{9}$ from similar observations of planetary nebulae." + More receutly. determuned a nass of 2.7 ALL.," More recently, determined a mass of $2.7$ $5.1\times10^{9}$ $_\odot$." + We use the velocity dispersion frou to caleulate an Eimasto halo virial nass of L5«1019 NE. with a mass 5.7<10° AL. within a tidal radius of 19 kpe.," We use the velocity dispersion from to calculate an Einasto halo virial mass of $4.5\times10^{10}$ $_\odot$, with a mass $5.7\times10^9$ $_\odot$ within a tidal radius of $19$ kpc." + These masses. somewhat above the caleulated masses at sinaller radii. also approximate the mass of a theorized conimnion halo between the LMC aud SAIC2008).. without the problemi of a large sublialo within sublalo Gu conunon halo models. the LMC halo will have a mass of ~20% of the conunon halo).," These masses, somewhat above the calculated masses at smaller radii, also approximate the mass of a theorized common halo between the LMC and SMC, without the problem of a large subhalo within subhalo (in common halo models, the LMC halo will have a mass of $\sim$$20\%$ of the common halo)." + These lnasses are also simular to those calculated from the stellar 1nass content aud seen muside N-body simulations2011)., These masses are also similar to those calculated from the stellar mass content and seen inside -body simulations. +. For the model. we ignore any mass loss from the clouds (in particular that due to tidal stripping). although mass loss is known to occur as evidenced bv the απο! of matter iu the Maeellauic Stream:2003).," For the model, we ignore any mass loss from the clouds (in particular that due to tidal stripping), although mass loss is known to occur as evidenced by the amount of matter in the Magellanic Stream;." +. Dynamical friction between the LMC. SMC and the Galaxy is also calculated usine the tidal masses of the LAIC aud SAIC.," Dynamical friction between the LMC, SMC and the Galaxy is also calculated using the tidal masses of the LMC and SMC." + The effect of dynamical friction between the LAIC and SAIC is ignored. although LAIC/SAIC binary svstenis are still found in N-body simmlations that include ανασα friction suggesting that the survival time with dynamical friction is sufficiently large for the Maegcllanic svstem to reach the present dav witli these masses.," The effect of dynamical friction between the LMC and SMC is ignored, although LMC/SMC binary systems are still found in -body simulations that include dynamical friction suggesting that the survival time with dynamical friction is sufficiently large for the Magellanic system to reach the present day with these masses." + The LMC/SMC and chwart galaxy orbits occur withiu the potential of the Calaxy and are altered due to interactions between the Magellanic clouds themsclyes and the dwarts., The LMC/SMC and dwarf galaxy orbits occur within the potential of the Galaxy and are altered due to interactions between the Magellanic clouds themselves and the dwarfs. + There has been uch work on Galaxy poteutial models «λος the first basic schemes devised by aud(1976)., There has been much work on Galaxy potential models since the first basic schemes devised by and. +. Iu contrast to other recent approaches that take simple spherical Galactic poteutials. we choose to model the Calaxy usine a multicomiponeut model of the Calactic poteutial developed by(1996).," In contrast to other recent approaches that take simple spherical Galactic potentials, we choose to model the Galaxy using a multicomponent model of the Galactic potential developed by." +.. This model matches known Galactic parameters such as the rotation curve. local disk density and disk scale-leneth to high accuracy.," This model matches known Galactic parameters such as the rotation curve, local disk density and disk scale-length to high accuracy." + The lack of noticeable warp iu the Magellanic Stream either in position on the sky or radial velocitysce sueeests that the disk aud bulec likely has sinall or negligible effects ou the orbits of the clouds. allowing us to have greater faith im our derived orbits.," The lack of noticeable warp in the Magellanic Stream either in position on the sky or radial velocity—see —suggests that the disk and bulge likely has small or negligible effects on the orbits of the clouds, allowing us to have greater faith in our derived orbits." + In this model the Galactic potential (R.i) is even iu cylindrical coordinates. where A8 is the planar CGalactocentrie radius. and 2 is the distance above the plane of the disk.," In this model the Galactic potential $\Phi(R,z)$ is given in cylindrical coordinates, where $R$ is the planar Galactocentric radius, and $z$ is the distance above the plane of the disk." + The total potential ® is modelled by the stun of the three different potentials: the dark halo Py. a central component Be. aud a disk $5: thus The poteutial of the dark halo ®y is assmmed to be spherical aud of the form where r ds the Galactoceutric radius (7= 27).," The total potential $\Phi$ is modelled by the sum of the three different potentials: the dark halo $\Phi_H$, a central component $\Phi_C$, and a disk $\Phi_D$; thus The potential of the dark halo $\Phi_H$ is assumed to be spherical and of the form where $r$ is the Galactocentric radius $(r^2=R^2+z^2)$ ." +" Thed potentialH has a core radiusH ry. and V,- is- the cirenlar velocity at large r."," The potential has a core radius $r_0$ , and $V_h$ is the circular velocity at large $r$ ." + The poteutial of the central componcnt δε is modelled by two spherical components. represcuting the bulge/stellarhalo and inuer core componcuts:," The potential of the central component $\Phi_C$ is modelled by two spherical components, representing the bulge/stellar-halo and inner core components:" +"than 1.4 arcsec, and airmasses were below 1.2, in order to achieve a good AO correction.","than 1.4 arcsec, and airmasses were below 1.2, in order to achieve a good AO correction." + The targets were used as AO reference stars., The targets were used as AO reference stars. + The data were reduced with the ESO pipeline reduction software., The data were reduced with the ESO pipeline reduction software. +" Dark frames were produced on the same night as each observing run, with the same integration time as the science frames produced."," Dark frames were produced on the same night as each observing run, with the same integration time as the science frames produced." +" Dark subtraction and flat field division were performed on individual frames by the pipeline, before the 5 calibrated frames were combined."," Dark subtraction and flat field division were performed on individual frames by the pipeline, before the 5 calibrated frames were combined." +" This was performed by automated detection of the target, determining an offset then registering and stacking of the 5 calibrated frames."," This was performed by automated detection of the target, determining an offset then registering and stacking of the 5 calibrated frames." +" All the raw data were checked by eye to investigate whether the automated detection did not produce artificial binaries, and in the few ambiguous cases, the add procedure was done by hand."," All the raw data were checked by eye to investigate whether the automated detection did not produce artificial binaries, and in the few ambiguous cases, the shift-and-add procedure was done by hand." +" Typical FWHM of 0.09 arcsec were measured on the final images, and the point spread function’s (PSF) widths ranged between nightly averages of 0.07 arcsec and 0.10 arcsec."," Typical FWHM of 0.09 arcsec were measured on the final images, and the point spread function's (PSF) widths ranged between nightly averages of 0.07 arcsec and 0.10 arcsec." +" The typical Strehl ratio achieved was30%.,, with a fair fraction of targets having Strehl ratios exceeding45%."," The typical Strehl ratio achieved was, with a fair fraction of targets having Strehl ratios exceeding." +. The raw and reduced images were visually inspected to identify binary companions., The raw and reduced images were visually inspected to identify binary companions. + The binary separation and position angle were determined by fitting a two-dimensional Gaussian function to both spatial profiles., The binary separation and position angle were determined by fitting a two-dimensional Gaussian function to both spatial profiles. +" The maximum separation found was 8 arcsec, which is about the maximum possible value in our square view with the target source positioned in the centre of the frame and a 5 arcsec jitter offset."," The maximum separation found was 8 arcsec, which is about the maximum possible value in our square field-of-view with the target source positioned in the centre of the frame and a 5 arcsec jitter offset." +" The minimum separation was 0.1 arcsec, at the resolution limit of our data."," The minimum separation was 0.1 arcsec, at the resolution limit of our data." +" The difference in K-band magnitude (or rather, difference in magnitude of the light through the NB filter, Am, which is very close to ΔΙ), was measured for each component using either aperture photometry, or by estimating the total counts of each star, based on a two-dimensional Gauss fit of the spatial profiles."," The difference in -band magnitude (or rather, difference in magnitude of the light through the NB filter, $\Delta +m$, which is very close to $\Delta K$ ), was measured for each component using either aperture photometry, or by estimating the total counts of each star, based on a two-dimensional Gauss fit of the spatial profiles." +" In the case of more than one companion, the magnitude difference was found for the closest of the companions."," In the case of more than one companion, the magnitude difference was found for the closest of the companions." + In two cases it proved difficult to obtain a reliable estimate of the secondary’s magnitude and the entry is left blank in Table 2.., In two cases it proved difficult to obtain a reliable estimate of the secondary's magnitude and the entry is left blank in Table \ref{restarget}. +" The largest, reliable, magnitude difference was 7.3 magnitude."," The largest, reliable, magnitude difference was 7.3 magnitude." +" As a check on the quality of our data, the parameters of known binaries were compared to our results."," As a check on the quality of our data, the parameters of known binaries were compared to our results." +" For example, HR 3745 was identified as having a separation of 0.55 arcsec and a PA of 275° by ?.."," For example, HR 3745 was identified as having a separation of 0.55 arcsec and a PA of $275^\circ$ by \cite{hartkopf_1996}." + Here we find a separation of 0.55 arcsec., Here we find a separation of 0.55 arcsec. + and a PA of 281°., and a PA of $281^\circ$. + The difference in position angle is not necessarily an indication of the errors involved as it can well be due to orbital motion., The difference in position angle is not necessarily an indication of the errors involved as it can well be due to orbital motion. + We find 14 visual binaries among the 40 observed B stars and 13 visual binaries among the 39 Be stars observed., We find 14 visual binaries among the 40 observed B stars and 13 visual binaries among the 39 Be stars observed. +" Taken at face value, the binary fractions are therefore 35 + 8 confidence interval) for the B stars and 33 + 8 for Be stars."," Taken at face value, the binary fractions are therefore 35 $\pm$ 8 confidence interval) for the B stars and 33 $\pm$ 8 for Be stars." + The numbers for B and Be stars are very close and comfortably within the uncertainties., The numbers for B and Be stars are very close and comfortably within the uncertainties. +" However, it is not obvious that these concern physical binaries or visual binaries whose secondaries have no physical association with the target stars."," However, it is not obvious that these concern physical binaries or visual binaries whose secondaries have no physical association with the target stars." + The best way to test this is to obtain second epoch data to investigate whether the secondary stars share a common proper motion or not., The best way to test this is to obtain second epoch data to investigate whether the secondary stars share a common proper motion or not. +" In the absence of such data, we consider the following as a proxy for the possibility of finding a chance binary companion."," In the absence of such data, we consider the following as a proxy for the possibility of finding a chance binary companion." +" The chances of finding one star in the field of view are large when the stellar density is at, or exceeds, a value of 1 star per field of view of 13.6 x 13.6 arcsec” corresponding to 19.5 stars per arcmin”."," The chances of finding one star in the field of view are large when the stellar density is at, or exceeds, a value of 1 star per field of view of 13.6 $\times$ 13.6 $^{2}$ corresponding to 19.5 stars per $^{2}$." + A measure of the local surface density of stars was obtained from the 2MASS Point Source Catalogue (?) by tallying up all objects within a circle with a radius of 1 arcmin from the target objects., A measure of the local surface density of stars was obtained from the 2MASS Point Source Catalogue \citep{2mass} by tallying up all objects within a circle with a radius of 1 arcmin from the target objects. + The frequency of the stellar densities towards the sample is shown in Fig. 2.., The frequency of the stellar densities towards the sample is shown in Fig. \ref{densfig}. +" Most objects lie in regions with stellar densities less than 10 per arcmin”, less than half a star per field of view, and the presence of a chance companion is not expected."," Most objects lie in regions with stellar densities less than 10 per ${^2}$, less than half a star per field of view, and the presence of a chance companion is not expected." +παπά jp=3.9«10.7? erg/s/em?/LI ffor UGC 732] and UGC 1281 respectively. or only lower and higher than the linear prediction.,"and $\mu=3.9\times 10^{-20}$ $^2$ for UGC 7321 and UGC 1281 respectively, or only lower and higher than the linear prediction." + We conclude that the selection of co-added fibers based on the nominal UVB strength has negligible impact our final limit., We conclude that the selection of co-added fibers based on the nominal UVB strength has negligible impact our final limit. + There are several potential sources of systematic error to the presented spectra., There are several potential sources of systematic error to the presented spectra. + We have already discussed the uncertainties in the model-based conversion of Ha surface brightness to UVB strength in §2.2.., We have already discussed the uncertainties in the model-based conversion of $\alpha$ surface brightness to UVB strength in \ref{sec_mod_fit}. + The uncertainty in the absolute spectral flux calibration due to the applied atmospheric extinction curve is discussed in §3.., The uncertainty in the absolute spectral flux calibration due to the applied atmospheric extinction curve is discussed in \ref{sec_data}. + The uncertainty in the absolute spectral flux calibration due to the standard. star observations is discussed in 83.1.., The uncertainty in the absolute spectral flux calibration due to the standard star observations is discussed in \ref{sec_flux}. + We now analyze a final systematic regarding the relative error determinations in the Ha spectra., We now analyze a final systematic regarding the relative error determinations in the $\alpha$ spectra. + We observe that the propagation of the errors from the data’s original read noise and shot noise does not fully account for the variation in sky subtracted spectra., We observe that the propagation of the errors from the data's original read noise and shot noise does not fully account for the variation in sky subtracted spectra. + This is especially true under bright skylines., This is especially true under bright skylines. + We discuss three possible causes with a focus on the variation. of spectral resolution across different fibers., We discuss three possible causes with a focus on the variation of spectral resolution across different fibers. + In any of the cases. the form of the systematic error will be to add à small percentage of the continuum subtracted sky background spectrum applied linearly with the random error.," In any of the cases, the form of the systematic error will be to add a small percentage of the continuum subtracted sky background spectrum applied linearly with the random error." + First. the instrumental spectral resolution varies by at most 1n different fibers due to small but detectable optical distortions in. the. camera.," First, the instrumental spectral resolution varies by at most in different fibers due to small but detectable optical distortions in the camera." + We further measure from. arc lamp exposures that the variation 1s between the sky and science fibers in the UGC 7321 data and in the UGC 1281 data., We further measure from arc lamp exposures that the variation is between the sky and science fibers in the UGC 7321 data and in the UGC 1281 data. + These factors are presented in column 3 of Table A2. and scaled by the background subtracted sky spectrum and applied as systematic errors i the spectra presented in Figures A4 and A5.., These factors are presented in column 3 of Table \ref{tab_lims} and scaled by the background subtracted sky spectrum and applied as systematic errors in the spectra presented in Figures \ref{fig_UGC7321_spec} and \ref{fig_UGC1281_spec}. + This form of the systematic. as the fractional error in the dispersion times the background subtracted sky spectrum. can be derived simply by taking the first order expansion of a Gaussian function near the line center.," This form of the systematic, as the fractional error in the dispersion times the background subtracted sky spectrum, can be derived simply by taking the first order expansion of a Gaussian function near the line center." + Second. the fiber-to-fiber throughput can vary slightly between flat field calibrations.," Second, the fiber-to-fiber throughput can vary slightly between flat field calibrations." + The relative fiber-to-fiber throughput is calibrated with sky flats taken at dawn and dusk., The relative fiber-to-fiber throughput is calibrated with sky flats taken at dawn and dusk. + This relative throughput has been measured to be stable to 100.," Since $_{10}$ is likely greater than unity, the propeller mechanism could only dominate if the star presently goes through a high density medium with $n$ $\geq$ 100." + The possibility that bbreaks through propeller interaction with a fallback disk seems as wel nulikely., The possibility that breaks through propeller interaction with a fallback disk seems as well unlikely. + First. απλαιο tat the entire N-ray buuinositv is due to friction leads to much lareer P than observe (Zane ct al. 200521:," First, assuming that the entire X-ray luminosity is due to friction leads to much larger $\dot{\rm P}$ than observed (Zane et al. \cite{zane02};" + kaplan et al. 2002))., Kaplan et al. \cite{kaplan02}) ). + Tn the framework oft je. blackbody interpretation of the N-vav data. a part of the optical flux excess above the extrapolation of this uodel to the low-energv haa (sce 5 and Fie. 5Γ))," In the framework of the blackbody interpretation of the X-ray data, a part of the optical flux excess above the extrapolation of this model to the low-energy band (see 5 and Fig. \ref{rxj0720_spec}) )" + colld arise from the reprocessing of N-ravs du a residual «lebris disk when the neutron sar has eutered the pulsar phase., could arise from the reprocessing of X-rays in a residual debris disk when the neutron star has entered the pulsar phase. + Euüssion spectra of such discs were coniputec by Perna et al. (2000)), Emission spectra of such discs were computed by Perna et al. \cite{perna00}) ) + in a uuuber of configurations., in a number of configurations. +" Although iu the B baud. the aitive flux excess above the N-ray blackody extrapolation could be due to a rounait dsc. its cuerey distribution is rather red aud would fail to accouit for the U baud excess,"," Although in the B band, the entire flux excess above the X-ray blackbody extrapolation could be due to a remnant disc, its energy distribution is rather red and would fail to account for the U band excess." + Tverefore. it is not clear wlictjer such discs can significantly coutiibute to the overal optical ciission ofL-3125.," Therefore, it is not clear whether such discs can significantly contribute to the overall optical emission of." + Similar conclusiois are reached by Ἱναρ]αι e . (20033)., Similar conclusions are reached by Kaplan et al. \cite{kaplan03}) ). + The lack uallax and radial velocity as well as the laree ταιο)ainty on the parauieters of the proper motion do not allow to euess the birth place of aas it is apALCithy possible for ((Walter Lattimer 20 32)., The lack of parallax and radial velocity as well as the large uncertainty on the parameters of the proper motion do not allow to guess the birth place of as it is apparently possible for (Walter Lattimer \cite{wl2002}) ). + The star is moving away frou the Galactic pale but is still very close to d., The star is moving away from the Galactic plane but is still very close to it. + Its general directiou of motion is cousiste iowdth a birh eiticr iu the Sco OD2 cou-ex. as for ((1n0r6 pro blvii the lower Ceutaurus Cr regio rat D ~LTO pe). or in tre Vela OD2 Tiuupler | association (D=5360[10 pc).," Its general direction of motion is consistent with a birth either in the Sco OB2 complex, as for (more probably in the lower Centaurus Crux region at D $\sim 170$ pc), or in the Vela OB2 + Trumpler 10 association $=360-410$ pc)." + Iu Doth cases. he fligh time from birth place is of he oxer of pp VT with a rather larec uncertaintw.," In both cases, the flight time from birth place is of the order of $^{6}$ yr with a rather large uncertainty." + This value is eencerallv consistewt with the spin down aoσο ειxd wit1 cooli1ο due (see Zane et al. 20023)., This value is generally consistent with the spin down age and with cooling time (see Zane et al. \cite{zane02}) ). + Several high magneic field (B Lat? (G) radio pulsus with loug perios have been recently discovered (Camilo et al. 2000: , Several high magnetic field (B $\geq$ $^{13}$ G) radio pulsars with long periods have been recently discovered (Camilo et al. \cite{camilo00}; ; +Morris ot al. 20023)., Morris et al. \cite{morris02}) ). + Oue of these radio pulsars. PSR J18301135 las properties astouishinely close to frose of (P," One of these radio pulsars, PSR J1830–1135 has properties astonishingly close to those of (P =" +to see the early-type galaxies dominate over the late types.,to see the early-type galaxies dominate over the late types. + ‘This is a factor 2 shorter in wavelength than found for the SCDAM model., This is a factor 2 shorter in wavelength than found for the SCDM model. + At anv wavelength shorter than this (excep in the near-Hli. between 2-5 microns) quiescent galaxies. orming up to z 10 of stars per vear dominate the tota ight emission (as well as the counts).," At any wavelength shorter than this (except in the near-IR between 2-5 microns), quiescent galaxies, forming up to $\approx$ 10 $_\odot$ of stars per year dominate the total light emission (as well as the counts)." + Again that was no he case with the SCDAL model where spirals dominated the whole background below 400 microns., Again that was not the case with the SCDM model where spirals dominated the whole background below 400 microns. + Finally. we point ou hat the model marginallyC underestimates the backgroumn> measured by DIRE.," Finally, we point out that the model marginally underestimates the background measured by DIRBE." + This cliserepaney could stem from a raction of the population of galaxiesὃν havingo slightlyIn cdilleren spectral shapes in the far-H3 than our model assumes they iive. or having a redshift. distribution cdillerent [from the one we predict they have.," This discrepancy could stem from a fraction of the population of galaxies having slightly different spectral shapes in the far-IR than our model assumes they have, or having a redshift distribution different from the one we predict they have." + Of. course. as these cllects are degenerate. in the sense that colder sources will emit more lux in the far-LR than hot ones. as will closer sources. it is very dillieult to decide which is the dominant. without urther euidance from the observations.," Of course, as these effects are degenerate, in the sense that colder sources will emit more flux in the far-IR than hot ones, as will closer sources, it is very difficult to decide which is the dominant without further guidance from the observations." + ligure 7. shows the comoving luminosity densities predicted bv the model from the far-UV (bottom curves) to the far-LR (top curves)., Figure \ref{figcolum} shows the comoving luminosity densities predicted by the model from the far-UV (bottom curves) to the far-IR (top curves). + His clear from the figure that there already is à greater amount of evolution in the [ar-115 than in the LUV-optical. and we know from the counts (fig. 3))," It is clear from the figure that there already is a greater amount of evolution in the far-IR than in the UV-optical, and we know from the counts (fig. \ref{figinf}) )" + that the submm must show still more evolution than the far-Ilt., that the submm must show still more evolution than the far-IR. + The agreement with the data is quite satisfactory both at low and high redshift and for all wavelengths., The agreement with the data is quite satisfactory both at low and high redshift and for all wavelengths. + This tells us that the luminosity budget of our galaxies is relatively accurate., This tells us that the luminosity budget of our galaxies is relatively accurate. + Another striking feature is that the model predicts that the spheroid (whether in dusty starburst phase or not) contribution to the comoving luminosity density becomes approximately equal to that of late type galaxies. around redshift 2 in the UV. and. [ας and. as early ar 2~ lin the D-band., Another striking feature is that the model predicts that the spheroid (whether in dusty starburst phase or not) contribution to the comoving luminosity density becomes approximately equal to that of late type galaxies around redshift 2 in the UV and far-IR and as early ar $z \sim$ 1 in the B-band. + As a result. the peak of the comoving luminosity density shows a complex behavior as a function of wavelength. starting around z 3in the UV and falling to 2o» 2 in the optical near-Ht. before shifting again to higher redshifts ες7 3) in the far-Ht and submum.," As a result, the peak of the comoving luminosity density shows a complex behavior as a function of wavelength, starting around $z \sim$ 3 in the UV and falling to $z \sim$ 2 in the optical near-IR, before shifting again to higher redshifts $z\sim$ 3) in the far-IR and submm." + This reflects the fact that both UV. and Far-H1t luminosities are extremely sensitive to the instantaneous SER. whereas the optical ancl near-LR. luminosities are tainted by a non-negligible contribution coming from an old stellar population.," This reflects the fact that both UV and far-IR luminosities are extremely sensitive to the instantaneous SFR, whereas the optical and near-IR luminosities are tainted by a non-negligible contribution coming from an old stellar population." + The predicted: comoving star formation rate is compared o observations on figure &.., The predicted comoving star formation rate is compared to observations on figure \ref{figsfr}. + Empty svmbols are observed estimates at various redshifts. uneorrected ος dust extinetion.," Empty symbols are observed estimates at various redshifts, uncorrected for dust extinction." + Filled svmbols inecicate dust corrected estimates ollowing the prescription of Somerville. Primack Faber )01) based on a luminosity dependant correction likely to x more realistic than the traditional scaling by a constant actor (between 3 to 5) at all z (see their [figure 9).," Filled symbols indicate dust corrected estimates following the prescription of Somerville, Primack Faber (2001) based on a luminosity dependant correction likely to be more realistic than the traditional scaling by a constant factor (between 3 to 5) at all $z$ (see their figure 9)." + The illecl point at z~0.7 is derived. from infrared. ISOCAM observations (Flores et al., The filled point at $z \sim 0.7$ is derived from infrared ISOCAM observations (Flores et al. + 1999)., 1999). + The comoving SER is in [fair agreement with the, The comoving SFR is in fair agreement with the +ave far from complete.,are far from complete. + Detailed investigations will certainly adel more members to the present-day open cluster census., Detailed investigations will certainly add more members to the present-day open cluster census. + Besides. since voung clusters ave rather easy to identifv even at large distances. the overlooked clusters are expected to be of the evolved/old age range.," Besides, since young clusters are rather easy to identify even at large distances, the overlooked clusters are expected to be of the evolved/old age range." + Thus. works like the present one are important not only because reliable astrophysical parameters are derived. for a sample of unstudied: clusters.," Thus, works like the present one are important not only because reliable astrophysical parameters are derived for a sample of unstudied clusters." + Perhaps the main importance lies in the unambiguous characterisation of open clusters with ages bevond several 107 vvr., Perhaps the main importance lies in the unambiguous characterisation of open clusters with ages beyond several $10^8$ yr. + A better statistics on the population of these - and older - clusters can be used to investigate cluster formation rates and to constrain the time scale of cluster dissolution in the Galaxy., A better statistics on the population of these - and older - clusters can be used to investigate cluster formation rates and to constrain the time scale of cluster dissolution in the Galaxy. + We thank the reviewer. Dr. AL. Molfat. for. interesting comments ancl suggestions.," We thank the reviewer, Dr. A.F. Moffat, for interesting comments and suggestions." + We acknowledge support [from the Brazilian Institution CNPq., We acknowledge support from the Brazilian Institution CNPq. + Εις publication makes use of cata products. from the Pwo Micron. All Sky Surveys. which is a joint project of the University of Alassachusetts and the Infrared. Processing ancl Analysis Contre/California Institute. of Technology. funded. by. the National Aeronautics and Space Administration and the National Science Foundation.," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Centre/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation." + This research. has mace use of the WEBDA database. operated. at. the Institute for Astronomy of the University of Vienna.," This research has made use of the WEBDA database, operated at the Institute for Astronomy of the University of Vienna." +A hot Jupiter setiles into a Cassini state on the same ~10° vr time scale as spin-orbit svuchronization.,A hot Jupiter settles into a Cassini state on the same $\sim$$10^6$ yr time scale as spin-orbit synchronization. + Whether it ends up in state 1 or 2 depends upon its initial obliquity aud orbital inclination. and upon e. which may be a fuuctiou of time.," Whether it ends up in state 1 or 2 depends upon its initial obliquity and orbital inclination, and upon $\epsilon$, which may be a function of time." + Reasous for variations in e include the disappearance of the protoplanetary disk. migration of the planet or its fellow planets or satellites. contraction of the planet. aud spin alteration of the planet or star.," Reasons for variations in $\epsilon$ include the disappearance of the protoplanetary disk, migration of the planet or its fellow planets or satellites, contraction of the planet, and spin alteration of the planet or star." + If € varies slowly compared to gt anda! (7adiabaticallv?). and slowly compared to the Cassini state settling time. then the obliquity tracks the evolving Cassini obliquity.," If $\epsilon$ varies slowly compared to $g^{-1}$ and $\alpha^{-1}$ (“adiabatically”), and slowly compared to the Cassini state settling time, then the obliquity tracks the evolving Cassini obliquity." + Abrupt changes in € or 0 may cause the planet to leave a Cassini state. aud ultimately to switch states if it lauds tn the basin of attraction of the other stable state.," Abrupt changes in $\epsilon$ or $\theta$ may cause the planet to leave a Cassini state, and ultimately to switch states if it lands in the basin of attraction of the other stable state." + Fig., Fig. + 2 suggests three ways in which a hot Jupiter cau maintain a siguilicaut obliquity., 2 suggests three ways in which a hot Jupiter can maintain a significant obliquity. + First. if €>eet. the only stable state is 2. for which the obliquity is nonzero (05— fas e— x).," First, if $\epsilon > \epsilon_{\rm crit}$, the only stable state is 2, for which the obliquity is nonzero $\theta_2\rightarrow I$ as $\epsilon \rightarrow \infty$ )." + This is the case for the Moon., This is the case for the Moon. + Secoud. it could be in state Lor 2 with e~esq. but this requires a specia coincidence. since e involves both plauet-specilic properties Guamely. its moments of inertia) aix seeminely unrelated quantities specific to other bodies (which determine the precessional torques).," Second, it could be in state 1 or 2 with $\epsilon\sim \epsilon_{\rm crit}$, but this requires a special coincidence, since $\epsilon$ involves both planet-specific properties (namely, its moments of inertia) and seemingly unrelated quantities specific to other bodies (which determine the precessional torques)." + Third. the planet could be in state 2 with €3 for two reasons."," However, its impact on the specific angular momentum is less significant than that of the cold, dense filamentary accretion at $z\ge3$ for two reasons." +" First, whilst cold, dense filamentary gas flows into the central gaseous disc rapidly, gas gravitationally bound to a satellite galaxy orbits around it for the time it takes dynamical friction to drag the satellite galaxy down."," First, whilst cold, dense filamentary gas flows into the central gaseous disc rapidly, gas gravitationally bound to a satellite galaxy orbits around it for the time it takes dynamical friction to drag the satellite galaxy down." +" As a consequence of this process, satellite gas angular momentum in Fig."," As a consequence of this process, satellite gas angular momentum in Fig." + 4 cannot be directly converted into the actual angular momentum of gas accreted onto the central galaxy., \ref{fig:jz} cannot be directly converted into the actual angular momentum of gas accreted onto the central galaxy. +" Second, satellites are accreted along the gas filaments, and hence part of the gas which we conservatively assigned to these satellite galaxies could be regarded as filamentary gas."," Second, satellites are accreted along the gas filaments, and hence part of the gas which we conservatively assigned to these satellite galaxies could be regarded as filamentary gas." +" In this sense, the estimate of ffor the satellites in Fig."," In this sense, the estimate of for the satellites in Fig." +" 4 should be considered as an upper bound, and thus the actual fraction of the total angular momentum contributed by cold filamentary accretion may be slightly higher than reported in this work."," \ref{fig:jz} should be considered as an upper bound, and thus the actual fraction of the total angular momentum contributed by cold filamentary accretion may be slightly higher than reported in this work." +" Finally, the contribution to the total angular momentum from other phases (hot diffuse, hot dense, and cold diffuse) to turns out to be minor for a DM halo of this mass in this redshift range."," Finally, the contribution to the total angular momentum from other phases (hot diffuse, hot dense, and cold diffuse) to turns out to be minor for a DM halo of this mass in this redshift range." +" Whilst the lack of angular momentum of the cold diffuse phase originates from its almost isotropic accretion which cancels its angular momentum to a high degree, that of the hot gas phase is mainly caused by the small mass fraction of hot gas."," Whilst the lack of angular momentum of the cold diffuse phase originates from its almost isotropic accretion which cancels its angular momentum to a high degree, that of the hot gas phase is mainly caused by the small mass fraction of hot gas." +" This result is not very surprising because our DM halo, at 10!!Mc is not massive enough to sustain a stable virial shock (??),, and supernova explosions in simulations are notoriously inefficient at ejecting "," This result is not very surprising because our DM halo, at $10^{11}\msun$ is not massive enough to sustain a stable virial shock \citep{birnboim03,keres05}, and supernova explosions in simulations are notoriously inefficient at ejecting large amounts of hot gas in the halo \citep{maclow99,dubois08,powell11,faucher-giguere11b}." +"Gas infall at low redshift is dominated by smooth accretion rather than mergers (e.g.?,andreferencestherein)..", Gas infall at low redshift is dominated by smooth accretion rather than mergers \citep[e.g.][ and references therein]{fakhouri10}. +" As discussed by ?,, this substantially increases the specific angular momentum of gas in the halo by delivering it coherently via the large-scale cosmic web."," As discussed by \citet{pichon11}, this substantially increases the specific angular momentum of gas in the halo by delivering it coherently via the large-scale cosmic web." +" This is illustrated in Fig. 5,,"," This is illustrated in Fig. \ref{fig:jtot_Co}," + where the time evolution of ffor the Milky Way-like halo in the simulation is found to be smoother at low redshift (z« 3) than at high redshift (z> 3)., where the time evolution of for the Milky Way-like halo in the simulation is found to be smoother at low redshift $z < 3$ ) than at high redshift $z > 3$ ). +" The final mass of the dark matter halo in the NutCO run is Mos~4x10!Mo (see Appendix for a time evolution of the halo mass), which is thought to be close to the mass at which gas accretion transitions from the cold mode to the hot mode (??).."," The final mass of the dark matter halo in the NutCO run is $\mvir\simeq 4\times 10^{11}\msun$ (see Appendix for a time evolution of the halo mass), which is thought to be close to the mass at which gas accretion transitions from the cold mode to the hot mode \citep{birnboim03,ocvirk08}." +" Since the NutCO run does not include supernova feedback, satellite galaxies can retain more gas than in the NutFB run, and thus the resulting relative contribution from the cold, dense filamentary gas to the total angular momentum gets slightly smaller at high redshift than the feedback run (Fig. 6))."," Since the NutCO run does not include supernova feedback, satellite galaxies can retain more gas than in the NutFB run, and thus the resulting relative contribution from the cold, dense filamentary gas to the total angular momentum gets slightly smaller at high redshift than the feedback run (Fig. \ref{fig:jz_Co}) )." +" However, there is little difference in oor bbetween these two runs for z>3."," However, there is little difference in or between these two runs for $z \geq 3$." + Fig., Fig. +" 6 also shows that at zS2, the diffuse hot gas phase becomes the dominant reservoir of angular momentum in the region (0.1€r/Ri,X 1)."," \ref{fig:jz_Co} also shows that at $z\lesssim2$, the diffuse hot gas phase becomes the dominant reservoir of angular momentum in the region $0.1\leq r/\rvir \leq1$ )." +" Considering that the post-shock temperature for an isothermal sphere of mass a few times 10!!Mc (our halo reaches 2 x10!!Mc at z~2, c.f."," Considering that the post-shock temperature for an isothermal sphere of mass a few times $10^{11}\msun$ (our halo reaches 2 $\times 10^{11}\msun$ at $z +\simeq 2$, c.f." +" Appendix A) is Tinockea©31v/8~10°K (2), where Tas(c35.9x(V./[km-s!])?) is the virial temperature, and that the run does not feature any SN feedback, we can safely identify the halo hot gas at z<2 with material which has been shock-heated."," Appendix A) is $ T_{\rm shocked} \gtrsim 3 T_{\rm vir}/8 \sim 10^5 {\rm K}$ \citep{dekel06}, , where $T_{\rm vir}(\simeq35.9 \times (V_{\rm c}/[\rm{km \cdot s}^{-1}])^2)$ is the virial temperature, and that the run does not feature any SN feedback, we can safely identify the halo hot gas at $z\lesssim 2$ with material which has been shock-heated." +" Nonetheless, as Fig."," Nonetheless, as Fig." +" 5 clearly demonstrates, this hot halo gas"," \ref{fig:jtot_Co} clearly demonstrates, this hot halo gas" +"by: where /,CMas./) is (he rate of birth of stars per unit galactic mass al a given (ime wilh main sequence mass. Mas. and 9(M3ijs) is the mass function of main sequence stars wilhmass between Wa,5 and λές+dM3gjs.","by: where $\dot{n_{\star}}(M_{MS},t)$ is the rate of birth of stars per unit galactic mass at a given time with main sequence mass, $M_{MS}$ and $\Phi(M_{MS})$ is the mass function of main sequence stars withmass between $M_{MS}$ and $M_{MS}+dM_{MS}$." +" The observable star formation rate per unit galaxy Diss. s,. al a given lime is then given by: where s,(/) has units of solar masses per vear per unit galaxy mass for έως nieasured in solar masses."," The observable star formation rate per unit galaxy mass, $s_{\star}$, at a given time is then given by: where $s_{\star}(t)$ has units of solar masses per year per unit galaxy mass for $M_{MS}$ measured in solar masses." + Assigning a probabilitv. PpCVas). that a star of main sequence mass. Mars. forms a black hole. the black hole birth rate. μμ). and mass formation rate. μμ). per unil ealaxv mass are given bv: and where MpguoGMs.f.Z) is the initial mass of the black hole (hat can. in principle. be a [function of the main sequence mass of its progenitor. the epoch it was born. the metallicity of the progenitor star and perhaps other parameters.," Assigning a probability, $P_{BH}(M_{MS})$, that a star of main sequence mass, $M_{MS}$ , forms a black hole, the black hole birth rate, $r_{BH}(t)$, and mass formation rate, $s_{BH}(t),$ per unit galaxy mass are given by: and where $M_{BH,0}(M_{MS},t,Z)$ is the initial mass of the black hole that can, in principle, be a function of the main sequence mass of its progenitor, the epoch it was born, the metallicity of the progenitor star and perhaps other parameters." + Black holes in binary svstems have measured masses in the range 6 - 20 citepCO06.SSMINI1..," Black holes in binary systems have measured masses in the range 6 - 20 \\citep{C06,SSMK11}." + For simplicitv. we will scale our results with the assumption that all black holes are born with a single mass. Mg;=10M.," For simplicity, we will scale our results with the assumption that all black holes are born with a single mass, $M_{BH,0} = 10$." +.. We can then define the efficiency. éepy. lo make black holes [rom a generation of stars al time / as: and (hus. The parameter ἐμμ 1s proportional to the fraction ofall stars that are O-twpe stars. fo. and the fraction of O-tvpe starsthat eventually turn into black holes. fj.," We can then define the efficiency, $\epsilon_{BH}$ to make black holes from a generation of stars at time $t$ as: and thus, The parameter $\epsilon_{BH}$ is proportional to the fraction ofall stars that are O-type stars,$f_{O\star}$ , and the fraction of O-type starsthat eventually turn into black holes,$f_{BH}$ ." + This efficiency, This efficiency +stellar evolution theory predicts that both slow expansions and contractions of stars within the pulsation instability strip are expected. that result in increasing and decreasing periods. respectively.,"Stellar evolution theory predicts that both slow expansions and contractions of stars within the pulsation instability strip are expected, that result in increasing and decreasing periods, respectively." + Since the pioneering work of Martin(1938) on the period changes of the RR Lyrae stars in w Centauri. numerous studies have been carried. out with à view to detect period. changes connected. with stellar evolution.," Since the pioneering work of \citet{ma38} on the period changes of the RR Lyrae stars in $\omega$ Centauri, numerous studies have been carried out with a view to detect period changes connected with stellar evolution." + The best targets for such investigations are elobular clusters (GC). as most of them have a Large lid Lvrae population. and their observations have extended utherto over 50100 vears.," The best targets for such investigations are globular clusters (GC), as most of them have a large RR Lyrae population, and their observations have extended hitherto over 50–100 years." + The investigations. however. jwe revealed that. contrary to all expectations. both fast »riod increase and decrease. repeated: irregular and/or abrupt changes may also characterize the period. changes.," The investigations, however, have revealed that, contrary to all expectations, both fast period increase and decrease, repeated irregular and/or abrupt changes may also characterize the period changes." + This kind of period-change behaviour could not at all »* reconciled with evolutionary models., This kind of period-change behaviour could not at all be reconciled with evolutionary models. + Moreover. the observations have shown peculiar cases. too.," Moreover, the observations have shown peculiar cases, too." + Some stars showed: regular or quasi-regular periocd-change behaviour or more than half à century. then suddenly. continuous or abrupt irregular changes in their period have taken place.," Some stars showed regular or quasi-regular period-change behaviour for more than half a century, then suddenly, continuous or abrupt irregular changes in their period have taken place." +amplitude for this to be readily. visible.,amplitude for this to be readily visible. + There arc also examples where the DNO aniplituce is unallectecd by the QPO modulation., There are also examples where the DNO amplitude is unaffected by the QPO modulation. + The examples are too numerous to show in their entirety. but reference. to Fig.," The examples are too numerous to show in their entirety, but reference to Fig." + 14 illustrates some of these points.," \ref{dnobup} + illustrates some of these points." + Some of the QPO maxima have large DNO amplitude associated: with them in a way that indicates the growth and decay of DNO amplitude over the QPO maximum., Some of the QPO maxima have large DNO amplitude associated with them – in a way that indicates the growth and decay of DNO amplitude over the QPO maximum. + Other QPO maxima have DNOs of low amplitude., Other QPO maxima have DNOs of low amplitude. + DNOs of large amplitude can be seen midwavy between the fourth and fifth ΟΡΟ maxima., DNOs of large amplitude can be seen midway between the fourth and fifth QPO maxima. + DNOs of nearly constant amplitude through a ΟΡΟ evcle can be seen in the inset to Fig. L.., DNOs of nearly constant amplitude through a QPO cycle can be seen in the inset to Fig. \ref{vwfig1}. + To expand on these cxamples. and those given in WB and RW. we show detailed analyses in Figs.," To expand on these examples, and those given in WB and RW, we show detailed analyses in Figs." + 20. and 21., \ref{omcex1} and \ref{omcex2}. +. Fig., Fig. + 20. shows a positive correlation between DNOs and QPOs. in the sense that the DNOs have maximum amplitude at the peaks of the QPOs.," \ref{omcex1} shows a positive correlation between DNOs and QPOs, in the sense that the DNOs have maximum amplitude at the peaks of the QPOs." + During this run the DNO phases appear largely independent. of the amplitude variations., During this run the DNO phases appear largely independent of the amplitude variations. + In Fig., In Fig. + 21 the DNO amplitudes appear relatively uncorrelated with the large QPO modulation. but there is an overall anticorrelation between DNO phase and QPO.," \ref{omcex2} the DNO amplitudes appear relatively uncorrelated with the large QPO modulation, but there is an overall anticorrelation between DNO phase and QPO." + It is certainly noticeable that the time scale of the moculations of DNO phase is similar to that of the QPO variations., It is certainly noticeable that the time scale of the modulations of DNO phase is similar to that of the QPO variations. + ‘This study was stimulated by the lighteurve of VW Ivi at the end of outburst. obtained in February 2000 (Fig. 12).," This study was stimulated by the lightcurve of VW Hyi at the end of outburst, obtained in February 2000 (Fig. \ref{vwfig1}) )." + At first sight this light curve looks typical of the fHickering seen ina CV late in outburst., At first sight this light curve looks typical of the flickering seen in a CV late in outburst. + But closer inspection shows that there is almost no IHlickering present the light curve is mace up of an orbital modulation plus variable amplitude DNOs and QPOs., But closer inspection shows that there is almost no flickering present – the light curve is made up of an orbital modulation plus variable amplitude DNOs and QPOs. + The evolution of the DNO and QPO periods in us light curve has assisted in selection among the various models of DNOs and QPOs that have been proposed., The evolution of the DNO and QPO periods in this light curve has assisted in selection among the various models of DNOs and QPOs that have been proposed. + Lt is vident that QPOs are more common than realised. — their gaiort coherence time results in a broad and noisy signal in the Fourier transform. where (as originally pointed out bv Patterson et al.," It is evident that QPOs are more common than realised – their short coherence time results in a broad and noisy signal in the Fourier transform, where (as originally pointed out by Patterson et al." + 1977). they are easily overlooked even rough they may be obvious to the eve in the light. curve.," 1977), they are easily overlooked even though they may be obvious to the eye in the light curve." + There is a need for an operational definition of QPOs. which can be applied objectively to the light curves of CVs.," There is a need for an operational definition of QPOs, which can be applied objectively to the light curves of CVs." + We thank the American Association of Variable Star Observers and. the Roval Astronomical Society of New Zealand. specifically and. respectively Janet Mattei and Frank Bateson. for supplving magnitudes of VW Livi from their archives.," We thank the American Association of Variable Star Observers and the Royal Astronomical Society of New Zealand, specifically and respectively Janet Mattei and Frank Bateson, for supplying magnitudes of VW Hyi from their archives." + We thank also Darragh ODonoghue lor the use of his EAGLE program., We thank also Darragh O'Donoghue for the use of his EAGLE program. + Many. observers. have contributed to the VW. Ivi archive. in particular. anc at," Many observers have contributed to the VW Hyi archive, in particular, and at" +We also note that the stellar profile is contracted similarly to the dark mater profile. because gas accretion is faster than star formation.,"We also note that the stellar profile is contracted similarly to the dark mater profile, because gas accretion is faster than star formation." + All of the simulations considered here indicate some degree of enhancement of the dark matter profile., All of the simulations considered here indicate some degree of enhancement of the dark matter profile. + Not a single case indicates halo expansion rather than contraction., Not a single case indicates halo expansion rather than contraction. + Figure 6 combines the resulting constraints on the parameters A and w of Equation (3)).," Figure \ref{fig:aw} + combines the resulting constraints on the parameters $A$ and $w$ of Equation \ref{eq:rrave}) )." + The models do not fill all the available parameter space. but instead concentrate in a fairly narrow region in which A and w are strongly correlated.," The models do not fill all the available parameter space, but instead concentrate in a fairly narrow region in which $A$ and $w$ are strongly correlated." + The original MAC model suggested by ? falls right in the middle of the new distribution., The original MAC model suggested by \citet{gnedin_etal04} falls right in the middle of the new distribution. + It is interesting to determine which combination of the parameters A and w yields the same amount of contraction., It is interesting to determine which combination of the parameters $A$ and $w$ yields the same amount of contraction. + Given the radial dependence of the mass enhancement factor Fi (Equation 7)). the solution to this problem varies with radius.," Given the radial dependence of the mass enhancement factor $F_M$ (Equation \ref{eq:fm}) ), the solution to this problem varies with radius." + However. we can remove most of the radial dependence by defining the enhancement factor relative to the SAC model: and evaluating it at some inner radius where the linear approximation for the contraction factor v(7) is valid.," However, we can remove most of the radial dependence by defining the enhancement factor relative to the SAC model: and evaluating it at some inner radius where the linear approximation for the contraction factor $y(r)$ is valid." + We take r20.005ri. which corresponds to about 1 kpe for the Milky Way galaxy.," We take $r = 0.005 \, r_{\rm +vir}$, which corresponds to about 1 kpc for the Milky Way galaxy." +" The exact value of r affects the resulting value of parameter w (for a given A) only logarithmically. as long as ror,. Which we take again to be p,=0.05A."," The exact value of $r$ affects the resulting value of parameter $w$ (for a given $A$ ) only logarithmically, as long as $r \ll r_e$, which we take again to be $r_e = 0.05 \, r_{\rm vir}$." + Lines in Figure 6 show the relation between A and w corresponding to three values of fij., Lines in Figure \ref{fig:aw} show the relation between $A$ and $w$ corresponding to three values of $f_M$. + All simulations but three fall below the level of contraction. predicted by the SAC model., All simulations but three fall below the level of contraction predicted by the SAC model. + At the same time. no simulation falls below the level of fy=0.3.," At the same time, no simulation falls below the level of $f_M = 0.3$." + Therefore. the MAC model is well constrained to be able to reliably predict the amount of dark matter in the inner regions of galaxies and clusters.," Therefore, the MAC model is well constrained to be able to reliably predict the amount of dark matter in the inner regions of galaxies and clusters." + The isocontours of constant fj become even more horizontal at smaller +., The isocontours of constant $f_M$ become even more horizontal at smaller $r$. + This suggests that parameter w may be more important than A in deseribing the amount of contraction., This suggests that parameter $w$ may be more important than $A$ in describing the amount of contraction. + It is also desirable to describe the strength of the, It is also desirable to describe the strength of the +ALEW. acknowledges the generous help of the Cirinvace Scholarship from the University of Melbourne for assistance in carrying out this research.,MTW acknowledges the generous help of the Grimwade Scholarship from the University of Melbourne for assistance in carrying out this research. + Phanks are due to Frank Masci. [or use of some of the polarisation data prior to publication. and to the referee. (Patrick Leahy) for some very useful comments which improved the paper.," Thanks are due to Frank Masci, for use of some of the polarisation data prior to publication, and to the referee (Patrick Leahy) for some very useful comments which improved the paper." +The assembly of rich. X-ray luminous galaxy clusters is such that the largest baryonic mass fraction of the system is oecupied by hot7~10110° K gas pervading the intracluster medium (ICM).,"The assembly of rich, X-ray luminous galaxy clusters is such that the largest baryonic mass fraction of the system is occupied by hot$T\sim10^7-10^8$ K gas pervading the intracluster medium (ICM)." +" In the central-- regions (7=10—100 kpe) of many clusters. the time scale for this gas to cool to T<10+ K can be shorter than the cluster lifetime (e.g.. Cowie&Binney1977:Fabian&NulsenEdgeetal. 1992)). giving rise to a subsonic. pressure-driven cooling flow that deposits mass onto the luminous and massive cD elliptical galaxy at the cluster center. """," In the central regions $r\la10-100$ kpc) of many clusters, the time scale for this gas to cool to $T\la 10^4$ K can be shorter than the cluster lifetime (e.g., \citealt{cowie77,fabian77,edge92}) ), giving rise to a subsonic, pressure-driven cooling flow that deposits mass onto the luminous and massive cD elliptical galaxy at the cluster center. “" +"Cool core"" clusters such as these often exhibit intense optical emission line nebulae associated with these central brightest cluster galaxies (BCGs).",Cool core” clusters such as these often exhibit intense optical emission line nebulae associated with these central brightest cluster galaxies (BCGs). + The nebulae exhibit extended Lyo emission (Hu1992) and far UV continuum emission (O'Deaetal.2004)., The nebulae exhibit extended $\alpha$ emission \citep{hu92} and far UV continuum emission \citep{odea04}. +. A previous study of two BCGs. Abell 1795 and Abell 2597 (0Ῥεαetal.2004).. found that the nebula exhibited both a diffuse component of Lye and more compact features such as knots and filaments.," A previous study of two BCGs, Abell 1795 and Abell 2597 \citep{odea04}, found that the nebula exhibited both a diffuse component of $\alpha$ and more compact features such as knots and filaments." + The Ενα emission was closely tied to the radio morphology suggesting that star formation and associated ionization. was present at the edges of radio lobes., The $\alpha$ emission was closely tied to the radio morphology suggesting that star formation and associated ionization was present at the edges of radio lobes. + That work demonstrated how Lya and far-ultraviolet continuum observations provide unique constraints on the physical properties of the nebulae in clusters., That work demonstrated how $\alpha$ and far-ultraviolet continuum observations provide unique constraints on the physical properties of the nebulae in clusters. + The far-UV continuum together with optical and infrared observations constrain the star formation history and the properties of young stars associated with the nebula., The far-UV continuum together with optical and infrared observations constrain the star formation history and the properties of young stars associated with the nebula. + The Εναto Ho or Hj flux ratio is a diagnostic of ionization. metal and dust content (Ferland&Osterbrock1985:Binetteetal. 1993).," The $\alpha$to $\alpha$ or $\beta$ flux ratio is a diagnostic of ionization, metal and dust content \citep{ferland85,binette93}." +. Previous optical and UV observations have found evidence for significant star formation in some BCGs in cool core clusters (Johnstone&Fabian1987;RomanishinΜο-2009;Pipinoetal. 2009)..," Previous optical and UV observations have found evidence for significant star formation in some BCGs in cool core clusters \citep{johnstone87,romanishin87, mcnamara89, mcnamara93, mcnamara04, + mcnamara04b, hu92, crawford93, hansen95, allen95, smith97, + cardiel98, hutchings00, oegerle01, + mittaz01,odea04,hicks05,rafferty06,bildfell08,loubser09,pipino09}." + Nearly all BCGs with young stellar populations are in cooling flows (Bildfelletal.2008;Loub-seretal. 2009).," Nearly all BCGs with young stellar populations are in cooling flows \citep{bildfell08,loubser09}. ." +. However. some BCGs in cooling flows do not have significant star formation (Quillenetal.2008:Loub- 2009)..," However, some BCGs in cooling flows do not have significant star formation \citep{quillen08,loubser09}. ." + Hence BCGs exhibiting elevated rates of star, Hence BCGs exhibiting elevated rates of star +EOS A (Paudharipande 1971) models the interaction of ueutrons at high densities with a Reid soft-core potential.,EOS A (Pandharipande 1971) models the interaction of neutrons at high densities with a Reid soft-core potential. + EOS D is inodel V of Bethe Johuson (1971)., EOS D is model V of Bethe Johnson (1974). + In EOS L the nucleon interaction is modeled iu terms of a mica scalar field (Pandharipande. Simith 1975).," In EOS L the nucleon interaction is modeled in terms of a mean scalar field (Pandharipande, Smith 1975)." + Both EOSs UT (Winuga. Fiks. Fabrocini 1988) aud FPS (Lorenz. Ravenhall. Pethick. 1993) are moderu versious of au emlier EOS proposed by Friedisui Paucdharipande (19813. which curplovs both two-body (UL1) aud Έποςbody uucleou interactions (TNT).," Both EOSs UT (Wiringa, Fiks, Fabrocini 1988) and FPS (Lorenz, Ravenhall, Pethick, 1993) are modern versions of an earlier EOS proposed by Friedman Pandharipande (1981), which employs both two-body (U14) and three-body nucleon interactions (TNI)." + EOS UT improves the treatment of matter at ligh densities. while FPS describes the iuteractious in terius of a Skyrme model.," EOS UT improves the treatment of matter at high densities, while FPS describes the interactions in terms of a Skyrme model." + EOS APR (Akinal. Pandharipaude. Ravenhall 1998) adopts a modern two-nucleou interaction (ALS) together with boost corrections. as well as the UIN. threc-nucleou potential.," EOS APR (Akmal, Pandharipande, Ravenhall 1998) adopts a modern two-nucleon interaction (A18) together with boost corrections, as well as the UIX three-nucleon potential." + The different descriptions of uucleou interactions affect the EOS only at high densities., The different descriptions of nucleon interactions affect the EOS only at high densities. +" At low deusities (pyX10eem 7). the EOSs eiiploy. the Fevumean. Metropolis. Teller (1919) EOS. joining outo the Bavin. Pethick. Sutherland (1971) EOS up to neutron drip at pyzc1ottecm Ὁ, "," At low densities $\rho_0 \lesssim 10^4 \mbox{g +cm}^{-3}$ ), the EOSs employ the Feynman, Metropolis, Teller (1949) EOS, joining onto the Baym, Pethick, Sutherland (1971) EOS up to neutron drip at $\rho_0 \approx 4 \times 10^{11} \mbox{g cm}^{-3}$ ." +Above neutrou drip. EOSs A. D join onto the Bavin. Bethe. Pethick (1971) EOS. while EOSs L and UT join outo the Negele Vautherin (1973) EOS.," Above neutron drip, EOSs A, D join onto the Baym, Bethe, Pethick (1971) EOS, while EOSs L and UT join onto the Negele Vautherin (1973) EOS." + EOS APR joius outo the FPS EOS below a umber density of 0.1 fin. (py=126ς10!iean 7). All EOSs are read iuto our numerical code in tabular fori. listing rest-density py. the total energw densitv €. and the pressure2 at discreet points.," EOS APR joins onto the FPS EOS below a number density of 0.1 $^{-3}$ $\rho_0 = 1.26 \times 10^{14} \mbox{g cm}^{-3}$ ), All EOSs are read into our numerical code in tabular form, listing rest-density $\rho_0$ , the total energy density $\epsilon$, and the pressure$P$ at discreet points." + Iutermecdiate values, Intermediate values +a color term of 3.3.? a value that is appropriate for nearby Calactic stars with simular colors to Cepheids (Madore&Freediiau1991).,"a color term of 3.3, a value that is appropriate for nearby Galactic stars with similar colors to Cepheids \citep{mf91}." +". ReeOn best estimates for scene to cleLLYASS values ranging frou 2 to 3H (oy56accordingtoOpolsà&Ciurla 1981). πο a otter approach would at first sight sec to be f) lc rTegreosskn techniques ou Cepheid samples to establish. au Clupirica correction term (ο,ον,Stit1990:Majaessetal.20 js)..."," Recent best estimates for seem to encompass values ranging from 2 to 3 \citep[or $\sim$5--6 according to][]{op84}, so a better approach would at first sight seem to be to use regression techniques on Cepheid samples to establish an empirical correction term \citep[e.g.,][]{st90,ma08}." + Madore&Freedina1{1991) aYSue otherwise., \citet{mf91} argue otherwise. + Receut Caibratious typically iuclude such a parameter base| upon the Wesenhleit formulation with luvwenitudes from t16 Lear-infrared spectral regio 1(e.e..Denedicetal.2007:vauLcemvenetal.2007:Fouquéct2007 ).," Recent calibrations typically include such a parameter based upon the Wesenheit formulation with magnitudes from the near-infrared spectral region \citep[e.g.,][]{bn07,vl07,fo07}." +. Use of ucar-infrared magnitudes las f1ο »oeueficial effect of reducing. or oossiblv chininatine. effects that migit be linked ο redching leUe variations or to metallicity cüffereuces betwecuu Cepheids. which affect colors (niünlv 1 the ο]οσα] region) corresponding to προς effective feniperares as well as the degree ο| peneration mto the strip for post-lyvdrogen-burning stars.," Use of near-infrared magnitudes has the beneficial effect of reducing, or possibly eliminating, effects that might be linked to reddening law variations or to metallicity differences between Cepheids, which affect colors (mainly in the optical region) corresponding to specific effective temperatures as well as the degree of penetration into the strip for post-hydrogen-burning stars." + But he question renas as to whether or not differences iu the properties of interstelar dst from) one galaxy to another are inporaut οnoueh o affect the caibration used to establish Cepheie clistances., But the question remains as to whether or not differences in the properties of interstellar dust from one galaxy to another are important enough to affect the calibration used to establish Cepheid distances. + Cunon the imporance he PL calibration of reddeimg and the koratio :iudivkual Cepheids wihin the iustability str» dt Is of interest to COusider the οher hods of calibration availade.," Given the importance to the PL calibration of reddening and the location of individual Cepheids within the instability strip, it is of interest to consider the other methods of calibration available." + The reddqiue he Cepjekds studied by Benedictetal.{2002.20OT) was established frou he colors anc specra types (jM stars line along the same Ines of sight. Le. frou their space reddeings.," The reddening of the Cepheids studied by \citet{bn02,bn07} was established from the colors and spectral types of stars lying along the same lines of sight, i.e. from their space reddenings." + Di uaa cases the stais also have independent estimates for trein reddenines frou spectroscopic iiethods (Lovvukhetal.2008) or their equivalents (Turnerctal.1987).. aud some frou indepenudenutIv-denvec space redadenuimgs (Turner2001:Laney&Callwell2007).. πο it Is iuforulative ο compare the results.," In many cases the stars also have independent estimates for their reddenings from spectroscopic methods \citep{ko08} or their equivalents \citep{te87}, and some from independently-derived space reddenings \citep{tu01,lc07}, so it is informative to compare the results." + Talle 1 colpares availade estimates for he reddening ofHST Cepheids. as obtained frou space reddevines (Turner198da.1981:Bouceictetal. 2007).. spectroscopic inethods (Turneretal.1957:Iovtvukhetal. 2008).. aud conrpilatious stanceuxlizec to space reddenings (Turner20VL:Laney&Cadwell 2007).," Table \ref{tab1} compares available estimates for the reddening of Cepheids, as obtained from space reddenings \citep{tu80a,tu84,bn07}, spectroscopic methods \citep{te87,ko08}, and compilations standardized to space reddenings \citep{tu01,lc07}." +. Photometric reddeuigsoe were omitted from such a comparison because they do not ahwavs account for the natural dispersion witlin the instability strip of Cepheids at coustaut perio (seealsoMadore&Frecedauan1991)., Photometric reddenings were omitted from such a comparison because they do not always account for the natural dispersion within the instability strip of Cepheids at constant period \citep[see also][]{mf91}. +. An attempt was made to deduce independent field reddenines for the Cepheids using 2MÁSS data (Cutiietal.2003).. as illstrated by Thrueretal.(2008).. but relatively arge uncertaiities du the infrared data. iu conibiration with the small reddeuings and distances to the stars. produced results with very large uncertainties. nimc1 larecr than values typically no larecr than +£0.02 flat apply to the color excesses cied in Table 1..," An attempt was made to deduce independent field reddenings for the Cepheids using 2MASS data \citep{cu03}, as illustrated by \citet{te08}, but relatively large uncertainties in the infrared data, in combination with the small reddenings and distances to the stars, produced results with very large uncertainties, much larger than values typically no larger than $\pm0.02$ that apply to the color excesses cited in Table \ref{tab1}." +" The differences between he inecependent reddening estimates i Table isi no larecr than the associated uncertaitics 1n the values. so a simple ""unsveiehted average Was onued for t10 Weald reddening of cach Cepheid."," The differences between the independent reddening estimates in Table \ref{tab1} is no larger than the associated uncertainties in the values, so a simple unweighted average was formed for the mean reddening of each Cepheid." + TThose values are lised in cohunn 8 of Table |. and he inferred iutriusic colors are plotted as a funcion of logarithm of pulsation period iu Fie. 1..," Those values are listed in column 8 of Table \ref{tab1}, and the inferred intrinsic colors are plotted as a function of logarithm of pulsation period in Fig. \ref{fig1}." + Also plotted are lutrinsic colors for Cehies of mown reddeniig. based upon an uupubished comxlation of photometric and spectroscopic redenines tied to a space reddening scale (Tinlor2001).. and intrinsic COors for Cepheids iu the salle salupieo lat rave the largest ampitudes for tici pulsation xxiod.," Also plotted are intrinsic colors for Cepheids of known reddening, based upon an unpublished compilation of photometric and spectroscopic reddenings tied to a space reddening scale \citep{tu01}, and intrinsic colors for Cepheids in the same sample that have the largest amplitudes for their pulsation period." + The latter shotId concentrae towards 1ο iter of the iustability strip or shehtly bhewarcl. where Cepheids of hugest lieht uplitide are ocatoc.," The latter should concentrate towards the center of the instability strip or slightly blueward, where Cepheids of largest light amplitude are located." + The scatter sποπ] closely represent the actual scatter Causec bv the natural dispersiou of C'epieids wilin tjo 1stability strip. fjus the »oluts correspoxdiug otιο Cepheids xovide direct information aON where frey lie din the iustabiitv strip.," The scatter should closely represent the actual scatter caused by the natural dispersion of Cepheids within the instability strip, thus the points corresponding to the Cepheids provide direct information about where they lie in the instability strip." + Aninspectio iof the data leads to the unexpected result that he three Cepheids of longest period in heLST sauple are redder (cooler) than Cepheids ving near the ceuter ofthe mstabilitv strip. while he bluest (hotest) Cepheids in the erotp are all at the short period end.," An inspection of the data leads to the unexpected result that the three Cepheids of longest period in the sample are redder (cooler) than Cepheids lying near the center of the instability strip, while the bluest (hottest) Cepheids in the group are all at the short period end." + The three longest period Cepheids iu theMST sauple are therefore less unminous tiu Cepheids vius near strip center. so any attempt to construct a poriod-absolute naguitude relation based solely ou the siuuple would pathologically unclerestimate the true slope of the relationship.," The three longest period Cepheids in the sample are therefore less luminous than Cepheids lying near strip center, so any attempt to construct a period-absolute magnitude relation based solely on the sample would pathologically underestimate the true slope of the relationship." + Just such an effect was noted in the Benedictetal.(2007) study. without the natural explanation evident frou the intrinsic," Just such an effect was noted in the \citet{bn07} study, without the natural explanation evident from the intrinsic" +Although the hardness ratios can be used to detect spectral variations m a model independent way. they cannot identify. unambiguously. which are the model components that are actually responsible for the observed variations.,"Although the hardness ratios can be used to detect spectral variations in a model independent way, they cannot identify, unambiguously, which are the model components that are actually responsible for the observed variations." + The main reason is that in most cases (including our work) the hardness ratios are computed using light curves over energy bands which are quite broad. and as a result multiple components contribute to the observed count rate.," The main reason is that in most cases (including our work) the hardness ratios are computed using light curves over energy bands which are quite broad, and as a result multiple components contribute to the observed count rate." + In order to investigate ma quantitative way the constrains that the colour-flux diagrams in Fig., In order to investigate in a quantitative way the constrains that the colour–flux diagrams in Fig. +" + can impose on current theoretical models. we performed the ""experiment"" we describe below."," \ref{colourflux} can impose on current theoretical models, we performed the “experiment"" we describe below." + Thermal Comptonization is the most commonly accepted mechanism for the X-ray emission from radio-quiet AGN., Thermal Comptonization is the most commonly accepted mechanism for the X–ray emission from radio-quiet AGN. +" We therefore considered a ""Comptonization plus reflection"" model. similar to what has been used numerous times in the past to fit the high energy X-ray spectra of many AGN. and we computed the expected hardness ratios when the model parameters were allowed to vary."," We therefore considered a “Comptonization plus reflection"" model, similar to what has been used numerous times in the past to fit the high energy X–ray spectra of many AGN, and we computed the expected hardness ratios when the model parameters were allowed to vary." +" We produced ""theoretical"" colour-flux diagrams. which we then compared with the observed diagrams."," We produced “theoretical"" colour–flux diagrams, which we then compared with the observed diagrams." + We first used the thermally comptonized continuum model in XSPEC (Zdziarskietal..1996:Zyckial..1999).," We first used the thermally comptonized continuum model in XSPEC \cite{zdziarski96,zycki99}." +". The main model parameters are the spectral slope. Γ. and the electron temperature. &7,. which determines the high energy rollover of the spectrum."," The main model parameters are the spectral slope, $\Gamma$, and the electron temperature, $kT_{e}$, which determines the high energy rollover of the spectrum." +" As for the other model parameters. we assumed a seed photon temperature of 10 eV. and a ""diskblackbody"" input spectrum (none of these two parameters affect significantly the shape of the spectrum at energies >20 keV)."," As for the other model parameters, we assumed a seed photon temperature of 10 eV, and a “diskblackbody"" input spectrum (none of these two parameters affect significantly the shape of the spectrum at energies $> 20$ keV)." + We then added a reflection component using the code (Magdziarz&Zdziarski..1995)... available in. XSPEC. which calculates the reflected exponentially cutoff power law spectrum from neutral material.," We then added a reflection component using the code \cite{magdziarz95}, available in XSPEC, which calculates the reflected exponentially cutoff power law spectrum from neutral material." + In all cases. we assumed an inclination angle of 45°. Fyexray=Ungthconp- and Epexray=kTenthcomp- ," In all cases, we assumed an inclination angle of $^{\rm o}$ , $\Gamma_{\tt pexrav} =\Gamma_{\tt Nthcomp}$, and $E_{c, {\tt pexrav}}=kT_{e,\tt Nthcomp}$ ." +The reflection amplitude. R. was negative so that would output the reflection component only.," The reflection amplitude, $R$, was negative so that would output the reflection component only." + We also considered neutral absorption with Ny=107em- (similar to what has been measured in the past for the objects we studied - see Appendix). although it does not affect significantly the spectrum at the energies we consider.," We also considered neutral absorption with ${\rm N}_{\rm H}=10^{23}\,{\rm cm}^{-2}$ (similar to what has been measured in the past for the objects we studied - see Appendix), although it does not affect significantly the spectrum at the energies we consider." +" We downloaded the diagonal BAT survey instrumentresponse fromHEASARC"".. and we used XSPEC to simulate Swift/BAT spectra for [=1.3.1.7.2.1."," We downloaded the diagonal BAT survey instrumentresponse from, and we used XSPEC to simulate /BAT spectra for ${\Gamma}=1.3,1.7,2.1$." +" For each T. we considered three different electron temperatures: AT,=50.100.200 kkeV. For each pair of (E.K7,) values. we used five different normalization values ofNthcomp. from A=| to A=9 photon em s7!. in steps of AA=2 photon em s. to produce 5 simulated spectra (our results do not depend on the particular choice of the model normalization. as long as the ratio (Aja/Aiii) Is large enough to reproduce the observed maximum-to-minimum flux variations)."," For each $\Gamma$, we considered three different electron temperatures: $kT_e=50,100,200$ keV. For each pair of $\Gamma, kT_e$ ) values, we used five different normalization values of, from $A=1$ to $A=9$ photon $^{-2}$ $^{-1}$, in steps of $\Delta A=2$ photon $^{-2}$ $^{-1}$ , to produce 5 simulated spectra (our results do not depend on the particular choice of the model normalization, as long as the ratio $A_{\rm max}/A_{\rm min}$ ) is large enough to reproduce the observed maximum–to–minimum flux variations)." + Finally. we added to the model a flux reflection component’.," Finally, we added to the model a flux reflection ." +".. The flux of the reflection component was such that Rican0.3.1. and 1.5 in the case of the ""mean"" spectrum (1.e. the continuum when Aa,5 photon em s! )."," The flux of the reflection component was such that $R_{\rm mean}=0.3, 1,$ and 1.5 in the case of the “mean"" spectrum (i.e. the continuum when $A_{\rm mean}=5$ photon $^{-2}$ $^{-1}$ )." + Obviously. it this scenario. R should decrease with increasing flux. since the flux of the reflection component remains constant.," Obviously, in this scenario, $R$ should decrease with increasing flux, since the flux of the reflection component remains constant." + For each model spectrum we computed the (20-100). (20— and (50-100) keV count rates. and the respective hardness ratio: (50-100 keV count rate)/(20-50 keV count rate).," For each model spectrum we computed the (20–100), (20--50) and (50–100) keV count rates, and the respective hardness ratio: (50–100 keV count rate)/(20–50 keV count rate)." + As a result. we were able to produce model colowflux diagrams. which we plot in Fig. 6..," As a result, we were able to produce model colour–flux diagrams, which we plot in Fig. \ref{bat_sim}." +" The x-axis in these plots indicate the 20-100 keV model count rate normalized to the same band count rate of the ""mean"" spectrum."," The x–axis in these plots indicate the 20–100 keV model count rate normalized to the same band count rate of the “mean"" spectrum." + The spread of the values along the x-axis is similar in the model and in the observed colour-flux diagrams (Fig. 4))., The spread of the values along the x–axis is similar in the model and in the observed colour–flux diagrams (Fig. \ref{colourflux}) ). +" The results for the R4,=0.3.I. and 1.5 case are plotted in the left. middle and right panel in this figure. respectively."," The results for the $R_{\rm mean}=0.3, 1,$ and 1.5 case are plotted in the left, middle and right panel in this figure, respectively." + Within each panel. the top. middle and bottom curves indicate the model colour-flux diagrams for Γ=1.3.1.7 and 2.1. respectively.," Within each panel, the top, middle and bottom curves indicate the model colour–flux diagrams for $\Gamma=1.3, 1.7$ and $2.1$, respectively." + Clearly. even the mean ZR value of a source can provide information about its Ανω and E.," Clearly, even the mean $HR$ value of a source can provide information about its $R_{\rm mean}$ and $\Gamma$." +" For example. the expected ZR value should be ~0.2 if Ru,=15 and [~2.1. as opposed to ~0.55—0.7 in the case of Ryan0.3 andP- 1.3."," For example, the expected $HR$ value should be $\sim 0.2$ if $R_{\rm mean}=1.5$ and $\Gamma\sim 2.1$, as opposed to $\sim 0.55-0.7$ in the case of $R_{\rm mean}=0.3$ and$\Gamma\sim 1.3$ ." + This is of course the case for the widest separatedparameter values we considered.and the difference between the model colour-flux diagrams becomes less pronounced for smaller differences in the input physical parameters.," This is of course the case for the widest separatedparameter values we considered,and the difference between the model colour–flux diagrams becomes less pronounced for smaller differences in the input physical parameters." + Nevertheless. a comparison between the observed," Nevertheless, a comparison between the observed" +is only one source iiaser with a peak flix deusity less than 1 Jy. which could be due to the nature of the methanol niaser huninosity fiction.,"is only one source maser with a peak flux density less than 1 Jy, which could be due to the nature of the methanol maser luminosity function." + This paper preseuts only the results of the survey., This paper presents only the results of the survey. + In a separate paper. we will report ou the enviroment around the masers. using ongoing observations of various molecules such as CO. IICO!. CN and NIT.," In a separate paper, we will report on the environment around the masers, using ongoing observations of various molecules such as CO, $^+$, CN and $_3$." + This will also elucidate on uv differences between faint masers aud their bright counterpart:Z, This will also elucidate on any differences between faint masers and their bright counterparts. + This will also elucidate on uv differences between faint masers aud their bright counterpart:ZF, This will also elucidate on any differences between faint masers and their bright counterparts. +,$t'>3/2$. + We show in Figure 4.2. (top) that this radial (hermal inversion causes (he migration rate to be positive (the migration time becomes negative - bottom panel). - i.e. planets migrateoulward in (he region with the positive temperature gradient.," We show in Figure \ref{fig3} (top) that this radial thermal inversion causes the migration rate to be positive (the migration time becomes negative - bottom panel), - i.e. planets migrate in the region with the positive temperature gradient." + The physical explanation of (his behavior is Chat the increasing; function of disk temperature changes the disk's pressure distribution which in turn causes the position of the outer Lindblad resonances to be further [rom the planet than the inner ones (Artvmowiez1993)., The physical explanation of this behavior is that the increasing function of disk temperature changes the disk's pressure distribution which in turn causes the position of the outer Lindblad resonances to be further from the planet than the inner ones \citep{a93}. +. This results in outer torques that are much weaker than the inner Figure 4.2. (bottom) also shows that the planets verv slowly enter the region of torque reversal - as seen bv (he strong positive spike in (he migration time.," This results in outer torques that are much weaker than the inner Figure \ref{fig3} (bottom) also shows that the planets very slowly enter the region of torque reversal - as seen by the strong positive ""spike"" in the migration time." + These regions correspond (o radii al which dTdr~0 (see Figure 2)) al (he inner and outer resonances. making the torque difference between (hem very small.," These regions correspond to radii at which $dT/dr\simeq 0$ (see Figure \ref{fig2}) ) at the inner and outer resonances, making the torque difference between them very small." + We emphasize that the positive temperature gradient arising [rom (he wall-like dust structure is achieved [or the case of a finite transition region.Ar«10h. in the value of turbulent a (IIP10).," We emphasize that the positive temperature gradient arising from the wall-like dust structure is achieved for the case of a finite transition region,$\bigtriangleup r \leq 10 h$, in the value of turbulent $\alpha$ (HP10)." + Although. for simplicity. we adopt a sharp spatial transition [rom the active to the dead zone in this Letter (Ar= 0). our above results are valid for the case of Arcfh. since the positions of Lindblad resonances are (vpically offset [rom (he planets bv 25/3 (Artvimowiez1993).," Although, for simplicity, we adopt a sharp spatial transition from the active to the dead zone in this Letter $\bigtriangleup r=0$ ), our above results are valid for the case of $\bigtriangleup r \simeq h$, since the positions of Lindblad resonances are typically offset from the planets by $2h/3$ \citep{a93}." +. In addition. it is interesting that the migration timescale of the M star svstem is similar to that of CTTS (~10!—10 vears) for the other two cases (well mixed. dust settling).," In addition, it is interesting that the migration timescale of the M star system is similar to that of CTTS $\sim 10^4-10^5$ years) for the other two cases (well mixed, dust settling)." + This is because the tidal torqueis scaled by X(h/r) ?., This is because the tidal torqueis scaled by $\Sigma(h/r)^{-2}$ . + A more detailed discussion of them is presented in Iasegawa&ποναί (2010b).., A more detailed discussion of them is presented in \citet{hp10}. . +where / is the moment of inertia of the pulsar. /? is its radius. P is the spin period. /? is the rate ofchange of the spin period. à is the angle between the magnetic axis and the spin axis of the pulsar and ὁ is the speed of light.,"where $I$ is the moment of inertia of the pulsar, $R$ is its radius, $P$ is the spin period, $\dot{P}$ is the rate of change of the spin period, $\alpha$ is the angle between the magnetic axis and the spin axis of the pulsar and $c$ is the speed of light." +" Dp, is defined as ", $B_{fac}$ is defined as $\sqrt{\frac{3c^3}{8 \pi^2}\frac{I}{R^6 sin^2 \alpha}}$. +Due to the lack of knowledge about the actual value of a. it isusually assumed to be 90°., Due to the lack of knowledge about the actual value of $\alpha$ it isusually assumed to be $90^{~\circ}$ . + Moreover. / is taken as 107emcm. 2 as 10km.," Moreover, $I$ is taken as $10^{45}~ \rm{gm~ cm^2}$, $R$ as $10~\rm{km}$." + Then {δει becomes 3.2 ., Then $B_{fac}$ becomes $3.2 \times 10^{19}~ \rm { gm^{1/2} cm^{-1/2} sec^ {-3/2}}$ . + Sooo Eqn. CL) , So Eqn. \ref{eq:def_mag}) ) +can be rewritten. as: Eqn. (2)), can be rewritten as: Eqn. \ref{eq:conv_mag}) ) +" is commonly used to calculate D, by putting the values of P? in seconds and J? as dimensionless.", is commonly used to calculate $B_s$ by putting the values of $P$ in seconds and $\dot{P}$ as dimensionless. + The problem with the above simplification is that. for a fixed value of the pulsar mass (AJ). one gets widely different values ofI and / by using different EsoS for the matter.," The problem with the above simplification is that, for a fixed value of the pulsar mass $M$ ), one gets widely different values of$R$ and $I$ by using different EsoS for the matter." + Moreover all the pulsars do not have the same value of A., Moreover all the pulsars do not have the same value of $M$. + In short. 7? and £ can be different from their canonical values 10kin and 107gmenm? respectively depending upon 1 and the choice of the EoS. a can also have any value between ϱ and 907," In short, $R$ and $I$ can be different from their canonical values $10~\rm{km}$ and $10^{45}~ \rm{gm~ cm^2}$ respectively depending upon $M$ and the choice of the EoS. $\alpha$ can also have any value between $0^{~\circ}$ and $90^{~\circ}$." +" a result the value of By... will be different from 3.2«.10emtOS7embesec77,the", As a result the value of $B_{fac}$ will be different from $3.2 \times 10^{19}~ \rm { gm^{1/2} cm^{-1/2} sec^ {-3/2}}$. +" For this reason. it is worthwhile to study the dependence of value of 5, on AJ. à and EsoS. In the next section. we study this dependence using Eqn. CH)."," For this reason, it is worthwhile to study the dependence of the value of $B_s$ on $M$, $\alpha$ and EsoS. In the next section, we study this dependence using Eqn. \ref{eq:def_mag}) )." + Among numerous EsoS for the dense matter. we have chosen five - three for the neutron matter (namely EoS BPALI2. EoS APR and EoS MFTI7 which are also used by Bejger efa£. 2005) and two for the strange quark matter.," Among numerous EsoS for the dense matter, we have chosen five - three for the neutron matter (namely EoS BPAL12, EoS APR and EoS MFT17 which are also used by Bejger $et~al.$ 2005) and two for the strange quark matter." + sone in the latter group is is EoS A (SSA) of Bagehi οἱαἲ. vol(2006) and the other one is from Bag model (BAGI) with parameters as : 0.17 where £ is the Bag parameter. mi. my. nm; are masses of s. u and d quarks respectively.," One in the latter group is is EoS A (SSA) of Bagchi $et~al.$ (2006) and the other one is from Bag model (BAG1) with model parameters as : $B = 60.0~{\rm MeV / fm^3},~m_s = 150.0~ {\rm MeV},~ m_u = m_d = 0,~\alpha_c = 0.17$ where $B$ is the Bag parameter, $m_s$, $m_u$, $m_d$ are masses of s, u and d quarks respectively." + These five EsoS are of widely different stiffness and so they give significantly different plots in the AJ// plane., These five EsoS are of widely different stiffness and so they give significantly different plots in the $M - R$ plane. + In Fig. |.," In Fig. \ref{fig:mass_all_eos}," + we show the AJ)/? and ALL plots for all of these EsoS. It is worthwhile to mention here that MFT!7 strongly overestimate and BPALI2 strongly underestimate the nuclear matter incompressibilitv., we show the $M - R$ and $M - I$ plots for all of these EsoS. It is worthwhile to mention here that MFT17 strongly overestimate and BPAL12 strongly underestimate the nuclear matter incompressibility. + APR seems to be a better description of the nuclear matter., APR seems to be a better description of the nuclear matter. + We have calculated the moment of inertia using the formalism derived by Kalogera e£a7. (1999) within Hartle's approximations., We have calculated the moment of inertia using the formalism derived by Kalogera $et~al.$ (1999) within Hartle's approximations. +" We have found that even for an angular rotational frequency (Q) as high as 5000 sec. Bj, obtained by this method differs maximum upto LOM from that obtained by using proper codes for rotating stars (RNS code. written by Nikolaos Stergioulas and available at httpz/Awww.gravity.phys.uwm.edu/rns/) for all of the EsoS mentioned earlier: the lower is the value of ©. the smaller is the difference."," We have found that even for an angular rotational frequency $\Omega$ ) as high as 5000 $\rm sec^{-1}$, $B_{fac}$ obtained by this method differs maximum upto $ \sim 10 \%$ from that obtained by using proper codes for rotating stars (RNS code, written by Nikolaos Stergioulas and available at http://www.gravity.phys.uwm.edu/rns/) for all of the EsoS mentioned earlier; the lower is the value of $\Omega$, the smaller is the difference." +" So for radio pulsars. the use of Hartle's approximations is well justified as the fastest pulsar yet discovered is PSR JI748 with ©=4500sec! (see ATNF pulsar catalog at httpzAwww.atnf.esiro.au/research/pulsar/psreat/).We have plotted the variation of Dy, with pulsar masses for different EsoS (Figs."," So for radio pulsars, the use of Hartle's approximations is well justified as the fastest pulsar yet discovered is PSR J1748 $-$ with $\Omega~=~ 4500~\rm sec^{-1}$ (see ATNF pulsar catalog at http://www.atnf.csiro.au/research/pulsar/psrcat/).We have plotted the variation of $B_{fac}$ with pulsar masses for different EsoS (Figs." + 2. and 3») and different values of a (17. 30 60 and 90°).," \ref{fig:bfac_all_alpha_strange} and \ref{fig:bfac_all_alpha_neutron}) ) and different values of $\alpha$ $1^{\circ}$, $30^{\circ}$ , $60^{\circ}$ and $90^{\circ}$ )." + We find that for the strange quark matter EsoS. the value of {λος is usually large.," We find that for the strange quark matter EsoS, the value of $B_{fac}$ is usually large." +" For EoS SSA. the value of Dy, is greater than 3.2.107 values of à and Ad."," For EoS SSA, the value of $B_{fac}$ is greater than $3.2 \times 10^{19}$ for any values of $\alpha$ and $M$." +" For EoS BAG]. εως is gre: than for:nDany107"" fora«90 Lata= 907. Dp=32' te"," For EoS BAG1, $B_{fac}$ is greater than $3.2 \times 10^{19}$ for $\alpha<90^{\circ}$, at $\alpha=90^{\circ}$ , $B_{fac} \simeq 3.2 \times 10^{19}$." +"rFor the neutron matter EsoS. £y, is greater than 3.2.102° for very low values of à. e.g. a=1 (for any values of AZ)."," For the neutron matter EsoS, $B_{fac}$ is greater than $3.2 \times 10^{19}$ for very low values of $\alpha$, $e.g.$ $\alpha=1^{\circ}$ (for any values of $M$ )." +" But fora.z307. Dj, can be smallerthan 3.2.10+? depending upon the values of A7."," But for $\alpha \geqslant 30^{\circ}$, $B_{fac}$ can be smaller than $3.2 \times 10^{19}$ depending upon the values of $M$." +" For EoS BPALI2. at à=O30. Brae3.2LO PAL 3.2 \times 10^{19}$ for any values of $M$, at $\alpha=60^{\circ}$, $B_{fac}<3.2 \times 10^{19}$ if $M<1.6 M_{\odot}$, at $\alpha=90^{\circ}$, $B_{fac}<3.2 \times 10^{19}$ if $M<1.8 M_{\odot}$." +" For EoS MFT17. ata—3 DfaeM.3}10? if AZ<15A.. but fora=60' anda— 907. -Dj,c3.2107"" for any values of AJ."," For EoS MFT17, at $\alpha=30^{\circ}$, $B_{fac}<3.2 \times 10^{19}$ if $M<1.5 M_{\odot}$, but for $\alpha=60^{\circ}$ and $\alpha=90^{\circ}$ , $B_{fac}<3.2 \times 10^{19}$ for any values of $M$." +" Fora Ad and fixed EoS. ἐς is minimum (23pac. iin) fora=90. giving the minimum possible value of the D, (Benin) i£. AZ and that EoS. From Figs."," For a fixed $M$ and fixed EoS, $B_{fac}$ is minimum $B_{fac,~min}$ ) for $\alpha=90^{\circ}$, giving the minimum possible value of the $B_s$ $B_{s,~min}$ ) for that $M$ and that EoS. From Figs." +12 2. and 3.. it is clear that P757for5the2jB.μυ”ic Al⋅⋅ giving DiuNdi," \ref{fig:bfac_all_alpha_strange} and \ref{fig:bfac_all_alpha_neutron},, it is clear that $B_{fac,~min}^{SSA} > B_{fac,~min}^{BAG1}>B_{fac,~min}^{BPAL12}>B_{fac,~min}^{APR}>B_{fac,~min}^{MFT17}$ for $M \leqslant 1.5~M_{\odot}$ giving $B_{s,~min}^{BAG1}>B_{s,~min}^{BPAL12}>B_{s,~min}^{APR}>B_{s,~min}^{MFT17}$ ." +"l ange in 2, for a change Aa ina keeping AZ fixed and 4£,|,, denote thechange in D, for a change AM in M keeping à ", Let $\Delta B_{s}\vert_{M} $ denote the change in $B_s$ for a change $\Delta \alpha$ in $\alpha$ keeping $M$ fixed and $\Delta B_{s}\vert_{\alpha} $ denote thechange in $B_s$ for a change $\Delta M$ in $M$ keeping $\alpha$ fixed. +"Then for À/ being in therangeSEE of OSL5 AL. ABIESAB.sed.NUsAB gor quues of AM for each vinSSA‘equalvaluesnuofBoforPALISs for Aa each EoS. study clearly reveals that in addition to constraining the EoS. the knowledgeof a is also necessary to know the actual value of D, for apulsar of known Al. /?andD."," Then for $M$ being in therange of $0.8 - 1.5~M_{\odot}$ , $\Delta B_{s}\vert_{\alpha}^{BPAL12}> \Delta B_{s}\vert_{\alpha}^{SSA} >\Delta B_{s}\vert_{\alpha}^{APR}>\Delta B_{s}\vert_{\alpha}^{MFT17}>\Delta B_{s}\vert_{\alpha}^{BAG1} $ for equal values of $\Delta M$ for each EoS;$\Delta B_{s}\vert_{M}^{SSA} > \Delta B_{s}\vert_{M}^{BAG1}> \Delta B_{s}\vert_{M}^{BPAL12}> \Delta B_{s}\vert_{M}^{APR}>\Delta B_{s}\vert_{M}^{MFT17} $ for equal values of $\Delta \alpha$ for each EoS. Our study clearly reveals that in addition to constraining the EoS, the knowledgeof $\alpha$ is also necessary to know the actual value of $B_s$ for apulsar of known$M$ , $P$and$\dot{P}$." + Sometimes it is possibleto determine a observationally., Sometimes it is possibleto determine $\alpha$ observationally. + As an example. for PSR 3039A. Jenet Ransom (2004) have foundà to be either 1.6x1.3| or 142x2 by careful study of the pulse profile while," As an example, for PSR $-$ 3039A, Jenet Ransom (2004) have found$\alpha$ to be either $1.6^{~\circ} \pm 1.3^{~\circ}$ or $14^{~\circ} \pm 2^{~\circ}$ by careful study of the pulse profile while" +available from astrometry sits at a comfortable level.,available from astrometry sits at a comfortable level. + Spectroscopic accuracy of S2 at θαs| is not vet available. but seenis plausible with future observations.," Spectroscopic accuracy of S2 at $10\kms$ is not yet available, but seems plausible with future observations." + For the late-tvpe star $35. which has a more favorable spectrum. a fit error of 10kinx| has been achieved (Cullessenetal.2009a).," For the late-type star S35, which has a more favorable spectrum, a fit error of $10\kms$ has been achieved \citep{gillessen}." +. We remark that auv systematic errors that do not change between observations are harmlessly absorbed iuto my., We remark that any systematic errors that do not change between observations are harmlessly absorbed into $\nu_0$. + Naturally. in detecting relativistic effects ou S Stars. data during pericenter passage are of greatest value.," Naturally, in detecting relativistic effects on S Stars, data during pericenter passage are of greatest value." + With iustruuentation currently available. au observation progran conceutrated over two mouths during S28 next periceuter passage (2018) will prove to be sufficient as a test for the Eiustei Equivalence Principle.," With instrumentation currently available, an observation program concentrated over two months during S2's next pericenter passage (2018) will prove to be sufficient as a test for the Einstein Equivalence Principle." + Figure (2)) argues that a small handful of spectral iicasuremenuts of S2 at θα around periccuter inuplv a lo accuracy on à of 0.3., Figure \ref{fig:recovering_alpha}) ) argues that a small handful of spectral measurements of S2 at $10\kms$ around pericenter imply a $1\sigma$ accuracy on $\alpha$ of $\sim 0.3$. + The approach we have taken above focuses on the essentials., The approach we have taken above focuses on the essentials. + The degeneracy between a. Fo and α for spectroscopy has been lifted by using astrometiy only to provide £.," The degeneracy between $\alpha$, $I$ and $a$ for spectroscopy has been lifted by using astrometry only to provide $I$." + Iu practice. all the parameters are fitted simultaneously to both astrometry aud spectroscopy.," In practice, all the parameters are fitted simultaneously to both astrometry and spectroscopy." + We have doue simulations to verify that when this is done. the degeneracy is implicitly broken bv the mechanism lughlielted in this Letter.," We have done simulations to verify that when this is done, the degeneracy is implicitly broken by the mechanism highlighted in this Letter." + Relativistic effects can be expected to become oeicreasinely oeuportant as corrections in other astroplivsics relating to the S stars., Relativistic effects can be expected to become increasingly important as corrections in other astrophysics relating to the S stars. + Three areas where this can be expected are the following., Three areas where this can be expected are the following. + Tn exploring 1 aud 2. oue cannot easily avoid relativity siuplv by considering stars with larger orbit sizes (3). as perturbations due to the enclosed mass become a xoblemi.," In exploring 1 and 2, one cannot easily avoid relativity simply by considering stars with larger orbit sizes (3), as perturbations due to the enclosed mass become a problem." + With the accuracy regime that spectroscopy will cuter in the the comine decade. oue of two types of perturbations to the redshift must be faced: cither hose from the extended mass distribution. for S Stars with large orbits. or. those perturbations from relativity. or stnaller orbit S Stars.," With the accuracy regime that spectroscopy will enter in the the coming decade, one of two types of perturbations to the redshift must be faced: either those from the extended mass distribution, for S Stars with large orbits, or, those perturbations from relativity, for smaller orbit S Stars." + While effects due to relativity are well understood iu principle. and can be casily reated. the constitution of the exteuded system is poorly understood. and expects a iore erueliug troatineut.," While effects due to relativity are well understood in principle, and can be easily treated, the constitution of the extended system is poorly understood, and expects a more grueling treatment." + Iu he coniuge decade. spectroscopic S Star accuracy at ~Olas+ is expected to be available. aud the discovery of stars closer in to the black hole is auticipated.," In the coming decade, spectroscopic S Star accuracy at $\sim 10\kms$ is expected to be available, and the discovery of stars closer in to the black hole is anticipated." + Iu ight of these prospects. so that coustraiuts on these quantities nav be miproved. relativistic perturbations. while interesting in their own right. cau no longer be ignored aud iust be faced.," In light of these prospects, so that constraints on these quantities may be improved, relativistic perturbations, while interesting in their own right, can no longer be ignored and must be faced." +" The authors thank ο, Callessen. €. Εν R. Ellis. aud J.-P. Uzan for discussion aud conmunuaeuts."," The authors thank S. Gillessen, G. F. R. Ellis, and J.-P. Uzan for discussion and comments." +orthogonal to the earlier value.,orthogonal to the earlier value. + This lrequeney-depencent EVPA separation has eontinued to the present time., This frequency-dependent EVPA separation has continued to the present time. + We attribute (his behavior to a change in the opacity in the em-band emitting region indicated by the evolution of the spectrum of the total Εικ density: the source has become transparent. so that al 14.5 GlIIz we are looking further into the jet al a region which has a substantially different magnetic field alignment.," We attribute this behavior to a change in the opacity in the cm-band emitting region indicated by the evolution of the spectrum of the total flux density: the source has become transparent, so that at 14.5 GHz we are looking further into the jet at a region which has a substantially different magnetic field alignment." + The 43 GlIIz map of Lister.(2001a) at epoch 1999.26 shows four jet components. significant. polarization from both the core and jet. and a total polarization of4.," The 43 GHz map of \citet{lis01a} at epoch 1999.26 shows four jet components, significant polarization from both the core and jet, and a total polarization of." +9%.. While our data represent a sum of contributions rom all components and extended structure. (he change we see is consistent with increased enission Irom the strongest VLBI component.," While our data represent a sum of contributions from all components and extended structure, the change we see is consistent with increased emission from the strongest VLBI component." + Both the millimeter VLBA data and single dish centimeter-band data suggest (hat at this epoch the predominant magnetic [field is aligned nearly perpendicular to the flow direction of ~111., Both the millimeter VLBA data and single dish centimeter-band data suggest that at this epoch the predominant magnetic field is aligned nearly perpendicular to the flow direction of $\sim 111^{\circ}$. + The behavior illustrates that significant changes in the characteristic behavior of (he polarization. as well as of the flux. ean occur with time scales of a decade and that shorter time windows (even a decade lone) max not adequately sample the range of variability.," The behavior illustrates that significant changes in the characteristic behavior of the polarization, as well as of the flux, can occur with time scales of a decade and that shorter time windows (even a decade long) may not adequately sample the range of variability." + Ii (he case of this source. a sequence of well-sampled high Irequency VLDP maps would have been helpful in unraveling the complicated behavior we find.," In the case of this source, a sequence of well-sampled high frequency VLBP maps would have been helpful in unraveling the complicated behavior we find." + In Paper I we compared the preferred intrinsic EVPAs at 14.5 and 4.8 Gllz with the flow direction (8) in the jet indicated by (he morphology apparent in the 5 GlIz Pearson-Reaclheac maps (PearsonandReacdheacd1983)., In Paper I we compared the preferred intrinsic EVPAs at 14.5 and 4.8 GHz with the flow direction $\theta$ ) in the jet indicated by the morphology apparent in the 5 GHz Pearson-Readhead maps \citep{per88}. +. We found that at 14.5 GlIz there was no difference in distribution of |[EVPA-0| between QSOs and BL Lacs. but at 4.8 GHz the values of |EVPA- 0| clustered near 0-20° for the BL Lacs ancl were preferentially near 90° for the QSOs: the clisparity between the shape of the distribution of orientation cilference for each optical class based on our data at 4.8 Gllz suggested intrinsic differences in the physical properties of the objects such that the magnetic field in the emitting region is prelerentially along the flow direction in QSOs and perpendicular to it in DL Lacs.," We found that at 14.5 GHz there was no difference in distribution of $\vert$ $\theta \vert$ between QSOs and BL Lacs, but at 4.8 GHz the values of $\vert$ $\theta\vert$ clustered near $^\circ$ for the BL Lacs and were preferentially near $^\circ$ for the QSOs; the disparity between the shape of the distribution of orientation difference for each optical class based on our data at 4.8 GHz suggested intrinsic differences in the physical properties of the objects such that the magnetic field in the emitting region is preferentially along the flow direction in QSOs and perpendicular to it in BL Lacs." + This result supports the class differences in EVPA orientation identified in Cawthorneetal.(1993).. but differs from it in concept since that result is based on single-epoch observations of jet (aud core) components wilh tvpical Lletimes of order a decade (Gabuzdaοἱal.1994). while our result is based on longterm data. which may include many different evolving components (which might have exhibited quite different EVPAs from event to event). and with the major contribution to the integrated polarization mostlikely arising in (he core and innermost jel components al 14.5 GHz (IIomanοἱal... 2002)..," This result supports the class differences in EVPA orientation identified in \citet{caw93}, but differs from it in concept since that result is based on single-epoch observations of jet (and core) components with typical lifetimes of order a decade \citep{gab94} + while our result is based on longterm data, which may include many different evolving components (which might have exhibited quite different EVPAs from event to event), and with the major contribution to the integrated polarization mostlikely arising in the core and innermost jet components at 14.5 GHz \citep{hom02}. ." +field.,field. + It should be noticed that the spectra of almost all the Fermi-LAT pulsars including MSPs. except very voung pulsars like the Crab pulsar. can be explained in terms of CR mechanism.," It should be noticed that the spectra of almost all the Fermi-LAT pulsars including MSPs, except very young pulsars like the Crab pulsar, can be explained in terms of CR mechanism." + Other models even (μον can fit the Fermi data of globular clusters equally well but they cannot be accepted as alternative models unless they have other new predictions and are supported by observations., Other models even they can fit the Fermi data of globular clusters equally well but they cannot be accepted as alternative models unless they have other new predictions and are supported by observations. +" Nevertheless with all these observational hints. we propose that there may be an alternative/aclditional emission mechanism to produce the observed gamma-ravs detected by Fermi-LAT aud explore the new predictions [rom this model,"," Nevertheless with all these observational hints, we propose that there may be an alternative/additional emission mechanism to produce the observed gamma-rays detected by Fermi-LAT and explore the new predictions from this model." + Becnarek Sitarek (2007) analvzecl gamma-ray emission of electrons accelerated al shock waves originated in collisions of the pulsar winds and/or inside (he pulsar maegnetospheres when eamunaravs are generated by (he inverse Compton (IC) scattering of ultra-relativistic electrons of relic ancl stellar photons., Bednarek Sitarek (2007) analyzed gamma-ray emission of electrons accelerated at shock waves originated in collisions of the pulsar winds and/or inside the pulsar magnetospheres when gamma-rays are generated by the inverse Compton (IC) scattering of ultra-relativistic electrons of relic and stellar photons. + Both of these models can give reasonable explanation for the eamuna-rav emission from 47 Tuc., Both of these models can give reasonable explanation for the gamma-ray emission from 47 Tuc. + It should be noticed {hat both of these models predict that gamma-rays are emitted [rom core region of GC. ie. <1 pe. where most MSPs are located.," It should be noticed that both of these models predict that gamma-rays are emitted from core region of GC, i.e. $<1$ pc, where most MSPs are located." + The kev difference between these (vo classes of models is that (he inverse Compton model predicts the existence of very hieh energy gaanma-rays. which can be detected by MAGIC: and IIESS.," The key difference between these two classes of models is that the inverse Compton model predicts the existence of very high energy gamma-rays, which can be detected by MAGIC and HESS." + In this paper we also study the inverse Compton scattering between (he relativistic electrons/positron in the pulsar wind and the background soft photons., In this paper we also study the inverse Compton scattering between the relativistic electrons/positron in the pulsar wind and the background soft photons. + We adopt the pulsar wind model proposed by Cheng. Taam Wane (2004. 2006).," We adopt the pulsar wind model proposed by Cheng, Taam Wang (2004, 2006)." + To generalize the soft photon field in our investigation. in addition to relie photons aud the star light photons in (he GC. we include (he background soft photons [rom the galactic disk including the inlrared photons and star lisht photons of the galactic disk.," To generalize the soft photon field in our investigation, in addition to relic photons and the star light photons in the GC, we include the background soft photons from the galactic disk including the infrared photons and star light photons of the galactic disk." + Our caleulations do not restrict (he inverse Compton scattering only in (he core instead we extend our ealeulation to several hundred pe from the core of GC., Our calculations do not restrict the inverse Compton scattering only in the core instead we extend our calculation to several hundred pc from the core of GC. + By the fitting the observed data of 47 Tuc and Terzan, By the fitting the observed data of 47 Tuc and Terzan +"Accreting black holes are thought to be the enugimes powering nost of the emissiou from active galactic uuclei (AGN) aud some X-ray binaries (δν),",Accreting black holes are thought to be the engines powering most of the emission from active galactic nuclei (AGN) and some X-ray binaries (XRBs). + Associated relativistic jets also contribute significantly to the overall spectrum over a wide range o ewwaveleneths.," Associated relativistic jets also contribute significantly to the overall spectrum, over a wide range of wavelengths." +" The current accretion paradieni is based ou the early success of standard. optically thick accretion disk models (Shalkura&Suuvaev1973) which correctly predicted the soft N-ray enission im stellar mass black holes CARBs) aud the ""big due bump” in quasars (Sandersetal.1989:Sun&Malla1989) and other AGN by scaling mass aud accretion rate."," The current accretion paradigm is based on the early success of standard, optically thick accretion disk models \cite{ShakuraSunyaev1973} which correctly predicted the soft X-ray emission in stellar mass black holes (XRBs) and the “big blue bump” in quasars \cite{SandersPhinneyNeugebauer1989,SunMalkan1989} and other AGN by scaling mass and accretion rate." + Scaling laws for the radio ciission of jet cores aud lobes inve also cen developed (Falcke&Dieriiauu1995:Ikaiser&Alexander1997:Heinz2002) and successtullv applied o NRBs and AGN (Faleke.Malkau.&Bicrinann1995:Faleke&Biermann1996: 1999).," Scaling laws for the radio emission of jet cores and lobes have also been developed \cite{FalckeBiermann1995,KaiserAlexander1997,Heinz2002} and successfully applied to XRBs and AGN \cite{FalckeMalkanBiermann1995,FalckeBiermann1996,FalckeBiermann1999}." +. The most miportant parameters of accreting black roles are probably the mass and the accretion rate. voth of which can vary over many orders of magnitudo.," The most important parameters of accreting black holes are probably the mass and the accretion rate, both of which can vary over many orders of magnitude." + Additional parameters which likely impact the observable characteristics of black holes are the spin and the inclination anele of thei spin axes., Additional parameters which likely impact the observable characteristics of black holes are the spin and the inclination angle of their spin axes. + Iucliuation-based unified schemes of ACN inerge appareutly different objects based on the augle between the spin axis aud the line of sieht (see e.e.. Antonucci1993:Urry.&Padovani 1995)).," Inclination-based unified schemes of AGN merge apparently different objects based on the angle between the spin axis and the line of sight (see e.g., \citeNP{Antonucci1993,UrryPadovani1995}) )." + The success of this scheme supports the evidence for dependent obscuration and relativistic beamine., The success of this scheme supports the evidence for angle-dependent obscuration and relativistic beaming. + However. the exact effect that changes iu the accretion rate have on the appearance of their associated black holes systems is a matter of ongoing debate.," However, the exact effect that changes in the accretion rate have on the appearance of their associated black holes systems is a matter of ongoing debate." +" À eood understaudie of this is crucial for modeling the cosmic evolution of black holes aud for cdiscutaneling the different sonrcee Classes,", A good understanding of this is crucial for modeling the cosmic evolution of black holes and for disentangling the different source classes. + A uuuber of recent results suggest that the trausitio- YOU a Hel acerction rate black hole o a low accretio- rate one ds uot sinooth. but rather accompanied by a “phase ransition.," A number of recent results suggest that the transition from a high accretion rate black hole to a low accretion rate one is not smooth, but rather accompanied by a “phase transition”." + Iu the low-power pase. the optically hick disk Clussion is either (Louieved by clissio- YOU ali optically thin corona. comletely reduced. teCc a radiaively inefficient inflow. or is runcated and au optically thin iuner raciatively inefiicicut flow exists closer o the compact object (see Poutanen1998. for a review of the various models).," In the low-power phase, the optically thick disk emission is either dominated by emission from an optically thin corona, completely reduced to a radiatively inefficient inflow, or is truncated and an optically thin inner radiatively inefficient flow exists closer to the compact object (see \citeNP{Poutanen1998} for a review of the various models)." +" For NRBs Esin.MeClintock.&Naravan(1997) estimate that this transition occurs once he accretion rate for a black hole of mas A‘, drops o less than a critical value (~10% of the EddiugOl accretion rate. Maa2(AL/I05AL.)Ms for Mgaaο "," For XRBs \citeN{EsinMcClintockNarayan1997} estimate that this transition occurs once the accretion rate for a black hole of mas $M_\bullet$ drops to less than a critical value $\sim$ of the Eddington accretion rate, $\dot M_{\rm Edd}\simeq2\times(M_\bullet/10^8 +M_\odot)\;M_\odot\,{\rm yr}^{-1}$ for $\dot M_{\rm Edd}=L_{\rm +Edd}/0.1 c^2$ )." +More recent work suggests that his transition could already occur around Egg (Maccaroue2003)... and that there is a lysteresis iu the critical accretion rate value depending on which direction he transition is goiug alone (Alaccarone&Coppi2003).," More recent work suggests that this transition could already occur around $\dot L_{\rm Edd}$ \cite{Maccarone2003}, and that there is a hysteresis in the critical accretion rate value depending on which direction the transition is going along \cite{MaccaroneCoppi2003}." +. Regardless of the exact details. a crucial point for this operis a phasc-transition as a function of black hole mass and accretion power.," Regardless of the exact details, a crucial point for this paper is a phase-transition as a function of black hole mass and accretion power." + We have previously sugecsted that the coutributiou of jets and outflows on the spectral cucrey distribution (SED) of black holes can be significant iu supermassive as well as stellar nass black roles (Faleke&/MarkoffMarkoff.&Faleke2002) and that the jet coutribution may in fact donünate the disk cimission in a JDAF a jet-domunated accretion flow.," We have previously suggested that the contribution of jets and outflows on the spectral energy distribution (SED) of black holes can be significant in supermassive as well as stellar mass black holes \cite{FalckeMarkoff2000,MarkoffFalckeFender2001,Fender2001,YuanMarkoffFalcke2002} and that the jet contribution may in fact dominate the disk emission in a JDAF – a jet-dominated accretion flow." + Jets are Inherentlv broad-baud. siuce they remain sclfsinular OVOY lay orders of magnitude im spatial scale aud produce nou-thermal particle distributions ranging over many orders of magnitude in cucrey.," Jets are inherently broad-band, since they remain self-similar over many orders of magnitude in spatial scale and produce non-thermal particle distributions ranging over many orders of magnitude in energy." + For this reason they should always be considered as potential coutributors at every wavelength., For this reason they should always be considered as potential contributors at every wavelength. + This concept of jet domination has now been empiricallv demonstrated for XRDs. where below παc10+ the kinetic euerev output through rachatively inefiicicut jets," This concept of jet domination has now been empirically demonstrated for XRBs, where below $L_{\rm Edd}\approx 10^{-4}$ the kinetic energy output through radiatively inefficient jets" +Fundamental Plane. any deviation of individual clusters from the FP may be a direct. rellection of its recent dynamical evolution.,"Fundamental Plane, any deviation of individual clusters from the FP may be a direct reflection of its recent dynamical evolution." + This would be true if the thickness of the plane would be entirely due to the merger history of the clusters., This would be true if the thickness of the plane would be entirely due to the merger history of the clusters. + I is certainly a viable implication of our conclusion that the Fundamental Plane's definition | the average plane ofa Large sample of clusters is nearly unassailable while we find strong Luctuations and deviations from the average FP in small samples of actively evolving clusters., It is certainly a viable implication of our conclusion that the Fundamental Plane's definition – the average plane of a large sample of clusters – is nearly unassailable while we find strong fluctuations and deviations from the average FP in small samples of actively evolving clusters. + In practice. it might mean that one could take samples of clusters in dillerent. redshift bands anc reliably average them in each band to use the resulting Fundamental Plane to study redshift evolution of observed. samples.," In practice, it might mean that one could take samples of clusters in different redshift bands and reliably average them in each band to use the resulting Fundamental Plane to study redshift evolution of observed samples." + Lt would also mean that within cach redshift band vou know which ones have Πας active lives., It would also mean that within each redshift band you know which ones have had active lives. + The Fundamental Plane is a direct. rellection of the viria heorem which. uncer particular assumptions. relates the averaged. velocity dispersion and radius of a system directhy oits mass.," The Fundamental Plane is a direct reflection of the virial theorem which, under particular assumptions, relates the averaged velocity dispersion and radius of a system directly to its mass." +" All ""virialised"" objects will lie on a plane defines in the space of those three variables.", All “virialised” objects will lie on a plane defined in the space of those three variables. + Phere is not even any [reeciom in the parameters for that plane: its slope ane ocation are fixed for all virialised objects., There is not even any freedom in the parameters for that plane: its slope and location are fixed for all virialised objects. + There are complications when assigning data to a Fundamental Plane., There are complications when assigning data to a Fundamental Plane. + Firstly. in its simplest form. the viria heoreni assumes that the νακο objects are. isolate spherical systems and. importantly. that they are stationary.," Firstly, in its simplest form, the virial theorem assumes that the virialised objects are isolated spherical systems and, importantly, that they are stationary." + The systems we study are not spherical and. they are certainly not stationary: they are generally in a state of dynamical evolution., The systems we study are not spherical and they are certainly not stationary: they are generally in a state of dynamical evolution. + The possible exception to this mieh be the largest most isolated. svstenis., The possible exception to this might be the largest most isolated systems. + Secondly. observec data does not have direct. knowledge of the svstenr mass except through interpreting the light that is observed.," Secondly, observed data does not have direct knowledge of the system mass except through interpreting the light that is observed." + The universality of the Fundamental Plane allows us to turn the problem around and. determine the dependence of light on mass in order that svstems should fit on the Fundamental Plane., The universality of the Fundamental Plane allows us to turn the problem around and determine the dependence of light on mass in order that systems should fit on the Fundamental Plane. + The simplest approach to this is to assume that the mass to light ratio in the observed waveband. is. directly related to mass., The simplest approach to this is to assume that the mass to light ratio in the observed waveband is directly related to mass. + There are further issues., There are further issues. +" For example. what do we mean when we refer to. ""averages"" of quantities?"," For example, what do we mean when we refer to “averages” of quantities?" + Using a dillerent. averaging process vields a diferent Fundamental Plane., Using a different averaging process yields a different Fundamental Plane. + There is also the fact that astrophysical svstenis are observed only in projection., There is also the fact that astrophysical systems are observed only in projection. + Llaving said that. we can express the Virial Theorem in terms of the variable we have used. here to describe the Fundamental Plane.," Having said that, we can express the Virial Theorem in terms of the variable we have used here to describe the Fundamental Plane." +" With the notation that a virialised system of mass AZ has a velocity dispersion. V. hall mass radius rj, and harmonic radius rj, we have. up to normalising constants: where X ds the projected. (surface) mass density."," With the notation that a virialised system of mass $M$ has a velocity dispersion $V$, half mass radius $r_{half}$ and harmonic radius $r_h$ we have, up to normalising constants: where $\Sigma$ is the projected (surface) mass density." + Eliminating AJ from these ancl taking logs viclds an expression for the Fundamental Plane: where we have transformed the surface mass density ὃν into logarithmic astro-units via We have explicith written equation (33)) in such a way as to expose the different roles of harmonic ancl hall-mass (ecometric) radii., Eliminating $M$ from these and taking logs yields an expression for the Fundamental Plane: where we have transformed the surface mass density $\Sigma$ into logarithmic astro-units via We have explicitly written equation \ref{eq:virial}) ) in such a way as to expose the different roles of harmonic and half-mass (geometric) radii. + The relationship between these radii in our models is illustrated in Fig., The relationship between these radii in our models is illustrated in Fig. + 2 and in Fig. 17..," \ref{fig:radiusrel} and in Fig. \ref{fig:mergaccrad}," + the latter dilferentiating between merging ancl quiescentIy. accreting halos., the latter differentiating between merging and quiescently accreting halos. + lt is important to understand. why the coellicients of he model Funcamental Plane might cliffer from the expectations based on the use of the virial theorem., It is important to understand why the coefficients of the model Fundamental Plane might differ from the expectations based on the use of the virial theorem. + Luminosity is not involved. here so we cannot appeal ο a varving mass-to-light ratio., Luminosity is not involved here so we cannot appeal to a varying mass-to-light ratio. + Moreover. the model Fundamental Plane is well defined. and so we cannot sav mt this is merely a question of fitting.," Moreover, the model Fundamental Plane is well defined and so we cannot say that this is merely a question of fitting." + The are at least two possible sources for this systematic ilference between the model and the virial theorem., The are at least two possible sources for this systematic difference between the model and the virial theorem. + The inst is to blame the HOP technique ancl assert that it systematically unclerestimates the cluster masses., The first is to blame the HOP technique and assert that it systematically underestimates the cluster masses. + The second is to sav that the internal cluster properties. (like velocity. distribution) vary svstematically with mass an so 16 normalisation of the virial plane is mass dependent., The second is to say that the internal cluster properties (like velocity distribution) vary systematically with mass an so the normalisation of the virial plane is mass dependent. + Either way. we shall model in a mass depencencey ancl consider this in relation to the LOD technique.," Either way, we shall model in a mass dependency and consider this in relation to the HOP technique." + Lhe process for the variable virial normalisation is analogous., The process for the variable virial normalisation is analogous. + The samples of clusters derived from these simulations are all based. on the LOP technique., The samples of clusters derived from these simulations are all based on the HOP technique. + Phere may well be a systematic bias in the assignment of particles to clusters (see section 2.1))., There may well be a systematic bias in the assignment of particles to clusters (see section \ref{sec:haloid}) ). + As a consequence. the radii and. velocity dispersion. derived for a HOD. selected. cluster will also be biased.," As a consequence, the radii and velocity dispersion derived for a HOP selected cluster will also be biased." + Clearly the bias will be more significant for smaller svstcHis., Clearly the bias will be more significant for smaller systems. + 1n this subsection we seek to account for systemic ellects of using HOP for identifving cluster membership. and derive a renormalisation procedure taking account of this and matching the dataset to the expected. virial theorem. Fundamental Plane (equation 33)).," In this subsection we seek to account for systemic effects of using HOP for identifying cluster membership, and derive a renormalisation procedure taking account of this and matching the dataset to the expected virial theorem Fundamental Plane (equation \ref{eq:virial}) )." + The easiest way to model this bias is to assume that the mocdel-based. estimate (biased) for the mass. AZ. is related to the actual mass M by a simple scaling relationship for some exponent a.," The easiest way to model this bias is to assume that the model-based estimate (biased) for the mass, $M$, is related to the actual mass $\cal{M}$ by a simple scaling relationship for some exponent $\alpha$." + The virial expression for the mass then becomes where the right hand side now refers to quantities derived from the model., The virial expression for the mass then becomes where the right hand side now refers to quantities derived from the model. + We can eliminate AZ from this in terms of the model surface mass density X—Mr; to give ‘Taking logs and using y=2.5fog® Linally vielcls, We can eliminate $M$ from this in terms of the model surface mass density $\Sigma = M/{r_h^2}$ to give Taking logs and using $\mu = -2.5log \Sigma$ finally yields +" The acceleration parameter can be expressed in terms of a ratio of a post shock pressure P"" 10 the upfront density p.",K = _1^2 The acceleration parameter can be expressed in terms of a ratio of a post shock pressure $P'$ to the upfront density $\rho$. + For relativistic strong shocks with the ratio of specific heats 5=1/3 while foruou-relativistic shock with 5=5/3 ae roe LoLP» Expressing the relevant quantities iu terms of Jv we find v== cl Vi=:— Vi=rs - VmHog (αη V where tlie approximations assumestrongly relativistic motion.," For relativistic strong shocks with the ratio of specific heats $\hat{\gamma}=4/3$ K= while fornon-relativistic shock with $\hat{\gamma}=5/3$ V_1^2 = = Expressing the relevant quantities in terms of $K$ we find = V = 1- V_1= V_1' = = V_x' ( 1 - ), where the approximations assumestrongly relativistic motion." +Finally. expressing V. in terms of V4 we find v2 001 (,"Finally, expressing $V $ in terms of $V_1$ we find V^2 = =" +Finally. expressing V. in terms of V4 we find v2 001 (5,"Finally, expressing $V $ in terms of $V_1$ we find V^2 = =" +Finally. expressing V. in terms of V4 we find v2 001 (5),"Finally, expressing $V $ in terms of $V_1$ we find V^2 = =" +equilibrium may become inappropriate.,equilibrium may become inappropriate. + This might happen at very small number densities., This might happen at very small number densities. +" In this case, we have to follow the detailed evolution of the chemical network using the full set of non-IE equations (6))."," In this case, we have to follow the detailed evolution of the chemical network using the full set of non-IE equations\ref{eN}) )." +" In those simulations, omitting the assumption of IE, the results differ only marginally from that ones of the IE-simulations."," In those simulations, omitting the assumption of IE, the results differ only marginally from that ones of the IE-simulations." +" Effects on the hydrodynamic evolution, coupled by the cooling/heating function to the chemical network, cannot be detected at all."," Effects on the hydrodynamic evolution, coupled by the cooling/heating function to the chemical network, cannot be detected at all." +" However, the chemical composition shows a slight deviation with respect to the IE simulations, and this occurs only in the direct vicinity of the shock."," However, the chemical composition shows a slight deviation with respect to the IE simulations, and this occurs only in the direct vicinity of the shock." + In the bottom panels of Fig., In the bottom panels of Fig. + 4 we show the number density of and around the shock using IE and non-IE., \ref{fChem} we show the number density of and around the shock using IE and non-IE. + In the non-IE simulation a very small region adjacent to the shock exists where for the degree of ionization is lower than in the IE case., In the non-IE simulation a very small region adjacent to the shock exists where for the degree of ionization is lower than in the IE case. + The corresponding chemical timescale becomes comparable with the hydrodynamical timescale., The corresponding chemical timescale becomes comparable with the hydrodynamical timescale. +" In combination with the motion of the shock, this produces the observed behavior."," In combination with the motion of the shock, this produces the observed behavior." +" Since, at the temperatures present in that region, the chemical rates for the Hare higher, the chemical time-scales are short, and a similar feature is not observed."," Since, at the temperatures present in that region, the chemical rates for the are higher, the chemical time-scales are short, and a similar feature is not observed." +" Owing to the discussed offset between temperature shock and density shock, the chemical time scales at the shock region are even larger, which results in a more extended region of delayed ionization."," Owing to the discussed offset between temperature shock and density shock, the chemical time scales at the shock region are even larger, which results in a more extended region of delayed ionization." +" The same behavior can be observed fort, only at much lower number densities."," The same behavior can be observed for, only at much lower number densities." +" Though the effects of non-IE for the primordial composition are only marginal, the influence of the particular conditions at the shock position has been demonstrated."," Though the effects of non-IE for the primordial composition are only marginal, the influence of the particular conditions at the shock position has been demonstrated." +" Thus for a medium containing some fraction of heavy elements which have very low densities for usual abundances, the non-IE must be taken into consideration."," Thus for a medium containing some fraction of heavy elements which have very low densities for usual abundances, the non-IE must be taken into consideration." +" If omitting the assumption of uniform temperature for all fluid components (electrons, ions) the effects on non-IE maybe much larger (?).."," If omitting the assumption of uniform temperature for all fluid components (electrons, ions) the effects on non-IE maybe much larger \citep{Teyssier98}." + The initial conditions for cosmological simulations of large-scale structure formation are given by a spectrum of perturbations., The initial conditions for cosmological simulations of large-scale structure formation are given by a spectrum of perturbations. +" To take this into account, we add Gaussian random perturbations according to the cosmological power spectrum to the particular perturbation given by Eq. (12))."," To take this into account, we add Gaussian random perturbations according to the cosmological power spectrum to the particular perturbation given by Eq. \ref{ePancake}) )." +" Then, in a given spatial region of size L, a pronounced pancake structure will only form if the considered initial perturbation amplitude dominates over the neighboring perturbation amplitudes at comparable scale size."," Then, in a given spatial region of size $L$, a pronounced pancake structure will only form if the considered initial perturbation amplitude dominates over the neighboring perturbation amplitudes at comparable scale size." +" Thus, we subsequently include all perturbation modes up to a scale size of (1/8,1/4,1/2)xL."," Thus, we subsequently include all perturbation modes up to a scale size of $(1/8,1/4,1/2)\times L$." + In Fig., In Fig. + 5 we show the density and temperature profiles of simulations including the small scale perturbations and cooling and heating atz=0., \ref{fGef} we show the density and temperature profiles of simulations including the small scale perturbations and cooling and heating at$z=0$. +" Clearly, the perturbation modes at scales comparable with L have the most impact on the final profiles at z=0."," Clearly, the perturbation modes at scales comparable with $L$ have the most impact on the final profiles at $z=0$." + In this case most of the power from perturbation modes at neighboring scales will be added., In this case most of the power from perturbation modes at neighboring scales will be added. +" In particular, this leads to an enhancement of the core density."," In particular, this leads to an enhancement of the core density." + The modes smaller than the actual Jeans length are erased by the heating due to the UV background., The modes smaller than the actual Jeans length are erased by the heating due to the UV background. +" However, the magnitudes of the temperatures and their coarse profiles in the post-shock region are on the same order and the density profiles are nearly preserved."," However, the magnitudes of the temperatures and their coarse profiles in the post-shock region are on the same order and the density profiles are nearly preserved." + We conclude that it is sufficient to consider only the collapse of a single (large enough) mode in order to gain information on the thermal and chemical state of the structures of interest., We conclude that it is sufficient to consider only the collapse of a single (large enough) mode in order to gain information on the thermal and chemical state of the structures of interest. +" In addition, the conservation of the system's symmetry serves as a probe for the quality of the numerical treatment."," In addition, the conservation of the system's symmetry serves as a probe for the quality of the numerical treatment." + The non-radiative simulations reproduce the earlier found analytic results with high accuracy (?).., The non-radiative simulations reproduce the earlier found analytic results with high accuracy \citep{Shandarin89}. + In the preceding section it was shown that the length scale of the perturbation L is the determining parameter for the evolution of the pancake characteristics (temperature and density profiles)., In the preceding section it was shown that the length scale of the perturbation $L$ is the determining parameter for the evolution of the pancake characteristics (temperature and density profiles). + In Fig., In Fig. +" 6 we present the L-dependence of the final values of the central temperature T,, the central hydrogen number density ny-, the radius of the isothermal core Λο, and the temperature at its edge Το."," \ref{fScaling} we present the $L$ -dependence of the final values of the central temperature $T_c$, the central hydrogen number density $n_{\mathsc{H}c}$ , the radius of the isothermal core $\lambda_\circ$ , and the temperature at its edge $T_\circ$ ." +scaling and shift.,scaling and shift. + We first measured the aperture colour inside.- the κ 2lmag-arcsec 2 Johuson P isophote., We first measured the aperture colour inside the 24 $\cdot$ $^{-2}$ Johnson $B$ isophote. +: Then we also obtained the colour inside the isophote of radius the effective radius (as measured in the Bo baud)., Then we also obtained the colour inside the isophote of radius the effective radius (as measured in the $B$ band). + Again. the Galactic extinction correction was performed using the Burstein Illes (19521) maps.," Again, the Galactic extinction correction was performed using the Burstein Heiles \cite{Buh82}) ) maps." + Conversion are 3.98 in B and 2.51 in r: both values were interpolated from Fitzpatrick (1999) )., Conversion constants are 3.98 in $B$ and 2.51 in $r$; both values were interpolated from Fitzpatrick \cite{Fit99}) ). + Iu Table 3 we sununuarize all these results: appareut total and οι inaguitudes iu coluuuns (2) aud (3): effective radius dm kpce and arcsec m cohuuns (£1) and (5) respectively: absolute B magnitudes corrected from Galactic extinetion in colunuu (6) aud effective. and isophote 21aünagsresec ? Boor colours in coblunus (7) and (8)., In Table 3 we summarize all these results: apparent total and $B_{24}$ magnitudes in columns (2) and (3); effective radius in kpc and arcsec in columns (4) and (5) respectively; absolute $B$ magnitudes corrected from Galactic extinction in column (6) and effective and isophote 24 $\cdot$ $^{-2}$ $B-r$ colours in columns (7) and (8). + Colour information is only available for those ealaxies with Camu r maeuitude measured by Vitores et al. (1996a))., Colour information is only available for those galaxies with Gunn $r$ magnitude measured by Vitores et al. \cite{Alv96a}) ). +"In the lower panel we show the frequency-luminosity ""Sependence for the first and second overtones along the three isochrones.",In the lower panel we show the frequency-luminosity dependence for the first and second overtones along the three isochrones. + The frequencies were calculated with our onadiabatic pulsation code (Dziembowski. 1977) for envelope models with surface parameters from the The data suggest that b» and b; sequences may be connected with excitation of these two overtones.," The frequencies were calculated with our nonadiabatic pulsation code (Dziembowski, 1977) for envelope models with surface parameters from the The data suggest that $_2$ and $_3$ sequences may be connected with excitation of these two overtones." +" The frequency of the specified mode ts roughly proportional to the square of mean nmensity. hence v;xLO,"," The frequency of the specified mode is roughly proportional to the square of mean density, hence $\nu_k\propto +L^{-3/4}$." +" while v4,οςL.", while $\nu_{\rm max}\propto L$. +" Therefore. with —nereasing Iuminosity. the order of modes preferred by SLOs ""Secreases."," Therefore, with increasing luminosity, the order of modes preferred by SLOs decreases." + Another characteristic of the SLOs is the dependence of oscillation amplitudes on stellar parameters., Another characteristic of the SLOs is the dependence of oscillation amplitudes on stellar parameters. + Kjeldsen Bedding (1995) derived a semi-empirical expression for the bolometric luminosity amplitude of 0LxCMMS., Kjeldsen Bedding (1995) derived a semi-empirical expression for the bolometric luminosity amplitude of $\delta L\propto(L/M)T_{\rm eff}^{-2}$. + Samadi et al. (, Samadi et al. ( +2007) considered à more general dependence of amplitude on the luminosity-to-mass ratio. (L/My.,"2007) considered a more general dependence of amplitude on the luminosity-to-mass ratio, $(L/M)^s$ ." + They show that the s-exponent (sv in their notation) is determined by the eddy time-correlation form., They show that the $s$ -exponent $sv$ in their notation) is determined by the eddy time-correlation form. + The Gaussian form yields s=1. while the Lorentzian s=0.7.," The Gaussian form yields $s=1$, while the Lorentzian $s=0.7$." + They also show that the latter value fits the ground-based data on radial velocity amplitudes better., They also show that the latter value fits the ground-based data on radial velocity amplitudes better. + The space data on luminosity amplitudes yield discrepant conclusions., The space data on luminosity amplitudes yield discrepant conclusions. + Stello et al. (, Stello et al. ( +2010) find s=0.7 from theKepler data. while Mosser et al. (,"2010) find $s\approx0.7$ from the data, while Mosser et al. (" +2010) found s=0.89+0.02 from their analysis of the CoRoT data.,2010) found $s=0.89\pm0.02$ from their analysis of the CoRoT data. + The OGLE-III Catalog contains data on mode amplitudes in the / band., The OGLE-III Catalog contains data on mode amplitudes in the $I$ band. + We compared these data with the prediction from a generalized form of Eq.(8) of Kjeldsen Bedding (1995). which at t=800nm yields The plots in Fig.," We compared these data with the prediction from a generalized form of Eq.(8) of Kjeldsen Bedding (1995), which at $\lambda=800 {\mbox{ nm}}$ yields The plots in Fig." + 3 show that the OGLE data follow the relation with »=0.9., 3 show that the OGLE data follow the relation with $s\approx0.9$. + Both. s=0.7 and I are excluded.," Both, $s=0.7$ and 1 are excluded." + Our result is. thus. similar to that of Mosser et al. (," Our result is, thus, similar to that of Mosser et al. (" +2010).,2010). + Identification of modes responsible for the peaks in power spectra such as those shown in Fig., Identification of modes responsible for the peaks in power spectra such as those shown in Fig. + | would give us a basis for areal seismic probing of individual luminous red giants., 1 would give us a basis for a real seismic probing of individual luminous red giants. + The plots in the bottom panel of Fig., The plots in the bottom panel of Fig. + 2 suggest that the dominant peaks are associated with the first and second overtones of radial pulsation., 2 suggest that the dominant peaks are associated with the first and second overtones of radial pulsation. + If it was true. the frequency distance between the peaks would be the large separation. Av.," If it was true, the frequency distance between the peaks would be the large separation, $\Delta\nu$." + | lower luminosity stars. more than one large separation is measured.," In lower luminosity stars, more than one large separation is measured." + Their mean value. . is regarded as an important seismic parameter.," Their mean value, $<\!\Delta\nu\!>$, is regarded as an important seismic parameter." +" However. the mean distance between sequences b» and bs at specified W, cannot be identified with «Av because it does not refer to the same object."," However, the mean distance between sequences $_2$ and $_3$ at specified $W_I$ cannot be identified with $<\!\Delta\nu\!>$ because it does not refer to the same object." + We need the distance between the dominant peaks measured separately 1n individual power spectra., We need the distance between the dominant peaks measured separately in individual power spectra. + With this in mind. we selected over 100 objects with the most significant (S/N.> 7) peaks within the two sequences.," With this in mind, we selected over 100 objects with the most significant $S/N\ge7$ ) peaks within the two sequences." + Plots in Fig., Plots in Fig. + 4 show that. although the observed Witvs) relation is reproduced well by the models. the difference between the frequencies. which 15 shown in the lower panel. is by some 20 percent less than the calculated difference between the second and the first overtones.," 4 show that, although the observed $W_I(\nu_2)$ relation is reproduced well by the models, the difference between the frequencies, which is shown in the lower panel, is by some 20 percent less than the calculated difference between the second and the first overtones." + This difference is quite a robust quantity depending almost only on vi., This difference is quite a robust quantity depending almost only on $\nu_1$. + It is only weakly dependent on the isochrone parameters. as may be seen in the figure. as well as on details of modeling the envelope and its pulsation.," It is only weakly dependent on the isochrone parameters, as may be seen in the figure, as well as on details of modeling the envelope and its pulsation." + Our first suspicion has been that the discrepancy may be due to inadequacies of our codes. which were written long ago (Paezyrísski 1969: Dziembowski 1977).," Our first suspicion has been that the discrepancy may be due to inadequacies of our codes, which were written long ago (Paczyńsski 1969; Dziembowski 1977)." + The microphysics has been updated. but still the turbulent pressure and the Lagrangian perturbation of the convective flux are ignored in these codes.," The microphysics has been updated, but still the turbulent pressure and the Lagrangian perturbation of the convective flux are ignored in these codes." + Fortunately. Xiong Deng (2004) have provided a functional expression for radial mode frequencies calculated with their much more advanced treatment of convection for a sequenceof red giant envelope models.," Fortunately, Xiong Deng (2004) have provided a functional expression for radial mode frequencies calculated with their much more advanced treatment of convection for a sequenceof red giant envelope models." + Repeating calculation with our code for models with the same surface parameters.," Repeating calculation with our code for models with the same surface parameters," +an order of magnitude lower than those estimated for the total molecular gas component in both galaxies and local clouds. ancl (his remains true independent of any adjusiments to the ealaxv data.,"an order of magnitude lower than those estimated for the total molecular gas component in both galaxies and local clouds, and this remains true independent of any adjustments to the galaxy data." + As cliscussed earlier. instead of acljusting the star formation rates. we could have adjusted the GSO4 galaxy masses (downward) by (he same constant offset in log(M).," As discussed earlier, instead of adjusting the star formation rates, we could have adjusted the GS04 galaxy masses (downward) by the same constant offset in log(M)." + By not correcting {he mass estimates we are assuming that the molecular-line derived masses and the extinction derived masses accurately reflect the same cloud material. that is. Ape = Alea and Me; = Meg.," By not correcting the mass estimates we are assuming that the molecular-line derived masses and the extinction derived masses accurately reflect the same cloud material, that is, $M_{DG}$ $=$ $M_{HCN}$ and $M_{TG}$ $=$ $M_{CO}$." + To assess this possibility for the case of the total cloud masses. Mqgé;. we compared the exüncGon measurements with CO observations of a subset of the local cloud sample.," To assess this possibility for the case of the total cloud masses, $M_{TG}$, we compared the extinction measurements with CO observations of a subset of the local cloud sample." + We obtained CO data for five of the elouds from the archive of the CLA 1.2 m Telescope (Dame οἱ al., We obtained CO data for five of the clouds from the archive of the CfA 1.2 m Millimeter-wave Telescope (Dame et al. + 2001)., 2001). + The CO observations were averaged over the individual clouds and the integrated CO intensities were measured for each cloud., The $^{12}$ CO observations were averaged over the individual clouds and the integrated CO intensities were measured for each cloud. +" Applving the standard CO-to-Ll» conversion factor of 2 x 107"") em2? Ux km !) ! (Dame et al.", Applying the standard $_2$ conversion factor of 2 x $^{20}$ $^{-2}$ (K km $^{-1}$ $^{-1}$ (Dame et al. + 2001) to convert the integrated intensities to Il» column densities. we determined the mass of each cloud.," 2001) to convert the integrated intensities to $_2$ column densities, we determined the mass of each cloud." + We [ound these CO derived masses to all agree with the corresponding extinction (Ag > 0.1 mag) derived masses to better than1256.. indicating that the extinction (Ay: > 0.1 mag) and CO derived (otal masses both trace the same cloud material for local clouds.," We found these CO derived masses to all agree with the corresponding extinction $_K$ $\geq$ 0.1 mag) derived masses to better than, indicating that the extinction $_K$ $>$ 0.1 mag) and CO derived total masses both trace the same cloud material for local clouds." + This suggests that total masses derived [rom CO can be a good proxy for extinction derived total masses and (hus that the masses derived fom CO observations of galaxies can be compared directly with those of the local cloud sample. provided that the galaxy measurements (race the sumnied CO emission from a population of GMCs.," This suggests that total masses derived from CO can be a good proxy for extinction derived total masses and thus that the masses derived from CO observations of galaxies can be compared directly with those of the local cloud sample, provided that the galaxy measurements trace the summed CO emission from a population of GMCs." + If there is any diffuse CO emission Irom inert. non star-forming. molecular gas contributing to the galaxv-weraged CO measurements. (hen the CO masses derived [or galaxies overestimate the masses in star forming GMCs.," If there is any diffuse CO emission from inert, non star-forming, molecular gas contributing to the galaxy-averaged CO measurements, then the CO masses derived for galaxies overestimate the masses in star forming GMCs." + In such a case the CO derived masses for the galaxies would have to be adjusted downwards to compare to ihe local observations., In such a case the CO derived masses for the galaxies would have to be adjusted downwards to compare to the local observations. + A similar comparison of extinction and IICN derived masses is not possible lor the local clouds since the corresponding HCN observations of these clouds do not exist., A similar comparison of extinction and HCN derived masses is not possible for the local clouds since the corresponding HCN observations of these clouds do not exist. + This is unfortunate because (he ICN masses derived by GSO4 are likely upper limits to the true masses (Gao and Solomon 2004b)., This is unfortunate because the HCN masses derived by GS04 are likely upper limits to the true masses (Gao and Solomon 2004b). + For example. if the clouds are bound but not. virialized then the derived masses could be somewhat underestimated.," For example, if the clouds are bound but not virialized then the derived masses could be somewhat underestimated." + Thus. although it appears that the extragalactic CO derived masses can be directly. placed on the SER-Molecular Cloud Mass diagram without any svstematic adjustment. the situation is somewhat less certain lor the HCN masses derived by GS04.," Thus, although it appears that the extragalactic CO derived masses can be directly placed on the SFR-Molecular Cloud Mass diagram without any systematic adjustment, the situation is somewhat less certain for the HCN masses derived by GS04." + However. we note that the average ratio of dense gas (i.e. Ay > 0.8 mag) to total cloud mass (1.6... Ay > 0.1 mag.)," However, we note that the average ratio of dense gas (i.e., $_V$ $\geq$ 0.8 mag) to total cloud mass (i.e., $_V$ $\geq$ 0.1 mag.)" + ealeulated from the extinction data is = 0.10 + 0.06 for the sample of local clouds., calculated from the extinction data is $$ $=$ 0.10 $\pm$ 0.06 for the sample of local clouds. + For the G504 sample of galaxies we find = 0.16 d: 0.14 comparing the ICN and CO derived masses., For the GS04 sample of galaxies we find $$ = 0.16 $\pm$ 0.14 comparing the HCN and CO derived masses. + The relatively close correspondence of fpe; Ior these (vo samples is consistent. wilh the idea that the hieh, The relatively close correspondence of $f_{DG}$ for these two samples is consistent with the idea that the high +Following ? we adopt a forward modeling approach to fitting our ARCES A star observations.,Following \citet{blake2010} we adopt a forward modeling approach to fitting our ARCES A star observations. + We focus on order 3. spanning (he wavelength range 970 indicates convolution and | indicates ""evaluated at."," This scaled spectrum $T_{0}^{\tau}$ is convolved with an estimate of the spectrograph line spread function (LSF), interpolated onto a lower resolution pixel grid that is defined through a polynomial relation between pixel, $i$, and wavelength, $W(i)$, and finally normalized by a continuum function, $N(i)$ , that is a polynomial in $i$ where $\otimes$ indicates convolution and $\vert$ indicates “evaluated at”." + This model has 10 free parameters: Four lor N (7). three for WWΕν two for the LSE. which is a 1D Moffat. fanction (?).. and 7.," This model has 10 free parameters: four for $N(i)$ , three for $W(i)$, two for the LSF, which is a 1D Moffat function \citep{moffat1969}, and $\tau$." + We fit this model to order 3 of each A star spectrum using an implementation of the AMOEBA downhill simples method (?.. 7)) to minimize 47.," We fit this model to order 3 of each A star spectrum using an implementation of the AMOEBA downhill simplex method \citealt{nelder1965}, \citealt{press1992}) ) to minimize $\chi^{2}$." + Examples of fits are shown in Figures 2. and 3.., Examples of fits are shown in Figures \ref{fig1} and \ref{fig1b}. + The agreement between our theoretical telluric template and the observed. spectra is excellent overall. with typical residuals of 1—2% and no evidence for degradation of the fits in the centers of prominent absorption limes for regions where ihe line density is modest and the average line optical depth is <1.," The agreement between our theoretical telluric template and the observed spectra is excellent overall, with typical residuals of $1-2\%$ and no evidence for degradation of the fits in the centers of prominent absorption lines for regions where the line density is modest and the average line optical depth is $<1$." + As shown in Figure 3.. in regions having numerous lines with large optical depths (he overall quality of the fits deteriorates and the models exhibit sienilicant discrepancies in the cores of some absorption features.," As shown in Figure \ref{fig1b}, in regions having numerous lines with large optical depths the overall quality of the fits deteriorates and the models exhibit significant discrepancies in the cores of some absorption features." + In all cases the fit residuals do exhibit slowly varving features that we attribute to imperlect correction for CCD Bringing. which has a large amplitude at (hese red wavelengths.," In all cases the fit residuals do exhibit slowly varying features that we attribute to imperfect correction for CCD fringing, which has a large amplitude at these red wavelengths." + Uncertaintv in the IIITRAN Π.Ο line parameters maa also lead to discrepancies between our models and observations ancl we note that incorporating empirical lime parameters. such as (hose from ?.. may help to improve the overall quality of the fits.," Uncertainty in the HITRAN $_{2}$ O line parameters may also lead to discrepancies between our models and observations and we note that incorporating empirical line parameters, such as those from \citet{alekseeva2010a}, may help to improve the overall quality of the fits." + We [found (hat overall the best-fit modelparameters had very small scatter. demonstrating that ARCES is physically stable. with small shifts in wavelength solutions over the period of a vear.," We found that overall the best-fit modelparameters had very small scatter, demonstrating that ARCES is physically stable, with small shifts in wavelength solutions over the period of a year." + ILowever. we expect (hat 7. which is presumably related to zenith angle ancl the amount of IHl4O in the atmosphere. should varywidely since I4O. unlike CII; (7).. is distributed in the lower atmosphereand iis concentration can change by," However, we expect that $\tau$, which is presumably related to zenith angle and the amount of $_{2}$ O in the atmosphere, should varywidely since $_{2}$ O, unlike $_{4}$ \citep{blake2010}, , is in-homogeneously distributed in the lower atmosphereand its concentration can change by" +"Figure 1 shows results for photons with energy E=0.5 keV. The upper-left panel depicts the case of the “orthogonal rotator,” in which the magnetic dipole vector rotates in the yz plane, intersecting the line of sight at »=0.","Figure \ref{fig:pfracE5} shows results for photons with energy $E=0.5$ keV. The upper-left panel depicts the case of the “orthogonal rotator,” in which the magnetic dipole vector rotates in the $yz$ plane, intersecting the line of sight at $\psi = 0$." + The upper-right and lower-left panels show less extreme geometries., The upper-right and lower-left panels show less extreme geometries. +" In G2, most of the photons are emitted with 6=7/4, and in G3, the dipole vector intersects the line of sight when «»=0, but sweeps out a cone whose base is perpendicular to the NS rotation axis."," In G2, most of the photons are emitted with $\delta\ga \pi/4$, and in G3, the dipole vector intersects the line of sight when $\psi = 0$, but sweeps out a cone whose base is perpendicular to the NS rotation axis." + The lower-right panel shows the extreme case when the rotational and magnetic axes are aligned., The lower-right panel shows the extreme case when the rotational and magnetic axes are aligned. +" There are two key features of Figure 1 that highlight the interplay between the NS geometry, magnetic field strength, and vacuum polarization effects to produce the emitted polarization."," There are two key features of Figure \ref{fig:pfracE5} that highlight the interplay between the NS geometry, magnetic field strength, and vacuum polarization effects to produce the emitted polarization." +" These features include: (1) |II°™| is smaller for values of ~ close to 0 than for values of / close to 7/2; (2) at wzz0, the sign of II™ is positive for B=7x1015 G, and negative for stronger magnetic fields."," These features include: (1) $|\Pi^{\rm em}|$ is smaller for values of $\psi$ close to $0$ than for values of $\psi$ close to $\pi/2$; (2) at $\psi\approx 0$, the sign of $\Pi^{\rm em}$ is positive for $B=7\times 10^{13}$ G, and negative for stronger magnetic fields." + Result (1) can be understood by considering the variation in X and O mode opacities with angle between the photon propagation and magnetic field directions., Result (1) can be understood by considering the variation in X and O mode opacities with angle between the photon propagation and magnetic field directions. +" For the orthogonal rotator, small phases roughly correspond to emission along the observer line of sight and"," For the orthogonal rotator, small phases roughly correspond to emission along the observer line of sight and" +variations in eccentric systems caused by dramatic changes to the physical properties (density and temperature) of the WCR. which ts radiative (optically thick) at periastron. but adiabatic (optically thin) at apastron.,"variations in eccentric systems caused by dramatic changes to the physical properties (density and temperature) of the WCR, which is radiative (optically thick) at periastron, but adiabatic (optically thin) at apastron." + Here. we present analytic predictions of the thermal spectra at radio frequencies from different radiative models. which satisfy the condition of radiative shocks (y.< 1) for the WCR.," Here, we present analytic predictions of the thermal spectra at radio frequencies from different radiative models, which satisfy the condition of radiative shocks $\chi<$ 1) for the WCR." + We have estimated the value of y for each shock of the WCR., We have estimated the value of $\chi$ for each shock of the WCR. + In Table 1. we list the different scenarios of colliding wind binaries that we have studied in this paper.," In Table 1, we list the different scenarios of colliding wind binaries that we have studied in this paper." + In models BI-B3. we have assumed identical wind sources (B= 1) in order to investigate the effect of the binary separation.," In models B1-B3, we have assumed identical wind sources $\beta$ = 1) in order to investigate the effect of the binary separation." + On the other hand. in models B2. B4. and BS. we have assumed different wind momentum ratios for the binary systems. while the distance between the components is unchanged.," On the other hand, in models B2, B4, and B5, we have assumed different wind momentum ratios for the binary systems, while the distance between the components is unchanged." + In models B4 and B5. we indicate the highest value of y for each system.," In models B4 and B5, we indicate the highest value of $\chi$ for each system." + These models for the radio emission from colliding wind binaries are presented in Figure 4.., These models for the radio emission from colliding wind binaries are presented in Figure \ref{f4}. +. We show the optical depth of the WCR as functions of the impact parameter r (left panels) and the predicted flux density at radio frequencies (right panels)., We show the optical depth of the WCR as functions of the impact parameter $r$ (left panels) and the predicted flux density at radio frequencies (right panels). + Top panels show the results of our Models BI-B3. while the models B2. B4. and BS are shown in the bottom panels.," Top panels show the results of our Models B1-B3, while the models B2, B4, and B5 are shown in the bottom panels." + All models are calculated for an observer located in the orbital plane along the symmetry axis (which corresponds to a system with an inclination angle /=90° with the most powerful stellar wind in front. as shown in Figure 1).," All models are calculated for an observer located in the orbital plane along the symmetry axis (which corresponds to a system with an inclination angle $i=90^{\circ}$ with the most powerful stellar wind in front, as shown in Figure 1)." + First. we observe that. for low-impact parameters (7/Ro< 1). the optical depth tiyce(43GHz) of the WCR (Figure 4-- a) does not depend on the value of r and scales as D.," First, we observe that, for low-impact parameters $r/R_0 \leq 1$ ), the optical depth $\tau_{WCR} (\mbox{43 GHz})$ of the WCR (Figure \ref{f4}- -a) does not depend on the value of $r$ and scales as $D^{-3}$." + On the other hand. for higher impact parameters (r/Ro.>> 1). τινGHz)&77? in all models.," On the other hand, for higher impact parameters $r/R_0 \gg 1$ ), $\tau_{WCR} (\mbox{43 GHz}) +\propto r^{-5}$ in all models." + At a given line of sight. we also note that. in this limit. the optical depth increases as D.," At a given line of sight, we also note that, in this limit, the optical depth increases as $D^{2}$." + The transition impact parameter is 7xRo (see. also. Appendix A).," The transition impact parameter is $r\propto R_0$ (see, also, Appendix A)." + This behavior of the optical depth (and therefore of the emission measure of the thin shell) can be explained, This behavior of the optical depth (and therefore of the emission measure of the thin shell) can be explained +In a recent important paper Cussons. Langanke. Liolios (2002) have pointed out the potential resonant screening elfects on stellar PC + C reaction rates.,"In a recent important paper Cussons, Langanke, Liolios (2002) have pointed out the potential resonant screening effects on stellar $^{12}$ C + $^{12}$ C reaction rates." + The PC + PC [usion, The $^{12}$ C + $^{12}$ C fusion +"ave May=ποορ1)n3. where £p,,) is the average matter density in the Universe.","have $\Mvir=4\pi(340\langle \rho_m \rangle)R^3_{\rm vir}/3$, where $\langle \rho_m \rangle$ is the average matter density in the Universe." + For a halo with this profile. pxr Fas r>0¢ uid sectlly ‘all off as pXr? at the virial radius.," For a halo with this profile, $\rho \propto r^{-1}$ as $r\rightarrow0$ and smoothly fall off as $\rho \propto r^{-3}$ at the virial radius." + The concentration parameter weakly devends on the viral niass with a significaut scatter comparableo the systematic change in C over τος decades in Mag (Bullocketal.2001:Eke2001).," The concentration parameter weakly depends on the virial mass with a significant scatter comparableto the systematic change in $C$ over three decades in $\Mvir$ \citep{Bullock01,Eke01}." +. Later simulations paid most of atcution to the inner slope of the profiles., Later simulations paid most of attention to the inner slope of the profiles. + Some results favored a steeper profile than NFW deusiv cusp with pwr12 (Fulkushiee&Makino1997:Mooreetal.1998:Jing&Suto2000:Cohignaet 2000)..," Some results favored a steeper profile than NFW density cusp with $\rho \propto r^{-1.5}$ \citep{FM97,M98,JS00,G00}. ." +" More recent simulations of halos wiἩ iuillious of particles within B4, seen to indicat| that there is a scatter in the iuncr slope of he deαμ] profiles across a wide range of masses fiπι dwarfs to clusters.", More recent simulations of halos with millions of particles within $\Rvir$ seem to indicate that there is a scatter in the inner slope of the density profiles across a wide range of masses – from dwarfs to clusters. + The inner slope varies between these two shapes: the NEW with an asyotic slope of oue and the steeper Mooreetal.(1098) with slope 1.5 (seeKlvpiuetal.2001:ReedοἳFukushigeetal. 2001).," The inner slope varies between these two shapes: the NFW with an asyntotic slope of one and the steeper \citet{M98} with slope 1.5 \citep[see][]{Klypin01,Reed03,Navarro04,Diemand,WBO04,Tasitsiomi04,FKM04}." +. Different approximations for desity. profiles were suggested and tested in the literature., Different approximations for density profiles were suggested and tested in the literature. + Just as some other groups. we find that tιο 3D Séónrsic hree-paraieter approximation gives extremely eood fits for dark natter halos (Navarroctal.2001:Merrittetal. 2005).," Just as some other groups, we find that the 3D Sérrsic three-parameter approximation gives extremely good fits for dark matter halos \citep{Navarro04,Merritt2005}." +. We «lielitly inodifv his approxination wadding the average matter deusitv of the Universe (py); , We slightly modify this approximation by adding the average matter density of the Universe $\langle \rho_m \rangle$. +This term cau o0 neelected. if one fits the deusitv inside the virial radius.," This term can be neglected, if one fits the density inside the virial radius." + Yet. a larger distances. it elves al Huportant contribution.," Yet, at larger distances, it gives an important contribution." + The approximation ca- dC Written as where p ds the Sérrsic iudex., The approximation can be written as where $n$ is the Sérrsic index. + lu addition to all the nuuerical sinulatious. a siguificaut effort has been mace to compare the predictious of the ACDME model with the observations.," In addition to all the numerical simulations, a significant effort has been made to compare the predictions of the $\LCDM$ model with the observations." + This is the case of tli| nost recent set of high quality observations of large samples of rotation curves of galaxies or the strong eravitational leusiug studies which place au imniportaut upper liit on the amount of dark matter in ealaxies and clusters in the πιο few to tens of kiloparsec Gvithiu ο) where the necx«l of a cuspy density profile ds still subject of an exciting debate (e.g...Flores&Prunack.1991:Moore|cdemistctal.2OL.andreferencestherein )..," This is the case of the most recent set of high quality observations of large samples of rotation curves of galaxies or the strong gravitational lensing studies which place an important upper limit on the amount of dark matter in galaxies and clusters in the inner few to tens of kiloparsec (within $r_{\rm s}$ ) where the need of a cuspy density profile is still subject of an exciting debate \citep[e.g.,][and references +therein]{FP94,M94,dB03,S03,Rhee03,Keeton98, Keeton01,Bro04}." + Ou the tieoretical side. however. the origin of the sha]px| of the dark matter halo density xofile readis poorly uuderstood.," On the theoretical side, however, the origin of the shape of the dark matter halo density profile remains poorly understood." + It is eenerallv accepted tha the dark matter halos are assembled o hierarchicαἱ clustering as the result of halo ucreine and continuous accretion., It is generally accepted that the dark matter halos are assembled by hierarchical clustering as the result of halo merging and continuous accretion. + This mereiue scenario has motivated an iuterest in the analysis of the mass accretion lustory of the halos in conjunction with their structural properties (6.9..Wechsleretal. 2002).," This merging scenario has motivated an interest in the analysis of the mass accretion history of the halos in conjunction with their structural properties \citep[e.g.,][]{Wechsler}." +. The systematic study of the NEW cevusity fits to many simulated halos shows that their mass accretion history is closely correlated with the concentration parameter C acl. therefore. with the qass inside the scale radius ry (Wechsleretal.2002:Zhao2003:Tasitsionetal.200 D.," The systematic study of the NFW density fits to many simulated halos shows that their mass accretion history is closely correlated with the concentration parameter $C$ and, therefore, with the mass inside the scale radius $r_{\rm s}$ \citep{Wechsler,Zhao,Tasitsiomi04}." +. These results sugeestOO tha he formation process of the dark matter halos cau )e eenerallv understood by an carly phase of fas nass accretion and a late phase of slow accretion of mass., These results suggest that the formation process of the dark matter halos can be generally understood by an early phase of fast mass accretion and a late phase of slow accretion of mass. + In this scenario. the inner deuse regious of he halos are build up carly diving the fast phase of mass accretion when the halo mass increases with time much faster than the expansion rate of he Universe.," In this scenario, the inner dense regions of the halos are build up early during the fast phase of mass accretion when the halo mass increases with time much faster than the expansion rate of the Universe." + At later epochs. during the phase of slow mass accretion. the outer regions of the halo are built. while its inner regions stay almost intac (seeZhaoetal.2003).," At later epochs, during the phase of slow mass accretion, the outer regions of the halo are built, while its inner regions stay almost intact \citep[see][]{Zhao}." +. Despite to all this effor dedicated to. the understanding of the ccutral deuse regions of the dark natter halos. very little attention has been devoted to the study of thei outskirts. i.e. the regions bevoud the ornalb virial radius.," Despite to all this effort dedicated to the understanding of the central dense regions of the dark matter halos, very little attention has been devoted to the study of their outskirts, i.e. the regions beyond the formal virial radius." + The outer parts of the halos aid therefore their deusityv profiles exhihi in these regious large fiuctuatious which can be unuderstooc as the result of iufalliug dark iatter (Guchiding intalline samaller halos or substructure) or du fo major inmergers., The outer parts of the halos and therefore their density profiles exhibit in these regions large fluctuations which can be understood as the result of infalling dark matter (including infalling smaller halos or substructure) or due to major mergers. + Iu voth cases the infalling naterial has not reached he equilibriun with he rest of the halo (Seo X01).., In both cases the infalling material has not reached the equilibrium with the rest of the halo \citep[see][]{FM01}. . + Ou the coutrary.a considerable observation:d effort is being mace to neasure the nass distribition around ealaxies aud clusters at large distances using weak eravitational chsing," On the contrary,a considerable observational effort is being made to measure the mass distribution around galaxies and clusters at large distances using weak gravitational lensing" +Llaving derived the dynamical and. photometric properties of these galaxies. we are now in a position to look for clues to their origins.,"Having derived the dynamical and photometric properties of these galaxies, we are now in a position to look for clues to their origins." + Figure 13. shows the physical parameters of the rotation curves in these galaxies., Figure \ref{r0v0} shows the physical parameters of the rotation curves in these galaxies. + There are no signs of any correlation between ry and ery. suggesting that these galaxies occupy dark halos with the usual wide range of characteristics. although the sample is really too small to say anvthing definitive.," There are no signs of any correlation between $r_0$ and $v_0$, suggesting that these galaxies occupy dark halos with the usual wide range of characteristics, although the sample is really too small to say anything definitive." + Similarly. there is no cdiscernable correlation between the mass scale-length. ry. and the photometric scale-Iength. rr. (see Fig. 14))," Similarly, there is no discernable correlation between the mass scale-length, $r_0$, and the photometric scale-length, $r_d$, (see Fig. \ref{r0rd}) )" + avguing against Neistein a£s (1999) suggestion that these galaxies should be more dominated by the mass in their stellar disces than normal spiral ealaxies., arguing against Neistein 's (1999) suggestion that these galaxies should be more dominated by the mass in their stellar discs than normal spiral galaxies. + A more interesting result comes when we look at the Tullv-Fisher relation lor SO galaxies., A more interesting result comes when we look at the Tully-Fisher relation for S0 galaxies. +" Figure 15. shows the I-band relation derived by Neistein ((1999). which lie quite ‘lose to the Tullv-Fisher relation for normal spiral galaxies λογος ""Tully 1992). but with a laree RAIS scatter of 0.6 magnitudes about the best-fit line parallel to the gaancared Tullv-Fisher relation."," Figure \ref{IbandTFfig} shows the I-band relation derived by Neistein (1999), which lie quite close to the Tully-Fisher relation for normal spiral galaxies (Pierce Tully 1992), but with a large RMS scatter of $\sim 0.6$ magnitudes about the best-fit line parallel to the standard Tully-Fisher relation." + The galaxies in the current sample lie systematically below the Neistein galaxies., The galaxies in the current sample lie systematically below the Neistein galaxies. + With the galaxies all selected to have comparable rotation Es»eeds. we clearly cannot derive a slope for the Tullv-Fisher relation from these data. but fixing the slope to that of the relation for normal spiral galaxies. we find that the best-fit line is olfset from the spiral galaxy relation by ~15 magnitude. with a scatter about the best-fit. line of only 0.3 magnitudes.," With the galaxies all selected to have comparable rotation speeds, we clearly cannot derive a slope for the Tully-Fisher relation from these data, but fixing the slope to that of the relation for normal spiral galaxies, we find that the best-fit line is offset from the spiral galaxy relation by $\sim 1.8$ magnitude, with a scatter about the best-fit line of only $\sim 0.3$ magnitudes." + For this fit we exclude NCGC2612 as its distance estimate is rather uncertain due to a small svstenic velocity., For this fit we exclude NGC2612 as its distance estimate is rather uncertain due to a small systemic velocity. + Even allowing for the small sample size. this scatter is smaller than that in the Neistein sample at à statistically significant level.," Even allowing for the small sample size, this scatter is smaller than that in the Neistein sample at a statistically significant level." + ‘This small scatter ancl olfset in the Tullvy-Fisher relation by LS magnitude in ΕΕ is what one would expect i£ star formation had. been suddenly switchecl olf a few Cars ago so that these SO galaxies contained just. the old stellar population of a normal spiral galaxy.," This small scatter and offset in the Tully-Fisher relation by $\sim +1.8$ magnitude in the $I$ -band is what one would expect if star formation had been suddenly switched off a few Gyrs ago so that these S0 galaxies contained just the old stellar population of a normal spiral galaxy." + Using stellar population models (Charlot Bruzual 1991). assuming SOs galaxies had had similar formation histories as late-tvpe spirals until a few Cars ago. we would expect the stellar population of SOs to have faded significantly. resulting in à noticeable olfset in E magnitude. such as the one observed in our sample.," Using stellar population models (Charlot Bruzual 1991), assuming S0s galaxies had had similar formation histories as late-type spirals until a few Gyrs ago, we would expect the stellar population of S0s to have faded significantly, resulting in a noticeable offset in I magnitude, such as the one observed in our sample." + This point is mace even more dramatically if. we convert to estimated. Z/-band magnitudes using the prescription in Section ??.., This point is made even more dramatically if we convert to estimated $H$ -band magnitudes using the prescription in Section \ref{datasec}. + In this band. the luminosity is dominated. by the old stellar populations: as Fig.," In this band, the luminosity is dominated by the old stellar populations; as Fig." + 16 shows. the SOs in the current sample lie quite close to the Tullv-Fisher relation for later-tvpe spiral galaxies. suggesting that their old. stellar populations are rather similar.," \ref{HbandTFfig} shows, the S0s in the current sample lie quite close to the Tully-Fisher relation for later-type spiral galaxies, suggesting that their old stellar populations are rather similar." + In this paper. we have calculated the first detailed dynamical models for a small sample of cdge-on SO galaxies with small xilees.," In this paper, we have calculated the first detailed dynamical models for a small sample of edge-on S0 galaxies with small bulges." + In addition to producing clistribution functions for hese disc-dominated svstems. which look very much as one would. expect for normal disc systems. the analysis also returned. estimates for the parameters of their eravitationa »otentials.," In addition to producing distribution functions for these disc-dominated systems, which look very much as one would expect for normal disc systems, the analysis also returned estimates for the parameters of their gravitational potentials." + The interpretation ofthese data all points to a simple picture in which these svstems were formed. by the stripping of gas from normal spiral galaxies., The interpretation of these data all points to a simple picture in which these systems were formed by the stripping of gas from normal spiral galaxies. + The distribution functions are all well modeled. by unexceptional stellar disces. similar to those expected. in the old. stellar populations of spira ealaxies.," The distribution functions are all well modeled by unexceptional stellar discs, similar to those expected in the old stellar populations of spiral galaxies." + In addition. the galaxies obey a reasonably tigh ‘Tully-Fisher relation. which is olfset from the relation for normal spiral galaxies by the amount that one would expec if star formation had been shut olf a lew Cives ago. so tha all that remains in these svstems are the rather fainter ok stellar populations.," In addition, the galaxies obey a reasonably tight Tully-Fisher relation, which is offset from the relation for normal spiral galaxies by the amount that one would expect if star formation had been shut off a few Gyrs ago, so that all that remains in these systems are the rather fainter old stellar populations." + This result appears to conlliet with Neistein αν (1999) analysis. which showed a much greater scatter in the Tully-Fisher relation with less svstematic ollset.," This result appears to conflict with Neistein 's (1999) analysis, which showed a much greater scatter in the Tully-Fisher relation with less systematic offset." + Part. of. the dillerence may. be due to the lower signal-to-noise ratio of the data in their larger sample. which limited their ability to carry out detailed: dvnamical modeling. particularly for galaxies that lie very close to edge-on.," Part of the difference may be due to the lower signal-to-noise ratio of the data in their larger sample, which limited their ability to carry out detailed dynamical modeling, particularly for galaxies that lie very close to edge-on." +. However. there is also a systematic dillerence in the way that the samples were selected: the eclge-on galaxies in the analysis of this paper were specifically chosen to contain small bulges.," However, there is also a systematic difference in the way that the samples were selected: the edge-on galaxies in the analysis of this paper were specifically chosen to contain small bulges." + This selection criterion means that these galaxies are prime candidates to have formed from gas-strippec spiral galaxics., This selection criterion means that these galaxies are prime candidates to have formed from gas-stripped spiral galaxies. +" If as Neistein suggest. SOs are a mixed bag"" that [ormed in a variety of ways. it should come as no surprise that this particular subsample obey. a tight Tullv-Fisher relation that is not seen in the general population of SOs."," If, as Neistein suggest, S0s are a “mixed bag” that formed in a variety of ways, it should come as no surprise that this particular subsample obey a tight Tully-Fisher relation that is not seen in the general population of S0s." +field systems. the companion remains unknown. although the other known ELM white dwarfs are companions to pulsars (vanKerkwijketal..1996;Bassa 2006).. so it 1s possible that NLTT 11748 has a neutron star companion.,"field systems, the companion remains unknown, although the other known ELM white dwarfs are companions to pulsars \citep{van1996,bas2006}, so it is possible that NLTT 11748 has a neutron star companion." + If we assume that the companion is a neutron star. then the orbital period is expected to be ~10—15 hours based on the mass versus orbital period relations for He WDs with a neutron star companion (Benvenuto&DeVito.2005).," If we assume that the companion is a neutron star, then the orbital period is expected to be $\sim 10 - 15$ hours based on the mass versus orbital period relations for He WDs with a neutron star companion \citep{ben2005}." +. We searched for radio sources in the vicinity of NLTT 11748 using VizieR. and the nearest source is an unrelated object 2.5 areminutes away.," We searched for radio sources in the vicinity of NLTT 11748 using VizieR, and the nearest source is an unrelated object 2.5 arcminutes away." + It is also possible that the companion 1s another. more massive white dwarf.," It is also possible that the companion is another, more massive white dwarf." + Nelemansetal.(2001) present a population synthesis of double degenerate systems that experience at least two phases of mass transfer. out of which at least one will result in a commotr envelope allowing the white dwarfs to spiral in.," \citet{nel2001} present a population synthesis of double degenerate systems that experience at least two phases of mass transfer, out of which at least one will result in a common envelope allowing the white dwarfs to spiral in." + Their population synthesis also shows a correlation between the mass of the brighter white dwarf and the orbital period., Their population synthesis also shows a correlation between the mass of the brighter white dwarf and the orbital period. + Although unlikely. the companion could be a low-mass main sequence star.," Although unlikely, the companion could be a low-mass main sequence star." +" Our optical spectrum and 2MASS photometry do not show any evidence of a cool companion,", Our optical spectrum and 2MASS photometry do not show any evidence of a cool companion. + Since NLTT 11748 ts relatively faint. 2MASS_ provides a reliable measurement in the / band. with H having a considerably greater uncertainty.," Since NLTT 11748 is relatively faint, 2MASS provides a reliable measurement in the $J$ band, with $H$ having a considerably greater uncertainty." +" The absolute / magnitude for the white dwarf is M,=9.3.", The absolute $J$ magnitude for the white dwarf is $M_J = 9.3$. + A red dwarf of spectral type M5 to M5.5 would have equally contributed in this spectral region (Kirkpatrick&MeCarthy.1994). so a prospective late type companion would have to be much less massive than 0.2M...," A red dwarf of spectral type M5 to M5.5 would have equally contributed in this spectral region \citep{kir1994}, so a prospective late type companion would have to be much less massive than $0.2\ M_\odot$." + Evolutionary scenarios that produce ELM white dwarfs involve an episode of Roche lobe overflow and accretion onto adegenerate companion that stripped the envelope of the white dwarf progenitor leaving a low-mass degenerate core (Willems&Kolb.2002:BenvenutoDeVito.2005).," Evolutionary scenarios that produce ELM white dwarfs involve an episode of Roche lobe overflow and accretion onto a degenerate companion that stripped the envelope of the white dwarf progenitor leaving a low-mass degenerate core \citep{wil2002,ben2005}." +. Assuming halo membership. as suggested by the kinematics. the progenitor lifetime on the main-sequence would range from 7 to 9 Gyrs corresponding to a late G star with a mass of 0.87—0.93M. (Girardietal..2000).," Assuming halo membership, as suggested by the kinematics, the progenitor lifetime on the main-sequence would range from 7 to 9 Gyrs corresponding to a late G star with a mass of $0.87 - 0.93\ M_\odot$ \citep{gir2000}." +. Based on a model atmosphere analysis. we have shown that NLTT 11748 is an ELM white dwarf (M=0.167x0.005M...).," Based on a model atmosphere analysis, we have shown that NLTT 11748 is an ELM white dwarf $M = 0.167\pm0.005\ M_\odot$ )." + Although we assumed a pure hydrogen composition. near ultraviolet spectroscopy is required to confirm our. model atmosphere analysis and to constrain the abundance of heavy elements in the atmosphere.," Although we assumed a pure hydrogen composition, near ultraviolet spectroscopy is required to confirm our model atmosphere analysis and to constrain the abundance of heavy elements in the atmosphere." + NLTT 11748 is one of only a handful of low-gravity white dwarfs that allow a test of our models in these unique conditions., NLTT 11748 is one of only a handful of low-gravity white dwarfs that allow a test of our models in these unique conditions. + The star almost certainly belongs to a close binary system with a still unidentified degenerate companion., The star almost certainly belongs to a close binary system with a still unidentified degenerate companion. + A probable evolutionary scenario involves a post supernova system comprising an evolving G star and a neutron star that experienced Roche lobe overflow accretion. leaving an ELM degenerate star.," A probable evolutionary scenario involves a post supernova system comprising an evolving G star and a neutron star that experienced Roche lobe overflow accretion, leaving an ELM degenerate star." + This object would be part of a growing family of low-mass X-ray binary remnants., This object would be part of a growing family of low-mass X-ray binary remnants. + A radial velocity study should be able to constrain the orbital period and the mass function of the unseen companion., A radial velocity study should be able to constrain the orbital period and the mass function of the unseen companion. + On the other hand. the binary companion may be another white dwarf. thereby confirming another formation channel for ELM white dwarfs in parallel with the channel involving a neutron star.," On the other hand, the binary companion may be another white dwarf, thereby confirming another formation channel for ELM white dwarfs in parallel with the channel involving a neutron star." + Our kinematic study shows that the system belongs to the Galactic halo. which helps constrain the age. hence the mass of the donor star.," Our kinematic study shows that the system belongs to the Galactic halo, which helps constrain the age, hence the mass of the donor star." + Finally. studies of ELM white dwarfs are useful in exploring the effect of residual burning on white dwarf mass-radius relations.," Finally, studies of ELM white dwarfs are useful in exploring the effect of residual burning on white dwarf mass-radius relations." + In this regard an independent estimate of the radius using parallax measurements is essential., In this regard an independent estimate of the radius using parallax measurements is essential. + In addition. as more ELM white dwarts are discovered. a more comprehensive set of evolutionary tracks for these stars is required.," In addition, as more ELM white dwarfs are discovered, a more comprehensive set of evolutionary tracks for these stars is required." +at F1012NE: Table 5) was found. the primary tarect was asstumed to be single.,"at F1042M; Table 5) was found, the primary target was assumed to be single." + Convergent binary solutions were obtained for both 2N[ASS 2739AD and 2MASS 2952AB in both filters., Convergent binary solutions were obtained for both 2MASS $-$ 2739AB and 2MASS $-$ 2952AB in both filters. + Figure 5 shows the original and PSF-subtracted nuages for these pais., Figure 5 shows the original and PSF-subtracted images for these pairs. + Residuals frou the subtraction were less than of the peak flux. at the level of S10 (FSLIW) aud 23 (FI012M) times the backerouud noise.," Residuals from the subtraction were less than of the peak flux, at the level of 8–10 (F814W) and 2–3 (F1042M) times the background noise." + In both cases smaller residuals were obtained for fits to two sources rather than ai single source.," In both cases, smaller residuals were obtained for fits to two sources rather than a single source." + This validates the duplicity of 2ATASS 1531. 2052AT whose separation (07065-07007) is sualler than both the diffraction limit (0/08 at FSIIN. Q'11 at F1012AD) aud Nyquist sampling limit «07016 Q'09) of the instrument.," This validates the duplicity of 2MASS $-$ 2952AB, whose separation $\farcs$ $\pm$ $\farcs$ 007) is smaller than both the diffraction limit $\farcs$ 08 at F814W, $\farcs$ 11 at F1042M) and Nyquist sampling limit $\times$ $\farcs$ 046 = $\farcs$ 09) of the instrument." + We are able to overcome the former because our technique resolves even siguificaut overlap of two PSFs. particularly when they have nearly equal brightucss.," We are able to overcome the former because our technique resolves even significant overlap of two PSFs, particularly when they have nearly equal brightness." + The latter constraint is overcome by using a PSF senerated from multiple measurements. allowing us to subsample below the Nyquist luit.," The latter constraint is overcome by using a PSF generated from multiple measurements, allowing us to subsample below the Nyquist limit." + Table 1 lists the derived binary paralucters for these svstenis: no other sources cau be seen in the subtracted nuages., Table 4 lists the derived binary parameters for these systems; no other sources can be seen in the subtracted images. + PSF fitting of the other primary T dwarf targets revealed only oue poteutial fait companion to 2\TASS 1217 0311. detected at FI012M oulv.," PSF fitting of the other primary T dwarf targets revealed only one potential faint companion to 2MASS $-$ 0311, detected at F1042M only." + This possible detection. which may simply be a residual cosmic ray. is discussect in detail iu $5.2.," This possible detection, which may simply be a residual cosmic ray, is discussed in detail in $\S$ 5.2." +" No other colmpanious were identified around aux of the other tareet objects within U1. aud no other faint objects with T dwarf-like colors were identified in auv of the WPFC2 Huaces,"," No other companions were identified around any of the other target objects within $\farcs$ 1, and no other faint objects with T dwarf-like colors were identified in any of the WPFC2 images." + To obtain proper calibration and determine the uucertainties of our results. we ran the algoritlan described above on 20.000 simulated binaries constructed from the FstlW aud F1012M images of 221ÀSS L101.," To obtain proper calibration and determine the uncertainties of our results, we ran the algorithm described above on 20,000 simulated binaries constructed from the F814W and F1042M images of 2MASS $-$ 1404." + These test nuages sampled a range of separations 115 pixels (O%05 0769). all oricutations. aud flux ratios AM = 0 mag.," These test images sampled a range of separations 1–15 pixels $\farcs$ $\farcs$ 69), all orientations, and flux ratios $\Delta$ M = 0--7 mag." + Once processed through the PSF fitting routine. those test cases having output separations within 0.5 pixels and corrected fux ratios within 0.2 mae of the input values were cousicered recovered binaries.," Once processed through the PSF fitting routine, those test cases having output separations within 0.5 pixels and corrected flux ratios within 0.2 mag of the input values were considered recovered binaries." + Correctious and uncertainties to both positions aud flux ratios were then determined separately for 2\LASS 2739AB aud 2MASS 2952AB in each filter. usine only those recovered fest cases having similar input separatious (a>0723 aud a< OTL. respectively) aud flux ratios (1AM<2.5 and 0 0\farcs23$ and $a < 0\farcs14$ , respectively) and flux ratios $1 < {\Delta}M < 2.5$ and $0 < {\Delta}M < 1$ , respectively)." +" Typical flux ratio corrections were approximately — 0.10 mae 1νο,shifting the secondary. to brighter magnitudes) with lo uncertainties of 0.01 aud 0.3 mag for 2\LASS 27T39AD aud 2MÁÀSS 2952AB. respectively: separation lo uncertainties were 0.12 and 0.15 pixels (07005 and 07007). respectively. translating into position anele uncertainties of 77 aud 97."," Typical flux ratio corrections were approximately $-$ 0.10 mag (i.e.,shifting the secondary to brighter magnitudes) with $\sigma$ uncertainties of 0.04 and 0.3 mag for 2MASS $-$ 2739AB and 2MASS $-$ 2952AB, respectively; separation $\sigma$ uncertainties were 0.12 and 0.15 pixels $\farcs$ 005 and $\farcs$ 007), respectively, translating into position angle uncertainties of $\degr$ and $\degr$." + The values listed in Table Lreflect these corrections and uicertainties., The values listed in Table 4 reflect these corrections and uncertainties. + Our calibration simulations allowed us to derive limiting detection 1nagnitudes as a function of separation. as shown iu Fiewe 6.," Our calibration simulations allowed us to derive limiting detection magnitudes as a function of separation, as shown in Figure 6." + Around 2\TASS E101. faint secondaries (AAT z 3) were generally missed at separations closer than (715. while AMI z 5.5 (6) and 1.5 (5) could be obtained for wide separations at the )) confideuce level at FaliW and F1012ML respectively.," Around 2MASS $-$ 1404, faint secondaries $\Delta$ M$\gtrsim$ 3) were generally missed at separations closer than $\farcs$ 15, while M $\approx$ 5.5 (6) and 4.5 (5) could be obtained for wide separations at the ) confidence level at F814W and F1042M, respectively." + Iu general. ouly ucar-equalinaguitude companions with «=0709 could be recovered better than of the time. as is the case for 2MÁASS 2052AD.," In general, only near-equal-magnitude companions with $a \lesssim +0{\farcs}09$ could be recovered better than of the time, as is the case for 2MASS $-$ 2952AB." + For «zοι. S/N = 7 limits (Table 5) wield the aN seusitivitv for faint companions. raneine frou AM = 136.9 at FALIW to AM = 2.9LO at FLOAT We can convert these values to mass ratio (¢= Ab ΑΕ) limits Wing a masslundnositv power law frou: Burrowsctal.(2001).. L X APSt. aud asstimine for simplicity cocvality and neeheible variation in bolometric corrections over thesamplet.. such that These values are listed in Table 5. aud range from 1.00.1.," For $a \gtrsim 0{\farcs}4$ S/N = 7 limits (Table 5) yield the maximum sensitivity for faint companions, ranging from $\Delta$ M = 4.3–6.9 at F814W to $\Delta$ M = 2.9–4.9 at F1042M. We can convert these values to mass ratio $q \equiv$ $_2$ $_1$ ) limits using a mass-luminosity power law from \citet{bur01}, L $\propto$ $^{2.64}$, and assuming for simplicity coevality and negligible variation in bolometric corrections over the, such that These values are listed in Table 5, and range from 1.0–0.4." + Overall. our. sample is complete to g20.1 for a2| AU (assuming a mean distauce of 10 pc). with less sensitivity for siiall mass ratios to separations approaching a5] AU.," Overall, our sample is complete to $q \gtrsim 0.4$ for $a \gtrsim +4$ AU (assuming a mean distance of 10 pc), with less sensitivity for small mass ratios to separations approaching $a \approx 1$ AU." + 2MASS 2739AB ids clearly resolved into two nunequalinaenitude componeuts im our UST iuages., 2MASS $-$ 2739AB is clearly resolved into two unequal-magnitude components in our HST images. + The colors of these objects are siguificautlv different. with the füuter companion having an F1012M color similar to the Ts Cliese 570D (Durgasseretal.2000b).. while the color of the primary is consistent with the spectral type of the combined syvstein. T6.," The colors of these objects are significantly different, with the fainter companion having an $-$ F1042M color similar to the T8 Gliese 570D \citep{me00a}, while the color of the primary is consistent with the spectral type of the combined system, T6." +" The magnitude ratios of this pair. AM gaiip = 1.5940.0L aud AM yypay = L05+0.03. are ercater than the absolute magnitude ratios of Cliese 229B (T6.5) aud Cliese 570D. AM esipy = LOLEOLO and Άλεν, = O.STEOLZ. consistent with 2\DASS 2739A being earlier than type T6.5."," The magnitude ratios of this pair, $\Delta$ $_{F814W}$ = $\pm$ 0.04 and $\Delta$ $_{F1042M}$ = $\pm$ 0.03, are greater than the absolute magnitude ratios of Gliese 229B (T6.5) and Gliese 570D, $\Delta$ $_{F814W}$ = $\pm$ 0.10 and $\Delta$ $_{F1042M}$ = $\pm$ 0.12, consistent with 2MASS $-$2739A being earlier than type T6.5." + Based on these colors. we speculate that this svstei is comprised of a T6 and Ts pair. which should be coufinned with spatiallv-resolved spectroscopy.," Based on these colors, we speculate that this system is comprised of a T6 and T8 pair, which should be confirmed with spatially-resolved spectroscopy." + No parallax has been ucasured for this svstem vot. but the spectrophotometric distance of the secondary. compared to (ποιο 570D. is Poop = Eypoy = Ul pe," No parallax has been measured for this system yet, but the spectrophotometric distance of the secondary, compared to Gliese 570D, is $d^B_{F814W}$ = $d^B_{F1042M}$ = 11.1 pc." + At JFband. 4] = 10.5. ane dB = 11.9. based on the absolute J-baud -wenitudes of the T6 SDSS 1621)0029 (Dahuctal.2002.M;=15.3340.07) and Gliese S70D (Durgasseretal.2000b.M;= 16.1740.05)..," At J-band, $d^A_{J}$ = 10.8 and $d^B_{J}$ = 11.9, based on the absolute J-band magnitudes of the T6 SDSS 1624+0029 \citep[M$_J$ = +15.33$\pm$0.07]{dah02} and Gliese 570D \citep[M$_J$ = +16.47$\pm$0.05]{me00a}." + All of these distauce estimates combined vicld a inean d37 = LL2+0.5 pe and projected plivsical separation à = 3.1740.11 AU: note that the uncertainties do not represeut probable scatter in the absolute maecnitude/spectral type relation. not vot adequately measured for the T dwarfs.," All of these distance estimates combined yield a mean $d^{AB}$ = $\pm$ 0.5 pc and projected physical separation $a$ = $\pm$ 0.14 AU; note that the uncertainties do not represent probable scatter in the absolute magnitude/spectral type relation, not yet adequately measured for the T dwarfs." +" Adopting Ty, z 1000 and SOO IX for the twocomponcuts(Durgasseroet 2002d).. assmnuiug coevalitv. aud using the models of Burrowsctal. (1997).. we can derive component niasses for ages of 0.5.1.0. and 5.0 Cir. as listed in Table 6."," Adopting $_{eff}$ $\approx$ 1000 and 800 K for the twocomponents\citep{me02}, assuming coevality, and using the models of \citet{bur97}, , we can derive component masses for ages of 0.5,1.0, and 5.0 Gyr, as listed in Table 6." + The derived mass ratio of this svstem is 4 = 0.7.0.8. depending on its age.," The derived mass ratio of this system is $q$ = 0.7–0.8, depending on its age." +" Asstning that. on average. the seminajor axis of a binary svstoeni <04,,51.20«a> (Fischer& 1992).. we estimate orbital periods of 2110 vr:"," Assuming that, on average, the semimajor axis of a binary system $ = 1.26$ \citep{fis92}, , we estimate orbital periods of 24–40 yr;" +The luminosities of quasars and other active galactic nuclei (AGN) have been observed. to vary on. time-scales from hours to decades. and from. X-ray to radio wavelengths.,"The luminosities of quasars and other active galactic nuclei (AGN) have been observed to vary on time-scales from hours to decades, and from X-ray to radio wavelengths." + The majority of quasars exhibit continuum variability. on the order of 20 per cent on time-scales of months to vears (e.g.Hook.MeMabhon.Bovle.&Lewin1994:Van-cenBerketal. 2004).," The majority of quasars exhibit continuum variability on the order of 20 per cent on time-scales of months to years \citep[e.g.,][]{hook94,vandenberk04}." +. In fact. variability has long been used as a selection. criterion in creating quasar samples from photometric data (e.g..Ixoo.Ίντο.&Cuchvorth1986:]veziéetal.2004:Hengstorf 2004).," In fact, variability has long been used as a selection criterion in creating quasar samples from photometric data \citep[e.g.,][]{ koo86, ivezic04, rengstorf04}." +. Alany simple correlations between photometric variability ancl various physical parameters have been known for decades., Many simple correlations between photometric variability and various physical parameters have been known for decades. + These relationships are sumamarised by Helfandetal.(2001). and Civeonetal.(1999)., These relationships are summarised by \citet{helfand01} and \citet{giveon99}. +.. Numerous studies (e$.Hawkins2002:deVries.Decker.&White2003) have shown variability to correlate with time lag.," Numerous studies \citep[e.g.,][]{hawkins02,devries03} have shown variability to correlate with time lag." + AXnti-correlations have been found between variability and. luminosity and wavelength (e.g...Giveonetal.L999:Trevese.&Dunone 2001).. VandenBerketal.(2004.hereafter VBO4).," Anti-correlations have been found between variability and luminosity and wavelength \citep[e.g.,][] {giveon99, trevese01}. \citet[hereafter VB04]{vandenberk04}," +. using a sample of 25.000 quasars. confirmed these known correlations. and. parametrized relationships between variability and time lag. Iuminosity. rest-frame wavelength and. recishift.," using a sample of $\sim25,000$ quasars, confirmed these known correlations, and parametrized relationships between variability and time lag, luminosity, rest-frame wavelength and redshift." + Our understanding of the physies of t1e central. black hole in quasars and active galactic nuclol (AGN) has long been tied to the variability of the quasars Luminosity., Our understanding of the physics of the central black hole in quasars and active galactic nuclei (AGN) has long been tied to the variability of the quasar's luminosity. +" Intra-day variability in X-ray ane optical light (see.e.g.""Terrell1967:Winman1968:Bollerctal.1997) »oint towards a compact object. specifically a supermassive black hole. at the centre o Tan AGN."," Intra-day variability in X-ray and optical light \citep[see, e.g.,][]{terrell67, kinman68, boller97} point towards a compact object, specifically a supermassive black hole, at the centre of an AGN." + More recenIv. reverberation mapping techniques have been used to determine the radius of the broad line region ancl. indirectly. to measure the mass of the central black hole (Blandford&Melxee1982:Peterson1993:]xaspietal. 2000).," More recently, reverberation mapping techniques have been used to determine the radius of the broad line region and, indirectly, to measure the mass of the central black hole \citep{blandford82, peterson93,kaspi00}." +. Recently. Wold.Brotherton.&Shang(2007).. presented evidence sugeesting that the photometric variability of quasars is linked. to the mass of the central black hole.," Recently, \citet{wold07}, presented evidence suggesting that the photometric variability of quasars is linked to the mass of the central black hole." + Strength of variability for LOO quasars was approximated by finding the greatest sinele-epoch {ρα deviation from, Strength of variability for $\sim$ 100 quasars was approximated by finding the greatest single-epoch $R$ -band deviation from +"where we have defined: €,=cos2(a,0) and S,=sin2(a,8).",where we have defined: $C_{x}=\cos 2(\alpha_{x}-\theta)$ and $S_{x}=\sin 2(\alpha_{x}-\theta)$. + There are 9 unknown parameters and three independent equations., There are $9$ unknown parameters and three independent equations. + Thus in principle any one of S. Z.or £ can be determined from a knowledge of the other two.," Thus in principle any one of $\cal +S$, $\cal I$ ,or $\cal L$ can be determined from a knowledge of the other two." + In particular the source parameters. S. can be determined if one knows the image. Z. and lens parameters. Z.," In particular the source parameters, $\cal S$, can be determined if one knows the image, $\cal I$, and lens parameters, $\cal L$." + In what follows we shall also use the parameter e‘describing the source and image ellipticity defined in terms of A and AX as lt is interesting to note. before making anv attempt to solve these equations. that they have a number of scaling laws which should make their solution easier.," In what follows we shall also use the parameter $e$ describing the source and image ellipticity defined in terms of $\lambda$ and $\Delta\lambda$ as It is interesting to note, before making any attempt to solve these equations, that they have a number of scaling laws which should make their solution easier." +" Thus for a &given lens. if (A;:2NÀ;0;À,:2NÀ0.) is à solution. then so is GrA;rAd;0;wr.PAA.6) where wis a constant."," Thus for a given lens, if $(\lambda_{i}, +\Delta\lambda_{i}, \alpha_{i}, \lambda_{s}, \Delta\lambda_{s}, +\alpha_{s} )$ is a solution, then so is $(x\lambda_{i}, +x\Delta\lambda_{i}, \alpha_{i}, x\lambda_{s}, x\Delta\lambda_{s}, +\alpha_{s} )$ where $x$ is a constant." + Àloreover. for a given image. if (A.AAS.νι8.5.0) is a solution. then so also is (yÀ..JFAd.ousgsgn. 8).," Moreover, for a given image, if $(\lambda_{s}, \Delta\lambda_{s}, \alpha_{s}, \kappa, \gamma, +\theta)$ is a solution, then so also is $(y^2\lambda_{s}, +y^2\Delta\lambda_{s}, \alpha_{s}, y\kappa, y\gamma, \theta)$ ." + Vhis means that one can arbitrarily take A; and 5s to be 1., This means that one can arbitrarily take $\lambda_{i}$ and $\kappa$ to be $1$. +" Dividing the equations (6))-(9)) by &7À;2 Gr=L/2NAÀ; and y= L/&s) one obtains: where Al;=AA)λε fo=AL(ON) Alo=AX, "," Dividing the equations \ref{s1}) \ref{s4}) ) by $\kappa^2\lambda_{i}$ $x=1/\Delta\lambda_{i}$ and $y=1/\kappa$ ) one obtains: where $\Delta l_{i}= \Delta\lambda_{i}/\lambda_{i}$, $l_{s}=\lambda_{s}/(\kappa^2\lambda_{i})$, $\Delta l_s = +\Delta\lambda_{s}/(\kappa^2\lambda_{i})$, and $\Gamma=\gamma/\kappa$." +The determination of the lens parameters in the weak lensing regime is usually treated as à statistical problem., The determination of the lens parameters in the weak lensing regime is usually treated as a statistical problem. + In the absence of additional information arout the source galaxies. in particular about the orientation of their axis of symmetry (if they have one). most analyses to date have attented to solve the problem of inferring the lens parameters by matching he distribution of source parameters obtained. [rom the data by the lens model. with an assumed. distribution of source galaxy parameters e.g. a uniform distribution of orienations of source galaxies.," In the absence of additional information about the source galaxies, in particular about the orientation of their axis of symmetry (if they have one), most analyses to date have attempted to solve the problem of inferring the lens parameters by matching the distribution of source parameters obtained from the data by the lens model, with an assumed distribution of source galaxy parameters e.g. a uniform distribution of orientations of source galaxies." + One can obtain the best Gt model parameters w standard. routines such as X7- fittinge., One can obtain the best fit model parameters by standard routines such as $\chi^2$ - fitting. + This sort of approach has been adopted by. for example. (Bonnetctal.1993).," This sort of approach has been adopted by, for example, \cite{Bon93}." +. Although very. promising. the measurements are not Dcrood enough to provide accurate determination of the lens parameters.," Although very promising, the measurements are not good enough to provide accurate determination of the lens parameters." + The techniques developed to reconstruct the cluster xotential from. the lensecl galaxies are now extremly sophisticated anc jev treat with great. care the observational and statistical uncertainties ((Ixaiseretal.1995).. (Schneider1996).. 990)3).," The techniques developed to reconstruct the cluster potential from the lensed galaxies are now extremly sophisticated and they treat with great care the observational and statistical uncertainties \cite{Kai95}, \cite{Sch96}, \cite{Koch90}) )." + We do not intend in this paper to &o into suc1 details., We do not intend in this paper to go into such details. + Our point is just to show that polarimetric measurements put garong constraints on the cluster potential reconstuction., Our point is just to show that polarimetric measurements put strong constraints on the cluster potential reconstuction. + We have alreacky argued. that information about the source galaxies. and. in particular about their orientation on the sky. which could be obtained from measurement of ticir optical polarization. would. lead to more accurate determinations.," We have already argued that information about the source galaxies, and in particular about their orientation on the sky, which could be obtained from measurement of their optical polarization, would lead to more accurate determinations." + Polarization measurements would undoubtedly be «lifficult for such faint galaxies., Polarization measurements would undoubtedly be difficult for such faint galaxies. + However. even the measurement of 1 polarization for a small subset of the galaxy sample can considerably improve the accuracy on the lens parameters determination.," However, even the measurement of the polarization for a small subset of the galaxy sample can considerably improve the accuracy on the lens parameters determination." + In the following section we illustrate tiis numerically for simple lens models., In the following section we illustrate this numerically for simple lens models. + ‘To assess what can be obtained via polarization rneasurements. we have carried out a simple simulation.," To assess what can be obtained via polarization measurements, we have carried out a simple simulation." + First we decide on a lens mocel. which fixes the values of &(r.gy). 5Geey) and 66r.5) in the lens plane.," First we decide on a lens model, which fixes the values of $\kappa(x,y)$, $\gamma(x,y)$ and $\theta(x,y)$ in the lens plane." + We then randomly generate a sample of source galaxies from a given distribution with the Following properties., We then randomly generate a sample of source galaxies from a given distribution with the following properties. + We take the distribution ofthe symmetry axis to he line of sight to be isotropic. which implies a unifxm cistribution in orientation on the sky. ane uniform distribution in ellipticity. (if. we assume the galactic dises to be circidar).," We take the distribution of the symmetry axis to the line of sight to be isotropic, which implies a uniform distribution in orientation on the sky, and uniform distribution in ellipticity (if we assume the galactic discs to be circular)." +" Since we are going te> use the svstem (10))-(13)) which contains only ""reduced? variables. we do not need to know the! brightness of the galaxies at this stage."," Since we are going to use the system \ref{sr1}) \ref{sr4}) ) which contains only “reduced” variables, we do not need to know the brightness of the galaxies at this stage." +ποσο galaxies are then lensecl ov a &iven lens model to give a clistribution of images.,These galaxies are then lensed by a given lens model to give a distribution of images. + We then attempt to reestablish the lens parameters from the sample of images., We then attempt to reestablish the lens parameters from the sample of images. + We shall see below that the ellectiveness of th10 retrieval depends crucially on the assumptions about the model. and whether or not we usepotential information about the orientation of the source galaxies. which we argue. could be furnished w polarimetric data.," We shall see below that the effectiveness of the retrieval depends crucially on the assumptions about the model, and whether or not we usepotential information about the orientation of the source galaxies, which we argue, could be furnished by polarimetric data." +is a scale radius. contains most of the mass. and the compact nuclear component X44€8) contains a small fraction (Ferrareseetal.2006b:Wehner&Harris2006). of the mass.,"is a scale radius, contains most of the mass, and the compact nuclear component $\Sigma_{\rm nucl}(R)$ contains a small fraction \citep{Ferrarese:06b,Wehner:06} of the mass." + In spheroidals the Sérrsic index varies in the range n~1—2 (Ferrareseetal.2006a:Kormendy2009).," In spheroidals the Sérrsic index varies in the range $n\sim 1-2$ \citep{Ferrarese:06a,Kormendy:09}." +. We work within the paradigm in which a stellar nucleus is not preser= prior to the formation of most of the galactic stars., We work within the paradigm in which a stellar nucleus is not present prior to the formation of most of the galactic stars. + Thus. we consider an initial galaxy with a spherically-averaged stellar density profile that lacks a nucleus.," Thus, we consider an initial galaxy with a spherically-averaged stellar density profile that lacks a nucleus." + We employ an approximation to the deprojected Sérrsic profile in the form (Prugniel&Simien1997;LimaNetoetal.1999) pam 9/7102) where(i) p = | - 0.6097 + 0.0546377.," We employ an approximation to the deprojected Sérrsic profile in the form \citep{Prugniel:97,LimaNeto:99} + = ], where p = 1 - 0.6097 + 0.05463." +13 )Here. n is normalized to the total stellar mass of the galaxy May=Amdοι(Τά2paVIGAznRipin. and Pa)=fyr7e7 is the gamma function.," Here, $\rho_{\rm sph}$ is normalized to the total stellar mass of the galaxy $M_{\rm sph}=4\pi\int_0^\infty\rho_{\star,0}(r)r^2dr=4\pi n R_{\rm s}^3 +\rho_{\rm sph}\Gamma[(3-p)n]$, and $\Gamma(a)=\int_0^\infty +t^{a-1}e^{-t}dt$ is the gamma function." + Quantities pertaining to the initial galaxy.dt before migration and disruption of clusters. are denoted by the subseript 0.," Quantities pertaining to the initial galaxy, before migration and disruption of clusters, are denoted by the subscript $0$." +" The stellar mass within radius ris ο) = απο ""sy where οί.1)=fo17Vedt is the lower incomplete Gamma function."," The stellar mass within radius $r$ is (r) = n ^3 ], where $\gamma(a,x)=\int_0^x t^{a-1}e^{-t}dt$ is the lower incomplete Gamma function." + The total initial density profile includes a dark matter component ην ου where we model the dark matter density profile with the NFW law (Navarroetal.1997) ppyG)2pxflOtrfryY] with scale radius 74. virial radius η. and concentration c—rar ," The total initial density profile includes a dark matter component _0(r) = +, where we model the dark matter density profile with the NFW law \citep{Navarro:97} $\rho_{\rm DM}(r) = \rho_s/[(r/r_s)(1 + r/r_s)^2]$ with scale radius $r_{\rm s}$, virial radius $r_{\rm vir}$ , and concentration $c\equiv r_{\rm vir}/r_{\rm s}$." +Spherotdals are often dark matter dominated. especially at the low-luminosity end. Le. pom)Zpel).," Spheroidals are often dark matter dominated, especially at the low-luminosity end, i.e., $\rho_{\rm DM}(r)\gtrsim \rho_{\star,0}(r)$." + Thus. for galaxies in which dark matter is the dominant mass component. the cluster migration time can be modeled with the N-body-simulation-calibrated Chandrasekhar dynamical friction torque in equation (7)).," Thus, for galaxies in which dark matter is the dominant mass component, the cluster migration time can be modeled with the $N$ -body-simulation-calibrated Chandrasekhar dynamical friction torque in equation \ref{eq:torque_spheroid}) )." + The surface density profiles oflate-type spiral disk galaxies can be fit with a Sérrsic-type exponential stellar disk and a nuclear component (Bokeretal.2003) as in equation (11))., The surface density profiles oflate-type spiral disk galaxies can be fit with a Sérrsic-type exponential stellar disk and a nuclear component \citep{Boker:03} as in equation \ref{eq:stellar_profile}) ). + We consider “pure” disk galaxies. which do not have a stellar bulge or à preexisting pseudobulge: we do not consider cluster formation in the galactic halo.," We consider “pure” disk galaxies, which do not have a stellar bulge or a preexisting pseudobulge; we do not consider cluster formation in the galactic halo." +" The initial stellar density profile without the nuclear component is then (R)2 Á]. where “ais,=Mui[2znR2T(220)."," The initial stellar density profile without the nuclear component is then = ], where $\Sigma_{\rm disk} =M_{\rm disk} / +[2\pi n R_{\rm s}^2 \Gamma(2n)]$." +" The corresponding initial stellar mass profile is o(R) = ""| (O7)where Eta.)=I(a)—ία.) is 2zxnthe upper incomplete Gamma function."," The corresponding initial stellar mass profile is (R) = n ], where $\Gamma(a,x)=\Gamma(a)-\gamma(a,x)$ is the upper incomplete Gamma function." +" The total initial density profile is then RIES REN, where XpytR) is the projected NEW profile.", The total initial density profile is then _0(R) = + where $\Sigma_{\rm DM}(R)$ is the projected NFW profile. + Disk galaxies generally have an exponential profile with Sérrsic index η.»I., Disk galaxies generally have an exponential profile with Sérrsic index $n\sim1$. + The total density profiles in disk galaxies may be dominated either by dark matter or by luminous matter in their central regions., The total density profiles in disk galaxies may be dominated either by dark matter or by luminous matter in their central regions. + Even if dark matter dominates the spherically averaged density profile. in disk galaxies it provides a relatively smaller contribution to the torque if the cluster is embedded in the disk.," Even if dark matter dominates the spherically averaged density profile, in disk galaxies it provides a relatively smaller contribution to the torque if the cluster is embedded in the disk." + This is because the torque coupling of the migrating star cluster with the flattened disk. equation (8)). is more significant than that with the dark matter halo (see.e.g..Bekki2010).," This is because the torque coupling of the migrating star cluster with the flattened disk, equation \ref{eq:torque_disk}) ), is more significant than that with the dark matter halo \citep[see, e.g.,][]{Bekki:10}." +. We use the disk torque to model the migration time in disk galaxies., We use the disk torque to model the migration time in disk galaxies. +" We assume that all stars form in clusters with masses between Ming, and Mag, (SeeSection 2.1).", We assume that all stars form in clusters with masses between $M_{\rm min}$ and $M_{\rm max}$ (seeSection \ref{subsec:icmf}) ). + Let dn/dMdr be the number of clusters per unit initial cluster. mass per unit radius that have formed in a galaxy such that «Cn/dMdrydr is proportional to the ICMF in equation (1).Ju while MMM = We sample cluster masses and initial radii randomly according to thedistribution d7n/dMdr.," Let $d^2n/dMdr$ be the number of clusters per unit initial cluster mass per unit radius that have formed in a galaxy such that $\int_0^\infty\, +(d^2n/dMdr)\,dr$ is proportional to the ICMF in equation \ref{eq:icmf}) ), while MdM = We sample cluster masses and initial radii randomly according to thedistribution $d^2n/dMdr$." + The total number of clusters in. our simulations. 1s m(dn/dM)dM--Mi/[MgagyblntMaMal where Mu is the total stellar mass. My or Mai;. of the galaxy.," The total number of clusters in our simulations is $ +\int_{M_{\rm min}}^{M_{\rm max}} (dn/dM)\,dM \approx M_{\rm tot} / [M_{\rm min}\ln(M_{\rm max}/M_{\rm min})], +$ where $M_{\rm tot}$ is the total stellar mass, $M_{\rm sph}$ or $M_{\rm disk}$, of the galaxy." +" At first. we do not explicitly take into account the possible prompt. possibly mass-independent (foradiscussionofmass-dependentpromptdis-solutionsee.e.g..Goodwin2009) dissolution of clusters (""infantmortality.”see.e.g.Falletal.2005:Bastian&Lamers 2009)."," At first, we do not explicitly take into account the possible prompt, possibly mass-independent \citep[for a discussion + of mass-dependent prompt dissolution see, e.g.,][]{Goodwin:09} + dissolution of clusters \citep[``infant mortality,'' see, + e.g.,][]{Fall:05,Bastian:06,Chandar:06,Goodwin:06,Gieles:07,Pellerin:07,Whitmore:07,Baumgardt:08,deGrijs:08,Lamers:09}." +. Then. we briefly assess the effects of prompt dissolution.," Then, we briefly assess the effects of prompt dissolution." +" We proceed to model cluster migration and dissolution by following the cluster orbital decay di=-—— ~~(20)*and mass. loss IMT where μις and jj, are. respectively. the migration time scale in equation (6)) and dissolution time scale in equation (9))."," We proceed to model cluster migration and dissolution by following the cluster orbital decay = -, and mass loss =, where $t_{\rm mig}$ and $t_{\rm dis}$ are, respectively, the migration time scale in equation \ref{eq:time_mig}) ) and dissolution time scale in equation \ref{eq:time_dis}) )." + We calculate the orbital decay of clusters in the order of increasing migration time (see equation [6]]) evaluated before migration has occurred., We calculate the orbital decay of clusters in the order of increasing migration time (see equation \ref{eq:time_mig}] ]) evaluated before migration has occurred. + We keep track of the mass tidally stripped fromthe cluster with Ath shortest migration time by caleulating its contribution to the stellar density profile Gm , We keep track of the mass tidally stripped fromthe cluster with $k$ th shortest migration time by calculating its contribution to the stellar density profile = . ( +rgo where ο 15 the cluster formation radius and k21.2.... Mor.,"r_k< ) , where $r_{k,0}$ is the cluster formation radius and $k=1,2,\ldots +n_{\rm tot}$ ." + We assume that a cluster has been fully disrupted when its, We assume that a cluster has been fully disrupted when its +vieldiug Tere τσηPN is. the ummber of. iiergers. of. a halo wit.[um nass between UMB[DMo aud Af|4M» that merec with a halo of mass M4 in a redshift interval εἰ.,yielding Here $\left.\frac{d^2N}{dz dM_2}\right|_{M_1}$ is the number of mergers of a halo with mass between $M_2$ and $M_2 + dM_2$ that merge with a halo of mass $M_1$ in a redshift interval $dz$. + The assumption that barvous shock during cach uerecr is optinustic., The assumption that baryons shock during each merger is optimistic. + While ierecrs will certainly increase the virial euergv per barvon. it is uot clear row the barvous would be recvcled through virialization shocks ou multiple occasions if they cool efficiently. to xoduce the ionizing radiation.," While mergers will certainly increase the virial energy per baryon, it is not clear how the baryons would be recycled through virialization shocks on multiple occasions if they cool efficiently to produce the ionizing radiation." + Tf the baryous cool frou he halo iuto a disk. they will not be heated to the virial cluperature of the halo (unless they have been expelled hrough some feedback process) and so should not warticipate iu the halo shock of the next merecr.," If the baryons cool from the halo into a disk, they will not be heated to the virial temperature of the halo (unless they have been expelled through some feedback process) and so should not participate in the halo shock of the next merger." + Thus. our results from equation (1)) represent the masini ionisation flux possible from the fast accretion shock uechanisui sunmnunuiuisediuiequatiou (1)).," Thus, our results from equation \ref{merge}) ) represent the maximum ionisation flux possible from the fast accretion shock mechanism summarised in equation \ref{Ngamma}) )." +" Results are shown as a function of redshift in the upper right panel of Figure 1.. assuming minimum halo nasses corresponding to (i,10kms+ aud eg30kms (again with no discernible difference owing o the domiuauce of massive halos}."," Results are shown as a function of redshift in the upper right panel of Figure \ref{plot1}, , assuming minimum halo masses corresponding to $v_{\rm vir}=10~{\rm km~s^{-1}}$ and $v_{\rm vir}=30~{\rm km~s^{-1}}$ (again with no discernible difference owing to the dominance of massive halos)." + We fiud that the re-xocessiueg of barvous through shocks im multiple mergers increases the ionising photon output bv a factor of a few relative to the collapsed fraction calculation., We find that the re-processing of baryons through shocks in multiple mergers increases the ionising photon output by a factor of a few relative to the collapsed fraction calculation. + However. we still find that ouly ~5 percent of the ICAL is reionized by τονδν ας 9P( atcδι indicating that shocks provide insufficient iouisine buuinositv to reionize the ICAL," However, we still find that only $\sim5$ percent of the IGM is reionized by $z\sim6$, and $\sim1\%$ at $z\sim8$, indicating that shocks provide insufficient ionising luminosity to reionize the IGM." + The difference between our findings and the results oxeseuted in Dopitaetal.(2011) originates partly from. he adoption of large value of σς=0.9 in that work. aud yartly from an errora in the caleulatiou method (Lawrence Irauss private commuuication).," The difference between our findings and the results presented in \citet[][]{Dopita2011} originates partly from the adoption of a large value of $\sigma_8=0.9$ in that work, and partly from an error in the calculation method (Lawrence Krauss, private communication)." +" Dopitactal.(2011) calculate the accretion rate and corresponding ijonisiug ""uimositv as a function of halo mass (see their Figure 2).", \citet[][]{Dopita2011} calculate the accretion rate and corresponding ionising luminosity as a function of halo mass (see their Figure 2). + They then iuteerate over the niass-£function and redshift., They then integrate over the mass-function and redshift. + This procedure effectively sets the rate at which eas in he halo doubles to be equal to the inverse of dynamical iue at the virial radius (which is shorter bv an order of magnitude than the IIubble time). aud so does uot account for the duty-cvele of the shocks @vlich should ο only 0.1).," This procedure effectively sets the rate at which gas in the halo doubles to be equal to the inverse of dynamical time at the virial radius (which is shorter by an order of magnitude than the Hubble time), and so does not account for the duty-cycle of the shocks (which should be only $\sim0.1$ )." + As a result. the calculation in Dopitaetal.(2011) aceretes an order of magnitude more gas han available per halo.," As a result, the calculation in \citet[][]{Dopita2011} + accretes an order of magnitude more gas than available per halo." + Iun the previous subsections we have estimated the uunuber of iouising photons available per hydrogen iu the ICM., In the previous subsections we have estimated the number of ionising photons available per hydrogen in the IGM. + We next estimate the niaxianuuΠΟΟΙ of photous, We next estimate the maximumnumber of photons +Evidence for supersolar abundance raios of a-elements Oo iron,Evidence for supersolar abundance ratios of $\alpha$ -elements to iron +or vanishing polarization and that ‘cloud’ 1 covers both this unpolarized portion of the jet and that part of the jet which emits polarized radiation (the faint blue section of the source in Fig. 5):,or vanishing polarization and that `cloud' 1 covers both this unpolarized portion of the jet and that part of the jet which emits polarized radiation (the faint blue section of the source in Fig. ): + the covering factor for the latter is C3., the covering factor for the latter is $C_{3}$. +" We further assume the presence of a velocity gradient in ‘cloud’ 1 such that the velocity centroid for the Stokes J spectrum is again vj, but for the AP/P spectrum is now 13 (see Fig. 5)."," We further assume the presence of a velocity gradient in `cloud' 1 such that the velocity centroid for the Stokes $I$ spectrum is again $v_{1}$, but for the $\Delta P/P$ spectrum is now $v_{3}$ (see Fig. )." + As a result we now have where 73(v) is the 21 cm optical depth of the gas in cloud 1 toward the polarized portion of the jet source: the velocity centroid of this gas is v3., As a result we now have where $\tau_{3}(v)$ is the 21 cm optical depth of the gas in cloud 1 toward the polarized portion of the jet source: the velocity centroid of this gas is $v_{3}$. +" Therefore, this model has velocity components centered on v1,v2,andv3."," Therefore, this model has velocity components centered on $v_{1}, v_{2}, {\rm and} \ v_{3}$." +" Since we shall also consider fits to the Ax spectrum in this case, we develop a model for the intrinsic change in polarization position angle across the jet source."," Since we shall also consider fits to the $\Delta \chi$ spectrum in this case, we develop a model for the intrinsic change in polarization position angle across the jet source." + We assume that x=x4 for the polarized flux density emitted by the fraction of the jet source that isnot incident on the cloud and x=x; for the fraction of emitted polarized flux density that is incident on the cloud., We assume that $\chi=\chi_{a}$ for the polarized flux density emitted by the fraction of the jet source that is incident on the cloud and $\chi=\chi_{b}$ for the fraction of emitted polarized flux density that is incident on the cloud. +" Noting that the net flux densities in Stokes Q and U parameters are given by Q=Q,+Q, and U=U,4+Uj, and x=0.5 xarctan(U/Q), we find that Here r=< Το>p/< Το>a, where «Zo>, and » are the surface brightnesses in Stokes Q averaged over regions a and b respectively, and we make use of the relation Q,/Q,= C3)."," Noting that the net flux densities in Stokes $Q$ and $U$ parameters are given by $Q=Q_{a}+Q_{b}$ and $U=U_{a}+U_{b}$, and $\chi$ $U/Q$ ), we find that Here $r$ $<{\cal I}_{Q}>_{ b}$ $<{\cal I}_{Q}>_{a}$, where $<{\cal I}_{Q}>_{ a}$ and $<{\cal I}_{Q}>_{b}$ are the surface brightnesses in Stokes $Q$ averaged over regions a and b respectively, and we make use of the relation $Q_{b}/Q_{a}=rC_{3}/(1-C_{3})$ ." +" The observed value of x averaged over the unattenuated continuum source, χσοπι, is obtained from the last equation by setting 75-0."," The observed value of $\chi$ averaged over the unattenuated continuum source, $\chi_{\rm cont}$, is obtained from the last equation by setting $\tau_{3}$ =0." + The difference between x in the line and xcoy is denoted by Ax(v)zx—Xcont-, The difference between $\chi$ in the line and $\chi_{\rm cont}$ is denoted by $\Delta \chi(v)$$\equiv$$\chi-\chi_{\rm cont}$. + We fitted the three-component model parameters to the spectra as follows., We fitted the three-component model parameters to the spectra as follows. +" First, observing that the parameters characterizing the Stokes ΔΙ/Ι spectrum are independent of those characterizing the AP/P and Ax spectra, we did not fit the three Stokes spectra simultaneously."," First, observing that the parameters characterizing the Stokes $\Delta I/I$ spectrum are independent of those characterizing the $\Delta P/P$ and $\Delta \chi$ spectra, we did not fit the three Stokes spectra simultaneously." +" Rather, we fitted the AI/I spectrum alone, and then the AP/P and Ax spectra simultaneously."," Rather, we fitted the $\Delta I/I$ spectrum alone, and then the $\Delta P/P$ and $\Delta \chi$ spectra simultaneously." +" Given the reasonable fit to the ΔΙ/1 spectrum discussed in 8 5.2, we adopt the parameters found for clouds 1 and 2 in that fit here."," Given the reasonable fit to the $\Delta I/I$ spectrum discussed in $\S$ 5.2, we adopt the parameters found for clouds 1 and 2 in that fit here." +" Second, in the case of the polarized spectra we need to find the gas parameters, Ny13/T;3, vs, 7,3, and C3 as well as the source parameters r, χα, and Xp."," Second, in the case of the polarized spectra we need to find the gas parameters, $N_{{\rm HI},3}/T_{s,3}$, $v_{3}$, $\sigma_{v,3}$, and $C_{3}$ as well as the source parameters $r$, $\chi_{a}$, and $\chi_{b}$." +" For the source parameters, we are guided by single-dish polarization measurements above 1.5 GHz that show the continuum polarization position angle for the entire source Xcont=33+5° (Tabara and Inoue 1980)."," For the source parameters, we are guided by single-dish polarization measurements above 1.5 GHz that show the continuum polarization position angle for the entire source $\chi_{\rm cont}$ $\pm$ $^{o}$ (Tabara and Inoue 1980)." + The VLBI polarization data at 5 GHZ (Jiang 1996; Cotton 1997; Cotton 2010 priv., The VLBI polarization data at 5 GHZ (Jiang 1996; Cotton 1997; Cotton 2010 priv. + comm.), comm.) +" further show that along the jet x decreases with increasing distance from the core source (i.e., in the SW direction) by as much as ~ 40? which is consistent with the orientation of the polarization bars in Fig5."," further show that along the jet $\chi$ decreases with increasing distance from the core source (i.e., in the SW direction) by as much as $\approx$ $^{\circ}$ which is consistent with the orientation of the polarization bars in Fig." + We found that the Ax(v) line spectral feature is more naturally explained by a large central optical depth το combined with a low covering factor C3 rather than vice versa.," We found that the $\Delta \chi(v)$ line spectral feature is more naturally explained by a large central optical depth $\tau_{0,3}$ combined with a low covering factor $C_{3}$ rather than vice versa." +" By contrast, C4 and 793 are degenerate in the case of the AP/P spectrum (see eq. ("," By contrast, $C_{3}$ and $\tau_{0,3}$ are degenerate in the case of the $\Delta P/P$ spectrum (see eq. (" +8)).,8)). + We used mrqmin to determine the cloud 1 and cloud 2 parameters by fits to the AI/I spectrum., We used $mrqmin$ to determine the cloud 1 and cloud 2 parameters by fits to the $\Delta I/I$ spectrum. +" As discussed above, the parameters of the third velocity component were found by separately fitting the AP/P and Ax spectra together."," As discussed above, the parameters of the third velocity component were found by separately fitting the $\Delta P/P$ and $\Delta \chi$ spectra together." +" This is possible since the polarized spectra are independent of the cloud 1 and 2 while the Stokes A7/I spectrum is independent of the parameters,component 3 parameters."," This is possible since the polarized spectra are independent of the cloud 1 and 2 parameters, while the Stokes $\Delta I/I$ spectrum is independent of the component 3 parameters." + After trials with several values of C3 we fixed C3=0.2 at the outset., After trials with several values of $C_{3}$ we fixed $C_{3}$ =0.2 at the outset. + Larger values of C3 must be accompanied by smaller values in 793 ," Larger values of $ C_{3}$ must be accompanied by smaller values in $\tau_{0,3}$ " +Weak gravitational lensing refers to the coherent distortions of background galaxy images by mass structures along the line of sight.,Weak gravitational lensing refers to the coherent distortions of background galaxy images by mass structures along the line of sight. + Lensing measurements [rom imaging surveys have emerged as a powerlul probe of cosmology., Lensing measurements from imaging surveys have emerged as a powerful probe of cosmology. +" With planned surveys that will cover thousands of square degrees, the staustical errors on measured shear correlations will be extremely small (e.g. the Dark Energy Survey (Abbott 2005).. PanSTARRS (Kaiser2004)., LSST (Starrοἱαἱ.2002),, and SNAP (Lamptonοἱef. 2002)))."," With planned surveys that will cover thousands of square degrees, the statistical errors on measured shear correlations will be extremely small (e.g. the Dark Energy Survey \citep{DES}, PanSTARRS \citep{PanSTARRS}, LSST \citep{LSST}, and SNAP \citep{SNAP}) )." + However systematic errors may exceed the statistical errors and dominate the error budget on cosmological parameters., However systematic errors may exceed the statistical errors and dominate the error budget on cosmological parameters. +" To analyze the elTect of errors on lensing statistics, we will consider the two-point correlation functions of the shear £4(0).£4.(0) and the shear power spectrum C.(0)."," To analyze the effect of errors on lensing statistics, we will consider the two-point correlation functions of the shear $\xi_{\gamma+}(\theta), \xi_{\gamma-}(\theta)$ and the shear power spectrum $C_\gamma(\ell)$." +" Other statistics often used in lensing measurements, such às the aperture mass variance and the top-hat shear variance,can be obtained by"," Other statistics often used in lensing measurements, such as the aperture mass variance and the top-hat shear variance,can be obtained by" +1992).,. +. The wealth of data aud proxiudtv of IC 312 permit the connection of cloud properties with star formation at sub-CAIC spatial scales., The wealth of data and proximity of IC 342 permit the connection of cloud properties with star formation at sub-GMC spatial scales. + The data published here represent the first hieh resolution maps of IEC4N(5. 1) and IIC4N(16 15) im an external galaxy., The data published here represent the first high resolution maps of $_3$ N(5–4) and $_3$ N(16–15) in an external galaxy. + Excitation is a kev component in the interpretation of molecular line intensities., Excitation is a key component in the interpretation of molecular line intensities. + So these two laps are compared with the previously published lower resolution UC'sNC109) map mace with the Owens Valley Millimeter Array (Meier&Turner2005).. to coustrain the plivsical conditions of the densest component of the ISM in the ceuter of IC312. aud to correlate properties of the dense gas with star formation. diffuse eas. aud chemistry.," So these two maps are compared with the previously published lower resolution $_{3}$ N(10–9) map made with the Owens Valley Millimeter Array \citep[][]{MT05}, to constrain the physical conditions of the densest component of the ISM in the center of IC342, and to correlate properties of the dense gas with star formation, diffuse gas, and chemistry." + Aperture svuthesis observations of the Που J=h1 rotational line (15.190316. GIIz) towards IC 312 were made with the D coufieuration of the VLA on 2005 November 25 (VLA ID: AAI839)., Aperture synthesis observations of the $_{3}$ N J=5–4 rotational line (45.490316 GHz) towards IC 342 were made with the D configuration of the VLA on 2005 November 25 (VLA ID: AM839). + The svuthesized beam is 4.1755 (EWIIM): position angle(paj=-11., The synthesized beam is $\times $ 5 (FWHM); position angle. +2°.. Fifteen 1.5625 MIITIZ channels were used. for a velocity resolution 10.3 ni 1. coutered at epsg=30 kins |.," Fifteen 1.5625 MHz channels were used, for a velocity resolution of 10.3 km $^{-1}$, centered at $v_{LSR}=30$ km $^{-1}$." + The phase ofceuter is a(.J2000) = 034 [69 00: 8(72000) = L7700., The phase center is $\alpha$ (J2000) = $\rm ^{h}$ $\rm ^{m}$ 3; $\delta$ (J2000) = 0. +. Amplitude tracking aud pointing was done bv observing the quasar 0228|673 every {5 unites., Amplitude tracking and pointing was done by observing the quasar 0228+673 every 45 minutes. + Phases were tracked by fast switching between the source and the quasar 03011655 with 1308/50s cycles., Phases were tracked by fast switching between the source and the quasar 0304+655 with 130s/50s cycles. + Absolute flux calibration was done using 3C18 and 3CLLT aud is good to —5LO%.., Absolute flux calibration was done using 3C48 and 3C147 and is good to $\sim$. + Calibration aud analysis was performed with the NRAO Astronomical hnuage Processing Software (AIPS) package., Calibration and analysis was performed with the NRAO Astronomical Image Processing Software (AIPS) package. + The naturally weighted aud CLEANed datacube has au rims of 0.70 uy beam+., The naturally weighted and CLEANed datacube has an rms of 0.70 mJy $^{-1}$. + Correction for the primary beam (1 aat 15 CITIz) has not been applied., Correction for the primary beam $\sim$ at 45 GHz) has not been applied. +" The shortest baselines iu the dataset are 25 kA. correspouding to scales of ~40""7:: structures larger than this are not woellsuupled."," The shortest baselines in the dataset are $\simeq$ 5 $\lambda$, corresponding to scales of $\sim$; structures larger than this are not well-sampled." + The IIC4N(16-15) line emission in the center of IC 312 was observed with 5 autenuas using the new 2 nuu receivers of the IRAM. Plateau de Bure interferometer (PdBI) ou 2007 December 28 and 31 in C configuration with basclines rauging from 21 to L76uun. The phase center of the observations was set to a(J2000) = 035 169 105: à(J2000) = ssl," The $_{3}$ N(16-15) line emission in the center of IC 342 was observed with 5 antennas using the new 2 mm receivers of the IRAM Plateau de Bure interferometer (PdBI) on 2007 December 28 and 31 in C configuration with baselines ranging from 24 to m. The phase center of the observations was set to $\alpha$ (J2000) = $\rm +^{h}$ $\rm ^{m}$ 105; $\delta$ (J2000) = 84." + NRAOLISO and 0212|735 served as phase calibrators and were observed every 20 iin., 150 and 0212+735 served as phase calibrators and were observed every 20 min. + Flux calibrators were 3C€81 and 301513., Flux calibrators were 3C84 and 3C454.3. + The calibration was done in the CILDAS package following staudard xocedures., The calibration was done in the GILDAS package following standard procedures. +" The TIC3N(16-15) line at 115.560951CCz was observed assumndue a systemic velocity of erpsg= laos Pound a spectral resolution of MAMI (5.15 aus ο),", The $_{3}$ N(16-15) line at GHz was observed assuming a systemic velocity of $v_{LSR}=46$ km $^{-1}$ and a spectral resolution of MHz (5.15 km $^{-1}$ ). + The average 2 mim coutiunmii was obtainedz o» averaemeg line-fee channels blue and redward of he TIC3N and ILCO lines aud subtracted from the de datacube to obtain a σολαμαπου datacube., The average 2 mm continuum was obtained by averaging line-free channels blue and redward of the $_{3}$ N and $_{2}$ CO lines and subtracted from the $uv$ datacube to obtain a continuum-free datacube. +" The final naturally woeiehted and CLEANed datacube with + wide chanucls has a CLEAN beam of «1.557. 16"" aud ΕΤ/beum."," The final naturally weighted and CLEANed datacube with $^{-1}$ wide channels has a CLEAN beam of $\times$ 1.55”, $^o$ and an rms of 1.7 mJy/beam." + Sinele-dish observations of IIC4N(1615) find a peals brightuess temperature of 6.7 τα] in a 16.9 beam (Aladroοἳal.2011)., Single-dish observations of $_{3}$ N(16–15) find a peak brightness temperature of 6.7 mK in a $^{''}$ beam \citep[][]{AMMMB11}. + Couvolving our imap to this bean and sampling it at the same location viekds a peak brightness of 7.1 aids. which aerees within the nucertaitics of both datasets.," Convolving our map to this beam and sampling it at the same location yields a peak brightness of 7.1 mK, which agrees within the uncertainties of both datasets." + Therefore no fux is resolved out of the interferometer maps. as expected for these high deusity tracers.," Therefore no flux is resolved out of the interferometer maps, as expected for these high density tracers." + Continuum subtracted integrated iutensitv line maps ο Ν(5E). HCSN(IO. 9) and TC3N(16 15) in IC 312 are shown in Figure 1..," Continuum subtracted integrated intensity line maps of $_{3}$ N(5–4), $_{3}$ N(10–9) and $_{3}$ N(16–15) in IC 342 are shown in Figure \ref{inti}. ." + Figure 2 shows the 5.£ transition overlaid on the 7 nua coutimmiun image generated frou offline channels., Figure \ref{inticont} shows the 5–4 transition overlaid on the 7 mm continuum image generated from offline channels. +" Locations of giant molecular cloud (GAIC) cores (Downesetal.1992:Meier&Turner2001) and the optical clusters (οι,Schinnereretal.2003).. compared with TC3N(5 1) and IICN(I0) (Downes ave also shown in Figure 2.."," Locations of giant molecular cloud (GMC) cores \citep[][]{DRGGGM92,MT01} and the optical clusters \citep[e.g.][]{SBM03}, compared with $_{3}$ N(5–4) and HCN(1–0) \citep[][]{DRGGGM92}, are also shown in Figure \ref{inticont}." + Figure 3. shows the IIC4N(5. £) and WC3N (16-15) spectra (iu flux units) taken over the same aaperture centered on cach cloud., Figure \ref{spec} shows the $_3$ N(5–4) and $_{3}$ N(16-15) spectra (in flux units) taken over the same aperture centered on each cloud. + (ο.1) enmuüssion picks out most clouds secu iu other dense eas tracers (e$ ΠοΧα0)., $_{3}$ N(5–4) emission picks out most clouds seen in other dense gas tracers (e.g HCN(1–0). + The ouly labeled cloud not clearly detected in WCsN (5£D) is GAIC D. the cloud associated with the nuclear star-forming region (Lin simlo®Lo:1983).," The only labeled cloud not clearly detected in $_3$ N(5–4) is GMC B, the cloud associated with the nuclear star-forming region \citep[$\rm L_{IR} 10^8;." + Positions of the GAIC's measured in IIC3N are consistcut with those fitted in C4%O(2-1) to within a beam (AIcier&Turner 2001)., Positions of the GMCs measured in $_{3}$ N are consistent with those fitted in $^{18}$ O(2-1) to within a beam \citep{MT01}. + An additional GAIC is detected in Νο. 1) just south of CAIC C3. labeled Cl.," An additional GMC is detected in $_{3}$ N(5–4) just south of GMC C3, labeled C4." +" Unlike ΟΦ1) which peaks at GMC ο Που. 1) oeission peaks farther north. towards CAIC C1. at a distance of 105 pe from the nucleus sugeesting chauges iu excitation across GAIC C. GMCS. A and D are resolved into two conmpoueuts, but we do not discuss them as separate entities eiven the lower signal-to-noise. other than to state that there is a clear difference in velocity ceutroid between the 5 land 1615 transitions (c.g. Figure 3)) indicating there is likely an excitation gradient across GAIC A. No cinission >26e is detected towards the weal CO(2-1) feature associated with the central nuclear star cluster (Schiuncreretal.2003)."," Unlike $^{18}$ O(2-1) which peaks at GMC C2, $_{3}$ N(5–4) emission peaks farther north, towards GMC C1, at a distance of 105 pc from the nucleus suggesting changes in excitation across GMC C. GMCs A and D' are resolved into two components, but we do not discuss them as separate entities given the lower signal-to-noise, other than to state that there is a clear difference in velocity centroid between the 5–4 and 16–15 transitions (e.g. Figure \ref{spec}) ) indicating there is likely an excitation gradient across GMC A. No emission $>2\sigma$ is detected towards the weak CO(2-1) feature associated with the central nuclear star cluster \citep[][]{SBM03}." +. The J-21615 line intensity is comparable to that of J=hIL line in all clouds., The J=16–15 line intensity is comparable to that of J=5–4 line in all clouds. + While dominated by CAIC C. IIC4N(16-15) is detected (or teutatively detected) towards all clouds. except GAIC D. Fluxes for GMCS D aud D' are quite uncertain since they are just inside the half power poiut of the PdBI primary beam.," While dominated by GMC C, $_{3}$ N(16-15) is detected (or tentatively detected) towards all clouds, except GMC D. Fluxes for GMCs D and D' are quite uncertain since they are just inside the half power point of the PdBI primary beam." + That GAIC C as much. brielter than the other clouds iu both transitions mdicates large quantities of deuse gas are present iu this cloud., That GMC C is much brighter than the other clouds in both transitions indicates large quantities of dense gas are present in this cloud. + IIC4N(1615) favors C2 over Cl., $_{3}$ N(16–15) favors C2 over C1. +" GMC D. while still faint in an absolute ποσο, is Significantly brighter iu HIC4N(1615) relative to 5l."," GMC B, while still faint in an absolute sense, is significantly brighter in $_{3}$ N(16–15) relative to 5--4." + This is not unexpected eiven the higher excitation requirements of the 1615 line aud that both D aud C2 have 7 nuu contiuuun sources associated with massive star formation., This is not unexpected given the higher excitation requirements of the 16–15 line and that both B and C2 have 7 mm continuum sources associated with massive star formation. + Comparisons of the Που (5£j. (109) and (1615) laps provide a chance to establish the excitation of the densest iiolecular cloud gas component.," Comparisons of the $_3$ N (5–4), (10–9) and (16–15) maps provide a chance to establish the excitation of the densest molecular cloud gas component." + Line intensities from the J=h Laud J=16-15 lines were moasured overaapertures centered on cach of the GAICs., Line intensities from the J=5–4 and J=16-15 lines were measured overapertures centered on each of the GMCs. + Table 1 records gaussian fits to each spectrum along with, Table \ref{Tinti} records gaussian fits to each spectrum along with +The stellar parameterslogg and logΊοῃ determined in the previous section allow us to estimate the luminosity of the investigated objects by interpolating the tables by Schmidt-Kaler (1982).,The stellar parameters$\log g$ and $\log T_{\rm eff}$ determined in the previous section allow us to estimate the luminosity of the investigated objects by interpolating the tables by \citet{schmidt}. +. The result is reported in Table 3.., The result is reported in Table \ref{tab3}. + Note that the errors on the luminosity were evaluated through the same tables by taking into account the errors on logg and logTeg., Note that the errors on the luminosity were evaluated through the same tables by taking into account the errors on $\log g$ and $\log T_{\rm eff}$. +" refhrdiagr shows the HR diagram for the nineteen stars studied in this work, in comparison with the zero-age main sequence (ZAMS) (Pickles1998),, the observed instability strip for 6SSct stars (Breger&Pamyatnykh1998) as well as the empirical and theoretical red edge of the yDDor instability strip (Handler&Shobbrook2002;Warneretal. 2003).."," \\ref{hrdiagr} shows the HR diagram for the nineteen stars studied in this work, in comparison with the zero-age main sequence (ZAMS) \citep{pickles}, the observed instability strip for $\delta$ Sct stars \citep{breger98} as well as the empirical and theoretical red edge of the $\gamma$ Dor instability strip \citep{handler, guzik}." +" We note that the pulsating variables (see Section 6)) are in the expected position, i.e. inside the instability strip, except for 008583770 1189177) and 005296877 1199-27597) which are hotter and cooler than the instability strip, respectively."," We note that the pulsating variables (see Section \ref{kepler}) ) are in the expected position, i.e. inside the instability strip, except for 08583770 189177) and 05296877 199-27597) which are hotter and cooler than the instability strip, respectively." +" The former was discussed in the previous section, the latter is not a pulsating star (see Section 6))."," The former was discussed in the previous section, the latter is not a pulsating star (see Section \ref{kepler}) )." + We can use the values of logg and logTeg to estimate the mass of the target stars., We can use the values of $\log g$ and $\log T_{\rm eff}$ to estimate the mass of the target stars. + For this purpose we used the evolutionary tracks in theBaSTI which are based on the evolutionary code (Chieffi&Straniero 1989)., For this purpose we used the evolutionary tracks in the which are based on the evolutionary code \citep{chieffi}. . +". We retrieved tracks with both canonical and non-canonical (i.e. with convective overshooting: Aov= 0.2H;) physics in the mass range 1—6 MMo with [M/H]=0.058, Z=0.0198, Y=0.273, and mixing 11.913 (Pietrinfernietal. 2004).."," We retrieved tracks with both canonical and non-canonical (i.e. with convective overshooting: $\lambda_{OV}=0.2H_p$ ) physics in the mass range $1 - 6$ $_{\odot}$ with $[M/H]= 0.058$, $Z= 0.0198$, $Y= 0.273$, and mixing 1.913 \citep{pietri04}. ." + In Fig., In Fig. + 4 we show the logg— diagram., \ref{fig3} we show the $\log g - \log T_{\rm eff}$ diagram. +" The left and rightpanels in the figure show canonical and non-canonical tracks, respectively."," The left and rightpanels in the figure show canonical and non-canonical tracks, respectively." + The resulting masses for the two cases (listed in, The resulting masses for the two cases (listed in + , +"The arrival of the era of synoptic imaging surveys heralds the start of a new chapter of astrophysics, where the real-time processing of images taxes the capacity to transport the data from remote sites and pushes to the limit the computational capabilities at processing centers (e.g., Jurié&Ivezié 2011)).","The arrival of the era of synoptic imaging surveys heralds the start of a new chapter of time-domain astrophysics, where the real-time processing of images taxes the capacity to transport the data from remote sites and pushes to the limit the computational capabilities at processing centers (e.g., \citealt{2011EAS....45..281J}) )." +" More profoundly novel, however, is that the data volumes have begun to surpass what is possible to visually inspect by even large teams of astronomers and volunteer “citizen scientists."""," More profoundly novel, however, is that the data volumes have begun to surpass what is possible to visually inspect by even large teams of astronomers and volunteer “citizen scientists.”" + This necessitates an increasingly more central role of software and hardware frameworks to supplant the traditional roles of humans in the real-time loop., This necessitates an increasingly more central role of software and hardware frameworks to supplant the traditional roles of humans in the real-time loop. +" This abstraction of people away from the logistics of the scientific process has been progressing rapidly, starting with the acquisition process itself."," This abstraction of people away from the logistics of the scientific process has been progressing rapidly, starting with the acquisition process itself." +" Indeed, robotictelescopes!,, capable of taking data autonomously at remote sites, have become an increasingly common form of operation at the sub-meter- and meter-class level (cf. Castro-Tirado 2010))."," Indeed, robotic, capable of taking data autonomously at remote sites, have become an increasingly common form of operation at the sub-meter- and meter-class level (cf. \citealt{ca2010}) )." +" Many robotic systems use queuing algorithms that optimize nightly observing over several scientific programs and many are capable of being interrupted by external alerts to observe high-priority transients (e.g., Filippenkoetal.al.2008;Kubanek 2010))."," Many robotic systems use queuing algorithms that optimize nightly observing over several scientific programs and many are capable of being interrupted by external alerts to observe high-priority transients (e.g., \citealt{2001ASPC..246..121F,2002SPIE.4845..126V,2003PASP..115..132A,2006PASP..118.1396C,2006ASPC..351..751B,2008AN....329..321S,2010arXiv1002.0108K}) )." +" Data from such facilities can be automatically transported, processed, photometered, and ingested into databases without human intervention."," Data from such facilities can be automatically transported, processed, photometered, and ingested into databases without human intervention." +" Since imaging data has spurious sources of noise and artifacts that can mimic real astrophysical sources, in the absence of watchful trained eyes on the images themselves, autonomous discovery of transients and variable stars on synoptic imaging surveys is a significant challenge."," Since imaging data has spurious sources of noise and artifacts that can mimic real astrophysical sources, in the absence of watchful trained eyes on the images themselves, autonomous discovery of transients and variable stars on synoptic imaging surveys is a significant challenge." +" Threshold cuts on photometric quality, changes in apparent magnitudes, etc.,"," Threshold cuts on photometric quality, changes in apparent magnitudes, etc.," + are effective in discovering bona fide astrophysics sources (Drakeetal.2009;Sokolowski2010)..," are effective in discovering bona fide astrophysics sources \citep{2009ApJ...696..870D,2010AdAst2010E..54S}." +" However, multi-parameter thresholding tends to be suboptimal because it treats each parameter derived from a given candidate as an independent variable when clearly there can be correlations between parameters."," However, multi-parameter thresholding tends to be suboptimal because it treats each parameter derived from a given candidate as an independent variable when clearly there can be correlations between parameters." +" Matched filtering—looking for light curve trends that fit the scientific expectation from a certain class of variables (e.g., microlensing; Tomaney&Crotts1996;Belokurovetal. 2003))—can be a very effective tool to discover new events, but other sorts of variables and transients are not easily recovered from that view of the dataset."," Matched filtering—looking for light curve trends that fit the scientific expectation from a certain class of variables (e.g., microlensing; \citealt{1996AJ....112.2872T,2003MNRAS.341.1373B}) )—can be a very effective tool to discover new events, but other sorts of variables and transients are not easily recovered from that view of the dataset." +" Likewise, previous machine-learning based discovery (e.g., supernova discovery with the Supernova Factory; Baileyetal. 2007)) have been optimized on domain-specific discovery, leaving aside the multitude of other variables not of direct interest to that particular project."," Likewise, previous machine-learning based discovery (e.g., supernova discovery with the Supernova Factory; \citealt{2007ApJ...665.1246B}) ) have been optimized on domain-specific discovery, leaving aside the multitude of other variables not of direct interest to that particular project." + Discovery that a varying source is truly astrophysical does not mean that the origin of that, Discovery that a varying source is truly astrophysical does not mean that the origin of that +features in high-redshift star-forming galaxies (c.g. Lacy&Scott 2005)).,features in high-redshift star-forming galaxies (e.g. \citealt{Sajina05}) ). + The results of this fitting process are plotted in Fig. 4.., The results of this fitting process are plotted in Fig. \ref{fig:fits2}. + Notice that the 350yim data point provides almost all of the contraining power on 74.," Notice that the $350\,\mathrm{\mu m}$ data point provides almost all of the contraining power on $T_\mathrm{A}$." + We have measured the apparent temperatures (24) of an unbiased sample of SALCGs from SLLADIES., We have measured the apparent temperatures $T_\mathrm{A}$ ) of an unbiased sample of SMGs from SHADES. + We calculate the corresponding temperature distribution by adding a rectangle of area one centred on the source 7. with a width eiven by the errors in Jy to the histogram (see Fig. 5)).," We calculate the corresponding temperature distribution by adding a rectangle of area one centred on the source $T_\mathrm{A}$, with a width given by the errors in $T_\mathrm{A}$ to the histogram (see Fig. \ref{fig:tdist}) )." + Ht appears that our survey selects SMGs with Z42(S83) Why. We recover the intrinsic dust temperature (74) distribution of SAIGs by shifting cach SED into the restframoe. following a similar procedure to that described for the 7 distribution above (see Fig. 5)).," It appears that our survey selects SMGs with $T_\mathrm{A}\simeq(8\pm3)$ K. We recover the intrinsic dust temperature $T_\mathrm{d}$ ) distribution of SMGs by shifting each SED into the restframe, following a similar procedure to that described for the $T_\mathrm{A}$ distribution above (see Fig. \ref{fig:tdist}) )." + The T4 distribution appears to contain structure Cbumps! centred. around. 121xIx.. INIx.. and IxIx) and is broader compared with the Z4 distribution.," The $T_\mathrm{d}$ distribution appears to contain structure (`bumps' centred around K, K, and K) and is broader compared with the $T_\mathrm{A}$ distribution." +" ‘To investigate if this distribution is more structured and/or wider than we would expect by chance. we have simulated the temperature distribution we would expect to see if all SAIGs have Z4=28 XI. The simulated Z4 cistribution is constructed. by adding an error. box of area one. centred on d4=δι|rand6. where ""rand is a normallv-clistributed random. number. for each source."," To investigate if this distribution is more structured and/or wider than we would expect by chance, we have simulated the temperature distribution we would expect to see if all SMGs have $T_\mathrm{d}=28$ K. The simulated $T_\mathrm{d}$ distribution is constructed by adding an error box of area one centred on $T_\mathrm{d}=28\,\mathrm{K}+\mathrm{rand}\times\sigma$, where `rand' is a normally-distributed random number, for each source." +" Although the shape and width of the simulated. distribution. changes notably depending on the random. realisation. we reproduce the ""bumps. at Z4z12 or 35IxIx and we wediet more SMGs at a few Ix than we see."," Although the shape and width of the simulated distribution changes notably depending on the random realisation, we reproduce the `bumps' at $T_\mathrm{d}\simeq12$ or K and we predict more SMGs at a few K than we see." + Therefore we conclude that the underlying distribution is just a bit wider han we expect by chance and is consistent with most SMCs al Ab with a few hotter ones., Therefore we conclude that the underlying distribution is just a bit wider than we expect by chance and is consistent with most SMGs at K with a few hotter ones. + Now we compare the dust emperatures of our photometric and spectroscopic redshift sub-samples and to previous estimates for SMCs., Now we compare the dust temperatures of our photometric and spectroscopic redshift sub-samples and to previous estimates for SMGs. + We derive a mean Z4=(31x18)lx (median Z4=31 Ix) or our sample of SMCis with spectroscopic redshifts.," We derive a mean $T_\mathrm{d}=(31\,\pm{18})\,\mathrm{K}$ (median $T_\mathrm{d}=31\,\mathrm{K}$ ) for our sample of SMGs with spectroscopic redshifts." + For the sample of 21 350 jan-observed SMCGs with only photometric redshifts. the mean is Z4=(35cx17)W (median 14= 32011.," For the sample of 21 $350\,\mathrm{\mu m}$ -observed SMGs with only photometric redshifts, the mean is $T_\mathrm{d}=(35\,\pm{17})\,\mathrm{K}$ (median $T_\mathrm{d}=29\,\mathrm{K}$ )." + A two-sided. IWS test reveals that. the temperature distributions of the photometric redshift and spectroscopic redshift subsets are. consistent with being drawn from the same distribution., A two-sided KS test reveals that the temperature distributions of the photometric redshift and spectroscopic redshift subsets are consistent with being drawn from the same distribution. + Fherefore we quote a mean Z4=(34+17)I (median Z4— POW) for our entire sample of SMCis.," Therefore we quote a mean $T_\mathrm{d}=(34\,\pm{17})\,\mathrm{K}$ (median $T_\mathrm{d}=29\,\mathrm{K}$ ) for our entire sample of SMGs." + Our mean Zi; is consistent with previous estimates for SMCis from Chapmanetal.(2005). (Z4=361x for radio-identified SALCGs spectroscopic redshifts). from. Ixovaesctal.(2006) (2)= 3541). and from Popeetal.(2006). (Z4=30h for SMCs in the Creat Observatories Origins Deep," Our mean $T_\mathrm{d}$ is consistent with previous estimates for SMGs from \citet{Chapman2005} $T_\mathrm{d}=36\,\mathrm{K}$ for radio-identified SMGs spectroscopic redshifts), from \citet{Kovacs} $T_\mathrm{d}=35\,\mathrm{K}$ ), and from \citet{Pope2006} $T_\mathrm{d}\simeq30\,\mathrm{K}$ for SMGs in the Great Observatories Origins Deep" +where d ds the starobserver distance.,where $d$ is the star–observer distance. + Here we utilize he analysis of the previous section to determune the orbital location of λα flux ratio for a subset of he known exoplauets., Here we utilize the analysis of the previous section to determine the orbital location of maximum flux ratio for a subset of the known exoplanets. + InTable 1.. we report the true anomaly f (anele between the direction of periapsis aud he current position of the planet im the orbit). fiux ratio e. and projected separation AO at this location for he ost eccentric exoplaucts.," InTable \ref{phasetab}, we report the true anomaly $f$ (angle between the direction of periapsis and the current position of the planet in the orbit), flux ratio $\epsilon$ , and projected separation $\Delta +\theta$ at this location for the most eccentric exoplanets." + We additionally report if the maximum flux ratio occurs for an cdge-on or acc-on orientation., We additionally report if the maximum flux ratio occurs for an edge-on or face-on orientation. + Notice that for face-ou orbits the rue anomaly is zero at this location since the flux is completely driven by the event of periastron passage., Notice that for face-on orbits the true anomaly is zero at this location since the flux is completely driven by the event of periastron passage. + There is a bias against optimal separation since uaxinumn fux naturally occurs when the plauct is closest o the star., There is a bias against optimal separation since maximum flux naturally occurs when the planet is closest to the star. +" This is most severe for edge-«n orbits. such as ΠΟ sd606b where the angular separatjon is ~0"" at his location."," This is most severe for edge-on orbits, such as HD 80606b where the angular separation is $\sim 0\arcsec$ at this location." + For this reason. the locatio1 of mandi Hux ratio prefercutially corresponds to the location of nüuiuun angular separation.," For this reason, the location of maximum flux ratio preferentially corresponds to the location of minimum angular separation." + Cousicder the case of IID 39091b (Figure 3)) where the peak fux arises frou an edee-on oricutation., Consider the case of HD 39091b (Figure \ref{fluxsep}) ) where the peak flux arises from an edge-on orientation. +" The peak flux ratio correspouds to a 0.02” separation. whereas the plauct reaches a separation of ~0.26"" during the cutive phase."," The peak flux ratio corresponds to a $0.02\arcsec$ separation, whereas the planet reaches a separation of $\sim +0.26\arcsec$ during the entire phase." + By iatchine the aneular separations above the resolution criteria of the coronagraph. one can choose the optimal targets for monitoring divine predicted the orbital phase of peak flux ratio.," By matching the angular separations above the resolution criteria of the coronagraph, one can choose the optimal targets for monitoring during predicted the orbital phase of peak flux ratio." + The main targets for which phase signature detection will be attempted will likely be those for which a planet is already kuown to be preseut., The main targets for which phase signature detection will be attempted will likely be those for which a planet is already known to be present. + Iu this case. the purpose of the observations are chatacterization rather than discovery.," In this case, the purpose of the observations are chatacterization rather than discovery." +" As shown bv Equation 1.. the fux ratio is a function of the geometric albedo sly. the periastron aremuenut ο. the orbital inclination /. aud the plauctary radius &,,."," As shown by Equation \ref{fluxratio}, the flux ratio is a function of the geometric albedo $A_g$, the periastron argument $\omega$, the orbital inclination $i$, and the planetary radius $R_p$." + The challenge of discutangling those conrponeuts frou the phase shape and amplitude will vary depending upon what is already known for that svsteni., The challenge of disentangling those components from the phase shape and amplitude will vary depending upon what is already known for that system. + If the planet is known to trausit then oue can measure iow. aud δὲν.," If the planet is known to transit then one can measure $i$, $\omega$, and $R_p$." +" One cau thus determine A, which. along with models of the planctary structure based upou the size and mass. cau coustrain the properties of the atinosplere."," One can thus determine $A_g$ which, along with models of the planetary structure based upon the size and mass, can constrain the properties of the atmosphere." + Tf however the planet does not transit. then we can estimate the radius of the planet from the measured mass using theoretical models. such as those of Boclenheiueretal.(2003). aud Fortneyctal. (2007)... for which there will be associated uncertainties depending wpou factors such as the age of the plauct and the asstmed core model.," If however the planet does not transit, then we can estimate the radius of the planet from the measured mass using theoretical models, such as those of \citet{bod03} + and \citet{for07}, for which there will be associated uncertainties depending upon factors such as the age of the planet and the assumed core model." + In some cases. astrometric nieasurenients of the host star can performed. such as those carried out by Wanetal.(2001).," In some cases, astrometric measurements of the host star can performed, such as those carried out by \citet{han01}." +. This will have the simultancous result of resolving the inclination of the orbit. thus breaking the degeneracy of the flux ratio with the periastron aremuent. and determining the true mass of the secondary companion.," This will have the simultaneous result of resolving the inclination of the orbit, thus breaking the degeneracy of the flux ratio with the periastron argument, and determining the true mass of the secondary companion." + For nou-trausiting planets discovered using the racial velocity method. one has the advantage of a strong bias towards bright host stars relative to those discovered using the transit method.," For non-transiting planets discovered using the radial velocity method, one has the advantage of a strong bias towards bright host stars relative to those discovered using the transit method." + The increase in signal-to-noise will be a great assetfor au appropriate instrament to exploit. as described in Section ον," The increase in signal-to-noise will be a great assetfor an appropriate instrument to exploit, as described in Section \ref{feasibility}. ." + Possible contamination of the phase signature at, Possible contamination of the phase signature at +Based on (he previous section. one concludes that there is no significant dillerence in CN strength between M 31 and MW GCs with [Fe/H]S—0.4.,"Based on the previous section, one concludes that there is no significant difference in CN strength between M 31 and MW GCs with $[Fe/H] \simless -0.4$." + The same cannot be said about more metal-rich GCs. since the three data points at the highanetallicitv eid of the AIW GC sample have substantially lower CN indices than their M 31 counterparts at the same [|Fe/1l].," The same cannot be said about more metal-rich GCs, since the three data points at the high-metallicity end of the MW GC sample have substantially lower CN indices than their M 31 counterparts at the same [Fe/H]." +" Moreover. these three high-|Fe/II] MW points seem to depart significantly from the (Fe) CN, and (Fe) CVs trends established by the MW. GCs alone. by being displaced. towards low CN values."," Moreover, these three high-[Fe/H] MW points seem to depart significantly from the $\langle$ $\rangle$ $CN_1$ and $\langle$ $\rangle$ $CN_2$ trends established by the MW GCs alone, by being displaced towards low CN values." + The two highest (Fe) points are from two spectra of NGC 6528. obtained with different spatial extraction windows during data reduction (seeschiavonetal.2005.for detzils).. whereas the third point comes from the spectrum of NGC 6553 (a third metal-rich GC. NGC 5927. also seems to be marginally too weak in CN).," The two highest $\langle$ $\rangle$ points are from two spectra of NGC 6528, obtained with different spatial extraction windows during data reduction \citep[see][for details]{s05}, whereas the third point comes from the spectrum of NGC 6553 (a third metal-rich GC, NGC 5927, also seems to be marginally too weak in CN)." + Visual confirmation of this result is offered at the bottom panel of Figure 9 where the spectrum of NGC 6528 is compared to the average spectrum of its M 31 counterparts., Visual confirmation of this result is offered at the bottom panel of Figure \ref{spectra} where the spectrum of NGC 6528 is compared to the average spectrum of its M 31 counterparts. + CN bands seem {ο be stronger in the average M 31 GC spectrum than in that of NGC 6528., CN bands seem to be stronger in the average M 31 GC spectrum than in that of NGC 6528. + Moreover. all features bluer than 4000. A. and particularly Che Call IL and Ix lines. are weaker in NGC 6528.," Moreover, all features bluer than 4000 ${\rm\AA}$, and particularly the CaII H and K lines, are weaker in NGC 6528." + The GCs NGC 6553 and 6528 are also the (wo most metal-rich in the sample. and no such difference between these GCs and their M 31 counterparts is seen in their data (Figure 10)).," The GCs NGC 6553 and 6528 are also the two most metal-rich in the \cite{pu02} sample, and no such difference between these GCs and their M 31 counterparts is seen in their data (Figure \ref{puziadata}) )." +" According to Puziaetal.(2002).. the CN, index is stronger in the spectra of NGC 6528 and 6553 bv 0.035. and 0.096 mae than in (2005)."," According to \cite{pu02}, the $CN_1$ index is stronger in the spectra of NGC 6528 and 6553 by 0.035 and 0.096 mag than in \cite{s05}." +. One therefore is left wondering whether these differences may be due to different treatments of skv subtraction ancl spatial sampling adopted by Puziaetal.(2002) ancl schiavonetal.(2005)., One therefore is left wondering whether these differences may be due to different treatments of sky subtraction and spatial sampling adopted by \cite{pu02} and \cite{s05}. +. Let us recall that NGC 6528 and 6553 are located towards the Galactic bulge. in a region of the skv that is affected by strong background ancl foreground contamination. rendering (he task of skv subtraction rather (icky and prone to large," Let us recall that NGC 6528 and 6553 are located towards the Galactic bulge, in a region of the sky that is affected by strong background and foreground contamination, rendering the task of sky subtraction rather tricky and prone to large" +determine empirical oxygen abuucdauces aud associated statistical errors. aucl have listed them in Table 3.,"determine empirical oxygen abundances and associated statistical errors, and have listed them in Table 3." + Pilyugin (2000) has proposed a new calibration of the empirical oxygen abundance scale., Pilyugin (2000) has proposed a new calibration of the empirical oxygen abundance scale. + After rearranging terius. Pilyugiu's calibration is: We have added the oxyge1 abuudances calculaed via Pilvugin’s calibration to Table 3. ane compared. Pilyiwins calibratiou to that of MeCGaugh in Figure [.," After rearranging terms, Pilyugin's calibration is: We have added the oxygen abundances calculated via Pilyugin's calibration to Table 3, and compared Pilyugin's calibration to that of McGaugh in Figure 4." + There are two siguificant aix systematic cilfereuces between Ιese two elupirical caibrations for low metallicity HII regions., There are two significant and systematic differences between these two empirical calibrations for low metallicity HII regions. + First. uote that Pilyieins calibration extends to arger values of log (Ro;).," First, note that Pilyugin's calibration extends to larger values of log $_{23}$ )." + This solves the well know: problem that NIcGaugh’s grid «oes nol cover the highest values of log (Rox) that are observe., This solves the well known problem that McGaugh's grid does not cover the highest values of log $_{23}$ ) that are observed. + Because Pilvuegiws calibration is entirely euidirical. it covers the full rauge of the observatious.," Because Pilyugin's calibration is entirely empirical, it covers the full range of the observations." + The second significaut dilference is 1ie chiauge ir slopes of the calibrations at low values of log (032)., The second significant difference is the change in slopes of the calibrations at low values of log (O32). + Since most of he calibratiug observatious uxed by Pilvugin had high values of log (032). the extrapolation o ‘his calibration iito the lowe: log (032) regine is subject to very large uucertainty.," Since most of the calibrating observations used by Pilyugin had high values of log (O32), the extrapolation of his calibration into the lower log (O32) regime is subject to very large uncertainty." + Note. for example. the 0.31 ¢ex difference between M91 aud POO for he low excitation HII regious NGC 625 #221.," Note, for example, the 0.34 dex difference between M91 and P00 for the low excitation HII regions NGC 625 21." + Because MeGaugh’s calibration is guided by pliotoionization= models iu this regime. it is probably the more accuule there.," Because McGaugh's calibration is guided by photoionization models in this regime, it is probably the more accurate there." + Vieure 5 shows the dillerences between the three oxyeen deter:inatious given in Table 3 as a [uncion of Ixth log (032 and log (Re)., Figure 5 shows the differences between the three oxygen determinations given in Table 3 as a function of both log (O32) and log $_{23}$ ). + When the differences are plotted as a fuction of log (O32). the systematic difference between the MOI auc POO calibrations is obvious.," When the differences are plotted as a function of log (O32), the systematic difference between the M91 and P00 calibrations is obvious." + As could also be see iiu Figure [. in some cases. the empirical oxvgen abuncdauces are in good ag'eelueut. with the direct oxygen abuldaunces. wih a slight das Lowa‘ds lower oxygen abundances from the direct ineasurements ( Le.{je cliferences for ESO [71-G06 #22. ESO 173-CG2| 3:22. aud GC 625 15 all fall beween -0.03 aud -0.15 dex).," As could also be seen in Figure 4, in some cases, the empirical oxygen abundances are in good agreement with the direct oxygen abundances, with a slight bias towards lower oxygen abundances from the direct measurements (i.e., the differences for ESO 471-G06 2, ESO 473-G24 2, and NGC 625 18 all fall between -0.03 and -0.15 dex)." + The stuall jas coul be iudicative of any of auunber of cifTe‘elit effects., The small bias could be indicative of any of a number of different effects. + One possibility is the breakdownu o ‘the assumption of a uniform electron temperature (see Peimbert. Peinvert. Luricliana 2002 and references therein).," One possibility is the breakdown of the assumption of a uniform electron temperature (see Peimbert, Peimbert, Luridiana 2002 and references therein)." +" A second. related possibility is the presence of a secoud heating source iL addition to photoionization (Stasitisska. Schae‘er, Leitherer 2001"," A second related possibility is the presence of a second heating source in addition to photoionization (Stasińsska, Schaerer, Leitherer 2001)." + Athi1“cL possibility is that racliatic1 fields are inich softer that tlose considered by the models (M(Caugl 1991 shows that this is a 'elatively small effect - compare the two models plotted iu Figye κ)+) nut Stasitisska et 22001 caution that use of only zero-age stellar 1uodels does not probe a large enough range in input stelar 1uocels)., A third possibility is that radiation fields are much softer that those considered by the models (McGaugh 1991 shows that this is a relatively small effect - compare the two models plotted in Figure 3 - but Stasińsska et 2001 caution that use of only zero-age stellar models does not probe a large enough range in input stellar models). + However. here a'e sone truly discrepant results which do require attention.," However, there are some truly discrepant results which do require attention." + ESO 317-C17 ii55 shows a«illerence of —0.35 dex for the M9!| calibration and —0.32 for the POO calibration., ESO 347-G17 5 shows a difference of $-$ 0.35 dex for the M91 calibration and $-$ 0.32 for the P00 calibration. + On the other 116]. 1le empirical abuudauces for HIE region ESO 317-CL7 4110 diller by 0.31 dex.," On the other hand, the empirical abundances for HII region ESO 347-G17 10 differ by 0.31 dex." + The brighest two HII regious lor NGC 625 also show siguificant diserepaucies., The brightest two HII regions for NGC 625 also show significant discrepancies. + The empirical abundances for NGC 625 3555 and #99 are between —0.19 to —0.35 dex below those of the direct oxygen abuncdauces., The empirical abundances for NGC 625 5 and 9 are between $-$ 0.19 to $-$ 0.35 dex below those of the direct oxygen abundances. + It is interesting that the empirical oxygen abuudances for the three brightest, It is interesting that the empirical oxygen abundances for the three brightest +which is a result that very unambiguously connects the depth of a reflecting stratum with the time of the reception of a signal reflected from that stratum (under the simple condition that the velocity of the wave propagation 15 unchanging).,which is a result that very unambiguously connects the depth of a reflecting stratum with the time of the reception of a signal reflected from that stratum (under the simple condition that the velocity of the wave propagation is unchanging). + An important point to note here is that the simple appearance of Eq. (28)), An important point to note here is that the simple appearance of Eq. \ref{b6}) ) + has been possible because the plane wave allows only the vertically propagating mode in Eq. (10)).," has been possible because the plane wave allows only the vertically propagating mode in Eq. \ref{kayzed}) )," + and so by default this mode has been decoupled from all the other modes., and so by default this mode has been decoupled from all the other modes. + This decoupling should be impossible to achieve for higher spatial dimensions. as a look at Eq. (10))," This decoupling should be impossible to achieve for higher spatial dimensions, as a look at Eq. \ref{kayzed}) )" + ought to reveal., ought to reveal. + It shall be useful at this point to derive the traveltime equation for the 3D zero-offset case. using the method of the stationary phase approximation.," It shall be useful at this point to derive the traveltime equation for the $3D$ zero-offset case, using the method of the stationary phase approximation." + From Eq. (19)), From Eq. \ref{invfou}) ) + one can. making use of the extrapolation condition given by Eq. (11)).," one can, making use of the extrapolation condition given by Eq. \ref{solinteg}) )," + write à compact relation that will read as in which It is easy to see that the integrand in Eq. (29)), write a compact relation that will read as in which It is easy to see that the integrand in Eq. \ref{statio}) ) + has an amplitude part and a phase part., has an amplitude part and a phase part. + If the phase is to vary much more rapidly than the amplitude (a requirement that is consistent with the analysis presented in Section 4)). then the consequent rapid oscillations over most of the range of the integration. will result in an average value of nearly zero.," If the phase is to vary much more rapidly than the amplitude (a requirement that is consistent with the analysis presented in Section \ref{sec35}) ), then the consequent rapid oscillations over most of the range of the integration, will result in an average value of nearly zero." + The exception to this principle will occur when 6 is stationary (Mathews&Walker1970:ArfkenWeber2001 ).," The exception to this principle will occur when $\Phi$ is stationary \citep{mw, aw}. ." +. This (the stationary phase approximation) will be quantified by 6®=0. which is. in fact. a condition for the turning points of b.," This (the stationary phase approximation) will be quantified by $\delta \Phi = 0$, which is, in fact, a condition for the turning points of $\Phi$." + The major contribution to the integral will. therefore. come from the turning points of the phase function ®.," The major contribution to the integral will, therefore, come from the turning points of the phase function $\Phi$." + This condition will help in identifying the physical coordinates from where one will obtain the most significant contribution to the signal received on the surface of the earth., This condition will help in identifying the physical coordinates from where one will obtain the most significant contribution to the signal received on the surface of the earth. + Going by Eq. (10)).," Going by Eq. \ref{kayzed}) )," +" itis already known that &. has a dependence on /,. 4, and ο."," it is already known that $k_z$ has a dependence on $k_x$, $k_y$ and $\omega$." + Taking this inconjunction with Eq. (30)).," Taking this inconjunction with Eq. \ref{phi}) )," + and noting that the relevant variables for the integration in Eq. (29)), and noting that the relevant variables for the integration in Eq. \ref{statio}) ) +" are &,. ἂν and το. the required extremum condition for ® can consequently be stated as The result above can only mean that all the three partial derivatives in Eq. (31))"," are $k_x$, $k_y$ and $\omega$, the required extremum condition for $\Phi$ can consequently be stated as The result above can only mean that all the three partial derivatives in Eq. \ref{phivar}) )" + are to be individually set to zero., are to be individually set to zero. + This will give three distinct mathematical conditions. from which one will be able to derive the traveltime equation for the 3D zero-offset case as This defines a sphere of radius c£ in the ./—jg—-: coordinate space. and a hyperboloid in the 4—4-f space for a constant value of :.," This will give three distinct mathematical conditions, from which one will be able to derive the traveltime equation for the $3D$ zero-offset case as This defines a sphere of radius $vt$ in the $x$ $y$ $z$ coordinate space, and a hyperboloid in the $x$ $y$ $t$ space for a constant value of $z$." + Considering only the inline direction and the depth (le. when y= 0). Eq. (32))," Considering only the inline direction and the depth (i.e. when $y=0$ ), Eq. \ref{traveltime}) )" + will define a semi-circle below the surface of the earth in the ./—: plane. and the usual hyperbola in the ..—? plane.," will define a semi-circle below the surface of the earth in the $x$ $z$ plane, and the usual hyperbola in the $x$ $t$ plane." + However. a much more interesting result can be seen for the simplest possible case of the planar wavefront. 1.8 where the spatial dependence ts only on the : coordinate.," However, a much more interesting result can be seen for the simplest possible case of the planar wavefront, i.e where the spatial dependence is only on the $z$ coordinate." + In that case the result in Eq. (32)), In that case the result in Eq. \ref{traveltime}) ) + willconverge simply to the linear solution += ct. which is exactly what Eq. (28))," willconverge simply to the linear solution $z = vt$ , which is exactly what Eq. \ref{b6}) )" + also furnishes., also furnishes. + So effectively this simple special case establishes an intriguing mathematical equivalence between the stationary phase approximation method and the constant velocity Stolt migration algorithm., So effectively this simple special case establishes an intriguing mathematical equivalence between the stationary phase approximation method and the constant velocity Stolt migration algorithm. +model are given by η3.,model are given by $n-3$. + For the inner lobes we assume the jets are still supplying these structures with energy. ancl so he number of free parameters is reduced to three. because f=da ," For the inner lobes we assume the jets are still supplying these structures with energy and so the number of free parameters is reduced to three, because $t=t_{\rm j}$." +Also. we show below that the spectra of the inner obes are consistent with power-laws.," Also, we show below that the spectra of the inner lobes are consistent with power-laws." + Hence the individual lux measurements are not independent and they. can only ο counted as a single measurement., Hence the individual flux measurements are not independent and they can only be counted as a single measurement. + For the inner lobes we herefore have only two observational constraints. but three ree parameters. resulting in relations between pairs of free xwameters instead of specifie values for cach of them.," For the inner lobes we therefore have only two observational constraints, but three free parameters, resulting in relations between pairs of free parameters instead of specific values for each of them." + The spectra of the inner lobes. up to an observed frequency of at least S46 Cillz. are consistent with a power-law description.," The spectra of the inner lobes, up to an observed frequency of at least 8.46 GHz, are consistent with a power-law description." + This implies that the electrons. emitting in the observed frequency range have not. sullered. significant radiative energv losses since they were accelerated. to relativistic velocities., This implies that the electrons emitting in the observed frequency range have not suffered significant radiative energy losses since they were accelerated to relativistic velocities. + The inner lobes must therefore. be voung., The inner lobes must therefore be young. + This requirement is consistent with our assumption wt the inner lobes represent a second activity phase of the jet Dow from the central AGN after the energy supply to the puter lobes shuts down., This requirement is consistent with our assumption that the inner lobes represent a second activity phase of the jet flow from the central AGN after the energy supply to the outer lobes shuts down. + The inner lobes then cannot be older mana few. 10 vvears. because for an older age the emission rom the outer lobes should have faded below the Duxes at 1 higher observed. frequencies.," The inner lobes then cannot be older than a few $10^7$ years, because for an older age the emission from the outer lobes should have faded below the fluxes at the higher observed frequencies." + Given the observed. elose lienment of the inner lobes with the outer lobes. we assume at the inner lobes are fully contained within the volumes of the outer lobes.," Given the observed close alignment of the inner lobes with the outer lobes, we assume that the inner lobes are fully contained within the volumes of the outer lobes." + We propose two dilferent interpretations or the existence and observed. properties of the inner lobes., We propose two different interpretations for the existence and observed properties of the inner lobes. + The first interpretation is that the same model used. to explain the outer lobes also applies to the inner lobes., The first interpretation is that the same model used to explain the outer lobes also applies to the inner lobes. + The jet flows inflating the inner lobes must then end in strong shocks giving rise to particle acceleration., The jet flows inflating the inner lobes must then end in strong shocks giving rise to particle acceleration. + The observations do not show much evidence for luminous radio hotspots which could. be identified as the jet shocks., The observations do not show much evidence for luminous radio hotspots which could be identified as the jet shocks. + However. given our limited. spatial resolution. we cannot rule out the oesence of compact structures with an enhanced. racio uminositv at the end of the inner lobes.," However, given our limited spatial resolution, we cannot rule out the presence of compact structures with an enhanced radio luminosity at the end of the inner lobes." + The model is identical to that developed in Section 5. for he outer lobes., The model is identical to that developed in Section \ref{KDA} for the outer lobes. + However. we can make some simplifications.," However, we can make some simplifications." + We assume that the inner lobes are currently. still supplied with energv by active jets., We assume that the inner lobes are currently still supplied with energy by active jets. + Hence /;=/ and we reduce he number of free parameters by one., Hence $t_{\rm j} =t$ and we reduce the number of free parameters by one. + Also. the. power-aw shape of the observed. spectrum implies. insignificant radiative energy losses of the relativistic electrons.," Also, the power-law shape of the observed spectrum implies insignificant radiative energy losses of the relativistic electrons." + The inner obes will be well described by the much simplified model in he adiabatic regime as described in. 2.., The inner lobes will be well described by the much simplified model in the adiabatic regime as described in \citet{kb07}. +" For an observed obe length. of DD. ancl radio luminosity density L, dt is straightforward to derive relations between the source age of the inner lobes. ὃς and the jet. power. Q. where e. fi. ff, and e, are constants defined in ?.."," For an observed lobe length of $D$ and radio luminosity density $L_{\nu}$ it is straightforward to derive relations between the source age of the inner lobes, $t_{\rm i}$, and the jet power, $Q$, where $\epsilon$, $f_{\rm L}$, $f_{\rm p}$ and $c_1$ are constants defined in \citet{kb07}." + We can also find a relation between the scale. density. of the ambient gas. py. and the jet. power. The observed power-law spectra have steep slopes and in the absence of significant. radiative energy Losses these steep slopes translate into large values for the exponent of the initial energy distribution of the relativistic electrons. m.," We can also find a relation between the scale density of the ambient gas, $\rho_{\rm i}$, and the jet power, The observed power-law spectra have steep slopes and in the absence of significant radiative energy losses these steep slopes translate into large values for the exponent of the initial energy distribution of the relativistic electrons, $m$." + Large m may indicate inellicient particle acceleration. ancl they make the model results more dependent. on the Low energv cut-olf of the electron energy. distribution as most of the enerey of the electrons is stored in the low energy end of the distribution., Large $m$ may indicate inefficient particle acceleration and they make the model results more dependent on the low energy cut-off of the electron energy distribution as most of the energy of the electrons is stored in the low energy end of the distribution. + However. our conclusions below clo not depend on our assumptions for j44. In.," However, our conclusions below do not depend on our assumptions for $\gamma_{\rm min}$." + general. this standard. ΕΠΗ interpretation of the inner lobes requires substantially. higher mass densities within the outer lobes than can be explained. with the material transported. along the jet Hows that inllated these structures.," In general, this standard FRII interpretation of the inner lobes requires substantially higher mass densities within the outer lobes than can be explained with the material transported along the jet flows that inflated these structures." + For the alternative model we assume that the density and pressure inside the outer lobes is insullicient to significantly slow clown the advance of the inner. voung jets.," For the alternative model we assume that the density and pressure inside the outer lobes is insufficient to significantly slow down the advance of the inner, young jets." + In. this case. the inner jets do not develop very strong shocks at their ends or inflate substantial lobes.," In this case, the inner jets do not develop very strong shocks at their ends or inflate substantial lobes." + We do not expect the ends of the jets to be sites of ellicient. particle acceleration and so the inner jets do not directly. produce much. racio emission., We do not expect the ends of the jets to be sites of efficient particle acceleration and so the inner jets do not directly produce much radio emission. + Ilowever. they push aside and compress. the contents of the outer lobes.," However, they push aside and compress the contents of the outer lobes." + The advance of| the inner jets is last and therefore leads to the formation of a bow shock within the outer lobes., The advance of the inner jets is fast and therefore leads to the formation of a bow shock within the outer lobes. + This bow shock can re-energise he relativistic electrons and also strengthens the magnetic iclcl by compressing it., This bow shock can re-energise the relativistic electrons and also strengthens the magnetic field by compressing it. + We icentify the radio svnchrotron emission observed [from the inner lobes with that produced w the shocked. and. compressed. outer lobe material., We identify the radio synchrotron emission observed from the inner lobes with that produced by the shocked and compressed outer lobe material. +" The model was discussed. brielly in ιν, where we also point out hat bow shock structures. but no lobes directly inllated by jets. are found in numerical simulations of restarting jets (?).."," The model was discussed briefly in \citet{bks07}, where we also point out that bow shock structures, but no lobes directly inflated by jets, are found in numerical simulations of restarting jets \citep{cb91}." + Here. we expand. the description of the model to see whether our observations of 11450|333 ancl 11834|620 are consistent with its predictions., Here we expand the description of the model to see whether our observations of 1450+333 and 1834+620 are consistent with its predictions. + A very similar model was recently developed independently by 2? to describe the inner obes ofthe DORG PIXS 321., A very similar model was recently developed independently by \citet{ssb08} to describe the inner lobes of the DDRG PKS $-$ 321. + In the restlrame of the bow shock propagating through he outer lobe ahead of the inner jet. the momentunm supplied. by the jet. must. be balanced by. the momentum of the receding lobe material," In the restframe of the bow shock propagating through the outer lobe ahead of the inner jet, the momentum supplied by the jet must be balanced by the momentum of the receding lobe material." +" We must also. include the oessures of the jet material. p. and of the lobe material. pj. in this balance ancl so (e.g.7) where 3) is the speed of the bow shock in the rest[rame of the outer lobe material and. Jp, is the speed. of the jet material in the restfreume of the bow shock."," We must also include the pressures of the jet material, $p_{\rm j}$, and of the lobe material, $p_{\rm l}$, in this balance and so \citep[e.g.][]{gb94} + where $\beta_{\rm j}$ is the speed of the bow shock in the restframe of the outer lobe material and $\beta_{\rm jb}$ is the speed of the jet material in the restframe of the bow shock." +" 5j and στ, are the corresponding Lorentz factors.", $\gamma_{\rm b}$ and $\gamma _{\rm jb}$ are the corresponding Lorentz factors. + From the relativistic transformation of velocities we get 8.5EEη(23;hy., From the relativistic transformation of velocities we get $\gamma _{\rm jb}^{2} \beta_{\rm jb}^2 = \gamma_{\rm b}^2 \beta_{\rm j}^2 \left( \beta_{\rm j} - \beta _{\rm b} \right)^2$. +" The jet has a cross-section of 24, and the bow shock has a surface area of li.", The jet has a cross-section of $A_{\rm j}$ and the bow shock has a surface area of $A_{\rm b}$. + The relativistic enthalpy wis defined as =e|p with c the internal energy. density. including the restmass cnerey density. ancl p the pressure.," The relativistic enthalpy $w$ is defined as $w=e+p$ with $e$ the internal energy density, including the restmass energy density, and $p$ the pressure." + For simplicity we assume that all gases can be deseribed by an equation of state of the form, For simplicity we assume that all gases can be described by an equation of state of the form + For simplicity we assume that all gases can be deseribed by an equation of state of the form., For simplicity we assume that all gases can be described by an equation of state of the form +"Extrasolar planets show a huge diversity in their properties and this has important implications [or theories of planet formation, structure and evolution.","Extrasolar planets show a huge diversity in their properties and this has important implications for theories of planet formation, structure and evolution." +" Systems with high orbital inclinations, in which the planet transits across the face of the host star as seen from Earth, are extremely valuable as they allow us to precisely measure many fundamental planetary properuies, including radius, mass and density, which can be used to test these theories (?).."," Systems with high orbital inclinations, in which the planet transits across the face of the host star as seen from Earth, are extremely valuable as they allow us to precisely measure many fundamental planetary properties, including radius, mass and density, which can be used to test these theories \citep{Haswell10}." + The parameter space which we are able to explore with transiting planets is biased by instrumental and observational limitations., The parameter space which we are able to explore with transiting planets is biased by instrumental and observational limitations. +" However, many challenges faced by the current surveys are being overcome by the ability to decrease systematic noise and optimise follow-up strategies."," However, many challenges faced by the current surveys are being overcome by the ability to decrease systematic noise and optimise follow-up strategies." +" Although the majority of the ~ 100 wransiting planets thus far discovered are short period, Jupiter-sized objects, they show a remarkable variety in their physical and dynamical characteristies, such às the extreme eccentricity of HD 80606b (22???),, the ultra-short period of WASP-19b (2) and the puzzlingly low densities of WASP-17b (2) and Kepler-7b (?).."," Although the majority of the $\sim$ 100 transiting planets thus far discovered are short period, Jupiter-sized objects, they show a remarkable variety in their physical and dynamical characteristics, such as the extreme eccentricity of HD 80606b \citep{Naef01,Laughlin09, Moutou09,Fossey09,Garcia09}, the ultra-short period of WASP-19b \citep{Hebb10} and the puzzlingly low densities of WASP-17b \citep{Anderson10} and Kepler-7b \citep{Latham10}." +" Here we describe the properties of a new transiting planet discovered by the SuperWASP survey, WASP-37b."," Here we describe the properties of a new transiting planet discovered by the SuperWASP survey, WASP-37b." + SuperWASP has been a major contributor to the discovery of bright (9 2 \sigma$ from the /Tycho-2 proper motion for the cluster of $\mu_{\alpha}$, $\mu_{\delta}$ ) = $-30.0 \pm 0.3$, $27.8 \pm 0.3$ ) $^{-1}$ (Mamajek et al." +⇀↦ ↾↓∖↓↥∢⊾↓⋅∢⋅⊔↓⋜↧⊀↓⊔⊲↓⊔⋏∙≟≼∼⋜⋯∠∐∠⇂⋜⋯⋅↕⊳∖∐⊳∖↿⋯⇂⋜↧⊳∖⋅⇀∖⊔∪⊔∙∖⇁⊔↓∪⊔⊳∖↼⊲↓⊔ ↾↓∖⋜↧∣⋡↓⋖⊾↓⋡≱∖⋠↓⊔≼∙∢⋅∖∖⊽∢⊾⋅∪⊔⊔∠⇂⊔∪≼∼⋜∐⋜↧↓∪⋏∙," The remaining candidate is listed as `Anonymous' in Table 1, since we found no catalogue entries for this star." +≟⋯⋅∢⊾⊔↓⋰⊓⊾⊳∖⇂⋅∪↓⋅⇂↓↕⊀↓⊳∖⊳∖↿⋜⊔⋅⊳ 3oth the spectrum and the colours of the star are consistent with a M5 spectral classification., Both the spectrum and the colours of the star are consistent with a M5 spectral classification. + “Phe star has weak Hoa emission and no detectable AGTOT Li Labsorption line., The star has weak $\alpha$ emission and no detectable $\lambda 6707$ Li I absorption line. + Study of on-line scanned. plates used to compile the USNO-A2.0 catalogue (Monet et al., Study of on-line scanned plates used to compile the USNO-A2.0 catalogue (Monet et al. + 1998) show the star has high proper motion (see Fig., 1998) show the star has high proper motion (see Fig. + 6)., 6). + We determined the position of the star against USNO-A2.0 positions for several nearby stars at d epochs (1978.10. 1986.18. 1996.13 and. 2000.12): the first 3 from scanned plates available from USNO. and the last from analysis of a SAAO CCD image.," We determined the position of the star against USNO-A2.0 positions for several nearby stars at 4 epochs (1978.10, 1986.18, 1996.13 and 2000.12); the first 3 from scanned plates available from USNO, and the last from analysis of a SAAO CCD image." + From these positions we derived (yi.. fis) 22 (140. 360) +," From these positions we derived $\mu_{\alpha}$, $\mu_{\delta}$ ) $\approx$ $-140$, 360) $^{-1}$." + Phe USNO-A2.0 catalogue rejects stars with pio>300 masvyr.ο., The USNO-A2.0 catalogue rejects stars with $\mu > 300$ $^{-1}$. + The star is likely a dMle., The star is likely a dMe. + HU it is a main-sequence star. then the star has a distance of 50 pe ancl tangential velocity of ~100 kiss," If it is a main-sequence star, then the star has a distance of $\sim 50$ pc and tangential velocity of $\sim 100$ $^{-1}$." + Lillenbranel Mever (1999) examined the frequency of disks as à function of stellar age for nearby. clusters ancl star forming regions and found that disks disperse on a timescale of <10 Myr. with few clisks remaining at ages 215 Alve.," Hillenbrand Meyer (1999) examined the frequency of disks as a function of stellar age for nearby clusters and star forming regions and found that disks disperse on a timescale of $< 10$ Myr, with few disks remaining at ages $> 15$ Myr." + llaisch. Lada Lada (2001) conducted a similar study on the stellar population of several voung clusters. concluding that essentially all stars lose their cisks within 6 Myr.," Haisch, Lada Lada (2001) conducted a similar study on the stellar population of several young clusters, concluding that essentially all stars lose their disks within $\sim 6$ Myr." + The prevalence of disks in the =9 Myr-old η Cha cluster will be examined by Lyo et al. (, The prevalence of disks in the $\approx 9$ Myr-old $\eta$ Cha cluster will be examined by Lyo et al. ( +in preparation),in preparation). + An old and nearby CET star such as ECILA JOS43.37905 is therefore rare., An `old' and nearby CTT star such as ECHA J0843.3--7905 is therefore rare. + A similar star is the ςΕΕ star TW Ίνα. also nearby (α=56 pc) and zzLO Myr old. (Webb et al.," A similar star is the CTT star TW Hya, also nearby $d = 56$ pc) and $\approx 10$ Myr old (Webb et al." + 1999)., 1999). + Phe resolved pole-on disk surrounding TW. Hya has been a focus for study of disk structure (e.g. Frilling et al., The resolved pole-on disk surrounding TW Hya has been a focus for study of disk structure (e.g. Trilling et al. + 2001) and the early. planet. formation environment., 2001) and the early planet formation environment. + ECLILA JOSA43.3.7905 is of later spectral tvpe and lower mass than TW Liva., ECHA J0843.3–7905 is of later spectral type and lower mass than TW Hya. + Imaging studies of ECILA 0843.3.7905 could eive valuable insight into the nature of evolved. disks (and indirectly. planets) around. dwarf M stars., Imaging studies of ECHA J0843.3–7905 could give valuable insight into the nature of evolved disks (and indirectly planets) around dwarf M stars. + Like other CPT stars. ECILA 0843.3.7905. has a strong infrared. excess.," Like other CTT stars, ECHA J0843.3–7905 has a strong infrared excess." + Analysis of L-band imagine of the cluster core obtained with the South Pole Infrared. Explorer telescope during. L999 (Lyo ct al.," Analysis of $L$ -band imaging of the cluster core obtained with the South Pole Infrared Explorer telescope during 1999 (Lyo et al.," + in. preparation) found L—7.8 for this star., in preparation) found $L = 7.8$ for this star. + Assuming an M2 spectral type. the L-band excess is z2 mag.," Assuming an M2 spectral type, the $L$ -band excess is $\approx 2$ mag." + The FSC entry for the star indicates high-quality 25- ancl 60-77 Iuxes., The FSC entry for the star indicates high-quality 25- and $\mu$ m fluxes. + Our survey of &40% of the known extent of the ‘luster found. 2 new cluster members not detected by the ciscovery URL image of Alamajek et al. (, Our survey of $\approx 40$ of the known extent of the cluster found 2 new cluster members not detected by the discovery HRI image of Mamajek et al. ( +1999. 2000). thereby increasing the number of stellar. primaries to 15.,"1999, 2000), thereby increasing the number of stellar primaries to 15." + A survey of similar depth across the cluster mig= therefore find only several more new members., A survey of similar depth across the cluster might therefore find only several more new members. + This resul appears to be at odds. with Alamajek et al. (, This result appears to be at odds with Mamajek et al. ( +2000) who predicted. [rom consideration of the cluster IAIF that the stellar population was 204. the (then) known number of 13 primaries.,2000) who predicted from consideration of the cluster IMF that the stellar population was $2-4 \times$ the (then) known number of 13 primaries. + Our survey might indicate that the cluster extent is not constrained bv the HIE field anc wt 15/—40 primaries await discovery bevond the LUA »undarv., Our survey might indicate that the cluster extent is not constrained by the HRI field and that $15-40$ primaries await discovery beyond the HRI boundary. + Alternatively. if the cluster is constrained. by 1e LR eld. then the low success rate of our study sugeests the cluster may contain as few as zz20 primaries. C," Alternatively, if the cluster is constrained by the HRI field, then the low success rate of our study suggests the cluster may contain as few as $\approx 20$ primaries. (" +Ehese estimates do not address the brown dwarl population expected: to accompany the stellar members.),These estimates do not address the brown dwarf population expected to accompany the stellar members.) + Either of the above population scenarios indicates the X-ray survey must jwe been relatively complete at detecting cluster members within the LIRE field., Either of the above population scenarios indicates the X-ray survey must have been relatively complete at detecting cluster members within the HRI field. + This result confirms the unusual skewness of the ratio of X-ray luminosity to. bolometric uminosity noted by Mamajek et al. (, This result confirms the unusual skewness of the ratio of X-ray luminosity to bolometric luminosity noted by Mamajek et al. ( +1999. 2000). with most of the late-twpe RECN stars having a flux ratio near the saturation level of log Lx μη®3.,"1999, 2000), with most of the late-type RECX stars having a flux ratio near the `saturation' level of log $L_{X}$ $L_{\rm bol} \approx -3$." + We thank the SAAQ and AISSSO time allocation committees for telescope time during. L999 ancl 2000., We thank the SAAO and MSSSO time allocation committees for telescope time during 1999 and 2000. + WAL and LAC thank the stall of SAAO) for. their assistance., WAL and LAC thank the staff of SAAO for their assistance. + WAL ancl EEM thank the MSSSO support stall for their guidance in operating the 2.3-m telescope., WAL and EEM thank the MSSSO support staff for their guidance in operating the 2.3-m telescope. + WAL acknowledges financial support from the Australian Research Council Small Grant Scheme ancl University College Special Research Grants., WAL acknowledges financial support from the Australian Research Council Small Grant Scheme and University College Special Research Grants. + LAC is supported by a NRE Post-gracluate Scholarship., LAC is supported by a NRF Post-graduate Scholarship. + ELAL thanks the SUVPE Legacy Science Program for support., EEM thanks the SIRTF Legacy Science Program for support. + EDI’s research. is supported in-part by NASA contracts NASS-38252° and NAG5-8422., EDF's research is supported in-part by NASA contracts NAS8-38252 and NAG5-8422. + his research made use of the POSS-LL, This research made use of the POSS-II +system produces a constant emission spectrum throughout the orbit and that this spectrum is alfected by the absorption component at phases 0.350.55.,system produces a constant emission spectrum throughout the orbit and that this spectrum is affected by the absorption component at phases 0.35–0.55. + As we can see in Fig., As we can see in Fig. + d. the variations in the EWs are small. if we exclude the phases of eclipse.," 4, the variations in the EWs are small, if we exclude the phases of eclipse." + To obtain the absorption spectrum of this component. we decided to subtract the averaged spectrum of V348 Pup corresponding to the orbital phases outside eclipse ancl outside the absorption events to the averaged spectrum at the absorption phases.," To obtain the absorption spectrum of this component, we decided to subtract the averaged spectrum of V348 Pup corresponding to the orbital phases outside eclipse and outside the absorption events to the averaged spectrum at the absorption phases." + Before the subtraction. we corrected all the spectra for orbital motion. using the semi-amplitude of the radial velocity curve of 15 and the corresponding phase ollsct (see Table 2).," Before the subtraction, we corrected all the spectra for orbital motion, using the semi-amplitude of the radial velocity curve of $\beta$ and the corresponding phase offset (see Table 2)." +" This analysis has to »e taken with extreme caution. since the presence of another emission component in certain orbital phases would. prevent he subtraction of being the spectrum of the absorption The spectrum of the absorption component in. V348 ""up is shown in Fig."," This analysis has to be taken with extreme caution, since the presence of another emission component in certain orbital phases would prevent the subtraction of being the spectrum of the absorption The spectrum of the absorption component in V348 Pup is shown in Fig." + 10., 10. + The A4686 and the Bowen Blend remain in emission. indicating that the absorption component is absent. from these lines.," The $\lambda$ 4686 and the Bowen Blend remain in emission, indicating that the absorption component is absent from these lines." + Little emission. also remains in Ες and L9., Little emission also remains in $\gamma$ and $\delta$. + This is due to the dillerent. semi- anc phase olfsets of their radial velocity curves with respect to Lh. when correcting for orbital motion.," This is due to the different semi-amplitudes and phase offsets of their radial velocity curves with respect to $\beta$, when correcting for orbital motion." + ‘Taking this into account. the llux ratios in the absorption spectrum of V348 Pup are very similar to those in the spectrum of a BO V star (taken from Jacoby. Hunter Christian 1984). as can be seen in Fig.," Taking this into account, the flux ratios in the absorption spectrum of V348 Pup are very similar to those in the spectrum of a B0 V star (taken from Jacoby, Hunter Christian 1984), as can be seen in Fig." + 10., 10. + This may bea coincidence. but the fact that we get à spectral tvpe similar to the observed by Ciroot et al. (," This may be a coincidence, but the fact that we get a spectral type similar to the observed by Groot et al. (" +2001) is If the spectrum we obtained is actually the spectrum of the absorption component. and is similar to that of a D star. the absorption events in SW Sex stars could be produced by an atmosphere.,"2001) is If the spectrum we obtained is actually the spectrum of the absorption component, and is similar to that of a B star, the absorption events in SW Sex stars could be produced by an atmosphere." + Following the magnetic model proposed. by guez-Gil ct al. (, Following the magnetic model proposed by guez-Gil et al. ( +2001). we suggest that the absorption is produced. by a vertically extended: atmosphere.,"2001), we suggest that the absorption is produced by a vertically extended atmosphere." + The gas gaream from the secondary hits the disc. and. part of the stream. which is thicker than the cise itself. overflows it.," The gas stream from the secondary hits the disc, and part of the stream, which is thicker than the disc itself, overflows it." + iXnother shock is then produced. when the stream meets 10 magnetosphere of the white dwarf., Another shock is then produced when the stream meets the magnetosphere of the white dwarf. + We Lgugeest that an tmosphere is formed arouncl this shock., We suggest that an atmosphere is formed around this shock. + Phe temperature of 1¢ photospheres of BO-type stars ranges between 1900011$ year. + Four of such sources were monitored following the cessation of their long X-ray outburst. which revealed that the thermal X-ray emission decaved over the course of vears (e.g..Wijnands2010:Fridrikssonetal.2011:DiazTrigo 2011).," Four of such sources were monitored following the cessation of their long X-ray outburst, which revealed that the thermal X-ray emission decayed over the course of years \citep[e.g.,][]{wijnands2001,wijnands2003,cackett2008,cackett2010,degenaar2010_exo2,fridriksson2011,diaztrigo2011}." +. The observed decrease in thermal radiation has been interpreted as cooling of the neutron star crust. which became severely heated during the prolonged. accretion outburst (Itutledgeetal.2002:Wijnancds2004)..," The observed decrease in thermal radiation has been interpreted as cooling of the neutron star crust, which became severely heated during the prolonged accretion outburst \citep[][]{rutledge2002,wijnands04_quasip}." + Confronting the observed cooling curves with neutron star thermal evolution cocles eives information about the amount of heat release and thermal conductivity of the neutron star crust. as well as the properties of the stellar core (Shterninetal.2007:Brown&Cumming2009).," Confronting the observed cooling curves with neutron star thermal evolution codes gives information about the amount of heat release and thermal conductivity of the neutron star crust, as well as the properties of the stellar core \citep[][]{shternin07,brown08}." +. The unusually long outburst. duration of these quasi-persistent neutron star LAINBs provides. the necessary conditions to significantly lift the neutron star. crust temperature so that the thermal relaxation becomes observable., The unusually long outburst duration of these quasi-persistent neutron star LMXBs provides the necessary conditions to significantly lift the neutron star crust temperature so that the thermal relaxation becomes observable. + However. Brownetal.(1905) argued that in regular transients with outburst durations of weeks. the crust might also become significantly heated above the core temperature. provided that the outburst is bright enough compared to the quiescent base level.," However, \citet{brown1998} argued that in regular transients with outburst durations of weeks, the crust might also become significantly heated above the core temperature, provided that the outburst is bright enough compared to the quiescent base level." + As noted by Degenaar&Wijnancds (2011).. the high outburst luminosity ane relatively faint quiescent level of JIT4S would then make it à good candidate to search for neutron star crust cooling.," As noted by \citet[][]{deeg_wijn2011}, the high outburst luminosity and relatively faint quiescent level of J1748 would then make it a good candidate to search for neutron star crust cooling." + In this letter. we report on a DDirector's Discretionary Time (DDT) observation of the elobular cluster Terzan 5. obtained after the cessation of the 2010 outburst of the newly discovered 11 Hz X-ray pulsar.," In this letter, we report on a Director's Discretionary Time (DDT) observation of the globular cluster Terzan 5, obtained after the cessation of the 2010 outburst of the newly discovered 11 Hz X-ray pulsar." + The aim of this observation was to study the ellect of the bright accretion outburst on the thermal properties of the Dheutron star crust., The aim of this observation was to study the effect of the bright accretion outburst on the thermal properties of the neutron star crust. + ‘Terzan 5 was observed with the /ACIS-8 on 2011 February 17 [rom 09:0617:58 for 29.7 ks (ID 13225)., Terzan 5 was observed with the /ACIS-S on 2011 February 17 from 09:06–17:58 for 29.7 ks (ID 13225). + The cluster is positioned on the $8 chip and the data was obtained in the faint mode with the nominal frame time of 3.2 s. Figure 2. clisplavs an image of the 2011 delata. together with a 36.4 ks archival observation that was obtained on 2009 July 15.16 (seeDegenaar&Wijnands 2011).," The cluster is positioned on the S3 chip and the data was obtained in the faint mode with the nominal frame time of 3.2 s. Figure \ref{fig:ds9} displays an image of the 2011 data, together with a $\sim36.4$ ks archival observation that was obtained on 2009 July 15–16 \citep[see][]{deeg_wijn2011}." +. Ht is clear from these images that JITAS is brighter in 2011. ~2 months after the cessation of the 2010 activity. iui d was in 2009. —14 months prior to that outburst.," It is clear from these images that J1748 is brighter in 2011, $\sim2$ months after the cessation of the 2010 activity, than it was in 2009, $\sim 14$ months prior to that outburst." + For the purpose of directly comparing the new data with archival observations. we use the same analysis ancl reduction steps as outlined in Degenaar&Wijnancls (2011).. emploving the software tools (v. 4.2).," For the purpose of directly comparing the new data with archival observations, we use the same analysis and reduction steps as outlined in \citet{deeg_wijn2011}, employing the software tools (v. 4.2)." + No background. [lares occurred. during the observation. so all data was used in further analysis.," No background flares occurred during the observation, so all data was used in further analysis." +" Source count rates and lighteurves were extracted. from a 1"" circular. region. centred at the source position. using the tool DMENTRACT."," Source count rates and lightcurves were extracted from a $1''$ circular region, centred at the source position, using the tool ." +". Corresponding background events. were obtained from a circular region with a radius of 40"". positioned on a source- part of the CCD that was located ~14’ west of the cluster core."," Corresponding background events were obtained from a circular region with a radius of $40''$, positioned on a source-free part of the CCD that was located $\sim1.4'$ west of the cluster core." + Our target is detected. at a count rate of (6.5+10*countss. |. which is about 6 times higher than observed in archival data obtained in 2003 and 2009 (seeDe-eenaar&Wijnands 2011).," Our target is detected at a count rate of $(6.5\pm 0.5)\times10^{-3}~\cnts$ , which is about 6 times higher than observed in archival data obtained in 2003 and 2009 \citep[see][]{deeg_wijn2011}." +. A total of 192 net sourcephotons were collected., A total of 192 net sourcephotons were collected. + We obtained source and background spectra using the tool PSENTRACT.. and. generated redistribution matrices (rmf) and. ancillary response files (arf) withthe tasks and MKARE.. respectively.," We obtained source and background spectra using the tool , and generated redistribution matrices (rmf) and ancillary response files (arf) withthe tasks and , respectively." + We group the spectrum to contain a minimum of 20 photons per bin. and," We group the spectrum to contain a minimum of 20 photons per bin, and" +"time""), and the time elapsed since the last significant infall of blue galaxies with Ms"">1019 occurred.","time”), and the time elapsed since the last significant infall of blue galaxies with $M^{\mathrm{stars}} \ge 10^{10}~\mathrm{M}_{\sun}$ occurred." +" These two timescales need to be compared with Mg,the minimum life-time of the stellar populations characterizing old, passively evolving (i.e., red-sequence) galaxies, which is about 2Gyr."," These two timescales need to be compared with the minimum life-time of the stellar populations characterizing old, passively evolving (i.e., red-sequence) galaxies, which is about $2~\mathrm{Gyr}$." + We now make the conservative assumption that the two fossil groups behave as other similarly massive groups observed at z£&0.4., We now make the conservative assumption that the two fossil groups behave as other similarly massive groups observed at $z \approx 0.4$. +" Furthermore, we assume that groups do not significantly increase their total masses between 2£0.5 and 2&0.415, in a statistical sense."," Furthermore, we assume that groups do not significantly increase their total masses between $z \approx 0.5$ and $z \approx 0.4$, in a statistical sense." +" By doing so, we can use results on the median formation redshifts of DM haloes obtained from the simulations of Gao et al. ("," By doing so, we can use results on the median formation redshifts of DM haloes obtained from the simulations of Gao et al. (" +2004) or different analytical models that are reproduced in fig.,2004) or different analytical models that are reproduced in fig. + 5 of Giocoli et al. (, 5 of Giocoli et al. ( +2007).,2007). +" Groups of similar total masses to 0095951+0212.6 and 0095951--0140.8 have assembled half of their DM halo at typical redshifts of ~1.2 and ~ 1.05, respectively, which correspond to similar lookback times of z4Gyr at z=0.425 and z=0.372."," Groups of similar total masses to $+$ 0212.6 and $+$ 0140.8 have assembled half of their DM halo at typical redshifts of $\sim 1.2$ and $\sim 1.05$ , respectively, which correspond to similar lookback times of $\approx 4~\mathrm{Gyr}$ at $z = 0.425$ and $z = 0.372$." +" According to the computations of dynamical friction and galaxy-merging timescales of Boylan-Kolchin, Ma Quataert (2008, their fig."," According to the computations of dynamical friction and galaxy-merging timescales of Boylan-Kolchin, Ma Quataert (2008, their fig." +" 1), galaxies with M***s~101° or 1011 need about 6 or 3 Gyr, respectively, to reach an orbit Mg,with maximum group-centric distance of 0.5[ρου from Roo for the first time."," 1), galaxies with $M^{\mathrm{stars}} \sim 10^{10}$ or $10^{11}~\mathrm{M}_{\sun}$ need about 6 or 3 Gyr, respectively, to reach an orbit with maximum group-centric distance of $0.5 R_{200}$ from $R_{200}$ for the first time." + Since the simulations of Boylan-Kolchin et al., Since the simulations of Boylan-Kolchin et al. +" mostly neglect the effect of the central galaxy in a given DM halo on the merging time-scale, and do not take into account the drag on ambient gas (e.g. Ostriker 1999), we infer that the epoch of last infall of blue galaxies with M$t*'$>1019 can date back to less than 3-6 Gyr with respect to the Molookback time of either fossil group."," mostly neglect the effect of the central galaxy in a given DM halo on the merging time-scale, and do not take into account the drag on ambient gas (e.g. Ostriker 1999), we infer that the epoch of last infall of blue galaxies with $M^{\mathrm{stars}} \ge 10^{10}~\mathrm{M}_{\sun}$ can date back to less than 3–6 Gyr with respect to the lookback time of either fossil group." +" The time of last infall of blue, intermediate-mass galaxies is comparable to the median time elapsed since the assembly of a group as massive as either fossil group, if not longer than that."," The time of last infall of blue, intermediate-mass galaxies is comparable to the median time elapsed since the assembly of a group as massive as either fossil group, if not longer than that." +" Thus, we can assume that most of the galaxies with Λ.Α>101° have resided within the bound region of either fossil Mg,group since its assembly (cf."," Thus, we can assume that most of the galaxies with $M^{\mathrm{stars}} \ge 10^{10}~\mathrm{M}_{\sun}$ have resided within the bound region of either fossil group since its assembly (cf." + von Benda-Beckmann et al., von Benda-Beckmann et al. + 2008; Dariush et al., 2008; Dariush et al. + 2010)., 2010). +" This “residence” time of z4Gyr is longer than the typical life-times of young stars, which are responsible for the bulk luminosity shortward of the 4000 A--break, when they are present."," This “residence” time of $\approx 4~\mathrm{Gyr}$ is longer than the typical life-times of young stars, which are responsible for the bulk luminosity shortward of the 4000 -break, when they are present." +" Hence, there is enough time to observe a dominating population of old, passively evolving galaxies with ΛΑ1015Mg, in the two fossil groups at z~0.4."," Hence, there is enough time to observe a dominating population of old, passively evolving galaxies with $M^{\mathrm{stars}} \ge 10^{10}~\mathrm{M}_{\sun}$ in the two fossil groups at $z \approx 0.4$." +" This is even more the case if the assembly time of a fossil group is shorter than the median assembly time for all coeval, similarly massive groups (D'Onghia et al."," This is even more the case if the assembly time of a fossil group is shorter than the median assembly time for all coeval, similarly massive groups (D'Onghia et al." + 2005; Dariush et al., 2005; Dariush et al. + 2007)., 2007). +" Environmental effects could be responsible for the raise of the population of old, passively evolving galaxies in the two fossil groups (see Boselli Gavazzi 2006 for a recent review)."," Environmental effects could be responsible for the raise of the population of old, passively evolving galaxies in the two fossil groups (see Boselli Gavazzi 2006 for a recent review)." + Member-candidate galaxies of either fossil group which are satellites at z20.4 could reach stellar masses of 8x10?Mo before their star formation turned to its end., Member-candidate galaxies of either fossil group which are satellites at $z \approx 0.4$ could reach stellar masses of $8 \times 10^{10}~\mathrm{M}_{\sun}$ before their star formation turned to its end. +" We note that, in spite of the generally held, negative feedback exerted by the group environment on the star formation activity of member galaxies, such an activity seems to be enabled in of the galaxies with M*'?'*~10!!Mg, in groups at similar redshifts and with similar masses (Fig."," We note that, in spite of the generally held, negative feedback exerted by the group environment on the star formation activity of member galaxies, such an activity seems to be enabled in of the galaxies with $M^{\mathrm{stars}} \sim 10^{11}~\mathrm{M}_{\sun}$ in groups at similar redshifts and with similar masses (Fig." + 2; Giodini et al., 2; Giodini et al. + in preparation;see also Wilman et al., in preparation;see also Wilman et al. + 2008 and references therein)., 2008 and references therein). +" It is easy to understand how tidal (e.g.,"," It is easy to understand how tidal (e.g.," +and 187199 davs. with no clata-model mismatches.,"and 187–199 days, with no data-model mismatches." + The more extended. interval does not allow a solution if we require that all data points match the model., The more extended interval does not allow a solution if we require that all data points match the model. + However. if we allow one mismatch. periods of 159 days and 189199 days are allowed.," However, if we allow one mismatch, periods of 159 days and 189–199 days are allowed." + The entire dataset has a periodic solution only if we allow (wo observations to be mismatched. “, The entire dataset has a periodic solution only if we allow two observations to be mismatched. “ +Periods” with only (wo mismatches are 156 days. 159 days. 172174 days. and 193197 days.,"Periods” with only two mismatches are 156 days, 159 days, 172–174 days, and 193–197 days." +" Thus. no matter which data subset was chosen. ""periods around 159 and 195 days best matched the data. but the longer the interval. the more mismatches are required."," Thus, no matter which data subset was chosen, “periods” around 159 and 195 days best matched the data, but the longer the interval, the more mismatches are required." + We also imvestigated if a shorter period («60 davs) might match the data (Figure 2 right)., We also investigated if a shorter period $<$ 60 days) might match the data (Figure \ref{periods} right). + Using the Msec interval. the best solutions (45 days) would produce two mismatches. usually the 2004 July oobservation. and either the first or last observations of an extended set of oobservations.," Using the Msec interval, the best solutions (45 days) would produce two mismatches, usually the 2004 July observation, and either the first or last observations of an extended set of observations." + Longer fit intervals produced greater numbers of mismatches for their best fits., Longer fit intervals produced greater numbers of mismatches for their best fits. + This analvsis suggests that hhas a characteristic tme-scale., This analysis suggests that has a characteristic time-scale. + It is not a strict periodicity. alübough it is still possible (hat il is a superposition of periodic and random variations.," It is not a strict periodicity, although it is still possible that it is a superposition of periodic and random variations." +" We present the light curves of 2004 July and 2004 December/2005 January high states in rellejuly ancl rellcjan.. respectively,"," We present the light curves of 2004 July and 2004 December/2005 January high states in \\ref{lcjuly} and \\ref{lcjan}, respectively." + The high deeree of short-timescale variability is obvious., The high degree of short-timescale variability is obvious. + Moreover. we confirm our earlier results (Alukaietal.2003) of the energy dependence of the variability: iis more variable at higher energies.," Moreover, we confirm our earlier results \citep{Mea2003} + of the energy dependence of the variability: is more variable at higher energies." + The power densitv spectra (PDS) of disk-accreting X-ray binaries can be classified into two broad categories based on whether the Eddinegton accretion rate (/-2L/ Lj) is above or below the critical value. /.. regardless of the nature of the primary. (vanderlis1994).," The power density spectra (PDS) of disk-accreting X-ray binaries can be classified into two broad categories based on whether the Eddington accretion rate $l$ $L$ $L_{\rm Edd}$ ) is above or below the critical value, $l_c$, regardless of the nature of the primary \citep{vdK1994}." +. At low accretion rate (/« /.). X-ray binaries have a fractional rms amplitude of a few times and a broken power law PDS (hereafter Type A}; at high aceretion rate (/> /.). they have a [ractional rms of a few percent. with power law (index 11.5) PDS (Type DB).," At low accretion rate $l < l_c$ ), X-ray binaries have a fractional rms amplitude of a few times and a broken power law PDS (hereafter Type A); at high accretion rate $l > l_c$ ), they have a fractional rms of a few percent, with power law (index 1–1.5) PDS (Type B)." + Although (here is a well-known hysteresis effect (Alivamotoetal.1995) that complicates the situation. Barnardetal.(2004). have nevertheless obtained an empirical calibration of { of 0.1. in," Although there is a well-known hysteresis effect \citep{Mea1995} that complicates the situation, \citet{Bea2004} have nevertheless obtained an empirical calibration of $l_c$ of 0.1, in" +partially resolved.,partially resolved. + We illustrate in Fie., We illustrate in Fig. + l a prominent compact IRC (Object I8) related to an optical nebula. and in Fie.," 1 a prominent compact IRC (Object 48) related to an optical nebula, and in Fig." + 2 oue (Object 131) related to a radio nebula., 2 one (Object 131) related to a radio nebula. + A loose IRC (Object 111) is shown Fig., A loose IRC (Object 114) is shown Fig. + 3., 3. + À deeply embedded ii nebular cussion and/or reflection partially resolved IRC (Object 33) is given in Fie., A deeply embedded in nebular emission and/or reflection partially resolved IRC (Object 33) is given in Fig. + L, 4. + This object has a clear dust lane. aud remarkably resembles the wellknown NGC2021 cluster (c.g. Lada et al.," This object has a clear dust lane, and remarkably resembles the well-known NGC2024 cluster (e.g. Lada et al." + 1991)., 1991). + The infrared cluster (Object L16) in the radio nebula 3327.3-0.5 is also similar to that in 2202I., The infrared cluster (Object 146) in the radio nebula 327.3-0.5 is also similar to that in 2024. + IRCCs are probably clusters. but are esseutially unresolved. aud require higher resolution aud deeper images for definitive diagnostic (e.g. Object 8 Fig.," IRCCs are probably clusters, but are essentially unresolved, and require higher resolution and deeper images for definitive diagnostic (e.g. Object 8 – Fig." + 5)., 5). + Παν are less dense than IRCs. some are rather compact but little populated.," IRGrs are less dense than IRCs, some are rather compact but little populated." + Au IRCr (Object 20) is eiven in Fie., An IRGr (Object 20) is given in Fig. + 6., 6. +" IROCs (ο,ο, Ohject 6 Fig.", IROCs (e.g. Object 61 – Fig. + 7) have similar appearence to optical open clusters and relatively large angular size (2° or 1ioro)., 7) have similar appearence to optical open clusters and relatively large angular size $\approx$ 2' or more). + Iu Table 3 there occur 16 IRC. 6L Πλέι 7 IRCC and 6 IROC objects. while in Table £31 IRC. 8 IRCs. 16 IRCC and 1 IROC objects.," In Table 3 there occur 46 IRC, 64 IRGr, 7 IRCC and 6 IROC objects, while in Table 4 31 IRC, 8 IRGr, 16 IRCC and 1 IROC objects." + Distances are mostly based ou kinematical estimates for the nebulae (Wilson e al., Distances are mostly based on kinematical estimates for the nebulae (Wilson et al. + 1970. Caswell Tavues 1987). but include as well averages with estinates frou individual stars. when available (c.g. Ceoreclin et al.," 1970, Caswell Haynes 1987), but include as well averages with estimates from individual stars, when available (e.g. Georgelin et al." + 1973. Ceorechin ct al.," 1973, Georgelin et al." + 2000 aud references therein)., 2000 and references therein). +" The ucar/far distance ambiguity has been solved bw some of these authors CR). else we indicate both (R,, or Ry)."," The near/far distance ambiguity has been solved by some of these authors $R$ ), else we indicate both $R_n$ or $R_f$ )." + One prominent cluster (Object 23) appears to be related to the optical nebula 223) (Munich 1955). which is thus uot a planetary ucbula (1892-15 or PNG2625.2|00.1. Acker 1992).," One prominent cluster (Object 23) appears to be related to the optical nebula 23 (Munch 1955), which is thus not a planetary nebula (K2-15 or PNG263.2+00.4, Acker 1992)." + Likewise. the uebulae ESOS313*N10 and ESOII28EN25 ire not planetary ucbulae.," Likewise, the nebulae 313*N10 and 128EN25 are not planetary nebulae." + We included in Table 2 uew embedded: clusters. and stellar groups related to optical reflection nebulae., We included in Table 3 new embedded clusters and stellar groups related to optical reflection nebulae. + This object type has been discussed in Dutra Bica (2001)., This object type has been discussed in Dutra Bica (2001). + They appear to be less massive clusters or stellar eroups where no ionizing star was formed., They appear to be less massive clusters or stellar groups where no ionizing star was formed. + The objects in Soares Bica (2002) are of this type or close to its limit towards lonizing stars., The objects in Soares Bica (2002) are of this type or close to its limit towards ionizing stars. + Detected objects in the van den Berel Ierbst (1975) reflection nebulae are probably of this type., Detected objects in the van den Bergh Herbst (1975) reflection nebulae are probably of this type. +" The 338a complex (often designated 33576 ILIII region) is well-known for harbouring the prominent 33576 IB. cluster (ο. ο, Persi ct al.", The 38a complex (often designated 3576 II region) is well-known for harbouring the prominent 3576 IR cluster (e. g. Persi et al. + 1991)., 1994). + Towever. we realized that it is not related to the original NCC cataogue sinall nebula 33576 (Table 1).," However, we realized that it is not related to the original NGC catalogue small nebula 3576 (Table 1)." + The massive cluster is related to the radio ITT region 2291.251-0.7123. and its optical counterpart 33581.," The massive cluster is related to the radio II region 291.284-0.713, and its optical counterpart 3581." + The sinall nebula 33576 itself has a less massive cluster (Object 65 Table 3). thus {ως a cluster pair in that complex.," The small nebula 3576 itself has a less massive cluster (Object 65 – Table 3), thus forming a cluster pair in that complex." + Owing to the usual 33576 designation in the literature for the massive cluster. woe suggest 33576A for the small one.," Owing to the usual 3576 designation in the literature for the massive cluster, we suggest 3576A for the small one." + Some of the BRC nebulae (oright-rimaimed clouds) have been shown to be sites of star cluster formation (Sugitani et al., Some of the BRC nebulae (bright-rimmed clouds) have been shown to be sites of star cluster formation (Sugitani et al. + 1995)., 1995). + The oeseut study medieates £ BRC clouds with new clusters or stellar exoups (Table 3). together with 2 BRCs in Table 1.," The present study indicates 4 BRC clouds with new clusters or stellar groups (Table 3), together with 2 BRCs in Table 1." + We compare in Fie., We compare in Fig. + 8 the angular distributions in galactic coordinates of the preseut samples (Tables 3 aud 1) with those of previously known infrared objects (Table 1) and the optical open clusters (Table 2)., 8 the angular distributions in galactic coordinates of the present samples (Tables 3 and 4) with those of previously known infrared objects (Table 1) and the optical open clusters (Table 2). + The merease of the sample of known clusters and stellar eroups is overwheliuimg., The increase of the sample of known clusters and stellar groups is overwhelming. + Note that the infrared clusters comune from the radio nebulae are münlv located between 3007«(«c 3507. corresponding to iuternal arius aud where absorption in the Calaxv shows a pronounced iucrease (Dutra Bica 20000).," Note that the infrared clusters coming from the radio nebulae are mainly located between $300^{\circ} < \ell < 350^{\circ}$ , corresponding to internal arms and where absorption in the Galaxy shows a pronounced increase (Dutra Bica 2000b)." + The distance istoerauus for the objects obtained from he optical and radio nebula samples (Tables 3 and L. respectively) are shown in Fie.," The distance histograms for the objects obtained from the optical and radio nebula samples (Tables 3 and 4, respectively) are shown in Fig." + 9., 9. + For the objects with distance anubiguitv we assunied their near distance as a ower limit for the histogram analysis., For the objects with distance ambiguity we assumed their near distance as a lower limit for the histogram analysis. + The histograms show that objects comine from radio nebulae are ou the average more distant peaking at 21.5 kpe. while those roni optical nebulae peak at 222.5 kpc.," The histograms show that objects coming from radio nebulae are on the average more distant peaking at $\approx$ 4.5 kpc, while those from optical nebulae peak at $\approx$ 2.5 kpc." + We conclude that a typical emibedded: cluster discovered with 2MASS is rot very far din the Galaxy. but the sampled clusters are ocated iu more internal arms than Sagittarius-Carina. ike Scutuni-Crux aud bevoud (c.g. Georeclin Goorgeliu 1976).," We conclude that a typical embedded cluster discovered with 2MASS is not very far in the Galaxy, but the sampled clusters are located in more internal arms than Sagittarius-Carina, like Scutum-Crux and beyond (e.g. Georgelin Georgelin 1976)." + The linear size histograms for the objects obtained roni the optical and radio nebula samples (Tables 3 aud respectively) are shown in Fie.," The linear size histograms for the objects obtained from the optical and radio nebula samples (Tables 3 and 4, respectively) are shown in Fig." + 10., 10. + The objects with, The objects with +Upon establishing1.2.6 the existence of a lag. it mst be characterized by not oulv the magnitude of the eracicut in the z direction but also its radial variation.,"Upon establishing the existence of a lag, it must be characterized by not only the magnitude of the gradient in the z direction but also its radial variation." + Using the one-component warped model. by diagrams further frou the center of the ealaxy than those shown in Figure 6 are exanuned in Figure 10 using a 3042307 cube. with the range of slice locations correspondiug to the dashed lines in Figure 2..," Using the one-component warped model, bv diagrams further from the center of the galaxy than those shown in Figure \ref{fig6} are examined in Figure \ref{fig10} using a $\times$ 30"" cube, with the range of slice locations corresponding to the dashed lines in Figure \ref{fig2}." + The data. at these large radii. are best fit with a shallowing lag. decreasing in magnitude to 52 lan s| t+ and [42 kins | | onear a radius of SN (10 kpc) in the approaching aud receding halves respectively.," The data, at these large radii, are best fit with a shallowing lag, decreasing in magnitude to $-5\pm$ 2 km $^{-1}$ $^{-1}$ and $-4\pm$ 2 km $^{-1}$ $^{-1}$ near a radius of 8' (10 kpc) in the approaching and receding halves respectively." +" This is most easily secu ou the teriinal side of cach panel: the data contours become more rounded. whereas the model with a constant lag retains too much of a ""V shape. with data contours extending further ou the terminal side than those of the mocel."," This is most easily seen on the terminal side of each panel: the data contours become more rounded, whereas the model with a constant lag retains too much of a “V"" shape, with data contours extending further on the terminal side than those of the model." + However. ou the negative offset side of the receding half. the shallowing of the lag may be even more substantial.," However, on the negative offset side of the receding half, the shallowing of the lag may be even more substantial." + This is best seen in panels corresponding to -8.2 and -9., This is best seen in panels corresponding to -8.2' and -9'. + Uncertaiufies are still quantified. based on the positive offset side. aud — [dan 3 ο is ouly given as an upper limit for the side with negative offset.," Uncertainties are still quantified based on the positive offset side, and $-4$ km $^{-1}$ $^{-1}$ is only given as an upper limit for the side with negative offset." + There is no strong indication of auv radial variation at radi smaller than approxinatelv 65 (7.5 kpc)., There is no strong indication of any radial variation at radii smaller than approximately 6' (7.5 kpc). + It should be noted that the S/N of the data is considerably cdinuuished at aud, It should be noted that the S/N of the data is considerably diminished at and +Moviel.,Movie1.— + Shows the orbits of individual stars near (he Galactic Center as measured with high resolution infrared observations., Shows the orbits of individual stars near the Galactic Center as measured with high resolution infrared observations. + The movie runs [rom 1992. the initial date of the observations. to the present. aud is extrapolated a few vears into the future.," The movie runs from 1992, the initial date of the observations, to the present, and is extrapolated a few years into the future." + Time is shown at top left., Time is shown at top left. + The star whose track (vith error bars) traces a complete ellipse was fitted by Sehodeel et al. (, The star whose track (with error bars) traces a complete ellipse was fitted by Schöddel et al. ( +2002) to a hiehly elliptical Keplerian orbit.,2002) to a highly elliptical Keplerian orbit. + From the fit thev calculated the mass of the supermassive DII to be 3.7d:1.5xLO°AL..," From the fit they calculated the mass of the supermassive BH to be $3.7 \pm 1.5 \times 10^6 +M_\odot$." + The inferred position ol the DII is shown bv the red cross. (, The inferred position of the BH is shown by the red cross. ( +Movie courtesy Reinharcl Genzel) Movie2.,Movie courtesy Reinhard Genzel) Movie2.— + Shows the orbits of individual stars near (he Galactic Center as measured with high resolution infrared observations., Shows the orbits of individual stars near the Galactic Center as measured with high resolution infrared observations. + The movie runs [rom 1995. the initial date of the observations. to the present. aud is extrapolated a few vears into the future.," The movie runs from 1995, the initial date of the observations, to the present, and is extrapolated a few years into the future." + Time is shown at top left., Time is shown at top left. + By combining the stellar positions with Doppler radial velocity nmeasurenienis and fitting to Ixepleriain orbits. Ghez et al. (," By combining the stellar positions with Doppler radial velocity measurements and fitting to Keplerian orbits, Ghez et al. (" +2004) refined the mass estimate of the supermassive DII to 3.720.2x109...,2004) refined the mass estimate of the supermassive BH to $3.7 \pm 0.2 \times 10^6 M_\odot$. + The inferred position of the BI is shown by the stationary. * at the center. (, The inferred position of the BH is shown by the stationary * at the center. ( +Movie courtesy Andrea Ghez) Movie3.|,Movie courtesy Andrea Ghez) Movie3.— + Shows the simulated image of a turbulent accretion disk heated by (he dissipation of magnetic fields., Shows the simulated image of a turbulent accretion disk heated by the dissipation of magnetic fields. + The dark region al the center is the inner boundary of the simulation. which is at a radius of 27254.," The dark region at the center is the inner boundary of the simulation, which is at a radius of $2R_S$." + During the sequence the view changes from almost [ace-on (7= 1°) to nearly edge-on (/= 307). as indicated by the bar on the right (0° is at the bottom of the bar and 90° at the top).," During the sequence the view changes from almost face-on $i=1\degr$ ) to nearly edge-on $i=80\degr$ ), as indicated by the bar on the right $0\degr$ is at the bottom of the bar and $90\degr$ at the top)." + As the inclination increases. note how the emission becomes enhanced to the left of the DII because of Doppler boost.," As the inclination increases, note how the emission becomes enhanced to the left of the BH because of Doppler boost." + Also. even though the disk is perfectly flat. it appears to be warped upward behind the DII.," Also, even though the disk is perfectly flat, it appears to be warped upward behind the BH." + This is because of the deflection of light ravs by the gravity of the BIL. (, This is because of the deflection of light rays by the gravity of the BH. ( +Based on Armitage Bevnolds 2003: movie courtesy the authors) Movied.,Based on Armitage Reynolds 2003; movie courtesy the authors) Movie4.— + Shows the variation of the fluorescent iron line in (he simulation shown in Movie3., Shows the variation of the fluorescent iron line in the simulation shown in Movie3. + The inclination is held fixecl at 80° and the line emissivity is taken to be proportional to the local energy generation rate., The inclination is held fixed at $80\degr$ and the line emissivity is taken to be proportional to the local energy generation rate. + Note that the line profile. shown at bottom right. varies rapidly as a result of turbulent [Iuctuations in the disk.," Note that the line profile, shown at bottom right, varies rapidly as a result of turbulent fluctuations in the disk." + The line extends from about 4 keV on the left to about 7 keV on the right. with a peak at around 6 keV. (Basecl on Armitage Reynolds 2003: movie courtesy the authors) Movieb.—," The line extends from about 4 keV on the left to about 7 keV on the right, with a peak at around 6 keV. (Based on Armitage Reynolds 2003; movie courtesy the authors) Movie5.—" + Simulation of a magnetic [hix tube accreting onto a maximally rotating BIL (Semenov et al., Simulation of a magnetic flux tube accreting onto a maximally rotating BH (Semenov et al. + 2004)., 2004). + The light circle at the center represents the event horizon, The light circle at the center represents the event horizon + ⊺∐∪∐∏≻↴∖↴∪∐↸∖↑⋜↧↕∙↕∩≝↭∶↕≧⋜↧↸⊳↕⊔⋜⋯∐↸∖↑⋜↧∙⊇∩∩∩∶≋∐∐⋅↕↸∖∙↖⇁↸∖↑ 1: voa ↸⊳≺≻⋯↻∏↑⋜↧⊓∪∐⋜↧∐⋅↖↽∐∐∖⊼↻↸∖∐↴∖↴↕↖↽↸∖∙⋜↧↴∖↴↕↑↥⋅↸∖≺∣∏∐⋅↸∖↴∖↴∐∪∐↑↑∐↕∶↴⋁ but only au estimate of the size of the uuiforii deusitv inner region: if is Iuscusitive to geeoinetryv or amount of magnetic support: and the quantitv it estimates is a vorv nauportaut one. as the central density sets the evolutionary stage of the core within the coutext of any specific dynamical model. aud so it ix necessary for further core modeling. iucludiug velocity. magnetic field. aud chemistrv.," but only an estimate of the size of the uniform density inner region; it is insensitive to geometry or amount of magnetic support; and the quantity it estimates is a very important one, as the central density sets the evolutionary stage of the core within the context of any specific dynamical model, and so it is necessary for further core modeling, including velocity, magnetic field, and chemistry." +" The reason for the robustuess of this technique is the siuplicditv of the physics behiud the formation of the nuiform. density inner region: the fat immer region of the density profile corresponc"" to the scale within which columthermal pressure CLASCS suy inliomogeneities, the thermal lenethscale."," The reason for the robustness of this technique is the simplicity of the physics behind the formation of the uniform density inner region: the flat inner region of the column density profile corresponds to the scale within which thermal pressure erases any inhomogeneities, the thermal lengthscale." + It is therefore expected to be present independently of the details of the core dynamics. such as the amount of magnetic support of the core.," It is therefore expected to be present independently of the details of the core dynamics, such as the amount of magnetic support of the core." +" Its exteut is also uot very sensitive to geometry effects or Core orientation.:⋅ as n⋅ the nuifor» density ⋅⋅incr region⋅ pressure Forces. even qm nuagnefie cores, are uci nore pronunent than further out. aud the geometry is closer to spherical than at larger scales."," Its extent is also not very sensitive to geometry effects or core orientation, as in the uniform density inner region pressure forces, even in magnetic cores, are much more prominent than further out, and the geometry is closer to spherical than at larger scales." +" Since the thermal lenethscaleomt dependsfopeuce oulv on the value of the ceutral: voluue density of the core aud the core temperature. which. in prestellar cores, typically varies by only a factor M ~ aa ""lednkaut of the {hem knetsak nme th ""d 9 i Thner Peston h ni Conal ‘in an"" d theNICK n N me contra DW WIT COMPALiAIC 1 OF ~ 2) accuracy,"," Since the thermal lengthscale depends only on the value of the central volume density of the core and the core temperature, which, in prestellar cores, typically varies by only a factor of $\sim 2$, a measurement of the thermal lengthscale, using the size of the flat inner region of the column density profile, can yield an estimate of the central volume density with comparable (factor of $\sim 2$ ) accuracy." + WI1 COHSITYindependent estimates of the aCTOrkinetic temperature (6.8. froma line intensity ratios). the accuracy on the central ↖↽∪↕⋯⊔↸∖≼∐∖∐↴∖↴↕↑⋅↖⇁↸⊳⋜⋯↕∐∏∐⋅∪↖↽↸∖↕," With independent estimates of the kinetic temperature (e.g. from line intensity ratios), the accuracy on the central volume density can improve further." +"↿∐⋅↑∐↸∖↥⋅∙ Iu S77.877ape wesvo d(describe. onr ⋅⋅propost"" dtC.∖ in. d""tail.. we demonstrateits sonerobustuess. and we discuss chainsystematic een the ana M density mc Hus eeW"," In \ref{method} we describe our proposed technique in detail, we demonstrate its robustness, and we discuss systematic uncertainties entering estimates of the central volume density obtained this way." +AY. Η Sii WU APPIN ΟΙΗ netlod and obtain obtiestimates for the central volume clensity of starloss ∖↴⊲⋅∖↴↴⊳⋅∖ Bosone NSaud the prestellar ∖⋅∖↴∖⋠⋅⊳⋅∖core L1511., In \ref{application} we apply our method and obtain estimates for the central volume density of the starless core B68 and the prestellar core L1544. +.-∙ We⇁∖thediscuss our findingsCOYC iu, Wediscuss our findings in \ref{discussion}. . +"[Fe/H]=—1.0 and to ~0.03 for [Fe/H]=—2.0, with colors at lower logg becoming redder (differences in other colors aresmaller).","$\FeoH=-1.0$ and to $\sim0.03$ for $\FeoH=-2.0$, with colors at lower $\log g$ becoming redder (differences in other colors aresmaller)." +" While the new Τομ-]ορg relations provided in Tables 4 and 5 cover the effective temperature range of Teg=3500—4900 KK, it should be taken into account that uncertainties in these relations will likely be larger below Teg 3900KK and above Teg~ 4600KK, where an extrapolation was used to compensate for a lack of stars in certain cluster groups."," While the new $T_{\rm eff}$ $\log g$ relations provided in Tables \ref{tabGGCscales} and \ref{tabGGCscalesfits} cover the effective temperature range of $T_{\rm eff}=3500-4900$ K, it should be taken into account that uncertainties in these relations will likely be larger below $T_{\rm eff}\sim3900$ K and above $T_{\rm eff}\sim4600$ K, where an extrapolation was used to compensate for a lack of stars in certain cluster groups." + It should also be noted that the new T.g—logg relations provided in Tables 2 and 4 are representative for RGB stars., It should also be noted that the new $T_{\rm eff}$ $\log g$ relations provided in Tables \ref{empGGCscales} and \ref{tabGGCscales} are representative for RGB stars. +" Appropriate care should thus be taken if these relations are used at higher (or lower) temperatures, where RGB stars may be mixed up with stars on the horizontal branch (or AGB stars)."," Appropriate care should thus be taken if these relations are used at higher (or lower) temperatures, where RGB stars may be mixed up with stars on the horizontal branch (or AGB stars)." +" The new T.g-logg relations that were discussed in the previous Section allow us to derive a set of new T.g- log g-color relations based on the synthetic photometric colors, and to compare them with various 7.g-color and color-color relations available from the literature."," The new $T_{\rm eff}$ $\log g$ relations that were discussed in the previous Section allow us to derive a set of new $T_{\rm +eff}$ $\log g$ –color relations based on the synthetic photometric colors, and to compare them with various $T_{\rm +eff}$ –color and color–color relations available from the literature." + This comparison is done separately for [M/H]=—1.0and —2.0., This comparison is done separately for $\MoH=-1.0$and $-2.0$. + Below we focus on the details of these steps., Below we focus on the details of these steps. + The new T.g-log g-color relations were constructed usingthe T.g-logg relations derived in Sect., The new $T_{\rm eff}$ $\log g$ –color relations were constructed usingthe $T_{\rm eff}$ $\log g$ relations derived in Sect. +" 4.2 and broad-band photometric colors calculated with thePHOENIX,MARCS, and stellar model atmospheres (Sect. 2))."," \ref{TGrelations} and broad-band photometric colors calculated with the, and stellar model atmospheres (Sect. \ref{synthcolors}) )." + These relations are provided in Table 6 (scales employing and colors) and Table 7 (relation based on colors)., These relations are provided in Table \ref{TGPhoenixMarcs} (scales employing and colors) and Table \ref{TGAtlas} (relation based on colors). +" They are based on the new empirical Te —logg scales corresponding to the metallicity scale of KI03 (Sect. 4.2)),"," They are based on the new empirical $T_{\rm eff}$ $\log g$ scales corresponding to the metallicity scale of KI03 (Sect. \ref{TGrelations}) )," +" and are delivered at four metallicities, [M/H]=—0.5, —1.0, —1.5, and —2.0."," and are delivered at four metallicities, $\MoH=-0.5$, $-1.0$, $-1.5$, and $-2.0$." +" Note that the limiting gravities in and grids of synthetic colors are logg=0.5 and logg=0.0, respectively; to extend the coverage in logg,and colors were linearly extrapolated to ~0.2 ddex below these values."," Note that the limiting gravities in and grids of synthetic colors are $\log g = 0.5$ and $\log g = 0.0$, respectively; to extend the coverage in $\log g$,and colors were linearly extrapolated to $\sim0.2$ dex below these values." + Uncertainties in the new T.g—log g-color relations are governed by the uncertainties in the empirical Tyg—log g scales that were obtained as best fits to the observed T.g— logg sequences of late-type giants in Galactic globular clusters (Sect. 4.2)).," Uncertainties in the new $T_{\rm eff}$ $\log g$ –color relations are governed by the uncertainties in the empirical $T_{\rm +eff}$ $\log g$ scales that were obtained as best fits to the observed $T_{\rm eff}$ $\log g$ sequences of late-type giants in Galactic globular clusters (Sect. \ref{TGrelations}) )." +" The typical RMS residual of the fitting procedure is ~ 0.15ddex in logg, or ~100KK in Tog."," The typical RMS residual of the fitting procedure is $\simeq0.15$ dex in $\log g$, or $\simeq100$ K in $T_{\rm eff}$." +" At Toe= 4400KK, logg= 1.5, and [M/H]= —1.0, the uncertainty ΔΤος=100 KK will be equivalent to changes in photometric colors A(B—V)~A(V 0.05, A(V—K)50.13, and A(J—K)~0.04 (correspondingly, to 0.07, 0.05, 0.14, and mmag, at logg—1.0 and [M/H]— —2.0)."," At $T_{\rm +eff}=4400$ K, $\log g=1.5$ , and $\MoH=-1.0$ , the uncertainty $\Delta T_{\rm eff}=100$ K will be equivalent to changes in photometric colors $\Delta(B-V)\simeq\Delta(V-I)\simeq0.05$ , $\Delta(V-K)\simeq0.13$, and $\Delta(J-K)\simeq0.04$ (correspondingly, to 0.07, 0.05, 0.14, and mag, at $\log +g=1.0$ and $\MoH=-2.0$ )." + The effect on photometric colors will increase slightly with decreasing gravity., The effect on photometric colors will increase slightly with decreasing gravity. +" Note, however, that photometric colors are less sensitive to uncertainties in logg: at Teg= 4400KK, logg=1.5, and [M/H]=—1.0, Alogg=0.15 will correspond to A(B—V)c 0.02, with differences in other colors at the level of mmag or lower."," Note, however, that photometric colors are less sensitive to uncertainties in $\log +g$: at $T_{\rm eff}=4400$ K, $\log g=1.5$, and $\MoH=-1.0$, $\Delta \log g =0.15$ will correspond to $\Delta(B-V)\simeq0.02$ , with differences in other colors at the level of mag or lower." +" While we will further quote +100 KK asa representative uncertainty of the new T.g— log g-color relations,it is rather obvious that this may only represent a lower limit on the true uncertainties, which include varioussystematical effects inherent in the spectroscopic derivations of Te, logg, and [Fe/H], limitations of the current stellar atmosphere models, and so forth."," While we will further quote $\pm100$ K asa representative uncertainty of the new $T_{\rm eff}$ $\log g$ –color relations,it is rather obvious that this may only represent a lower limit on the true uncertainties, which include varioussystematical effects inherent in the spectroscopic derivations of $T_{\rm eff}$ , $\log +g$ , and $\FeoH$ , limitations of the current stellar atmosphere models, and so forth." +to à 5-20 Mos. Pow ~1.8 day configuration.,"to a 5+20 $_{\odot}$ , $_{\rm orb} \sim$ 1.8 day configuration." + Comparable simulations with improved input physics are currently being undertaken by Schneider et al. (, Comparable simulations with improved input physics are currently being undertaken by Schneider et al. ( +in prep.).,"in prep.)," +" and although a higher mass progenitor would be required (e.g. 40M,. for a current mass of «10 M4) such à common envelope pathway appears promising for the production of W239.", and although a higher mass progenitor would be required (e.g. $_{\odot}$ for a current mass of $\sim$ 10 $_{\odot}$ ) such a common envelope pathway appears promising for the production of W239. + As an alternative to binary evolution driven by mass transfer we also highlight theCase Mixing). evolutionary, As an alternative to binary evolution driven by mass transfer we also highlight theCase M(ixing) evolutionary +bbv using equation (5)).,by using equation \ref{eq:polyfit}) ). + The estimated. column. density distribution is shown in Figure S., The estimated column density distribution is shown in Figure \ref{fig:column_dist}. + An average and a standard deviation of the distribution are. found.to. be (logINgTonNy)=(22.8.0.5).," An average and a standard deviation of the distribution are foundto be $(\overline{\log N_H}, \sigma_{\log N_H}) += (22.8, 0.5)$." + Marconietal.(2009). has suggested that radiation pressure does play an important role in DLIt gas dvnamües if column densities of BLR clouds have intrinsic dispersion such as (logNa.io.Xy)=(23.0. 0.5).," \citet{mar09} has suggested that radiation pressure does play an important role in BLR gas dynamics if column densities of BLR clouds have intrinsic dispersion such as $(\overline{\log N_H}, +\sigma_{\log N_H}) = (23.0, 0.5)$ ." + Our results support that the assumption adopted in Marconietal.(2009) is appropriate and that the radiation pressure plays an important role in BLIt clouds., Our results support that the assumption adopted in \citet{mar09} is appropriate and that the radiation pressure plays an important role in BLR clouds. + Figure 9 shows the relation between aand the Eddington ratio., Figure \ref{fig:fefe_edd} shows the relation between and the Eddington ratio. + X positive correlation is seen., A positive correlation is seen. + Linear regression analysis. using an LOL procedure “PITENYpro (c£. Pressetal. 19922).," Linear regression analysis, using an IDL procedure “FITEXY.pro” (cf. \citealt{pre}) )," + gives the relation as: The Spearman's rank correlation for assessing the nonlinear correlation is rs=0.58., gives the relation as: The Spearman's rank correlation for assessing the nonlinear correlation is $r_S=0.58$. + This means— the probability of the null hypothesis that there is no correlation is less than 1005., This means the probability of the null hypothesis that there is no correlation is less than $10^{-13}$. + μιας the correlation between aand the EddingtonU V)ratio is real., Thus the correlation between and the Eddington ratio is real. + This implies that the column density increases with the Exdcdington ratio. because lincreases with the columnV) density.," This implies that the column density increases with the Eddington ratio, because increases with the column density." + As was recently suggested. by Dongetal.(2009).. under the condition where the BLR clouds are subject to the radiation pressure. low-columm-cdensity clouds would be blown away by relatively large radiation pressure at large hoafleaeo. so that only high-column-density clouds would be able to be gravitationally bound.," As was recently suggested by \citet{dong}, under the condition where the BLR clouds are subject to the radiation pressure, low-column-density clouds would be blown away by relatively large radiation pressure at large $L_{bol}/L_{Edd,0}$, so that only high-column-density clouds would be able to be gravitationally bound." + Figure 9 is a supportive evidence for their suggestion., Figure \ref{fig:fefe_edd} is a supportive evidence for their suggestion. + Figure 10. plots our samples on NyLoafLiao plane.," Figure \ref{fig:edd_permit} plots our samples on $N_H - +L_{bol}/L_{Edd,0}$ plane." + Each line represents Lior=Levant. so that the lower region of the line corresponds super-Eddington area.," Each line represents $L_{bol}=L_{Edd,rad}$, so that the lower region of the line corresponds super-Eddington area." + Ll we adopt ionizing photon fraction @=0.6 (i.e... thick solid line in Figure 10)). which is an average value for AGINs calculated by Marconietal. (2008)... almost all of our samples become super-IExdcdington.," If we adopt ionizing photon fraction $a=0.6$ (i.e., thick solid line in Figure \ref{fig:edd_permit}) ), which is an average value for AGNs calculated by \citet{mar08}, , almost all of our samples become super-Eddington." + This result can be interpreted in two wavs: (i) the conversion from, This result can be interpreted in two ways: (i) the conversion from +flat fielding) with4/47 packagecedred?.,flat fielding) with package. +". ""Ehe subsequent echelle data reduction was also carried out withARAL and ds package.", The subsequent echelle data reduction was also carried out with and its package. + The RVs from the ‘ThaAr wavelength calibrated spectra were calculated with our own implementation of the two climensional cross-correlation technique (Zucker&Alazeh1994.TODCOLn)., The RVs from the ThAr wavelength calibrated spectra were calculated with our own implementation of the two dimensional cross-correlation technique \citep[TODCOR]{zuc94}. + As templates. we used. svnthetic spectra computed. with the ATLASS and ATLASI2 codes (IlXurucz1992).," As templates, we used synthetic spectra computed with the ATLAS9 and ATLAS12 codes \citep{kur92}." +. With the exception of two svstems. AL Phe CASAS. J010034). and UN Alen CASAS. 053003). the RVs from the iocine cell calibrated spectra were also computed with TODCOR.," With the exception of two systems, AI Phe (ASAS J010934) and UX Men (ASAS J053003), the RVs from the iodine cell calibrated spectra were also computed with TODCOR." + The details of such a procedure are described by Ixonacki(2005).., The details of such a procedure are described by \cite{Konacki:05::}. + For the two brightest binaries in this sample. UX Men (ASAS 053003) and AL Phe CASAS J010934) we were able to collect eight spectra taken with the iodine cell each.," For the two brightest binaries in this sample, UX Men (ASAS J053003) and AI Phe (ASAS J010934) we were able to collect eight spectra taken with the iodine cell each." + This is about the smallest number of spectra still sullicicnt to carry out a tomographic disentangling to obtain observed component spectra of a binary., This is about the smallest number of spectra still sufficient to carry out a tomographic disentangling to obtain observed component spectra of a binary. + Qur disentangling procedure is described bv Konacki(2009).., Our disentangling procedure is described by \cite{Konacki:09::}. + Lt essentially allows one to derive the component spectra [rom the observed. composite spectra and then use them to compute the RVs., It essentially allows one to derive the component spectra from the observed composite spectra and then use them to compute the RVs. + As we have shown this approach is capable of providing RVs of the components of double-line binary stars with a precision reaching 5 ms1 (IKonacki2009)..," As we have shown this approach is capable of providing RVs of the components of double-line binary stars with a precision reaching 5 $\,$ $^{-1}$ \citep{Konacki:09::}." + While the spectra of UN Alen and Al Phe from the ANT/UCLES are characterized by a much lower οἱ of 740-100 compared to those used in Wonacki(2000).. the disentangling still can be carried out.," While the spectra of UX Men and AI Phe from the AAT/UCLES are characterized by a much lower $SNR$ of $\sim$ 40-100 compared to those used in \cite{Konacki:09::}, the disentangling still can be carried out." + This procedure has resulted in a higher IV. precision compared to the standard. Thr approach or the iocline cell based approach combined with PODCOIt for these two argets., This procedure has resulted in a higher RV precision compared to the standard ThAr approach or the iodine cell based approach combined with TODCOR for these two targets. + In particular. the best-fitting RY solution for Al Phe is characterized by an rms of 62 and 24 ms for the »wimary and secondary respectively.," In particular, the best-fitting RV solution for AI Phe is characterized by an $rms$ of 62 and 24 $\,$ $^{-1}$ for the primary and secondary respectively." + Phe UN Alen rns is not nearly as good (210 and 270 ms 1) but this is due to the very wide spectral lines of its components.," The UX Men $rms$ is not nearly as good (210 and 270 $\,$ $^{-1}$ ) but this is due to the very wide spectral lines of its components." + “Phere is no doubt hat the RV. precision for AL Phe would be much higher if ugher SNA spectra were available., There is no doubt that the RV precision for AI Phe would be much higher if higher $SNR$ spectra were available. + It should be noted here hat the IV. precision from the iodine cell spectra used in his paper is not representative for this technique as most of he time we were dealing with an SNA [ar too low for what is required. to obtain a high RY precision., It should be noted here that the RV precision from the iodine cell spectra used in this paper is not representative for this technique as most of the time we were dealing with an $SNR$ far too low for what is required to obtain a high RV precision. + Still the results are quite satisfactory precision-wisc., Still the results are quite satisfactory precision-wise. + For seven svstems the iodine cell based. solution was substantially better than the PhaAr based one., For seven systems the iodine cell based solution was substantially better than the ThAr based one. + For four other targets we used RVs based on both methods of wavelength calibration since the number of iodine cell based spectra was too small or the resulting rms was comparable., For four other targets we used RVs based on both methods of wavelength calibration since the number of iodine cell based spectra was too small or the resulting $rms$ was comparable. + The seven remaining targets. (including the two observed. at S.AO) have their solutions based on the Thr calibrated. spectra only., The seven remaining targets (including the two observed at SAAO) have their solutions based on the ThAr calibrated spectra only. + The RV measurements together with their errors and O—€ are collected. in Table Al in the Appendix. ??.., The RV measurements together with their errors and $O-C$ are collected in Table \ref{tab_allrv} in the Appendix \ref{sec_allrv}. + As it turned out the formal errors συ computed from the scatter between the echelle orders used in the analysis were somewhat underestimated., As it turned out the formal errors $\sigma_0$ computed from the scatter between the echelle orders used in the analysis were somewhat underestimated. + Hence to obtain 47=1 for our, Hence to obtain $\chi^2 \simeq 1$ for our +for Epstein drag and for Stokes drag.,for Epstein drag and for Stokes drag. + Regime 2.1 progresses (o regime 2.2 when (he grains grow large enough that (he turbulence can no longer prevent settling al which point (34)) applies., Regime $2.1$ progresses to regime $2.2$ when the grains grow large enough that the turbulence can no longer prevent settling at which point \ref{hd2}) ) applies. + This occurs Lor Accorclingly Regime 2.1 ends for with Epstein drag ancl for with Stokes drag., This occurs for Accordingly Regime $2.1$ ends for with Epstein drag and for with Stokes drag. + Regime 2.2 occurs for dust grains large enough {ο begin settling out of the gas disc and also have their collisional velocity determined by coupling to non-mininnm scale turbulence., Regime $2.2$ occurs for dust grains large enough to begin settling out of the gas disc and also have their collisional velocity determined by coupling to non-minimum scale turbulence. + llere we use 6.5 from (24)) and £145 from (34)) while approximating 1—€Ὁτ=I.," Here we use $v_{c, 2}$ from \ref{vc2}) ) and $H_{d,2}$ from \ref{hd2}) ) while approximating $1-e^{-\frac{1}{\Omega \tau}} = 1$." + Solving (3)) we lind the time scale to grow from some o; to oy in (his regime to be with Epstein drag aud, Solving \ref{main}) ) we find the time scale to grow from some $\phi_i$ to $\phi_f$ in this regime to be with Epstein drag and +The residuals could be diminished partially by observing several calibrators.,The residuals could be diminished partially by observing several calibrators. + But not to a full extent: the noise would average only on the calibrator dataset., But not to a full extent: the noise would average only on the calibrator dataset. + 992945 HHya) is a K1.5V star located at 22 Ρο., 92945 Hya) is a K1.5V star located at 22 pc. + (?) estimated an age of 100 Myr using lines., \citep{2004ApJ...614L.125S} estimated an age of 100 Myr using lines. +" If we suppose 992945 to be part of the AB Dor group (?),, this age agrees with the estimation of z 50-120 Myr for the whole group (??).."," If we suppose 92945 to be part of the AB Dor group \citep{2006ApJ...643.1160L}, this age agrees with the estimation of $\approx 50$ -120 Myr for the whole group \citep{2004ApJ...613L..65Z,2005ApJ...628L..69L}." + Observations with the Advanced Camera for Surveys (ACS) on the have revealed it is surrounded by a disk ~30? from on with a 77-thick bright ring at 57 ffrom the star., Observations with the Advanced Camera for Surveys (ACS) on the have revealed it is surrounded by a disk $\sim30\degr$ from edge-on with a 7-thick bright ring at 57 from the star. + The disk has a diffuse component detected between 55 and 170 citep20071yot.confE..46G.., The disk has a diffuse component detected between 55 and 170 \\citep{2007lyot.confE..46G}. + The inner gap may well be the result of resonances with a stellar or planetary companion (?).. 11, The inner gap may well be the result of resonances with a stellar or planetary companion \citep{1994ApJ...421..651A}. +41569 is a 5+ 3-Myr-old (?) pre-main sequence B9.5V star in a triple system., 141569 is a $5\pm3$ -Myr-old \citep{2004A&A...419..301M} pre-main sequence B9.5V star in a triple system. + It is located 99-pc away, It is located 99-pc away +"observed flux (fops) and the estimated continuum flux (feont) at 9.85um,, and we obtained the silicate strength S,i; as For sourcesfco with a silicate absorption feature, S, can be interpreted as the negative of the apparent silicate optical depth (To; ym).","observed flux $f_{\rm obs}$ ) and the estimated continuum flux $f_{\rm cont}$ ) at 9.85, and we obtained the silicate strength $S_{\rm sil}$ as For sources with a silicate absorption feature, $S_{\rm sil}$ can be interpreted as the negative of the apparent silicate optical depth $\tau_{9.7~\mu \rm m}$ )." + The of Fig., The of Fig. + 8 shows the IRS/SL map (rebinned to the 2x2 aperture of the IRS/SH map) of the apparent silicate strength at 9.85µπι., \ref{fig:Tau-sil} shows the IRS/SL map (rebinned to the $\times$ 2 aperture of the IRS/SH map) of the apparent silicate strength at 9.85. +". The silicate-based extinction Ay(9.85um), shown in the of Fig. 8,,"," The silicate-based extinction $A_{\rm V}(9.85~\mu \rm m)$, shown in the of Fig. \ref{fig:Tau-sil}," +" was estimated from the average visual extinction to silicate optical depth ratio Αν/τ(9.7um)=18, which is appropriate for the local ISM (??).."," was estimated from the average visual extinction to silicate optical depth ratio $A_{\rm V}/\tau(9.7~\mu \rm m)=18$, which is appropriate for the local ISM \citep{roche84, rieke85}." + The spatial distribution of the silicate- extinction is similar to that of the stellar light-based extinction Ay(H—K) estimated in Sec.??.., The spatial distribution of the silicate-based extinction is similar to that of the stellar light-based extinction $A_{\rm V}(H-K)$ estimated in \ref{sec:Av-HK}. . +" The peak extinction is found at the same relative position (AR.A.«2"", 1""), at about 2.3"" (one pixel) northeast of the HO mega maser (?).."," The peak extinction is found at the same relative position $\Delta~\rm R.A.\approx2''$, $\Delta~\rm Dec\approx1''$ ), at about $2.3''$ (one pixel) northeast of the $_2$ O mega maser \citep{greenhill97}." + It is known that extinction estimated from optical or near-infrared observations generally underestimates the actual extinction if the environment probed is optically thick at the emission lines observed., It is known that extinction estimated from optical or near-infrared observations generally underestimates the actual extinction if the environment probed is optically thick at the emission lines observed. +" If most of the emitting region is obscured, as in the case of the nucleus of NGC 4945, the visual extinction estimate is representative of the surface of the obscured region and not the region itself."," If most of the emitting region is obscured, as in the case of the nucleus of NGC 4945, the visual extinction estimate is representative of the surface of the obscured region and not the region itself." +" This effect is reflected in the different extinctions derived from the H—K optical image and the silicate-based estimate derived from our mid-IR observations, where the peak extinction is about 7 times stronger than that of the stellar light-based estimate."," This effect is reflected in the different extinctions derived from the $H-K$ optical image and the silicate-based estimate derived from our mid-IR observations, where the peak extinction is about 7 times stronger than that of the stellar light-based estimate." + Our Av(9.85um) is also a factor ~1.7 higher than the extinction (Ay=36718 mag) previously inferred from ISO observations of the 18.7/33.5 lline ratio (?).., Our $A_{\rm V}(9.85~\mu \rm m)$ is also a factor $\sim$ 1.7 higher than the extinction $A_{\rm V}=36^{+18}_{-11}~\rm mag$ ) previously inferred from ISO observations of the 18.7/33.5 line ratio \citep{spoon00}. +" This difference can be explained by the peaked nature of the silicate-based extinction map and because of the larger apertures (>14""x 20"") of the ISO observations, which averages out the extinction to the lower value."," This difference can be explained by the peaked nature of the silicate-based extinction map and because of the larger apertures $\ge14''\times20''$ ) of the ISO observations, which averages out the extinction to the lower value." +" Figure 9 shows the contour lines of the IRS/SH surface brightness map of the molecular hydrogen line HyS(1) 17.0 left)), Hz S(2) 12.3 right)), the IRS/SL map of the Hz S(3) 9.7 lline left)), and the 12.81"," Figure \ref{fig:AvSil-H2} shows the contour lines of the IRS/SH surface brightness map of the molecular hydrogen line $_2$S(1) 17.0 ), $_2$ S(2) 12.3 ), the IRS/SL map of the $_2$ S(3) 9.7 line ), and the 12.81" +"and the ratio of injected particle numbers at p=1GeVcl is Kop=A,JAyM1«107.",and the ratio of injected particle numbers at $p = 1~{\rm GeV}~c^{-1}$ is $K_{ep} \equiv A_e/A_p = 1.1 \times 10^{-4}$. + In the case of the leptonic scenario. the electron index is constrained by the LAT spectrum.," In the case of the leptonic scenario, the electron index is constrained by the LAT spectrum." +" Considering inverse Compton seattering in the Thomson regime. the photon index and electron index are related as s,=20-1."," Considering inverse Compton scattering in the Thomson regime, the photon index and electron index are related as $s_e = 2 \Gamma - 1$." +" Therefore. the observed photon index [=1.85 leads to a soft electron spectrum with s,22.7."," Therefore, the observed photon index $\Gamma = 1.85$ leads to a soft electron spectrum with $s_e = 2.7$." + With this index. however. the synchrotron spectrum cannot fit the radio and X-ray data well.," With this index, however, the synchrotron spectrum cannot fit the radio and X-ray data well." + We then adopted a harder electron spectrum within the range allowed by the statistical and systematic errors in the LAT spectral points., We then adopted a harder electron spectrum within the range allowed by the statistical and systematic errors in the LAT spectral points. +" The model curves in Figure 3 (b) are the leptonic models when the electron index is s,=2.15 with a high-energy cutoff at pa,=25TeV col. The magnetic field strength can be determined as B=12μα so that the synchrotron to inverse Compton flux ratio matches the data.", The model curves in Figure \ref{fig:modeling} (b) are the leptonic models when the electron index is $s_e = 2.15$ with a high-energy cutoff at $p_{0e} = 25~{\rm TeV}~c^{-1}$ The magnetic field strength can be determined as $B = 12~\mu{\rm G}$ so that the synchrotron to inverse Compton flux ratio matches the data. + Electron bremsstrahlung emission is calculated to be ΕΙΝ/dE~107ergems! at 10 GeV. Both models face difficulty., Electron bremsstrahlung emission is calculated to be $E^2 dN/dE \sim 10^{-13}~{\rm erg}~{\rm cm}^{-2}~{\rm s}^{-1}$ at 10 GeV. Both models face difficulty. + The leptonic model requires a weak magnetic field of order 10j/G. which may be contradicted by observations.," The leptonic model requires a weak magnetic field of order $10~\mu{\rm G}$, which may be contradicted by observations." + The width of the filaments in the shell observed in gives a magnetic field estimate of ~100μμ (Bambaetal.2005:Berezhko 2009).," The width of the filaments in the shell observed in gives a magnetic field estimate of $\sim 100~\mu{\rm G}$ \citep{bamba05, berezhko09}." +. These results are interpreted as evidence for magnetic field amplification by streaming of accelerated cosmic rays., These results are interpreted as evidence for magnetic field amplification by streaming of accelerated cosmic rays. + However. It is possible to reconcile a high magnetic field with the leptonic model if GeV gamma rays are radiated not only from the filamentary structures seen by. but also from other regions in the SNR where the magnetic field may be weaker.," However, It is possible to reconcile a high magnetic field with the leptonic model if GeV gamma rays are radiated not only from the filamentary structures seen by, but also from other regions in the SNR where the magnetic field may be weaker." + The hadronic model requires an unrealistically large energy of protons if the gas density is as small as <0.01em”., The hadronic model requires an unrealistically large energy of protons if the gas density is as small as $\lesssim 0.01~{\rm cm}^{-3}$. + Thermal X-ray emission works as a probe to estimate the gas density., Thermal X-ray emission works as a probe to estimate the gas density. + However. no clear detection is reported so far mainly because it is difficult to separate emission from RX J0852.0—4622 and the Vela SNR. which overlaps and emits strong thermal emission in the soft X-ray band.," However, no clear detection is reported so far mainly because it is difficult to separate emission from RX $-$ 4622 and the Vela SNR, which overlaps and emits strong thermal emission in the soft X-ray band." + Using data. Slaneetal.(2001) placed an upper limit of n3.3«107(4/750pey7.£77em? where d and f are the distance to RX J0852.0—4622 and the filling factor of a sphere taken às the emitting volume in the region used for their spectral analysis. respectively.," Using data, \cite{slane01} placed an upper limit of $n < 3.3 \times 10^{-2}~(d/750~{\rm pc})^{-1/2}~f^{-1/2}~{\rm cm}^{-3}$, where $d$ and $f$ are the distance to RX $-$ 4622 and the filling factor of a sphere taken as the emitting volume in the region used for their spectral analysis, respectively." +" In the case of f20.4. which corresponds to rather thick shell width. the upper limit becomes problematic since W, reaches >10°! erg. which means that almost all the kinetic energy released by the supernova explosion must go to accelerated protons."," In the case of $f \gtrsim 0.4$, which corresponds to rather thick shell width, the upper limit becomes problematic since $W_p$ reaches $\gtrsim 10^{51}$ erg, which means that almost all the kinetic energy released by the supernova explosion must go to accelerated protons." +" The proton index s,=1.8 for the hadronic scenario is somewhat harder than s,=2.0 expected in the test-particle regime of diffusive shock acceleration with a strong shock.", The proton index $s_p = 1.8$ for the hadronic scenario is somewhat harder than $s_p = 2.0$ expected in the test-particle regime of diffusive shock acceleration with a strong shock. + Therefore. it is of interest to compare the observed spectrum with models based on diffusive shock acceleration in the nonlinear regime.," Therefore, it is of interest to compare the observed spectrum with models based on diffusive shock acceleration in the nonlinear regime." + Berezhkoetal.(2009)— has calculated the broadband SED of RX J0852.0—4622 with their time-dependent. nonlinear kinetic theory for cosmic-ray acceleration.," \cite{berezhko09} has calculated the broadband SED of RX $-$ 4622 with their time-dependent, nonlinear kinetic theory for cosmic-ray acceleration." + Their model curve is plotted over the GeV-to-TeV data in Figure 2.., Their model curve is plotted over the GeV-to-TeV data in Figure \ref{fig:spec}. + The LAT spectrum appears harder than the model by Berezhkoetal.(2009).. but the discrepancy tis not large compared to the systematic errors.," The LAT spectrum appears harder than the model by \cite{berezhko09}, but the discrepancy is not large compared to the systematic errors." + It is of interest to compare the gamma-ray spectrum of RX JO852.0-4622 with that of the similar object. RX J1713.7-3946 (Abdoetal.2011).," It is of interest to compare the gamma-ray spectrum of RX $-$ 4622 with that of the similar object, RX $-$ 3946 \citep{fermi1713}." + The two SNRs have roughly the same age. size. and similar radio. X-ray and TeV gamma-ray spectra.," The two SNRs have roughly the same age, size, and similar radio, X-ray and TeV gamma-ray spectra." + The LAT spectrum of RX J1713.7-3946 has a photon index of T=1.540.1. which is more preferable for leptonic models than hadronie models (Abdoetal.2011)..," The LAT spectrum of RX $-$ 3946 has a photon index of $\Gamma =1.5 \pm 0.1$, which is more preferable for leptonic models than hadronic models \citep{fermi1713}." + In the case of RX JO852.0-4622. a leptonic origin. is not favored but cannot be ruled out.," In the case of RX $-$ 4622, a leptonic origin is not favored but cannot be ruled out." + The difference between the two similar objects could be an important point to be further explored to probe gamma-ray production in SNRs., The difference between the two similar objects could be an important point to be further explored to probe gamma-ray production in SNRs. + Further studies of LAT data in the future. particularly better modeling of the Galactic diffuse emission. will reduce the uncertainties and will allow us to probe particle acceleration in the SNR in greater detail.," Further studies of LAT data in the future, particularly better modeling of the Galactic diffuse emission, will reduce the uncertainties and will allow us to probe particle acceleration in the SNR in greater detail." + The LAT Collaboration acknowledges support from a number of agencies and institutes for both development and the operation of the LAT as well as scientific data analysis., The LAT Collaboration acknowledges support from a number of agencies and institutes for both development and the operation of the LAT as well as scientific data analysis. + These include NASA and DOE in the United States. CEA/Irfu and IN2P3/CNRS in France. ASI and INFN in Italy. MEXT. KEK. and JAXA in Japan. and the K.A. Wallenberg Foundation. the Swedish Research Council and the National Space Board in Sweden.," These include NASA and DOE in the United States, CEA/Irfu and IN2P3/CNRS in France, ASI and INFN in Italy, MEXT, KEK, and JAXA in Japan, and the K.A. Wallenberg Foundation, the Swedish Research Council and the National Space Board in Sweden." + Additional support from INAF in Italy and CNES in France for science analysis during the operations phase is also gratefully acknowledged., Additional support from INAF in Italy and CNES in France for science analysis during the operations phase is also gratefully acknowledged. +clubedded system viewed along a different sightliue.,embedded system viewed along a different sightline. + The effects are more severe at shorter wavelengths., The effects are more severe at shorter wavelengths. + Although the case of YSOs is very different to the ealactic environment studied here it should be noted hat the effects of à ο) galactic ISM. containiug optically thick clouds may cause viewing angle aud inclination dependent effects on the MIR. SEDs that are rot accounted for in our steoth. axisvuuuetric models.," Although the case of YSOs is very different to the galactic environment studied here it should be noted that the effects of a clumpy galactic ISM containing optically thick clouds may cause viewing angle and inclination dependent effects on the MIR SEDs that are not accounted for in our smooth, axisymmetric models." + Ciuveutle we lave no observations of our ealaxy saluple in the sub-nunu waveleneth reguue., Currently we have no observations of our galaxy sample in the sub-mm wavelength regime. + It has been ound that in order to reproduce the sub-nua cussion observed froin some TSB galaxies it becomes necessi o iuclude an additional dust mass (2??)..," It has been found that in order to reproduce the sub-mm emission observed from some HSB galaxies it becomes necessary to include an additional dust mass \citep{popescu_modelling_2000,misiriotis_modeling_2001,popescu_modelling_2011}." + The additional dust mass reveals its prescuce through swb-muu emissiou hat is underestimated by models which can adequately describe the SED to — 1004422. Further evidence for au extended cold dust componcut is also found by? based on attenuation of backeround galaxies in face-on disks., The additional dust mass reveals its presence through sub-mm emission that is underestimated by models which can adequately describe the SED to $\sim 100\mu$ m. Further evidence for an extended cold dust component is also found by \citet{holwerda_opacity_2005} based on attenuation of background galaxies in face-on disks. + τι. observations become available we cannot be sure that our models do not lack cold dust emissiou at longer waveleugtlis., Until observations become available we cannot be sure that our models do not lack cold dust emission at longer wavelengths. + We have utilized woultinwwaveleneth imaging and photometry iu conjunction with sophisticated Monte Carlo radiation trauster codes to investigate the structure of three edge-on. LSB disk galaxies.," We have utilized multi-wavelength imaging and photometry in conjunction with sophisticated Monte Carlo radiation transfer codes to investigate the structure of three edge-on, LSB disk galaxies." + The galaxies lave been chosen to span a range iu central. optical surface brightness and molecular hydrogen masses.," The galaxies have been chosen to span a range in central, optical surface brightness and molecular hydrogen masses." + We have been able to reproduce the elobal. optical appearance of all three galaxies using sooth cuussivity and dust distributions.," We have been able to reproduce the global, optical appearance of all three galaxies using smooth emissivity and dust distributions." + We find that the composition and size distribution of dust erains adopted. which are based ou Milkv Way extinction. provide a good match to the observed properties iu our sample of LSB disk galaxies.," We find that the composition and size distribution of dust grains adopted, which are based on Milky Way extinction, provide a good match to the observed properties in our sample of LSB disk galaxies." + Our models also reproduce thetotal cnussion at 70 and 160422 for all three galaxies., Our models also reproduce thetotal emission at $70$ and $160\mu$ m for all three galaxies. + However. the FIR inorpholoey of our models appears more ceutrally concentrated than the more diffuse distribution sugeested by the data.," However, the FIR morphology of our models appears more centrally concentrated than the more diffuse distribution suggested by the data." + We find that the dust mass appears to be distributed iu au exponential disk with a scale height comparable to or exceeding that of the stellar disk., We find that the dust mass appears to be distributed in an exponential disk with a scale height comparable to or exceeding that of the stellar disk. + This is iu coutrast to the Ποιος for ISB ealaxies where the dust disk is found have a vertical scale height of approximately half the stellar disk (?).., This is in contrast to the findings for HSB galaxies where the dust disk is found have a vertical scale height of approximately half the stellar disk \citep{xilouris_are_1999}. + The comparable scale heights in the dust aud stellar disks is likely associated with the increased stability of the ISM in LSB disks against vertical collapse (2). aud the hin nature of the stellar disks. which suggests minimal lwnamiuical heating.," The comparable scale heights in the dust and stellar disks is likely associated with the increased stability of the ISM in LSB disks against vertical collapse \citep{dalcanton_formation_2004} and the thin nature of the stellar disks, which suggests minimal dynamical heating." + The dust masses aud distributions derived sugeest dust nasses in the range L1162.38&LO°AL.. corvespouding to ace on. V. baud. optical depths between τε=0.031Y. 106.," The dust masses and distributions derived suggest dust masses in the range $1.16-2.38\times 10^{6} M_{\odot}$ corresponding to face on, $V$ band, optical depths between $\tau_{face}=0.034-0.106$ ." + Iu future work we hope to develop our radiatiou rausfer models to include sinall scale uon-axisviuiuietric structures which may shed further light ou the structure of the ISAL and star formation processes in LSB disk galaxies., In future work we hope to develop our radiation transfer models to include small scale non-axisymmetric structures which may shed further light on the structure of the ISM and star formation processes in LSB disk galaxies. + The inclusion of radial variations in the dust disk scale heights. associated with a flaring gas disk. could also prove important.," The inclusion of radial variations in the dust disk scale heights, associated with a flaring gas disk, could also prove important." + Additional sub-uuu observations may also allow us to uncover cold dust that is associated with the extended Τι. as has been found in some TSB ealaxies (?7)..," Additional sub-mm observations may also allow us to uncover cold dust that is associated with the extended , as has been found in some HSB galaxies \citep{popescu_first_2003,hinz_extended_2006}. ." +VLBI data.,VLBI data. + The flux of 2.1 ταν reported by Παν ct àl., The flux of 2.1 mJy reported by Ray et al. + is below the flux level of around 3 Jv which we beleve is resolved out by the VLBI (see Table 2))., is below the flux level of around 3 mJy which we believe is resolved out by the VLBI (see Table \ref{fluxes}) ). + A source size of around 100 mas corresponds to a scale around IATA iu the visibility funetious., A source size of around 100 mas corresponds to a scale around $\lambda$ in the visibility functions. + The shortest VEDI baseline is VLA-Pie Town (PT). with a distance of around LATA.," The shortest VLBI baseline is VLA-Pie Town (PT), with a distance of around $\lambda$." + We see no evidence of resolution within the set of VLA-PT scans. or between these aud the next shortest baseline.," We see no evidence of resolution within the set of VLA-PT scans, or between these and the next shortest baseline." + Nevertheless. our visibility aniplitudes are poorly sampled ou these scales. and we cannot rule out that there is structure there which is nüssed.," Nevertheless, our visibility amplitudes are poorly sampled on these scales, and we cannot rule out that there is structure there which is missed." + It should of course be borne iu iindthat these sources are intrinsically variable., It should of course be borne in mind that these sources are intrinsically variable. + Furthermore. iu the case of blobs of outflowing material. sienificant motion iav well have occurred on time scales of vears between the observations.," Furthermore, in the case of blobs of outflowing material, significant motion may well have occurred on time scales of years between the observations." + circular polarizationn. is ⋅⋅⋅⋅↴an indicator of⋅⋅ mthe presence :ασια] of magnetic fields., The circular polarization is an indicator of the presence of magnetic fields. +" The short timescale variability. frou iours down to secouds. together with⋅ MIN.- brightuess: eniperatures. polnuts to a o""vrosvuchrotron omission : ⊔↸∖↸⊳∐⋜∐"," The short timescale variability, from hours down to seconds, together with MK brightness temperatures, points to a gyrosynchrotron emission mechanism." +∐↴∖↴⋯∙⊺↕∐∖↓∩∩↖⋰⋰⊲∣↥⋅↕∶↴∙⊾∐↑⊣⋜⋯≼↧↻∪↕⋜∐⋅↕∑↸∖≼↧↸∖∐∐↴∖↴↴∖↴↕∪∐⋅ ⋅ ⋅⋅ ⋅↴↴ ⋅ - ⋅↴⇁↴ ⋅⋡⋅ ↸∖↸∖↸⊳⊓⋅∪∐↸⊳↖↽↸⊳↕∪⊓⋅∪∐⋯⋜↧↴∖↴↸∖↥⋅∙⋜∐⋅↕↴∖↴∐↕≼↴∙⊾↥⋅∪⋯∐∐↕≼∐↖↽↥⋅↸∖↕⋜↕⊓∏↴∖↴⊓↸⊳ ↖↖↕∐↸⊳∐↕↴∖↴∖↸∖↸∖∐∐↕↑∐↸∖↴∖↻∐↘↸∖⋜↧↑∣∐≺∏∐↴∖↕↴⊺⋜⋯≼↧⋜↧∶↴∙⋜↧⋯ nore prominently⋅ at ↽↽Ll hours UTUp probably arises: from: ⋅ ⋅↔⋅ ⋅↴ ⋜↧⋯↕∐∖∐∖∐↑↸∖∐∐↴∖↴∖↕∪∐, The right-hand polarized emission which is seen in the spike at 7 hours UT and again more prominently at 11 hours UT probably arises from a coherent emission mechanism. +⋯↸∖↸⊳∐⋜∐∐↴∖⋯∙↽∕∏∐↴∖↕↴∖≼∐↴∖↸⊳∏↴∖↴∖↸∖≼⊔↿∐↑∐↸∖↥ ↴⋅↴⋅↴⋅↴↔ ⋅⋯∐↸∖≼∐↴∖↴⊓⋅↕↴⋝∏↑↕∪∐⋯⋜∏⇁≼∐∖↖⇁↸∖↕∪↻∙↽∕∏∐↴∖↴↸∖∐∐↴∖∷∖↴↕∪∐↖↖↽∪∏↕≼↧∪↸⊳↸⊳↿∐⋅ yelow., This is discussed further below. + The: initialA luminosity. was around qe1«10 ere + |, The initial luminosity was around $1.3\times10^{17}$ erg $^{-1}$ $^{-1}$. +" The extra fux appearing after UTS7 hows has a huninosity of L1«1019 eres ! 1, ", The extra flux appearing after UT=7 hours has a luminosity of $4.4\times10^{16}$ erg $^{-1}$ $^{-1}$ . +The luminosity of the extra flux appearing after UT=11 hows was Ls«1010 ere + +., The luminosity of the extra flux appearing after UT=11 hours was $1.8\times10^{16}$ erg $^{-1}$ $^{-1}$. + Both the steady huninosity and the extra flux appearing during the variations have similar values to those observed in typical RS CVu systems. whose flux is also believed to arise in large-scale powerful maguetic fields (e.g. Beng Cuicdelel. 199L).," Both the steady luminosity and the extra flux appearing during the variations have similar values to those observed in typical RS CVn systems, whose flux is also believed to arise in large-scale powerful magnetic fields (e.g. Benz Güddel, 1994)." + The polarization flip which was observed with the first flux increase at UT=7 hours is similar to that seen bv Skinner Brown (1991)., The polarization flip which was observed with the first flux increase at UT=7 hours is similar to that seen by Skinner Brown (1994). + They attributed this to à chauge in the relative optical depth of the o- aud ο ποιος during the brighter phase. after a model for RS CVu systems developed by Morris. Mutel Su (1990).," They attributed this to a change in the relative optical depth of the $o$ - and $x$ -modes during the brighter phase, after a model for RS CVn systems developed by Morris, Mutel Su (1990)." + Since. these modes correspond (unmider normal conditions) to different observed polarizations. such changes can lead to the observed polarization fip.," Since these modes correspond (under normal conditions) to different observed polarizations, such changes can lead to the observed polarization flip." + Another possible explanation would be a change in the seuse of the dominant maenuetic field along the line of sight. perhaps caused by stellar rotation. or the addition of new lighly-polarized emission.," Another possible explanation would be a change in the sense of the dominant magnetic field along the line of sight, perhaps caused by stellar rotation, or the addition of new highly-polarized emission." + The size of the compact source is constrained from the VLBI data to be less than 0.£8 mas., The size of the compact source is constrained from the VLBI data to be less than 0.48 mas. + At a distance of 110 pc this corresponds to a scale of 115 BR..., At a distance of 140 pc this corresponds to a scale of 14.5 $_{\odot}$. + This leneth scale is not ach leer than the expected size of a T Tauri imagnetosplere., This length scale is not much larger than the expected size of a T Tauri magnetosphere. + We can also estimate the size of the varviuso region from the timescale of the flux variations., We can also estimate the size of the varying region from the timescale of the flux variations. + The most rapid variations take place on timescales down to our fundamental inteeration time of 10 s. This indicates a source of length scale at most ονdt=3<10? an. or L3 R..," The most rapid variations take place on timescales down to our fundamental integration time of 10 s. This indicates a source of length scale at most $c\times dt= +3\times10^{9}$ m, or 4.3 $_{\odot}$." + This vields a lower limit to the briehtuess temperature of around 2.10? EK for the region involved in the rapidly varving emission in cach case;, This yields a lower limit to the brightness temperature of around $2\times10^9$ K for the region involved in the rapidly varying emission in each case. + Whilst this is PAill consistent with evrosvuchrotron cussion from nou-electrous. the 10054 polarization of the The extra flux seen after the second flux increase is not.," Whilst this is still consistent with gyrosynchrotron emission from non-thermal electrons, the right-hand polarization of the extra flux seen after the second flux increase is not." + A coherent eissiou οςλα‘πάσι must be invoked. to ]produce Us strong polarization., A coherent emission mechanism must be invoked to produce this strong polarization. + ↽∕∏∐∖⋯∪↴∖↴↑↕∐↽↸∖↕↖↽↸⊳∪∐↸∖↥⋅↸∖∐↑↸∖⊔↕↴∖↴↴∖↴↕∪∐≻↥⋅∪↸⊳↸∖↴∖↴↴∖↴↕↴∖↴⋪⋯ ∙∙ mE qa ⋅ . d ectrous Drap]trapped in magneticco flux tubes. where ‘a loss- us ..," The most likely coherent emission process is an electron cyclotron maser, arising from mildly relativistic electrons trapped in magnetic flux tubes, where a loss-cone distribution may develop." +" D at the fundamental or a low harmonic of the evrofrequency S.LCHz emission at the fundamental or first harmonic therefore implies a maguetic field streneth B~Lb 3kC. Although the cimission from au individual maser pulse would typically have a frequency range of aboutσεν, the eussson arises from a range of locations in the flux tube auc hence a rauge of uaguetic field strengths. which gives rise to a much broader spectrum."," This emission would occur at the fundamental or a low harmonic of the gyrofrequency 8.4GHz emission at the fundamental or first harmonic therefore implies a magnetic field strength $B \sim 1.5-3$ kG. Although the emission from an individual maser pulse would typically have a frequency range of about, the emission arises from a range of locations in the flux tube and hence a range of magnetic field strengths, which gives rise to a much broader spectrum." + The main alternative to this would be a plasima maser. with the cmuitting frequency being the fiudameutal or first harmonic ol the plasiia frequency or at the hybrid frequency (vs|0201/7 with My>Ms," The main alternative to this would be a plasma maser, with the emitting frequency being the fundamental or first harmonic of the plasma frequency or at the hybrid frequency $(\nu_p^2 + \nu_c^2)^{1/2}$ with $\nu_p > +\nu_c$." +" This would imply 5,—10770. 7", This would imply $n_e \sim 10^{12}$ $^{-3}$. +" For in accreting T. Tauri star. a possible hypothetical cutting region ix a iaenetically confined accretion fuuncl which would provide aregion with maguetic field. high density aud high density gradients at the stream οσο,"," For an accreting T Tauri star, a possible hypothetical emitting region is a magnetically confined accretion funnel, which would provide aregion with magnetic field, high density and high density gradients at the stream edge." + For a typical T Tau accretion rate of, For a typical T Tauri accretion rate of +truck contiuues to cover equal are length ou the globe it covers larger αμα larger clistances ou the map.,truck continues to cover equal arc length on the globe it covers larger and larger distances on the map. + Thus. the center of the segmeut fat 15° latitude) is uot centered on the segment on the map.," Thus, the center of the segment (at $45^\circ$ latitude) is not centered on the segment on the map." + We define tle skewuess: Taking the explicit case of the Mercator projection. we fiuc: and thus: so the skewness (for a vector pointed N-5) is simply: The skewuess at 15° is 1. showing a lopsicleduess toward the north.," We define the skewness: Taking the explicit case of the Mercator projection, we find: and thus: so the skewness (for a vector pointed N-S) is simply: The skewness at $45^\circ$ is 1, showing a lopsidedness toward the north." + Given this relation. the skewness is positive (nortliward lopsideduess) iu tle northeru hemisphere aud negative (soutliwar lopsidecluess) in the southern hemisphere.," Given this relation, the skewness is positive (northward lopsidedness) in the northern hemisphere and negative (southward lopsidedness) in the southern hemisphere." + Conskler a geodesic through a poiut iu the northeru bemisphere tipped at an azimuth augle ol 0 with respect to north., Consider a geodesic through a point in the northern hemisphere tipped at an azimuth angle of $\theta$ with respect to north. + The only thing increasing the speed of the truck is the gradient. of the scale factor as one moves northward. so the amplitude of the parallel acceleration is equal to the maximum acceleration (obtained golug straight. north) times cos 0.," The only thing increasing the speed of the truck is the gradient of the scale factor as one moves northward, so the amplitude of the parallel acceleration is equal to the maximum acceleration (obtained going straight north) times cos $\theta$." + To get the average of the absolute value of the skewness for all geodesics through that point at all random angles 0. oue simply integrates over 0: As with the flexion. we can integrate this over all points on the sphere to produce the average skewness over tlie whole sphere 5. Similarly. we fiuc: Notice that this is exactly the same value as the average flexion. F. lor the Mercator.," To get the average of the absolute value of the skewness for all geodesics through that point at all random angles $\theta$ , one simply integrates over $\theta$: As with the flexion, we can integrate this over all points on the sphere to produce the average skewness over the whole sphere S. Similarly, we find: Notice that this is exactly the same value as the average flexion, $F$, for the Mercator." + We will iud that for conformal projectious. the average absolute value of the skewness and flexion at a eiven polut and over the whole globe are always equal. (," We will find that for conformal projections, the average absolute value of the skewness and flexion at a given point and over the whole globe are always equal. (" +This is only true for conformal projections. or general projections tle skewuess aud flexion cau be different. as illustrated by the enomouic srojection which bas zero flexion but nou-zero skewness.),"This is only true for conformal projections, for general projections the skewness and flexion can be different, as illustrated by the gnomonic projection which has zero flexion but non-zero skewness.)" + Iu the Mercator projection. at the equator. the skewiess is zero. as we would expect from symaunetry cousideratious.," In the Mercator projection, at the equator, the skewness is zero, as we would expect from symmetry considerations." + ludeed. because any geodesic crossing equator has a syminetric shape in the northern and southern hemisphere. the skewuess s = 0 lor any geodesic line evaluated at a j»oiut on the equator.," Indeed, because any geodesic crossing equator has a symmetric shape in the northern and southern hemisphere, the skewness s = 0 for any geodesic line evaluated at a point on the equator." + Likewise. the [lexiou is zero for any geodesic line evaluated at a point ou the equator.," Likewise, the flexion is zero for any geodesic line evaluated at a point on the equator." + 50 the Mercator map has perfect local shapes along the equator. uniform scale alongthe," So the Mercator map has perfect local shapes along the equator, uniform scale alongthe" +orbital motion could be detectable with a high-resolution spectrum.,orbital motion could be detectable with a high-resolution spectrum. +" The possible detection of a change in the separation can also be interpreted as evidence for slightly different proper motion between HR 7329 A and companion candidate, namely a difference of 0.66+0.57 km/s (namely 2.91+ mas/yr in 47.741.5 pc)."," The possible detection of a change in the separation can also be interpreted as evidence for slightly different proper motion between HR 7329 A and companion candidate, namely a difference of $0.66 \pm 0.57$ km/s (namely $2.91 \pm 2.41$ mas/yr in $47.7 \pm 1.5$ pc)." + This value is comparable to the typical velocity dispersion in £8 Pic and other young associations (see Sect., This value is comparable to the typical velocity dispersion in $\beta$ Pic and other young associations (see Sect. +" 2), so that one can still not exclude that the two objects called HR. 7329 A and B are two independent members of the 8 Pic association."," 2), so that one can still not exclude that the two objects called HR 7329 A and B are two independent members of the $\beta$ Pic association." + The probability for this possibility is very low., The probability for this possibility is very low. + Orbital motion with curvature is not yet detected., Orbital motion with curvature is not yet detected. +" The most precise measurement of the separation between A and is the long 22.5 min exposure in Aug 2008 with NACO (4194B+16 mas), even though HR 7329 A issaturated?;; the position of A was determined with MIDAS center/moment, the position of B after subtraction of the PSF of A with MIDAS center/gauss; with a distance of 47.7+1.5 pc, the projected physical separation between HR 7329 A and B is then 200+16 AU, the semi-major axis for a pole-on circular orbit; for a uniform eccentricity distribution (e=0 to 1) and a random viewing angle, we correct this value by a factor of 1.10085 (Torres 1999, Allers et al."," The most precise measurement of the separation between A and B is the long 22.5 min exposure in Aug 2008 with NACO $4194 \pm 16$ mas), even though HR 7329 A is; the position of A was determined with MIDAS center/moment, the position of B after subtraction of the PSF of A with MIDAS center/gauss; with a distance of $47.7 \pm 1.5$ pc, the projected physical separation between HR 7329 A and B is then $200 \pm 16$ AU, the semi-major axis for a pole-on circular orbit; for a uniform eccentricity distribution (e=0 to 1) and a random viewing angle, we correct this value by a factor of $1.10 ^{+0.91} _{-0.36}$ (Torres 1999, Allers et al." + 2009) and obtain 220*21 AU., 2009) and obtain $220 ^{+214} _{-84}$ AU. + We use 2.2+0.1 Mo as mass of the star HR 7329 (Tetzlaff et al., We use $2.2 \pm 0.1$ $_{\odot}$ as mass of the A0-type star HR 7329 (Tetzlaff et al. +" 2011) and as total mass of HR 7329 A+B. For this system mass, the orbital period would then be ~1900 yrs for a semi-major axis of 200 AU (2200 yrs for 220 AU, 345 to 6100 yrs for 64 to 434 AU)."," 2011) and as total mass of HR 7329 A+B. For this system mass, the orbital period would then be $\sim 1900$ yrs for a semi-major axis of 200 AU (2200 yrs for 220 AU, 345 to 6100 yrs for 64 to 434 AU)." + Smith et al. (, Smith et al. ( +2009) directly detected the debris disk around HR 7329 A with an outer radius being 24 AU.,2009) directly detected the debris disk around HR 7329 A with an outer radius being 24 AU. +" From the very existence (and, hence, stability) of this debris disk (and its outer radius), we can constrain the eccentricity of HR 7329 B even further: Its eccentricity cannot be too large, otherwise it would fly through the disk."," From the very existence (and, hence, stability) of this debris disk (and its outer radius), we can constrain the eccentricity of HR 7329 B even further: Its eccentricity cannot be too large, otherwise it would fly through the disk." + Our deep imaging (Fig., Our deep imaging (Fig. + 3 below) shows that there is no additional companion outside of 24 AU (or between 24 and 200 AU) with a mass larger than ~20 Mjup., 3 below) shows that there is no additional companion outside of 24 AU (or between 24 and 200 AU) with a mass larger than $\sim 20$ $_{jup}$. +" If we further assume that HR 7329 B is responsible for shaping the debris disk and thereby fixing its outer radius, we can constrain the eccentricity as follows: Following Pichardo et al. ("," If we further assume that HR 7329 B is responsible for shaping the debris disk and thereby fixing its outer radius, we can constrain the eccentricity as follows: Following Pichardo et al. (" +"2005), for the masses given here for HR 7329 A and B, and assuming that HR 7329 B has its apocenter at 200 AU (see above) and is responsible for the outer disk radius at 24 AU, we determine the pericenter distance of HR 7329 B to be 71 AU, its semi-major axis to be 136 AU, and, hence, its eccentricity to be e=0.47.","2005), for the masses given here for HR 7329 A and B, and assuming that HR 7329 B has its apocenter at 200 AU (see above) and is responsible for the outer disk radius at 24 AU, we determine the pericenter distance of HR 7329 B to be 71 AU, its semi-major axis to be 136 AU, and, hence, its eccentricity to be $e=0.47$." +" Then, the orbital period would be ~10° yrs."," Then, the orbital period would be $\sim 10^{3}$ yrs." +" In the three deepest images (1998 HST NICMOS, 2009 VLT NACO, and 2008 VLT NACO H-band), no additional companion candidates were detected up to <9"" separation."," In the three deepest images (1998 HST NICMOS, 2009 VLT NACO, and 2008 VLT NACO H-band), no additional companion candidates were detected up to $\le 9 ^{\prime \prime}$ separation." +" Companions with 12 Myrs age with 13 Jup (or 1 Jup) masses, would have a luminosity of log(L/Lco)—4 (or -5.9) (Burrows et al."," Companions with 12 Myrs age with 13 Jup (or 1 Jup) masses, would have a luminosity of $\log (L/L_{\odot}) \simeq -4$ (or -5.9) (Burrows et al." +" 1997), hence a magnitude difference of ~10 mag (or 14.7 mag) to HR 7329 A, they would just be detectable at >1"" (or >3"", respectively) with NICMOS and NACO (Fig."," 1997), hence a magnitude difference of $\sim 10$ mag (or 14.7 mag) to HR 7329 A, they would just be detectable at $\ge 1 ^{\prime \prime}$ (or $\ge 3 ^{\prime \prime}$, respectively) with NICMOS and NACO (Fig." + 3)., 3). +" At ~10"" separation (~500 AU), two companion candidates are detected in the HST images with J=21 and H=17.5 mag, which would be in the planetary mass regime, but probably are background; they are outside the NACO S13 field (~9"" radius around HR 7329 A), too faint and/or blue for the NACO L-band L27 field (AL~10 mag), and too close and/or faint for the ISAAC and Sofl images, and are, hence, detected only once."," At $\sim 10 ^{\prime \prime}$ separation $\sim 500$ AU), two companion candidates are detected in the HST images with J=21 and H=17.5 mag, which would be in the planetary mass regime, but probably are background; they are outside the NACO S13 field $\sim 9 ^{\prime \prime}$ radius around HR 7329 A), too faint and/or blue for the NACO L-band L27 field $\Delta L \simeq 10$ mag), and too close and/or faint for the ISAAC and SofI images, and are, hence, detected only once." + We determined the dynamic range for all images by measuring the 3c above the background noise for any pixel (or group of 3 or 9 or 49 pixels) in all co-added images and compared this background flux to the flux of the central star HR. 7329 A. The flux ratio between background and HR 7329 is plotted in Fig., We determined the dynamic range for all images by measuring the $3 \sigma$ above the background noise for any pixel (or group of 3 or 9 or 49 pixels) in all co-added images and compared this background flux to the flux of the central star HR 7329 A. The flux ratio between background and HR 7329 A is plotted in Fig. +" 3 for the images with the best dynamic Aranges, i.e. where the closest and faintest companions could be detected."," 3 for the images with the best dynamic ranges, i.e. where the closest and faintest companions could be detected." +Nozawa et al. (,Nozawa et al. ( +1991) that was utilized by Ryden (1996) has the same qualitative feature of a narrow peak near p=0.5 and a steep decline toward pel.,1991) that was utilized by Ryden (1996) has the same qualitative feature of a narrow peak near $p=0.3$ and a steep decline toward $p=1$. + Hever et al. (, Heyer et al. ( +2001) note that he vast majority of their clouds (all but the largest clouds. which we loosely label GAIC’s) are no selt-exavitatiug.,"2001) note that the vast majority of their clouds (all but the largest clouds, which we loosely label GMC's) are not self-gravitating." + They are cither transient features or are held together by external pressure., They are either transient features or are held together by external pressure. + Hf these clouds are iudeed brought together by large scale turbulence πι the interstellar απο (or even confined for some time by an anisotropic rau. pressure) we nuelt expect that they have au elongated. fluneutairyv shape (sec e.g.àY Nagai. lIuutsuka. Alivama 1998: Balsara. Wiurd-Thonmpson. Crutcher 2001: review by Shu et al.," If these clouds are indeed brought together by large scale turbulence in the interstellar medium (or even confined for some time by an anisotropic ram pressure) we might expect that they have an elongated, filamentary shape (see e.g., Nagai, Inutsuka, Miyama 1998; Balsara, Ward-Thompson, Crutcher 2001; review by Shu et al." + 1999)., 1999). + Since even the largest clouds secu to have these shapes. we sunmuise that all clouds iav be brought together by external forcing (due to shock waves or turbulent motions for example). with only the largest clouds or densest regions within smaller clouds able to become sclteravitating.," Since even the largest clouds seem to have these shapes, we surmise that all clouds may be brought together by external forcing (due to shock waves or turbulent motions for example), with only the largest clouds or densest regions within smaller clouds able to become self-gravitating." +" This ties in with the eeneral picture of a rapid formation of molecular clouds due to external trigecrs (see ο,ος, Πατπα, Dallesteros-Paredes. Bereiu 2001: Pringle. Allen. Lubos 2001)."," This ties in with the general picture of a rapid formation of molecular clouds due to external triggers (see e.g., Hartmann, Ballesteros-Paredes, Bergin 2001; Pringle, Allen, Lubow 2001)." +together).,together). +" The density of the outflow at raclius r is Let the jets encounter the surroundingOs gas residinge within a distance r, and havinee a typical density ps.", The density of the outflow at radius $r$ is Let the jets encounter the surrounding gas residing within a distance $r_s$ and having a typical density $\rho_s$. + The head of each jet. proceedsal a speed vj given by the balance of pressures on its (wo sides., The head of each jet proceedsat a speed $v_h$ given by the balance of pressures on its two sides. +" Assuming supersonic motion this equality reads D=py(tg—Uny. whieh ean be solved for vy, ∖∖⊽∐≼↲↕⋅≼↲↥∐⊔∐↲⊳∖⇁≼↲≺∢∪↕∐⇂≼↲≺⇂∏≀↧↴∐↥∡∖↽⊔∐↲∐↓≀↧↪∖⊽⊳∖⇁↕∐∐∪∖∖↽↕⋅≀↧↴∥↲⇀⊔⊽∖≃∔⊼∕↗⊽∖⊔∣⋮⋮∶⋟⋃↽≻⋡∖↽≀↧↴⊳∖⊽⊳∖⇁∏∐↓↕↽≻∐∪∐≺≨↕⋟⋅↥⋯⊔∖⇁ been substituted."," Assuming supersonic motion this equality reads $\rho_s v_h^2 = \rho_f (v_f-v_h)^2$, which can be solved for $v_h$ where in the second equality the mass inflow rate $\dot M_s \simeq 4 \pi \rho_s \sigma r_s^2$ (by assumption 6), has been substituted." + The time required for the jets to cross the surrounding gas aud break out ol il is given bv where in the last equality (he same values as in equation (4)) have been used., The time required for the jets to cross the surrounding gas and break out of it is given by where in the last equality the same values as in equation \ref{eq:vh1}) ) have been used. + If there are no changes5 in the relative 5geometry of the SMDILI and inflowing5 mass. the jets will rapidly penetrate the surrounding gas ancl expand to large distances.," If there are no changes in the relative geometry of the SMBH and inflowing mass, the jets will rapidly penetrate the surrounding gas and expand to large distances." + In this case the jets will not deposit their energy in (he inflowing gas., In this case the jets will not deposit their energy in the inflowing gas. + For an ellicient deposition of energv to the inflowing gas. we require that there will be a relative motion between the SAIBIT and the inflowing gas. such that the jets continnously encounter fresh mass.," For an efficient deposition of energy to the inflowing gas, we require that there will be a relative motion between the SMBH and the inflowing gas, such that the jets continuously encounter fresh mass." + The relevant time is the time that the transverse motion of the jet crosses it width τι=ΙωoDijo. as bv our assumption 5 the relative velocity is e76.," The relevant time is the time that the transverse motion of the jet crosses it width $\tau_s \equiv D_j/v_{\rm rel} \simeq D_j/\sigma$, as by our assumption 5 the relative velocity is $v_{\rm rel} \simeq \sigma$." +" The width of the jet αἱ a distance r. from its source is D;=2r,sina. whereà is the half opening angle of the jet."," The width of the jet at a distance $r_s$ from its source is $D_j=2 r_s \sin \alpha$, where$\alpha$ is the half opening angle of the jet." +" For a narrow jet The demand for efficient energy deposition. τε5 /,. reads then"," For a narrow jet $\sin \alpha \simeq \alpha \simeq (2 \delta)^{1/2}$ , and The demand for efficient energy deposition, $\tau_s \la t_p$ , reads then" +of superhuninal motions and detection of significant X-ray variability at kpc-scales in Ms7 (Birettaetal.1999:ILarris2003.respectivelv).,"of superluminal motions and detection of significant X-ray variability at kpc-scales in M87 \citep[respectively]{bir99,har03}." +. In fact. a small number of detected optical ancl X-ray jets in radio galaxies as compared to the number of known radio jels suggests (hat relativistic beaming effects can play an iniportant role in these sources (Sparksοἱal.1995:ScarpaandUrry2002:Jester2003).," In fact, a small number of detected optical and X-ray jets in radio galaxies as compared to the number of known radio jets suggests that relativistic beaming effects can play an important role in these sources \citep{spa95,sca02,jes03}." +". Including a relativistic correction. the electron break Lorentz factor is where Voyti=raus/LOM Hz and B.,=B/10! G. Let us consider a kpc-scale (ο.1—4 kpc) relativistic jetwith the bulk Lorentz factor D. inclined at an angle 8 to the line of sight."," Including a relativistic correction, the electron break Lorentz factor is where $\nu_{syn, \, 14} \equiv \nu_{syn, \, br} / 10^{14}$ Hz and $B_{-4} \equiv + B / 10^{-4}$ G. Let us consider a kpc-scale $r \sim 1 - 4$ kpc) relativistic jetwith the bulk Lorentz factor $\Gamma$, inclined at an angle $\theta$ to the line of sight." +" The energy densitv of the jet magnetic field. as measured in the emitting region rest [rame. is As mentioned above. the equipartition value inferred from the svnehrotron emission of FR I knot regions is typically Be,eLO} G. and it corresponds to the observed synchrotron jet Iuminosities Ly,~LOM10—107? erg/s. The observed Iuminosity is related to the total emitted synchrotron power £5,/by the relation where in cases of a continuousjel (cj) or a moving single racdiating blob (i5). respectively (Sikoraetal.1997.seealsoAppendixA)"," The energy density of the jet magnetic field, as measured in the emitting region rest frame, is As mentioned above, the equipartition value inferred from the synchrotron emission of FR I knot regions is typically $B_{eq} \sim 10^{-4}$ G, and it corresponds to the observed synchrotron jet luminosities $L_{syn} \sim 10^{40} + - 10^{42}$ erg/s. The observed luminosity is related to the total emitted synchrotron power $L'_{syn}$by the relation where in cases of a continuousjet $cj$ ) or a moving single radiating blob $mb$ ), respectively \citep[see also Appendix A]{sik97}." +" ""Below we analvse these both possibilities as limiting models [ον the kpc-scale knots.", Below we analyse these both possibilities as limiting models for the kpc-scale knots. +" We also consider svnchrotron luminosity al a given break freeueney 75, pe. hereafter denoted as |. [pL,o. vather than the bolometric one L.,,"," We also consider synchrotron luminosity at a given break frequency $\nu_{syn, \, br}$ , hereafter denoted as $[\nu + L_{\nu}]_{syn, \, br}$ , rather than the bolometric one $L_{syn}$ ," +" We also consider svnchrotron luminosity al a given break freeueney 75, pe. hereafter denoted as |. [pL,o. vather than the bolometric one L.,,."," We also consider synchrotron luminosity at a given break frequency $\nu_{syn, \, br}$ , hereafter denoted as $[\nu + L_{\nu}]_{syn, \, br}$ , rather than the bolometric one $L_{syn}$ ," +released curing their formation.,released during their formation. + They may also possess leavy-element atimospleres if siguilficant light-element fallback cid uot occur during the supernova: their short lifetimes aud stroug magnetic fields render unlikely a significant accumulation of hydrogen rich mate‘ial from the interstellar umedium that could suppress atomic lines., They may also possess heavy-element atmospheres if significant light-element fallback did not occur during the supernova: their short lifetimes and strong magnetic fields render unlikely a significant accumulation of hydrogen rich material from the interstellar medium that could suppress atomic lines. + In the case of bursters. the heavy elements in their atmospheres are continually replenished by accretion aud the thermouuceal Hashes provide large amounts of thermal energy.," In the case of bursters, the heavy elements in their atmospheres are continually replenished by accretion and the thermonuclear flashes provide large amounts of thermal energy." + Both types of neutron stars that are prime cauclicates for the detectio1 of spectral lines are fast 'otators. (, Both types of neutron stars that are prime candidates for the detection of spectral lines are fast rotators. ( +We do not cousider here maguetars. for which the presence of τιtrastroug iagnetic fields introduces large uncertainties in calculating the rest energies of atonuie lires).,"We do not consider here magnetars, for which the presence of ultrastrong magnetic fields introduces large uncertainties in calculating the rest energies of atomic lines)." + The spin frequencies of known pulsars with ages <104 vr is between ~5—65 Hz (see. e.g.. Becker Pavlov 2002): the inferred spin frequencies of bursters is between ~270—620 Hz (Strohmayer 2001).," The spin frequencies of known pulsars with ages $<10^4$ yr is between $\simeq 5-65$ Hz (see, e.g., Becker Pavlov 2002); the inferred spin frequencies of bursters is between $\simeq +270-620$ Hz (Strohmayer 2001)." + These high spin recquencies iutroduce several relativistic effects such as Doppler boosts. stroug self-leusiig. aud differential gravitational redshift arisiug [rom the stellar oblateness.," These high spin frequencies introduce several relativistic effects such as Doppler boosts, strong self-lensing, frame-dragging and differential gravitational redshift arising from the stellar oblateness." + All of hem alter he line profiles observed at iufinity., All of them alter the line profiles observed at infinity. + Iu thisLetter. we show the effects of relativistic Doppler boosts aud strong gravitatioial leusiug ou the width and asyiumetry of line profiles originating [rom the surfaces of rotating neutron stars.," In this, we show the effects of relativistic Doppler boosts and strong gravitational lensing on the width and asymmetry of line profiles originating from the surfaces of rotating neutron stars." + We theu investigate the systematic uncertainties introduced by these effects in iuerring the comipactuess of μοιήτοι stars., We then investigate the systematic uncertainties introduced by these effects in inferring the compactness of neutron stars. + Neutron stars are the most Compact stelar objects aud the rotational velocities at their surfaces can reach au appreciable [fraction of the 5;»eed of light., Neutron stars are the most compact stellar objects and the rotational velocities at their surfaces can reach an appreciable fraction of the speed of light. + Therefore. rotaional eTects on the lie proliles originating from the neutron star surface are qualitatively different. [r«4u the case of a Newtonian slowly spiunuing star.," Therefore, rotational effects on the line profiles originating from the neutron star surface are qualitatively different from the case of a Newtonian slowly spinning star." + lu particuar. the siape of the liue profies obseved at iulinity S altered in four ways.," In particular, the shape of the line profiles observed at infinity is altered in four ways." + First. relativistic Doppler boosts give ‘Ise 10 al asVlunetry iu the 5»ectral iue profiles whie broadening them.," First, relativistic Doppler boosts give rise to an asymmetry in the spectral line profiles while broadening them." +" Secoucl. stroug gravitatioial lensing of swlace emission by theteutron star alte"" the relative contribution of surface elements with cil‘ereul line-o[-siglit. veocities o the line profile."," Second, strong gravitational lensing of surface emission by the neutron star alters the relative contribution of surface elements with different line-of-sight velocities to the line profile." + Third. frame crageing in the rotating spacetime of the neτοι star allecs the ploton trajectories aud thus the observable surface elements.," Third, frame dragging in the rotating spacetime of the neutron star affects the photon trajectories and thus the observable surface elements." + Finally. tje stelar oblateuess caused by the rapid spin introduces a difference in the gravitational redshifts of Liles that are generated at the rotational equator and at the poles.," Finally, the stellar oblateness caused by the rapid spin introduces a difference in the gravitational redshifts of lines that are generated at the rotational equator and at the poles." + Iu thisLetter. we use the numerical methods described in Mino. Ozzel. Chakrabarty (2009 aic references therein) to caleulate spectral liue profiles taking into account the first two of the effects discussed above.," In this, we use the numerical methods described in Muno, Özzel, Chakrabarty (2002 and references therein) to calculate spectral line profiles taking into account the first two of the effects discussed above." + Investigatiug the latter two requires the caleulatiou of numerical spacetimes for specific equations of state of ueutrou-star matter aud will be addressed iu a forthcoming paper., Investigating the latter two requires the calculation of numerical spacetimes for specific equations of state of neutron-star matter and will be addressed in a forthcoming paper. + The results presented here are accurate for spin frequeucies sinaller tliau the Ixepleriau frequeucy, The results presented here are accurate for spin frequencies smaller than the Keplerian frequency +eive a precise (wo dimensional location (without highly precise timing information). bul merely confines the source to a one-dimensional locus.,"give a precise two dimensional location (without highly precise timing information), but merely confines the source to a one-dimensional locus." + We were able. however. (o use the two observations together to obtain a two-dimensional localization as each observation was taken at a slightly dillerent orientation. with roll angles differing by 4.87.," We were able, however, to use the two observations together to obtain a two-dimensional localization as each observation was taken at a slightly different orientation, with roll angles differing by $4.8\degr$." + When the positions are combined. (here is a quasi-elliplical region of overlap centered αἱ Oogg)—1950714362. day)=FOO?LO'20V04. with statistical l-o error contours of major axis (728. semi-minor axis 001. and position angle 73° East of North (Figure 6)).," When the one-dimensional positions are combined, there is a quasi-elliptical region of overlap centered at $\alpha_{2000}=19^{\rm h}07^{\rm m}14\fs362$ , $\delta_{2000}=+09\degr19\arcmin20\farcs04$, with statistical $\sigma$ error contours of semi-major axis $0\farcs28$, semi-minor axis $0\farcs01$, and position angle $73\degr$ East of North (Figure \ref{fig:pos}) )." +" This position has absolute systematic uncertainties of up to 1"" due to aaspect errors. bul we have here assumed (hat the separate observations possess internally consistent aspect determinations to SO.) aresec: this agrees with our experience with other delata sets."," This position has absolute systematic uncertainties of up to $1\arcsec$ due to aspect errors, but we have here assumed that the separate observations possess internally consistent aspect determinations to $\lesssim$ 0.1 arcsec; this agrees with our experience with other data sets." + Under (his assumption. the resulting localization is consistent (o a high degree ol accuraey with the position of the August 1998 radio transient. (Frailetal.1999).," Under this assumption, the resulting localization is consistent to a high degree of accuracy with the position of the August 1998 radio transient \citep{fkb99}." +. Frailetal.(1999). reported. following the giant August. 1998 flare. the detection of a transient radio source.," \citet{fkb99} reported, following the giant August 1998 flare, the detection of a transient radio source." + Their observations covered the time interval from 1 week to 1 month alter the burst., Their observations covered the time interval from 1 week to 1 month after the burst. + The source was detected in the first observation. | week alter the burst. and then declined over the course of the following four observations (930 d: Fig. 1)).," The source was detected in the first observation, 1 week after the burst, and then declined over the course of the following four observations (9–30 d; Fig. \ref{fig:radio}) )." + Thus a least for this giant flare — the radio source appears to have peaked about a week alter the burst ancl subsequently undergone a power-law decay., Thus – at least for this giant flare – the radio source appears to have peaked about a week after the burst and subsequently undergone a power-law decay. + For the April 2001 flare. we undertook VLA observations beginning 0.17 d alter the event aud ending almost two weeks later (Table 4)).," For the April 2001 flare, we undertook VLA observations beginning 0.17 d after the event and ending almost two weeks later (Table \ref{tab:vla}) )." + Despite our prompt radio observations. we did not detect a radio source comparable in strength to the August 1993 flare at anv of our five epochs of observation.," Despite our prompt radio observations, we did not detect a radio source comparable in strength to the August 1998 flare at any of our five epochs of observation." + The fluence of the August 1998 Hare was LO7 ((here we include (he contribution to the fluence [rom the initial hard spike and (he subsequent softer afterglow: Ferocietal. 2001))., The fluence of the August 1998 flare was $10^{-2}$ (here we include the contribution to the fluence from the initial hard spike and the subsequent softer afterglow; \citealt{fhd+01}) ). + In contrast. the f[Inence of the April 2001 flare was 26x10| citeph-4-01..," In contrast, the fluence of the April 2001 flare was $2.6\times 10^{-4}$ \\citep{h+01}." + The inferred peak flux of the transient radio source for the August 1993 flare was about 400 j(Jv in the 8.46-GIHIz band., The inferred peak flux of the transient radio source for the August 1998 flare was about 400 $\mu$ Jy in the 8.46-GHz band. + If the radio flux is proportional to the energy releasedbv the flare then we would expect a peak radio flux of 10 jiJy in the same band some time, If the radio flux is proportional to the energy releasedby the flare then we would expect a peak radio flux of 10 $\mu$ Jy in the same band some time +The Red Reetanele (RR). an enigmatic protoplanctary nebula. reveals iself to be a complex astrophysical svstem.,"The Red Rectangle (RR), an enigmatic protoplanetary nebula, reveals itself to be a complex astrophysical system." + The central source of the nebula. Lb 44179. is a binary system in a sxνου evolutionary phase between the asymptotic gian branch (AGB) and the planetary nebula stage of its life. when the expelled circumstellar matter is ionized by the ;ohotometric primary star's hot. core (van 2007)..," The central source of the nebula, HD 44179, is a binary system in a short-lived evolutionary phase between the asymptotic giant branch (AGB) and the planetary nebula stage of its life, when the expelled circumstellar matter is ionized by the photometric primary star's hot core \citep{vanwinckel2003, vanwinckel2007, siodmiak2008, szczerba2007}." + The dom.inant features of the nebulas form are the eshaped biconical oulllows (nearby. on the skv (Cohenetal. 2004))) theu emerge from⋅ a dise. (roughly on the sky). seen nearly edge-on.," The dominant features of the nebula's form are the $\times$ -shaped biconical outflows (nearly on the sky \citep{cohen2004}) ) that emerge from a disc (roughly on the sky), seen nearly edge-on." + We will refer to this optically-thick disc. resolved. by (Cohenctal.200:4) and ground-basecl interferometric Observations (Rockeleretal.1995:Osterbart.Langer.&Weigelt’ LOOT). as he circumbinary disc.," We will refer to this optically-thick disc, resolved by \citep{cohen2004} and ground-based interferometric observations \citep{roddier1995, osterbart1997}, as the circumbinary disc." + Phe dimensions assumed for the cireumbinary disc are based on a distance of 710 pc (Menshehikovetal.2002)., The dimensions assumed for the circumbinary disc are based on a distance of 710 pc \citep{menshchikov2002}. +. While the distance is uncertain. we adopt this value in our investigation.," While the distance is uncertain, we adopt this value in our investigation." + Bujarrabalctal.(2005) raced the outer radius of the cireumbinary disc to. 1.550 AU.," \citet{bujarrabal2005} traced the outer radius of the circumbinary disc to 1,850 AU." + The radius of the central cavity in the circumbinarydisc is obtained from a scattered light model to be 14 AU (Menshehikovetal. 2002)., The radius of the central cavity in the circumbinarydisc is obtained from a scattered light model to be 14 AU \citep{menshchikov2002}. +. The thickness of the circumbinary disc has been determined to be between 90 XU (Alenshejkovetal.2002) and 107 AU (Dujarrabaletal.2005:Rocleerct1995: LOOT). ," The thickness of the circumbinary disc has been determined to be between 90 AU \citep{menshchikov2002} and 107 AU \citep{bujarrabal2005, roddier1995, osterbart1997}. ." +Lt is worth noting that, It is worth noting that +proper motion errors will bias our pericentre measurement to be systematically high (see further discussion below).,proper motion errors will bias our pericentre measurement to be systematically high (see further discussion below). + Our derived apocentre distances Τα are lower than the mean of the 29° VL1 subhalos (the 50 most massive at redshift z— 0)., Our derived apocentre distances $r_a$ are lower than the mean of the $z^{50}_0$ VL1 subhalos (the 50 most massive at redshift $z=0$ ). + This discrepancy cannot be explained by a bias due to the proper motion errors as this effect has the wrong sign., This discrepancy cannot be explained by a bias due to the proper motion errors as this effect has the wrong sign. +" Interestingly, however, the mean of our 210 sample (the 50 most massive before redshift z— 10) agrees well with the observed mean."," Interestingly, however, the mean of our $z^{50}_{10}$ sample (the 50 most massive before redshift $z=10$ ) agrees well with the observed mean." + This lends further support to the idea that the Milky Way's dwarfs formed early before reionisation (e.g.???7????77)..," This lends further support to the idea that the Milky Way's dwarfs formed early before reionisation \citep[e.g.][]{1992MNRAS.256P..43E,1999ApJ...523...54B,2000ApJ...539..517B,2002MNRAS.333..177B,2004ApJ...609..482K,2005ApJ...629..259R,2005MNRAS.364..367D,2006ApJ...645.1054G,2006MNRAS.368..563M,2009arXiv0903.4681M}." + The mean recovered Yp/Ta ratio is higher in the dwarfs than in the simulation samples., The mean recovered $r_p/r_a$ ratio is higher in the dwarfs than in the simulation samples. + Analogously to the pericenter distribution this can be explained by the bias from the proper motion errors., Analogously to the pericenter distribution this can be explained by the bias from the proper motion errors. +" We investigate the influence of satellite depletion by a disc using the three subsets [ρίγα= 10kkpc), zjQ(ra= l5kkpc) and z1ü(ra= 20kkpc) (50 most massive before z—10, without orbits that have pericenters rp 15kkpc is too extreme., This suggests that the total destruction of satellites within $r_d > 15$ kpc is too extreme. + With the current data it is not possible to discriminate between the 270 sample without disc depletion and zj9(ra=10 kkpc) with a disc with radius τα=10 kkpc., With the current data it is not possible to discriminate between the $z^{50}_{10}$ sample without disc depletion and $z^{50}_{10}(r_d=10$ kpc) with a disc with radius $r_d=10$ kpc. + We have evaluated how well we can recover the orbits of Milky Way satellites in the light of measurement errors and model limitations., We have evaluated how well we can recover the orbits of Milky Way satellites in the light of measurement errors and model limitations. +" To do this, we compared orbits in a high resolution cosmological simulation of a Milky Way analogue with similar orbits integrated in a fixed background potential."," To do this, we compared orbits in a high resolution cosmological simulation of a Milky Way analogue with similar orbits integrated in a fixed background potential." + We find that:, We find that: +chussion have been made by Ehlnouttie et al. (,emission have been made by Elmouttie et al. ( +1997). Ables et al. (,"1997), Ables et al. (" +1987). aud Dalilem et al. (,"1987), and Dahlem et al. (" +1993) respectively.,1993) respectively. + The larec-scale radio emission has a steep (frequeucyv) spectral index (a = l.l1zx0.1): 10 enussjon was detected from the halo., The large-scale radio emission has a steep (frequency) spectral index $\alpha$ = $\pm$ 0.1); no emission was detected from the halo. + was found to extend out to F--dπι with au additional feature at 10 in the south-west., was found to extend out to $r=7'$ with an additional feature at $10'$ in the south-west. + eolussion is concentrated towards the nucleus. but was also detected out to 6 aloue the major axis.," emission is concentrated towards the nucleus, but was also detected out to $6'$ along the major axis." + We have used the Compact Array of the Australia Telescope National Facility (ATCA. see Frater Brooks 1992) to study the ο conutinuau aud emission of [1915 in more detail and with higher sensitivity than before.," We have used the Compact Array of the Australia Telescope National Facility (ATCA, see Frater Brooks 1992) to study the GHz continuum and emission of 4945 in more detail and with higher sensitivity than before." + We have also observed comission using the Swedish-ESO Subimillimetre Telescope (SEST: sce Booth et al., We have also observed emission using the Swedish-ESO Submillimetre Telescope (SEST; see Booth et al. + 1989)., 1989). + In addition to the results preseutec by Dahlem et al. (, In addition to the results presented by Dahlem et al. ( +1993) our ddata are not confined to the nuclear region but also include significant parts of the disk.,1993) our data are not confined to the nuclear region but also include significant parts of the disk. + aud CO have beeu compared at à conunon auegulaur resolution of ~237 and channel spacines of 6.6 andlius. corresponding to velocity resolutious of ~7.25 audkms. respectively: 237 corresponds to 750ppc at D = MMpe (sce also Table 1.. footnote τα).," and CO have been compared at a common angular resolution of $\sim$ $''$ and channel spacings of 6.6 and, corresponding to velocity resolutions of $\sim$ 7.25 and, respectively; $''$ corresponds to pc at $D$ = Mpc (see also Table \ref{TAB.INT.PROPERTIES}, footnote `a')." + The ATCA was equipped with 20cem FET receivers with system temperatures between 33 and KE. Four observing rus provided 53 independent baselines. aud were sufficient to achieve a satisfactory siuupliug of the plane.," The ATCA was equipped with cm FET receivers with system temperatures between 33 and K. Four observing runs provided 53 independent baselines, and were sufficient to achieve a satisfactory sampling of the plane." + Table 2. displavs the configurations. observiug dates and baseline ranges.," Table \ref{TAB.HI.CONFIG} displays the configurations, observing dates and baseline ranges." + A standard correlator setup of 512 channels was used with a total baudwidth of MMIIZz ceutred on MMIIz., A standard correlator setup of 512 channels was used with a total bandwidth of MHz centred on MHz. +" The chiiunel separation was KXIIz Κιν),", The channel separation was kHz ). + An observing cycle was adopted in which a lUO-1unute tracking of 11915. was bracketed by alternate 5-inin observations of the phase-calibratiou SOULCOS 11215-157 aud 11320-L16., An observing cycle was adopted in which a 40-minute tracking of 4945 was bracketed by alternate 5-min observations of the phase-calibration sources 1215-457 and 1320-446. + The radio ealaxy 1193L638 was observed as a flux-deusity calibrator (assunied to have a flux density of 16.1JJw at L.LGGITIZ)., The radio galaxy 1934-638 was observed as a flux-density calibrator (assumed to have a flux density of Jy at GHz). + Observations of the quasar 00107-6585 provided antenna spectral baudpasses., Observations of the quasar 0407-658 provided antenna spectral bandpasses. +" The data were reduced with an ATNF ioclificatiou (xilleen 1992) of the Astronomical Tage Processing System. CAIPS) of the US National Radio ΑποώμώινἩ Observatory,", The data were reduced with an ATNF modification (Killeen 1992) of the Astronomical Image Processing System (AIPS) of the US National Radio Astronomy Observatory. + A single data file of correlatedwr data was produced for each observing run., A single data file of correlated data was produced for each observing run. + After editing of the spectra obtained with individual baseliues for phase and amplitude errors. the central 100 spectral channels were selected.," After editing of the spectra obtained with individual baselines for phase and amplitude errors, the central 400 spectral channels were selected." + Those channels free of contributions were averaged in the domain for each of the four data sets and immersed iuto a continmuni-enüssiou data base., Those channels free of contributions were averaged in the domain for each of the four data sets and merged into a continuum-emission data base. +" Standard niagiug aud CLEAN (IHóseboni 1971) procedures were then applied. aud the resultant mage was further improved by a standard selfcalibration routine (οιο, Cormwvell Fomalout 1989)."," Standard imaging and CLEAN (Höggbom 1974) procedures were then applied, and the resultant image was further improved by a standard self-calibration routine (e.g. Cornwell Fomalont 1989)." + Iun the ολα. the coutiuuuni onidssion Was removed. frou channels containiug the data usine the line-free chanucls (van Laugevelde Cotton 1990: Coruwell et al.," In the domain, the continuum emission was removed from channels containing the data using the line-free channels (van Langevelde Cotton 1990; Cornwell et al." + 1992)., 1992). + After concatenation of the four data sets. the gain solutions obtained roni the calibration of the coutimmun image were applied.," After concatenation of the four data sets, the gain solutions obtained from the self-calibration of the continuum image were applied." + To facilitate subsequent analysis two sets of images were produced for roth he οσοπα and line data., To facilitate subsequent analysis two sets of images were produced for both the continuum and line data. +" ""Natural weighting’ provided a resolution of 19” (R.A) « 25"" (Dec.) for investigation of the extended structure: uniform weighting’ vielded a restore beam of 3.2% 4 |"" to reveal the spatial fue structure of the unclear region.", `Natural weighting' provided a resolution of $''$ (R.A.) $\times$ $''$ (Dec.) for investigation of the extended structure; `uniform weighting' yielded a restoring beam of $''$ $\times$ $''$ to reveal the spatial fine structure of the nuclear region. + A correction for the gain variation across the ATCA 22-1 antenna beams was applied to all images., A correction for the gain variation across the ATCA 22-m antenna beams was applied to all images. + The rus noise of the final coutimmiun inage is ~2.012Jy L, The rms noise of the final continuum image is $\sim2.0$ . +o Faint artifacts are esent with flux densities of ο +., Faint artifacts are present with flux densities of 5 –. +. The peak-to-noise ratio for the nucleus is of the order of 800., The peak-to-noise ratio for the nucleus is of the order of 800. + Averaging two contiguous channels (this vields a chamnel spacing of kms)) the ins noise of the data becomes, Averaging two contiguous channels (this yields a channel spacing of ) the rms noise of the data becomes. + To investigate the possibility of nüssius extended chussion. the galaxy was also nappe with the Parkes 6Lin telescope.," To investigate the possibility of missing extended emission, the galaxy was also mapped with the Parkes 64-m telescope." + The total flix density integrated over velocity for Parkes observations ) was found to be lower than its ATCA counterpart }., The total flux density integrated over velocity for Parkes observations ) was found to be lower than its ATCA counterpart ). + The difference is within the uucertaiutv of at least for cach estimate., The difference is within the uncertainty of at least for each estimate. + In παν. the comparison of fis. densities provides πο evidence for missing enuüssion m the ATCA observations.," In summary, the comparison of flux densities provides no evidence for missing emission in the ATCA observations." +" Ine cussion was mapped during 1993 March 110 with the SEST at an augular resolution of 227,", line emission was mapped during 1993 March 1–10 with the SEST at an angular resolution of $''$. + A GIIz SIS receiver vielded system temperatures of 600 TI ou an autenua temperature (£33) scale., A GHz SIS receiver yielded system temperatures of 600 – K on an antenna temperature ) scale. + The main beau efficiency was 0.16., The main beam efficiency was 0.46. + An acousto«optical spectrometer. with 1110 channels aud a total bandwidth of CGIIz. provided," An acousto-optical spectrometer, with 1440 channels and a total bandwidth of GHz, provided" +The compact groups and their member galaxies. presented in Tables 1 4 were selected. based on their photometric properties. in particular their projected. positions aud r band apparent magnitudes.,"The compact groups and their member galaxies presented in Tables 1 – 4 were selected based on their photometric properties, in particular their projected positions and $r-$ band apparent magnitudes." + However. the SDSS DRG has spectroscopic information available for 1.27 million objects. of which 679733 objects are galaxies which satisfy our criteria for inclusion in Catalogue D. Thus. many of the galaxies we identify as being à member of a compact group will have spectroscopic data available.," However, the SDSS DR6 has spectroscopic information available for 1.27 million objects, of which 679733 objects are galaxies which satisfy our criteria for inclusion in Catalogue B. Thus, many of the galaxies we identify as being a member of a compact group will have spectroscopic data available." + For the groups identified in Catalogue A. 4131 out of 9713 member galaxies which vielcl a reliable. redshift σονzc 0.7).," For the groups identified in Catalogue A, 4131 out of 9713 member galaxies which yield a reliable redshift $z_{conf} \ge 0.7$ )." + For groups identified in Catalogue B. 16566 out of 313508. member ealaxios have spectroscopic information available. 16405 of which vield a reliable redshift.," For groups identified in Catalogue B, 16566 out of 313508 member galaxies have spectroscopic information available, 16405 of which yield a reliable redshift." + The fraction of galaxies with spectroscopic data available in C'atalogue B is significantly less than in Catalogue A. since Catalogue A contains a," The fraction of galaxies with spectroscopic data available in Catalogue B is significantly less than in Catalogue A, since Catalogue A contains a" +"We now consider the barred, dwarf Sa (dSa) simulation in GalMer (RC of 100 km/s and initial scale-length of 1.3 kpc) to see how its density distribution and metallicity gradient evolve with time.","We now consider the barred, dwarf Sa (dSa) simulation in GalMer (RC of 100 km/s and initial scale-length of 1.3 kpc) to see how its density distribution and metallicity gradient evolve with time." + In Fig., In Fig. +" 5 we plot the time development of the stellar disk density and metallicity profiles, for the same time steps as in Fig. Ι.."," \ref{fig:dSa} we plot the time development of the stellar disk density and metallicity profiles, for the same time steps as in Fig. \ref{fig:gSa}." +" We can clearly see that the stellar disk extends to more than ten scale-lengths in ~3 Gyr, while preserving its exponential mass density profile."," We can clearly see that the stellar disk extends to more than ten scale-lengths in $\sim3$ Gyr, while preserving its exponential mass density profile." +" At the same time, the metallicity gradient becomes flat in less than 1 Gyr."," At the same time, the metallicity gradient becomes flat in less than 1 Gyr." +" Both the predicted extent of the stellar disk and the flattening of the metallicity profile are consistent with what has been recently reported for nearby low-mass galaxies (Bland-Hawthornetal.,2005;Vlajié 2009)."," Both the predicted extent of the stellar disk and the flattening of the metallicity profile are consistent with what has been recently reported for nearby low-mass galaxies \citep{bland05,vlajic09}." + We have examined the redistribution of angular momentum in galactic disks by means of Tree-SPH and high-resolution pure N-body simulations., We have examined the redistribution of angular momentum in galactic disks by means of Tree-SPH and high-resolution pure N-body simulations. + We have found that resonance overlap of multiple patterns (such as bar + SS or SS + SS) induces strong exchange of angular momentum throughout the disks in agreement with the predictions of MF10., We have found that resonance overlap of multiple patterns (such as bar + SS or SS + SS) induces strong exchange of angular momentum throughout the disks in agreement with the predictions of MF10. +" Since in the self-consistent simulations analyzed in this work spirals may be transient, we should also expect a contribution from the SBO2 radial migration mechanism."," Since in the self-consistent simulations analyzed in this work spirals may be transient, we should also expect a contribution from the SB02 radial migration mechanism." +" However, for the short timescales considered here (<1 Gyr) transients would simply have a very small effect (5Β02)."," However, for the short timescales considered here $<1$ Gyr) transients would simply have a very small effect (SB02)." +" The resonance overlap mechanism is clearly identified by a bimodality in the changes in angular momentum, AL, caused by the bar’s corotation and 2:1 OLR (Fig. 1,, 2,,"," The resonance overlap mechanism is clearly identified by a bimodality in the changes in angular momentum, $\Delta L$, caused by the bar's corotation and 2:1 OLR (Fig. \ref{fig:gSa}, \ref{fig:gSall}," + and 4))., and \ref{fig:nbody}) ). +" We have contrasted this to a simulation lacking a stable central bar (Fig. 2,,"," We have contrasted this to a simulation lacking a stable central bar (Fig. \ref{fig:gSall}," +" rightmost panel), where the AL distribution is smooth."," rightmost panel), where the $\Delta L$ distribution is smooth." +" The effect is especially strong when a gaseous component is present, as a result of the exchange of L between gas and stars coming from the gravity torques, the phase shift between the two components, and the gas dissipation (BournaudandCombes,2002)."," The effect is especially strong when a gaseous component is present, as a result of the exchange of $L$ between gas and stars coming from the gravity torques, the phase shift between the two components, and the gas dissipation \citep{bournaud02}." +". Depending on the amount of gas and the strength of the bars and spirals, the metallicity gradients can flatten in less than 1 Gyr (Fig. 1))."," Depending on the amount of gas and the strength of the bars and spirals, the metallicity gradients can flatten in less than 1 Gyr (Fig. \ref{fig:gSa}) )." +" This is in drastic contrast to the current understanding that galactic disks need a Hubble time for sufficient mixing (5Β02, Roskaretal. 2008))."," This is in drastic contrast to the current understanding that galactic disks need a Hubble time for sufficient mixing (SB02, \citealt{roskar08}) )." + How can we tie our results to galactic disk evolution?, How can we tie our results to galactic disk evolution? +" BournaudandCombes(2002) followed the detailed processes of bar formation, bar destruction, and bar re-formation, while varying the disk to bulge ratio."," \cite{bournaud02} followed the detailed processes of bar formation, bar destruction, and bar re-formation, while varying the disk to bulge ratio." + These authors identified three bar formation episodes in a Hubble time., These authors identified three bar formation episodes in a Hubble time. +" In this model, we can regard a given GalMer isolated galaxy simulation as one such episode of a gas accretion event during a galaxy lifetime."," In this model, we can regard a given GalMer isolated galaxy simulation as one such episode of a gas accretion event during a galaxy lifetime." +" In this scenario, the rapid flattening of the metallicity gradients expected from the vigorous migration (Fig. 1))"," In this scenario, the rapid flattening of the metallicity gradients expected from the vigorous migration (Fig. \ref{fig:gSa}) )" + would be followed by a gas enrichment at each bar re-formation event resulting in the rebuilding of the gradient., would be followed by a gas enrichment at each bar re-formation event resulting in the rebuilding of the gradient. +" The presence or not of a metallicity gradient, or its intensity, would then be an indicator of the bar/accretion phase of the galaxy."," The presence or not of a metallicity gradient, or its intensity, would then be an indicator of the bar/accretion phase of the galaxy." + The mechanism described in this study works even in mass galaxies (Fig. 5)), The mechanism described in this study works even in low-mass galaxies (Fig. \ref{fig:dSa}) ) + and can thus for the first time provide an explanation for the extended disk profiles observed in galaxies with V.~100 km/s (Bland-Hawthorn," and can thus for the first time provide an explanation for the extended disk profiles observed in galaxies with $V_c\sim100$ km/s \citep[Minchev et al. 2011, in +preparation]{bland05}. ." +needs additional modelling work that is bevoud the scope of this paper.,needs additional modelling work that is beyond the scope of this paper. + A quite different approach nay be considered on the basis of recent analysis of spectral variability of another Be/X-rav binary.63.. made hy Nemerucla aud Okazaki (20 H)) aud Neeucrucla et al. (200] D.," A quite different approach may be considered on the basis of recent analysis of spectral variability of another Be/X-ray binary, made by Negueruela and Okazaki \cite{NO2000}) ) and Negueruela et al. \cite{Neg2001}) )." + These authors argue that inCas.. just as in sinular De/N-rav systems. a major role is plaved by an equaorial decretiou disc arouid the optical star.," These authors argue that in, just as in similar Be/X-ray systems, a major role is played by an equatorial decretion disc around the optical star." + This disc is £1ncated as a result of tidal/resonaif interaction with the neutron star companion., This disc is truncated as a result of tidal/resonant interaction with the neutron star companion. + At a certain stage the disc becoues uustable. tilts aud war25 Alle starts fo process.," At a certain stage the disc becomes unstable, tilts and warps and starts to precess." + Later ou. the disc is disrupted die to interacion with the orbiting ueutron star. aud a ejut X-rav outburst may occur.," Later on, the disc is disrupted due to interaction with the orbiting neutron star, and a giant X-ray outburst may occur." + The above authors argue hat tus model does iot nuplv anv subsautial chanee in the optica stars nass oss rate., The above authors argue that this model does not imply any substantial change in the optical star's mass loss rate. +" Following that idea. we nnv suppose that the oecessine body in the system, is a ilted/warpecd dise. around Be star."," Following that idea, we may suppose that the precessing body in the system is a tilted/warped disc around Be star." + Then it is natural hat the C»eravitatioiab pull of the neutron star iu the noments of closest approach of the NS o the disc causes its distortion. which is observed 1i the Πο curve as Μου. variations supoerposed on global changes.," Then it is natural that the gravitational pull of the neutron star in the moments of closest approach of the NS to the disc causes its distortion, which is observed in the light curve as minor variations superposed on global changes." +" Due to recession. the times of closest approac1 do not coimcide with the periastron onssages, but raher precede thei by zsRI aftey each orbi."," Due to precession, the times of closest approach do not coincide with the periastron passages, but rather precede them by $\approx8^\mathrm{d}$ after each orbit." + Within this approach we cau fud a jatural explanation for the fact tlat chiving the ascending outs of the elobal light curve. no 1day modulation is seen: it only means that the decretion disc is not large enough to become tilted aux lwar»ed and therefore it docs iof precess.," Within this approach we can find a natural explanation for the fact that during the ascending parts of the global light curve, no 103-day modulation is seen: it only means that the decretion disc is not large enough to become tilted and warped and therefore it does not precess." + To further develop this model. one nav speculate hat the changes of projectlon o the warped disc to f1ο aue of sky. as seen from t16 optical comxuüon. slotld cause substautial variations with the precessional period. superposed on the optical helt curve.," To further develop this model, one may speculate that the changes of projection of the warped disc to the plane of sky, as seen from the optical companion, should cause substantial variations with the precessional period, superposed on the optical light curve." + Tus 19 actualv observed: we noted before he coincidence of t10 110yl time-scale of global optica varlafions wihn almost the sane period of precession whicjever the precessiug , This is actually observed: we noted before the coincidence of the $1400^\mathrm{d}$ time-scale of global optical variations with almost the same period of precession – whichever the precessing body. +One might expect that he N-raw activitv witn it system may also m some way be connected with t1ο pLaASCS of the 1EY! period., One might expect that the X-ray activity within the system may also in some way be connected with the phases of the $103^\mathrm{d}$ period. +" We tried o check whether such a correlatioj exists, and found that out of 18 docunented outbursts with amplitude =0.2 Crab (Cüovanurelli Ciraziati. 1992 and Fiuger et al. 1996))."," We tried to check whether such a correlation exists, and found that out of 18 documented outbursts with amplitude $\geq 0.2$ Crab (Giovannelli Graziati, \cite{GG92} and Finger et al. \cite{F96}) )," +" 7 peaked iu a Εν nuerval 0.60.7 of the optical phase curve (Ετος,", 7 peaked in a narrow interval $0.6-0.7$ of the optical phase curve (Figs. + and 5)). where phase 0.0 refers to optical minimμι (see Table 1)).," \ref{fig:4} + and \ref{fig:5}) ), where phase 0.0 refers to optical minimum (see Table \ref{tab:1}) )." + This means that N-rav outbursts “prefer” to inppou 1) to 20 davs after the optical maxima., This means that X-ray outbursts “prefer” to happen 10 to 20 days after the optical maxima. + If we consider this iu the context of the warped decretion disc uodel. this delay can be explained as the time needed or the matter captured from the De-star disc disturbed diving the close passage of the neutron star — to travel to the vicinity of the NS companion. where it loses its aneular monienutuui falls outo the NS. and causes an X-rav flare.," If we consider this in the context of the warped decretion disc model, this delay can be explained as the time needed for the matter captured from the Be-star disc – disturbed during the close passage of the neutron star – to travel to the vicinity of the NS companion, where it loses its angular momentum, falls onto the NS, and causes an X-ray flare." + It can be noted tiat during5 2001- two occasions of favourable X-ray aud optical phases will be in June and October: unluckilv. the former date corresponds to the seasonal gap of observations.," It can be noted that during 2001 two occasions of favourable X-ray and optical phases will be in June and October; unluckily, the former date corresponds to the seasonal gap of observations." +" An analvsis of the uniform photometric data set obtained in the period 19831998 has allowed a confident separation of the periodic coustitueut iu the Πο curve of the hieh-uass X-ray binary, À053526/V725 Tau."," An analysis of the uniform photometric data set obtained in the period 1983–1998 has allowed a confident separation of the periodic constituent in the light curve of the high-mass X-ray binary, A0535+26/V725 Tau." + The parameters of this periodic couponuent and its link with the phase of activity of the optical compoicut allow us to suggest of oxecessionaui accretioudisc around the neutron star ora warped equatorial ¢isc around a Be star as the most likely uechauisuis., The parameters of this periodic component and its link with the phase of activity of the optical component allow us to suggest precession of an accretion disc around the neutron star or a warped equatorial disc around a Be star as the most likely mechanisms. + At this poiut. both nodels seein to be viable: ucamwhile. the analysis ο already existiie spectral data could be helpful. iu the seuse that if the warxed disce node does reflect reality. oue might exyect fo see V/R aud EW variations correspouding to the precessioial motion.," At this point, both models seem to be viable; meanwhile, the analysis of already existing spectral data could be helpful, in the sense that if the warped disc model does reflect reality, one might expect to see $V/R$ and $EW$ variations corresponding to the precessional motion." + Within both models. we do not expect that substautia AX-aav outbursts to occur during ascending parts of the large-scale optical light curve.," Within both models, we do not expect that substantial X-ray outbursts to occur during ascending parts of the large-scale optical light curve." + Moreover. X-ray outbursts tend to oceur af specific phases of the 10515 optica helt curve.," Moreover, X-ray outbursts tend to occur at specific phases of the $102\fd 8$ optical light curve." +" Taken together. both effects cau explain the ""ndssune outburst” phenomenon."," Taken together, both effects can explain the “missing outburst” phenomenon." + Our results lead us to suppose that in other similar svstenis we cau hope to distinguish the ΠΟ plotometric, Our results lead us to suppose that in other similar systems we can hope to distinguish the minor photometric +"We usually find that the disc settles into a state of constant o, in which case a single value of a can be assigned to the disc.","We usually find that the disc settles into a state of constant $\alpha$, in which case a single value of $\alpha$ can be assigned to the disc." +" At first, the disc is almost axisymmetric with no source of heating."," At first, the disc is almost axisymmetric with no source of heating." + It will therefore start to cool at constant surface density towards Q=1., It will therefore start to cool at constant surface density towards $Q=1$. +" Since f is constant, the cooling time scale ἴεοοι is shortest in the inner parts of the disc, and this is where Q—1 is reached first."," Since $\beta$ is constant, the cooling time scale $t_\mathrm{cool}$ is shortest in the inner parts of the disc, and this is where $Q=1$ is reached first." +" Therefore, gravito-turbulence starts from the inside and moves outward as more and more of the disc cools towards Q—1."," Therefore, gravito-turbulence starts from the inside and moves outward as more and more of the disc cools towards $Q=1$." +" At this point, the initial temperature distribution has been washed out completely, and the temperature is now given by the requirement of Qz1 with the initial surface density."," At this point, the initial temperature distribution has been washed out completely, and the temperature is now given by the requirement of $Q \approx 1$ with the initial surface density." + It is easy to see that this gives H/Rος for our initial surface density profile., It is easy to see that this gives $H/R \propto R$ for our initial surface density profile. +" The number of grid cells per scale height is therefore an increasing function of R, which means that the outer parts of the disc are better resolved."," The number of grid cells per scale height is therefore an increasing function of $R$, which means that the outer parts of the disc are better resolved." +" The disc then starts to evolve viscously towards a steady-state accreting solution with constant values of Q and a, which has XοςΗ5/2 and constant temperature(2."," The disc then starts to evolve viscously towards a steady-state accreting solution with constant values of $Q$ and $\alpha$, which has $\Sigma \propto R^{-3/2}$ and constant temperature." +10)).. We choose 8=7.5 and evolve the discs for 2000 orbits at R=1., We choose $\beta=7.5$ and evolve the discs for 2000 orbits at $R=1$. +" According to and?,, these discs should not fragment for y=5/3, though previous results may be affected by resolution(?)."," According to and, these discs should not fragment for $\gamma=5/3$, though previous results may be affected by resolution." +. We increase the resolution to see if convergence can be reached., We increase the resolution to see if convergence can be reached. + The resulting surface density distributions are shown in Fig. 1.., The resulting surface density distributions are shown in Fig. \ref{figdens}. +" In our lowest resolution case (left panels), we resolve H outside R=3, but we find no evidence of fragmentation (defined such that the fragment is more than 2 orders of magnitude denser than the surroundings)."," In our lowest resolution case (left panels), we resolve $H$ outside $R\approx 3$, but we find no evidence of fragmentation (defined such that the fragment is more than 2 orders of magnitude denser than the surroundings)." +" In the final gravito-turbulent state (bottom left panel), we measure a total stress that compares very well with the predicted value of equation (3))."," In the final gravito-turbulent state (bottom left panel), we measure a total stress that compares very well with the predicted value of equation \ref{eqabc}) )." +" When increasing the resolution by a factor of 2 (middle panels of Fig. 1)),"," When increasing the resolution by a factor of 2 (middle panels of Fig. \ref{figdens}) )," + we do find fragmentation., we do find fragmentation. +" After 150 orbits at R=1, clumps start to form at the outer edge of the turbulent region (top middle panel in Fig. 1))."," After 150 orbits at $R=1$, clumps start to form at the outer edge of the turbulent region (top middle panel in Fig. \ref{figdens}) )." +" The first surviving clumps form atRz15, even though the necessary length scale (H) isresolved much further in."," The first surviving clumps form at$R\approx 15$, even though the necessary length scale $H$ ) isresolved much further in." +" When increasing the resolution even further (right panels of Fig. 1)),"," When increasing the resolution even further (right panels of Fig. \ref{figdens}) )," + the disc fragments further in (R=6)., the disc fragments further in $R\approx 6$ ). +" After 300 orbits at R=1, several of the fragments have merged or migrated off the computational domain?),, leaving five distinct clumps."," After 300 orbits at $R=1$, several of the fragments have merged or migrated off the computational domain, leaving five distinct clumps." +" We have observed fragmentation for values of 9 as high as 15, at a resolution that is 8 times the lowest resolution presented in this Letter."," We have observed fragmentation for values of $\beta$ as high as 15, at a resolution that is 8 times the lowest resolution presented in this Letter." +" Interestingly, for 6=15, clumps form in the outer disc, migrate inward and leave the computational domain."," Interestingly, for $\beta=15$, clumps form in the outer disc, migrate inward and leave the computational domain." +" After all clumps have left the disc, no further fragmentation is observed."," After all clumps have left the disc, no further fragmentation is observed." + This already suggests that the fragmentation for high ϐ is transient and related to the initial state of the disc., This already suggests that the fragmentation for high $\beta$ is transient and related to the initial state of the disc. +" We can obtain further insight into why these discs fragment by looking at radial profiles of Q, H/R and X."," We can obtain further insight into why these discs fragment by looking at radial profiles of $Q$, $H/R$ and $\Sigma$." + Figure 2 shows the results for the low resolution simulation with B—1.5., Figure \ref{figedge} shows the results for the low resolution simulation with $\beta=7.5$. +" From the top panel, we see that the disc cools down to a state of constant Q~1.5."," From the top panel, we see that the disc cools down to a state of constant $Q\approx 1.5$." +" Note that inside R=3, the scale height the disc tries to cool towards can not be resolved at this resolution, so we should only consider R>3."," Note that inside $R=3$, the scale height the disc tries to cool towards can not be resolved at this resolution, so we should only consider $R>3$." +" As mentioned before, the disc becomes turbulent from the inside out."," As mentioned before, the disc becomes turbulent from the inside out." +" The outer edge of the turbulent region can be clearly identified as a location where Q<1 in the top panel, and as bumps in the temperature and surface density profiles in the bottom two panels."," The outer edge of the turbulent region can be clearly identified as a location where $Q<1$ in the top panel, and as bumps in the temperature and surface density profiles in the bottom two panels." + Such a global feature makes the set-up no longer scale-free., Such a global feature makes the set-up no longer scale-free. + A region where Q<1 is of course unstable to gravitational instabilities., A region where $Q<1$ is of course unstable to gravitational instabilities. +" Moreover, sharp transitions in temperature and surface density are notoriously unstable??717)."," Moreover, sharp transitions in temperature and surface density are notoriously unstable." +". It is conceivable that the formation of clumps at the outer edge of the turbulent region, as observed for example in the top middle panel of Fig. 1,"," It is conceivable that the formation of clumps at the outer edge of the turbulent region, as observed for example in the top middle panel of Fig. \ref{figdens}," + is related to this sharp transition between the turbulent and laminar parts of the disc., is related to this sharp transition between the turbulent and laminar parts of the disc. +" 'To test to what extent the smooth initial conditions, and the related formation of this global feature, play a part in disc fragmentation, we employ the same strategy as?."," To test to what extent the smooth initial conditions, and the related formation of this global feature, play a part in disc fragmentation, we employ the same strategy as." +". We start off with a disc that has 6=30, and let gravito-turbulence fully develop over 1000 orbits at R=1."," We start off with a disc that has $\beta=30$, and let gravito-turbulence fully develop over 1000 orbits at $R=1$." +" For this value of 8, we have not observed fragmentation for any of the resolutions."," For this value of $\beta$, we have not observed fragmentation for any of the resolutions." +" Over the next 500 orbits, we then linearly decrease 8 to a value of 15,7.5,5,4 or 2.5."," Over the next 500 orbits, we then linearly decrease $\beta$ to a value of $15, 7.5, 5, 4$ or $2.5$." +" This way, we avoid any initial transients affecting the possible fragmentation of the disc, and restore the scale-free nature of the set-up."," This way, we avoid any initial transients affecting the possible fragmentation of the disc, and restore the scale-free nature of the set-up." + We then evolve the discs with £ held fixed at the desired values for another 2000 orbits (or until they fragment)., We then evolve the discs with $\beta$ held fixed at the desired values for another 2000 orbits (or until they fragment). +" In Fig. 3, "," In Fig. \ref{figdens2}, ," +we show the resulting surface density structures after 2000 orbits for the same three resolutions and the same final 6=7.5 as in Fig. 1.., we show the resulting surface density structures after 2000 orbits for the same three resolutions and the same final $\beta=7.5$ as in Fig. \ref{figdens}. . +" This time, we do"," This time, we do" +lutersteHar racliation fiek|UC=0.1.1.10.100.,"interstellar radiation field$\UMW=0.1,1,10,100$." + The parameter CuJiνι specifies the streugth of the interstellar radiation1 field iu uuüits of the radiation field of the Milky Wav a 10004: Aye=105 photons cu7 sol τον + (Draine1978:Mathisctal. 1983).. ," The parameter $\UMW=J/J_{\rm MW}$ specifies the strength of the interstellar radiation field in units of the radiation field of the Milky Way at : $J_{\rm MW}=10^6$ photons $^{-2}$ $^{-1}$ $^{-1}$ $^{-1}$ \citep{1978ApJS...36..595D, 1983A&A...128..212M}. ." +We continued oue of our snuulatious CZ/Z..=1. Up= 1) for additional OO Myr and found no siguificaut changes in the relation.," We continued one of our simulations $Z/Z_\odot=1$, $\UMW=1$ ) for additional 400 Myr and found no significant changes in the relation." + This iudicates that the predictions of our smmlatious should also hold for redshifts ;Xm3. at least unless/uutil ISAL properties change radically.," This indicates that the predictions of our simulations should also hold for redshifts $z\lesssim{}3$, at least unless/until ISM properties change radically." + αι total. we rau a set of 18 simulations Gucluding 2 ruus for a resolution study) iu order to explore the effect of varving metallicitv. radiation field aud deusity threshold ou the relation (see Table 1)).," In total, we ran a set of 18 simulations (including 2 runs for a resolution study) in order to explore the effect of varying metallicity, radiation field and density threshold on the relation (see Table \ref{tab:sims}) )." +" The 1nolecular livdrogen fraction fj, is coumputed sclconsistently. iucludiug a chemical network comprised of 6 species and radiative transter of the UV continui aud he LEyiiuu-Merner bands (Cuecdiun&Ixravtsov2011).."," The molecular hydrogen fraction $f_\H2{}$ is computed self-consistently, including a chemical network comprised of 6 species and radiative transfer of the UV continuum and the Lyman-Werner bands \citep{2010arXiv1004.0003G}." + If he average deusitv m a smnmlatiou cell is ualle than he density typical for molecular clouds. we have to revert o a uberid iuterpretation of the ITI» fraction.," If the average density in a simulation cell is smaller than the density typical for molecular clouds, we have to revert to a `subgrid' interpretation of the $\H2$ fraction." +" In this case. we asstie that the fraction fj, corresponds to the (miass) fraction «ft liverogen iu individual (unresolved) nolecular clouds."," In this case, we assume that the fraction $f_\H2$ corresponds to the (mass) fraction of hydrogen in individual (unresolved) molecular clouds." + Cousequently. the star formation timescale is giveu ΠΠ of (1) the free-fall time corresponding o the average deusity iu the cell and (2) the free-fall ine correspoudirie to the niunuuni density of molecular clouds that form STATS fie. Le. We stress that for densities smaller than ο the relation between SER and II abundance is linear. while it becomes non-linear for larger dqsities. because Te)NHg1/2 “," Consequently, the star formation timescale is given by the minimum of (1) the free-fall time corresponding to the average density in the cell and (2) the free-fall time corresponding to the minimum density of molecular clouds that form stars $\nc$, i.e., We stress that for densities smaller than $\nc$ the relation between SFR and $\H2$ abundance is linear, while it becomes non-linear for larger densities, because $\tau_\mathrm{ff}(n_{\rm H})\propto{}n_{\rm H}^{-1/2}$." +LA j0n-liucar steepenius of the relation at Xi>100AZ.. 7 is motivated by heoretical studies (0.9...5 Ixumholzetal. 20095)). but rot vet confirms by obscqvatious.," A non-linear steepening of the relation at $\Sigma_\H2>100 M_\odot$ $^{-2}$ is motivated by theoretical studies (e.g., \citealt{2009ApJ...699..850K}) ), but not yet confirmed by observations." +" We therefore explore he case in which He—50 5 7.Le. close to the typical average density o molecular clouds (~100200 2). mt also discuss t16 possilylity of mach larger thresholds such as se=LO? ci3 and p,=10° em P."," We therefore explore the case in which $\nc=50$ $^{-3}$, i.e., close to the typical average density of molecular clouds $\sim{}100-200$ $^{-3}$ ), but also discuss the possibility of much larger thresholds such as $\nc=10^3$ $^{-3}$ and $\nc=10^6$ $^{-3}$ ." + Since our sumuulations do ixit capture densities of 2107 cni7. a lhyeshold above lis value effectively correspouds to a tly linear SER - II» relation ou small scales.," Since our simulations do not capture densities of $\gtrsim{}10^5$ $^{-3}$, a threshold above this value effectively corresponds to a fully linear SFR - $\H2$ relation on small scales." + lustautaneous SFRs are computed directly using equations (1)) aid (2))., Instantaneous SFRs are computed directly using equations \ref{eq:SFR}) ) and \ref{eq:tau}) ). + Our simmlationus use spp/= 1005., Our simulations use $\epsilon_\mathrm{SFR}=0.005$ . + This vali which is consistent with small-scale observatious (Ixuuholz&Tan2007).. eusures that the jorimalization of he relation ou kpe scales is similar in simulations and observations.," This value, which is consistent with small-scale observations \citep{2007ApJ...654..304K}, ensures that the normalization of the relation on kpc scales is similar in simulations and observations." + Tinc-averaged SFRs (over tine T) are calculated by counting the iuuber of stars in a cube of given scale with ages below T., Time-averaged SFRs (over time $T$ ) are calculated by counting the number of stars in a cube of given scale with ages below $T$. + Uuless otherwise noted. we use £=20 My. but we wave explicitly checked that our results do not change sienificautly if larger averaging times are used (up to JI-—200 My).," Unless otherwise noted, we use $T=20$ Myr, but we have explicitly checked that our results do not change significantly if larger averaging times are used (up to $T=200$ Myr)." +" SER estimates based on observations of UV huninositices iu the wavelength range 1250-2800 correspond to an averaging time of ~ 100 Myr. which based on uchulay emission lines. such as 776. typically correspoud to Z7~10 Myr. and estimates based on the FIR οσοπα (c.g.. 21 nui) correspond to a range (~ 10 - 100 My) of averaging time scales: sec. οσοι,1998a)."," SFR estimates based on observations of UV luminosities in the wavelength range 1250-2800 correspond to an averaging time of $\sim$ 100 Myr, which based on nebular emission lines, such as $H\alpha$, typically correspond to $T\sim{}10$ Myr, and estimates based on the FIR continuum (e.g., 24 $\mu{}$ m) correspond to a range $\sim{}$ 10 - 100 Myr) of averaging time scales; see, e.g.,." +". The slope aud intercept of the relation are obtained with a bisector reeression in log-logVy, space (Isobeetal.1990)..", The slope and intercept of the relation are obtained with a bisector regression in log-log space \citep{1990ApJ...364..104I}. + Although tie use of bisector regression cannot be rigorously justified iu general (see kelly2007:Ίοσοaoetal.2010)). the bisector method is sufficient Kkyr our purposes as we porkqin regression on tightly correlated data without error bars.," Although the use of bisector regression cannot be rigorously justified in general (see \citealt{2007ApJ...665.1489K, 2010arXiv1008.4686H}) ), the bisector method is sufficient for our purposes as we perform regression on tightly correlated data without error bars." + We estimate scatter about the best-fit relation. as the root mean square Of loglO of the spatially averaged star formation rate density relative to its value ou the regression line with the sale density. see also equation (B1)).," We estimate scatter about the best-fit relation, as the root mean square of log10 of the spatially averaged star formation rate density relative to its value on the regression line with the same density, see also equation \ref{eq:sigSFR}) )." + We estimate errors fcx the slope. intercept. aud scatter using the standard bootstrap imiethod (Efron1979) with a sample size of 2UU.," We estimate errors for the slope, intercept, and scatter using the standard bootstrap method \citep{Efron1979} with a sample size of 200." + In Fig., In Fig. + 1 woe plot aud conipare the relation or (1) solar metallicitv aud yyy=04. and (2) Z/Z.=O11 and (yyy=100.," \ref{fig:KSrelation} we plot and compare the relation for (1) solar metallicity and $\UMW=0.1$, and (2) $Z/Z_\odot=0.1$ and $\UMW=100$." + Measured over the range 10<3c. of which Cowie et al. (," Of the 31 candidates, 24 are detected at $> 3\sigma$, of which Cowie et al. (" +2009) classified one as an AGN and 17 as galaxies.,2009) classified one as an AGN and 17 as galaxies. + In COSMOS. both a deep and a wide survey has been published in the MIPS band.," In COSMOS, both a deep and a wide survey has been published in the MIPS band." + The deep survey covers ~33% of the field surveyed for Lyc. and the sensitivities reached in the two surveys were 7] and 150 yy. respectively (5c).," The deep survey covers $\sim 33$ of the field surveyed for $\alpha$, and the sensitivities reached in the two surveys were 71 and 150 $\mu$ Jy, respectively $\sigma$ )." + To find counterparts to the Lya emitters. the public catalogues were," To find counterparts to the $\alpha$ emitters, the public catalogues were" +background region matched recent ISO observations of the zodiacal light. giving us confidence in our subtracted fIuxes(Reachetal.1996)..,"background region matched recent ISO observations of the zodiacal light, giving us confidence in our subtracted fluxes\citep{Ref G}." + Furthermore the profile measurements will nol be affected as no background is subtracted. and. in fact. little to no increase in [lux over is seen in (he profiles of the background region.," Furthermore the profile measurements will not be affected as no background is subtracted, and, in fact, little to no increase in flux over is seen in the profiles of the background region." + The CVF data were processed independent of the CAM because the exposure times are significantly shorter. filters much narrower and fIux levels reduced from the narrow band passes.," The CVF data were processed independent of the CAM because the exposure times are significantly shorter, filters much narrower and flux levels reduced from the narrow band passes." + These conditions [force different methods to be used (o eliminate (he instruments electronic signature., These conditions force different methods to be used to eliminate the instrument's electronic signature. + Data reduction followed the standard CVF steps: alter dark current subtraction we cleaned the data cube of elitehes using both the CLA multiresolution median transform aleorithim and a manual deglitching method by inspecting the cube [frame bv frame., Data reduction followed the standard CVF steps: after dark current subtraction we cleaned the data cube of glitches using both the CIA multiresolution median transform algorithm and a manual deglitching method by inspecting the cube frame by frame. + As with the narrow band filter data. the Fouks-Schubert model for stabilization was used and a jitter correction was applied to the data.," As with the narrow band filter data, the Fouks-Schubert model for stabilization was used and a jitter correction was applied to the data." + We devoted effort to producing an appropriate and accurate flat. field ancl background subtraction procedure., We devoted effort to producing an appropriate and accurate flat field and background subtraction procedure. + Sav light from extended sources may. affect the flat-field quality of the reduced images., Stray light from extended sources may affect the flat-field quality of the reduced images. + The standard flat-field procedure on CVF images produces a deviation [rom flatness. around on average. which limits the ability to image [aint extended sources and to perform accurate photometry.," The standard flat-field procedure on CVF images produces a deviation from flatness, around on average, which limits the ability to image faint extended sources and to perform accurate photometry." + The characteristic double-Ilobe CWF flat field pattern. present in the data after this flat field process. can be removed by subtracting an appropriate zodiacal light cube and by using a more appropriate detector Hat field.," The characteristic double-lobe CVF flat field pattern, present in the data after this flat field process, can be removed by subtracting an appropriate zodiacal light cube and by using a more appropriate detector flat field." + ILowever. since neither cube was available from (he standard distribution of calibration files at time of reductions. we produced our own.," However, since neither cube was available from the standard distribution of calibration files at time of reductions, we produced our own." + We constructed a zodiacal light eube lor both sectors of the lw CVF by interpolation. starting from 8 zodiacal images taken with (he CVF at fixed wavelengths within the wavelength range of our measurements.," We constructed a zodiacal light cube for both sectors of the lw CVF by interpolation, starting from 8 zodiacal images taken with the CVF at fixed wavelengths within the wavelength range of our measurements." + We were able to remove the zodiacal light bv scaling (his cube to the intensity measured off-source in our observations. and subtracting the scaled cube from. our measurements.," We were able to remove the zodiacal light by scaling this cube to the intensity measured off-source in our observations, and subtracting the scaled cube from our measurements." + Flat fields were produced in a similar manner., Flat fields were produced in a similar manner. + From (he data archive we recovered all the detector flats which were taken with the FOV through narrow band filters., From the data archive we recovered all the detector flats which were taken with the FOV through narrow band filters. + These flats were interpolated to match the wavelengths in our data set. and each data frame was divided by the appropriate flat.," These flats were interpolated to match the wavelengths in our data set, and each data frame was divided by the appropriate flat." + The final calibrated. data for NGC 1404. the only galaxy observed by ΟΝΕ in our program. is reported as the background subtracted fux within within radius in Table 4..," The final calibrated data for NGC 1404, the only galaxy observed by CVF in our program, is reported as the background subtracted flux within within radius in Table \ref{fluxes-cvf}. ." +might play a crucial role in explaining these systems’ apparent under-luminosities.,might play a crucial role in explaining these systems' apparent under-luminosities. +" A possible explanation for the under-luminosity of the HR 8799 planets has recently been proposed by ? and ?,, who demonstrate that a thick cloud atmosphere can reproduce the observed photometry of the HR 8799 planets."," A possible explanation for the under-luminosity of the HR 8799 planets has recently been proposed by \citet{2011arXiv1101.1973C} and \citet{2011arXiv1102.5089M}, who demonstrate that a thick cloud atmosphere can reproduce the observed photometry of the HR 8799 planets." +" The thick cloud models might be applicable to 2MASS 1207 b as well, and we investigate this possibility in Section ??.."," The thick cloud models might be applicable to 2MASS 1207 b as well, and we investigate this possibility in Section \ref{atmosphere}." +" The problems in invoking a near edge-on disk to explain the under-luminosity of 2MASS 1207 b stem from the fact that near edge-on and edge-on disks tend to produce more variability than is detected for 2MASS 1207 b, along with the fact that such a configuration is unlikely and cannot be used to explain other under-luminous brown dwarfs."," The problems in invoking a near edge-on disk to explain the under-luminosity of 2MASS 1207 b stem from the fact that near edge-on and edge-on disks tend to produce more variability than is detected for 2MASS 1207 b, along with the fact that such a configuration is unlikely and cannot be used to explain other under-luminous brown dwarfs." +" In this section, we model the system with a shell of dust, which should provide the extinction of a near edge-on disk, without the geometric effects that we have previously deemed unlikely."," In this section, we model the system with a shell of dust, which should provide the extinction of a near edge-on disk, without the geometric effects that we have previously deemed unlikely." +" Such a shell may have been discovered around the brown dwarf, G 196-3 B (?).."," Such a shell may have been discovered around the brown dwarf, G 196-3 B \citep{2010ApJ...715.1408Z}." +" We use the dust shell radiative transfer code, YPIT] (?),, to model dust shells at a variety of distances from the central brown dwarf."," We use the dust shell radiative transfer code, $^{,}$ \citep{1999astro.ph.10475I}, to model dust shells at a variety of distances from the central brown dwarf." +" In all cases, the central source is the DUSTY'?] brown dwarf model atmosphere (7.5—1,600 K, R5 —0.16 R and log(g)=4.5) used as a model of 2MASS 1207 b in Section B.1|."," In all cases, the central source is the ] brown dwarf model atmosphere $T_{\rm eff}$ =1,600 K, $R_{\rm b}$ =0.16 $R_{\rm \odot}$ and $log(g)$ =4.5) used as a model of 2MASS 1207 b in Section \ref{RADMC Disk Models}." +" We also use the same dust grain size distribution and optical properties as described in 3.1], with a minimum grain-size of @yin=1um (although varying this does not significantly change our results)."," We also use the same dust grain size distribution and optical properties as described in \ref{RADMC Disk Models}, with a minimum grain-size of $a_{min}$ $\micron$ (although varying this does not significantly change our results)." +" For each shell, we fix the J-band extinction to be 2.5 magnitudes (the observed under-luminosity of 2MASS 1207 b)."," For each shell, we fix the J-band extinction to be 2.5 magnitudes (the observed under-luminosity of 2MASS 1207 b)." +" The shells are geometrically thin, and placed at Ry, 5Ry, 10R,, 20R,, 40R,, and 80R,."," The shells are geometrically thin, and placed at $R_{\rm b}$, $R_{\rm b}$, $R_{\rm b}$, $R_{\rm b}$, $R_{\rm b}$, and $R_{\rm b}$ ." +" Since DUSTY is a 1-D code, we are forced to choose a slab geometry (which we use for the FE, model), and a shell geometry (which we use for all others), where the central source is assumed to be a point-source."," Since DUSTY is a 1-D code, we are forced to choose a slab geometry (which we use for the $R_{\rm b}$ model), and a shell geometry (which we use for all others), where the central source is assumed to be a point-source." +" The difference between using a slab and point-like approximation is small compared to the bulk SED shape,so we consider the point-like approximation reasonable"," The difference between using a slab and point-like approximation is small compared to the bulk SED shape,so we consider the point-like approximation reasonable" + Henryetal.(2000).. Fortney&Nettelmann(2010) (see.e.g..Perna Mandushevetal.(2007).. Torresetal.(2008)," \citet{2000ApJ...529L..45C} \citet{2000ApJ...529L..41H}, \citet{2010SSRv..152..423F} \citep[see, e.g.,][]{2010ApJ...724..313P, 2010ApJ...714L.238B}. \citet{2007ApJ...667L.195M}," + Sozzettietal.(2009).., \citet{2008ApJ...677.1324T} \citet{2009ApJ...691.1145S}. + Here. we present 5 new light curves for this system.," Here, we present 5 new light curves for this system." + We also present 2 new light curves for WASP-12. a planet that was discovered by Hebbetal.(2009) and for which occultation photometry has been used to characterize the planet’s atmosphere and orbit (Campoetal.husudhanetal. 2011)..," We also present 2 new light curves for WASP-12, a planet that was discovered by \citet{2009ApJ...693.1920H} and for which occultation photometry has been used to characterize the planet's atmosphere and orbit \citep{2010arXiv1003.2763C,2011Natur.469...64M}." +" Our third target. HAT-P-3 (Torres 2007).. is 1n the opposite category of planets that are ""too small.”"," Our third target, HAT-P-3 \citep{2007ApJ...666L.121T}, is in the opposite category of planets that are “too small.”" + Gibsonetal.(2010) have published 7 high-quality light curves of the system., \citet{2010MNRAS.401.1917G} have published 7 high-quality light curves of the system. + We present six new light curves. and provide independent estimates of the planetary and stellar parameters.," We present six new light curves, and provide independent estimates of the planetary and stellar parameters." + Almost all the observations were conducted at the Fred Lawrence Whipple Observatory (FLWO) located on HHopkins. Arizona. using the 1.2m telescope and KeplerCam detector.," Almost all the observations were conducted at the Fred Lawrence Whipple Observatory (FLWO) located on Hopkins, Arizona, using the 1.2m telescope and KeplerCam detector." + The KeplerCam is a 40967 CCD with a field of view of 23/1«23/1., The KeplerCam is a $4096^2$ CCD with a field of view of $23\farcm1 \times 23\farcm1$. + The pixels were binned 2« on the chip for faster readout., The pixels were binned $2 \times 2$ on the chip for faster readout. + The binned pixels subtend 0768 on a side., The binned pixels subtend $0\farcs68$ on a side. + Observations were made through Sloan / and z filters., Observations were made through Sloan $i$ and $z$ filters. + One of the WASP-12 transits was observed with the Nordic Optical Telescope (NOT) located in the Canary Islands. using the ALFOSC detector.," One of the WASP-12 transits was observed with the Nordic Optical Telescope (NOT) located in the Canary Islands, using the ALFOSC detector." + The ALFOSC detector is a 20487 CCD with a field of view of 6/4«614. corresponding to 0719 per pixel.," The ALFOSC detector is a $2048^2$ CCD with a field of view of $6\farcm4 \times 6\farcm4$, corresponding to $0\farcs19$ per pixel." + The observation was made through a Johnson V filter., The observation was made through a Johnson $V$ filter. + On each night we attempted to observe the entire transit. with at least an hour before ingress and an hour after egress. but the weather did not always cooperate.," On each night we attempted to observe the entire transit, with at least an hour before ingress and an hour after egress, but the weather did not always cooperate." + We performed overscan correction. trimming. bias subtraction and flat-field division withIRAF?.," We performed overscan correction, trimming, bias subtraction and flat-field division with." +. To generate— the light curves. we performed aperture photometry on the target star and all the comparison stars with similar brightnesses to the target star (within about afactorof two).," To generate the light curves, we performed aperture photometry on the target star and all the comparison stars with similar brightnesses to the target star (within about afactorof two)." + We tried many different choices for the photometric aperture, We tried many different choices for the photometric aperture +"(DZ;—&kD;—)/o(By) less than dp""1.0. 2.0. at IN. J and L are (22424. 2832 2030.}.. 22588,bat) MN 2055. 90. 2867. 2047.","$|(DZ_i- \kappa B_i-DZ(0))/\sigma(B_i)|$ less than 0.5, 1.0, 2.0, at K, J and L are 2424, 2832, 2930, 2588, 2873, 2955, and 2590, 2867, 2947,." +29055}.. ForBD with a -- distribution. one would expect these numbers to be. (11138.We 2028.have 2835.20631: more spread out in the center. with fewer pixels in the extended tail.," For comparison, with a Gaussian distribution, one would expect these numbers to be, 1138, 2028, 2835,; more spread out in the center, with fewer pixels in the extended tail." + usecl a robust. least sum of absolute values fitting method. which. by its nature. is insensitive to the few pixels with large residual/o.," We have used a robust, least sum of absolute values fitting method, which, by it's nature, is insensitive to the few pixels with large $\sigma$." + Thus the fit lor the slopes. &. and intercepts. DZ(Q). is {ο the pixels in the narrow core. aud (he outliers have little effect on their derived values and the final CIRB values.," Thus the fit for the slopes, $\kappa$, and intercepts, DZ(0), is to the pixels in the narrow core, and the outliers have little effect on their derived values and the final CIRB values." + Figures 7.. 8. and 9 show the derived intercepts. DZ(0). vs. ecliptic latitude in the three bands.," Figures \ref{kecl}, \ref{jecl}, and \ref{lecl} show the derived intercepts, DZ(0), vs. ecliptic latitude in the three bands." + We see here (he same trends with eclipüe latitude as were apparent in Wright (2001)., We see here the same trends with ecliptic latitude as were apparent in \citet{wrj01}. +. While the Ix-band intercepts appear reasonably independent of ecliptic latitude. there is a strong trend in J anda slight negative trend in the L-band.," While the K-band intercepts appear reasonably independent of ecliptic latitude, there is a strong trend in J and a slight negative trend in the L-band." + The zodiacal light is fainter at 3.5 pom than at 2.2 jan. and so the stronger dependence on 2at L than al IX may seem surprising.," The zodiacal light is fainter at 3.5 $\mu$ m than at 2.2 $\mu$ m, and so the stronger dependence on $\beta$at L than at K may seem surprising." + Llowever. al 3.5 jan. we begin to see thermal emission [rom (he interplanetary dust. in addition to (he scattered sunlight.," However, at 3.5 $\mu$ m, we begin to see thermal emission from the interplanetary dust, in addition to the scattered sunlight." + From a modeling standpoint. (his gives another free. parameter which should provide a better fit. but. from a physical standpoint. we are likely seeing a more complicated emission/scattering pattern on the sky which is more clifficult to model correctly.," From a modeling standpoint, this gives another free parameter which should provide a better fit, but from a physical standpoint, we are likely seeing a more complicated emission/scattering pattern on the sky which is more difficult to model correctly." + There is an overall scaling [actor at each wavelength in (he zodiacal light model. but the parameters that determine the physical shape of the dust cloud. were fit simultaneously to observations at 8 DIRBE bands.," There is an overall scaling factor at each wavelength in the zodiacal light model, but the parameters that determine the physical shape of the dust cloud were fit simultaneously to observations at 8 DIRBE bands." + The trend. with ecliptic alitude at both J and L indicates a problem with the modeled shape of the cloud., The trend with ecliptic latitude at both J and L indicates a problem with the modeled shape of the cloud. + The better fit to the scattered sunlight at IX could be a coincidence. rather than better modeling in that band.," The better fit to the scattered sunlight at K could be a coincidence, rather than better modeling in that band." + This remaining solar svstem dependence in (wo of the three analvzed bands is evidence of a problem with (he zodiacal light model that still prevents us [rom claiming a detection at 1.25 yam. The model may be improved by requiring (hat these DIRBE minus 2NLASS intercepts be ecliptie independent simultaneously in all three bands., This remaining solar system dependence in two of the three analyzed bands is evidence of a problem with the zodiacal light model that still prevents us from claiming a detection at 1.25 $\mu$ m. The model may be improved by requiring that these DIRBE minus 2MASS intercepts be ecliptic independent simultaneously in all three bands. + Improvements to the zodiacal light model will be addressed in future work., Improvements to the zodiacal light model will be addressed in future work. + The results of this DIRBE minus 2MASS subtraction in these 40 regionsof the skv give astatistically significant isotropic background of 14.69 + 4.49 kJv fat 2.2 pim and 15.62 +3.34kJvsv tat 3.5 yan where the uncertainty has not been significantly reduced since the dominant sources of error are svstematic., The results of this DIRBE minus 2MASS subtraction in these 40 regionsof the sky give astatistically significant isotropic background of 14.69 $\pm$ 4.49 kJy $^{-1}$ at 2.2 $\mu$ m and 15.62 $\pm$ 3.34 kJy $^{-1}$ at 3.5 $\mu$ m where the uncertainty has not been significantly reduced since the dominant sources of error are systematic. + These results are consistent with earlier results. summarized in Table 4... including Gorjianοἱal. (2000).. Wright&Reese (2000).. Wright and with the 13 similarly analvzed regions from Wright&Johnson(2001) all of," These results are consistent with earlier results, summarized in Table \ref{oldvals} , including \citet{gor00}, , \cite{wrr00}, , \citet{elw01} and with the 13 similarly analyzed regions from \citet{wrj01} all of" +quantitatively unclear.,quantitatively unclear. + Let's discuss however the predicted QDO [requencies., Let's discuss however the predicted QPO frequencies. + The first (vo strongest resonances for each azimuthal mode mi with the plus sign used in Eq. (3)), The first two strongest resonances for each azimuthal mode $m$ with the plus sign used in Eq. \ref{eq:ResPara2}) ) + match with good accuracy the frequencies observed in GRS 19154-105., match with good accuracy the frequencies observed in GRS 1915+105. +. We find namely. in decreasing order. 163.0. 112.0. 56.0. 42.0 and 28.0 Hz for the most significant values.," We find namely, in decreasing order, 168.0, 112.0, 56.0, 42.0 and 28.0 Hz for the most significant values." + Some other resonances should also appear. but due to their higher order (higher integer 7) ancl location farther away from (he inner edge of the disk (thus an amplitude of the excitation decreasing with radius). they do not possess a significant growth rate relevant for the study presented here.," Some other resonances should also appear, but due to their higher order (higher integer $n$ ) and location farther away from the inner edge of the disk (thus an amplitude of the excitation decreasing with radius), they do not possess a significant growth rate relevant for the study presented here." + Moreover. X-ray emission from the outer part of the disk is fainter. thus more difficult to detect.," Moreover, X-ray emission from the outer part of the disk is fainter, thus more difficult to detect." + As a consequence. (he parametric resonance induced by a spiral density wave passing through an accretion disk predicts the HF-QPOs as well as some LF-QPOs (as a byproduct) in a single unilied picture.," As a consequence, the parametric resonance induced by a spiral density wave passing through an accretion disk predicts the HF-QPOs as well as some LF-QPOs (as a byproduct) in a single unified picture." + However. there is no wav (o predict the angular momentum of the black hole without some knowledge of its mass.," However, there is no way to predict the angular momentum of the black hole without some knowledge of its mass." + We can derived a relation mass-spin to constrain the DII parameters. see next section.," We can derived a relation mass-spin to constrain the BH parameters, see next section." +" Nevertheless. it is worthwhile noting that Kato(2004) was lead to the conclusion that (his microquasar is well described by a non rotating geometry indicating that the black hole angular momentum is weak. apy24.5\,{\rm mag/arcsec^2}$ leads to no difference in the clustering signal." + However eliminating objects with maga«22mag/arcesec? reduces the amplitude of clustering at large scales by large factors even though they represents a small percentage of the total sample.," However eliminating objects with $mag_{50}<22\,{\rm mag/arcsec^2}$ reduces the amplitude of clustering at large scales by large factors even though they represents a small percentage of the total sample." + Hence. we next discuss the motivation for this cut in more detail (further evidence for this cllect is given in Sec. 3.2)).," Hence, we next discuss the motivation for this cut in more detail (further evidence for this effect is given in Sec. \ref{sec:starcontamination}) )." + The distribution of meagso is given in Fig. L1. , The distribution of $mag_{50}$ is given in Fig. \ref{fig:histomu50}. . +lt is well concentrated: around. megzo~23nmag/aresec? but. shows long tails due to objects contaminating the LRG sample.," It is well concentrated around $mag_{50}\sim 23\,{\rm mag/arcsec^2}$ but shows long tails due to objects contaminating the LRG sample." + This contamination is more clearly depicted in the petror vs. mrtgso diagram in Fig. 2.., This contamination is more clearly depicted in the $petror$ vs. $mag_{50}$ diagram in Fig. \ref{fig:histomu501}. +" Top panel corresponds. to our photometric sample and shows a cilferent trend. for mag,«22mag/arcsec? and meagso24.5mag/zwesec. with the core of LRGs [lying in between."," Top panel corresponds to our photometric sample and shows a different trend for $mag_{50}<22\,{\rm mag/arcsec^2}$ and $mag_{50}>24.5\,{\rm mag/arcsec^2}$, with the core of LRGs lying in between." + Bottom panel shows the same ciagram but for the SDSS DRT spectroscopic sample. after imposing the selection in I5q. (12).," Bottom panel shows the same diagram but for the SDSS DR7 spectroscopic sample, after imposing the selection in Eq. \ref{eq:selection1}) )." + This panel nicely shows that all objects with πα<22mag/arcesec could be contaminated. by stars.," This panel nicely shows that all objects with $mag_{50}<22\,{\rm mag/arcsec^2}$ could be contaminated by stars." + In addition. Fig.," In addition, Fig." + 3. shows a histogram of number. of objects per. pixel. (here the pixel size is O.0ldeg?) asa function. of galactic latitude and dillerent. recshilt bins. including (solid) or excluding (dashed) galaxies with low meagsu.," \ref{fig:galacticlatitude} shows a histogram of number of objects per pixel (here the pixel size is $0.01 {\rm deg}^2$ ) asa function of galactic latitude and different redshift bins, including (solid) or excluding (dashed) galaxies with low $mag_{50}$ ." + Objects with low memgso clearly concentrate at. low galactic latitudes introducing, Objects with low $mag_{50}$ clearly concentrate at low galactic latitudes introducing +phase.,phase. + In the asvinplolic limit of that phase. the initial mass of (he plasmold becomes negligible.," In the asymptotic limit of that phase, the initial mass of the plasmoid becomes negligible." + Although we have properly included the effect of radiative cooling ancl radiation drag. in most cases (he fraction of swept-up energy which is translerred to ultrarelativistic electrons and can therefore be racdiated away elliciently. will be small.," Although we have properly included the effect of radiative cooling and radiation drag, in most cases the fraction of swept-up energy which is transferred to ultrarelativistic electrons and can therefore be radiated away efficiently, will be small." + Therefore. we can approximate the equation of motion of the plasmoid by. an acdiabatie solution. as and where we have. for simplicity. assumed a constant external density p(r)=pos const.," Therefore, we can approximate the equation of motion of the plasmoid by an adiabatic solution, as and where we have, for simplicity, assumed a constant external density $\rho(r) \equiv \rho_{\rm ext} =$ const." +" This svstem has a sell-similar solution of the form Assuming for the purpose of an analvtie estimate that we are looking right down thejel (Bun,= 0). Che observers time as a function of Lorentz [actor can be expressed as The observed steep spectral index of the optical svnchrotron emission from ος 279 1.7). indicates p4.4."," This system has a self-similar solution of the form Assuming for the purpose of an analytic estimate that we are looking right down thejet $\theta_{\rm obs} = 0$ ), the observer's time as a function of Lorentz factor can be expressed as This yields a solution for the Lorentz factor as a function of observer's time: The observed steep spectral index of the optical synchrotron emission from 3C 279 $\alpha \sim 1.7$ ), indicates $p \sim 4.4$." +" This. in turn. signifies that the svstem is in the fast cooling regime since otherwise a cooling break would not produce à pF, peak at the svnchrotron Irequency corresponding (ο ο."," This, in turn, signifies that the system is in the fast cooling regime since otherwise a cooling break would not produce a $\nu F_{\nu}$ peak at the synchrotron frequency corresponding to $\gamma_b$." + Furthermore. electrons svuchrotron radiating al optical frequencies are most likely bevond the break energy. i.e... 5> 544.," Furthermore, electrons synchrotron radiating at optical frequencies are most likely beyond the break energy, i.e., $\gamma > \gamma_{\rm min}$ ." + For the prediction of svnchrotron lieht, For the prediction of synchrotron light +to provide estimates of the (internal) uncertainties. of the presented atomic data.,to provide estimates of the (internal) uncertainties of the presented atomic data. + In the case of PI. we also compared our cross sections to recent experimental measurements to gauge the accuracy of our results.," In the case of PI, we also compared our cross sections to recent experimental measurements \citep{esteves09, esteves10, sterling11} to gauge the accuracy of our results." + We find that the direct PI cross sections for these Se tons are uncertain by30-5067. while the RR rate coefficients are uncertain by «106€ except for Se (-30—409c)).," We find that the direct PI cross sections for these Se ions are uncertain by, while the RR rate coefficients are uncertain by $<$ except for $^+$ $\sim$ )." + Our DR rate coefficients exhibit larger uncertainties (from up to 1—2 orders of magnitude) due to the unknown energies of low-lying autoionizing states., Our DR rate coefficients exhibit larger uncertainties (from up to 1–2 orders of magnitude) due to the unknown energies of low-lying autoionizing states. + The importance of uncertainties in DR rate. coefficients. is highlighted by the dominance of DR over RR near 10+ Κ. the typical temperature of photoionized nebulae such as PNe.," The importance of uncertainties in DR rate coefficients is highlighted by the dominance of DR over RR near $^4$ K, the typical temperature of photoionized nebulae such as PNe." + These uncertainties are significant. particularly for the near-neutral cases.," These uncertainties are significant, particularly for the near-neutral cases." + Low-charge Se tons are complex systems for which no comprehensive photoionization or recombination data existed prior to our study., Low-charge Se ions are complex systems for which no comprehensive photoionization or recombination data existed prior to our study. + Thus one of our goals in. addition to producing these data was to provide realistic internal uncertainties. whose effect on nebular abundance determinations can be quantified through numerical simulations of tonized nebulae.," Thus one of our goals in addition to producing these data was to provide realistic internal uncertainties, whose effect on nebular abundance determinations can be quantified through numerical simulations of ionized nebulae." + Such an investigation will help to target the tonic systems and atomic processes that require further analysis., Such an investigation will help to target the ionic systems and atomic processes that require further analysis. + This paper ts the first in a series to present atomic data for the photoionization and recombination properties of trans-iron elements., This paper is the first in a series to present atomic data for the photoionization and recombination properties of trans-iron elements. + In future papers. we will present similar data for Kr and Xe. as well as charge exchange rate coefficients for low-charge ions of several --capture elements.," In future papers, we will present similar data for Kr and Xe, as well as charge exchange rate coefficients for low-charge ions of several -capture elements." + The ultimate goal of this study is to produce atomic data for --capture elements that are suitable to be incorporated into photoionization codes that numerically simulate the ionization and thermal structure of astrophysical nebulae., The ultimate goal of this study is to produce atomic data for -capture elements that are suitable to be incorporated into photoionization codes that numerically simulate the ionization and thermal structure of astrophysical nebulae. + Photoionization codes ean be used to derive analytical ionization corrections needed to accurately estimate the abundances of unobserved Se. Kr. and Xe tons in ionized nebulae — and hence to determine their elemental abundances.," Photoionization codes can be used to derive analytical ionization corrections needed to accurately estimate the abundances of unobserved Se, Kr, and Xe ions in ionized nebulae – and hence to determine their elemental abundances." + Our efforts will thus enable much more accurate and robust abundance determinations of trans-iron elements in astrophysical nebulae than was previously possible., Our efforts will thus enable much more accurate and robust abundance determinations of trans-iron elements in astrophysical nebulae than was previously possible. + SSterling acknowledges support from an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0901432 and from NASA grant 06-APRA206-0049., Sterling acknowledges support from an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0901432 and from NASA grant 06-APRA206-0049. + We are grateful to BBadnell for several helpful discussions regarding the use of AUTOSTRUCTURE. and for his careful reading of this manuscript.," We are grateful to Badnell for several helpful discussions regarding the use of AUTOSTRUCTURE, and for his careful reading of this manuscript." +as described in retsec:evolution..,as described in \\ref{sec:evolution}. + These stars will arrive on the IB spinuing critically. causing the peak at loeeyo:/Coir)7 0.," These stars will arrive on the HB spinning critically, causing the peak at $\log(v_\mathrm{rot}/v_\mathrm{crit}) \approx 0$ ." + Since WB stars do not expand caamatically during the core helium burning phase. many will be found at or close to this peak.," Since HB stars do not expand dramatically during the core helium burning phase, many will be found at or close to this peak." + When these stars continue their evolutiou on the AGB. they will be spinning less rapidly than they were on the RGB at a similar radius; because of angular 1ionmenutun loss during the iuterveniug evolution.," When these stars continue their evolution on the AGB, they will be spinning less rapidly than they were on the RGB at a similar radius, because of angular momentum loss during the intervening evolution." + Iu fact. of the three sub-populations. there exists a relatively sharp lower cutoff in rotational velocities at 0.1ee for both the TB aud RGB stars. whereas the AGB irstars show a sinmeoth decline.," In fact, of the three sub-populations, there exists a relatively sharp lower cutoff in rotational velocities at $0.1\,v_\mathrm{crit}$ for both the HB and RGB stars, whereas the AGB stars show a smooth decline." + For merged thatobjects that are ACD y.ars at the preseut epoch. we find were formed through a merger ou the RGB. while only actually iucreed on the AGB.," For merged objects that are AGB stars at the present epoch, we find that were formed through a merger on the RGB, while only actually merged on the AGB." +" To facilitate comparison with observations. Figure 2aa shows the same distributions as τοποντοD, erger.butwiththeecloeitgeonvertedtoaphgusicalequat"," To facilitate comparison with observations, Figure \ref{fig:vrot-kms_1D_merger}a a shows the same distributions as \\ref{fig:vrot_1D_merger}, but with the velocity converted to a physical equatorial velocity in $^{-1}$ ." +seen since (Callowayetal.2007.2008b).,"seen since \citep{gal07, gal08b}." +. There is sole evidence that the intermittent appearance of the aceretion-powered pulsations may be related to the occurrence of Type I A-vay bursts. but whether the relatiouship is causal is not clear (Gallowayetal.2007).," There is some evidence that the intermittent appearance of the accretion-powered pulsations may be related to the occurrence of Type I X-ray bursts, but whether the relationship is causal is not clear \citep{gal07}." +. From 2005-2008 several N-rav bursts froi this source were detected byHETE-IT. theFEvplorer (RATE) andSWIFT?.," From 2005-2008 several X-ray bursts from this source were detected by, the (RXTE) and." +. We searched the bursts detected by both the RATE Proportional Counter Array (PCA) and the SWIFT Dist Alert Telescope (BAT) for oscillations. but found no significant signal.," We searched the bursts detected by both the RXTE Proportional Counter Array (PCA) and the SWIFT Burst Alert Telescope (BAT) for oscillations, but found no significant signal." + Durus routine monitoring on April 2 2009 (08:57 UTC). RNTE detected another N-rav. burst with both he PCA and theEcperiment (IIENTE).," During routine monitoring on April 2 2009 (08:57 UTC), RXTE detected another X-ray burst with both the PCA and the (HEXTE)." + Timing analysis was conducted using 125 js ine resolution PCA event mode data from the two active PCUs (PCUs 0 and 2)., Timing analysis was conducted using 125 $\mu$ s time resolution PCA event mode data from the two active PCUs (PCUs 0 and 2). + We used all photous in the 2-30 seV range. the baud where burst cussion exceeds the ersistent level.," We used all photons in the 2-30 keV range, the band where burst emission exceeds the persistent level." + The data were barvceutered using the JPL DE105 aud spacecraft ephemerides. with the source xosition of Fox(2005).," The data were barycentered using the JPL DE405 and spacecraft ephemerides, with the source position of \citet{fox05}." +. Some event mode data overruuis occurred iu the burst peak. leading to short data gaps.," Some event mode data overruns occurred in the burst peak, leading to short data gaps." + A dynamical power spectruuu (Figure 1)) reveals stroug must oscillations curing the initial decay of the burst. drifting upwards by about 1 Uz.," A dynamical power spectrum (Figure \ref{dps}) ) reveals strong burst oscillations during the initial decay of the burst, drifting upwards by about 1 Hz." + Selecting only data after =ls in Fieure 1. (the time of the final data eap). we fiud πιαια Lealy power in the range 17-30 in 1 independent consecutive 2 s bins. an extremely robust detection.," Selecting only data after t=4 s in Figure \ref{dps} (the time of the final data gap), we find maximum Leahy power in the range 17-30 in 4 independent consecutive 2 s bins, an extremely robust detection." + The oscillations fall below the detectability lreshold when the frequency is 376.3 Tz. 1 ITz below the spin frequency.," The oscillations fall below the detectability threshold when the frequency is 376.3 Hz, 1 Hz below the spin frequency." + Although the binary cplemeris caunot © extended to 2009 (the errors are too large). orbital Doppler effects will shift the spinfrequency by 0.009 IIz at most (saaretetal.2006).. so the offset im the yequencies is secure.," Although the binary ephemeris cannot be extended to 2009 (the errors are too large), orbital Doppler effects will shift the spinfrequency by 0.009 Hz at most \citep{kaa06}, so the offset in the frequencies is secure." + A search for accretion-powered oilsatious duriug the observation revealed no siguificaut signals., A search for accretion-powered pulsations during the observation revealed no significant signals. + To study the pulse profile. we extracted data for he ὃν period (f = L[-12s in Fievure 1)) duriug which he oscillations are detected.," To study the pulse profile, we extracted data for the 8s period (t = 4-12s in Figure \ref{dps}) ) during which the oscillations are detected." + We fitted the frequency dift using a polynomial model to maximise the power., We fitted the frequency drift using a polynomial model to maximise the power. + Power was dasximised with a quadratic drift model. vem|moltty)volt f9Y.cwvdüch yielded a Lealiy jorimialised power of SS (the best near model gave a uaximuuni power of 758. and higher order polvuonuüal erms vielded uo further mniproveient).," Power was maximised with a quadratic drift model, $\nu = \nu_0 + +\dot{\nu}_0 (t-t_0) + \ddot{\nu}_0(t-t_0)^2$, which yielded a Leahy normalised power of 88 (the best linear model gave a maximum power of 78, and higher order polynomial terms yielded no further improvement)." + Using the best fit quadratic οποιος model we generated a folded xofile., Using the best fit quadratic frequency model we generated a folded profile. + The persistent (uou-burst) cussion level was estimated from the ~— 10008 prior to the burst., The persistent (non-burst) emission level was estimated from the $\sim 1000$ s prior to the burst. + The resulting pulse profile is shown iuset in Figure 1.., The resulting pulse profile is shown inset in Figure \ref{dps}. + The xilse is well fit with a simple simisoid (4? /dof 31/37). with no need for harmonic content.," The pulse is well fit with a simple sinusoid $\chi^2$ = 31/37), with no need for harmonic content." + The fractional= amplitude was (3.52:0.3) Root Mean Square (RMS) (2-30 keV). comparable to the fractional amplitude of the aceretion-powered pulsatious of this source.," The fractional amplitude was $3.5 \pm 0.3$ Root Mean Square (RMS) (2-30 keV), comparable to the fractional amplitude of the accretion-powered pulsations of this source." + The wuplitude of the pulsations rises with cnerey from (3.1 £0.5) RAIS in the 2-5 keV. baud to (5.5+O0.6)% RAMS in the 10-20 keV band.," The amplitude of the pulsations rises with energy, from $3.1 \pm 0.5$ ) RMS in the 2-5 keV band to $5.5\pm 0.6$ RMS in the 10-20 keV band." + The burst oscillations also show hard lags. with the 10-20 keV pulse lageing the 2-5 keV pulse by (0.09 = 0.03) eveles.," The burst oscillations also show hard lags, with the 10-20 keV pulse lagging the 2-5 keV pulse by (0.09 $\pm$ 0.03) cycles." + To study the characteristics of the burst containing he oscillations. we extracted spectra every 0.258 from he PCA Event data (E_1125us_66118). which las Gl energv chanucls between 2 and 60 keV. We usec data from PCUs 0 aud 2. which were operating at the iue of the N-rav burst.," To study the characteristics of the burst containing the oscillations, we extracted spectra every 0.25s from the PCA Event data 1s), which has 64 energy channels between 2 and 60 keV. We used data from PCUs 0 and 2, which were operating at the time of the X-ray burst." + We eeuerated an iustrument response nati for each spectrum. aud fitted if using NSPEC version 11.5.2.," We generated an instrument response matrix for each spectrum, and fitted it using XSPEC version 11.3.2." + We added a systematic error o the spectra aud restricted the spectral fits to the enerev range 2-25 keV. We extracted spectra using l100s seemients before and after the burst. aud used hei as background in our fits.," We added a systematic error to the spectra and restricted the spectral fits to the energy range 2-25 keV. We extracted spectra using 100s segments before and after the burst, and used them as background in our fits." + We fitted each of the 4258 spectra with an absorbed(wabs. Morrisou&AIc-Canunon 1983)) Dlackbody model model in NSPEC).," We fitted each of the 0.25s spectra with an absorbed, \citealt{mor83}) ) blackbody model model in XSPEC)." + This approach. which takes (bbodyradaccount of the contribution of the persistent aud. backeround cussion. is standard procedure in X-ray burst analysis (see for example I&uulkersetal. 2003)).," This approach, which takes account of the contribution of the persistent and background emission, is standard procedure in X-ray burst analysis (see for example \citealt{kuu03}) )." + The simple blackbody ft described above vields a peak uninositv of 3.6«1075 eyes s+. aud a total cucrey of τς10° cres.," The simple blackbody fit described above yields a peak luminosity of $3.6 \times 10^{38}$ ergs $^{-1}$, and a total energy of $2.7 \times 10^{39}$ ergs." +" The evolution ofthe temperature Z5, aud he radius A5, eiven by the spectral fits resembles those seen dn plotospheric radius expansion (PRE) bursts (Callowayetal.2008a).. e. G) Ry, reaches a local ΜΑΧΗ close to the time of the peal flax. (i) lower values of Rp, were measured following the flux maxi and (n) there is a local miuiumumu in {η at the same inue as the παπα in Ry."," The evolution of the temperature $T_{bb}$ and the radius $R_{bb}$ given by the spectral fits resembles those seen in photospheric radius expansion (PRE) bursts \citep{gal08a}, i.e. (i) $R_{bb}$ reaches a local maximum close to the time of the peak flux, (ii) lower values of $R_{bb}$ were measured following the flux maximum and (iii) there is a local minimum in $T_{bb}$ at the same time as the maximum in $R_{bb}$ ." + This conclusion (aud the, This conclusion (and the +This trend. reflects the hierarchical or nearly hierarchica spatial clistribution that the multiple’s components acdop for the svstem to attain dynamical stability note tha this trend is better visualised if the svmbols rather than the primary masses are considered.,This trend reflects the hierarchical or nearly hierarchical spatial distribution that the multiple's components adopt for the system to attain dynamical stability – note that this trend is better visualised if the symbols rather than the primary masses are considered. + The second trend. reflects a shortcoming of these simulations: binaries have à mean e of z]O020 AU. e. the vast majority are close binaries anc only 1 svstem has à e>50 AU.," The second trend reflects a shortcoming of these simulations: binaries have a mean $a$ of $\approx 10-20$ AU, i.e. the vast majority are close binaries and only 1 system has a $a > 50$ AU." + This result is independent of primary mass., This result is independent of primary mass. + In other words. wide binaries of any mass are underproduced by our simulations and although closer binaries might have been formed if the softening length use to smooth the gravitational force at short distances had been chosen smaller. no wider system can be produced with these initial conditions.," In other words, wide binaries of any mass are underproduced by our simulations and although closer binaries might have been formed if the softening length used to smooth the gravitational force at short distances had been chosen smaller, no wider system can be produced with these initial conditions." + The reason for this is (wololcl: first. there is not enoughglobal angular momentum initially to form many wide binaries or a very wide binary. ancl second. but. most importantly. the high density of the sub-clusters prevents wide systems from surviving.," The reason for this is twofold: first, there is not enough angular momentum initially to form many wide binaries or a very wide binary, and second but most importantly, the high density of the sub-clusters prevents wide systems from surviving." + Fhis conclusion can be applied to all the systems underprocluced by our simulations: Low mass. low mass ratio and wide binaries share the property of having low binding energies.," This conclusion can be applied to all the systems underproduced by our simulations: low mass, low mass ratio and wide binaries share the property of having low binding energies." + Only the systems with larec binding energies can survive the dynamical encounters taking place in the dense clusters formed. by collapsing turbulent [ows., Only the systems with large binding energies can survive the dynamical encounters taking place in the dense clusters formed by collapsing turbulent flows. + Figure 2. shows pictorial sketches of the spatial distribution of the components of a representative sample of the multiple svstems resulting from our simulations. ie. à binary star orbiting a triple (top left). a binary quadruple (a binary orbiting another binary) with someou/liers (bottom left). a binary orbiting a binary quacdruple plus someoulliers (top right). and aplancfary quadruple (bottom right).," Figure 2 shows pictorial sketches of the spatial distribution of the components of a representative sample of the multiple systems resulting from our simulations, i.e. a binary star orbiting a triple (top left), a binary quadruple (a binary orbiting another binary) with some (bottom left), a binary orbiting a binary quadruple plus some (top right), and a quadruple (bottom right)." + The first three classes of multiples comprise about of all the multiples formed in the simulations., The first three classes of multiples comprise about of all the multiples formed in the simulations. + Phe remaining (the fourth class) is comprised by multiples. svstems in which companions are not members of binary/triple systems other than the multiple itself:," The remaining (the fourth class) is comprised by multiples, systems in which companions are not members of binary/triple systems other than the multiple itself." + One of the noteworthy features of the mass distribution within the multiples is how similar he masses of all components are. either taking an individual star or a binary star as the funcamental unit.," One of the noteworthy features of the mass distribution within the multiples is how similar the masses of all components are, either taking an individual star or a binary star as the fundamental unit." + In other words. most binaries have very high q and. when two binaries are ind. to cach other. the ratio of the total mass of one to he total mass of the other is also close to unity (in the range Q.5r1: see left panel of Figure 1).," In other words, most binaries have very high $q$ and, when two binaries are bound to each other, the ratio of the total mass of one to the total mass of the other is also close to unity (in the range $0.5-1$; see left panel of Figure 1)." + This pattern of mass distribution— can be readily related to the hierarchical ormation mechanism that produces the multiples: at each sub-cluster merging event. the binary stars involved interact strongly. exchanging components. and. accreting material with higher specific angular momentum than that of their orbit.," This pattern of mass distribution can be readily related to the hierarchical formation mechanism that produces the multiples: at each sub-cluster merging event, the binary stars involved interact strongly, exchanging components, and accreting material with higher specific angular momentum than that of their orbit." + These two processes invariably favour the formation of bound. pairs (a pair of stars in the first small cluster. a pair a binaries after a merging event. ete.)," These two processes invariably favour the formation of bound pairs (a pair of stars in the first $N$ cluster, a pair a binaries after a merging event, etc.)" + with high values oq (2 0.5).," with high values of $q$ $> +0.5$ )." + Whereas the first star formation episodes. involve. as much direct fragmentation in filaments as disc fragmentation (BBB). at later times. star formation is dominated by instabilities in circum-hbinary/quacruple cises.," Whereas the first star formation episodes involve as much direct fragmentation in filaments as disc fragmentation (BBB), at later times, star formation is dominated by instabilities in circum-binary/quadruple discs." + Most objects ormed at late times are unlikely to remain bound. since rev start oll with a mass close to the opacity limit for ragmentation whereas the other sub-cluster members have en acercting for some time and have much larger masses. rverefore being better. positioned. for future. threc-body encounters.," Most objects formed at late times are unlikely to remain bound since they start off with a mass close to the opacity limit for fragmentation whereas the other sub-cluster members have been accreting for some time and have much larger masses, therefore being better positioned for future three-body encounters." + The net result is that only those objects formed curing the first [rec-Fall times after star formation begins are ikelv to remain in bound structures. whilst the rest will be ejected at. very. large separations (oulfiers) or simply escape rom the cloud.," The net result is that only those objects formed during the first free-fall times after star formation begins are likely to remain in bound structures, whilst the rest will be ejected at very large separations ) or simply escape from the cloud." + This is rellected in Table 1 in the penultimate column (Nau). which gives the number of objects that being bound o à given multiple system. have a mass much lower than he other components. and highly eccentric orbits.," This is reflected in Table 1 in the penultimate column $_{\rm +out}$ ), which gives the number of objects that being bound to a given multiple system, have a mass much lower than the other components, and highly eccentric orbits." + Outliers, Outliers +Alo. whose values are in the ratio qd. where1.,"$M_2$, whose values are in the ratio $q$, where." +. We assume. for simplicity. that the two are ou circular orbits about their mutual center of mass. with radii ay aud ao.," We assume, for simplicity, that the two are on circular orbits about their mutual center of mass, with radii $a_1$ and $a_2$ ." + For a given Aqu. q. audae.. the fixed separation between the stars. we want to ascertain the perturbing effect of the binary ou the surrounding cloud. taken to have a background density of pp.," For a given $M_{\rm tot}$, $q$, and, the fixed separation between the stars, we want to ascertain the perturbing effect of the binary on the surrounding cloud, taken to have a background density of $\rho_0$." + Iu the linear approximation valid [or small clisturbauces. momeutuim couservation reads Here we lave denoted the simall tucucecl velocity by wy. aud have used the same subscript [or perturbatious in the density aud potential.," In the linear approximation valid for small disturbances, momentum conservation reads Here we have denoted the small induced velocity by $\bu_1$, and have used the same subscript for perturbations in the density and potential." + The superscript on Φ{ emphasizes that this perturbation is due to the stars alone., The superscript on $\Phi_1^\ast$ emphasizes that this perturbation is due to the stars alone. + The equation of mass continuity is. again to linear order. In practice. the term Vej in the momentum equation is relatively sinall.," The equation of mass continuity is, again to linear order, In practice, the term $\bnabla\Phi_1^\ast$ in the momentum equation is relatively small." + However. this term canuot be discarded. as it is the stellar gravity that ultimately drives the outogiug waves.," However, this term cannot be discarded, as it is the stellar gravity that ultimately drives the outoging waves." + The next step is to manipulate equations (3) aud (1) to obtain a wave equation for pj., The next step is to manipulate equations (3) and (4) to obtain a wave equation for $\rho_1$. + Taking the divergeuce of equation (3) yields But the time derivative of equation (1) may be cast as Combining the last two equations yields where pj obeys Polssou's equation: Equation (5) is the inhomogeneous wave equation used in all analyses of dynamical friction in gases., Taking the divergence of equation (3) yields But the time derivative of equation (4) may be cast as Combining the last two equations yields where $\rho_1^\ast$ obeys Poisson's equation: Equation (5) is the inhomogeneous wave equation used in all analyses of dynamical friction in gases. + In. previous studies. however. the source density pp tracks the straight-line trajectory of the gravitating mass.," In previous studies, however, the source density $\rho_1^\ast$ tracks the straight-line trajectory of the gravitating mass." + After solving the wave equation for py. the retarcing force on that mass is calculated.," After solving the wave equation for $\rho_1$, the retarding force on that mass is calculated." + Iu our case. the gravitating mass is spatially confiued.," In our case, the gravitating mass is spatially confined." + It is thus more convenieut to idealize py as being non-zero ouly at the origin., It is thus more convenient to idealize $\rho_1^\ast$ as being non-zero only at the origin. +" This singular. ""equivalent density” for the binary is derived in the next section. wherewe also show how its temporal change generates acoustic waves."," This singular, “equivalent density” for the binary is derived in the next section, wherewe also show how its temporal change generates acoustic waves." +Severa studies. both observational and computational. of groups and clusters of galaxies. and of individual earlv-type galaxies. assume that the hot gas present in these svstems is in hyvedrostatie equilibrium (115) in he overall &ravitational potential.,"Several studies, both observational and computational, of groups and clusters of galaxies, and of individual early-type galaxies, assume that the hot gas present in these systems is in hydrostatic equilibrium (HE) in the overall gravitational potential." + Even {rough the assumption of HIS has been shown to not always be robust (e.g. Bertin et al., Even though the assumption of HE has been shown to not always be robust (e.g. Bertin et al. + 1993: Ciotti Pellegrini 2004). ΕΙΗΝ iuithors assume an isothermal z;-mocdel (Cavaliere Fusco-Fenuano 1976. 1978) to ¢lescribe the eas xrofile. but several solve t1e LUE equation. assuming that the gravitational potential is due to dark matter (DAL) only. usualv described by an NEW (Navarro. Frenk White 1996). profile (c.g. Makino. Sasaki Suto 1998: Suto. Szwaki Alakino 1998). or bv more genera profiles (c.g. Ciotti Pellegrini 2008).," 1993; Ciotti Pellegrini 2004), many authors assume an isothermal $\beta$ -model (Cavaliere Fusco-Femiano 1976, 1978) to describe the gas profile, but several solve the HE equation, assuming that the gravitational potential is due to dark matter (DM) only, usually described by an NFW (Navarro, Frenk White 1996) profile (e.g. Makino, Sasaki Suto 1998; Suto, Sasaki Makino 1998), or by more general profiles (e.g. Ciotti Pellegrini 2008)." + While this approximation seems to hold for groups and custers of galaxies. it is well known that stars contribute a significant fraction of the total and. barvonic mater of elliptical galaxies. and even become the dominant. part at the very. centre of these svstems (c.g. Ferreras. Saha MAlians 200!5).," While this approximation seems to hold for groups and clusters of galaxies, it is well known that stars contribute a significant fraction of the total and baryonic matter of elliptical galaxies, and even become the dominant part at the very centre of these systems (e.g. Ferreras, Saha Williams 2005)." + T10 purpose of this paper is to extend. previous calculations that only considered the potential due to DAL. ane explicitly inclue the contribution of the stellar component as well.," The purpose of this paper is to extend previous calculations that only considered the potential due to DM, and explicitly include the contribution of the stellar component as well." + Providing an analvtic fitting Formula for caleulating the response of the gas to the presence of the gravitational potential induced by both DM and stars is useful while interpreting X-ray data of earlv-tvpe galaxies and lor setting u» initial conditions for simulations of elliptical galaxies., Providing an analytic fitting formula for calculating the response of the gas to the presence of the gravitational potential induced by both DM and stars is useful while interpreting X-ray data of early-type galaxies and for setting up initial conditions for simulations of elliptical galaxies. + The derivation of an equilibrium gas profile neglecting the elfects of stars leads to an, The derivation of an equilibrium gas profile neglecting the effects of stars leads to an +minor (a couple of percent) time-dependent variation iun the caleulated fiux deusity.,minor (a couple of percent) time-dependent variation in the calculated flux density. + Typically between 70 aud. 80 baselines contribute to the final datum., Typically between 70 and 80 baselines contribute to the final datum. +" With an ris noise of 20nuuJwv per bascline. for 1205s averaging times in one LOMA baud. we then achieved. a typical. noise. iu. the source fux density- about 0, for ⊽⋡∙⋅⋡↽cight 10. ΔΠ bauds↴ aud↽⊀↔ 1203s ↸≻↸≻↸≻⋉≨−−≻↓"," With an rms noise of mJy per baseline, for s averaging times in one MHz band, we then achieved a typical noise in the source flux density of about mJy for eight 10 MHz bands and s integration." +↓⋀−- This⋅qe 0.8. is the thermal noise error., This mJy is the thermal noise error. +990-0;-02.— For Q90-0T-05 L telescope (6 baseline). 60⋅ observations frou. μα.9909-07-17 2000 onwards the error is about 1 tines largcr. 0990-07-30 2 12Jx.," For the 4 telescope (6 baseline), MHz observations from December 2000 onwards the error is about 4 times larger, i.e. 3 mJy." + The final flux density error is the quadrature 9990-08-13 of a thermal eror aud the slowly varving eain error 9090-08-27 above which we estimated at typically 1%., The final flux density error is the quadrature sum of a thermal error and the slowly varying gain error discussed above which we estimated at typically. +.With 9099-09-07 median flux density of the source of 200 wJ¥ we cau see 9009-00-21. the overall error is never more than about 1., With a median flux density of the source of 200 mJy we can see that the overall error is never more than about. +5%...9099-00-25 light curves. iu periods of slow scintillations. show ρουμ] our estimations of. the- errors are realistic.," The smooth light curves, in periods of slow scintillations, show that our estimations of the errors are realistic." + OO-00€C the lOunuuinute observations. couducted. in sie990-]]-28 1999.2300€ .aud iu two. periodsHN in September 032000 999-12-11].9909-12-30 November 2000. the errors in the absolute fiux deusity 200-0121 outweigh the errors frou the fit.," For the minute observations conducted in September 1999, and in two periods in September 2000 and November 2000, the errors in the absolute flux density scale outweigh the errors from the fit." + 2000-02-05 spectrum was obtained at the WIT on 17 July 2000 2000-04-14 three 1800 exposures., A spectrum was obtained at the WHT on 17 July 2000 in three sec exposures. + The ISIS spectrograph was 2000-05-13 to obtain a spectimm from 3000A SOOOA..with 2000-05-28 resolution of. 0.56 and 2.9 and∙⊽⊽↜↴⋟⊓ ceutral waveleneths of 2000-07-261195À OncD-DO the blueLG aud 00.red amis ncrespective].," The ISIS spectrograph was used to obtain a spectrum from to, with a resolution of 0.86 and /pixel and central wavelengths of and on the blue and red arms respectively." + spective;Ix..standard 2000-0g-2T=arGT: proceduresou were followed ALISusing the utr package., Standard reduction procedures were followed using the software package. + The combined spectu. with the blue /— boxcar smoothed over 9 pixels. is shown in Fig 1 2000-11-25," The combined spectrum, with the blue arm boxcar smoothed over 9 pixels, is shown in Fig \ref{fig:spectrum}." + fitted Hue properties are shown in Table 2.. 2000-12-10, The fitted line properties are shown in Table \ref{tab:opt}. + calculate the redshift to be 0.5334:0.001., We calculate the redshift to be $\pm$ 0.001. +. The 2000-12-26 is fairly typical for quasars. both in the 2001-01-07. widths of the lines aud iu the continui 2001-01-12. (seo. ce. Miller oet. al.," The spectrum is fairly typical for quasars, both in the equivalent widths of the lines and in the continuum emission (see e.g. Miller et al." + 1992: .Daker 1997: —2001-03-17H et al., 1992; Baker 1997; Brotherton et al. + 2001)., 2001). + We note. the relatively small2 o decrement ΤΠ). possibly indicating little 200 sight2 obscuration .by - dust.," We note the relatively small Balmer decrement $\beta$ $\gamma$ ), possibly indicating little line of sight obscuration by dust." +The [OT bIuniuosity 2001-06-10. —.− inoderatelv 2powerful FRU 2001-07-02 183 520021 2 at 151 MMIIz of 107 2n , The [OII] luminosity $^{34}$ W) is comparable to moderately powerful FRII radio sources (radio luminosities at MHz of $10^{25}-10^{27}$ \citep{wil99}. +no »(Willott Wehaveobserved- - — . oncethis twice.- Ot On€ botlbo," \nocite{bro01,mil92,bak97} We have observed the optical spectrum of this source twice." +"th€ month dav—h numbersH2112145216 917duration was sourcesimilar. aud twice,the lines 1099-00252 251256"," On both ocassions the optical magnitude was similar, and the lines strong and broad." +261262 strong and these observations it does nof — 2.LL.245.246.247.248.249.— Bmin 3815 i$ rather a typical 2000-10, From these observations it does not seem that is a typical OVV; rather a typical quasar. + 2000-00— 300., In Fig. +301.302.303.306 Amin ↕∐⊟∶↴∙⊾∙⊇↖↖↽↸∖↻↥⋅↸∖↴∖↴↸∖∐↑↑↕∐∖∐∶↴∙⊾∐↑⊣⊳↿∐⋅↖↽↸∖↴∖↴∪↖↽↸∖↥⋅⋯∪↥⋅↸∖↑↕↓⋜, \ref{fig:res} we present the light-curves over more than two years of observations. +⋯↑↖↖↽∪↴⋝↸∖↑↖↖⇁↸∖↸∖∐↑↖↖↽∪↕≯↥⋅↸∖≺∣⋯∖∐↸⊳↕↸∖↴∖↴⋜∏⋝∪∏↑↕∪∩≓⊇∩∩⋀∖∐↕∑⋜⋯⋜∐⋅↑∙ ⋅↖⇁↸∖⋜∐⋅↴∖↴∪↕⋡∪↴⋝↴∖↴↸∖↥⋅↖↽⋜↧↑↕∪∐↴∖↴∙↕∐↑∐↸∖↴∖↴↸∖↸⊳↿∐⋅↖⇁↸∖↴∖↴↑∐↸∖⋜↧↖↽↸∖↥⋅⋜↧∶↴," In these curves the averaging time is s. For the analysis we used sec integrations, an example of which is plotted in Fig. \ref{fig:example}." +∙⊾↕∐∶↴⋁↑↕∐⊔∖ ↕↴∖↴⊔∩↴∖∷∖↴∙⊟≻↥⋅↑↕∐∖⋜⋯⋜↧↕⋅↖⇁↴∖↴↕↴∖↴↖↖↽↸∖∏↴∖↴↸∖≼↧∶≩∩↴∖↴↴∖↴↸∖↸⊳↕∐↑↸∖∶↴∙⊾↥⋅⋜↧↑↕∪∐↴∖↴∙⋜⋯↴⋝⋅↖⇁⋅↧∏∐↸∖⋜↕∐≼↧⋅↧↿∏⋅↖↽∪↕≯↑∐⋜↧↑⋅↖⇁, The interrupted light curves in September 1999 are due to rapid switching between two frequencies about 100-200 MHz apart. +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ��⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ��∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻��↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐��," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖���↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖⇁," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖⇁↸," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖⇁↸∖," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖⇁↸∖↥," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +↸∖⋜∐⋅∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑∐↸∖⋅↖⇁↴∖↴↑↕∐⋜∐⋅↸∖↕≯⋜↧↴∖↴↑ . ⋅⋅ ↸∖⊼⋜⋯∏≻↕↸∖∪↕↖↖⇁↕∐↸⊳∐↕↴∖↴↻↕∪↑↑↸∖≼↧↕∐⊡∶↴⋁∙⋅≻∙∙⊺∐↸∖↕∐↑↸∖↥⋅↥⋅∏↻↑↸∖≼↧⋅ ⋅∙≻ ↴⋅ . ⋅↴ ↽⊓⋠ ↕↕∶↴∙⊾∐↑↸⊳↿∐⋅↖↽↸∖↴∖↴⋯∺↸∖↻↑↸∖⋯↴⋝↸∖↥⋅↓⊳⋟⊳⋟≽⋟⋜∐⋅↸∖≺↧⋯∖↑∪↥⋅⋜⋯↕≼↧↴∖↴↖↖↽↕↑↸⊳↕∐∐∶↴⋁ ⋅ ⋅⋅ ↽∕∏∐∖↥⋅⋜∏≻↕≼↧↖⇁⋜∐⋅↕⋜↕↑↕∪∐↴∖↴∪↕⋟⋀∖↕⋜↧⋅↖↽↕∩∩∩∐⋜↧↖↽↸∖↴⋝↸∖≼⊳∪∐∐∖↴∖↴↕∪↖↖⇁↸∖↥⋅," The rapid variations of May 1999 have become slower by June and July of that year, although they still are fast" +In local molecular clouds. molecular hydrogen (115) forms primarily on the surface of dust grains: (wo hvdrogen. atoms are. adsorbed onto the surface of the grain and react to form Ils. which subsequently escapes back into the interstellar medium.,"In local molecular clouds, molecular hydrogen $\mHt$ ) forms primarily on the surface of dust grains: two hydrogen atoms are adsorbed onto the surface of the grain and react to form $\mHt$, which subsequently escapes back into the interstellar medium." + However. II5 can also form in the gas-pliase. primarily through the reactions although some also forms via (he slower reactions," However, $\mHt$ can also form in the gas-phase, primarily through the reactions although some also forms via the slower reactions" +Iu the last decade or so. the presence of stellar streams within nearby galaxies has provided direct evidence for he contimal buikl-up of massive galaxies frou the accretion of small uciglibouriug satellites.,"In the last decade or so, the presence of stellar streams within nearby galaxies has provided direct evidence for the continual build-up of massive galaxies from the accretion of small neighbouring satellites." + For example. iu he Milky Way. there is the Sagittarius stream (Ibatactetal. 2001.. and im MS. a nuniber of new streams iive been recently discovered within the (MeCounachiectal.2009).. adding o the list of already kuown streams in that ealaxy (Ibataetal. 20015).," For example, in the Milky Way, there is the Sagittarius stream \citep{ibata97,ibata01a,majewski03,martinez04}, and in M31, a number of new streams have been recently discovered within the \citep{mcconnachie09}, adding to the list of already known streams in that galaxy \citep{ibata01b}." +.. As well. striking surface brightucss features racing a galaxv merger historv have been fouud iu NGC 5907 (Shanesetal. 1998).. NGC 891 (Mouhlcineetal. 20108)... NCC 1093 (Martínez-Doleadoetal.2008).. NGC 253 and NGC 5236 (Malin&Tadley1997)... ancl UGC 10211 (Forbesetal.2003).. among others.," As well, striking surface brightness features tracing a galaxy's merger history have been found in NGC 5907 \citep{shang98}, , NGC 891 \citep{mouhcine10a}, NGC 4093 \citep{martinez08}, NGC 253 and NGC 5236 \citep{malin97}, and UGC 10214 \citep{forbes03}, among others." + Past work lias shown that it is possible to use streams to trace the orbits of satellite. galaxies that had merece with the Milkv. Wavy in the distant past (Lyuden-Bell&Lyuden-Bell1995:YoonLee 2002).," Past work has shown that it is possible to use streams to trace the orbits of satellite galaxies that had merged with the Milky Way in the distant past \citep{lyndenbell95,yoon02}." +. Lyudeu-Bell&Lyudeu-Bell(1995) sugeest that the tidal debris stripped out of the pareut dwarf ealaxy will not undergo much dynamical friction because its mass content will be low., \cite{lyndenbell95} suggest that the tidal debris stripped out of the parent dwarf galaxy will not undergo much dynamical friction because its mass content will be low. + Iu systems bevoud our Local Group galaxies. fiudiug hese traces of old satellites becomes more ditiicult.," In systems beyond our Local Group galaxies, finding these traces of old satellites becomes more difficult." + But in addition to the field-star integrated light. it is possible o use other old objects such as globular clusters (CC's) and planetary nebulae (PNe) to provide evidence for the xesenuce of accretion reninants (Forteetal.L982:Muzzio1987:Cotéetal.1998.2002:IHilker 1999).," But in addition to the field-star integrated light, it is possible to use other old objects such as globular clusters (GCs) and planetary nebulae (PNe) to provide evidence for the presence of accretion remnants \citep{forte82,muzzio87,cote98,cote02,hilker99}." +".. As well. Pipinoctal.(2007) have modelled. the GC assembly Ustorv m massive galaxies. and showed that iu acdcdition o the GCs that formed along with the ealaxy halo. a senificaut fraction of GCs that make up the massive ealaxies were likely accreted from dwarf galaxies in later ines,"," As well, \cite{pipino07} have modelled the GC assembly history in massive galaxies, and showed that in addition to the GCs that formed along with the galaxy halo, a significant fraction of GCs that make up the massive galaxies were likely accreted from dwarf galaxies in later times." + Indeed. suberoups of CCS have been fouud within he halos of massive galaxies or have been associated with he stellar streams in the Sagittarius stream (Ibata1991.1997:Bellazzinietal.2003). aud in M31(Perrett 2010).," Indeed, subgroups of GCs have been found within the halos of massive galaxies or have been associated with the stellar streams in the Sagittarius stream \citep{ibata94,ibata97,bellazzini03} and in M31\citep{perrett03,mackey10}." +. For our search. we target NCC 5128 (Centaurus Aj which is a giant elliptical galaxy at a distance of 3.830.1 Alpe (ILurisetal.2010a).," For our search, we target NGC 5128 (Centaurus A) which is a giant elliptical galaxy at a distance of $3.8\pm0.1$ Mpc \citep{harris10}." +. While this galaxy is quite close. it is still bevond the distance where detecting fait stellar streams Is an casy task.," While this galaxy is quite close, it is still beyond the distance where detecting faint stellar streams is an easy task." + Previous work bv Maliuetal.(1985). and Pengeetal.(2002). las shown there are existing complex shell structures that surrounds the ceutral regions out to 15 kpe. attributed to the many accretions of inall neiglibouriug galaxies.," Previous work by \cite{malin83} and \cite{peng02} has shown there are existing complex shell structures that surrounds the central regions out to 15 kpc, attributed to the many accretions of small neighbouring galaxies." + There are also III shells (Schimunevichetal.1991) and a prominent warped dust lane (Cwalam1979) is also present in the central region where eas and dust is settling iuto the ceutral potential well., There are also HI shells \citep{schiminovich94} and a prominent warped dust lane \citep{graham79} is also present in the central region where gas and dust is settling into the central potential well. + There is evidence for young stars (Rejkubaetal.2001.2002:Ellis2009) that may be aligned with the radio jet.," There is evidence for young stars \citep{rejkuba01,rejkuba02,ellis09} that may be aligned with the radio jet." + This starformation could have con. trieecred by the jet collidingwith cloud material ought iuto NGC 5128 via a mereiug event., This starformation could have been triggered by the jet collidingwith cloud material brought into NGC 5128 via a merging event. + NGC 5128 has 607 confirmed GCs. via radial velocity ueasuremieuts (vaudeuBerghetal.1981:TesserWoodleyetal.2010a.b) and/or resolved images frou he ou the Telescope (Ilarvisctal.2006:Mouhlcineet 2010b)..," NGC 5128 has 607 confirmed GCs, via radial velocity measurements \citep{vandenbergh81,hesser84,hesser86,harris92,peng04b,woodley05,rejkuba07,beasley08,woodley10a,woodley10b} and/or resolved images from the on the \citep{harris06,mouhcine10b}. ." + Of these. 563 havemeasured radial velocities.," Of these, 563 havemeasured radial velocities." +Their hean mcasurement uucertaintv is 42 kin |. though he range of uncertainties is quite laree. with 96% of he population having uncertaiutics «200 kins +.,"Their mean measurement uncertainty is 42 km $^{-1}$ , though the range of uncertainties is quite large, with $96\%$ of the population having uncertainties $< 200$ km $^{-1}$ ." + The, The +For the data we exposure-correct using exposure maps Weighted in energy using a 3.75 keV 0.97. spectral model (2)... with an appropriate redshift and Galactic absorption.,"For the data we exposure-correct using exposure maps weighted in energy using a 3.75 keV $0.5\Zsun$ spectral model \citep{SmithApec01}, with an appropriate redshift and Galactic absorption." + For we produce exposure maps in the range of 0.8 to 1.6 keV around the Fe-L peak., For we produce exposure maps in the range of 0.8 to 1.6 keV around the Fe-L peak. + As can be seen in these images. the cluster has a smooth surface brightness profile.," As can be seen in these images, the cluster has a smooth surface brightness profile." + There is. however. structure within | aremin radius of the central nucleus and there is a jump in the surface brightness profile around 3 aremin to the west of the nucleus.," There is, however, structure within 1 arcmin radius of the central nucleus and there is a jump in the surface brightness profile around 3 arcmin to the west of the nucleus." + This western edge ean be seen in unsharp-masked images of the cluster in this band (Fig., This western edge can be seen in unsharp-masked images of the cluster in this band (Fig. + 2. left panels: marked as Edge) and may be a poorly defined cold front., \ref{fig:unsharp} left panels; marked as Edge) and may be a poorly defined cold front. + To understand the effect of turbulence and gas motions. we would ideally want to study the 3D distribution of properties of he ICM. including the velocity. pressure. density. temperature and metallieity.," To understand the effect of turbulence and gas motions, we would ideally want to study the 3D distribution of properties of the ICM, including the velocity, pressure, density, temperature and metallicity." + Unfortunately these quantities are difficult to measure., Unfortunately these quantities are difficult to measure. + The most accessible quantity is the X-ray surface brightness., The most accessible quantity is the X-ray surface brightness. + The bolometric surface brightness closely probes the square of the density. integrated along the line of sight.," The bolometric surface brightness closely probes the square of the density, integrated along the line of sight." + The energy band used is sensitive to different plasma oroperties. however.," The energy band used is sensitive to different plasma properties, however." + The 3.5 to 7.5 keV energy band can be used o examine variations in the integral of pressure-squared along the ine of sight for temperatures between | and 3 keV (2)., The 3.5 to 7.5 keV energy band can be used to examine variations in the integral of pressure-squared along the line of sight for temperatures between 1 and 3 keV \citep{FormanM8707}. + In Fig., In Fig. + 3 we show the temperature dependence of the surface brightness in his hard band and in a wider 0.6 to 5 keV band. at constant density and at constant pressure.," \ref{fig:ctdepend} we show the temperature dependence of the surface brightness in this hard band and in a wider 0.6 to 5 keV band, at constant density and at constant pressure." + For this simulation we used a Galactic column density of 107!em7 and the redshift of 77., For this simulation we used a Galactic column density of $10^{21} \psqcm$ and the redshift of 7. + Typical emperatures in 77 are 3.75 keV. so the 3.5 to 7.5 keV energy band is not purely sensitive to the pressure-squared and has some temperature dependence.," Typical temperatures in 7 are $\sim 3.75$ keV, so the 3.5 to 7.5 keV energy band is not purely sensitive to the pressure-squared and has some temperature dependence." + It is a better pressure-squared proxy han the 0.6 to 5 keV band., It is a better pressure-squared proxy than the 0.6 to 5 keV band. + The 0.6 to 5 keV band has very little emperature dependence at constant density and is a good proxy for he integral of the density-squared along the line of sight., The 0.6 to 5 keV band has very little temperature dependence at constant density and is a good proxy for the integral of the density-squared along the line of sight. + We show in Fig., We show in Fig. + 2. (right panels) an unsharp-masked image in the 3.5 to 7.5tA keV band. with the same smoothing parameters as or the 0.6 to UA5 keV images in the left panels.," \ref{fig:unsharp} (right panels) an unsharp-masked image in the 3.5 to 7.5 keV band, with the same smoothing parameters as for the 0.6 to 5 keV images in the left panels." + The signal to noise ratio is lower in this image and we do not clearly see the edge that is apparent in the 0.6 to 5 keV unsharp-masked image., The signal to noise ratio is lower in this image and we do not clearly see the edge that is apparent in the 0.6 to 5 keV unsharp-masked image. + This may be because this band is more pressure sensitive and features are contact discontinuities. or that the signal to noise ratio is lower.," This may be because this band is more pressure sensitive and features are contact discontinuities, or that the signal to noise ratio is lower." + We would like to remove the large scale cluster structure to examine the surface brightness variation on smaller scales., We would like to remove the large scale cluster structure to examine the surface brightness variation on smaller scales. + As the cluster is rather smooth and free of substructure. we model it by fitting ellipses to contours of surface brightness. spaced logarithmically.," As the cluster is rather smooth and free of substructure, we model it by fitting ellipses to contours of surface brightness, spaced logarithmically." + Model values for each pixel are calculated by linearly interpolating in logarithmic space the surface brightness between the two neighbouring ellipses., Model values for each pixel are calculated by linearly interpolating in logarithmic space the surface brightness between the two neighbouring ellipses. + We mask out point sources and excluded sky regions in both the real and model images., We mask out point sources and excluded sky regions in both the real and model images. + We call the ellipse fitting and interpolation model ELFIT for the purposes of this paper., We call the ellipse fitting and interpolation model ELFIT for the purposes of this paper. + The ELFIT model removes both azimuthally symmetric emission and edge features. such as cold fronts. making it ideal to model the more subtle differences due to turbulent motions.," The ELFIT model removes both azimuthally symmetric emission and edge features, such as cold fronts, making it ideal to model the more subtle differences due to turbulent motions." + In the 0.6 to 5 keV band we tit 20 ellipses and 10 in the 3.5 to 7.5 keV band., In the 0.6 to 5 keV band we fit 20 ellipses and 10 in the 3.5 to 7.5 keV band. + To construct a simulated image of the cluster. we take the smooth ELFIT model of the cluster. multiply by the exposure map of the real observation. add a background component and then make a Poisson realisation.," To construct a simulated image of the cluster, we take the smooth ELFIT model of the cluster, multiply by the exposure map of the real observation, add a background component and then make a Poisson realisation." + This model dataset can be compared to the observed surface-brightness by background-subtracting and exposure-correcting., This model dataset can be compared to the observed surface-brightness by background-subtracting and exposure-correcting. + We can include fluctuations in the ELFIT model to see how detectable they are., We can include fluctuations in the ELFIT model to see how detectable they are. + Fig., Fig. + 4. shows the real background-subtracted. exposure-corrected 0.6 to 5 keV and a simulated image of the cluster without any additional fluctuations.," \ref{fig:sim_vs_real} shows the real background-subtracted, exposure-corrected 0.6 to 5 keV and a simulated image of the cluster without any additional fluctuations." + The ELFIT model does a good job at reproducing the overall surface brightness distribution., The ELFIT model does a good job at reproducing the overall surface brightness distribution. + In our images we use a binning of four detector pixels (1.968 arcsec) per image pixel., In our images we use a binning of four detector pixels (1.968 arcsec) per image pixel. + In Section 4.2. we will examine histograms of the surface brightness of the cluster as a function of radius., In Section \ref{sect:histo} we will examine histograms of the surface brightness of the cluster as a function of radius. + To do this we divide the cluster into regions using ellipses fitted to contours in surface brightness., To do this we divide the cluster into regions using ellipses fitted to contours in surface brightness. + The contours are placed logarithmically in surface brightness in the 0.5 to 7 keV band., The contours are placed logarithmically in surface brightness in the 0.5 to 7 keV band. + They are shown in Fig., They are shown in Fig. + 4 and their properties are listed in Table 3.., \ref{fig:sim_vs_real} and their properties are listed in Table \ref{tab:ellipses}. + Images of the model ELFIT surface brightness can be compared to the real surface brightness in detail., Images of the model ELFIT surface brightness can be compared to the real surface brightness in detail. + We smooth both the real data and the model with a Gaussian of @—8 pixels (15.7 arcsec}. and then compute the fractional difference between the smoothed data and smoothed model.," We smooth both the real data and the model with a Gaussian of $\sigma = 8$ pixels (15.7 arcsec), and then compute the fractional difference between the smoothed data and smoothed model." + We also make a simulated dataset using the ELFIT model and compute the fractional difference between this simulation and the model., We also make a simulated dataset using the ELFIT model and compute the fractional difference between this simulation and the model. + Fig., Fig. + 5 shows the fractional differences in the 0.6 to 5 keV band of the real data (top-left panel) and simulated data to the model (bottom-left panel)., \ref{fig:sim_div_real} shows the fractional differences in the 0.6 to 5 keV band of the real data (top-left panel) and simulated data to the model (bottom-left panel). + There are features in the real data which are not visible in the simulated dataset., There are features in the real data which are not visible in the simulated dataset. + The most easily seen features are the apparent radial depressions from the south east of the image to the centre. and from the south-south-east to the centre.," The most easily seen features are the apparent radial depressions from the south east of the image to the centre, and from the south-south-east to the centre." + The 3.5 to 7.5 keV band is sensitive to the pressure-squared., The 3.5 to 7.5 keV band is sensitive to the pressure-squared. + Fig., Fig. + 5 (top-right panel) shows the fractional residuals between the real data and its ELFIT model. smoothing both by a Gaussian of the same size us was used for 0.6 to 5 keV. We also plot the fractional differences between a realisation of this model and itself in the lower-right panel.," \ref{fig:sim_div_real} (top-right panel) shows the fractional residuals between the real data and its ELFIT model, smoothing both by a Gaussian of the same size as was used for 0.6 to 5 keV. We also plot the fractional differences between a realisation of this model and itself in the lower-right panel." + These data are much noisier than the 0.6 to 5 keV band. due to the lower number of counts and the higher level of background.," These data are much noisier than the 0.6 to 5 keV band, due to the lower number of counts and the higher level of background." + Many of the features appear to match the wider band data. however. in particular the troughs in surface brightness in the south.," Many of the features appear to match the wider band data, however, in particular the troughs in surface brightness in the south." + The strength of the variations also appears to match the 0.6 to 5 keV band and is larger than the variations in the simulated data., The strength of the variations also appears to match the 0.6 to 5 keV band and is larger than the variations in the simulated data. + It should be noted. however. that the two bands are not completely independent. although the wider band is dominated by softer emission.," It should be noted, however, that the two bands are not completely independent, although the wider band is dominated by softer emission." + We can also examine the data to check that, We can also examine the data to check that +"must be much longer. in the range Ao,72«LO to 3.10° s. From the detailed discussion given by Desch Connolly (2002) it is evident that the cooling histories of choncdrules are. varied. ancl complicated.","must be much longer, in the range $t_{\rm cool} \sim 2 \times 10^5$ to $3 +\times 10^6$ s. From the detailed discussion given by Desch Connolly (2002) it is evident that the cooling histories of chondrules are varied and complicated." + Scott οἱ al. (, Scott et al. ( +"1996) argue that the cise must have been “a maclstron where temperatures [uctuated: through ~1000 Ix. and solids experienced: ""multiple eveles of melting. evaporation. recondensation. crystallization and aggregation"" '","1996) argue that the disc must have been `a maelstrom where temperatures fluctuated through $\sim 1000$ K and solids experienced “multiple cycles of melting, evaporation, recondensation, crystallization and aggregation” '." + Desch Connolly (2002) give. characteristic cooling times in the range of hours to days., Desch Connolly (2002) give characteristic cooling times in the range of hours to days. + Εις is in line with the recent. review by Connolly et al. (, This is in line with the recent review by Connolly et al. ( +2006). who quote the majority of chondules as having a cooling time of ~107 s. Thus we adopt a typical cooling time for chondrules of The fact that. Fossafant has implications for the energeties of the formation process.,"2006), who quote the majority of chondules as having a cooling time of $\sim 10^5$ s. Thus we adopt a typical cooling time for chondrules of The fact that $t_{\rm cool} \gg t_{\rm rad}$ has implications for the energetics of the formation process." + Not only must the chondrules be heated to melting temperature. but they must so be kept near that temperature for a prolonged. period.," Not only must the chondrules be heated to melting temperature, but they must also be kept near that temperature for a prolonged period." + στhis implies that not just the protochondrules but. also 1ο material surrounding them must be heated to around 2000 IX (Masson. 1996).," This implies that not just the proto–chondrules but also the material surrounding them must be heated to around 2000 K (Wasson, 1996)." + In addition. a sullicient volume of surrounding material must be heated so that its optical thickness implies a cooling time of hours to days.," In addition, a sufficient volume of surrounding material must be heated so that its optical thickness implies a cooling time of hours to days." + The specilic energy required. to heat the surrounding gas to a temperature of 7=2000 Ix is given approximately by where we have assumed that the typical σας particle is à hivelrogen molecule. and have ignored molecular dissociation.," The specific energy required to heat the surrounding gas to a temperature of $T = 2000$ K is given approximately by where we have assumed that the typical gas particle is a hydrogen molecule, and have ignored molecular dissociation." + This implies. again assuming that on average cach element of chondritie material is subject to one melting event. that a fraction f of around of the available accretion energy. must. be used. to provide the heating process for choncdrule formation.," This implies, again assuming that on average each element of chondritic material is subject to one melting event, that a fraction $f$ of around of the available accretion energy must be used to provide the heating process for chondrule formation." + Desch Connolly (2002) have proposed a detailed: moce for the formation of chondrules in terms of the therma processing of particles in shocks in the accretion disc., Desch Connolly (2002) have proposed a detailed model for the formation of chondrules in terms of the thermal processing of particles in shocks in the accretion disc. + Their model takes account of the properties of the shock. including gasdrag heating of particles which are not slower instantly in the gas shock. radiative processes. inclucling the disassociation ancl recombination of Hl». and particle evaporation.," Their model takes account of the properties of the shock, including gas–drag heating of particles which are not slowed instantly in the gas shock, radiative processes including the disassociation and recombination of $_2$, and particle evaporation." + Their model can account for the. therma history of chondrules., Their model can account for the thermal history of chondrules. + Fheir canonical model. requires à shock velocity of For their assumed. ambient. disc temperature of Z=300 Ix and. so sound speed: (assuming a gas of molecular hvdrogen) of e;=1.5«10 em this corresponds to a Mach number of around A4~5.," Their canonical model requires a shock velocity of For their assumed ambient disc temperature of $T = 300$ K and so sound speed (assuming a gas of molecular hydrogen) of $c_s = 1.5 \times +10^5$ cm $^{-1}$, this corresponds to a Mach number of around ${\cal + M} \sim 5$." + llowever. while this model can successfully account for chondrule formation in the disce. the proposed. source. of the shocks within the disc does not sit easily with our current understanding of protostellar disc evolution.," However, while this model can successfully account for chondrule formation in the disc, the proposed source of the shocks within the disc does not sit easily with our current understanding of protostellar disc evolution." + Desch Connolly (2002) propose that the shocks are caused. by eravitational instabilities within the disc., Desch Connolly (2002) propose that the shocks are caused by gravitational instabilities within the disc. + But this proposal is at odds with their disc properties. which are taken from a disc model by Bell et al. (," But this proposal is at odds with their disc properties, which are taken from a disc model by Bell et al. (" +1997) with an accretion rate of AJ—10M. L and a viscous. parameter of a=10.,1997) with an accretion rate of $\dot{M} = 10^{-8} M_\odot$ $^{-1}$ and a viscous parameter of $\alpha = 10^{-4}$. +" The disc density is assumed to be p—19"" & em7. which for a disc semi-thickness 449(0,/M,I?4.107 em gives a disc surface density of X—SOOO & 7."," The disc density is assumed to be $\rho = +10^{-9}$ g $^{-3}$, which for a disc semi-thickness $H = +(c_s/V_\phi) R \approx 4 \times 10^{12}$ cm gives a disc surface density of $\Sigma = 8000$ g $^{-2}$." +" This gives rise to a strongly. unstable disc with a Toomre parameter (Loonie 1964) of Such strong gravitational instability is indeed requirect if this process is to give rise to shocks with velocities V,sOAV.", This gives rise to a strongly unstable disc with a Toomre parameter (Toomre 1964) of Such strong gravitational instability is indeed required if this process is to give rise to shocks with velocities $V_s \approx 0.4 V_\phi$. + Lt is. however. hard to envisage how all the material in such an strongly unstable clise would manage to avoid being strongly shocked a large number of times.," It is, however, hard to envisage how all the material in such an strongly unstable disc would manage to avoid being strongly shocked a large number of times." + Such a dise has à mass at racii around 3 AU (sav. at radi 2 4 AU) of Such a violently unstable disc. is. likely either. to fragment into predominantly gaseous. bodies. gathering any chondrules with them. or. if not. give rise to strong eravitational torquing and so a high accretion rate.," Such a disc has a mass at radii around 3 AU (say, at radii 2 – 4 AU) of Such a violently unstable disc is likely either to fragment into predominantly gaseous bodies, gathering any chondrules with them, or, if not, give rise to strong gravitational torquing and so a high accretion rate." + The accretion rate expected for a disc which is so gravitationally unstable is the same as one for which the viscous parameter has à50.06 (Lodato Rice. 2004: luce. Lodato Armitage. 2005) which would give (Pringle 1981) Such a high aceretion rate would give a midplane disc temperature of ~2000 Ix (Bell et al.," The accretion rate expected for a disc which is so gravitationally unstable is the same as one for which the viscous parameter has $\alpha \approx 0.06$ (Lodato Rice, 2004; Rice, Lodato Armitage, 2005) which would give (Pringle 1981) Such a high accretion rate would give a midplane disc temperature of $\sim 2000$ K (Bell et al.," + 1997) and contracdicts the original assumption that AZ=10“AL. +., 1997) and contradicts the original assumption that $\dot{M} = 10^{-8} M_\odot$ $^{-1}$. + Lt would eive a local clise lifetime of around. MuiAl=600 MEL, It would give a local disc lifetime of around $M_{\rm disc}/\dot{M} \approx 600$ yr. + Similar considerations of accretion rates ancl lifetimes apply to the more detailed. models. presented. by Boss Durisen (2005: see also Boss. 2007).," Similar considerations of accretion rates and lifetimes apply to the more detailed models presented by Boss Durisen (2005; see also Boss, 2007)." + These authors stress that early times. when infall onto the disc was high. the disc is most likely massive enough for self-gravity to play a major role in angular momentum redistribution (see. for example. Lin Pringle. 1990).," These authors stress that early times, when infall onto the disc was high, the disc is most likely massive enough for self-gravity to play a major role in angular momentum redistribution (see, for example, Lin Pringle, 1990)." + Thev also note. however. that the local. non-axisvmmetric instabilities driven by scll-eravity occur al. or around. co-rotation (see also the cliscussion in C'ossins et al.," They also note, however, that the local, non-axisymmetric instabilities driven by self-gravity occur at, or around, co-rotation (see also the discussion in Cossins et al.," + 2009)., 2009). + Εις implies that in general. throughout the bulk of the disc. the relative velocity between dise gas and the (unstable. but transient) spiral pattern is small.," This implies that in general, throughout the bulk of the disc, the relative velocity between disc gas and the (unstable, but transient) spiral pattern is small." + Thus. in general. the short-lived. localised. transient spirals which are the typical outcomes of selíi-gravitational disc instabilities. are not able to drive strong shocks in the gas.," Thus, in general, the short-lived, localised, transient spirals which are the typical outcomes of self-gravitational disc instabilities, are not able to drive strong shocks in the gas." + In this context. in the numerical simulations of Boss Durisen (2005) it is only close to (within a few eric cells of) the inner &rid boundary that they are able to achieve a shock pattern which is neither close to co-rotation. nor strongly trailing.," In this context, in the numerical simulations of Boss Durisen (2005) it is only close to (within a few grid cells of) the inner grid boundary that they are able to achieve a shock pattern which is neither close to co-rotation, nor strongly trailing." + Thus while we cannot rule out the possibility that a massive disc and sclleravitational instabilities are able to provide the necessary shocks. we now consider what we regard as more likely properties for the disc in the region ofa few AU at the time of chondrule formation.," Thus while we cannot rule out the possibility that a massive disc and self-gravitational instabilities are able to provide the necessary shocks, we now consider what we regard as more likely properties for the disc in the region of a few AU at the time of chondrule formation." +evolution of the CSM free-free opacity at different frequencies.,evolution of the CSM free-free opacity at different frequencies. + The flux density. of the supernova was reproduced by the RAMSES model without FFA. multiplied by exp(—7). where τ is the opacity shown in Fig. 3..," The flux density of the supernova was reproduced by the RAMSES model without FFA, multiplied by $\exp{(-\tau)}$, where $\tau$ is the opacity shown in Fig. \ref{CSMTemp}." + We note that 1t is impossible to model the data shown in Fig., We note that it is impossible to model the data shown in Fig. + 3 using a simple model for the CSM temperature. as that deseribed in Appendix AppendixA:..," \ref{CSMTemp} using a simple model for the CSM temperature, as that described in Appendix \ref{RAMSESApp}." + The presence of inhomogeneities (1.e.. clumps) in the CSM of a given radial distribution could help us to model the data.," The presence of inhomogeneities (i.e., clumps) in the CSM of a given radial distribution could help us to model the data." + Indeed. the two earliest data points (the first one at GGHz and the second one at GGHz) in Fig.," Indeed, the two earliest data points (the first one at GHz and the second one at GHz) in Fig." + 3 do not follow the same general trend as the remainder of the data., \ref{CSMTemp} do not follow the same general trend as the remainder of the data. + The opacity at these two epochs is larger than expected from the backward extrapolation of the general trends., The opacity at these two epochs is larger than expected from the backward extrapolation of the general trends. + A possible explanation of these large opacities at very early epochs (earlier than 10 days after explosion) could be the presence of strong inhomogeneities (clumps) in the CSM close to the explosion centre., A possible explanation of these large opacities at very early epochs (earlier than 10 days after explosion) could be the presence of strong inhomogeneities (clumps) in the CSM close to the explosion centre. + A rapid evolution of the CSM temperature to explain these large opacity changes Is less realistic. since it would imply a sudden of the CSM during the first days after the shock breakout.," A rapid evolution of the CSM temperature to explain these large opacity changes is less realistic, since it would imply a sudden of the CSM during the first days after the shock breakout." + The break time fitted around day 360 after explosion may be due to an evolution in the structure index. 5. of the ejected material. to a change in the structure index. s. of the CSM. or to a combination of both.," The break time fitted around day 360 after explosion may be due to an evolution in the structure index, $n$, of the ejected material, to a change in the structure index, $s$, of the CSM, or to a combination of both." + In any case. tf a value »=2 is assumed after this early break. the parameter 7? after the break takes an effective value of 9.7. according to the Chevalier model.," In any case, if a value $s = 2$ is assumed after this early break, the parameter $n$ after the break takes an effective value of 9.7, according to the Chevalier model." + This relatively low value of 5» implies that there has been an enhancement of X-ray luminosity originating in. the shocked ejecta region (although the X-ray emission could still be dominated by the circumstellar shock) given that the shock is more adiabatic and. therefore. its opacity becomes smaller (see Fransson. Lundqvist Chevalier 1996)).," This relatively low value of $n$ implies that there has been an enhancement of X-ray luminosity originating in the shocked ejecta region (although the X-ray emission could still be dominated by the circumstellar shock) given that the shock is more adiabatic and, therefore, its opacity becomes smaller (see Fransson, Lundqvist Chevalier \cite{Fransson1996}) )." + When Fransson. Lundqvist Chevalier (1996)) and Immler. Aschenbach Wang (2001 )) estimated that s~1.6-1.7 from their X-ray data. they did not consider the effect of a greater X-ray luminosity from the shocked ejecta due to n~10.," When Fransson, Lundqvist Chevalier \cite{Fransson1996}) ) and Immler, Aschenbach Wang \cite{Immler2001}) ) estimated that $s \sim 1.6 - 1.7$ from their X-ray data, they did not consider the effect of a greater X-ray luminosity from the shocked ejecta due to $n \sim 10$." + Mioduszewski. Dwarkadas Ball (2001)) simulated radio Images and the radio light curves of 11993J and also claimed that s~1.7 provides the best fit to the data. although these authors did not take the electron radiative cooling into account.," Mioduszewski, Dwarkadas Ball \cite{Miodus2001}) ) simulated radio images and the radio light curves of 1993J and also claimed that $s \sim 1.7$ provides the best fit to the data, although these authors did not take the electron radiative cooling into account." + More recently. Nymark. Chandra Fransson (2009)) reported on a fit to the 11993J X-ray data using a model with y=2 and an X- emission dominated by the reverse shock.," More recently, Nymark, Chandra Fransson \cite{Nymark2009}) ) reported on a fit to the 1993J X-ray data using a model with $s=2$ and an X-ray emission dominated by the reverse shock." + Chandra et al. (2009)), Chandra et al. \cite{Chandra2009}) ) + were also successful in modelling the X-ray data using 5-2., were also successful in modelling the X-ray data using $s=2$. + It is worth noticing that the wide fractional shell reported in Mareaide et al. (2009a)), It is worth noticing that the wide fractional shell reported in Marcaide et al. \cite{Marcaide2009}) ) + and Paper I (7305€ of the outer radius) is incompatible with 5«2 in the frame of the Chevalier model (Chevalier 1982a)). because for s=2 this shell implies that n6 (see Table | of Chevalier 19924).," and Paper I $\sim$ of the outer radius) is incompatible with $s < 2$ in the frame of the Chevalier model (Chevalier \cite{Chevalier1982a}) ), because for $s = 2$ this shell implies that $n \sim 6$ (see Table 1 of Chevalier \cite{Chevalier1982a}) )." + A lower value of s would imply an even lower value of η. which must be larger than 5 for a self-similar expansion (Chevalier 1982a)).," A lower value of $s$ would imply an even lower value of $n$, which must be larger than 5 for a self-similar expansion (Chevalier \cite{Chevalier1982a}) )." + On the other hand. the combination of 5~6 and s=2 translates into wm=0.75. a value much smaller than a=0.87 (our fitted value after the break) although closer to the expansion index fitted to the GGHz data Gv~0.8. see Marcaide et al.," On the other hand, the combination of $n \sim 6$ and $s = 2$ translates into $m = 0.75$, a value much smaller than $m = 0.87$ (our fitted value after the break) although closer to the expansion index fitted to the GHz data $m \sim 0.8$, see Marcaide et al." + 2009a and Paper D. which. in our interpretation. does not describe the true supernova expansion.," \cite{Marcaide2009} and Paper I), which, in our interpretation, does not describe the true supernova expansion." + Could the similarity between the expansion index at 5GGHz and the theoretical value derived from a fractional shell-width indicate that the true expansion curve (r.e.. that corresponding to the forward shock) is traced by the high-frequency data?," Could the similarity between the expansion index at GHz and the theoretical value derived from a fractional shell-width indicate that the true expansion curve (i.e., that corresponding to the forward shock) is traced by the high-frequency data?" + In this case. the evolution of the ejecta opacity would have been the opposite of that proposed in Mareaide et al. (2009a))," In this case, the evolution of the ejecta opacity would have been the opposite of that proposed in Marcaide et al. \cite{Marcaide2009}) )" + and Paper I. That is. the ejecta would have been transparent to the radio emission at all frequencies and early epochs. and would have become increasingly opaque to the GGHz radiation after day 1500.," and Paper I. That is, the ejecta would have been transparent to the radio emission at all frequencies and early epochs, and would have become increasingly opaque to the GHz radiation after day 1500." + We rule out this possibility. since in this case the fit to the," We rule out this possibility, since in this case the fit to the" +(Schramm Turner The second interest is to compare the ?Li abundance observed. with the predictions of recent theoretical models of spallation (Duncan ct al. 1992..,"(Schramm Turner The second interest is to compare the $^6$ Li abundance observed, with the predictions of recent theoretical models of spallation (Duncan et al. \cite{DLL92}," + Vaugioui-Flun ct al. 1991.. 1996.. 1997)).," Vangioni-Flam et al. \cite{VLC94}, \cite{VCOF96}, \cite{VFCR97}) )." + These models have heen triggered x the discovery that the other spallative nuclei Be aud D rave an abundance increasing lincarily with iietallicity im uetal-poor stars (Bocseaard aud Ning 1993.. Doesgaard 1996.. Duncan et al. 1997..," These models have been triggered by the discovery that the other spallative nuclei Be and B have an abundance increasing linearily with metallicity in metal-poor stars (Boesgaard and King \cite{BK93}, , Boesgaard \cite{B96}, Duncan et al. \cite{DPR97}," + Molaro et al. 1997))., Molaro et al. \cite{MBCP97}) ). + More specifically. Vaugioui-Flam et al. (1997))," More specifically, Vangioni-Flam et al. \cite{VFCR97}) )" + have predicted he production ratios of °Li. “Li. De. ?D. HD by low- accelerate uuclei frou SNe IL nupiugiug the overwhehlingly dominant species in the carly iuterstellar natter: IL aud “He.," have predicted the production ratios of $^6$ Li, $^7$ Li, Be, $^9$ B, $^{11}$ B by low-energy accelerated nuclei from SNe II, impinging the overwhelmingly dominant species in the early interstellar matter: H, and $^4$ He." + Receuthy Lemoime et al. (1997)), Recently Lemoine et al. \cite{lemoine}) ) + rave dieussed extensively the abundance found earlier for ( Li in the light of these recent predictions., have dicussed extensively the abundance found earlier for $^6$ Li in the light of these recent predictions. + Let us recall hat the dominant processus of formation for 9Li is the a| reaction. now believed to occur between α particles accelerated. in SNe II ejecta colliding with à particles of the surrounding interstellar matter (ISMD.," Let us recall that the dominant processus of formation for $^6$ Li is the $\alpha+\alpha$ reaction, now believed to occur between $\alpha $ particles accelerated in SNe II ejecta colliding with $ \alpha $ particles of the surrounding interstellar matter (ISM)." + This sets the initial abundance of °Li iu the star formed from this ISAL and part of this °Li may be subsequeutly burnt in the star. where we observe the remaining fraction.," This sets the initial abundance of $^6$ Li in the star formed from this ISM, and part of this $^6$ Li may be subsequently burnt in the star, where we observe the remaining fraction." +" The initial Li can be iufoxred from the predicted ""Li /Be ratio. aud the observed Be abundance (see also Molaro al. 19973)."," The initial $^6$ Li can be inferred from the predicted $^6$ Li /Be ratio, and the observed Be abundance (see also Molaro et al. \cite{MBCP97}) )." + The depletion of πο within the star. by the ?Li(p.? Πο process. is obtained by comparing the observed abundance of Li iu he star to the estimated initial abundance.," The depletion of $^6$ Li within the star, by the $^6$ $^3$ $^4$ He process, is obtained by comparing the observed abundance of $^6$ Li in the star to the estimated initial abundance." + Models of lithimm burning by mining of the convective zoue with deeper lavers allow to translate the estimated depletion level of Li into some maxi epletion level for ‘Li in the same star (Cavrel ct al. L998), Models of lithium burning by mixing of the convective zone with deeper layers allow to translate the estimated depletion level of $^6$ Li into some maximum depletion level for $^7$ Li in the same star (Cayrel et al. \cite{CLM99}) +). For this approach it is of course vital to be sure that SLi is present., For this approach it is of course vital to be sure that $^6$ Li is present. + It is why we have decided to observe IID S1937 with a higher S/N ratio. aud to observe two other halo stars. not vot analysed for Li. BD |12 2667 and BD [36 2165.," It is why we have decided to observe HD 84937 with a higher S/N ratio, and to observe two other halo stars, not yet analysed for $^6$ Li, BD +42 2667 and BD +36 2165." + Actually BD | 12 2667 was also analysed in parallel by Sinith et al. (1998)).," Actually BD + 42 2667 was also analysed in parallel by Smith et al. \cite{Smith98}) )," + but we were not aware of the work at the time of our observations., but we were not aware of the work at the time of our observations. +" The observations aud the method of data reduction are described in section 2,", The observations and the method of data reduction are described in section 2. + Section 3 is the analvsis of the data ii terms of ‘Li aud Li respective abundauces., Section 3 is the analysis of the data in terms of $^7$ Li and $^6$ Li respective abundances. + Section Lis devoted to tle case of ITD &1937., Section 4 is devoted to the case of HD 84937. + Section 5 eives the results for the two other stars., Section 5 gives the results for the two other stars. + Section 6 diseusses the significance of our results. aud section 7 sununuarizes our conclusions.," Section 6 discusses the significance of our results, and section 7 summarizes our conclusions." + Spectra of the three selected stars WD 81957. DD|12 2667 aud BD136 2165 were obtained with the spectrograph GECKO at the 3.6m CFHT telescope in Iia (the log-book of the observations is given iu Table 1).," Spectra of the three selected stars HD 84937, BD+42 2667 and BD+36 2165 were obtained with the spectrograph GECKO at the 3.6m CFHT telescope in Hawaii (the log-book of the observations is given in Table 1)." + The resolving power of the spectrograph. measured on thori lines. is R=100000.," The resolving power of the spectrograph, measured on thorium lines, is $R=100 000$." + The spectra were centered at 6710 A.. the region of the Lill resonance doublet.," The spectra were centered at 6710 , the region of the I resonance doublet." + Ta order to check the width of the lines. a spectru of ID 81937 was also obtained in the region of the stronger. well defined. calcium line at A=6162A.," In order to check the width of the lines, a spectrum of HD 84937 was also obtained in the region of the stronger, well defined, calcium line at $\lambda =6162$." +. The detector was a 20ISX20Ls Lynn pixels CCD fabricated by Loral., The detector was a 2048x2048 $\mu$ m pixels CCD fabricated by Loral. + The nominal gain of this CCD is 2.30 tf. the read out noise 5.3 ¢ i. totally neelieible on our well exposed spectra.," The nominal gain of this CCD is 2.3 e $^{-1}$, the read out noise 5.3 e $^{-1}$, totally negligible on our well exposed spectra." + The pixel width was 155/00. corresponding to .69sA.," The pixel width was $15\mu$ m, corresponding to .02698." + All the spectra have been reduced wiLa nunatie code specially developed at Observatoire de Paxis-Meudon (Spite 19903)., All the spectra have been reduced with a matic code specially developed at Observatoire de Paris-Meudon (Spite \cite{Spi}) ). + It performs the optimal extraction of the spectrum. the flat fielding iux je waveleneth calibration from the comparison lamp cApectrun.," It performs the optimal extraction of the spectrum, the flat fielding and the wavelength calibration from the comparison lamp spectrum." + The wavelength calibration was performed with a argon-jioriuni lamp., The wavelength calibration was performed with a argon-thorium lamp. + The laboratory weveleugis were faken in Palmer Enechnan (1983)) for thorimm. aud iu Ixaufuiuun Edlen (1971)) for Árgou.," The laboratory wevelengths were taken in Palmer Engelman \cite{PE83}) ) for thorium, and in Kaufmann Edlen \cite{KE74}) ) for Argon." + The rims of a urd order polynomial fit. corrected for the actual umber of freedliom. is of the order ofO.," The rms of a third order polynomial fit, corrected for the actual number of freedhom, is of the order of." +"003A. It ust be noted lat. as it is well known (and e. ο, it is alluded to by IIlobbs Thorburn 1997)). the collimation angle between je beam of the luup aud the stellar beam. makes that iere is a possible zero point shift between the wave lenethl scale of the stellar aud lamp spectra. and this shift may change diving the night."," It must be noted that, as it is well known (and e. g. it is alluded to by Hobbs Thorburn \cite{HT97}) ), the collimation angle between the beam of the lamp and the stellar beam, makes that there is a possible zero point shift between the wave length scale of the stellar and lamp spectra, and this shift may change during the night." + Moreover. asx the exposure time or the thorimm Lbuup is about oue hour. the calibration spectra were usually taken at the begiuuiug and at the eud of the üeht.," Moreover, as the exposure time for the thorium lamp is about one hour, the calibration spectra were usually taken at the beginning and at the end of the night." + The thorium spectra are quite usable or establishiug a calibration curve. but their zero-point can be slightly shifted a the time of the stellar exposure.," The thorium spectra are quite usable for establishing a calibration curve, but their zero-point can be slightly shifted at the time of the stellar exposure." + This is why we have determined he zero-point of the wavelength scale mostly from the position of the calcium A=6717 lline. taken iu the same exposure as the ithiui feature.," This is why we have determined the zero-point of the wavelength scale mostly from the position of the calcium $\lambda =6717$ line, taken in the same exposure as the lithium feature." + Although the calcium A=6162.172A Mine Qvaveleneth from Sugar Corliss 1982)) is better defined. it was not used for wawcleneth zero-point determination. or the reason explained above.," Although the calcium $\lambda = 6162.172 $ line (wavelength from Sugar Corliss \cite{SC82}) ) is better defined, it was not used for wawelength zero-point determination, for the reason explained above." + Flat-ticlding has con done using a quartz lamp aud a rapidlv rotatiug hot star., Flat-fielding has been done using a quartz lamp and a rapidly rotating hot star. + Due to simall fringesy. produced x the CCD in the spectral range of the lithium line. the rot star has been finally. preferred for flat fielding.," Due to small fringes produced by the CCD in the spectral range of the lithium line, the hot star has been finally preferred for flat fielding." + The stellar broadening (2:15 iis clearly larecr hau the spectrograph resolution (7 OF Aj). aud thus he stellar profiles are dominated by intrinsic stellar xoadeniue.," The stellar broadening $\approx $ .15 is clearly larger than the spectrograph resolution $\approx $ .07 ), and thus the stellar profiles are dominated by intrinsic stellar broadening." + The models used iu the analysis of the stars have been interpolated im the eri defined by Edvardsson et al. (1993)), The models used in the analysis of the stars have been interpolated in the grid defined by Edvardsson et al. \cite{Edva93}) ) + computed with au updated version of the MIARCS code of Castafssonct al. (1975)), computed with an updated version of the MARCS code of Gustafssonet al. \cite{Gus75}) ) + with tuproved UV line blauketing (see also Edvardssou et al. 1991))., with improved UV line blanketing (see also Edvardsson et al.\cite{Edva94}) ). + The physical paraicters of the models have been taken from the literature aud are given in table 2., The physical parameters of the models have been taken from the literature and are given in table 2. +shows a clear trend toward increasing overall blueshift as the penumbral filaments become more aligned with the slit. which in turn is nearly parallel to the direction toward disk center.,"shows a clear trend toward increasing overall blueshift as the penumbral filaments become more aligned with the slit, which in turn is nearly parallel to the direction toward disk center." + In fact. their bisectors at the relative intensity level are all blueshifted. consistent with the maximum of the Evershed flow in the deeper layers.," In fact, their bisectors at the relative intensity level are all blueshifted, consistent with the maximum of the Evershed flow in the deeper layers." + To study the effect of the finite viewing angle on the bisector velocities from the simulated sunspot and to directly compare with the results of Bellot Rubio et al. (, To study the effect of the finite viewing angle on the bisector velocities from the simulated sunspot and to directly compare with the results of Bellot Rubio et al. ( +2010). we have redone the line synthesis assuming an inclination. by 5.4 to the vertical.,"2010), we have redone the line synthesis assuming an inclination by $5.4^\circ$ to the vertical." + Figure 11. shows the resulting spatially and spectrally degraded maps of the LOS velocity determined from the bisector shifts of the 7090 line at relative intensity level., Figure \ref{fig:figure11} shows the resulting spatially and spectrally degraded maps of the LOS velocity determined from the bisector shifts of the 7090 line at relative intensity level. + Disk center is located to the left or to the right. respectively. in the direction of the slit indicated in the right panel.," Disk center is located to the left or to the right, respectively, in the direction of the slit indicated in the right panel." + In the latter case (right panel). which corresponds to the observations of Bellot Rubio et al. (," In the latter case (right panel), which corresponds to the observations of Bellot Rubio et al. (" +2010). the penumbral region ts clearly dominated by blueshifts due to the projection of the Evershed flow.,"2010), the penumbral region is clearly dominated by blueshifts due to the projection of the Evershed flow." + LOS velocity profiles along the slit are shown in Figure 12.., LOS velocity profiles along the slit are shown in Figure \ref{fig:figure12}. + As the penumbral filaments become more aligned with the slit. the projection of the Evershed flow onto the direction of the LOS increases.," As the penumbral filaments become more aligned with the slit, the projection of the Evershed flow onto the direction of the LOS increases." +" In addition, the Evershed flow speed increases outward in the penumbra."," In addition, the Evershed flow speed increases outward in the penumbra." + The combination of these two effects leads to the downward tilt in the velocity profile shown 1n. Figure 12... which ts more conspicuous for the bisectors in the line wings since the Evershed flow is faster in the deeper layers.," The combination of these two effects leads to the downward tilt in the velocity profile shown in Figure \ref{fig:figure12}, which is more conspicuous for the bisectors in the line wings since the Evershed flow is faster in the deeper layers." + These properties of the velocity profiles are very similar to those found by Bellot Rubio et al. (, These properties of the velocity profiles are very similar to those found by Bellot Rubio et al. ( +2010. see their Figure 3). strongly suggesting that both result from the projection of the Evershed flow.,"2010, see their Figure 3), strongly suggesting that both result from the projection of the Evershed flow." + The difference between the bisector shifts of line core and line wings 1s even higher in the observations. indicating that the increase of the Evershed flow with depth is possibly underestimated in the simulations.," The difference between the bisector shifts of line core and line wings is even higher in the observations, indicating that the increase of the Evershed flow with depth is possibly underestimated in the simulations." + We have found that the observational detection of overtuming convection in penumbral filaments can be severely compromised by a) finite spatial and. spectral resolution and b) projection of the Evershed flow for observations not taken exactly at disk center., We have found that the observational detection of overturning convection in penumbral filaments can be severely compromised by a) finite spatial and spectral resolution and b) projection of the Evershed flow for observations not taken exactly at disk center. + In particular. the relatively weak downflows along the edges of filaments suffer most from these effects: for the 7090 line their observational signal is reduced to values below— the limit of reliable detection around ~200 ms! in most cases.," In particular, the relatively weak downflows along the edges of filaments suffer most from these effects: for the 7090 line their observational signal is reduced to values below the limit of reliable detection around $\sim200$ $\,$ $^{-1}$ in most cases." + The situation is somewhat better for the 5380 line. which samples the higher velocity amplitudes in the deeper layers of the atmosphere owing to its lower formation height.," The situation is somewhat better for the 5380 line, which samples the higher velocity amplitudes in the deeper layers of the atmosphere owing to its lower formation height." + Obviously. better spatial resolution also significantly improves the detectability of these small-scale velocity structures.," Obviously, better spatial resolution also significantly improves the detectability of these small-scale velocity structures." + The combined effects of degradation and. projection probably also affected the observations of Bellot Rubio et al. (, The combined effects of degradation and projection probably also affected the observations of Bellot Rubio et al. ( +2010). so that they cannot decide upon the presence or absence of overturning convection in the penumbra exceeding the stated detection limit of ~50 ms.,"2010), so that they cannot decide upon the presence or absence of overturning convection in the penumbra exceeding the stated detection limit of $\sim50$ $\,$ $^{-1}$." + Note that our analysis did not take into account the effects of noise (due to the detector. electronics. and seeing) in the observed line profiles.," Note that our analysis did not take into account the effects of noise (due to the detector, electronics, and seeing) in the observed line profiles." + Comparing our degraded images with Figure 2 of Bellot Rubio et al. (, Comparing our degraded images with Figure 2 of Bellot Rubio et al. ( +2010) also indicates that the spatial degradation of the observations is stronger than in our attempts to simulate them.,2010) also indicates that the spatial degradation of the observations is stronger than in our attempts to simulate them. + However. the main problem of these observations is the projection of the Evershed flow onto the LOS.," However, the main problem of these observations is the projection of the Evershed flow onto the LOS." + Even at a heliocentric angle of only 5.47. up to of the Evershed flow is projected onto the LOS and the trend in the velocity profile shown in Figure 3 of Bellot Rubio et al. (," Even at a heliocentric angle of only $^\circ$, up to of the Evershed flow is projected onto the LOS and the trend in the velocity profile shown in Figure 3 of Bellot Rubio et al. (" +2010) clearly indicates that this effect is present in their observations.,2010) clearly indicates that this effect is present in their observations. + This is also supported by the fact that the bisector blueshifts are significantly stronger for the line wings than for the line core., This is also supported by the fact that the bisector blueshifts are significantly stronger for the line wings than for the line core. + Therefore. the absence of redshifted bisectors must not necessarily be taken as evidence for the absence of downflows since the projection of the Evershed effect easily leads to a blueshift of several 100 ms7!.," Therefore, the absence of redshifted bisectors must not necessarily be taken as evidence for the absence of downflows since the projection of the Evershed effect easily leads to a blueshift of several 100 $\,$ $^{-1}$." + On the other hand. the non-detection of overturning penumbral convection by Franz Schlichenmaier (2009) and Bellot Rubio et al. (," On the other hand, the non-detection of overturning penumbral convection by Franz Schlichenmaier (2009) and Bellot Rubio et al. (" +2010) could. in principle. indicate that these flows are overestimated 1n. current. numerical simulations owing to insufficient spatial resolution. artificial boundary conditions. and possibly insufficient domain depth.,"2010) could, in principle, indicate that these flows are overestimated in current numerical simulations owing to insufficient spatial resolution, artificial boundary conditions, and possibly insufficient domain depth." + We have recently carried out a convergence study (Rempel. in preparation) for which the numerical grid spacing was varied between 96 km and 16 km in the horizontal directions and between 32 km and 12 km in the vertical direction.," We have recently carried out a convergence study (Rempel, in preparation) for which the numerical grid spacing was varied between 96 km and 16 km in the horizontal directions and between 32 km and 12 km in the vertical direction." + The simulation analyzed here is the second best resolved simulation of that series (the highest resolution case has not yet been computed with non-grey radiative transfer)., The simulation analyzed here is the second best resolved simulation of that series (the highest resolution case has not yet been computed with non-grey radiative transfer). + Since we do not use explicit diffusivities. changes in the grid spacing directly affect the overall numerical dissipation (which scales at least linearly with grid spacing near discontinuities but with a higher order in well-resolved regions).," Since we do not use explicit diffusivities, changes in the grid spacing directly affect the overall numerical dissipation (which scales at least linearly with grid spacing near discontinuities but with a higher order in well-resolved regions)." + The amount of overturning motions (characterized by the vertical rms velocity at 7= 1) turns out to be robust: it is directly tied to the penumbral brightness. which does not change significantly with resolution.," The amount of overturning motions (characterized by the vertical rms velocity at $\tau=1$ ) turns out to be robust: it is directly tied to the penumbral brightness, which does not change significantly with resolution." + On the other hand. we find that the average width of filaments decreases somewhat with increasing resolution since they are still only marginally resolved.," On the other hand, we find that the average width of filaments decreases somewhat with increasing resolution since they are still only marginally resolved." + It 1s therefore well conceivable that our current simulation overestimates the visibility of convective motions in the penumbra., It is therefore well conceivable that our current simulation overestimates the visibility of convective motions in the penumbra. + We also investigated the influence of the top boundary condition., We also investigated the influence of the top boundary condition. + While the overall extent of the penumbra depends on the choice of the boundary condition. the detailed structure of the magnetoconvection is mostly unaffected.," While the overall extent of the penumbra depends on the choice of the boundary condition, the detailed structure of the magnetoconvection is mostly unaffected." + We find an approximate relationship of the form 7x\/v!MS(r=1) between the azimuthally averagec bolometric intensity and the rms vertical velocity defined through the azimuthal average at each radial position of the spot (Rempel 2011). independent of the boundary conditioi and extent of the penumbra.," We find an approximate relationship of the form $I\propto\sqrt{v_{z}^{\mbox{rms}}(\tau=1)}$ between the azimuthally averaged bolometric intensity and the rms vertical velocity defined through the azimuthal average at each radial position of the spot (Rempel 2011), independent of the boundary condition and extent of the penumbra." + As a consequence. the predictior that the vertical rms velocity in the penumbra should be about half of the value found in quiet sun ts fairly robust: it 1s also consistent with the width of distribution functions. inferrec from Hinode data by Franz Schlichenmaier (2009).," As a consequence, the prediction that the vertical rms velocity in the penumbra should be about half of the value found in quiet sun is fairly robust; it is also consistent with the width of distribution functions inferred from Hinode data by Franz Schlichenmaier (2009)." + Our analysis indicate that the 5380 line is better suited than the 7090 line for the detection of overturning convection in penumbral filaments., Our analysis indicate that the 5380 line is better suited than the 7090 line for the detection of overturning convection in penumbral filaments. + This 1s due to the fact that this line originates deeper in the atmosphere (and thus samples higher velocities) and that the shorter wavelength affords a higher spatial resolution for a given telescope., This is due to the fact that this line originates deeper in the atmosphere (and thus samples higher velocities) and that the shorter wavelength affords a higher spatial resolution for a given telescope. + Even after degradation. many downflows exceed a typical detectability threshold of 200 ms7!.," Even after degradation, many downflows exceed a typical detectability threshold of 200 $\,$ $^{-1}$." + Results on penumbral up- and downflows using the 5380 line were obtained by Schlichenmaier Schmidt (1999) with the German VTT on Tenerife and recently by Joshi et al. (, Results on penumbral up- and downflows using the 5380 line were obtained by Schlichenmaier Schmidt (1999) with the German VTT on Tenerife and recently by Joshi et al. ( +2011) and Scharmer et al. (,2011) and Scharmer et al. ( +2011) with the CRISP 2D spectropolarimeter at the SST.,2011) with the CRISP 2D spectropolarimeter at the SST. + In any case. such observations are complicated by the projection of the strong Evershed flows and therefore should be carried as near as possible to disk center— assuming that the Evershed," In any case, such observations are complicated by the projection of the strong Evershed flows and therefore should be carried as near as possible to disk center – assuming that the Evershed" +performed the standard full aperture extraction (same weights within a defined spatial range of pixels).,performed the standard full aperture extraction (same weights within a defined spatial range of pixels). + We tested (hat an aperture of 10 pixels gives the best S/N in the extracted speclra., We tested that an aperture of 10 pixels gives the best S/N in the extracted spectra. + The wavelength calibration is done using the collapsed background frame in spatial direction and cross-correlating il (o a svnthetic model spectrum of the atmosphere., The wavelength calibration is done using the collapsed background frame in spatial direction and cross-correlating it to a synthetic model spectrum of the atmosphere. + The dispersion relation is well approximated by a first order polynomial and the offset in pixels (hat maximizes the cross-correlation can be determined to an accuracy of 0.01 pixel (20.03 kkin/s)., The dispersion relation is well approximated by a first order polynomial and the offset in pixels that maximizes the cross-correlation can be determined to an accuracy of 0.01 pixel $\sim$ km/s). + We also compared (he observed peak position of (wo photospheric absorption lines in the spectra of the Ix5 HII star HD. 136422 and the Ix4.5 HI star ID 139127 to those from the MARCS model atmosphere of a Ix5 IHE star (II. 6705. the spectrum was kindly provided by L. Decin).," We also compared the observed peak position of two photospheric absorption lines in the spectra of the K5 III star HD 136422 and the K4.5 III star HD 139127 to those from the MARCS model atmosphere of a K5 III star (HR 6705, the spectrum was kindly provided by L. Decin)." + This exercise demonstrates that the peak centroids can be determined to a lower accuracy. varving from a tenth up to a few km/s (see also Table 3).," This exercise demonstrates that the peak centroids can be determined to a lower accuracy, varying from a tenth up to a few km/s (see also Table 3)." + All steps described above have been applied to both the target and its standard star(s)., All steps described above have been applied to both the target and its standard star(s). + In addition. the extracted spectrum of the standard star was flux calibrated using (he model flux from the VISIR. standard star catalog.," In addition, the extracted spectrum of the standard star was flux calibrated using the model flux from the VISIR standard star catalog." + From (his calibration we created a spectral response [unetion (Jv/ADU) that we applied at the wavelengths of the science spectrum., From this calibration we created a spectral response function (Jy/ADU) that we applied at the wavelengths of the science spectrum. + This step removes most of the “fringine” present in the VISIR spectra (the flux modulation with wavelengths has a peak to peak amplitude of about and is found to be stable over periods of at least several months. vanBoekeletal.2009) ).," This step removes most of the “fringing” present in the VISIR spectra (the flux modulation with wavelengths has a peak to peak amplitude of about and is found to be stable over periods of at least several months, \citealt{van09}) )." + As mentioned in Sect. ??..," As mentioned in Sect. \ref{sect:targets}," + all targets except HD 34700 have published low- (R105) or high-resolution (R600) IRS spectra with bright emission lines (see Table 1))., all targets except HD 34700 have published low- $\sim$ 105) or high-resolution $\sim$ 600) IRS spectra with bright emission lines (see Table \ref{table:prop}) ). + We have reduced the archival IRS low-resolution spectrum of IID 34700 and report no detection of the, We have reduced the archival IRS low-resolution spectrum of HD 34700 and report no detection of the +In the flavor basis. the mass matrix for Majorana neutrinos contains nine physical parameters including three mass eigenvalues. (μου mixing angles and three CP violating phases.,"In the flavor basis, the mass matrix for Majorana neutrinos contains nine physical parameters including three mass eigenvalues, three mixing angles and three CP violating phases." +" Two squared mass differences (Nar, and Am.) and two mixing angles (845 and 05;) have been measured in solar. atmospheric and reactor experiments."," Two squared mass differences $\Delta m^2_{12}$ and $\Delta +m^2_{13}$ ) and two mixing angles $\theta_{12}$ and $\theta_{23}$ ) have been measured in solar, atmospheric and reactor experiments." + The third mixing angle 944 and the Dirac type CP violating phase 6 are expected to be measured in the forthcoming neutrino oscillation experiments., The third mixing angle $\theta_{13}$ and the Dirac type CP violating phase $\delta$ are expected to be measured in the forthcoming neutrino oscillation experiments. + Possible neasurement of effective Majorana mass in neutrinoless double beta decay searches will provide an additional constraint on the remaining (iree neutrino parameters viz., Possible measurement of effective Majorana mass in neutrinoless double beta decay searches will provide an additional constraint on the remaining three neutrino parameters viz. + neutrino mass scale and (wo Majorana type CP violating phases., neutrino mass scale and two Majorana type CP violating phases. + While (he neutrino mass scale will be independently determined rom the direct beta decay. searches and cosmological observations. (he (vo Majorana phases. will 100 be uniquely determined from the measurement of effective Majorana mass even if the overall jeutrino mass scale is known.," While the neutrino mass scale will be independently determined from the direct beta decay searches and cosmological observations, the two Majorana phases will not be uniquely determined from the measurement of effective Majorana mass even if the overall neutrino mass scale is known." + Thus. it is not possible to filly reconstruct (he neutrino mass matrix rom (he observations from feasible experiments.," Thus, it is not possible to fully reconstruct the neutrino mass matrix from the observations from feasible experiments." + Under (he circumstances. it is natural to emplov other theoretical inputs for the reconstruction of the neutrino mass matrix.," Under the circumstances, it is natural to employ other theoretical inputs for the reconstruction of the neutrino mass matrix." + The possible forms of (hese additional theoretical inputs are constrained by (he existing neutrino data., The possible forms of these additional theoretical inputs are constrained by the existing neutrino data. + Several proposals have been made in literature to restrict the possible forms of the neutrino mass matrix by reducing the number of free parameters which include the presence of texture zeros |l.2.3].. the requirement of zero determinant [4].. the zero trace condition [5] amongst others.," Several proposals have been made in literature to restrict the possible forms of the neutrino mass matrix by reducing the number of free parameters which include the presence of texture zeros \cite{1,2,3}, the requirement of zero determinant \cite{4}, the zero trace condition \cite{5} amongst others." + There have been munerous attempts aimed at understanding the pattern of the lermion masses aud mixings bv introducing Abelian and non-Abelian flavor svimmetries some of which lead to texture zeros in (he fermion mass matrices., There have been numerous attempts aimed at understanding the pattern of the fermion masses and mixings by introducing Abelian and non-Abelian flavor symmetries some of which lead to texture zeros in the fermion mass matrices. + Furthermore. as discussed earlier. it is not possible to [fully reconstruct the neutrino mass matrix solely [rom the results of presently feasible experiments and the introcuction of texture zeros is an extra ingredient aimed at reducing the number of [ree parameters.," Furthermore, as discussed earlier, it is not possible to fully reconstruct the neutrino mass matrix solely from the results of presently feasible experiments and the introduction of texture zeros is an extra ingredient aimed at reducing the number of free parameters." + llowever. some sets of these texture zeros can be obtained by suitable weak basis transformations and have no physical meaning as such.," However, some sets of these texture zeros can be obtained by suitable weak basis transformations and have no physical meaning as such." + Hlowever. a large class of sets of leptonic texture zeros considered in the literature imply (he vanishing of certain CP-odd weak basis invariants ancl one can. thus. recognize a lepton flavor model] in which the texture zeros are not. explicitly present but which corresponds (to a particular texture structure in a certain weak basis.," However, a large class of sets of leptonic texture zeros considered in the literature imply the vanishing of certain CP-odd weak basis invariants and one can, thus, recognize a lepton flavor model in which the texture zeros are not explicitly present but which corresponds to a particular texture structure in a certain weak basis." + The presence of texture zeros. in general. leads to a decrease in the number of independent CP violating phases.," The presence of texture zeros, in general, leads to a decrease in the number of independent CP violating phases." + A particular texture zero structure gives rise to definite relationships between different CP. violating phases [2]., A particular texture zero structure gives rise to definite relationships between different CP violating phases \cite{2}. +. Such exact relations in closed form were obtained in Ref. |6]., Such exact relations in closed form were obtained in Ref. \cite{6}. +". Correlations between Dirac and Majorana CP violating phases lor a particular texture zero scheme were studied in detail in Ref, [ο].", Correlations between Dirac and Majorana CP violating phases for a particular texture zero scheme were studied in detail in Ref. \cite{2}. + It is. therefore. important to examine the interrelationships between the CP odd weak basis Invariants which are required to vanish as a necessary and sufficient condition for CP conservation.," It is, therefore, important to examine the interrelationships between the CP odd weak basis invariants which are required to vanish as a necessary and sufficient condition for CP conservation." + It is the purpose of the present work (o examine svstematically such interrelationships in ternis οἱ the weak basis invariants constructed from the elements of the neutrino mass malrix., It is the purpose of the present work to examine systematically such interrelationships in terms of the weak basis invariants constructed from the elements of the neutrino mass matrix. + The texture zeros are nol weak basis (WB) invariants [7)., The texture zeros are not weak basis (WB) invariants \cite{7}. +. This means that a given set of texture zeros Which arise in a certain WD mav not be present or may appear in different entries in anolLer WD., This means that a given set of texture zeros which arise in a certain WB may not be present or may appear in different entries in another WB. + A large class of sets of leptonic texture zeros considered in the literature imply the vanishing of certain CP-odd weak-basis invariants [7]., A large class of sets of leptonic texture zeros considered in the literature imply the vanishing of certain CP-odd weak-basis invariants \cite{7}. +. Thus. we can recognize a lepton mass model in which the texture zeros are not explicitly present and which corresponds (ο a particular texture scheme in," Thus, we can recognize a lepton mass model in which the texture zeros are not explicitly present and which corresponds to a particular texture scheme in" + unfortunately cannot provide anv firm conclusion about the variability of that object. as these were obtained with different instrumental setups aud spectroscopic resolutions.,"— unfortunately cannot provide any firm conclusion about the variability of that object, as these were obtained with different instrumental setups and spectroscopic resolutions." + To investigate whether GD 323 is a spectroscopic variable. Koesteretal.(1994) obtained a series of nine high signal-to-noise ratio. ~7 rresolution spectra of that object. all secured with an essentially identical instrumental setup.," To investigate whether GD 323 is a spectroscopic variable, \citet{kls94} obtained a series of nine high signal-to-noise ratio, $\sim7$ resolution spectra of that object, all secured with an essentially identical instrumental setup." +" Six ol these spectra. were obtained within a single night. while the other three had. been secured. respectively 15 days. 9 months. aud one vear earlier,"," Six of these spectra were obtained within a single night, while the other three had been secured, respectively 15 days, 9 months, and one year earlier." + On the basis of (hese spectra. the best ancl most homogeneous set secured at that time for a DAB star. [Xoester concluded that the available data were compatible with the assumption that the star showed no variability.," On the basis of these spectra, the best and most homogeneous set secured at that time for a DAB star, \citet{kls94} concluded that the available data were compatible with the assumption that the star showed no variability." + ILowever. it was our feeling al the time that. while it was clear that there were no large-aumplitude changes in the line spectrum of GD 323. the sienal-to-noise ratio of the individual spectra was good enough to permit a much more detailed. analysis ol the variability of GD 323 than had been carried out by Koesteretal.(1994).," However, it was our feeling at the time that, while it was clear that there were no large-amplitude changes in the line spectrum of GD 323, the signal-to-noise ratio of the individual spectra was good enough to permit a much more detailed analysis of the variability of GD 323 than had been carried out by \citet{kls94}." +.. This was undertaken independently by Wesemaeletal.(1995).. but their analvsis could not «quite meet the burden of proof associated wilh an investigation of this type: while (heir results suggested (he presence of small spectroscopic variations in the Koesteretal.(1994). data. their analysis also showed that a well-planned observing strategy. and a rigorous data reduction process would both be required in order to build a stronger case for the presence. or absence. of low-level spectroscopic variations in GD 323.," This was undertaken independently by \citet{wesemaeletal95}, but their analysis could not quite meet the burden of proof associated with an investigation of this type: while their results suggested the presence of small spectroscopic variations in the \citet{kls94} data, their analysis also showed that a well-planned observing strategy and a rigorous data reduction process would both be required in order to build a stronger case for the presence, or absence, of low-level spectroscopic variations in GD 323." + Recently. the opportunity. arose (o revisit (he issue of spectroscopic variability of GD 323. ancl we succeeded in securing new data for that object Chat sheds some light on this 20 vear-old problem.," Recently, the opportunity arose to revisit the issue of spectroscopic variability of GD 323, and we succeeded in securing new data for that object that sheds some light on this 20 year-old problem." + This paper summarizes (he results of (lis revamped investigation., This paper summarizes the results of this revamped investigation. + As discussed above. the study of Koesteretal.(1994). reveals that the spectroscopic variations in GD 323. if anv. are quite small.," As discussed above, the study of \citet{kls94} reveals that the spectroscopic variations in GD 323, if any, are quite small." + IIence a careful observing strategv had to be developed to ensure that any. variabilitw detected in a series of (me-resolved spectroscopic observations is intrinsic to Che star and not the result of changes in the atmospheric transparency or of an artifice introduced in the data reduction., Hence a careful observing strategy had to be developed to ensure that any variability detected in a series of time-resolved spectroscopic observations is intrinsic to the star and not the result of changes in the atmospheric transparency or of an artifice introduced in the data reduction. + The best wav to assess the variability of GD 323 is (o secure in parallel spectroscopic observations of a constant comparison star., The best way to assess the variability of GD 323 is to secure in parallel spectroscopic observations of a constant comparison star. + Ideally. one would like to have both GD 323 and the comparison star on the slit in order to obtain simultaneous time-resolved spectroscopic observations.," Ideally, one would like to have both GD 323 and the comparison star on the slit in order to obtain simultaneous time-resolved spectroscopic observations." + Unfortunately. no such star could be found in (he vicinity of GD 323. and we had (o rely on a different strategy.," Unfortunately, no such star could be found in the vicinity of GD 323, and we had to rely on a different strategy." +The abundances of metals were derived [rom the usual method of equivalent width measurements and curves of growth caleulated in LTE.,The abundances of metals were derived from the usual method of equivalent width measurements and curves of growth calculated in LTE. + Ca. V. Fe and Ti abundances were used as a measure of the metallicity of the stars.," Ca, V, Fe and Ti abundances were used as a measure of the metallicity of the stars." + The [M/II] value shown in Table 4 is (he mean metallicity obtained. from these elements., The [M/H] value shown in Table 4 is the mean metallicity obtained from these elements. + The upper limits in Table 3 were not considered when deriving [M/II]., The upper limits in Table 3 were not considered when deriving [M/H]. + In the star Y CVn we were not able to identify any metallic line useful for abundance analvsis., In the star Y CVn we were not able to identify any metallic line useful for abundance analysis. + For (his purpose. we adopted the metallicity obtained by Lambertetal.(1986). [rom several Fe and Ca lines.," For this purpose, we adopted the metallicity obtained by \citet{lam86} from several Fe and Ca lines." + Note that the number of metallic lines analyzed per star is rather low: minimum. one and maximum. eight [or WZ Cas.," Note that the number of metallic lines analyzed per star is rather low: minimum, one and maximum, eight for WZ Cas." + This star is the only one for which a reasonable statistic with Fe lines (five) can be performed., This star is the only one for which a reasonable statistic with Fe lines (five) can be performed. + We found a mean dispersion of zü.1 dex around the mean iron abundance derived. which is compatible with (he error introduced by the uncertainty in (he equivalent width measurements.," We found a mean dispersion of $\pm 0.1$ dex around the mean iron abundance derived, which is compatible with the error introduced by the uncertainty in the equivalent width measurements." + On the other hand. the elements having isotopes formed by neutron captures have verv few useful lines in the visible spectra of J-stars.," On the other hand, the elements having isotopes formed by neutron captures have very few useful lines in the visible spectra of J-stars." + Note the significant nunmber of empty entries or upper limits in Table 3., Note the significant number of empty entries or upper limits in Table 3. + WZ Cas is again the sole star where il is possible to detect a significant number of heavy element lines., WZ Cas is again the sole star where it is possible to detect a significant number of heavy element lines. + A resolving power of ~10? is needed to perform an accurate analvsis of these stus., A resolving power of $\sim 10^5$ is needed to perform an accurate analysis of these stars. + This means that abundance analyses in C-stars based. on intermediate-low resolution spectra and/or on (he visual intensity οἱ spectral lines can lead to important errors., This means that abundance analyses in C-stars based on intermediate-low resolution spectra and/or on the visual intensity of spectral lines can lead to important errors. + For example. the A4607.34 Sr LI. A4554.04 Da IL and AG709.49 La I features. used by Dominyv(1985). to define an abundance index of these s-process elements. appeared in our spectra as very crowded blends.," For example, the $\lambda 4607.34$ Sr I, $\lambda 4554.04$ Ba II and $\lambda 6709.49$ La I features, used by \citet{dom85} to define an abundance index of these s-process elements, appeared in our spectra as very crowded blends." + At these wavelengths many CN and Cs lines contribute significantly in C-stars., At these wavelengths many CN and $_2$ lines contribute significantly in C-stars. + Therefore. a high intensity of such lines does not necessarily mean an enhancement of Sr. Ba or La. This kind ol analvsis is only useful in relative abundance studies between stars. not to derive absolute abundances.," Therefore, a high intensity of such lines does not necessarily mean an enhancement of Sr, Ba or La. This kind of analysis is only useful in relative abundance studies between stars, not to derive absolute abundances." +accuracy lor 132:6 )) or 364:9 )) jxanets lor a five-year mission or ~19+1 or 55£11 (56%)) planets for a ten-year mission. depending on tie single measurement precision.,"accuracy for $\sim13\pm6$ ) or $36\pm9$ ) planets for a five-year mission or $\sim19\pm7$ ) or $55\pm11$ ) planets for a ten-year mission, depending on the single measurement precision." + Ten-year radial velocity surveys of the sarje slars with 3+) aceuracy would determii accurate orbits (but not ii and sin; separately) lor ~6dτLor 2:t8 planets. for the 1 au2 n Tm target [iss.," Ten-year radial velocity surveys of the same stars with 3 accuracy would determine accurate orbits (but not $m$ and $\sin i$ separately) for $\sim6\pm4$ or $25\pm8$ planets, for the 1 and 2 $\mu$ as target lists." + Depeudiug on the single measurement accuracy. SIN would be expected to determine Inasses and orbits with accuracy for 10 or 26 new plajets. (i.e... planets withou orbi5 determined by the radial velocity survey) for a five vear niissiol aud ~16 or LL new planets for a len-year iuission.," Depending on the single measurement accuracy, SIM would be expected to determine masses and orbits with accuracy for $\sim10$ or $26$ new planets (i.e., planets without orbits determined by the radial velocity survey) for a five year mission and $\sim16$ or $44$ new planets for a ten-year mission." + The two-tier strateey could iueasure masses aud orbits with accuαςv lor ~66τίj and 106+20 planets for a five or ten year mission. respectively. of which 31 aud 60 would no be meastred by a radial velocity survey of the same stars.," The two-tier strategy could measure masses and orbits with accuracy for $\sim66\pm16$ and $106\pm20$ planets for a five or ten year mission, respectively, of which $\sim31$ and $60$ would not be measured by a radial velocity survey of the same stars." + Of course. racial velocity iueasurements alone cannot determine the mass aud inclination separately.," Of course, radial velocity measurements alone cannot determine the mass and inclination separately." + [If additional SIM obse‘vine time were to be allocated to planet surveys. then SIN. woud be expected to detect a larger uumber of planets.," If additional SIM observing time were to be allocated to planet surveys, then SIM would be expected to detect a larger number of planets." + However. he adiional target stars would {οιxl to be more distant thau the sars already targeted. so the 1lease is less than linear in the nuuber OL target stars.," However, the additional target stars would tend to be more distant than the stars already targeted, so the increase is less than linear in the number of target stars." + lu relPlauetsVsStars we plot he number of planets detecte (solid |ines) as a fuuction of the uuuber ol targe stars. assuming a1 equal uutaber of EF. G. Il. axl AL stars are targeted aud oy=Ίμας.," In \\ref{PlanetsVsStars} + we plot the number of planets detected (solid lines) as a function of the number of target stars, assuming an equal number of F, G, K, and M stars are targeted and $\sigma_d = 1\mu$ as." + The dash-clotted lines are for measuretrents of he mass aud orbial parameters with accuacy. tlie dashed lines are lor neasurelvents of the lass with accuacy. aud the dotted lines are for meast=‘ements of the mass aud orbital parameters witli accuracy.," The dash-dotted lines are for measurements of the mass and orbital parameters with accuracy, the dashed lines are for measurements of the mass with accuracy, and the dotted lines are for measurements of the mass and orbital parameters with accuracy." + The top row shows the 1umber of detection 150 planets of all inasses. the middle row shows the 1tunber of detectious of platels less1jassive thau 2047. aud the bottom row shows the jiunber1 of¢etectious of planets witli masses less tIan 34M .," The top row shows the number of detections of planets of all masses, the middle row shows the number of detections of planets less massive than $20 M_\oplus$, and the bottom row shows the number of detections of planets with masses less than $3 M_\oplus$." +Tie left columna is for a five year mission with 21 twe» dimensional observations o Peach ta‘eet star. aid the right columu is for a ten year nisslon witL {5 two dimensional observations of each target star.," The left column is for a five year mission with 24 two dimensional observations of each target star, and the right column is for a ten year mission with 48 two dimensional observations of each target star." +" Cearly. searches for Eartli-uiass jxanets could greatly benelit from a siguiicant lnc'ease in the nunbe το, μίας targeted."," Clearly, searches for Earth-mass planets could greatly benefit from a significant increase in the number of stars targeted." + Iu priuiple. a len year 1ussion could be subcdivided iuto two five vear surveys wilth separate target lists.," In principle, a ten year mission could be subdivided into two five year surveys with separate target lists." + While WὉ separate five year surveys would cdeect a larger total number o{ planets (by ~ 35%). asingle e year survey would be expected to detect slightly more platels nith 1asses less than 2047 (by s 10%) aud a few times as many planets with masses less than BA1," While two separate five year surveys would detect a larger total number of planets (by $\sim35\%$ ), a single ten year survey would be expected to detect slightly more planets with masses less than $20 +M_\oplus$ (by $\sim10\%$ ) and a few times as many planets with masses less than $3 M_\oplus$." + JTjerefore. if the SIM 1uission Cal be extended beyond the planued five vears. planet searcLes slOUd continue observing Slars ΑΗ]ch were part of the original target list.," Therefore, if the SIM mission can be extended beyond the planned five years, planet searches should continue observing stars which were part of the original target list." + An alternative strategvOe for a ten-vear mission is to doule the 1number of target stars aud stil make ouly 21 two dimeusional observations of each targe star., An alternative strategy for a ten-year mission is to double the number of target stars and still make only 24 two dimensional observations of each target star. + Tus would result iu an evelli ereater number of planets detected and characterized to withi1 fo “both planets of all masses, This would result in an even greater number of planets detected and characterized to within for both planets of all masses +conditions.,conditions. + The first condition is based on thresholds on the relative jumps in pressure., The first condition is based on thresholds on the relative jumps in pressure. + For example. for a given cell (4.jA). the relative jump along x-direction is defined as minílp;Wi)lg./val.pinus][νιsagas.," For example, for a given cell $(i,j,k)$, the relative jump along x-direction is defined as $\displaystyle{\frac{|p_{i+1,j,k}-p_{i-1,j,k}|} +{min(|p_{i+1,j,k}|,|p_{i-1,j,k}|)}}$." + WsThe jumps: are computed. for⋅ the three directions (x.v.z) and the condition is satisfied if any of them is larger than the given threshold.," The jumps are computed for the three directions (x,y,z) and the condition is satisfied if any of them is larger than the given threshold." + The second condition is similar to the previous one but is applied to the density., The second condition is similar to the previous one but is applied to the density. + The third condition looks at relative variations of right and left velocity derivatives., The third condition looks at relative variations of right and left velocity derivatives. +" For a given cell (7.7.4) we define⋅ the right⋠ derivate. as 2,=npDiach and the left clerivate as J)=dult."," For a given cell $(i,j,k)$, we define the right derivate as $\displaystyle{D_r=\frac{v^{i+1,j,k}_x-v^{i,j,k}_x}{\Delta x}}$ and the left derivate as $\displaystyle{D_l=\frac{v^{i,j,k}_x-v^{i-1,j,k}_x}{\Delta x}}$." +" Phe relative variation of the derivate is defined as méntlD,J.μοιD)|"," The relative variation of the derivate is defined as $\displaystyle{\frac{|D_r-D_l|}{min(|D_r|,|D_l|)}}$." + The first. condition is specially suited for. identifving shocks. seconcl condition. finds shocks ancl contact iscontinuities. whereas the third. condition can track the jo and teil of the rarefaction waves.," The first condition is specially suited for identifying shocks, second condition finds shocks and contact discontinuities, whereas the third condition can track the head and tail of the rarefaction waves." + The code has automatically allocated and callocatec the numerical patches needed: to. integrate re hvdrodynamies equations., The code has automatically allocated and deallocated the numerical patches needed to integrate the hydrodynamics equations. + La order to illustrate the refinement. structure. the location of cach patch and their evel are plotted in the bottom panel of Figure 2..," In order to illustrate the refinement structure, the location of each patch and their level are plotted in the bottom panel of Figure \ref{tube}." + Vhe main features of the analvtical solution are recovered. well., The main features of the analytical solution are recovered well. + The shock is sharply resolved. in. two or hree cells of the highest level., The shock is sharply resolved in two or three cells of the highest level. + The contact discontinuity and the rarefactions wave are also well described., The contact discontinuity and the rarefactions wave are also well described. + Tinv oscillations are visible in the velocity associated. with the contact cüscontinuitv., Tiny oscillations are visible in the velocity associated with the contact discontinuity. + 3ertschinger (1985) presented the solution for the evolution of a single mass perturbation in a Lat Einstein-de Sitter Universe with no cosmological constant., Bertschinger (1985) presented the solution for the evolution of a single mass perturbation in a flat Einstein-de Sitter Universe with no cosmological constant. +" Given a perturbation of radius £2; at time /; with overdensity 0;=op p. shells of matter surrounding the perturbation start to decelerate anc eventually decouple from the Hubble flow at some ""turn around time which is related to the parameters defining the perturbation as follows: For a collisional fluid. the infalling matter produces an increase in pressure which eventually results in a strong shock wave propagating outwards."," Given a perturbation of radius $R_i$ at time $t_i$ with overdensity $\delta_i=\delta\rho/\rho$ , shells of matter surrounding the perturbation start to decelerate and eventually decouple from the Hubble flow at some ""turn around"" time which is related to the parameters defining the perturbation as follows: For a collisional fluid, the infalling matter produces an increase in pressure which eventually results in a strong shock wave propagating outwards." + According to Bertschinger’s solution. the position of the shock is given by Ax=πα," According to Bertschinger's solution, the position of the shock is given by $\lambda_s=r_s/r_{ta}$." + Due to the self-similar of the characterproblem. the solution is Lully characterized. by. the dimensionless functions V. D and P:," Due to the self-similar character of the problem, the solution is fully characterized by the dimensionless functions V, D and P:" +Over the last [ew years. the Lyman break techuique has been used extensively to isolate salaxies at zzz 2-[ (e.g. Steklel et al.,"Over the last few years, the Lyman break technique has been used extensively to isolate galaxies at $z\approx$ 2–4 (e.g., Steidel et al." + 1996: Lowenthal et al., 1996; Lowenthal et al. + 1997: Dickinson 1998: Steilel et al., 1997; Dickinson 1998; Steidel et al. + 1999) observed near the peak in the cosinic star-Lormatiou rate (e.g.. Blain et al.," 1999) observed near the peak in the cosmic star-formation rate (e.g., Blain et al." + 1999)., 1999). + These objects often exhibit stellar aud interstellar absorption lines characteristic of local starburst ealaxies., These objects often exhibit stellar and interstellar absorption lines characteristic of local starburst galaxies. + Their morphologies are varied. with multiple knots of emission aud diffuse wispy tails that suggest uourelaxed systems.," Their morphologies are varied, with multiple knots of emission and diffuse wispy tails that suggest nonrelaxed systems." + The Lyman break echuique has found particular application in the Hubble Deep Field North (HDF-N: Williams οἱ al., The Lyman break technique has found particular application in the Hubble Deep Field North (HDF-N; Williams et al. + 1996) due to the excellent imagine data aud photometry available there. although the small aigular size of the HDE-N liuits somewhat its utility for statistical stuclies (e.e.. Dickinson 1995).," 1996) due to the excellent imaging data and photometry available there, although the small angular size of the HDF-N limits somewhat its utility for statistical studies (e.g., Dickinson 1998)." + We have recently completed an zz1 Ms observaion of the HDF-N ancl its envirous with the (Weisskopl et al., We have recently completed an $\approx 1$ Ms observation of the HDF-N and its environs with the (Weisskopf et al. + 2000): he, 2000): the +We wish to acknowledge C. C. Steidel and K. A. Adelberger for helpful discussions.,We wish to acknowledge C. C. Steidel and K. A. Adelberger for helpful discussions. + This work was supported by grants NSF-AST-0096023.. NSF-AST-9900866. NSF-AST-9618537. NASA NAGS-8506. and DoE DE-FG03-92-ER40701.," This work was supported by grants NSF-AST-0096023, NSF-AST-9900866, NSF-AST-9618537, NASA NAG5-8506, and DoE DE-FG03-92-ER40701." +According to our aforementioned discussion. we construct a iiodoel with nouconumuutative o-field in the range of {ryΑη|A} and with commutative o-fields in the ranges of (ü.riA} aud Uu|Δ.Ε) where £ veprescuts an infrared cutoff iu the model.,"According to our aforementioned discussion, we construct a model with noncommutative $\phi$ -field in the range of $\{r_H -\Delta, +r_H + \Delta\}$ and with commutative $\phi$ -fields in the ranges of $\{0, r_H -\Delta \}$ and $\{r_H +\Delta,L\}$ where $L$ represents an infrared cutoff in the model." + It is essential that A should be au intrinsic quantity of the model. which characterizes the boundary between the noncommutative space-time range and the commutative space-time ranges. and should be determined bv dvuamics of the model.," It is essential that $\Delta$ should be an intrinsic quantity of the model, which characterizes the boundary between the noncommutative space-time range and the commutative space-time ranges, and should be determined by dynamics of the model." +" Surprisingly. this expectation can be realized aud a divergence-free QET with horizon but without the ""brick wall” can be constructed by the ollowiug consideration: 1) Starting with a simplest nonconuuutative o-field action witlin uetric of black hole. the equation of motion of o can be derived exactly: 2) This equation of notion iun noncommutative field theory should be. of course. quite different from the ordinary Wlein-Gordon equation of o-field within black hole metric."," Surprisingly, this expectation can be realized and a ultraviolet-divergence-free QFT with horizon but without the ""brick wall"" can be constructed by the following consideration: 1) Starting with a simplest noncommutative $\phi$ -field action within metric of black hole, the equation of motion of $\phi$ can be derived exactly; 2) This equation of motion in noncommutative field theory should be, of course, quite different from the ordinary Klein-Gordon equation of $\phi$ -field within black hole metric." + This fact implies o-fields should ο Moved in a curve space with a new effective metric g/7:3) Remarkably. it will be shown lose that g/ has two new sineularities besides the the original one at r=ry. one is outside he horizon and auother is iu inside.," This fact implies $\phi$ -fields should be moved in a curve space with a new effective metric $\widetilde{g}^{\mu\nu}$ ;3) Remarkably, it will be shown below that $\widetilde{g}^{tt}$ has two new singularities besides the the original one at $r=r_H$, one is outside the horizon and another is in inside." + Denoting their locations as ry+A’ respectively. we will fud. A’ isdependent on the euereies of the noucommuitative fields wie. A’=A'(w).," Denoting their locations as $r_H +\pm \Delta '$ respectively, we will find $\Delta '$ isdependent on the energies of the noncommutative fields $\omega$ ,i.e., $\Delta '=\Delta '(\omega)$ ." + This means that to the fields with euerev i. o(uw.r). its red-shifting ou the (egdA(uw) surface is iufiuite due to ger—rydAG)=0. aud then we have 1j We arene that the fact that the uoucomuuttative fields vauish at r=rg|Aw) lncans we can think the space-time coordinates ou the surface of r—ry|Atte) to be comuutative. be. [ένΕν4patie?=0.," This means that to the fields with energy $\omega$, $\phi(\omega,r)$, its red-shifting on the $(r_H\pm \Delta '(\omega))$ surface is infinite due to $\widetilde{g}_{tt}(r=r_H\pm\Delta' (\omega))=0$, and then we have 4) We argue that the fact that the noncommutative fields vanish at $r=r_H + \Delta '(\omega)$ means we can think the space-time coordinates on the surface of $r=r_H + \Delta '(\omega)$ to be commutative, i.e., $[t,r]|_{r=r_H + \Delta '(\omega)}=0$." + and then we will further have Comparing eq.(9)) with eq.(7)). we ect Because of the appearance the new infinite red-shifting surface. an observer outside this surface can not detect the o4 atr10% flux variation over a period of about 20 years. while this fraction is only for GPS galaxies.,"2009) found that GPS quasars exhibit $>10\%$ flux variation over a period of about 20 years, while this fraction is only for GPS galaxies." + Detailed analysis of our monitoring data will be presented in another paper. and here we just list the measured flux densities of our VLBI targets in Table 3. for epochs during which most of the VLBI targets are observed.," Detailed analysis of our monitoring data will be presented in another paper, and here we just list the measured flux densities of our VLBI targets in Table \ref{table:flux} for epochs during which most of the VLBI targets are observed." + The flux densities were determined using antenna slews with “cross - scans? in azimuth and elevation. fourfold in each coordinate.," The flux densities were determined using antenna slews with `cross - scans' in azimuth and elevation, fourfold in each coordinate." + This enabled us to check the pointing offsets in both coordinates., This enabled us to check the pointing offsets in both coordinates. + After applying a correction for small pointing offsets. the amplitudes of both AZ and EL were averaged.," After applying a correction for small pointing offsets, the amplitudes of both AZ and EL were averaged." + We then corrected the measurements for the elevation-dependent antenna gain and the remaining systematic time-dependent effects. using a number of steep spectrum and non-vartable secondary calibrators.," We then corrected the measurements for the elevation-dependent antenna gain and the remaining systematic time-dependent effects, using a number of steep spectrum and non-variable secondary calibrators." + Finally. we related our observations to the absolute flux density by using the scale 7.5 Jy at 5 GHz of the primary calibrator 3C286 (Ott et al.," Finally, we related our observations to the absolute flux density by using the scale 7.5 Jy at 5 GHz of the primary calibrator 3C286 (Ott et al." + 1994)., 1994). + Given the flux variation measured in 2007-2009. and by comparing with PKS90 5 GHz data in Table 3.. we find that the flux densities are quite stable or slightly different for the majority of the twelve GPS sources. but that two core-jet sources J1203+0414 (quasar) and J1600—0037 exhibit considerable variability in total flux density.," Given the flux variation measured in 2007–2009, and by comparing with PKS90 5 GHz data in Table \ref{table:flux}, we find that the flux densities are quite stable or slightly different for the majority of the twelve GPS sources, but that two core-jet sources J1203+0414 (quasar) and $-$ 0037 exhibit considerable variability in total flux density." + It is interesting to estimate the jet-viewing angles for the confirmed CSOs. because the CSOs are assumed to lie nearly within the plane of sky.," It is interesting to estimate the jet-viewing angles for the confirmed CSOs, because the CSOs are assumed to lie nearly within the plane of sky." + From the VLBI morphologies and the steep spectral indices of the VLBI components between 1.6 GHz and 5 GHz. we can firmly classify JO210+0419. 135-0021. and J2058+0540 as CSOs that display symmetric double lobes/tails.," From the VLBI morphologies and the steep spectral indices of the VLBI components between 1.6 GHz and 5 GHz, we can firmly classify J0210+0419, $-$ 0021, and J2058+0540 as CSOs that display symmetric double lobes/tails." +" We estimated the jet viewing angle from the flux ratio between the approaching and receding lobes with the formula (Taylor Vermeulen 1997) where $,/S, is the flux ratio between the approaching (stronger) lobe and receding lobe. O is the viewing angle (degree) of the jet axis. « is the source spectral index ay in optical-thin regime (from paper D. &=2o0rk =3 for a continuous or discrete jet. respectively. and f is the lobe velocity as a fraction of the speed of light."," We estimated the jet viewing angle from the flux ratio between the approaching and receding lobes with the formula (Taylor Vermeulen 1997) where $S_{a}/S_{r}$ is the flux ratio between the approaching (stronger) lobe and receding lobe, $\Theta$ is the viewing angle (degree) of the jet axis, $\alpha$ is the source spectral index $\alpha_{h}$ in optical-thin regime (from paper I), $k=2$ or $k=3$ for a continuous or discrete jet, respectively, and $\beta$ is the lobe velocity as a fraction of the speed of light." + By adopting B=0.1. which is à mean lobe velocity in CSOs (Giroletti Polatidis 2009). we estimated the jet viewing angles for the continuous and discrete cases at 1.6 GHz and 5 GHz. respectively. in Table 4..," By adopting $\beta=0.1$, which is a mean lobe velocity in CSOs (Giroletti Polatidis 2009), we estimated the jet viewing angles for the continuous and discrete cases at 1.6 GHz and 5 GHz, respectively, in Table \ref{table:angle}." + From a rough estimation. we could say that the three CSOs have relatively large jet-viewing angles.," From a rough estimation, we could say that the three CSOs have relatively large jet-viewing angles." + Although this method should be used only for the confirmed double-lobe sources. and the lobes of the CSOs are not resolved too much in the VLBI images.," Although this method should be used only for the confirmed double-lobe sources, and the lobes of the CSOs are not resolved too much in the VLBI images." + Nevertheless the najority of sources in our sample have double-lobe like morphologies and total flux densities that are quite stable over 20 years., Nevertheless the majority of sources in our sample have double-lobe like morphologies and total flux densities that are quite stable over 20 years. + Only 3 CSOs are firmly classified according to Fanti’s suggestion that a CSO should have a core in-between double lobes or have double lobes with twin jets/tails rf a core is not detected (Fanti 2009)., Only 3 CSOs are firmly classified according to Fanti's suggestion that a CSO should have a core in-between double lobes or have double lobes with twin jets/tails if a core is not detected (Fanti 2009). + Furthermore. 6 sources are resolved-out in the 5 GHz VLBI images by more than of total flux density. probably due to diffuse emission in the sources.," Furthermore, 6 sources are resolved-out in the 5 GHz VLBI images by more than of total flux density, probably due to diffuse emission in the sources." + This may lead to quite steep spectral indices of VLBI components of the sources as listed in Table 1.., This may lead to quite steep spectral indices of VLBI components of the sources as listed in Table \ref{table:gps}. + The spectra of all 6 sources peak at 0.4 GHz (Snellen et al., The spectra of all 6 sources peak at 0.4 GHz (Snellen et al. + 2002). which is close to the spectral peaks of CSS sources (see e.g.. Fanti et al.," 2002), which is close to the spectral peaks of CSS sources (see e.g., Fanti et al." + 1995). so some of the sources are probably not GPS but rather CSS sources if they have larger-scale diffuse emission that is resolved-out by the VLBI observations.," 1995), so some of the sources are probably not GPS but rather CSS sources if they have larger-scale diffuse emission that is resolved-out by the VLBI observations." +"not strictly isothermal, as the models assume.","not strictly isothermal, as the models assume." +" Furthermore, correcting for the systematic offset due to using the virial temperature rather than the X-ray temperature would bring the models into better agreement with the data."," Furthermore, correcting for the systematic offset due to using the virial temperature rather than the X-ray temperature would bring the models into better agreement with the data." +" The hot gas fraction (Mnot/Mvir), however, shows no dependence with temperature over this range, unlike the data, which shows a mild trend of increasing baryonic fraction with temperature."," The hot gas fraction $M_{\mathrm{hot}}/M_{\mathrm{vir}}$ ), however, shows no dependence with temperature over this range, unlike the data, which shows a mild trend of increasing baryonic fraction with temperature." + This behavior was already seen in S08 (their Figure 8)., This behavior was already seen in S08 (their Figure 8). +" ? have shown that if “radio mode” AGN feedback not only prevents the cooling of gas but is also allowed to eject some of the hot gas out of the halo, lower-mass clusters in the models will also show lower gas fractions."," \citet{bower08} have shown that if “radio mode” AGN feedback not only prevents the cooling of gas but is also allowed to eject some of the hot gas out of the halo, lower-mass clusters in the models will also show lower gas fractions." + It is worth noting that our models agree well with themean gas fraction of the entire data sample and that model haloes below 1 keV (My~ 10'*Mo) show a sudden drop in the predicted gas fraction (see S08 Figure 8)., It is worth noting that our models agree well with the gas fraction of the entire data sample and that model haloes below 1 keV $M_{\mathrm{vir}}\sim 10^{12}\msun$ ) show a sudden drop in the predicted gas fraction (see S08 Figure 8). +" This step-like behaviour in the gas fraction is common to other models (??),, and is due to the rapid transition from infall-limited cooling (sometimes called “cold mode"") to cooling-time limited cooling (“hot mode"")."," This step-like behaviour in the gas fraction is common to other models \citep{deLucia04,menci06}, and is due to the rapid transition from infall-limited cooling (sometimes called “cold mode”) to cooling-time limited cooling (“hot mode”)." +" Our model proved in Paper I to be successful at reproducing the metallicity- and [a/Fe]-mass relations of local early-type galaxies, as well as the SN rates as a function of SSFR."," Our model proved in Paper I to be successful at reproducing the metallicity- and $\alpha$ /Fe]-mass relations of local early-type galaxies, as well as the SN rates as a function of SSFR." + We now examine the iron content and the abundance ratios between different elements and Fe in the ICM to further test its accuracy., We now examine the iron content and the abundance ratios between different elements and Fe in the ICM to further test its accuracy. + We use an ensemble of X-ray cluster surveys for this purpose., We use an ensemble of X-ray cluster surveys for this purpose. +" From ? we take Si and Fe; from ? we take O, Ne, Mg, Si and Fe; from ? we take Fe ?); and from ? we take Si, S, Ar, Ca, Fe and Ni."," From \citet{fukazawa98} we take Si and Fe; from \citet{peterson03} we take O, Ne, Mg, Si and Fe; from \citet{degrandi04} we take Fe \citep[see + also][]{ettori02}; and from \citet{dePlaa07} we take Si, S, Ar, Ca, Fe and Ni." +" Galaxy clusters often show metallicity gradients for some elements, with increasing abundances towards the cluster centre."," Galaxy clusters often show metallicity gradients for some elements, with increasing abundances towards the cluster centre." +" These clusters, known ascore (CC) clusters, are mostly relaxed systems and the central metal enhancement is associated with feedback from the BCG."," These clusters, known as (CC) clusters, are mostly relaxed systems and the central metal enhancement is associated with feedback from the BCG." +" In contrast,core (NCC) clusters have almost flat abundance profiles and show signatures of recent merging events (?).."," In contrast, (NCC) clusters have almost flat abundance profiles and show signatures of recent merging events \citep{DGM01}." +" Since the observational data are measured near the clusters centres and our models predict global abundances averaged over an entire cluster, it is necessary to correct the observations for gradients, but only for those clusters tagged as havingcores."," Since the observational data are measured near the clusters centres and our models predict global abundances averaged over an entire cluster, it is necessary to correct the observations for gradients, but only for those clusters tagged as having." +" We do this by following the procedure of A,, who used the results of ? on Fe gradients to convert the measured central Fe abundance to global average values."," We do this by following the procedure of \citet[Appendix + A]{Nagashima05a}, who used the results of \citet{degrandi04} on Fe gradients to convert the measured central Fe abundance to global average values." +" Elements like Si, S, Ar, Ca and Ni are known to have the same gradients as Fe, and are corrected by the same factor."," Elements like Si, S, Ar, Ca and Ni are known to have the same gradients as Fe, and are corrected by the same factor." +" On the other hand, O, Ne and Mg do not show gradients even in CC clusters, so we assume that the global abundance is equal to the central measurement (??).."," On the other hand, O, Ne and Mg do not show gradients even in CC clusters, so we assume that the global abundance is equal to the central measurement \citep{tamura01,tamura04}." +" We have also renormalised the abundances to the solar values of ?,, as in the models."," We have also renormalised the abundances to the solar values of \citet{Grevesse}, as in the models." +" In Figure 2,, we examine the elemental abundance of iron ([Fe/H])."," In Figure \ref{iron}, we examine the elemental abundance of iron ([Fe/H])." + We pay particular attention to Fe because it is the ICM element most precisely measured and most extensively studied., We pay particular attention to Fe because it is the ICM element most precisely measured and most extensively studied. +" Both the data and the models show a flat behaviour with temperature, an effect also seen in the"," Both the data and the models show a flat behaviour with temperature, an effect also seen in the" +171488 have similar levels of cliflerential rotation to that of Tau Boo and do not show any evidence of a magnetic polarity reversal over ~8 vears of observations (Jelfersetal.2010:Marsdenct 2010).,"171488 have similar levels of differential rotation to that of Tau Boo and do not show any evidence of a magnetic polarity reversal over $\sim$ 3 years of observations \citep{JeffersSV:2010, MarsdenSC:2010}." +. So the role of differential rotation in the magnetic eveles (or i£ they even have eveles) of these voung stars is still unknown., So the role of differential rotation in the magnetic cycles (or if they even have cycles) of these young stars is still unknown. + Fie., Fig. + 3 shows the structure of the coronal magnetic field and also of the X-ray emission for the 2007. 2009 and 2010 magnetic maps (see Paper D).," \ref{cor_struct} shows the structure of the coronal magnetic field and also of the X-ray emission for the 2007, 2009 and 2010 magnetic maps (see Paper I)." + Open field lines that would o9 X-ray clark and carrying the stellar wind. are shown Xue. while the N-rav bright. closecl field. lines are shown white.," Open field lines that would be X-ray dark and carrying the stellar wind, are shown blue, while the X-ray bright closed field lines are shown white." + On the largest scales. the dipole component. of the ield dominates and the field structure for the three vears is similar.," On the largest scales, the dipole component of the field dominates and the field structure for the three years is similar." + The main change on the large scale is in the tilt of he dipole component of the field., The main change on the large scale is in the tilt of the dipole component of the field. + Between 2007 ancl 2009. he dipole axis shifted from a Latitude of 68 degrees to down ο 52 degrees and then back to 72 in 2010.," Between 2007 and 2009, the dipole axis shifted from a latitude of 68 degrees to down to 52 degrees and then back to 72 in 2010." + On smaller scales. he field structure is very dillerent between the three vears and this is rellectecl in the N-ray images.," On smaller scales, the field structure is very different between the three years and this is reflected in the X-ray images." + As noted in Paper L the poor quality of the 2009 dataset may play some role in hese results., As noted in Paper I the poor quality of the 2009 dataset may play some role in these results. + Fig., Fig. + 7 shows the magnitude and rotational modulation of the emission. measure. and also the emission-measure weighted density: These are calculated for a range of values of the base pressure. since without simultaneous X-ray observations we are unable to determine the value of base pressure.," \ref{cor_mod} shows the magnitude and rotational modulation of the emission measure, and also the emission-measure weighted density: These are calculated for a range of values of the base pressure, since without simultaneous X-ray observations we are unable to determine the value of base pressure." + We also show in the left. column results for à model with a small corona (extending to 2.5 )) and in the right column results for a model with a larger corona (extending to 4.8 ))., We also show in the left column results for a model with a small corona (extending to 2.5 ) and in the right column results for a model with a larger corona (extending to 4.8 ). + Lt is immediately apparent that the small changes introduced by uncertainty in the size of the corona are small compared to the changes between vears., It is immediately apparent that the small changes introduced by uncertainty in the size of the corona are small compared to the changes between years. + We expect the corona to be populated with loops at à range of temperatures. but. in the absence of any. data that might. determine the nature of the differential emission. measure. we take the simplest approach.," We expect the corona to be populated with loops at a range of temperatures, but in the absence of any data that might determine the nature of the differential emission measure, we take the simplest approach." + We therefore show results for a corona at a uniform temperature of Ix and for comparison. one at a uniform temperature of Kx. As might be expected. at the lower temperatures. the densities and. emission measures reach higher values.," We therefore show results for a corona at a uniform temperature of $\times$ K and for comparison, one at a uniform temperature of $\times$ K. As might be expected, at the lower temperatures, the densities and emission measures reach higher values." + The range of values is however similar for all vears and is typical for active stars., The range of values is however similar for all years and is typical for active stars. + The largest dillerence between the three vears is in the arrangement of the smaller scale loops and this results in a significant increase in the X-ray rotational mocdulation in 2009., The largest difference between the three years is in the arrangement of the smaller scale loops and this results in a significant increase in the X-ray rotational modulation in 2009. + We have extrapolated the coronal magnetic field of 11D 141943 from surface magnetograms acquirecl at 3 epochs., We have extrapolated the coronal magnetic field of HD 141943 from surface magnetograms acquired at 3 epochs. + These show that the large-scale field structure of the corona is dominated. by a dipole component with the axis of the dipole shifting between the 3 epochs., These show that the large-scale field structure of the corona is dominated by a dipole component with the axis of the dipole shifting between the 3 epochs. + The small scale structure shows an increase in the mocdellect rotational X-ray moculation in 2009. compared to the other epochs (2007 2010).," The small scale structure shows an increase in the modelled rotational X-ray modulation in 2009, compared to the other epochs (2007 2010)." + The surface differential rotation of LID 141943 has been measured at 2 epochs (2007 2010) from both the surface brightness features (2010) and the surface magnetic features (2007 2010)., The surface differential rotation of HD 141943 has been measured at 2 epochs (2007 2010) from both the surface brightness features (2010) and the surface magnetic features (2007 2010). + The differential rotation measured from the magnetic features is one of the highest values of dillerential rotation measured on a voung solar-tvpe star and is similar in level to the other voung carly-G star LID 171488. but higher than the more swollen star LD LO6506.," The differential rotation measured from the magnetic features is one of the highest values of differential rotation measured on a young solar-type star and is similar in level to the other young early-G star HD 171488, but higher than the more swollen star HD 106506." + We thus conclude that the depth of the stellar convective zone plavs a strong role in the level of surface dillerential rotation seen on solar-type stars. with a large increase in cillerential rotation seen for star's with convective zone depths shallower than ~0.2 1t...," We thus conclude that the depth of the stellar convective zone plays a strong role in the level of surface differential rotation seen on solar-type stars, with a large increase in differential rotation seen for star's with convective zone depths shallower than $\sim$ 0.2." + The 2010 dataset for LID 141948 shows a laree increase in the level of differential rotation measured from magnetic features to that measured from. brightness features., The 2010 dataset for HD 141943 shows a large increase in the level of differential rotation measured from magnetic features to that measured from brightness features. + This is similar to that seen on earlv-Ix. stars but with a much &ereater dillerence and is in contrast to the results from other early-G stars which show Littlor no dilflerence between the differential rotation measured from brightness and magnetic features., This is similar to that seen on early-K stars but with a much greater difference and is in contrast to the results from other early-G stars which show little or no difference between the differential rotation measured from brightness and magnetic features. + Our results. only find tentative evidence for temporal evolution in the dilferential rotation of LID 141943., Our results only find tentative evidence for temporal evolution in the differential rotation of HD 141943. + These results when combined with those from the early-C star LID 171488 (which shows no evidence of temporal evolution in dilferential rotation) imply that carly-G stars do not undergo large-scale. evolution in their dillerential rotation., These results when combined with those from the early-G star HD 171488 (which shows no evidence of temporal evolution in differential rotation) imply that early-G stars do not undergo large-scale evolution in their differential rotation. + However. the errors in our measurements are too large to rule out small scale evolution in cillerential rotation similar to that seen on carly-l stars.," However, the errors in our measurements are too large to rule out small scale evolution in differential rotation similar to that seen on early-K stars." + HD) 141943 and stars of similar spectral tvpe warrant further observations to determine what elfect a shallow convective zone has on the cillerential rotation levels of such stars and indeed if they do show temporal evolution of their cillerential rotation as seen on carly-ly stars., HD 141943 and stars of similar spectral type warrant further observations to determine what effect a shallow convective zone has on the differential rotation levels of such stars and indeed if they do show temporal evolution of their differential rotation as seen on early-K stars. + The observations in this paper were obtained with the Anelo-Australian telescope., The observations in this paper were obtained with the Anglo-Australian telescope. + We would like to thank the technical stall of the Anelo-Australian Observatory (now the Australian Astronomical Observatory) for their. as usual. exccllent assistance curing these observations.," We would like to thank the technical staff of the Anglo-Australian Observatory (now the Australian Astronomical Observatory) for their, as usual, excellent assistance during these observations." + We would also like to thank the anonymous referee who helped improve this paper., We would also like to thank the anonymous referee who helped improve this paper. + This project is supported. bv the Commonwealth of Xustralia under the International Science Linkages program., This project is supported by the Commonwealth of Australia under the International Science Linkages program. +the LESS sources with fluxes above 6 percent of the Crab flux were taken into account (o derive the number-intensity relation.,the HESS sources with fluxes above 6 percent of the Crab flux were taken into account to derive the number-intensity relation. + The contribution of unresolved IIESS-like sources to the dilfuse emission measured by Milagro was also estimated., The contribution of unresolved HESS-like sources to the diffuse emission measured by Milagro was also estimated. + Using the logN-logS relation for the HESS sample of Galactic 5-rav. emitters at least LO per cent of the diffuse emission al TeV energies is estimated to be due to the contribution of unresolved I1ESS-like sources., Using the logN-logS relation for the HESS sample of Galactic $\gamma$ -ray emitters at least 10 per cent of the diffuse emission at TeV energies is estimated to be due to the contribution of unresolved HESS-like sources. + This result is a lower limit for such a contribution because we have taken into account only sources delected above 6 per cent of the Crab fIux ancl because LESS sensitivity gels worse for extended sources. meaning that some extended sources might have been missed by HESS.," This result is a lower limit for such a contribution because we have taken into account only sources detected above 6 per cent of the Crab flux and because HESS sensitivity gets worse for extended sources, meaning that some extended sources might have been missed by HESS." + Using the logN-log$ relation we have also predicted the number of WESS-like sources which VERITAS. LESS and ILAWC should detect. during their survey of the sky.," Using the logN-logS relation we have also predicted the number of HESS-like sources which VERITAS, HESS and HAWC should detect during their survey of the sky." + An alternative procedure to evaluate the contribution of unresolved IIESS-like sources to Milagro diffuse enission gives (he diffuse flux due to unresolved sources comparable to the diffuse emission itself., An alternative procedure to evaluate the contribution of unresolved HESS-like sources to Milagro diffuse emission gives the diffuse flux due to unresolved sources comparable to the diffuse emission itself. + We finally constrained the slope of the luminosity function., We finally constrained the slope of the luminosity function. + The main uncertainty ol this ealeulation consists in assuming that the distribution of y-ray sources follows the distribution of either pulsars or SNRs observed in the radio., The main uncertainty of this calculation consists in assuming that the distribution of $\gamma$ -ray sources follows the distribution of either pulsars or SNRs observed in the radio. +" In particular. in order (to predict how many PSRs observed in the radio have a PWN and are possible eanma ray emitters we used the result that the spin-down energy loss dE/dl>/dl,=4x10erg/s [or a voung energetic pulsar to fori a PWN."," In particular, in order to predict how many PSRs observed in the radio have a PWN and are possible gamma ray emitters we used the result that the spin-down energy loss $dE/dt > dE/dt_c = 4 \times {10}^{36} erg/s$ for a young energetic pulsar to form a PWN." + In this respect we have ignored (he existence of pulsars. such as Geminga. which are y-ray loud. vet not observed in the radio.," In this respect we have ignored the existence of pulsars, such as Geminga, which are $\gamma$ -ray loud, yet not observed in the radio." + New observational results support the hvpothesis that a population of unresolved sources contribute significantly to the emission al very high energy., New observational results support the hypothesis that a population of unresolved sources contribute significantly to the emission at very high energy. + Milagro has recently reported the discovery of TeV gamma ray emission from the Cygnus Region of the Galaxy. which exceeds the predictions of conventional models of gamma -ray production (Abdoetal.2007) [rom the same region in the Galaxy where the Tibet Array has detected. an excess of cosmic ravs (Amenomorietal.," Milagro has recently reported the discovery of TeV gamma ray emission from the Cygnus Region of the Galaxy, which exceeds the predictions of conventional models of gamma -ray production \citep{Abdo:2006} from the same region in the Galaxy where the Tibet Array has detected an excess of cosmic rays \citep{Amenomori}." +2006).. Thanks to its improved sensitivity Milagro has also better imaged the whole Northern sky and discovered. [our sources and. [our source. candidates (Abdoetal.2007)., Thanks to its improved sensitivity Milagro has also better imaged the whole Northern sky and discovered four sources and four source candidates \citep{Abdo:2007}. +. HESS has seen very high οποιον emission spatially correlated. with eiat molecular clouds located in the Galactic Center (Aharonianetal.2006b)., HESS has seen very high energy emission spatially correlated with giant molecular clouds located in the Galactic Center \citep{Aharonian:nature}. +. The energy spectrum measured by ILESS close to the Galactic Center is E.77. significantly harder than the E.7 spectrum of the dilfuse emission and equal to the average spectrum of the IIESS source population.," The energy spectrum measured by HESS close to the Galactic Center is $E^{-2.3}$, significantly harder than the $E^{-2.7}$ spectrum of the diffuse emission and equal to the average spectrum of the HESS source population." + The emission [rom the Galactic Center might possibly unveil a cosmic rav accelerator., The emission from the Galactic Center might possibly unveil a cosmic ray accelerator. + To craw more definitive conclusions about the very high energy 5-ravy sky. new observations awe of fandamental importance.," To draw more definitive conclusions about the very high energy $\gamma$ -ray sky, new observations are of fundamental importance." + New hints will be provide bv both NAGIC and VERITAS. which already survev the Cvenus Region.," New hints will be provide by both MAGIC and VERITAS, which already survey the Cygnus Region." + Finally GLAST will investigate the window of enereyv between 10 MeV to 300 GeV. covering the enerev gap left between EGRET and the eround-hased low threshold gamma-ray observatories.," Finally GLAST will investigate the window of energy between 10 MeV to 300 GeV, covering the energy gap left between EGRET and the ground-based low threshold gamma-ray observatories." +of the emission [rom these objects.,of the emission from these objects. +" It could. be similar to that responsible for the ""giant. pulses! observed. [roni some pulsars (e.g.2)..", It could be similar to that responsible for the `giant pulses' observed from some pulsars \citep[e.g.][]{kbm+06}. + IH could also be that the sporadic emission is related to the fact that these objects are near the radio ‘deathline’ (e.g.77) and/or are examples of extreme nulling (e.g.?)..," It could also be that the sporadic emission is related to the fact that these objects are near the radio `death–line' \citep[e.g.][]{cr93,zgd07} and/or are examples of extreme nulling \citep[e.g.][]{rr09}." + The phenomenon has also been attributed to the presence of a circumstellar asteroid belt (?2) or a radiation belt as seen in planetary magnetospheres (?)..," The phenomenon has also been attributed to the presence of a circumstellar asteroid belt \citep{li06,cs08} or a radiation belt as seen in planetary magnetospheres \citep{lm07}." + Or. perhaps. some are transient. N-rav magnetars (e.g.?)..," Or, perhaps, some are transient X-ray magnetars \citep[e.g.][]{wkg+05}." + Another idea is that their properties lie at the extreme end of the population of normal radio pulsars., Another idea is that their properties lie at the extreme end of the population of normal radio pulsars. + Weltevrecde et al. (, Weltevrede et al. ( +2006) show that PSR. | 1H. a nearby micelleaged pulsar which emits pulses with energies many times its mean pulse energy. would be discovered as a RIUXE source if it were farther away.,"2006) \nocite{wsrw06} show that PSR $+$ 14, a nearby middle--aged pulsar which emits pulses with energies many times its mean pulse energy, would be discovered as a RRAT source if it were farther away." + RATS may also be considered as an extreme case of mode changing (see.e.g..72) where the on state is less than or about one pulse period.," RRATs may also be considered as an extreme case of mode changing \citep[see, e.g.,][]{wmj+06} where the on state is less than or about one pulse period." + Furthermore. ? have recently shown that many. pulsars exhibit a twostate phenomenon in which varving pulse profile shapes are correlated with variations in spindown rates and implied changes in magnetospheric particle density.," Furthermore, \citet{lhk+10} have recently shown that many pulsars exhibit a two--state phenomenon in which varying pulse profile shapes are correlated with variations in spin–down rates and implied changes in magnetospheric particle density." + These changes re quasiperiodic. with timescales ranging from one month o many vears.," These changes are quasi–periodic, with timescales ranging from one month to many years." + It could be that the RRATs are similar twoate svstems. in which the profile changes are so dramatic o make them undetectable in the more common state.," It could be that the RRATs are similar two--state systems, in which the profile changes are so dramatic to make them undetectable in the more common state." + Determining the time variability and/or perioclicity of j)0 RAAT pulses is therefore an important. cliagnostic of 1 RRAP emission. mechanism., Determining the time variability and/or periodicity of the RRAT pulses is therefore an important diagnostic of the RRAT emission mechanism. + While the pulse profile. and. in most cases. pulse intensity changes of ? are quasi»eriodic. the pulse intensity distributions of normal pulsars and giantpulsing pulsars are believed. to be random. over ime.," While the pulse profile, and, in most cases, pulse intensity changes of \citet{lhk+10} are quasi--periodic, the pulse intensity distributions of normal pulsars and giant–pulsing pulsars are believed to be random over time." + On the other hand. nulling pulsars in general show on and olf timescales of more than one consecutive pulse. indicating larecly nonrandom distributions (see.e.g..?)," On the other hand, nulling pulsars in general show on and off timescales of more than one consecutive pulse, indicating largely non–random distributions \citep[see, e.g.,][]{rr09}." + tadio emitting neutron stars often show transient spindown phenomena as well., Radio emitting neutron stars often show transient spin--down phenomena as well. + For instance. glitches. or sudden increases in the spin frequency. have been observed. from voung pulsars and one RILVE ( PSR JISIO. 1458).," For instance, glitches, or sudden increases in the spin frequency, have been observed from young pulsars and one RRAT ( PSR $-$ 1458)." + One of the elitehes from PSIUJISI9 1458 was accompanied by a 3.5 increase in the pulse detection rate (7)..., One of the glitches from PSR $-$ 1458 was accompanied by a $\sigma$ increase in the pulse detection rate \citep{lmk09}. + Racliative events do not normally accompany the elitches of normal radio pulsars. but are quite common for magnetars (?)..," Radiative events do not normally accompany the glitches of normal radio pulsars, but are quite common for magnetars \citep{dkg08}." + ‘This. along with the high magnetic field of PSR 145s. hints at a relationship with magnetars ancl also provides additional motivation to examine the pulse rate variations with time for all RRATS.," This, along with the high magnetic field of PSR $-$ 1458, hints at a relationship with magnetars and also provides additional motivation to examine the pulse rate variations with time for all RRATs." + We search for periodicities and quantify the randomness of the detected. RRAT pulses in several different wavs., We search for periodicities and quantify the randomness of the detected RRAT pulses in several different ways. + We first. search Lor periodicities in the pulse arrival times on minutesvear long time scales and pulse detection rates on monthvear long time scales using a LombScarele analysis., We first search for periodicities in the pulse arrival times on minutes–year long time scales and pulse detection rates on month–year long time scales using a Lomb–Scargle analysis. + We then quantify the rancominess of the RAAT pulse arrival times using Ixolmogorov.Smirnov tests on seconds.vear long time scales., We then quantify the randomness of the RRAT pulse arrival times using Kolmogorov–Smirnov tests on seconds–year long time scales. + The observations are described in Section 2. the methods and results in Section 3.λ and the conclusions and plans for future work in Section 4.," The observations are described in Section 2, the methods and results in Section 3, and the conclusions and plans for future work in Section 4." + All cight sources. discussed. in this paper were discovered bv AleLaughlin ct al. (, All eight sources discussed in this paper were discovered by McLaughlin et al. ( +2006) in à reanalvsis of data from the Parkes Multibeam Pulsar Survey (PAIRS).,2006) in a re–analysis of data from the Parkes Multi–beam Pulsar Survey (PMPS). + We have ignored three of the original 11 RIRATS as their pulse detection rates are too low to perform this analvsis., We have ignored three of the original 11 RRATs as their pulse detection rates are too low to perform this analysis. + The discovery data were taken between Jan 1998 and Feb 2002 and followup observations began in Aue 2003 and are ongoing using the αι Parkes telescope., The discovery data were taken between Jan 1998 and Feb 2002 and follow–up observations began in Aug 2003 and are ongoing using the 64-m Parkes telescope. + Most of the observations used. the central beam. of the multibeam receiver with a central frequeney of 1.4. Cllz and a bandwidth of 256 MlIIZ., Most of the observations used the central beam of the multi--beam receiver with a central frequency of 1.4 GHz and a bandwidth of 256 MHz. + A few observations used. other [requencies: we ignore these in our analysis to ensure uniformity of pulse detection rates., A few observations used other frequencies; we ignore these in our analysis to ensure uniformity of pulse detection rates. + The sources have been observed at between 27 and SO epochs at 1.4 Gllz. with each observation 2 hr in duration (see Table 1).," The sources have been observed at between 27 and 89 epochs at 1.4 GHz, with each observation $-$ 2 hr in duration (see Table 1)." + One important consideration in our analvsis is the influence of the interstellar medium. on. our. observed pulses., One important consideration in our analysis is the influence of the interstellar medium on our observed pulses. + For all of these sources. the predicted: dillractive scintillation bandwidths at. 1.4 CGlLlz are less than 1 MllIz (?).. making modulation. due to dillractive scintillation unimportant.," For all of these sources, the predicted diffractive scintillation bandwidths at 1.4 GHz are less than 1 MHz \citep{cl02}, making modulation due to diffractive scintillation unimportant." + In. Table. 1. we list the predicted: timescales [or refractive. seintillation at our. observing frequency of LA Cllz. estimated from the precieted dillractive scintillation timescales ancl bancdwidths [rom 7? (sec. og. ?)).," In Table 1, we list the predicted timescales for refractive scintillation at our observing frequency of 1.4 GHz, estimated from the predicted diffractive scintillation timescales and bandwidths from \citet{cl02} (see, e.g., \citet{lk05}) )." + These timescales range from 21 to 197 clays., These timescales range from 21 to 197 days. + However. the actual timescales could vary significantly from those predicted.," However, the actual timescales could vary significantly from those predicted." + The predicted: modulation indices due to refractive scintillation (7). range from 0.09 to 0.17. meaning these are expected to be relatively minor contributions to pulse rate variations.," The predicted modulation indices due to refractive scintillation \citep{lk05} range from 0.09 to 0.17, meaning these are expected to be relatively minor contributions to pulse rate variations." + Pulse detection is performed by dedispersing the data at the ispersion measure (DM) of the ΜΑ and at a DM of zero., Pulse detection is performed by dedispersing the data at the dispersion measure (DM) of the RRAT and at a DM of zero. + Then pulses are searched for in both time series above a Sa weshold. using the pulsar processing packageSIGPROCH., Then pulses are searched for in both time series above a $\sigma$ threshold using the pulsar processing package. +. Pulses which are brighter at the DM of the RRA are likely o be from the source., Pulses which are brighter at the DM of the RRAT are likely to be from the source. + We inspect the pulses visually by ‘heeking for pulse shape ancl pulse phase consistency. to » certain Of their astrophysical nature., We inspect the pulses visually by checking for pulse shape and pulse phase consistency to be certain of their astrophysical nature. + For some epochs which have large amounts of radio frequency interference. we applied the above procedure but with multiple trial DAIs as described by 2..," For some epochs which have large amounts of radio frequency interference, we applied the above procedure but with multiple trial DMs as described by \citet{mlk+09}." + ΗΕ more than one pulse is detected within an observation. a second check based on the known period of the source can be mace by requiring that all pulses have arrival times which diller by integer multiples of the period.," If more than one pulse is detected within an observation, a second check based on the known period of the source can be made by requiring that all pulses have arrival times which differ by integer multiples of the period." + For the sources with phaseconnected timing solutions. we check that the pulse arrival time is consistent with the solution.," For the sources with phase–connected timing solutions, we check that the pulse arrival time is consistent with the solution." + In Table 1. we list the number of epochs for which pulses were detected. for all sources.," In Table 1, we list the number of epochs for which pulses were detected for all sources." + In. Table 3 we list the total number of pulses detected. within the entire time span of observations., In Table 3 we list the total number of pulses detected within the entire time span of observations. + The LombSearele test (Scargle et al., The Lomb–Scargle test (Scargle et al. + 1982) is a statistical procedure for uncovering perioclic signals hidden in noise., 1982) is a statistical procedure for uncovering periodic signals hidden in noise. + We use this technique in our analysis as our cata are unevenly sampled. thereby making. standard Fourier," We use this technique in our analysis as our data are unevenly sampled, thereby making standard Fourier" +thickness of the “warm” inner edge is ~ 5 AU at R = 45 AU.,thickness of the “warm” inner edge is $\sim$ 5 AU at R = 45 AU. +" For the UV to penetrate the inner 5 AU, the disk would need contain <0.1% of the original amount of dust present to erase the effect seen here."," For the UV to penetrate the inner 5 AU, the disk would need contain $< 0.1\%$ of the original amount of dust present to erase the effect seen here." +" We note however that settling will also increase the depth of the heating, which will in principle thicken the extent of the wall further."," We note however that settling will also increase the depth of the heating, which will in principle thicken the extent of the wall further." + An underlying feature of the disk model is the presence of a large inner “void” which allows the outer disk to be heated directly by the star.," An underlying feature of the disk model is the presence of a large inner “void"" which allows the outer disk to be heated directly by the star." + One possible explanation for such a void is grain growth into rocky planetesimals (e.g.Skrutskieetal., One possible explanation for such a void is grain growth into rocky planetesimals \citep[e.g.][]{skrutskie1990}. + Thus one could postulate the existence of a small 1990)..amount of undetected dust inside the gap., Thus one could postulate the existence of a small amount of undetected dust inside the gap. +" Indeed some models infer the presence of some moderate mass of silicates in the inner disk (e.g.Calvetetal.2002,2005;Espaillat"," Indeed some models infer the presence of some moderate mass of silicates in the inner disk \citep[e.g.][]{calvet2002,calvet2005,espaillat2007}." +" This material has the ability to shadow the outer disk 2007)..from being directly heated by the star, thus inhibiting the presence of the warm molecular interface presented here."," This material has the ability to shadow the outer disk from being directly heated by the star, thus inhibiting the presence of the warm molecular interface presented here." +" First, one can ask how much dust would be required to cause the gap to become opaque (rt~ 10) to optical radiation peaking at A=0.7 wm as is the case for a T=4000 K star."," First, one can ask how much dust would be required to cause the gap to become opaque $\tau \sim10$ ) to optical radiation peaking at $\lambda=0.7$ $\mu$ m as is the case for a $T=4000$ K star." +" If we assume that the inner disk follows the outer disk profile (Section 2.1)), a total mass in dust of Mau=107? Mmoon inside 45 AU would be sufficient for the midplane to become optically thick to the optical heating radiation."," If we assume that the inner disk follows the outer disk profile (Section \ref{sec:mod}) ), a total mass in dust of $_{\rm dust}=10^{-3}$ $_{\rm Moon}$ inside 45 AU would be sufficient for the midplane to become optically thick to the optical heating radiation." +" While this is an extremely small amount of dust, one must also consider the scale height of the material: h~3.3 AU at R — 45 AU."," While this is an extremely small amount of dust, one must also consider the scale height of the material: $h\sim3.3$ AU at R = 45 AU." +" Therefore, this cross-section is very small, and would not strongly shadow the outer disk and would not significantly alter the heating and UV irradiation except at the very central midplane region."," Therefore, this cross-section is very small, and would not strongly shadow the outer disk and would not significantly alter the heating and UV irradiation except at the very central midplane region." +" Furthermore, because the original model from Sauteretal.(2009) included scattered light images, the presence of small grains at or above the disk scale height would significantly contribute to the amount of scattered light and not match observations for the model presented"," Furthermore, because the original model from \citet{sauter2009} included scattered light images, the presence of small grains at or above the disk scale height would significantly contribute to the amount of scattered light and not match observations for the model presented" +width limit.,width limit. + Llowever. if the anti-correlation of bias and equivalent width holds. then we would expect our bias measurement to lic slightly above the generated curve. as it does.," However, if the anti-correlation of bias and equivalent width holds, then we would expect our bias measurement to lie slightly above the generated curve, as it does." + Ehe amplitude of the disagreement between the model curves and our bias measurement. may thus appear slightly exaggerated in this plot., The amplitude of the disagreement between the model curves and our bias measurement may thus appear slightly exaggerated in this plot. + Still. any correction due to the equivalent width limit of the absorber sample is unlikely to change the measurement bevond the Large. stated errors.," Still, any correction due to the equivalent width limit of the absorber sample is unlikely to change the measurement beyond the large, stated errors." + As detailed in Tinker&Chen(2010).. the physical implications of a model without mass evolution require the σας radius of a halo. A. to expand relative to the halo virial radius. fogg. with increasing recdshift as ὃν this calculation. the gas radius should be just greater than half of the DAL halo virial radius at z=1.," As detailed in \citet{TC10}, the physical implications of a model without mass evolution require the gas radius of a halo, $R_{g}$, to expand relative to the halo virial radius, $R_{200}$, with increasing redshift as By this calculation, the gas radius should be just greater than half of the DM halo virial radius at z=1." + This implies the eas radius increases by ~40% in units of Boog between z=0.6 and z=1., This implies the gas radius increases by $\sim$ in units of $R_{200}$ between z=0.6 and z=1. + The errors. on our absorber bias measurement are larecly due to the small number of absorbers available in this analysis. Thus. better constraining the absorber-ealaxy cross-correlation measurement requires a deeper survey in the DEEP? region.," The errors on our absorber bias measurement are largely due to the small number of absorbers available in this analysis, Thus, better constraining the absorber-galaxy cross-correlation measurement requires a deeper survey in the DEEP2 region." + The size of the sample in this work is largely determined by the SNR ancl resolution of SDSS spectra., The size of the sample in this work is largely determined by the SNR and resolution of SDSS spectra. + Follow-up observations of the same SDSS quasars with either. higher SNR. higher resolution. or both might casily increase the number of detections. without necessitating the detection. of many additional quasars within the survey footprint.," Follow-up observations of the same SDSS quasars with either higher SNR, higher resolution, or both might easily increase the number of detections, without necessitating the detection of many additional quasars within the survey footprint." + Increasing he size of the DEEP? data by a factor of 4. might reduce the error on the bias measurement sulliciently o rule out the Tinkeretal.(2010) model of evolving halo miss.," Increasing the size of the DEEP2 data by a factor of 4, might reduce the error on the bias measurement sufficiently to rule out the \citet{Tinker10} model of evolving halo mass." + A lareer dataset may also enable measurements of bias or subsets of the absorber sample split by any number of observable parameters (c.g. equivalent width).," A larger dataset may also enable measurements of bias for subsets of the absorber sample split by any number of observable parameters (e,g, equivalent width)." + Since o»evious analyses have identified a weak anti-correlation of absorber equivalent width ancl bias (Bouchéctal.2006:Lundgrenetal.2000:Gauthier 2009).. it would. be interesting to examine whether any evolution in this trend is evident higher redshift.," Since previous analyses have identified a weak anti-correlation of absorber equivalent width and bias \citep{B06,L09,Gauthier09}, it would be interesting to examine whether any evolution in this trend is evident higher redshift." + Due to the limitations of our current data. we must save this potentially interesting analysis for future work.," Due to the limitations of our current data, we must save this potentially interesting analysis for future work." + The average elfective gas radius of ~ 40 hi. κρο estimated rom our measurement of the Alell covering fraction in DEEP? galaxies. agrees well with the speculated extent of star-formation driven outllows from the same galaxy population. which was estimated. by Weineretal.(2009) o be 20-50 tkpe.," The average effective gas radius of $\sim$ 40 $^{-1}$ kpc, estimated from our measurement of the MgII covering fraction in DEEP2 galaxies, agrees well with the speculated extent of star-formation driven outflows from the same galaxy population, which was estimated by \citet{Weiner09} to be 20-50 $^{-1}$ kpc." +" Early studies found that galaxies with Luminosities of —L* have a Wh???OSA covering [raction of y25zf, l on scales less than 60 !kpe (Bergeron 1994)."," Early studies found that galaxies with luminosities of $\sim L^{*}$ have a $_{r}^{\lambda2796}>$ covering fraction of $\leq f_{c} \leq$ 1 on scales less than 60 $^{-1}$ kpc \citep{BB91, Bechtold92, Steideletal94}." +". These analysesrevealed an anti-correlation between R,/ and AXPTOWY? which: was also shown to scale slightly.S with. galaxy luminosity."," These analysesrevealed an anti-correlation between $R_{g}$ and $^{\lambda2796}_{r}$, which was also shown to scale slightly with galaxy luminosity." + An additional dependence of f[. and Hi on galaxy type was suggested by the results of(1905).. who used LIST to examine z«0.2 galaxies in close proximity to quasar sightlines.," An additional dependence of $f_{c}$ and $R_{g}$ on galaxy type was suggested by the results of, who used HST to examine $<$ 0.2 galaxies in close proximity to quasar sightlines." + Bowenetal.(1995) measured an cllective gas radius Ry50 tkpe for galaxies with late-type or interacting morphologies and found an absence of absorption for the earlv-tvpe galaxies with high confidence., \citet{Bowen95} measured an effective gas radius $R_{g}\sim$ 50 $^{-1}$ kpc for galaxies with late-type or interacting morphologies and found an absence of absorption for the early-type galaxies with high confidence. + Due to the varving choices of galaxy samples. Wer? andmecian absorber redshift. the covering fraction measurements from numerous other studies are dillicult to directly. compare with our results.," Due to the varying choices of galaxy samples, $^{\lambda2796}_{r,min}$ , andmedian absorber redshift, the covering fraction measurements from numerous other studies are difficult to directly compare with our results." + Tripp&Bowen(2005) measured a covering [fraction of fi.~0.5 within 60h. kpe for WMESOLA absorbers at z~0.5., \citet{TB05} measured a covering fraction of $f_{c}\sim$ 0.5 within 60 $^{-1}$ kpc for $_{r}^{\lambda2796}>$ absorbers at $\sim$ 0.5. +" Chen&Tinker(2008) measure a covering fraction of {~0.8 within 69 tkpe for WP?50,3. which seems to follow a scaling relation dependent on halo mass ancl galaxy Luminosity."," \citet{CT08} measure a covering fraction of $f_{c}\sim$ 0.8 within 69 $^{-1}$ kpc for $_{r}^{\lambda2796}>$, which seems to follow a scaling relation dependent on halo mass and galaxy luminosity." + In. general agreement with this result. Barton&Cooke(2009) measure 0.25«f.«0.4 for N270.3 within ~ 755 tkpe at z0.1.," In general agreement with this result, \citet{BC09} measure $$ within $\sim$ 75 $^{-1}$ kpc at $\sim$ 0.1." + A number of additional. measurements. motivated by recent LELRG correlation results. have investigated the cold gas content of LARC haloes (Chenetal.2010:Gau-thieretal.2010:Bowen&Chelouche 2011).. which generally agree on a WTSO.6AMell coveringe fraction of ~10% for 100 tkpe radii aroundὃν LRGs.," A number of additional measurements, motivated by recent -LRG correlation results, have investigated the cold gas content of LRG haloes \citep{Chen10, Gauthier10, BC11}, which generally agree on a $_{r}^{\lambda2796}\geq$ covering fraction of $\sim$ for 100 $^{-1}$ kpc radii around LRGs." + However. Bowen&Chelouche(2011). find little evidence to support a trend: of decreasing. f; on the same scales with increasing galaxy luminosity. which might be expected if gas is evaporated at larger radii in the haloes of more luminous ealaxies (Vinker&Chen2008:etal.2010).," However, \citet{BC11} find little evidence to support a trend of decreasing $f_{c}$ on the same scales with increasing galaxy luminosity, which might be expected if gas is evaporated at larger radii in the haloes of more luminous galaxies \citep{TC08, Chen10}." +. Unfortunately. the number of absorber-galaxy pairs at separations less than 100 tkpe in this work is insullicient to make a strong argument in support of anv of these previous analvses.," Unfortunately, the number of absorber-galaxy pairs at separations less than 100 $^{-1}$ kpc in this work is insufficient to make a strong argument in support of any of these previous analyses." + Future surveys to obtain quasar spectra with higher resolution and. SNR would be most. useful in order to better constrain the eas covering fraction of these ealaxies., Future surveys to obtain quasar spectra with higher resolution and SNR would be most useful in order to better constrain the gas covering fraction of these galaxies. + Such a dataset would also allow for the extension of the covering fraction measurement to lower absorber equivalent width limits., Such a dataset would also allow for the extension of the covering fraction measurement to lower absorber equivalent width limits. + We measure the two-point cross-correlation function of 21 absorbers detected in the SDSS DIU. with 732.000 spectroscopic galaxies from the DEEP2 galaxy survey in the redshift range O.77/2 one can expand integraud in CÀ2)) in terms ofz/2 O0) axe illminated ouly bydistant parts of disk.equation οsindith so that 1c9 (while resin’~1) in equation (À2)).," Polar regions of the star $\theta\to 0$ ) are illuminated only bydistant parts of the disk, $R\gtrsim R_\star/\sin\theta\gg +R_\star$, so that $x\gg 1$ (while $x\sin\theta\sim 1$ ) in equation \ref{eq:1D}) )." + Also. far from the star one can safely use (6)) w f=1 to Παν arrive at the equation (11)) with This result is independent of the structure of the boundary laver near the stellar surface since the polar regions of the star do not have direct sight lines to the boundary laver.," Also, far from the star one can safely use equation \ref{eq:vis_dissip}) ) with $f=1$ to finally arrive at the equation \ref{eq:as}) ) with This result is independent of the structure of the boundary layer near the stellar surface since the polar regions of the star do not have direct sight lines to the boundary layer." + This is evidenced by the convergence at ϐ»0 of the two curves iu Figure 2. calculated assuming f(A)= land f(R)=1ΠΠ).," This is evidenced by the convergence at $\theta \to 0$ of the two curves in Figure \ref{fig:irr_flux} + calculated assuming $f(R)=1$ and $f(R)=1-(R_\star/R)^2$ ." + To determine the validity μπιτςof the 1D solution for the structure of the radiative zone found ‘ousidoringin 8.1. we evaluate the magnitude of the corrections arising when the latitudinal radiative transter is accounted for., To determine the validity limits of the 1D solution for the structure of the radiative zone found in \ref{sect:1D} we evaluate the magnitude of the corrections arising when the latitudinal radiative transfer is accounted for. + € LD solution (35)) as a zeroth-order approximation we plug it iuto the full equation (21)) aud carefully expand all 0-derivatives. romieniberime that P is almost indepeudent of 0 (latitudinal pressure eradicuts are small).," Considering 1D solution \ref{eq:PofT}) ) as a zeroth-order approximation we plug it into the full equation \ref{eq:rad_tran}) ) and carefully expand all $\theta$ -derivatives, remembering that $P$ is almost independent of $\theta$ (latitudinal pressure gradients are small)." +" Tuteerating the resultant expression once over kr We again arrive at tlhe equation (26)) but with £i,>δει in the left-hand side. where Lory OZ(PjdP. ere Z(GP)— Lis a weak function of pressure (varving by at most a factor ~ 1) and Ly is a characteristic scale of latitudinal variation of Ly=ΠινΠατν0]! "," Integrating the resultant expression once over $r$ we again arrive at the equation \ref{eq:flux}) ) but with $F_{in}\to F_{in}+\delta F_{in}$ in the left-hand side, where )^2 Z(P)dP. Here $Z(P)\sim 1$ is a weak function of pressure (varying by at most a factor $\sim 1$ ) and $L_\theta$ is a characteristic scale of latitudinal variation of $T_{ph}$, $L_\theta = R_\star(\partial \ln T_{ph}/\partial \theta)^{-1}$." +"Ow DID approximation is justified if the correction to tlhe 1D result à, is iuall comparedTjj. to £j, given by equation (55)).", Our 1D approximation is justified if the correction to the 1D result $\delta F_{in}$ is small compared to $F_{in}$ given by equation \ref{eq:F_in}) ). +" Tuteeral iu (D2)) attaius its highest value at P~Py, (latitudiual radiation transfer is easiest in tlie upper. low density lavers of the star just below the photosphere) aud one can casily find using equations (52)). (53)). aud (55)) that where is a photospheric scale height."," Integral in \ref{eq:deltaF}) ) attains its highest value at $P\sim P_{ph}$ (latitudinal radiation transfer is easiest in the upper, low density layers of the star just below the photosphere) and one can easily find using equations \ref{eq:P_ph}) ), \ref{eq:nab_ph}) ), and \ref{eq:F_in}) ) that , where $H_{ph}$ is a photospheric scale height." + This result makes it clear that the 1D solution for tho structure of the radiative [1οzone should be reasonableas loug as the coucition (67)) is fulfilled., This result makes it clear that the 1D solution for the structure of the radiative zone should be reasonableas long as the condition \ref{eq:validity}) ) is fulfilled. +p3.,$p_3$. +" This is because, with the exception of X Per, the observed pa are higher than the curve resulting from the steady state of fformation and destruction using these DR rate coefficients, and inclusion of the ++ rreaction further drives p3 toward the value expected for thermal equilibrium."," This is because, with the exception of X Per, the observed $p_3$ are higher than the curve resulting from the steady state of formation and destruction using these DR rate coefficients, and inclusion of the + reaction further drives $p_3$ toward the value expected for thermal equilibrium." +" Recent storage ring experiments by Tometal.(2009) and Kreckeletal.(2005) both saw an increased DR cross-section when iis produced fromp-H2;; however, recentHi imaging results presented in Kreckeletal.(2010) suggest that the ions in these experiments have been heated during extraction from the ion sources, and the difference between the aand mmay p-Hjtherefore have been overestimated."," Recent storage ring experiments by \cite{tom2009} and \cite{kreckel2005} both saw an increased DR cross-section when is produced from; however, recent imaging results presented in \cite{kreckel2010} suggest that the ions in these experiments have been heated during extraction from the ion sources, and the difference between the and may therefore have been overestimated." +" Further experimental work is clearly needed to pin down the enhancement in DDR, and confirmation(ifany) of p-Hjthe theoretical predictions would also be quite helpful."," Further experimental work is clearly needed to pin down the enhancement (if any) in DR, and confirmation of the theoretical predictions would also be quite helpful." +" To summarize, according to our models the reaction of iis expected to effectively thermalize the nuclear spin configurations of aat steady-state, provided that sufficient collisions occur within the lifetime of an."," To summarize, according to our models the reaction of with is expected to effectively thermalize the nuclear spin configurations of at steady-state, provided that sufficient collisions occur within the lifetime of an." +". In diffuse molecular clouds, however, the average numberH3 of reactive collisions with ssuffered by an iis small, indicatingHj that the formation and destruction rates of the two nuclear spin species should be important."," In diffuse molecular clouds, however, the average number of reactive collisions with suffered by an is small, indicating that the formation and destruction rates of the two nuclear spin species should be important." + A more complete model which takes these factors into account reaches reasonable agreement with observations in 4 of 5 sight lines provided $4 is on the order of 0.9 and aand aare destroyedp-Hj at equal rates owing to DR., A more complete model which takes these factors into account reaches reasonable agreement with observations in 4 of 5 sight lines provided $S^{id}$ is on the order of 0.9 and and are destroyed at equal rates owing to DR. +" Reconciling the observations with the spin-dependent theoretical rates of dosSantosetal.(2007) is difficult, and accurate experimental measurements of the spin-dependent DR rates of aat low temperature are needed."," Reconciling the observations with the spin-dependent theoretical rates of \cite{santos2007} is difficult, and accurate experimental measurements of the spin-dependent DR rates of at low temperature are needed." +" While all evidence seems to suggest that Jo, inferred from ultraviolet spectroscopy of Hz accurately reflects the kinetic temperature of diffuse molecular clouds, the observed excitation temperature of Hj is clearly non-thermal in 4 of the 5 measured sight lines."," While all evidence seems to suggest that $T_{01}$ inferred from ultraviolet spectroscopy of $_2$ accurately reflects the kinetic temperature of diffuse molecular clouds, the observed excitation temperature of $_3^+$ is clearly non-thermal in 4 of the 5 measured sight lines." +" Based on the microcanonical model of Park&Light(2007),, we have constructed a steady state model to predict the para-Hj fraction (ps) if reactive collisions between Hi and H» control the spin modifications of Hj."," Based on the microcanonical model of \cite{park2007}, we have constructed a steady state model to predict the $_3^+$ fraction $p_3$ ) if reactive collisions between $_3^+$ and $_2$ control the spin modifications of $_3^+$." +" Those results show pa slightly below the limit expected for full thermalization, and far from the observations."," Those results show $p_3$ slightly below the limit expected for full thermalization, and far from the observations." +" However, à steady state model that incorporates both the Hj + He reaction as well as the Hj formation (following cosmic-ray ionization) and destruction (by electron recombination) can reproduce the observed ps if the reactive collision rate is somewhat slow and the dissociative recombination rates for aand HHj are comparable."," However, a steady state model that incorporates both the $_3^+$ + $_2$ reaction as well as the $_3^+$ formation (following cosmic-ray ionization) and destruction (by electron recombination) can reproduce the observed $p_3$ if the reactive collision rate is somewhat slow and the dissociative recombination rates for and $_3^+$ are comparable." +" Our interpretation, given the currently available data, is that Hi suffers relatively few spin-changing collisions with H» in its lifetime, and is thus incompletely equilibrated by this reaction."," Our interpretation, given the currently available data, is that $_3^+$ suffers relatively few spin-changing collisions with $_2$ in its lifetime, and is thus incompletely equilibrated by this reaction." + The observed para-Hj fraction therefore lies between the nascent fraction and the nearly-thermal fraction that would be reached with sufficient reactive collisions., The observed $_3^+$ fraction therefore lies between the nascent fraction and the nearly-thermal fraction that would be reached with sufficient reactive collisions. +" If our model is correct (and the spin-dependent DR rates of aare nearly equal at low temperature), this marks Hithe first determination of the reactive rate coefficient of the H3++ rreaction, and suggests a value on the order of 10:10 cm? g-1,"," If our model is correct (and the spin-dependent DR rates of are nearly equal at low temperature), this marks the first determination of the reactive rate coefficient of the + reaction, and suggests a value on the order of $10^{-10}$ $^{-3}$ $^{-1}$." +" Fully quantum reactive scattering calculations of the Hj + Hg reaction would be highly desirable, as they would pin down the state-to-state rate coefficients needed to predict the interstellar para-Hij fraction."," Fully quantum reactive scattering calculations of the $_3^+$ + $_2$ reaction would be highly desirable, as they would pin down the state-to-state rate coefficients needed to predict the interstellar $_3^+$ fraction." + Further experiments and theoretical calculations to elucidate the dependence (if any) of the dissociative recombination on, Further experiments and theoretical calculations to elucidate the dependence (if any) of the dissociative recombination on +The result is the emission of two photons below the Lya energy that do not further interact with the eas.,The result is the emission of two photons below the $\alpha$ energy that do not further interact with the gas. + With our code we are therefore able το separate the Lye radiation from the less energetic photons produced. by cascade from the 2s level., With our code we are therefore able to separate the $\alpha$ radiation from the less energetic photons produced by cascade from the $2s$ level. + ALL Lye photons that result from collisional excitations (or Lye photons. see Chen Miralda Eseudé 2004) have a cooling elect on the gas.," All $\alpha$ photons that result from collisional excitations (or $\alpha$ photons, see Chen Miralda $\acute{e}$ 2004) have a cooling effect on the gas." + We also included. in our caleulation processes. that can produce continuum photons. such as recombinations ancl Bremsstrahlune free-[ree interactions of electrons with ionized atoms.," We also included in our calculation processes that can produce continuum photons, such as recombinations and Bremsstrahlung free-free interactions of electrons with ionized atoms." + This effects. are negligible as we will see and there is virtually no production. of photons between the Lyman-a and the Lyman-:? resonances., This effects are negligible as we will see and there is virtually no production of photons between the $\alpha$ and the $\beta$ resonances. + This radiation would. have redshifted into Lya photons (Chen Maralda Escudé 2004) and by entering the Lvo resonance from its blue wing would have heated the gas instead., This radiation would have redshifted into $\alpha$ photons (Chen Miralda $\acute{e}$ 2004) and by entering the $\alpha$ resonance from its blue wing would have heated the gas instead. + Electron-clectron collisions between secondary and thermal electrons were implemented. as in SVSs5 and S79 and we treated similarly the energy clistribution of secondary electrons following collisional ionization of LE and He. so we refer the reader to those works for a detailed explanation.," Electron-electron collisions between secondary and thermal electrons were implemented as in SVS85 and S79 and we treated similarly the energy distribution of secondary electrons following collisional ionization of H and He, so we refer the reader to those works for a detailed explanation." + As we mentioned. earlier. we found that 1000. Monte Carlo realizations gave Consistent results.," As we mentioned earlier, we found that 1000 Monte Carlo realizations gave consistent results." + Going from 50 to 1000 realizations changed the averaged. values by less than 5% and the respective σ by less than 104., Going from 50 to 1000 realizations changed the averaged values by less than $5\%$ and the respective $\sigma$ by less than $10\%$. + We deemed. 1000 to be a sullicient number of realizations for a stable result., We deemed 1000 to be a sufficient number of realizations for a stable result. + The 1o values of the energy deposition fractions are included in Table 1 and they vary from 2.9% to 0.044 of the total amount of energy., The $1 \sigma$ values of the energy deposition fractions are included in Table 1 and they vary from $2.9 \%$ to $0.04\%$ of the total amount of energy. + For our purposes it is important to use the best available cross sections for the several interaction channels considered in our calculation., For our purposes it is important to use the best available cross sections for the several interaction channels considered in our calculation. + We will give here the references rather than entering in the description of physical details that are not the scope of this Letter., We will give here the references rather than entering in the description of physical details that are not the scope of this Letter. + The cross sections (ση) for collisional ionization of HI. He. He]. were taken from Wim Ruelel (1994). Shah et al. (," The cross sections $\sigma_i$ ) for collisional ionization of H, He, He+ were taken from Kim Rudd (1994), Shah et al. (" +LOST). Shah οἱ al. (,"1987), Shah et al. (" +1988).,1988). +" A simple functional fit is where c=Lj,ο is the ratio between the incoming electron energy and the binding cnerey of the atomic electron.", A simple functional fit is where $x=E_{kin}/E_B$ is the ratio between the incoming electron energy and the binding energy of the atomic electron. + Thecollisional excitation cross sections of 11 ancl Le are from Ixim Desclaux (2002). while for the excitation to the 2s level of LE we used the work from Branscen Noble (1976).," Thecollisional excitation cross sections of H and He are from Kim Desclaux (2002), while for the excitation to the $2s$ level of H we used the work from Bransden Noble (1976)." +" The cross section for Coulomb collisions between electrons is [rom Spitzer Scott (1969) as in STO. while the free-Lree cross section (σε) for electrons interacting with protons is given by the DBethe-Lleitler quantum-mechanical )orn approximated. result. (see c.g. Llaug 1997). where m, the electron mass. ro is the classical electron radius. à is the fine structure constant. &=me is the photon energy in units of er. and. p;. pr are the momenta of the incident and scattered. electron respectively. again in units of mec."," The cross section for Coulomb collisions between electrons is from Spitzer Scott (1969) as in S79, while the free-free cross section $\sigma_{ff}$ ) for electrons interacting with protons is given by the Bethe-Heitler quantum-mechanical Born approximated result (see e.g. Haug 1997), where $m_e$ the electron mass, $r_0$ is the classical electron radius, $\alpha$ is the fine structure constant, $k = E/mc^2$ is the photon energy in units of $m_e c^2$, and $p_i$, $p_f$ are the momenta of the incident and scattered electron respectively, again in units of $m_e c^2$." + As mentioned before we also included the recombination cross section σε. to take into account a process which could produce. continuum photons.," As mentioned before we also included the recombination cross section $\sigma_r$, to take into account a process which could produce continuum photons." +" Neglecting helium recombinations. we have that lor hydrogen where gy, is the Gaunt [actor of O(1). v is the emitted radiation frequency. 0, dis the electron velocity ancl η is the level at which the electron recombines."," Neglecting helium recombinations, we have that for hydrogen where $g_{fb}$ is the Gaunt factor of $\cal O$ (1), $\nu$ is the emitted radiation frequency, $v_e$ is the electron velocity and $n$ is the level at which the electron recombines." + Our results are summarized in Table 1 and in Figure 1., Our results are summarized in Table 1 and in Figure 1. + In the Table we report the fraction of the energy of a 10 keV primary electron. which. for cillerent values of the gas ionized fraction wr. is deposited into heat. Lye excitations. ionizations and photons with 47« 10.2 eV. The errors correspond the standard deviation from the mean calculated over 1000 Monte Carlo realizations of the experiment.," In the Table we report the fraction of the energy of a 10 keV primary electron which, for different values of the gas ionized fraction $x_e$, is deposited into heat, $\alpha$ excitations, ionizations and photons with $E <$ 10.2 eV. The errors correspond the standard deviation from the mean calculated over 1000 Monte Carlo realizations of the experiment." + The energy fraction that goes into heating grows rapidly as the gas ionization fraction. value becomes higher., The energy fraction that goes into heating grows rapidly as the gas ionization fraction value becomes higher. + The physical reason of this behavior is that clectron-clectron interactions become dominant for high values of wr., The physical reason of this behavior is that electron-electron interactions become dominant for high values of $x_e$. + We performed our calculations for 3 keV and. 10 keV o account for the theoretically aad observationally inferred range of the sterile neutrino mass and found that the energy paction distributed among the dilferent. processes remains similar in both cases (within 2! ), We performed our calculations for 3 keV and 10 keV to account for the theoretically and observationally inferred range of the sterile neutrino mass and found that the energy fraction distributed among the different processes remains similar in both cases (within $2\%$ ). + This is in agreement with he results from $79 and SVSS5 that find thatfor primary electrons with energies higher than 100 eV. the fractional energy depositions converge rapidly to a common behavior., This is in agreement with the results from S79 and SVS85 that find thatfor primary electrons with energies higher than 100 eV the fractional energy depositions converge rapidly to a common behavior. + We eave à convenient. functional form to our results by fitting the data in Table 1 with an accuracy greater than Jom De)., We gave a convenient functional form to our results by fitting the data in Table 1 with an accuracy greater than $3.5$ . + Our results ciller sensibly from those of the past stuclics. with cillercnces as high as for some values of wr.," Our results differ sensibly from those of the past studies, with differences as high as for some values of $x_e$ ." + The, The +spectral type box. including à + 1000 K uncertainty for the spectral type definitions.,"spectral type box, including a $\pm$ 1000 K uncertainty for the spectral type definitions." + Note that the models show considerable jumps in and luminosity when the stars enter the Wolf-Rayet phase., Note that the models show considerable jumps in and luminosity when the stars enter the Wolf-Rayet phase. + These extremely fast crossings (< 50000 years) through the HR-Diagram are not included in the tables., These extremely fast crossings $<$ 50000 years) through the HR-Diagram are not included in the tables. + The newly arrived spectral-type-mass relation. as arrived in Sect. 2..," The newly arrived spectral-type-mass relation, as arrived in Sect. \ref{sec:def}," + is now compared with the dynamical (Table 2)) and literature spectroscopic (Table 3)) mass estimates from Sect. 3.., is now compared with the dynamical (Table \ref{tab:dyn}) ) and literature spectroscopic (Table \ref{tab:spec}) ) mass estimates from Sect. \ref{sec:data}. + In Fig., In Fig. + 3 the results of the MSHOS and the present-day mass estimates are compared with dynamical mass estimates for massive stars from the literature as shown in. Table 2.., \ref{fig:comp_all} the results of the MSH05 and the present-day mass estimates are compared with dynamical mass estimates for massive stars from the literature as shown in Table \ref{tab:dyn}. + Inspecting the mass estimates. a very large spread is noticeable.," Inspecting the mass estimates, a very large spread is noticeable." + The results of a linear correlation analysis for all stars in the sample are shown in Table 8.., The results of a linear correlation analysis for all stars in the sample are shown in Table \ref{tab:corr}. + For each sample. the slope and offset of a best-fitting linear relation are given together with the correlation coefficients.," For each sample, the slope and offset of a best-fitting linear relation are given together with the correlation coefficients." +" The ""MSHI"" column provides the mass estimates from MSHO05 using the theoretical T. scale. whilst ""MSH2"" gives the results from their observational Ty scale. """," The “MSH1” column provides the mass estimates from MSH05 using the theoretical $T_\mathrm{eff}$ scale, whilst “MSH2” gives the results from their observational $T_\mathrm{eff}$ scale. “" +"ini"". ""evo. “start” and ""end"" are the results arrived at here. with “int” marking the results for initial masses of the models. “start” the evolved mass when a star enters a spectral type and ""end"" the one when he leaves it. ""","ini”, “evol”, “start” and “end” are the results arrived at here, with “ini” marking the results for initial masses of the models, “start” the evolved mass when a star enters a spectral type and “end” the one when he leaves it. “" +"Evol"" is the mean mass computed from {μμ and oyu.",Evol” is the mean mass computed from $m_\mathrm{start}$ and $m_\mathrm{end}$. + With an offset very close to 0. a slope of nearly |. and a correlation coefficient of ~0.9. the ΜΡΗΟΣ mass estimates agree very well with the dynamical masses of the sample.," With an offset very close to 0, a slope of nearly 1, and a correlation coefficient of $\sim$ 0.9, the MSH05 mass estimates agree very well with the dynamical masses of the sample." + Furthermore. the minimal. maximal. and evolutionary present-day masses. which are calibrated on the theoretical Το scale of MSHOS. agree very well with the observed dynamical masses.," Furthermore, the minimal, maximal, and evolutionary present-day masses, which are calibrated on the theoretical $T_\mathrm{eff}$ scale of MSH05, agree very well with the observed dynamical masses." + The offsets are quite close to 0 and, The offsets are quite close to 0 and +o dust extending into the adopted fitting range. a limiter number of data points to produce reliable fits. etc.).,"to dust extending into the adopted fitting range, a limited number of data points to produce reliable fits, etc.)." + In Fig., In Fig. + 1 we show the resulting colour gradients as a 'unction of projected galactocentric distance for our larges colour baseline. (2£). for all of our sample galaxies. anc sorted by (revised) Llubble type.2 (indicated in the upper eft-hand corner of cach panel).," \ref{totgrads.fig} we show the resulting colour gradients as a function of projected galactocentric distance for our largest colour baseline, $(B-I)$, for all of our sample galaxies, and sorted by (revised) Hubble type, (indicated in the upper left-hand corner of each panel)." + We have plotted the results obtained from either side of the plane hy open squares anc illecl circles. respectively.," We have plotted the results obtained from either side of the plane by open squares and filled circles, respectively." + The latter represent. the side of he ealactic plane least alfected by extinction. if appropriate.," The latter represent the side of the galactic plane least affected by extinction, if appropriate." + However. one should note that in some cases the cust [ane corrugates. in which case the least (most) clusty side refers o the side where the majority of the profiles are least (most) alfected by extinction.," However, one should note that in some cases the dust lane corrugates, in which case the least (most) dusty side refers to the side where the majority of the profiles are least (most) affected by extinction." + Tho fits were done in the range 1.05.<|z|X3.0h.: for exceptions see the caption of Fig. ," The fits were done in the range $1.0 h_z \le +|z| \le 3.0 h_z$; for exceptions see the caption of Fig. \ref{totgrads.fig}." +Although one can immediately see that the detected vertical colour gradients are small ancl relatively constant as a [function of position along the galaxies. major axes. in most galaxies small-scale variations in the magnitude. and even the sense of the colour gradients are observed.," Although one can immediately see that the detected vertical colour gradients are small and relatively constant as a function of position along the galaxies' major axes, in most galaxies small-scale variations in the magnitude, and even the sense of the colour gradients are observed." + The differences are in the sense that the outer regions ecnerally display τοσο colours with increasing z-distance. whereas the opposite behaviour is often observed. in and near the ealactic centres.," The differences are in the sense that the outer regions generally display redder colours with increasing -distance, whereas the opposite behaviour is often observed in and near the galactic centres." + In addition. it is also clear that although the results obtained from cither side of the galactic planes agree reasonably well. small cilfercnces between both sides. are appreciated in almost all of our sample galaxies.," In addition, it is also clear that although the results obtained from either side of the galactic planes agree reasonably well, small differences between both sides are appreciated in almost all of our sample galaxies." + These are unlikely to be caused by incorrect background subtractions. since we are comparing colours within a given ealaxy.," These are unlikely to be caused by incorrect background subtractions, since we are comparing colours within a given galaxy." + In all cases. the background emission in the field of view of our programme galaxies could be well represented by a two-dimensional plane. determined by the (lux in regions sullicientlv. far away [rom the galaxies themselves in order not to be alleeted by residual galactic light.," In all cases, the background emission in the field of view of our programme galaxies could be well represented by a two-dimensional plane, determined by the flux in regions sufficiently far away from the galaxies themselves in order not to be affected by residual galactic light." + For the majority of our observations. these planes were closely approximated »v constant [ux values across the CCD field.," For the majority of our observations, these planes were closely approximated by constant flux values across the CCD field." + The remaining uncertainties in the background. contribution are due to »oisson noise statistics (see de Grijs 1997] for a detailed overview of the image reduction techniques used)., The remaining uncertainties in the background contribution are due to poisson noise statistics (see de Grijs [1997] for a detailed overview of the image reduction techniques used). + We will. cliseuss the observed colour behaviour for hose sample galaxies for which we have high-quality colour oofiles. and those that show significant point-to-point variations. either on one side or compared to the other side of the plane. or significant non-zero e&racients.," We will discuss the observed colour behaviour for those sample galaxies for which we have high-quality colour profiles, and those that show significant point-to-point variations, either on one side or compared to the other side of the plane, or significant non-zero gradients." + For the detailed (projected) two-dimensional behaviour of colour as a function of position within our sample galaxies. we refer the reader to the (£4.A) colour maps for 24 of the galaxies in de Cirijs et al. (," For the detailed (projected) two-dimensional behaviour of colour as a function of position within our sample galaxies, we refer the reader to the $(I-K)$ colour maps for 24 of the galaxies in de Grijs et al. (" +1997) and the (£3)£) colour maps for all galaxies in de Grijs (1997. Chapter 9).,"1997) and the $(B-I)$ colour maps for all galaxies in de Grijs (1997, Chapter 9)." + In Fig., In Fig. + 2. à number of the colour profiles. extracted at various positions along the galaxies major axes. are shown for a few of the most instructive sample galaxies.," \ref{colprofs.fig} a number of the colour profiles, extracted at various positions along the galaxies' major axes, are shown for a few of the most instructive sample galaxies." + The calibrated colours apply to the colour profiles at the galactic centres: olfsets in colour. in increments of ACB/)=x10 mag. have been applied to the other profiles for reasons of clarity and display purposes.," The calibrated colours apply to the colour profiles at the galactic centres; offsets in colour, in increments of $\Delta (B-I) = \pm 1.0$ mag, have been applied to the other profiles for reasons of clarity and display purposes." +the HID. were two distinct regions separated by a gap around 150 ets 11 are clearly visible (see Fig.,"the HID, were two distinct regions separated by a gap around 150 cts 1 are clearly visible (see Fig." + | here. right hand panel. and Fig.," \ref{hid} here, right hand panel, and Fig." + | in DSO2)., 1 in DS02). + These two HB regions were labeled as an upper HB and a lower HB (UHB and LHB. respectively) and two broad-band spectra were separately extracted for each of them.," These two HB regions were labeled as an upper HB and a lower HB (UHB and LHB, respectively) and two broad-band spectra were separately extracted for each of them." + We adopted the same approach., We adopted the same approach. + On the other hand. during the 1997 observation. the source was at NB: in this latter case. the NB structure is more evident using the CD.," On the other hand, during the 1997 observation, the source was at NB; in this latter case, the NB structure is more evident using the CD." + Three spectra were extracted by DSO2 for the NB. and they were labeled as an upper NB. a medium NB and a lower LNB (UNB. MNB and LNB. respectively. see Fig.," Three spectra were extracted by DS02 for the NB, and they were labeled as an upper NB, a medium NB and a lower LNB (UNB, MNB and LNB, respectively, see Fig." + 2. left panel. in DSO2).," 2, left panel, in DS02)." + We note however that the time-filters used by DSO2 to extract the source spectra of the three parts of the NB were almost rough as they actually simply considered three consecutive time intervals which were associated to the periods spent by the source at UNB. MNB and LNB stages. respectively.," We note however that the time-filters used by DS02 to extract the source spectra of the three parts of the NB were almost rough as they actually simply considered three consecutive time intervals which were associated to the periods spent by the source at UNB, MNB and LNB stages, respectively." + We realize however that this ts. in fact. not true as the source motion along the NB ts not strictly continuous and smooth but it actually behaves like a random motion: this is clearly evident plotting e.g. the SC as a function of time (see Fig. 2)).," We realize however that this is, in fact, not true as the source motion along the NB is not strictly continuous and smooth but it actually behaves like a random motion; this is clearly evident plotting e.g. the SC as a function of time (see Fig. \ref{cvstime}) )." + In order to produce spectra which correspond to UNB. MNB and LNB positions we thus used a more refined criterion: namely. we also divided the NB in three intervals (see Fig. 1..," In order to produce spectra which correspond to UNB, MNB and LNB positions we thus used a more refined criterion; namely, we also divided the NB in three intervals (see Fig. \ref{hid}," + left panel). according to SC-value being greater than 0.25 (UNB). between 0.23 and 0.25 (MNB) and lower than 0.23 (LNB). but temporal filter for spectral extraction were produced following the source behaviour as reported in Fig. 2..," left panel), according to SC-value being greater than 0.25 (UNB), between 0.23 and 0.25 (MNB) and lower than 0.23 (LNB), but temporal filter for spectral extraction were produced following the source behaviour as reported in Fig. \ref{cvstime}." + The shift of the source position in the CD/HID between our plots and those reported in DS02 has two reasons: first. we extracted light-curves (such as energy spectra) from a4’ region centered around the MECS source image (while in DSO2 no spatial selection was used for producing CD/HID). and second we used a better MECS instrumental channel-to-energy conversion law. required when extracting light curves in different energy bands.," The shift of the source position in the CD/HID between our plots and those reported in DS02 has two reasons: first, we extracted light-curves (such as energy spectra) from a $\arcmin$ region centered around the MECS source image (while in DS02 no spatial selection was used for producing CD/HID), and second we used a better MECS instrumental channel-to-energy conversion law, required when extracting light curves in different energy bands." + The time-filtered spectra of the high-energy (PDS:?) were produced using the package and we grouped the PDS channels in order to have S/N & 3., The time-filtered spectra of the high-energy \citep[PDS;][]{frontera97} were produced using the package and we grouped the PDS channels in order to have S/N $\ga$ 3. + Bin with lower threshold. as being statistically meaningless. were discarded and not included in the fit.," Bin with lower threshold, as being statistically meaningless, were discarded and not included in the fit." + A further comment is required for the use of the LECS response files (the energy redistribution matrix and the effective area files. RMF and ARF. respectively).," A further comment is required for the use of the LECS response files (the energy redistribution matrix and the effective area files, RMF and ARF, respectively)." + When the LECS count rate is about 50-60 ets +] or lower. the standard on-line available response matrix can be safely used. while for higher count rates specific observation-related RMF and ARF files should be produced.," When the LECS count rate is about 50-60 cts 1 or lower, the standard on-line available response matrix can be safely used, while for higher count rates specific observation-related RMF and ARF files should be produced." + The fits on the LECS Crab spectrum (~ 200 ets 1) clearly shows the presence of < residuals in the region 0.5-1 keV and 2-3keV., The fits on the LECS Crab spectrum $\sim$ 200 cts 1) clearly shows the presence of $\la$ residuals in the region 0.5-1 keV and 2-3. +. It is thus evident that when à source spectrum is very bright and far from being Crab-like. these instrumental features can be significantly enhanced.," It is thus evident that when a source spectrum is very bright and far from being Crab-like, these instrumental features can be significantly enhanced." + This problem may be critical for LMXBs as their low-energy X-ray emission is dominated by the presence of the soft — 0.5 keV BB-like component in addition to, This problem may be critical for LMXBs as their low-energy X-ray emission is dominated by the presence of the soft $\sim$ 0.5 keV BB-like component in addition to +By integrating the volume deusity along the line of sight to the compact object at apastron. we find a column deusity of 1.9107906.?«€Ny€L8x1076 at the 99% CL. which is much. higher than results fouud elsewhere in the literature (Watersetal.1988:MartiEspositoetal.,"By integrating the volume density along the line of sight to the compact object at apastron, we find a column density of $1.9\times10^{26}\rm{cm}^{-2}\leq N_H \leq 1.8\times 10^{28}\rm{cm}^{-2}$ at the $99\%$ CL, which is much higher than results found elsewhere in the literature \citep{waters, marti, esposito}." + 2007).. In particular. when we use tlie column deusity found by X-ray observations Ny=(5.7£0.3)x10?!ei.7 Espositoetal. (2007)... we find a reduced {7 of 3.06 (11 degrees of [reedom). corresponding to a 47 probability P(X?>3.06)=0.01%.," In particular, when we use the column density found by X-ray observations $N_H=(5.7\pm 0.3)\times 10^{21}\rm{cm}^{-2}$ \cite{esposito}, , we find a reduced $\tilde{\chi}^2$ of $3.06$ (11 degrees of freedom), corresponding to a $\chi^2$ probability $P(\tilde{\chi}^2\geq 3.06)=0.04\%$." + A rough estimate suggests that by including 105€ of helium. the columu density would be reduced by a factor of 2. which is uot sullictent to achieve compatibility with X-ray results.," A rough estimate suggests that by including $\sim 10\%$ of helium, the column density would be reduced by a factor of $\sim 2$, which is not sufficient to achieve compatibility with X-ray results." + Density profiles in Be stars typically have radial dependences of 1/727. where 2.3 interactions are insullicient to account for the orbital modulation. then the intrinsic uou-attenuated cdillerential spectrum is essentially tle same as the observed spectrum (apower law of spectral index —2.1).," Since, in the TeV range, the interaction with matter is approximately independent of the energy, and since, as figure \ref{light_curve_gamma_gamma} shows, $\gamma\gamma$ interactions are insufficient to account for the orbital modulation, then the intrinsic non-attenuated differential spectrum is essentially the same as the observed spectrum (apower law of spectral index $-2.4$ )." + However. the intrinsic TeV luminosity is several orders of magnitude higher than the measured luminosity.," However, the intrinsic TeV luminosity is several orders of magnitude higher than the measured luminosity." + Taking the cistauce to the, Taking the distance to the +eusenible of simnulatious to explore svstematic effects iu SNela.,ensemble of simulations to explore systematic effects in SNeIa. +" By relating p. to tego, aud Ni to ΔΡ]. our results support the observational finding that SNeIa from older stellar populations are svstematically dimuucr."," By relating $\cdens$ to $\tcool$ and $\Ni{56}$ to $\dmb$, our results support the observational finding that SNeIa from older stellar populations are systematically dimmer." + While a degeneracy between age and ietallicity iu the inteerated liebt of stellar populations exists. the observed dependence of 1ie22 briglituess of SNoIn ou mean stellar age ds apparently the stronger effect (Gallagherotal.2008:Howelletal. 2009).," While a degeneracy between age and metallicity in the integrated light of stellar populations exists, the observed dependence of mean brightness of SNeIa on mean stellar age is apparently the stronger effect \citep{GallagherEtAl08,howelletal+09}." +.. Accordingly. our choice to ueglect metallicity effects and consider only the effect of coutral density on 79Ni vield allows us to offer a theoretical explanation for this observed trend.," Accordingly, our choice to neglect metallicity effects and consider only the effect of central density on $\Ni{56}$ yield allows us to offer a theoretical explanation for this observed trend." + If we udditionally consider the effect of metallicity. we may see a slightly stronger treud of decreasing briglhtuess with Increasing age as has been previously suggested (Timunestal. 2003)..," If we additionally consider the effect of metallicity, we may see a slightly stronger trend of decreasing brightness with increasing age as has been previously suggested \citep{timmes.brown.ea:variations}." + Other effects besides progenitor central cusity and metallicity. such as progenitor main sequence lass. niav also contribute to this trend.," Other effects besides progenitor central density and metallicity, such as progenitor main sequence mass, may also contribute to this trend." + The iuseuxitivitv of the overall Fe-eroup vield to central density. aud therefore delay timc. along with the dependence of the 79N3 vield on central density. implics hat SNela of similar brightuesses (aud therefore simular PONT yield) from progenitors of different ages will not jiwe the same total Fe-eroup vield.," The insensitivity of the overall Fe-group yield to central density, and therefore delay time, along with the dependence of the $\Ni{56}$ yield on central density, implies that SNeIa of similar brightnesses (and therefore similar $\Ni{56}$ yield) from progenitors of different ages will not have the same total Fe-group yield." + Those from older »opulatious will. on average. have larger masses of stable species.," Those from older populations will, on average, have larger masses of stable species." + This may argue for a slight non-uuiforuüdtv iu he Philips relation based on cuviroument (Woosleyetal.2007:IIóflüich. 2010).," This may argue for a slight non-uniformity in the Phillips relation based on environment \citep{WoosleyEtAl07:LightCurves, +hoeetal2010}." +.. The resulting closely-related famuly of brightuess-decline time relations also xovides a physical motivation for intrinsic scatter in the Phillips relation as a result of combining populations with different mean stellar ages., The resulting closely-related family of brightness-decline time relations also provides a physical motivation for intrinsic scatter in the Phillips relation as a result of combining populations with different mean stellar ages. + In this picture the primary parameter is the deeree of expansion at DDT. determined by the morphology of the carly flame (aud the DDT deusitv. which we hold constant). aud the age acts a weaker secondary paramcter.," In this picture the primary parameter is the degree of expansion at DDT, determined by the morphology of the early flame (and the DDT density, which we hold constant), and the age acts a weaker secondary parameter." + Tn any case. the possibility of such an effect motivates further exploration of the impact of ceutral density ou the liebt curve itself.," In any case, the possibility of such an effect motivates further exploration of the impact of central density on the light curve itself." + This work was supported by the Departinent of Encrev through erauts DE-FG02-07ER11516. DE-FC02-OSER11570. and DE-EGO2-08ER11565. and by NASA through eraut NNNODADLOG. ACC also acknowledges support from the Departineut of Enerey under eraut DE-FGU2-87ER10317.," This work was supported by the Department of Energy through grants DE-FG02-07ER41516, DE-FG02-08ER41570, and DE-FG02-08ER41565, and by NASA through grant NNX09AD19G. ACC also acknowledges support from the Department of Energy under grant DE-FG02-87ER40317." + DMT received support from the Bart BBok fellowship at the University of Arizona for part of this work., DMT received support from the Bart Bok fellowship at the University of Arizona for part of this work. + The authors eratefiully acknowledge the generous assistance of Pierre Lesaffre. as well as fruitful diseussons with Alike Zineale. aud the use of NSE and weak reaction tables developed by Ivo Seitenzahl.," The authors gratefully acknowledge the generous assistance of Pierre Lesaffre, as well as fruitful discussions with Mike Zingale, and the use of NSE and weak reaction tables developed by Ivo Seitenzahl." +" The authors also acknowledge the hospitality of the KITDP. which is supported bv NSF erant PIIYO5-51161. during the programs “Accretion and Explosion: the Astrophysics of Degenerate Stars” and ""Stellar Death and Superuovae."," The authors also acknowledge the hospitality of the KITP, which is supported by NSF grant PHY05-51164, during the programs “Accretion and Explosion: the Astrophysics of Degenerate Stars” and “Stellar Death and Supernovae.”" + The software used in this work was du part developed bv the DOE-supported ASC/Alliances Center for Astrophysical Thermounclear Flashes at the University of Chicago., The software used in this work was in part developed by the DOE-supported ASC/Alliances Center for Astrophysical Thermonuclear Flashes at the University of Chicago. + This research utilized resources at the New York Center for Computational Scicuces at Stony Brook. University/Brookhaven National Laboratory which is supported by the DDepartincut of Euecrgy mucder Contract DDE-ACO2-98CTITOSSG aud bv the State of New York., This research utilized resources at the New York Center for Computational Sciences at Stony Brook University/Brookhaven National Laboratory which is supported by the Department of Energy under Contract DE-AC02-98CH10886 and by the State of New York. +We used ROSAT PSPC observations for analyzing surface brightness profiles.,We used $ROSAT$ PSPC observations for analyzing surface brightness profiles. + Although the PSPC has less angular resolution than the HRI. its background is much lower. which is important for extended faint objects such as clusters.," Although the PSPC has less angular resolution than the HRI, its background is much lower, which is important for extended faint objects such as clusters." + To account for the point spread function of the PSPC. which depends on the angle from the center of the instrument. we calculated the point spread function for each cluster depending on its location in the field of view (Markevitch Vikhlinin 1997).," To account for the point spread function of the PSPC, which depends on the angle from the center of the instrument, we calculated the point spread function for each cluster depending on its location in the field of view (Markevitch Vikhlinin 1997)." + The surface brightness distribution was fit to a J model convolved with the PSF with 3. Xo. o. and a constant background as free parameters using the method of maximum-likelihood.," The surface brightness distribution was fit to a $\beta$ model convolved with the PSF with $\beta$, $\Sigma_0$, $a$, and a constant background as free parameters using the method of maximum-likelihood." + We exclude from our sample all clusters which cannot be fit to a physically reasonable ./ model (e.g.. A851 and A1758. which are known to have significant substructure — see Schindler Wambsganss 1996. Rines et al.," We exclude from our sample all clusters which cannot be fit to a physically reasonable $\beta$ model (e.g., A851 and A1758, which are known to have significant substructure – see Schindler Wambsganss 1996, Rines et al." + 1998)., 1998). + The fit for A1576 1s shown in Figure |. and the resulting parameters for all clusters are given in Table 2. along with the resulting gas masses. gravitational masses. and gas fractions as calculated assuming qo = 0.0.," The fit for A1576 is shown in Figure 1, and the resulting parameters for all clusters are given in Table 2, along with the resulting gas masses, gravitational masses, and gas fractions as calculated assuming $q_0$ = 0.0." + The parameters for CLOO164-16 agree with those found by Neumann Bóhhringer (1997) for HRI observations., The parameters for CL0016+16 agree with those found by Neumann Böhhringer (1997) for HRI observations. + Zw3146 is known to contain a large cooling flow (Allen et al., Zw3146 is known to contain a large cooling flow (Allen et al. + 1996). so the temperature and hence the total mass ts likely underestimated.," 1996), so the temperature and hence the total mass is likely underestimated." + The observed gas fraction for this cluster would then be overestimated., The observed gas fraction for this cluster would then be overestimated. + With improved angular resolution. we could exclude the cooling flow region and determine a more accurate temperature for the cluster.," With improved angular resolution, we could exclude the cooling flow region and determine a more accurate temperature for the cluster." + Table 2 shows the gas fractions both at | Mpe and at F'soo. and we compare these gas fractions in Figure/r;! 3.," Table 2 shows the gas fractions both at 1 $h_{50}^{-1}$ Mpc and at $r_{500}$, and we compare these gas fractions in Figure 3." +" These two methods use equations 5 and 7. respectively to estimate M,,,: both use equations 2. and 3. to estimate the gas mass."," These two methods use equations \ref{mass2} and \ref{massev} + respectively to estimate $M_{tot}$; both use equations \ref{density} + and \ref{cendensity} to estimate the gas mass." + The values of f.(rso9) are lower on average than f.(Visi Mpc). but the difference lies mostly within the fractional errors taken from the confidence limits on the temperature.," The values of $f_g(r_{500})$ are lower on average than $f_g(1 h_{50}^{-1}$ Mpc), but the difference lies mostly within the fractional errors taken from the confidence limits on the temperature." + Evrard (1997) finds that in simulations. the gas fraction at rsqq 18 related to the gas fraction at ry. according to {ο=fury)<(soo/rx )!. where η = 0.13-0.17.," Evrard (1997) finds that in simulations, the gas fraction at $r_{500}$ is related to the gas fraction at $r_X$, according to $f_g(r_{500}) = f_g(r_X)\times (r_{500}/r_X)^{\eta}$ , where $\eta$ = 0.13-0.17." +" The above discussion of the effect of temperature profiles on f£, suggests that κος). 18 understimated by assuming isothermality.", The above discussion of the effect of temperature profiles on $f_g$ suggests that $f_g(r_{500})$ is understimated by assuming isothermality. + This effect 1s in the sense of reducing the difference between the simulations. of Evrard and the present results., This effect is in the sense of reducing the difference between the simulations of Evrard and the present results. + To determine the apparent change of the measured gas fraction. we must know the gas fraction at the present epoch.," To determine the apparent change of the measured gas fraction, we must know the gas fraction at the present epoch." + We do this by averaging the gas fractions found by JF98 for clusters observed with at z«0.1 with luminosities greater than 10% erg s! and accurate measurements of 3 and core radius e., We do this by averaging the gas fractions found by JF98 for clusters observed with at $z < 0.1$ with luminosities greater than $^{45}$ erg $^{-1}$ and accurate measurements of $\beta$ and core radius $a$. + Comparing to ASCA data introduces only a small error: Henry (1997) found that the fluxes determined from ASCA for EMSS clusters agree very well with those found byEinstein., Comparing to $ASCA$ data introduces only a small error; Henry (1997) found that the fluxes determined from $ASCA$ for $EMSS$ clusters agree very well with those found by. + For ten distant clusters. the mean flux ratio ts 0.998. with an rms scatter of 0.205.," For ten distant clusters, the mean flux ratio is 0.998, with an rms scatter of 0.205." + There is excellent agreement between the gas fractions calculated from the method of equations 2 and 3 and those in JF98. so the relative gas fractions of nearby versus distant clusters should be accurate.," There is excellent agreement between the gas fractions calculated from the method of equations \ref{density} + and \ref{cendensity} and those in JF98, so the relative gas fractions of nearby versus distant clusters should be accurate." + The luminosity selection ensures that the nearby clusters are similar to those we study at high redshift., The luminosity selection ensures that the nearby clusters are similar to those we study at high redshift. + For example. the well-known relation between X-ray luminosity and temperature (e.g.. David et al.," For example, the well-known relation between X-ray luminosity and temperature (e.g., David et al." + 1993. Mushotzky Scharf 1997) indicates that clusters with comparable luminosities also have comparable gas temperatures.," 1993, Mushotzky Scharf 1997) indicates that clusters with comparable luminosities also have comparable gas temperatures." + It has been found, It has been found +higher for the faintest targets. which have a brightness of about Hpx22.8.,"higher for the faintest targets, which have a brightness of about $m_R \approx 22.8$." + Converting measured velocities to velocity dispersion is not straightforward., Converting measured velocities to velocity dispersion is not straightforward. + Fluctuations may arise from inhomogeneous sampling or from small numbers of outliers (stemming e.g. from erroneous velocities or contaminants belonging to other galaxies—in our case 11404. which ts at a projected distance of 54 kpe).," Fluctuations may arise from inhomogeneous sampling or from small numbers of outliers (stemming e.g. from erroneous velocities or contaminants belonging to other galaxies—in our case 1404, which is at a projected distance of 54 kpc)." +" Therefore we define a ""safe"" subsample. where the objects have distances greater than 3 aremin to 11404. velocity uncertainties smaller than 75 km/s. and magnitudes in the range 21. 9)., In the following we concentrate on secure clusters (SpARCS flux $>9$ ). + This eusures that only massive systems and uo potential false positives are stacked., This ensures that only massive systems and no potential false positives are stacked. +" Selecting clusters in this wav leads to more massive systems at lieh redslift than at low redshift. but with the limited data set used here we cannot afford more physical cuts οσο, fixed richness)."," Selecting clusters in this way leads to more massive systems at high redshift than at low redshift, but with the limited data set used here we cannot afford more physical cuts (e.g. fixed richness)." + Within the overlap area between SpARCS aud CFIITLS we find 15 moderately rich aud securely detected clusters with redshifts between 2=0.2 aud 2=1.6., Within the overlap area between SpARCS and CFHTLS we find 48 moderately rich and securely detected clusters with redshifts between $z=0.2$ and $z=1.6$. + As sources we use Lyinanu-break galaxies (LBCs)} selected four catalogues of the CFIITLS. field W1 which overlaps with the SpARCS XMSLLSS Ποια by approximately six dee?., As sources we use Lyman-break galaxies (LBGs) selected from catalogues of the CFHTLS field W1 which overlaps with the SpARCS XMM-LSS field by approximately six $^2$. + The optical multi-color catalogues are created from stacked images (seeErbeu im a similar wav as described in Wildebrandtetal.(20093). with the notable exception that the PSF is brought to the same Cuussiuir shape over the whole pointing aud for all five images in the agri filters of cach field.," The optical multi-color catalogues are created from stacked images \citep[see][for details on the imaging + data reduction with the THELI pipeline]{2009A&A...493.1197E} in a similar way as described in \cite{2009A&A...498..725H} with the notable exception that the PSF is brought to the same Gaussian shape over the whole pointing and for all five images in the $ugriz$ -filters of each field." + This allows more accurate color measurements which results iu better photo-:« aud a more accurate Lymanu-break galaxy selection., This allows more accurate color measurements which results in better $z$ 's and a more accurate Lyman-break galaxy selection. + Details of this method will be presented iu a forthcoming paper., Details of this method will be presented in a forthcoming paper. + The LBC color selection of d-. ο aud r-dropouts is identical to the one preseuted in Wildebrandtetal.(20094) with the exception that «-dropouts are onlv required to have 4οc Linstead of wg>1.5.," The LBG color selection of $u$ -, $g$ -, and $r$ -dropouts is identical to the one presented in \cite{2009A&A...498..725H} with the exception that $u$ -dropouts are only required to have $u-g>1$ instead of $u-g>1.5$." + This relaxed αο cut is identical to the one used by vanderBureetal.(2010) whose huninositv function Düeasureinents we use for Iu the following we use 23 a-dropouts aud 2~1L g- The s-dropouts are the cleanest background saluple in terms of low-: contamination and they are he lowest-redshift dropout sample detectable from the eround which results iu the highest apparent brieltucss., This relaxed $u-g$ cut is identical to the one used by \cite{2010A&A...523A..74V} whose luminosity function measurements we use for In the following we use $z\sim3$ $u$ -dropouts and $z\sim4$ $g$ The $u$ -dropouts are the cleanest background sample in terms of $z$ contamination and they are the lowest-redshift dropout sample detectable from the ground which results in the highest apparent brightness. + The g-dropouts are fainter by ~O5imag due to their arecr distance resulting iu larecr photometric errors.," The $g$ -dropouts are fainter by $\sim0.5{\rm + mag}$ due to their larger distance resulting in larger photometric errors." + They also show larger low-: coutamination., They also show larger $z$ contamination. + Thus. we can onlv safely cross-correlate them to clusters with intermediate and high redshifts - bevoud the redshifts of he possible coutaminants.," Thus, we can only safely cross-correlate them to clusters with intermediate and high redshifts - beyond the redshifts of the possible contaminants." + In the following. we correct for he dilution due to contamination (boosting the signal by 1054) whenever g-dropouts are used.," In the following, we correct for the dilution due to contamination (boosting the signal by $10\%$ ) whenever $g$ -dropouts are used." + The CEITTLS-Wide data are too shallow to select r-dropouts iu siguificaut manuboers., The CFHTLS-Wide data are too shallow to select $r$ -dropouts in significant numbers. + The signal-to-noise ratio (S/N) of the leusine signal per background galaxy is generally lower for uagnification-based iiethods than for shearbased ILlowever. since magnitudes - the only requirement for naguification - are easier to measure than shapes. there are always more galaxies aud im particular hiebher-redslüft ealaxies available for magnification.," The signal-to-noise ratio (S/N) of the lensing signal per background galaxy is generally lower for magnification-based methods than for shear-based However, since magnitudes - the only requirement for magnification - are easier to measure than shapes, there are always more galaxies and in particular higher-redshift galaxies available for magnification." + Thus. there is a weak-even redshift for cach data set bevoud which nagnification becomes more powerful than shear because ellipticity measurements become impossible.," Thus, there is a break-even redshift for each data set beyond which magnification becomes more powerful than shear because ellipticity measurements become impossible." + This is the case with clusters at τς0.8 where very few galaxies with reliable shape measurements are available at redshift ueher than the cluster redshift. as it is very difficult to neasure shapes of ealaxies with +>1 from erouud-based data.," This is the case with clusters at $z\ga0.8$ where very few galaxies with reliable shape measurements are available at redshift higher than the cluster redshift, as it is very difficult to measure shapes of galaxies with $z>1$ from ground-based data." + Since the S/N is too low for a single high-: cluster to be detected via magnification with the backeround samples described above we rely ou stacking the signals of several clusters., Since the S/N is too low for a single $z$ cluster to be detected via magnification with the background samples described above we rely on stacking the signals of several clusters. + In this wav we can in principle estimate their average nass. a method very similar to what is usually done in galaxv-galaxy-leusiug.," In this way we can in principle estimate their average mass, a method very similar to what is usually done in galaxy-galaxy-lensing." + Therefore. the estimator is identical to Eq.," Therefore, the estimator is identical to Eq." + 11 in Hildebrandtetal.(2009h).. ic. we look for correlations in the positions of LBGs aud the SpARCS clusters.," 11 in \cite{2009A&A...507..683H}, i.e. we look for correlations in the positions of LBGs and the SpARCS clusters." + Those positions are correlated due to the magnification-bias effect of WL., Those positions are correlated due to the magnification-bias effect of WL. + The iiagnuificatiou signal scales linearly with the slope of the nuuber counts. αι Thus. it is important to select a source population with a large Jal.," The magnification signal scales linearly with the slope of the number counts, $\alpha$ Thus, it is important to select a source population with a large $\left|\alpha\right|$." + For the interpretation of the signal. a needs to be measured.," For the interpretation of the signal, $\alpha$ needs to be measured." + ILlowever. we cannot use the measured o directly because of incompleteness iu our catalogue.," However, we cannot use the measured $\alpha$ directly because of incompleteness in our catalogue." + Furthermore. the incompleteness changes over the magnitude biu audso does the a itself.," Furthermore, the incompleteness changes over the magnitude bin andso does the $\alpha$ itself." + In order to account for these effects we model the incompleteness as à function of iiagnuitude by comparing the LBC uuuber counts in the to the ones in the CEITTLS-Deep (ITildebrandt 2009a).., In order to account for these effects we model the incompleteness as a function of magnitude by comparing the LBG number counts in the CFHTLS-Wide to the ones in the CFHTLS-Deep \citep{2009A&A...498..725H}. +We then multiply the Iuniuosity function of the LBCs measured in vauderBureetal.(2010) with, .We then multiply the luminosity function of the LBGs measured in \cite{2010A&A...523A..74V} with +The origin of the energy required to power the solar wind. aud heat the Suus coroii fo its iulti-iullion degree teniperatures. reluains an elusive luvsteryv iu pliysies.,"The origin of the energy required to power the solar wind, and heat the Sun's corona to its multi-million degree temperatures, remains an elusive mystery in physics." + One promising mechanis1i involves the creation of waves neu the solar surface. which cau peuetrate upwards hrough the Suu's atiuosphiere with nüniual reflection or enuergv loss(Erdélyi&Fed2007 ).," One promising mechanism involves the creation of waves near the solar surface, which can penetrate upwards through the Sun's atmosphere with minimal reflection or energy loss\citep[][]{Erd07}." +. To date. there has been great controversy smroundius the iuterpretation of periodic transverse 1uotions αν Alfvénn waves.," To date, there has been great controversy surrounding the interpretation of periodic transverse motions as Alfvénn waves." +" Aagueto-bydrodvuamic (ATID) waves with a wave nunber. a=l. can he considered. “kink” oscillatious when structured by a cvlndrical svaveeuikο,"," Magneto-hydrodynamic (MHD) waves with a wave number, $m = 1$, can be considered “kink” oscillations when structured by a cylindrical waveguide." + The associated periodic transverse 1uotions nuply a deeree of plasma ποοτιΗτν. resulting in intrinsic ΑΠΟ waves exhibiting mined (i.c. slow. fast. Alfvónuu) properties depending on the local plasina parzuneters.," The associated periodic transverse motions imply a degree of plasma non-uniformity, resulting in intrinsic MHD waves exhibiting mixed (i.e. slow, fast, Alfvénn) properties depending on the local plasma parameters." +" As a result. the term ""Alfvénic was introduced to describe waves which are predominantly characterised by the signatures cisplaved bv pure Alfvénn waves."," As a result, the term ” was introduced to describe waves which are predominantly characterised by the signatures displayed by pure Alfvénn waves." +" Accordingly. transverse kiuk oscillations may often be considered iu nature providing they are incompressible, exhibit no iuteusity fluctuations along the structure. and display periodic displacements with the magnetic tension as the restoring force (Coosseusetal.2009)."," Accordingly, transverse kink oscillations may often be considered in nature providing they are incompressible, exhibit no intensity fluctuations along the structure, and display periodic displacements with the magnetic tension as the restoring force \citep[][]{Goo09}." +. Thus. bv definition. waves requie stroue maeuctic field concentrations. and/or a steep density eradieut with the external plasma. to act as waveguides (VanDoorssclacreetal.2008).. allowing οσον to be directly channelled through the solar atmosphere.," Thus, by definition, waves require strong magnetic field concentrations, and/or a steep density gradient with the external plasma, to act as waveguides \citep[][]{Van08}, allowing energy to be directly channelled through the solar atmosphere." + In the Sun's1 atmosphere. magnetic field lines chump together into1 tight bundles. forming flux tubes.," In the Sun's atmosphere, magnetic field lines clump together into tight bundles, forming flux tubes." + The ubiquitous nature of waves iu magnetic flux tubes las been demonstrated in a rauge of chromospleric aud coronal plasiuas (DePontieuetal.al. 2009).," The ubiquitous nature of waves in magnetic flux tubes has been demonstrated in a range of chromospheric and coronal plasmas \citep[][]{DeP07b, Tom07, Ban09, Jes09}." +. Spicules are dvneaic. straw-like maeguetie structures found in the solar chromosphere. aud can be divided into two distinct classes (Zaqarashvil& Exdébi 2009).," Spicules are dynamic, straw-like magnetic structures found in the solar chromosphere, and can be divided into two distinct classes \citep[][]{Zaq09}." +. We will focus on Type spicules. which are ubiquitous throughout the solar atimosphere. aud longer lived (10 iin) than their jet-like Type counterparts (DePonticuetal. 2007a)..," We will focus on Type spicules, which are ubiquitous throughout the solar atmosphere, and longer lived $\approx$ 10 min) than their jet-like Type counterparts \citep[][]{DeP07a}. ." + The narrow, The narrow +The observations presented here are the largest. near-Ii spectroscoplC Sμον of central cluster galaxies in cooling lows by a actor of four.,The observations presented here are the largest near-IR spectroscopic study of central cluster galaxies in cooling flows by a factor of four. + Seven of the objects presented have »reviousiv pubished. HX spectra., Seven of the objects presented have previously published IR spectra. + Hyvedra-X.. 1NJ0338|09. UNJOT47-19. NGC1275. A1795. A2020 and 2597 are included. in Faleke et al. (," Hydra-A, RXJ0338+09, RXJ0747-19, NGC1275, A1795, A2029 and A2597 are included in Falcke et al. (" +1998). Jaffe Bremer (1997). Ixrabbe et:d. (,"1998), Jaffe Bremer (1997), Krabbe et al. (" +2000) and Jalle. Bremer van der Wert (2001).,"2000) and Jaffe, Bremer van der Werf (2001)." +" “Phe| tine Dlluxes. are all within <30% of those oesented elsewvere when the dillerent slit widths are taken into account (e.g. Jalfe et 22001 use a slit width of 2"" to 1.2"" used in the paper).", The line fluxes are all within $<$ of those presented elsewhere when the different slit widths are taken into account (e.g. Jaffe et 2001 use a slit width of $''$ to $''$ used in the paper). + In addition. this study includes. for he first time. coverage of the Fell) line.," In addition, this study includes, for the first time, coverage of the [FeII] line." + 1n this section we discuss the limits this dataset Cah ace on reddening. trends in the Fell] line. detection. of veh ionization lines and column density constraints.," In this section we discuss the limits this dataset can place on reddening, trends in the [FeII] line, detection of high ionization lines and column density constraints." + One of the most important results from the optical spectroscopy of central. cluster galaxies is the significant reddening of the lines (Allen et 11995: Crawford. et 11999)., One of the most important results from the optical spectroscopy of central cluster galaxies is the significant reddening of the lines (Allen et 1995; Crawford et 1999). + “Phe values of E(B-V) derived imply column densities of lott em?., The values of E(B-V) derived imply column densities of $\times 10^{21}$ $^{-2}$. + There :ue several combinations of optical anc NI lines that can be used to gauge the reddening., There are several combinations of optical and NIR lines that can be used to gauge the reddening. + The mos straightforward. is Ila Paa. Figure 3 shows the ratio of lla Pao with recshift or our joint detections with Crawforc et al. (, The most straightforward is $\alpha$ $\alpha$ Figure 3 shows the ratio of $\alpha$ $\alpha$ with redshift for our joint detections with Crawford et al. ( +1909).,1999). + Phe majority of the points agree with the ratio expected from Case D. recombination. (Osterbrock 1959) but several outliers are obvious., The majority of the points agree with the ratio expected from Case B recombination (Osterbrock 1989) but several outliers are obvious. + The majority of the objects have ratios consistent with E(B-Vj<0. are consistent with the values derived. from H2 /LIa ratios by Crawford οἱ al. (, The majority of the objects have ratios consistent with $<0.1$ and are consistent with the values derived from $\beta$ $\alpha$ ratios by Crawford et al. ( +1999).,1999). + The objects with the coσι west ratios (ALOGS. ALS35 anc RAJOS21|07) have hieh E(CD-V) values in Crawford. et al. (," The objects with the three lowest ratios (A1068, A1835 and RXJ0821+07) have high E(B-V) values in Crawford et al. (" +1999) ΟυΠΟ στ 0.40οςUU and 572OLN respectively) which are consistent with those implied from this analysis (0.54-60.02. 0.540.038 anc O.71-E0.04).,"1999) $^{+0.07}_{-0.08}$ , $^{+0.06}_{-0.06}$ and $^{+0.33}_{-0.58}$ respectively) which are consistent with those implied from this analysis $\pm$ 0.02, $\pm$ 0.03 and $\pm$ 0.04)." + Llowever. there are several objects with similar values of E(D-V) in Crawford et al. (," However, there are several objects with similar values of E(B-V) in Crawford et al. (" +1999) where our estimate is substantially lower (e.g. ALGG4. (02 2250.05).,"1999) where our estimate is substantially lower (e.g. A1664, $^{+0.06}_{-0.07}$ vs. $\pm$ 0.05)." + Phis variation may be due in xwt to the extent of the emission lines and the fact that the optical and CCGS4 spectra. whie using the same slit. width (1.27). were nol »erformed a( the same position angle.," This variation may be due in part to the extent of the emission lines and the fact that the optical and CGS4 spectra, while using the same slit width $''$ ), were not performed at the same position angle." + Therefore there are several poins (e.g. ZwS216) where the optical line is brighter »v a factor ¢X two and others (e.g. RAJOS21|07) where the ine is faüinter by a factor of three., Therefore there are several points (e.g. Zw8276) where the optical line is brighter by a factor of two and others (e.g. RXJ0821+07) where the line is fainter by a factor of three. + The onv other reddening diagnostic line combination available wiin our CGS4 data are the Fell] lines at. 1.258 and 1.644542. We have data for two objects. Zw3146 and A2204. and 1e equivalent I(D-V) values. given the expected ratio of 1.258/n to 1.64457 of 1.3. are. «0.1 and 0.540.3 (as opposed ο 0.2 and 0.0 in Crawford et al.," The only other reddening diagnostic line combination available within our CGS4 data are the [FeII] lines at 1.258 and $\mu$ m. We have data for two objects, Zw3146 and A2204, and the equivalent E(B-V) values, given the expected ratio of $\mu$ m to $\mu$ m of 1.3, are $<$ 0.1 and $\pm$ 0.3 (as opposed to 0.2 and 0.0 in Crawford et al." + 1999).CGlven 1c restrictions on the current datasets. dUis not possible to make any quantitative statements about, 1999).Given the restrictions on the current datasets itis not possible to make any quantitative statements about +rpSPH retains all of these advantages while at the same time reducing the discretisation error., retains all of these advantages while at the same time reducing the discretisation error. + In the standard SPH approach there are in fact infinitely many possible choices for the discretisation of the pressure equation (equation3.5in?).., In the standard SPH approach there are in fact infinitely many possible choices for the discretisation of the pressure equation \citep[equation 3.5 in][]{1992ARA&A..30..543M}. + This is also true forr, This is also true for. +pSPH. E.g. P!VP?/p—PV1/p=opVP which suggest the discretisation d omoi ] for any σ different from zero.," E.g. $P^{\sigma-1} \nabla P^\sigma/\rho - P \nabla 1/\rho = +\sigma \rho^{-1} \nabla P$ which suggest the discretisation = - m_j ] for any $\sigma$ different from zero." + We have verified that many choices of c work for a variety of test problems., We have verified that many choices of $\sigma$ work for a variety of test problems. +" In the following, however, we restrict our attention to the case of o=1."," In the following, however, we restrict our attention to the case of $\sigma=1$." + Whether these theoretically advantageous properties ofrpSPH hold up in practice is assessed in a range of test problems in the following section., Whether these theoretically advantageous properties of hold up in practice is assessed in a range of test problems in the following section. + We will employ a Courant factor of 0.3 (i.e. 0.15 in Gadget where the kernel has a maximal radius of A)., We will employ a Courant factor of $0.3$ (i.e. $0.15$ in Gadget where the kernel has a maximal radius of $h$ ). +" We start with 50? particles on a periodic regular lattice with j—1.4, a sound speed and uniform density of unity and zero initial velocities."," We start with $50^2$ particles on a periodic regular lattice with $\gamma=1.4$, a sound speed and uniform density of unity and zero initial velocities." + The particles should stay at rest., The particles should stay at rest. +" However, as we can see in Figure 1 the total kinetic energy in the volume grows rapidly."," However, as we can see in Figure \ref{fig:ke-noise} the total kinetic energy in the volume grows rapidly." +" The total energy in the system is, however, conserved to better than 5x107? of the initial value for these tests at a value (»y(y—1))!e1.78."," The total energy in the system is, however, conserved to better than $5\times 10^{-5}$ of the initial value for these tests at a value $(\gamma (\gamma-1))^{-1}\approx 1.78$." + So the kinetic energy growth in the particle distribution only corresponds to about less than one in one thousand of the total., So the kinetic energy growth in the particle distribution only corresponds to about less than one in one thousand of the total. + I.e. the kinetic energy the particles obtain is taken from slightly decreasing internal energy allowing the total to be aconserved to high precision., I.e. the kinetic energy the particles obtain is taken from a slightly decreasing internal energy allowing the total to be conserved to high precision. + The lower the neighbour number the faster that growth., The lower the neighbour number the faster that growth. + The maximum noise reached is controlled by the artificial viscosity., The maximum noise reached is controlled by the artificial viscosity. + The noise also decreases only very slowly over time after reaching the maximum., The noise also decreases only very slowly over time after reaching the maximum. + This is one of the main reasons why particle settling is so important in SPH simulations., This is one of the main reasons why particle settling is so important in SPH simulations. + The slow decline also shows why in general settling procedures can be computationally quite intensive., The slow decline also shows why in general settling procedures can be computationally quite intensive. + In the same figure we also plot results usingrpSPH which dramatically reduces this spurious kinetic energy creation keeping it at zero to machine precision., In the same figure we also plot results using which dramatically reduces this spurious kinetic energy creation keeping it at zero to machine precision. + ? caution that it makes no sense in standard SPH to increase particle numbers while keeping the number of neighbors fixed., \cite{2000PThPS.138..609R} caution that it makes no sense in standard SPH to increase particle numbers while keeping the number of neighbors fixed. + Once a neighbour number is reached that keeps noise in the force calculation to a minimum we findrpSPH to be stable while only increasing the particle number., Once a neighbour number is reached that keeps noise in the force calculation to a minimum we find to be stable while only increasing the particle number. + Note that we also have ran these tests dramatically reducing the Courant factor without any improvement in the case of standard SPH., Note that we also have ran these tests dramatically reducing the Courant factor without any improvement in the case of standard SPH. + The thick solid line in Figure 1 uses 20 neighbors which seems optimal for this 2D calculation with the cubic spline kernel., The thick solid line in Figure \ref{fig:ke-noise} uses 20 neighbors which seems optimal for this 2D calculation with the cubic spline kernel. + Here one has enough neighbors to estimate the gradients more accurately while still having too few neighbours to show its pairing instability., Here one has enough neighbors to estimate the gradients more accurately while still having too few neighbours to show its pairing instability. +" So one may be tempted to dismiss the finding that one has the large velocity noise as long as one uses the ""correct? number of neighbours in one simulation.", So one may be tempted to dismiss the finding that one has the large velocity noise as long as one uses the “correct” number of neighbours in one simulation. +" Unfortunately, this best choice, however, is only applicable at the uniform density."," Unfortunately, this best choice, however, is only applicable at the uniform density." + To show this we perturb the x positions by a small amount so that the initially uniform xo positions are changed by adding sin(27σο)/25 to them which gives central densities that are about above the mean.," To show this we perturb the $x$ positions by a small amount so that the initially uniform $x_0$ positions are changed by adding $\sin(2\pi\,x_0)/25$ to them which gives central densities that are about above the mean." + We keep again the pressure to be exactly constant by setting the entropy of the gas only once the density has been estimated from kernel smoothing., We keep again the pressure to be exactly constant by setting the entropy of the gas only once the density has been estimated from kernel smoothing. + The thick long dashed line in Figure 1 gives the associated velocity noise., The thick long dashed line in Figure \ref{fig:ke-noise} gives the associated velocity noise. + It again is of order one percent of the sound speed and grew very rapidly., It again is of order one percent of the sound speed and grew very rapidly. + One may also be tempted to to dismiss this particle noise as irrelevant as it only contains less than a tenth of a percent of the total energy of the system., One may also be tempted to to dismiss this particle noise as irrelevant as it only contains less than a tenth of a percent of the total energy of the system. +" However, we will see in the following it is what leads to unphysical shear viscosity once one considers shear flows further below."," However, we will see in the following it is what leads to unphysical shear viscosity once one considers shear flows further below." + 'The velocity noise we just discussed is unfortunately not isotropic nor is it random., The velocity noise we just discussed is unfortunately not isotropic nor is it random. + It has a dominant component for velocities towards directions of other particles and is an effect that aids the pairing instability., It has a dominant component for velocities towards directions of other particles and is an effect that aids the pairing instability. + Our reasoning here is somewhat contrary to the explanations in the literature as, Our reasoning here is somewhat contrary to the explanations in the literature as +"In the present paper, we demonstrate that the main fraction of the ionized gas in NGC 7743 rotates on the orbits, considerably inclined to the main stellar disk of the galaxy, what may be a result of external gas accretion or tidal destruction of a gas-rich small companion.","In the present paper, we demonstrate that the main fraction of the ionized gas in NGC 7743 rotates on the orbits, considerably inclined to the main stellar disk of the galaxy, what may be a result of external gas accretion or tidal destruction of a gas-rich small companion." +" The paper is organized as follows: Section 2 gives an overview of the literature on the previous studies of the galaxy; Section 3 describes the spectroscopic observations and data reduction process; in Section 4 the properties of the stellar population (kinematics, age, and metallicity) are considered; Section 5 contains the study of gas kinematics and ionization state; Section 6 includes an overall discussion of the structure and kinematics of NGC 7743."," The paper is organized as follows: Section \ref{sec_n7743} gives an overview of the literature on the previous studies of the galaxy; Section \ref{sec_obs} + describes the spectroscopic observations and data reduction process; in Section \ref{sec_stars} the properties of the stellar population (kinematics, age, and metallicity) are considered; Section \ref{sec_gas} contains the study of gas kinematics and ionization state; Section \ref{sec_dis} includes an overall discussion of the structure and kinematics of NGC 7743." + NGC 7743 is a barred early-type galaxy (NED morphological type (R)SBO*(s)) with a total blue luminosity of Mg=—19.4 (according to the HyperLeda , NGC 7743 is a barred early-type galaxy (NED morphological type $^+$ (s)) with a total blue luminosity of $M_B=-19.4$ (according to the HyperLeda ). +"Following Jensenetal.(2003),, we adopt the distance to database!)).the galaxy to be 19.2 Mpc, that corresponds to a linear scale of 93pcαγορας”."," Following \citet{Jensen2003}, we adopt the distance to the galaxy to be 19.2 Mpc, that corresponds to a linear scale of $93\,\mbox{pc}\,\mbox{arcsec}^{-1}$." + A smooth two-armed spiral structure without any traces of star formation dominates in that optical images of the galaxy., A smooth two-armed spiral structure without any traces of star formation dominates in that optical images of the galaxy. + According to Hoetal.(1997) the galaxy has an active nucleus of the Sy2 type., According to \citet*{Ho1997} the galaxy has an active nucleus of the Sy2 type. +" Radio observations reveal a compact (under the beam=1""— 5"")) non-thermal source in the nucleus (Nagaretal.1993;Ho&Ulvestad 2001)."," Radio observations reveal a compact (under the $=1\arcsec -5$ ) non-thermal source in the nucleus \citep{Nagar1993,HoUlvestad2001}." +" However, the nuclear activity is not very high."," However, the nuclear activity is not very high." + Alonso-Herreroetal.(2000) using their spectral observations classified NGC 7743 as a ‘low-luminosity AGN’ and noted that the optical emission line ratios correspond to the boundary case between Sy2 and LINER activity types., \citet{Alonso-Herrero2000} using their spectral observations classified NGC 7743 as a `low-luminosity AGN' and noted that the optical emission line ratios correspond to the boundary case between Sy2 and LINER activity types. +" The X-ray observations by Terashimaetal.(2002) also suggest a relatively low activity of the NGC 7743 nucleus, compared with other Seyfert galaxies from their sample."," The X-ray observations by \citet{Terashima2002} also suggest a relatively low activity of the NGC 7743 nucleus, compared with other Seyfert galaxies from their sample." + The HST imagingreveals a complex structure of the circumnuclear region (R<300-500 pc) where emission knots and several curved dust lanes are observed (Regan&Mulchaey1999)., The HST imagingreveals a complex structure of the circumnuclear region $R<300$ $500$ pc) where emission knots and several curved dust lanes are observed \citep{ReganMulchaey1999}. +. Martinietal.(2003) described this structure as a loosely wound circumnuclear spiral., \citet{Martini2003} described this structure as a loosely wound circumnuclear spiral. + Moiseevetal.(2004) used integral-field spectroscopy to study the inner region morphology and kinematics., \citet*{Moiseev2004} used integral-field spectroscopy to study the inner region morphology and kinematics. +" They suggested that a turn of the innermost isophotes relates with the circumnuclear dust spiral, rather than with a triaxial bulge as it was claimed earlier."," They suggested that a turn of the innermost isophotes relates with the circumnuclear dust spiral, rather than with a triaxial bulge as it was claimed earlier." +" Despite the relatively low intensity of emission lines, they succeeded to obtain some conclusions about the ionized gas kinematics."," Despite the relatively low intensity of emission lines, they succeeded to obtain some conclusions about the ionized gas kinematics." +" In the region of R<2”—4” (190-370 pc) the gas motions agree in general with the stellar disk rotation, but non-circular ionized gas motions were also detected locally to the south from the nucleus."," In the region of $R<2\arcsec -4\arcsec$ (190–370 pc) the gas motions agree in general with the stellar disk rotation, but non-circular ionized gas motions were also detected locally to the south from the nucleus." + Moiseevetal.(2004) have as well supposed that the inner part of the galaxy disk can be tilted with respect to the outer disk., \citet{Moiseev2004} have as well supposed that the inner part of the galaxy disk can be tilted with respect to the outer disk. +" Radio observations by Duprie&Schneider(1996) demonstrated a very low content of neutral gas (near their detection limit) for the NGC 7743 disk through the ~3’ beam, whereas two separate HI clouds were discovered in the tight neighborhood of the galaxy having the masses of My;=64x107Mo and 4.5x105Mo."," Radio observations by \citet{DuprieSchneider1996} demonstrated a very low content of neutral gas (near their detection limit) for the NGC 7743 disk through the $\sim 3\arcmin$ beam, whereas two separate HI clouds were discovered in the tight neighborhood of the galaxy having the masses of $M_{HI}=6.4\times10^7\,M_\odot$ and $4.5\times10^8\,M_\odot$." +" Their systemic velocities are 1610kms-! and 1509kms~, respectively, that is close to the NGC 7743 systemic velocity which is 1710kms! (HyperLeda)."," Their systemic velocities are $1610\km$ and $1509\km$, respectively, that is close to the NGC 7743 systemic velocity which is $1710\km$ (HyperLeda)." +" According to the NED database these clouds are associated with the galaxies KUG 2341+097 and LSBC F750-04, with projected distances of 8/33 (46 and 11/00 kpc) from NGC 7743."," According to the NED database these clouds are associated with the galaxies KUG 2341+097 and LSBC F750-04, with projected distances of 3 (46 kpc) and 0 (61 kpc) from NGC 7743." + The absolute kpc)magnitude of the (61brighter galaxy (KUG 2341--097 ) is about Mp= —16., The absolute magnitude of the brighter galaxy (KUG 2341+097 ) is about $M_B=-16$ . + Therefore NGC 7743 is surrounded by the satellites which altogether contain the neutral hydrogen amount by one order larger compared to the disk of the main galaxy., Therefore NGC 7743 is surrounded by the satellites which altogether contain the neutral hydrogen amount by one order larger compared to the disk of the main galaxy. +" At the same time Maiolinoetal.(1997) from their radio observations with a beam=55"" detected the CO emission, corresponding to the total mass of the molecular gas in the galaxy disk of about 1.3x105Mo, which is comparable with the HI mass in the local environment of NGC 7743."," At the same time \citet{Maiolino1997} from their radio observations with a $=55\arcsec$ detected the CO emission, corresponding to the total mass of the molecular gas in the galaxy disk of about $1.3 \times10^8\,M_\odot$, which is comparable with the HI mass in the local environment of NGC 7743." +" In our study of NGC 7743, we have used two types of spectral data."," In our study of NGC 7743, we have used two types of spectral data." +" Long-slit spectroscopy with exposures long enough allows to reach very outer parts of the galaxy disk; however, one-dimensional character of the long-slit data restricts consideration of any azimuthal variations."," Long-slit spectroscopy with exposures long enough allows to reach very outer parts of the galaxy disk; however, one-dimensional character of the long-slit data restricts consideration of any azimuthal variations." + Integral-field spectroscopy (so called ‘3D spectroscopy’) provides two-dimensional mapping of the kinematical and stellar population characteristics but only for the central part of the galaxy due to the limited field-of-view., Integral-field spectroscopy (so called `3D spectroscopy') provides two-dimensional mapping of the kinematical and stellar population characteristics but only for the central part of the galaxy due to the limited field-of-view. + 'Thespectral observations were made at the prime focus of the SAO RAS 6-m telescope with the multimode focal reducer SCORPIO (Afanasiev&Moiseev 2005).., Thespectral observations were made at the prime focus of the SAO RAS 6-m telescope with the multimode focal reducer SCORPIO \citep{AfanasievMoiseev2005}. . + The slit had 6/11 in length and, The slit had 1 in length and +drier locations may provide a larger number of uscable observing nights per vear.,drier locations may provide a larger number of useable observing nights per year. + Based on its current NERD. it seems that PIUALPER could be used to study bright. high-mass. star-forming regions at higher angular resolution than has been previously possible at this wavelength.," Based on its current NEFD, it seems that THUMPER could be used to study bright, high-mass, star-forming regions at higher angular resolution than has been previously possible at this wavelength." + The TIIUMPER. Veam would like to acknowledge the assistance of the stall of the JCMIT throughout the planning and commissioning of ΓΕιο., The THUMPER Team would like to acknowledge the assistance of the staff of the JCMT throughout the planning and commissioning of THUMPER. + Phe ΙΟΝΤ is operated bv the Joint Astronomy Centre. Hawaii. on behalf of the Ulx Particle Physics and Astronomy Research Council (PPARC). the Netherlands Organization for Scientific Research (NWO). ancl the Canadian. National Research Council (NRC).," The JCMT is operated by the Joint Astronomy Centre, Hawaii, on behalf of the UK Particle Physics and Astronomy Research Council (PPARC), the Netherlands Organization for Scientific Research (NWO), and the Canadian National Research Council (NRC)." + PPARC are eratefully acknowledged: for erant Funding to build TIIUMPER., PPARC are gratefully acknowledged for grant funding to build THUMPER. +In this paper we study the impact of the relative velocities on the distribution of the star-forming halos at high redshift and on the redshift of formation of the very first star.,In this paper we study the impact of the relative velocities on the distribution of the star-forming halos at high redshift and on the redshift of formation of the very first star. + In. particular. we include an aspect of the relative velocity elfect that has not been previously. accounted. for. and which is critical for understanding the overall impact of the velocities on the distribution of star formation.," In particular, we include an aspect of the relative velocity effect that has not been previously accounted for, and which is critical for understanding the overall impact of the velocities on the distribution of star formation." + Recent small-scale numerical simulations (Stacyetal.2010:Greiletal.2011). found that the relative velocity. substantially increases the minimum halo mass in which stars can forni from σας that cools via molecular hydrogen. cooling (The elect of the velocities has also been simulated by al.(2011) and Naozetal. (2011))).," Recent small-scale numerical simulations \citep{Stacy:2011,Greif:2011} found that the relative velocity substantially increases the minimum halo mass in which stars can form from gas that cools via molecular hydrogen cooling (The effect of the velocities has also been simulated by \citet{Maio:2011} and \citet{Naoz:2011}) )." + This paper is organized as follows., This paper is organized as follows. + In. Section 2 we briclly review the results o£ Tseliakhovich&Lirata(2010) and Tseliakhovich.Barkana&Hirata(2010).., In Section 2 we briefly review the results of \citet{Tseliakhovich:2010} and \citet{Tseliakhovich:2010b}. + In Section 3 we summarize the results of recent simulations that include the nonlinear elfect of the relative velocity on the formation ofthe first stars via molecular cooling., In Section 3 we summarize the results of recent simulations that include the nonlinear effect of the relative velocity on the formation of the first stars via molecular cooling. + We use the simulation results to find the behavior of the minimal cooling mass versus redshift and magnitude of the relative velocity., We use the simulation results to find the behavior of the minimal cooling mass versus redshift and magnitude of the relative velocity. + In Section 4 we study in detail the probability distribution of he eas fraction in halos at high redshift. separating out and comparing the importance of the various elfects of the oulk velocity.," In Section 4 we study in detail the probability distribution of the gas fraction in halos at high redshift, separating out and comparing the importance of the various effects of the bulk velocity." + In Section 5r we then estimate the redshift of he very first star accounting for the relative velocity ellect., In Section 5 we then estimate the redshift of the very first star accounting for the relative velocity effect. +" Finally. in Section 6 we summarize our results and also give a complete list of dillerences compared. to three previous xipers: Fseliakhovich&1irata(2010).. Dalal.Pen&Seljak (2010).. and ""Fselialkkhovicl1i.Darkana&Hirata(2010).."," Finally, in Section 6 we summarize our results and also give a complete list of differences compared to three previous papers: \citet{Tseliakhovich:2010}, \citet{Dalal:2010}, and \citet{Tseliakhovich:2010b}." +" Our calculations. are carried out. in a flat CDAM universe with cosmologica parameters taken from the 7-vear WALAP results ff) maximum likelihood fit from Ixomatsuctal.( 2010))): the dark matter density today £4=0.2265. the xwvon density £254,=0.0455. the vacuum energy density Oy=0.728. the Hubble constant ily=10.4 kms * Alpe1 and the spectral index n;= 0.967."," Our calculations are carried out in a flat $\Lambda$ CDM universe with cosmological parameters taken from the 7-year WMAP results $H_0$ maximum likelihood fit from \citet{WMAP7}) ): the dark matter density today $\Omega_{c,0} = 0.2265$, the baryon density $\Omega_{b,0} = 0.0455$, the vacuum energy density $\Omega_\Lambda = +0.728$, the Hubble constant $H_0 =70.4$ km $^{-1}$ $^{-1}$, and the spectral index $n_s = 0.967$ ." + We normalize the DOWEL Spectrum to give a present value of e;=OSL (Ilxomatsuetal.2010).., We normalize the power spectrum to give a present value of $\sigma_8 = 0.81$ \citep{WMAP7}. +. We use the CAXMD-sources linear. perturbation code (Lewis&Challi-nor2007) to generate initial conditions at. recombination (specifically. at +=1020 and z=970 in order to obtain the needed derivatives)," We use the CAMB-sources linear perturbation code \citep{CAMBsources} to generate initial conditions at recombination (specifically, at $z = 1020$ and $z = 970$ in order to obtain the needed derivatives)." + In this section we brielle review the non-linear effect of the relative velocities between the barvons ancl dark matter. as ciscussed in Tseliakhovich&Llirata(2010). and Tseliakhovich.Barkana&LHirata(2010).. the latter of which we Closely follow in our subsequent calculations.," In this section we briefly review the non-linear effect of the relative velocities between the baryons and dark matter, as discussed in \citet{Tseliakhovich:2010} and \citet{Tseliakhovich:2010b}, the latter of which we closely follow in our subsequent calculations." + The initial conditions at recombination include significant relative velocities between the barvons ancl the cold dark matter (which we denote e)., The initial conditions at recombination include significant relative velocities between the baryons and the cold dark matter (which we denote $v\bc$ ). + Before the barvons kinematically clecouple from. the radiation. (around 2= 1100). they are carried. along with the photons. while the dark matter moves according to the gravitational growth of fluctuations which has been advancing since matter-radiation equality. (2~ 3200).," Before the baryons kinematically decouple from the radiation (around $z=1100$ ), they are carried along with the photons, while the dark matter moves according to the gravitational growth of fluctuations which has been advancing since matter-radiation equality $z \sim 3200$ )." + At decoupling. the barvonic speed of sound drops precipitously. and the relative velocity then becomes a substantial cllect.," At decoupling, the baryonic speed of sound drops precipitously, and the relative velocity then becomes a substantial effect." + In the standard: picture of Gaussian initial conditions (e.g. from a period of inflation). the density. ancl the components of relative velocity are Gaussian random variables.," In the standard picture of Gaussian initial conditions (e.g., from a period of inflation), the density and the components of relative velocity are Gaussian random variables." + The velocity and density are spatially correlated (at. dillerent. points) since the continuity equation relates the velocity divergence to the density., The velocity and density are spatially correlated (at different points) since the continuity equation relates the velocity divergence to the density. + Indeed. this equation eives an extra factor of L/h in the velocity (where & is the wavenumber). making the velocity Ποιά coherent on larger scales than the density.," Indeed, this equation gives an extra factor of $1/k$ in the velocity (where $k$ is the wavenumber), making the velocity field coherent on larger scales than the density." + Specifically. velocity Huctuations have significantD power over the range5 &0.010.5 7.," Specifically, velocity fluctuations have significant power over the range $k \sim 0.01-0.5$ $^{-1}$." + The relative velocity is thus coherent on scales smaller than ~3 comoving Alpe., The relative velocity is thus coherent on scales smaller than $\sim 3$ comoving Mpc. +" We therefore analyze probability distributions in such coherent patches. and refer to the uniform relative velocity within each patch as the ""bulk? or ""streaming"" velocity."," We therefore analyze probability distributions in such coherent patches, and refer to the uniform relative velocity within each patch as the “bulk” or “streaming” velocity." + The magnitude of the bulk velocity in each coherence patch at recombination is distributed according to à Maxwell-Doltzmann distribution function: (red = where σος30 km is the root-mean-square velocity at recombination., The magnitude of the bulk velocity in each coherence patch at recombination is distributed according to a Maxwell-Boltzmann distribution function: ) = ^2 where $\sigma_{ v\bc}\sim 30 $ km $^{-1}$ is the root-mean-square velocity at recombination. + Just like any peculiar velocity. the bulk velocity en. decays as (1|2) with the expansion of the universe.," Just like any peculiar velocity, the bulk velocity $v\bc$ decays as $(1+z)$ with the expansion of the universe." + In addition to the bulk velocity. within each patch there are small-scale peculiar velocities of the barvons and dark matter related to the evolution of perturbations (ancl formation of halos) within the patch.," In addition to the bulk velocity, within each patch there are small-scale peculiar velocities of the baryons and dark matter related to the evolution of perturbations (and formation of halos) within the patch." + As was shown in the above references. inside cach coherent region the linear evolution equations for density ancl velocity. perturbations are modified.," As was shown in the above references, inside each coherent region the linear evolution equations for density and velocity perturbations are modified." + For example. on small scales the nonlinear term in the continuity. equation hat couples the local density to the velocity. field. etv- Vo. is comparable to linear terms such as the velocity erm e.IVv.," For example, on small scales the nonlinear term in the continuity equation that couples the local density to the velocity field, $a^{-1} \bf{v} \cdot \bf{\nabla} \delta$ , is comparable to linear terms such as the velocity term $a^{-1}\bf{\nabla} \cdot +\bf{v}$." +" The leading contribution of the nonlinear erm comes from the bulk motion (aIvy,: VO) and this contribution is then linear in terms of the perturbations within the patch.", The leading contribution of the nonlinear term comes from the bulk motion $a^{-1} \bf{v}\bc \cdot \bf{\nabla} \delta$ ) and this contribution is then linear in terms of the perturbations within the patch. + As a result. the evolution equations for he perturbations inside a coherent patch are still linear but dependent on the bulk ene.," As a result, the evolution equations for the perturbations inside a coherent patch are still linear but dependent on the bulk $v\bc$." + Phe resulting velocitv-dependent erms were previously neglected but must be included when structure on small scales and at high redshifts is considered., The resulting velocity-dependent terms were previously neglected but must be included when structure on small scales and at high redshifts is considered. + The relative velocity effect. is. particularly important or the formation of the first. stars ancl galaxies., The relative velocity effect is particularly important for the formation of the first stars and galaxies. + As the irst barvonic objects try to form. they must do so in a moving background of the dark matter potential wells.," As the first baryonic objects try to form, they must do so in a moving background of the dark matter potential wells." + This relative motion means that the dark matter's gravity must work harder in order to trap the barvons., This relative motion means that the dark matter's gravity must work harder in order to trap the baryons. + As a result. the formation of the first bounded barvonie objects is delaved.," As a result, the formation of the first bounded baryonic objects is delayed." + The clleet. though. is less relevant for structure formation today. since the relative velocity decavs with time while the typical mass of galactic host. halos increases.," The effect, though, is less relevant for structure formation today, since the relative velocity decays with time while the typical mass of galactic host halos increases." + However. the relative motion may shift slightly the positions of the BAO peaks and. produce a unique signature in the bispectrum of galaxies (Yoo.Dalal&Seljak 2011)..," However, the relative motion may shift slightly the positions of the BAO peaks and produce a unique signature in the bispectrum of galaxies \citep{Yoo:2011}. ." +nudertaken new nulti-epoch observations of ITI absorption agalust a set of bright pulsars.,undertaken new multi-epoch observations of HI absorption against a set of bright pulsars. + We chose to observe the same sources as Frail et al. (, We chose to observe the same sources as Frail et al. ( +1991) in order to enbauce the mmuber of available time baselines for comparison.,1994) in order to enhance the number of available time baselines for comparison. + This paper sununuarizes our first results frou this project for three pulsus. BOS23|26. D1133|16 and B2016|25.," This paper summarizes our first results from this project for three pulsars, B0823+26, B1133+16 and B2016+28." + Iu Section 2 we sunmnunnuarize our observing and data processimg strategies., In Section \ref{s:obs} we summarize our observing and data processing strategies. + Results ou individual objects are presented im Section 3. and discussed in Section L.., Results on individual objects are presented in Section \ref{s:results} and discussed in Section \ref{s:discussion}. + We have used the Arecibo to obtain new multi-epoch III absorption measurements agalust six pulsars previously studied by Frail ot., We have used the Arecibo to obtain new multi-epoch HI absorption measurements against six pulsars previously studied by Frail et. + al. (, al. ( +1991).,1994). + The observiug and data processing procedures were verv simular to Stanimirovicctal.(2003).. who measured pulsu OTL absorption profiles at Arecibo with the same backend. the Caltech Baschand Recorder.," The observing and data processing procedures were very similar to \cite{Stanimirovic03}, who measured pulsar OH absorption profiles at Arecibo with the same backend, the Caltech Baseband Recorder." + We had four observing sessious: August 2000. December 2000. September 2001 auc November 2001. searching for IIT absorption profile variations over time intervals from less than a dav to 1.25 vears.," We had four observing sessions: August 2000, December 2000, September 2001 and November 2001, searching for HI absorption profile variations over time intervals from less than a day to 1.25 years." + Data from Frail et., Data from Frail et. + al. (, al. ( +1991) extend the time baselines to a decade.,1994) extend the time baselines to a decade. + Qur basic data product was a cube of pulsar intensity as a function of pulsar rotational phase aud radio frequency., Our basic data product was a cube of pulsar intensity as a function of pulsar rotational phase and radio frequency. +" ""Pulsu-onu aud μαιο spectra were then acciunulated w fiuding the pulsar pulse iu software. aud extracting spectra during the pulse aud between pulses. respectively."," `Pulsar-on' and `pulsar-off' spectra were then accumulated by finding the pulsar pulse in software, and extracting spectra during the pulse and between pulses, respectively." + The pulsar absorption spectim is created by generating he pulsar-on ~pulsav-off spectrum for cach scan. doing frequency switching to flatten the baseline. and accumulatiue all such spectra with a weight proportional ο Tre Where Tpug is the autenna temperature of the oilsar.," The pulsar absorption spectrum is created by generating the `pulsar-on' – `pulsar-off' spectrum for each scan, doing frequency switching to flatten the baseline, and accumulating all such spectra with a weight proportional to $T^2_{\rm PSR}$, where $T_{\rm PSR}$ is the antenna temperature of the pulsar." + >.Final absorption aud emission spectra have velocity resolution of 0.5|., Final absorption and emission spectra have velocity resolution of 0.5. +.. For cach pulsar. we display a single absorption spectrum ou top. and then one or more absorption spectra. differences between two epochs.," For each pulsar, we display a single absorption spectrum on top, and then one or more absorption spectrum differences between two epochs." + Overlaid atop each difference spectrum is a 2-0 sienificance cuvelope., Overlaid atop each difference spectrum is a $\pm$ $\sigma$ significance envelope. + To calculate this expected noise level. we have taken into account the following contributions. (," To calculate this expected noise level, we have taken into account the following contributions. (" +1) The skv vackeround contribution was estimated from the all kv survey at [08 MIIZ by Haslametal.(1982). aud scaled o Lt CIIz using a spectral index of 2.6. (,1) The sky background contribution was estimated from the all sky survey at 408 MHz by \cite{Haslam82} and scaled to 1.4 GHz using a spectral index of $-2.6$. ( +2) The contribution from ΠΠ was estimated from the ~pulsar-off spectra and scaled appropriately to match xeviouslv published observations.,2) The contribution from HI was estimated from the `pulsar-off' spectra and scaled appropriately to match previously published observations. + The effect of this contribution can be very significant., The effect of this contribution can be very significant. + For example the mus roise online is three times higher than the rius noise offine for the case of D2016|25. (, For example the rms noise on–line is three times higher than the rms noise off-line for the case of B2016+28. ( +3) The pulsar contin Cluission itself is also suffcieutlv strong to contribute an additional ~—1030% to the noise temperature.,3) The pulsar continuum emission itself is also sufficiently strong to contribute an additional $\sim10-30\%$ to the noise temperature. + Tuformation on individual objects aud au upper linüt on Ar are given in Table 1 Gaote that these are 2-0 limits at the III lino)., Information on individual objects and an upper limit on $\Delta\tau$ are given in Table 1 (note that these are $\sigma$ limits at the HI line). + Our measurements aud the Frail et al. (, Our measurements and the Frail et al. ( +1991) results are discussed for each source below.,1994) results are discussed for each source below. + Fig., Fig. + l (loft) compares HI absorption profiles obtained toward BOS23)26 in 2000.6 and 2000.9., \ref{f:0823_1133} (left) compares HI absorption profiles obtained toward B0823+26 in 2000.6 and 2000.9. + We fudne sjeuificaut change in absorption over this time span down to a Ar level of about 0.01., We find significant change in absorption over this time span down to a $\Delta\tau$ level of about 0.04. + The time iuterval of 0.3 vr translates to a transverse distance of about LO AU., The time interval of 0.3 yr translates to a transverse distance of about 10 AU. + Frail et al. (, Frail et al. ( +1991) also found almost no variations over a period of 0.6 vr. butdid report variations of about 0.07 over a period of 1.1 vr.,"1994) also found almost no variations over a period of 0.6 yr, but report variations of about 0.07 over a period of 1.1 yr." + See the Discussion section below for an cxanunation of the discrepancies between the two eroups? results., See the Discussion section below for an examination of the discrepancies between the two groups' results. + Fig., Fig. + l (Gieht) displays HI absorption profiles obtained toward BLl133|16 in 2000.6 and 2000.9., \ref{f:0823_1133} (right) displays HI absorption profiles obtained toward B1133+16 in 2000.6 and 2000.9. + We detectπο variations down to a Ar level of about 0.02., We detect variations down to a $\Delta\tau$ level of about 0.02. + Diving this period. the pulsar traveled 20 AU.," During this period, the pulsar traveled 20 AU." + Frail et al. (, Frail et al. ( +1991) also detected no significant variations on their 0.6 vr baseline ou this pulsar. butdid sco variations iu 7 of about 0.01over a period of L.1 xr.,"1994) also detected no significant variations on their 0.6 yr baseline on this pulsar, but see variations in $\tau$ of about 0.04 over a period of 1.1 yr." + We investigate the discrepaucy. below., We investigate the discrepancy below. +" For this pulsar. we have compared ITE absorption profiles on our usual short (~0.3 vr) timescale. but also on decadelong scales,"," For this pulsar, we have compared HI absorption profiles on our usual short $\sim0.3$ yr) timescale, but also on decade–long scales." + On the loft side of Fig. 2..," On the left side of Fig. \ref{f:2016}," + iu the nuddle pauel we show the usual difference of absorption spectra frou 2000.6 and 2000.9. while the bottom left panel investigates an 8.7 vear interval bv cdifferenciusg our 2000.6 spectrum and the Frail et al. (," in the middle panel we show the usual difference of absorption spectra from 2000.6 and 2000.9, while the bottom left panel investigates an 8.7 year interval by differencing our 2000.6 spectrum and the Frail et al. (" +1991). 1991.9 data.,1994) 1991.9 data. + The 0.3 vy baseline. corresponding to à —3 AU scale. shows ouly a miarginal 2.6-0 chauge in absorption. while the 8.7 vr. TO AU baseline. does not exhibit amy sienificant variations at all.," The 0.3 yr baseline, corresponding to a $\sim 3$ AU scale, shows only a marginal $\sigma$ change in absorption, while the 8.7 yr, $\sim 70$ AU baseline, does not exhibit any significant variations at all." + This result is very different from Frail et al. (, This result is very different from Frail et al. ( +199L). who found very large optical depth variations of Ar>1 over periods of 0.6 aud 1.7 vi.,"1994), who found very large optical depth variations of $\Delta\tau \ga 1$ over periods of 0.6 and 1.7 yr." + The Frail et al. (, The Frail et al. ( +1991) epoch 1991.9 absorption προςπι that we use for the above comparison is particularly different from the previous three Frail et al.,1994) epoch 1991.9 absorption spectrum that we use for the above comparison is particularly different from the previous three Frail et al. + epochs (their fist one actually being from Cliftonetal. (1988)))., epochs (their first one actually being from \cite{Clifton88}) ). + While it agrees well with our epoch ~2000 results. their 1991.9 profile exhibits siguificautlv shallower absorption in all four principal absorption features than do their earlier epoch results.," While it agrees well with our epoch $\sim2000$ results, their 1991.9 profile exhibits significantly shallower absorption in all four principal absorption features than do their earlier epoch results." + Frail et al. (, Frail et al. ( +1991) remarked that the apparent change could be caused by an incorrect normalization but were confident that was not the case.,1994) remarked that the apparent change could be caused by an incorrect normalization but were confident that was not the case. + We further investigate the apparent Frail ot al., We further investigate the apparent Frail et al. + variatious by comparing our 2000.6 spectra with the pre-1991.90 Frail et al., variations by comparing our 2000.6 spectrum with the pre-1991.9 Frail et al. + profiles., profiles. + As expected. there is a laree difference in the depth of all four absorption features. as shown in the top plot ou the right side of Fig.," As expected, there is a large difference in the depth of all four absorption features, as shown in the top plot on the right side of Fig." + 2. where the 1990.9 Frail et al., \ref{f:2016} where the 1990.9 Frail et al. + spectrum is plotted with a dot-dashed line and our 2000.6 result is plotted as a thin solid line., spectrum is plotted with a dot-dashed line and our 2000.6 result is plotted as a thin solid line. + Since all four lines exhibit the same trend. it seenis reasonable that the differences may result from a slight calibration problem in the Frail et al.," Since all four lines exhibit the same trend, it seems reasonable that the differences may result from a slight calibration problem in the Frail et al." + data., data. + To test this hvpothnesis. we performed a leastsquares fit for a suele scale factor (plus an offset) that would iiniuiüze the difference between these two spectra.," To test this hypothesis, we performed a least–squares fit for a single scale factor (plus an offset) that would minimize the difference between these two spectra." + After applviug this fitted scale factor (0.900+4 0.005) and a constant (0.097+ 0.001) to the 1990.9 spectrmim. we arrive at the “scaled” 1990.9 spectrum. shown as a thick," After applying this fitted scale factor $0.900\pm0.005$ ) and a constant $0.097\pm0.004$ ) to the 1990.9 spectrum, we arrive at the “scaled” 1990.9 spectrum, shown as a thick" +the effective temperature). the contribution of low luminosity white dwarfs to the optical depth decreases.,"the effective temperature), the contribution of low luminosity white dwarfs to the optical depth decreases." + Finally. the top right panel of Fig.," Finally, the top right panel of Fig." + 2 shows the relative contributions to the microlensing optical depth when model C ts considered., 2 shows the relative contributions to the microlensing optical depth when model C is considered. + As can be seen. the results are virtually indistinguishable of those obtained for model B. as one should expect given that the cooling tracks of Bergeron et al. (," As can be seen, the results are virtually indistinguishable of those obtained for model B, as one should expect given that the cooling tracks of Bergeron et al. (" +1995) are very similar to those of Benvenuto Althaus (1997).,1995) are very similar to those of Benvenuto Althaus (1997). + In all cases it is important to realize that as far as the entire population is concerned. there is a noticeable increase in the contribution to the optical depth. which ts exclusively due to the inclusion of non-DA white dwarfs in our calculations.," In all cases it is important to realize that as far as the entire population is concerned, there is a noticeable increase in the contribution to the optical depth, which is exclusively due to the inclusion of non-DA white dwarfs in our calculations." + Moreover. the global contribution of white dwarfs to the microlensing optical depth is very similar in both models — see the bottom panels of Fig.," Moreover, the global contribution of white dwarfs to the microlensing optical depth is very similar in both models — see the bottom panels of Fig." +" 2 — for magnitude cuts larger than zy~23"" which is a reasonable value for current surveys."," 2 — for magnitude cuts larger than $m_V\sim 23^{\rm mag,}$ which is a reasonable value for current surveys." + This value. roughly 30% of the observed optical depth obtained by the MACHO team. represents à 50% increment with respect to the value found in our previous studies. see Torres et al. (," This value, roughly $30\%$ of the observed optical depth obtained by the MACHO team, represents a $50\%$ increment with respect to the value found in our previous studies, see Torres et al. (" +2008).,2008). + A more detailed information can be obtained from our simulations., A more detailed information can be obtained from our simulations. + A summary Is presented in Table 1. where we show several interesting parameters for the three models under study as a function of the adopted magnitude cut.," A summary is presented in Table 1, where we show several interesting parameters for the three models under study as a function of the adopted magnitude cut." + In particular we show in this table the expected number of white dwarf microlensing events. the number of microlensing events produced by red dwarfs. the average mass of the microlenses for both the microlensing events produced by white and red dwarfs. the fraction of the white dwarf microlensing events produced by white dwarfs of the non-DA spectral type Gp over the total white dwarf microlensing events. the average proper motion. distance and tangential velocity of the lenses. the corresponding Einstein crossing times of the microlenses and finally the relative contribution to the microlensing optical depth.," In particular we show in this table the expected number of white dwarf microlensing events, the number of microlensing events produced by red dwarfs, the average mass of the microlenses for both the microlensing events produced by white and red dwarfs, the fraction of the white dwarf microlensing events produced by white dwarfs of the non-DA spectral type $\eta$ ) over the total white dwarf microlensing events, the average proper motion, distance and tangential velocity of the lenses, the corresponding Einstein crossing times of the microlenses and finally the relative contribution to the microlensing optical depth." + A close inspection of Table | reveals that all three models produce similar results except in one aspect. the fraction of microlensing events attributable to a non-DA white dwarf.," A close inspection of Table 1 reveals that all three models produce similar results except in one aspect, the fraction of microlensing events attributable to a non-DA white dwarf." + As can be seen. the expected number of microlensing events Is very small in all models. since in all cases no more than one microlensing event is expected to be found at the Io confidence level.," As can be seen, the expected number of microlensing events is very small in all models, since in all cases no more than one microlensing event is expected to be found at the $1\sigma$ confidence level." + Additionally. the average masses of the microlenses are around 0.6Me in the case of white dwarfs. while for red dwarfs it is ~0.2Me. the average distances to the microlenses are also very similar for both models. and there are no significant differences in the Einstein crossing times.," Additionally, the average masses of the microlenses are around $0.6\, M_{\sun}$ in the case of white dwarfs, while for red dwarfs it is $\sim 0.2\, M_{\sun}$, the average distances to the microlenses are also very similar for both models, and there are no significant differences in the Einstein crossing times." + The only relevant difference between the simulations is the spectral type of the white dwarf responsible for the simulated microlensing events., The only relevant difference between the simulations is the spectral type of the white dwarf responsible for the simulated microlensing events. + Whereas for model A the DA type prevails in ~73% of the cases. for model B this fraction drops to ~20% of the cases. while for model C we obtain a very similar value. ~21%.," Whereas for model A the DA type prevails in $\sim 73\%$ of the cases, for model B this fraction drops to $\sim +20\%$ of the cases, while for model C we obtain a very similar value, $\sim 21\%$." + This can be understood by the same reasoning employed before., This can be understood by the same reasoning employed before. + For models B and C. at low effective temperatures the fraction of hydrogen-rich white dwarfs is considerably smaller than for model A and. additionally. old DÀ white dwarfs are brighter than non-DAs.," For models B and C, at low effective temperatures the fraction of hydrogen-rich white dwarfs is considerably smaller than for model A and, additionally, old DA white dwarfs are brighter than non-DAs." + Thus for models B and C non-DA white dwarfs dominate at low lummosities anc produce most of the microlensing events., Thus for models B and C non-DA white dwarfs dominate at low luminosities and produce most of the microlensing events. + Additionally. from a detailed analysis of the data used to build Table 1. we have found that on average the microlenses produced by non-DA white dwarfs have slightly higher average masses (~0.61Mc; and ~0.56 Mo. respectively) and can be found at a smaller distances (~1.7 kpe and ~2.9 kpe. respectively) than those produced by the population of DA white dwarfs.," Additionally, from a detailed analysis of the data used to build Table 1, we have found that on average the microlenses produced by non-DA white dwarfs have slightly higher average masses $\sim 0.61\, +M_{\sun}$ and $\sim 0.56\, M_{\sun}$ , respectively) and can be found at a smaller distances $\sim 1.7$ kpc and $\sim 2.9$ kpc, respectively) than those produced by the population of DA white dwarfs." + That is again a consequence of the different cooling rates and colors of non-DA white dwarts., That is again a consequence of the different cooling rates and colors of non-DA white dwarfs. + As previously mentioned. non-DA white dwarfs cool faster and moreover. as they cool. they become substantially dimmer than their corresponding DA counterparts.," As previously mentioned, non-DA white dwarfs cool faster and moreover, as they cool, they become substantially dimmer than their corresponding DA counterparts." + Hence. the population of non-DA white dwarfs can produce microlenses at significantly smaller distances.," Hence, the population of non-DA white dwarfs can produce microlenses at significantly smaller distances." + Also. these values do not depend significantly on the model adopted for the evolution of the atmospheric composition of white dwarfs.," Also, these values do not depend significantly on the model adopted for the evolution of the atmospheric composition of white dwarfs." + Since the distribution of velocities does not depend on the spectral type. the final result is that the Einstein crossing times are on average different for the microlensing events produced by non-DA and DA white dwarfs (~40 and ~57days. respectively).," Since the distribution of velocities does not depend on the spectral type, the final result is that the Einstein crossing times are on average different for the microlensing events produced by non-DA and DA white dwarfs $\sim 40$ and $\sim 57$days, respectively)." +"For the sake of comparison. we also compute (he [ree magnetic enerev derived [rom the magnetic virial theorem. assuming a force-Iree field (ο.ο,Alv1989:Ixlimchuketal.1992:Metcalfοἱal.1995.2005:Wheatland& 2006).","For the sake of comparison, we also compute the free magnetic energy derived from the magnetic virial theorem assuming a force-free field \citep[e. +g.][]{aly89, kli92, met95, met05, whe06}." +.. Considering that the magnetic field can be decomposed into a potential part ancl a nonpotential one. B=Boe+h. then folowing Alv(1989) the [free magnetic energv (above potential) is: in (he hall-space above (he surface X.," Considering that the magnetic field can be decomposed into a potential part and a nonpotential one, $\vec B = \vec +B_{pot} + \vec b$ , then following \cite{aly89} the free magnetic energy (above potential) is: in the half-space above the surface $\Sigma$." + The free magnetic energy from the virial theorem only requires (he magnetic field distribution on the bottom boundary., The free magnetic energy from the virial theorem only requires the magnetic field distribution on the bottom boundary. + We compute Equ., We compute Eqn. + 6 from either the observed vector magnetic field (not necessarily lorce-[ree) or the reconstructed fielcl on the photosphere., \ref{eq:vir} from either the observed vector magnetic field (not necessarily force-free) or the reconstructed field on the photosphere. + It is important to note that the energy. values derived from ihe magnetic virial (heorem are strongly. influenced by the spatial resolution as mentioned in Whimehuketal.(1992)., It is important to note that the energy values derived from the magnetic virial theorem are strongly influenced by the spatial resolution as mentioned in \cite{kli92}. +. Considering (hat the field is the minimum energy state of the field. the energy values ean be sorted as follows: In Fig. L..," Considering that the field is the minimum energy state of the field, the energy values can be sorted as follows: In Fig. \ref{fig:nrj}," + we plot the free enerey values in the reconstructed magnetic configurations using the potential field as reference field for the fields ((viangles) ancl lields (crosses)., we plot the free energy values in the reconstructed magnetic configurations using the potential field as reference field for the fields (triangles) and fields (crosses). + The difference between the two values is the minimum [ree energv AE! according to Woltjers theorem., The difference between the two values is the minimum free energy $\Delta E_{lff}^{nlff}$ according to Woltjer's theorem. + Figure 1. clearly shows that the [ree magnetic energv can vary by ad least 2 orders of magnitude: the energy is stronglv influenced by the total magnetic fIux and the distribution of the polarities., Figure \ref{fig:nrj} clearly shows that the free magnetic energy can vary by at least 2 orders of magnitude: the energy is strongly influenced by the total magnetic flux and the distribution of the polarities. + By comparing the amount of [ree energv AE and the observed eruptive phenomena. we can conclude that AED eives a better estimate of the [ree energv.," By comparing the amount of free energy $\Delta E_{pot}^{nlff}$ and the observed eruptive phenomena, we can conclude that $\Delta E_{lff}^{nlff}$ gives a better estimate of the free energy." + For instance. AE is similar for ARS151 and ARS210 but AED is nearly three times larger for ARS?10.," For instance, $\Delta E_{pot}^{nlff}$ is similar for AR8151 and AR8210 but $\Delta E_{lff}^{nlff}$ is nearly three times larger for AR8210." + And the related eruptive phenomena are very different: a slow filament eruption without a flare for ARSI51 and a C-class flare for ARS210., And the related eruptive phenomena are very different: a slow filament eruption without a flare for AR8151 and a C-class flare for AR8210. +" For ARQOTT. still(Pe enough ∙to trigger an N-class∙ flare butUU certainly not theAE! N5.7 flare observedjs ⋅ ⋅ ⋅ νο »"" is sienilicantl"," For AR9077, $\Delta E_{lff}^{nlff}$ is still enough to trigger an X-class flare but certainly not the X5.7 flare observed prior to the time considered here." +"y ↽reduced. compared lo↕↽≻↕⋅↕∪↕⋅↥∪⊔∐↲⊔∐∐↲≺∢∪∐⋟∖⊽↕≺⇂≼↲↕⋅≼↲≺⊔∐↲↕⋅≼↲⋅⊟≻↕⋅↼∐≹∐⊔⋈≻⋅ AB"" 7but still enough to trigger powerful [lares which explains the high level of activity in this active region (Metcalfetal.2005)."," For AR10486, $\Delta E_{lff}^{nlff}$ is significantly reduced compared to $\Delta E_{pot}^{nlff}$ but still enough to trigger powerful flares which explains the high level of activity in this active region \citep{met05}." +. In Table 1.. we summarize the different. values of [ree magnetic energy. the magnetic enerey of the magnetic configurations νο and the relative magnetic helicitv.," In Table \ref{tab:nrj}, we summarize the different values of free magnetic energy, the magnetic energy of the magnetic configurations $E_m^{nlff}$ ) and the relative magnetic helicity." + We also mention (he a values used (o compute the fields satisfying Woltjers theorem., We also mention the $\alpha$ values used to compute the fields satisfying Woltjer's theorem. + We notice that the different. values of [ree οποιον are consistentanc increase when the, We notice that the different values of free energy are consistentand increase when the +accretion streams.,accretion streams. +" Iu steady state. the iuncer οσο of the disk coimekles with the eravitational ""X-ρολ! of the svsteii."," In steady state, the inner edge of the disk coincides with the gravitational ``X-point'' of the system." + Asstmine that the flows eunuiate roni this location allows a senmi-analvtie theory to be developed that vields a clear picture of outflows from ow-lnass YSOs., Assuming that the flows emanate from this location allows a semi-analytic theory to be developed that yields a clear picture of outflows from low-mass YSOs. + Shu et al. (, Shu et al. ( +1995) showed that the density asvinptotes to cvlindrical contours while the streamlines ecole racial.,1995) showed that the density asymptotes to cylindrical contours while the streamlines become radial. + The physical couditious (temperature and ionization) of the N-wind were developed by Shang et al. (, The physical conditions (temperature and ionization) of the X-wind were developed by Shang et al. ( +2002: henceforth SCSL) using Shanes (1998) seni-analytic represcutation of the flow.,2002; henceforth SGSL) using Shang's (1998) semi-analytic representation of the flow. + SGSL assumed the eas is atomic and heated bv shocks and ionized bw X-ravs., SGSL assumed the gas is atomic and heated by shocks and ionized by X-rays. + They were able to give a good account of the optical observatious of the forbidden lines aud of the radio continu observations of the Class I source. L1551 IRS 5 (Shane et al.," They were able to give a good account of the optical observations of the forbidden lines and of the radio continuum observations of the Class I source, L1551 IRS 5 (Shang et al." + 2001)., 2004). + Iu addition to N-xavs. SCSL considered a variety of jieating and ionization sources (listed in their Table 1).," In addition to X-rays, SGSL considered a variety of heating and ionization sources (listed in their Table 1)." + For example. iu addition to N-ravs. they treated IT xioto-detachineut. Balmer coutiuuunu photoionization xostellar photous aud by UV radiation from accretion- hot spots. as well as collisional ionization.," For example, in addition to X-rays, they treated $^-$ photo-detachment, Balmer continuum photoionization by stellar photons and by UV radiation from accretion-funnel hot spots, as well as collisional ionization." + The X-ravs can jonize a large part of the flow because radiative reconibiation occurs on a much longer time scale thui he flow time scale (Bacciotti ct al., The X-rays can ionize a large part of the flow because radiative recombination occurs on a much longer time scale than the flow time scale (Bacciotti et al. + 1995)., 1995). + Thus a sieuificaut level of ionizatiou generated near the source of the flow is frozen into the wind out to larger distances., Thus a significant level of ionization generated near the source of the flow is frozen into the wind out to larger distances. + Mauy heating aud cooling processes were considered. but SCSL found that the most iniportant were: (1) adiabatic cooling of the expaudiug flow aud (2) heating via the dissipation of internal velocity fluctuations aud shocks.," Many heating and cooling processes were considered, but SGSL found that the most important were: (1) adiabatic cooling of the expanding flow and (2) heating via the dissipation of internal velocity fluctuations and shocks." +" SCSL formulated the latter with a phenomenological formmla for the mechanical heating rate per unit volume. where p aud e are the local gas deusity and flow velocity. s is the distance the fluid clement has traveled along a streamline to the poiut of interest. aud à,«1 is a phenomenological coefficient that characterizes the magnitude of the mechanical heating."," SGSL formulated the latter with a phenomenological formula for the mechanical heating rate per unit volume, _h where $\rho$ and $v$ are the local gas density and flow velocity, $s$ is the distance the fluid element has traveled along a streamline to the point of interest, and $\alpha_h \ll 1$ is a phenomenological coefficient that characterizes the magnitude of the mechanical heating." + Equation (2)) may be considered a prescription for shock heating in the jet., Equation \ref{mech_heat}) ) may be considered a prescription for shock heating in the jet. +" Tf velocity fluctuations of magnitude de generate the mechanical heating. theu opc?~(δυο, and the values of ay, needed to heat the flow. 0.001. correspond to moderate velocity fluctuations. δυο~0.1."," If velocity fluctuations of magnitude $\delta v$ generate the mechanical heating, then $\alpha_h v^3\sim (\delta v)^3$, and the values of $\alpha_h$ needed to heat the flow, $\sim 0.001$, correspond to moderate velocity fluctuations, $\delta v/v \sim 0.1$." + Values of this order may be inferred frou high-spatial resolution observations of the optical forbidden lines emitted bv revealed jets (e.g... Bacciotti et al.," Values of this order may be inferred from high-spatial resolution observations of the optical forbidden lines emitted by revealed jets (e.g., Bacciotti et al." + 2000: Woitas ct al., 2000; Woitas et al. + 2002)., 2002). + We use the SGSL model to calculate the neon finc-structure Lue cussion for the wind aud jet of a low-mass YSO., We use the SGSL model to calculate the neon fine-structure line emission for the wind and jet of a low-mass YSO. + We also compare the neon Hines with À6300. representative of optically forbidden lines.," We also compare the neon lines with $\lambda6300$, representative of optically forbidden lines." + The results depeud ou the model parameters listed in Table 1., The results depend on the model parameters listed in Table 1. + The most important are the wind mass-loss rate and X-rav huninosity., The most important are the wind mass-loss rate and X-ray luminosity. +" In Table 1. both the wind mass-loss rate Ad, and the N-rav bhuuduositv Lx are for a onc-siced jet."," In Table 1, both the wind mass-loss rate $\Mw$ and the X-ray luminosity $\LX$ are for a one-sided jet." + The ummerical values iu the table are based ou prior studies of jets (SCSL: Shane et al., The numerical values in the table are based on prior studies of jets (SGSL; Shang et al. + 2001)., 2004). + We use the fiducial case of SGSL as our own reference model for a revealed T Tauri source that uudergoes active accretion and drives a bright optical jet., We use the fiducial case of SGSL as our own reference model for a revealed T Tauri source that undergoes active accretion and drives a bright optical jet. + SGSL conceived the N-rav enüssion to arise from a soft coronal source of climeusion £A. cuhauced by magnetic reconnection regions. one located im the disk iiidplaue aud two others associated with magnetic field Y-coufiguratious above and below the plane (see the schematic drawing of Shu et al.," SGSL conceived the X-ray emission to arise from a soft coronal source of dimension $R_\ast$ enhanced by magnetic reconnection regions, one located in the disk midplane and two others associated with magnetic field $Y$ -configurations above and below the plane (see the schematic drawing of Shu et al." + 1997)., 1997). + They idealized this complex as a series of point sources. one at the origin aud the others displaced alone the axis at heights +tRy with Rx~0.07 AAU.," They idealized this complex as a series of point sources, one at the origin and the others displaced along the axis at heights $\pm R_{\rm X}$ with $R_{\rm X}\sim 0.07 $ AU." + The total N-vayv luuinesity of cach source was composed equally of soft anc hard N-ravs. described respectively bv thermal spectra witli Ty= lkkeV and Ty=2kkeV. Because of the σα. dimensions of all of the sources (~Rx). they behave auch like a single coronal X-ray source with an additional hard component.," The total X-ray luminosity of each source was composed equally of soft and hard X-rays, described respectively by thermal spectra with $T_{\rm X} = 1$ keV and $T_{\rm X} = 2$ keV. Because of the small dimensions of all of the sources $\sim R_{\rm X}$ ), they behave much like a single coronal X-ray source with an additional hard component." + The soft N-rays are absorbed over relatively short distances for outflows with appreciable mass loss., The soft X-rays are absorbed over relatively short distances for outflows with appreciable mass loss. + According to the reference uodel. the more penetrating hard XN-ravs mradiate one remisphere with a huninositv. Ly~1.5«10ergs|.," According to the reference model, the more penetrating hard X-rays irradiate one hemisphere with a luminosity, $\LX \sim 1.5\times 10^{31}\, \ergps$." + The relatively lavee fiducial value of the Xaax tuinosity in Table 1 was adopted by SCSL to treat active YSOs with bright jets and large accretion rates., The relatively large fiducial value of the X-ray luminosity in Table 1 was adopted by SGSL to treat active YSOs with bright jets and large accretion rates. + Their conception of the N-rav properties of YSOs was onued by the results from the carly N-rav observatories. audROSAT. as discussed bv Shu et al. (," Their conception of the X-ray properties of YSOs was formed by the results from the early X-ray observatories, and, as discussed by Shu et al. (" +1997) who related the properties of the fluctuating N-wiud uodel aud the cmission of N-ravs.,1997) who related the properties of the fluctuating X-wind model and the emission of X-rays. + Following Shu ct al. (, Following Shu et al. ( +1997). SGSLE based their N-ray ποπέν onROSAT observations of embedded sources (Neuhauuser 1997). and tool iuto account the likely effect of absorbing material close to the YSO that reduces the fraction of X-raves that escape to be detected by an N-rvav observatory.,"1997), SGSL based their X-ray luminosity on observations of embedded sources (Neuhäuuser 1997), and took into account the likely effect of absorbing material close to the YSO that reduces the fraction of X-rays that escape to be detected by an X-ray observatory." + Potentially an even more important absorber than the wind is the svstem of acerction columnis. which carry a lareer mass flux (e.e.. Alexander et al.," Potentially an even more important absorber than the wind is the system of accretion columns, which carry a larger mass flux (e.g., Alexander et al." + 2004. 2005).," 2004, 2005)." + Screening by circumstellar material may help explain the lone-standing puzzle concerning the larecer N-rav huuinosities of weak-linedk versus classical TTSs. as suggested earlier by Cali (1980) and by Walter Iuli (1981).," Screening by circumstellar material may help explain the long-standing puzzle concerning the larger X-ray luminosities of weak-lined versus classical TTSs, as suggested earlier by Gahm (1980) and by Walter Kuhi (1981)." + Caeegory et al. (, Gregory et al. ( +"2007) have demoustrated this effect by calculating the propagation of N-ravs through the accretion fuunels of a maeuetospheric accretion model,",2007) have demonstrated this effect by calculating the propagation of X-rays through the accretion funnels of a magnetospheric accretion model. + As discussed above. the calculations reported here cover a wide range of X-ray huninosities aud not just the fiducial level iu Table 1: £x spaus 3 orders of naguitude. from Lx=3«107 to Ly=3«10 eres 1.," As discussed above, the calculations reported here cover a wide range of X-ray luminosities and not just the fiducial level in Table 1: $\LX$ spans 3 orders of magnitude, from $\LX=3\times 10^{29}$ to $\LX=3\times 10^{32}$ $\ps$ ." + This range mav be conrpared with the X-ray huuinosities in, This range may be compared with the X-ray luminosities in +3.9m).,m). +" The ice-cored dust absorption (A~2.8— 3.2um) feature is marginally found, but the bare carbonaceous dust absorption (A~ 3.4m) feature is hardly seen in Fig. 4.."," The ice-cored dust absorption $\lambda\sim2.8-3.2\mu$ m) feature is marginally found, but the bare carbonaceous dust absorption $\lambda\sim3.4\mu$ m) feature is hardly seen in Fig. \ref{specall}." +" Around A~ 3.4um, instead of absorption features, a marginal excess is seen, which we currently have no idea of."," Around $\lambda\sim3.4\mu$ m, instead of absorption features, a marginal excess is seen, which we currently have no idea of." + It is noted that the long-ward of 3.9 wm of the stacked spectrum shows values totally lower than the fit., It is noted that the long-ward of 3.9 $\mu$ m of the stacked spectrum shows values totally lower than the fit. +" The known line features at A> 4um is Bro emission (A= 4.054m) and CO» absorption (A= 4.26um), but those featuresare not clearly identified in Fig. 4.."," The known line features at $\lambda>4\mu$ m is $\alpha$ emission $\lambda=4.05\mu$ m) and $_2$ absorption $\lambda=4.26\mu$ m), but those featuresare not clearly identified in Fig. \ref{specall}." + The apparent ‘dip’ at À> 4um may be simply because the spectral fitting was carried out in the wavelength range of 2.5—3.9um. Figs., The apparent `dip' at $\lambda>4\mu$ m may be simply because the spectral fitting was carried out in the wavelength range of $2.5-3.9\mu$ m. Figs. +" 5 — 7 show the stacked spectra of the BEGs in bins of several properties: [Οτή-based specific star formation rate (SSFRjo1)), 9-l(u—r) color, and optical spectral type."," \ref{specsfr} – \ref{spectype} show the stacked spectra of the BEGs in bins of several properties: }]-based specific star formation rate $_{[\textrm{\protect\tiny OII}]}$ ), $^{0.1}(u-r)$ color, and optical spectral type." +" In the estimation of SSFR, the star formation rate (SFR) was estimated using the [Or] emission line that is known to be hardly contaminated by AGN emission (Ho2005;Kimetal. 2006),, and the stelar mass was derived using the 2MASS K, magnitude (seeLeeetal. 2010a).."," In the estimation of SSFR, the star formation rate (SFR) was estimated using the ] emission line that is known to be hardly contaminated by AGN emission \citep{ho05,kim06}, , and the stellar mass was derived using the 2MASS $K_s$ magnitude \citep[see][]{lee10a}. ." + When considering the objects with sufficient [Orr] emission signal (S/Njoy)>3; Fig., When considering the objects with sufficient ] emission signal $_{[\textrm{\protect\tiny OII}]}\ge3$; Fig. + bbb and Fig., \ref{specsfr}b b and Fig. +" 5cc), PAH emission is hardly found for low SSFRyjon) BEGs, whereas the stacked spectum of the high SSFRjo;; BEGs shows a clear PAH emission feature."," \ref{specsfr}c c), PAH emission is hardly found for low $_{[\textrm{\protect\tiny OII}]}$ BEGs, whereas the stacked spectum of the high $_{[\textrm{\protect\tiny OII}]}$ BEGs shows a clear PAH emission feature." + This confirms the fact that PAH emission reflects current SF., This confirms the fact that PAH emission reflects current SF. +" It is noted that 14 BEGs have low S/N(o,;, which may be mainly due to the difficulty in measuring the (3727À)) line using the SDSS spectroscopy (3800— 9200À)) at low redshift."," It is noted that 14 BEGs have low $_{[\textrm{\protect\tiny OII}]}$, which may be mainly due to the difficulty in measuring the ) line using the SDSS spectroscopy $3800-9200$ ) at low redshift." + Fig., Fig. +" 6 shows that the bluer BEGs have the stronger PAH emission feature, which indicates that the blue colors of BEGs reflect their SF activity."," \ref{specur} shows that the bluer BEGs have the stronger PAH emission feature, which indicates that the blue colors of BEGs reflect their SF activity." + This trend is also confirmed in Fig. 7::, This trend is also confirmed in Fig. \ref{spectype}: +" the PAH equivalent width (EW3,29) of SF BEGs (25.2 nm) is larger than those of Seyfert and LINER BEGs (17.2 nm and 13.4 nm, respectively)."," the PAH equivalent width $_{3.29}$ ) of SF BEGs (25.2 nm) is larger than those of Seyfert and LINER BEGs (17.2 nm and 13.4 nm, respectively)." +" In Fig. 7,,"," In Fig. \ref{spectype}, ," + itis noted that theEW3.29of Seyfert BEGs (17.2+0.8 nm) is larger than that of LINER BEGs (13.4+1.4 nm) by more than 2c., it is noted that the$_{3.29}$of Seyfert BEGs $\pm$ 0.8 nm) is larger than that of LINER BEGs $\pm$ 1.4 nm) by more than $\sigma$ . + Seyferts show, Seyferts show +the FR color aud the observation that V...RzRI (Bessel 1979).,the $V-R$ color and the observation that $V-R\approx R-I$ (Bessel 1979). + The rotation velocity is taken to be 210 kin sἘν σον a linear rise from the ceuter to 6 kpe.," The rotation velocity is taken to be 240 km $^{-1}$, with a linear rise from the center to 6 kpc." +" The bulge model is based ou the ""uall bulee™ of ναι (1989) with a central surface brightness of jp=11. and [=0.72 (Walterbos Iseunicutt 1988: Bessel 1979)."," The bulge model is based on the “small bulge” of Kent (1989) with a central surface brightness of $\mu_R = 14$, and $R-I=0.72$ (Walterbos Kennicutt 1988; Bessel 1979)." +Ro This is an axisvunuctric bulge with a roughly exp(793) falloff in volue density with au effective radius of approximately Lhkpe and axis ratio. ef~0.8.," This is an axisymmetric bulge with a roughly $\exp(-r^{0.4})$ falloff in volume density with an effective radius of approximately 1 kpc and axis ratio, $c/a \sim 0.8$." + Values of the bulge density are normalized to make Mig=[1019AZ...," Values of the bulge density are normalized to make $M_{\rm +bulge} = 4 \times 10^{10} M_\odot$." +" The velocity distribution of bulge stars is taken to be masxwellian GUN.ox»exp (2/2928), withl."," The velocity distribution of bulge stars is taken to be maxwellian $dN\propto\exp(-v^2/2\sigma^2)d^3\vec{v}$ ), with." +. These quautities are failv well known. aud unlikely to chauge the results by a large amount if revised.," These quantities are fairly well known, and unlikely to change the results by a large amount if revised." + We explore a paramcterized set of M31 halo models., We explore a parameterized set of M31 halo models. + Each model halo is an axisviunuetric cored “isothermal sphere” deteriuiued by two parameters: the axis ratio or flattening 4 (where gq=1 indicates no flattening) aud the core radius r.., Each model halo is an axisymmetric cored “isothermal sphere” determined by two parameters: the axis ratio or flattening $q$ (where $q=1$ indicates no flattening) and the core radius $r_c$. +" As we are concerned with lensing objects we also define the MACTIO fraction (f),) as the fraction of the halo mass that consists of lensing objects.", As we are concerned with lensing objects we also define the MACHO fraction $f_b$ ) as the fraction of the halo mass that consists of lensing objects. + Together. the mass density of halo lensing objects is then where ο=v1q? aud Vise)=210 lau loas taken from observations of the M31 disk.," Together, the mass density of halo lensing objects is then where $e=\sqrt{1-q^2}$ and $V_c(\infty)=240$ km $^{-1}$ is taken from observations of the M31 disk." + Thevelocity distribution of the halo is taken to be maxwelliau. with a circular velocity equal to Vitec). making o.=170 kin sod ," Thevelocity distribution of the halo is taken to be maxwellian, with a circular velocity equal to $V_c(\infty)$, making $\sigma=170$ km $^{-1}$ ." +We compute rate distributions with respect to two cüffereut timescales., We compute rate distributions with respect to two different timescales. +" The first is the more famuliar Einstein time (fg). defined as the time to cross the full Eiisteiu disk which has racius where wis the fractional distance to the leus in relation totie distauce to the source. D,."," The first is the more familiar Einstein time $(\tE)$, defined as the time to cross the full Einstein disk which has radius where $x$ is the fractional distance to the lens in relation to the distance to the source, $D_s$." + The Eiustein time is then eiven in terms of the lens velocity perpendicular to the ine of sieht. The secoud timescale we use is the fullwidth at half wasinnun of the liehteurve fij. much more easilv measured when the source &elds are crowded. as is tle case for M31.," The Einstein time is then given in terms of the lens velocity perpendicular to the line of sight, The second timescale we use is the full–width at half maximum of the lightcurve $\thalf$, much more easily measured when the source fields are crowded, as is the case for M31." + This timescale is simply related to the Einstein time and to the niuinimni inipact parameter) of the Ίος relative to the line of sight (taken in units of Re)., This timescale is simply related to the Einstein time and to the minimum impact parameter $\beta$ of the lens relative to the line of sight (taken in units of $R_E$ ). +" This timescale is given by (Condolo 1999) with the following definitions The function dw)~9 for all values of 3. with the limiting bchavior Thus. fq.~off. and ds more degenerate than fr due to the dependence on impact paramcter,"," This timescale is given by (Gondolo 1999) with the following definitions The function $w(\beta)\sim\beta$ for all values of $\beta$, with the limiting behavior Thus, $\thalf\sim\beta\tE$, and is more degenerate than $\tE$ due to the dependence on impact parameter." + A further complication is that fy is csscutially the measured timescale., A further complication is that $\thalf$ is essentially the measured timescale. + Determining thle Einstein time requires kuoxcledge of the magnification. which cau be very diffieult iu highly crowded fields as the source stars are lughly blended.," Determining the Einstein time requires knowledge of the magnification, which can be very difficult in highly crowded fields as the source stars are highly blended." + The Einstein time night be inferred with extra effort. either by high resolution Πασάς of the source star or statistically with the £5 technique (Baltz Silk 2000).," The Einstein time might be inferred with extra effort, either by high resolution imaging of the source star or statistically with the $t_\sigma$ technique (Baltz Silk 2000)." + It is crucial to evaluate if these are necessary., It is crucial to evaluate if these are necessary. + Using the model of refinodel. we compute he rate of detectable iicroleusing events.," Using the model of \\ref{model}, we compute the rate of detectable microlensing events." + Two computer codes written entirely independently aud described elsewhere (Baltz Silk 2000: Civ Crotts 2000) have been used. aud produce nearly identical results.," Two computer codes written entirely independently and described elsewhere (Baltz Silk 2000; Gyuk Crotts 2000) have been used, and produce nearly identical results." + A nuuerical iutegration of the rate is performed. over the positious and velocities of the source Qvith fixed. brightuess) aud leus (with fixed mass) along a given liue of sight.," A numerical integration of the rate is performed, over the positions and velocities of the source (with fixed brightness) and lens (with fixed mass) along a given line of sight." + The probability that an eveut with eiven parameters is detected is folded iuto the the integral., The probability that an event with given parameters is detected is folded into the the integral. + Mass functious for the leuses aud Iuninositv fictions for he sources are applied at the eud., Mass functions for the lenses and luminosity functions for the sources are applied at the end. +" With lines of sight spaced at one arciunute intervals. we construct contour maps of the event rate for ""solf chasing by stars (assumune no binary lenses. at lost a itributiou (Mao Paczyvüsskà 1991: Daltz Condolo 2001)) as well as for halo lensing."," With lines of sight spaced at one arcminute intervals, we construct contour maps of the event rate for “self lensing” by stars (assuming no binary lenses, at most a contribution (Mao Paczyńsski 1991; Baltz Gondolo 2001)) as well as for halo lensing." + The self eusine contribution has four logical contributions. talàug sources and lenses from the bulee aud disk.," The self lensing contribution has four logical contributions, taking sources and lenses from the bulge and disk." + The bulec-ποσο contribution dominates the self lensing near the center of the bulge. along the minor axis the bulgc-disk and disk-bulee contributions dominate. and far from the lee the disk-disk contribution dominates (thoueh this is alwavs dominated by the halo coutribution for fy20.05).," The bulge-bulge contribution dominates the self lensing near the center of the bulge, along the minor axis the bulge-disk and disk-bulge contributions dominate, and far from the bulge the disk-disk contribution dominates (though this is always dominated by the halo contribution for $f_b>0.05$ )." + These four iu sumi give a self leusimg rate that is ucarly sviunietrie about both the major aud minor axes of M31., These four in sum give a self lensing rate that is nearly symmetric about both the major and minor axes of M31. + The halo lensing coutributious arise from both the N31 aud Milky Way halos. kusing both disk aud buec stars.," The halo lensing contributions arise from both the M31 and Milky Way halos, lensing both disk and bulge stars." + The Milkv Way coutrilition has a nearly uniforiu optical depth., The Milky Way contribution has a nearly uniform optical depth. + The M31 coutrilition is stronely απλοric. with a senificautlv larecr rate from the far edge of tjo niiluor axis.," The M31 contribution is strongly asymmetric, with a significantly larger rate from the far edge of the minor axis." + We have assuned the following definition for a detectable microleusing event., We have assumed the following definition for a detectable microlensing event. + The MEGA survey most yequently cuyplovs the MDM 2.lin telescope. thus we use its capabilities in the following.," The MEGA survey most frequently employs the MDM 2.4m telescope, thus we use its capabilities in the following." + We asstime hat inteerations totaling three hows are taken twice weekly durius the \31 observing season: this is fairly conservative., We assume that integrations totaling three hours are taken twice weekly during the M31 observing season: this is fairly conservative. + We define an event as a deviation that has woconsecutive samples four standard deviations above he baseline., We define an event as a deviation that has twoconsecutive samples four standard deviations above the baseline. + We assune one arcsecoud secing. which is vpical: the median MDM ποσο is 0.95 arcsecond.," We assume one arcsecond seeing, which is typical: the median MDM seeing is 0.95 arcsecond." + To approximate the seusitivitv of MDM. we assume that a star of Rinaguitucde my=25.2 or J magnitude my=21.8 eives one photoclectron per second. aud furthermore that the noise in the damages is twice the photon counting noise (this is quite couscrvative).," To approximate the sensitivity of MDM, we assume that a star of $R$ magnitude $m_R=25.2$ or $I$ magnitude $m_I=24.8$ gives one photoelectron per second, and furthermore that the noise in the images is twice the photon counting noise (this is quite conservative)." + We take the distance modulus to M31 as D=215. a distance of 795 kpc.," We take the distance modulus to M31 as $D=24.5$, a distance of 795 kpc." + We assume a sky briehtuess of pg=21 aud p;=20 mae 7., We assume a sky brightness of $\mu_R=21$ and $\mu_I=20$ mag $^{-2}$ . + Tere we veiteratethe fact that the source stars are uot resolved. but iustead are typically highly blended.," Here we reiteratethe fact that the source stars are not resolved, but instead are typically blended." + Ouly with difference tuaging can the microlens variability be, Only with difference imaging can the microlens variability be +"The rms scatter in the final pure-emission. images: is about 0.035 electrons/pixel/second (or 2.519TUNVa7[or Fig,fae=138. the average value found byPhillips (1986). for a sample of normal ellipticals).","The rms scatter in the final pure-emission images is about 0.035 electrons/pixel/second (or $2.5 \times 10^{-20}\,{\rm W~m}^{-2}$for $F_{\rm{[N{\sc ii}]}_2}/F_{\rm{H}\alpha}=1.38$, the average value found byPhillips \cite{phi} for a sample of normal ellipticals)." + The Bool and | images were used to extract surface brightness. position angle and ellipticity profiles (see Figure 3)).," The B, R and I images were used to extract surface brightness, position angle and ellipticity profiles (see Figure \ref{surf_207}) )." +" The deviations of the isophotes from a pure elliptic shape were quantified by expanding the intensity. variation along an isophotal ellipse in a fourth order Fourier.series with cocllicicnts S4. S3. C', and C's :"," The deviations of the isophotes from a pure elliptic shape were quantified by expanding the intensity variation along an isophotal ellipse in a fourth order Fourierseries with coefficients $S_4$ , $S_3$ , $C_4$ and $C_3$ :" +for example: (?)) and (2)).,for example: \citealt{2005A&A...437..189V}) ) and ). + Some group II sources display the II] rA line in emission. for example: (?)).," Some group II sources display the I] \\r{A} line in emission, for example: \citealt{2005A&A...436..209A}) )." + These sources all seem to be in a transitional phase from a gas rich flaring disk to a gas poor self-shadowed disk., These sources all seem to be in a transitional phase from a gas rich flaring disk to a gas poor self-shadowed disk. + Roughly half of the known group II sources are in this phase. which means we can infer that half of the life time of the disk of a group II source is spent on the dispersal of the gas.," Roughly half of the known group II sources are in this phase, which means we can infer that half of the life time of the disk of a group II source is spent on the dispersal of the gas." + An for this time scale is the photoevaporation time scale. which is on the order of ~ 10° yyr (2)) A comprehensive study was performed to map the distribution of the gas and dust in the protoplanetary disk around 9953881.," An for this time scale is the photoevaporation time scale, which is on the order of $\sim$ $^6$ yr \citealt{2009ApJ...690.1539G}) ) A comprehensive study was performed to map the distribution of the gas and dust in the protoplanetary disk around 95881." + In we displayed a schematic representation of the disk. which puts all results in perspective.," In we displayed a schematic representation of the disk, which puts all results in perspective." + The AMBER K-band interferometry showed that there is an extended hot inner region. with emission coming from within the sublimation radius., The AMBER K-band interferometry showed that there is an extended hot inner region with emission coming from within the sublimation radius. + The detection of the [OIL] rA indicated that the disk has a flaring gas surface at large distances (from one to tens of AU) from the star., The detection of the I] \\r{A} indicated that the disk has a flaring gas surface at large distances (from one to tens of AU) from the star. + The finding of PAH features in the Spitzer and VISIR spectra confirmed the presence of an illuminated gas surface., The finding of PAH features in the Spitzer and VISIR spectra confirmed the presence of an illuminated gas surface. + The resolved VISIR spectrum traced this surface up to We used the radiative transfer code MCMax (?)) to create a model of the disks density and temperature structure., The resolved VISIR spectrum traced this surface up to We used the radiative transfer code MCMax \citealt{2009A&A...497..155M}) ) to create a model of the disk's density and temperature structure. + Our model satisfactorily reproduced all of our observations., Our model satisfactorily reproduced all of our observations. + The main conclusions that followed from our model are that the inner disk contains most of the (2))., The main conclusions that followed from our model are that the inner disk contains most of the \citealt{2007A&A...473..457D}) +In [5].. a generalization of the Ogawa type inequality [12]. to the parabolic framework has been shown.,"In \cite{Ibrahim09}, a generalization of the Ogawa type inequality \cite{Og03} to the parabolic framework has been shown." + Ogawa inequality cau. be considered as a generalized version iu the Lizorkin-Triebel spaces of the remarkable estimate of Brézzis-CGallouétt-Wainger [1.2). that holds in a limiting case of the Sobolev embedding theorem.," Ogawa inequality can be considered as a generalized version in the Lizorkin-Triebel spaces of the remarkable estimate of Brézzis-Gallouëtt-Wainger \cite{BG80, BW80} that holds in a limiting case of the Sobolev embedding theorem." + The inequality showed in [5.Theorem1.1] provides an estimate of the £* norm of a function in terms of its parabolic BALO norm. with the aid of the square root of the logarithmic dependency of a higher order Sobolev norm.," The inequality showed in \cite[Theorem 1.1]{Ibrahim09} provides an estimate of the $L^{\infty}$ norm of a function in terms of its parabolic $BMO$ norm, with the aid of the square root of the logarithmic dependency of a higher order Sobolev norm." + More precisely. for any vector-valued function f=Vg€Εμ ρρη Nog€L'(g?. ιο exists a constant C=Cí(im.n)>O such that: where Wis the parabolic Sobolev space (we refer to [11]. for the definition and further properties). aud BALO is the parabolic bounded mean oscillation space (defined. via parabolic balls instead of Euclidean ones [5.Definition 2.1])).," More precisely, for any vector-valued function $f=\nabla g \in W_{2}^{2m,m}(\R^{n+1})$, $g\in L^{2}(\R^{n+1})$ with $m,n\in \N^{*}$ , $2m>\frac{n+2}{2}$, there exists a constant $C=C(m,n)>0$ such that: where $W^{2m,m}_{2}$ is the parabolic Sobolev space (we refer to \cite{LSU} for the definition and further properties), and $BMO$ is the parabolic bounded mean oscillation space (defined via parabolic balls instead of Euclidean ones \cite[Definition~2.1]{Ibrahim09}) )." + The aboveinequality reflects a limiting case of Sobolev, The aboveinequality reflects a limiting case of Sobolev +"one invokes more plentiful mini-halos (Ti710 IX). reionization bv Pop II stars can produce a similar Az,= 0.0040.016.","one invokes more plentiful mini-halos $T_{\rm vir} \approx 10^3$ K), reionization by Pop II stars can produce a similar $\Delta \tau_e =$ 0.004–0.016." +" These results suggest that partial reionizalion al z> ccould be achieved by minihalos (Ti210* IX). with a mixture of Pop IHE and Pop IH stars. or a [ransition from one to the other. lor which the rising blue curve gives Az,z 0.020.06 [or foe= 0.10.4."," These results suggest that partial reionization at $z >$ could be achieved by minihalos $T_{\rm vir} \approx 10^3$ K), with a mixture of Pop III and Pop II stars, or a transition from one to the other, for which the rising blue curve gives $\Delta \tau_e \approx$ 0.02–0.06 for $f_{\rm esc} =$ 0.1–0.4." + For these minihalos. Pop II stars produce 7<0.02. but 722 0.020.06 can be achieved either by metal-poor stellar populations in halos with unusually high f; or fo... or by metal-free massive stars in halos with verv low star formation rates and/or escape fraction of ionizing radiation.," For these minihalos, Pop II stars produce $\tau \la 0.02$, but $\tau \approx$ 0.02–0.06 can be achieved either by metal-poor stellar populations in halos with unusually high $f_*$ or $f_{\rm esc}$, or by metal-free massive stars in halos with very low star formation rates and/or escape fraction of ionizing radiation." +" Another robust conclusion [rom Figure 1: is that mini-haloes (Ty,=10° IK) with Pop III stars can easily account lor theenfire additional optical depth (Av,x0.03) needed to explain the WMAP-5 data.", Another robust conclusion from Figure 1 is that mini-haloes $T_{\rm vir} = 10^3$ K) with Pop III stars can easily account for the additional optical depth $\Delta \tau_e \approx 0.03$ ) needed to explain the WMAP-5 data. + Even in cases where reionization begins with Pop ILL stars (elliciencies 10°) and then transitions to Pop II (efficiencies ~107). the optical depths are significant.," Even in cases where reionization begins with Pop III stars (efficiencies $\sim 10^3$ ) and then transitions to Pop II (efficiencies $\sim 10^2$ ), the optical depths are significant." + This is an important astrophivsical issue. since the duration of the epoch dominated by massive first stars is still uncertain. (EVS04).," This is an important astrophysical issue, since the duration of the epoch dominated by metal-free massive first stars is still uncertain (TVS04)." + These results mark a significant departure [rom similar curves for à WMAP-3 cosmology (left panel). owing to the increased small-scale power in à WAIAP-5 cosmology relative to WMAP-3.," These results mark a significant departure from similar curves for a WMAP-3 cosmology (left panel), owing to the increased small-scale power in a WMAP-5 cosmology relative to WMAP-3." +" Although the sealar spectral index ης did not change appreciably from vear 3 to vear 5 of WMADP. both the matter density. Q,,. and the normalization. o4. rose bv several percent. an appreciable effect for power available for low-mass halos."," Although the scalar spectral index $n_s$ did not change appreciably from year 3 to year 5 of WMAP, both the matter density, $\Omega_m$, and the normalization, $\sigma_8$, rose by several percent, an appreciable effect for power available for low-mass halos." +" Thus. reionization occurs slightly earlier in à WALAP-55Hr cosmology (Az,& 12) in our models. lowering the required efficiencies by a factor of about 4."," Thus, reionization occurs slightly earlier in a WMAP-5 cosmology $\Delta z_r \approx$ 1–2) in our models, lowering the required efficiencies by a factor of about 4." +" Our predictions can be compared to those of Haiman Bryan (2006). who used WMAP-3 dala to constrain the efficiency factor. ei. relative (o a ficucial value eg,4;=200 for minihalos."," Our predictions can be compared to those of Haiman Bryan (2006), who used WMAP-3 data to constrain the efficiency factor, $\epsilon_{\rm UV}$, relative to a fiducial value $\epsilon_{\rm mini} = 200$ for minihalos." + Thev argued that the efficiency for the production of ionizing photons must have been reduced bv an order of magnitude. in order to avoid overproducing the optical depth.," They argued that the efficiency for the production of ionizing photons must have been reduced by an order of magnitude, in order to avoid overproducing the optical depth." + However. (heir constraints were based on the WMADP-3 value. 7.z0.09. whereas half that scattering is accounted for from the fully ionized IGM at 2