ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 In addition. case (i) more easily. permits the identification of two planets in 111 resonant orbits.," In addition, case (ii) more easily permits the identification of two planets in 1:1 resonant orbits."3 We adopted approach (ii) in the current analysis., We adopted approach (ii) in the current analysis.4 All of the models considered in this paper incorporate an extra noise parameter. s. that can allow for any additional noise bevond the known measurement uncertaintiesS.," All of the models considered in this paper incorporate an extra noise parameter, $s$, that can allow for any additional noise beyond the known measurement uncertainties."51.4.).. We assume the noise variance is finite. and. adopt a Gaussian distribution with a variance s., We assume the noise variance is finite and adopt a Gaussian distribution with a variance $s^2$.6 Thus. the combination of the known errors and extra noise has a Gaussian distribution with variance —a?|so where e; is the standard deviation of the known noise for VENEi data point.," Thus, the combination of the known errors and extra noise has a Gaussian distribution with variance $= \sigma_i^2 + s^2$, where $\sigma_i$ is the standard deviation of the known noise for $^{\mbox{\tiny th}}$ data point."7. For. example. suppose that the star actually has two planets. and the model assumes only. one is present.," For example, suppose that the star actually has two planets, and the model assumes only one is present."8 In regard to the single jxanet model. the velocity," In regard to the single planet model, the velocity"9the literature (2277).,the literature .10.. Stellar masses of these objects are estimated. by fitting the multi-color photometry to moclel SEDs produced with stellar population svnthesis models resuling in uncertainties of ~0.2 ddexολ., Stellar masses of these objects are estimated by fitting the multi-color photometry to model SEDs — produced with stellar population synthesis models — resulting in uncertainties of $\sim 0.2$ dex.11 Acditionallv. due to the excellent depth and resolution of the JCMOS. images (pixel scale after. resampling of 0.71 and a point spread. funcion PSE] of 0.73 [ull width half maximum. EFWIIM]). we are able to estimate the Sérrsic indices ancl sizes of the objects using theGALFIT code(?).," Additionally, due to the excellent depth and resolution of the NICMOS images (pixel scale after resampling of $0.\arcsec121$ $^{-1}$, and a point spread function [PSF] of $0.\arcsec133$ full width half maximum [FWHM]), we are able to estimate the Sérrsic indices and sizes of the objects using the code."14. Average properies of the sources used in our analysis are listed in Table 1.., Average properties of the sources used in our analysis are listed in Table \ref{tab:source_properties}.15 The selection for the GNS galaxies is based. on mass and recshift. with Li<2.2s.," The selection for the GNS galaxies is based on mass and redshift, with $1.7 < z < 2.8$."16 7jese galaxies were located initially through color selection techniques. such as the Dzlx(2).. ERO and DRG criteria. and laer refined through spectroscopic and. photometric redshifts within the two GOODS fields.," These galaxies were located initially through color selection techniques, such as the BzK, ERO and DRG criteria, and later refined through spectroscopic and photometric redshifts within the two GOODS fields."17 »erform. several tests to ensure that the sample is complete., perform several tests to ensure that the sample is complete.18 A possible bias might. be that extremely custy galaxies. could. be missed. by this criteria due to attenuation. but the deep limiting //-bancl magnitude ereatly exceeds that of he expected upper boum for clusty SMGs?7).. so that we are confident that we are not missing the dustiest ealaxics duc to attenuation.," A possible bias might be that extremely dusty galaxies could be missed by this criteria due to attenuation, but the deep limiting $H$ -band magnitude greatly exceeds that of the expected upper bound for dusty SMGs, so that we are confident that we are not missing the dustiest galaxies due to attenuation."19 Lastly. it is expected that this selection of galaxies closely approximates the true ratio of red to blue galaxies in these mass and redshift ranges.," Lastly, it is expected that this selection of galaxies closely approximates the true ratio of red to blue galaxies in these mass and redshift ranges."20 For more details concerning the selection technique and. possible biases, For more details concerning the selection technique and possible biases.21 We use the publicly available Spizer//MIDPS. map at gun fromthieFarInfravedDecplcortragalactieLegacySurvey( E , We use the publicly available /MIPS map at $\micro$ m from the Far Infrared Deep Extragalactic Legacy Survey (FIDEL ).22The 5-0 point source sensitivity of this map is mms., The $\sigma$ point source sensitivity of this map is mJy.23 We use publicly available Herscehet£/PACS observations of the GOODS-South field from the PACS Evolutionary survey., We use publicly available /PACS observations of the GOODS-South field from the PACS Evolutionary survey.24 The data was re-processecl with he Processing Environment., The data was re-processed with the Processing Environment.256.0.2110).. PEP was designed o provide data in all three PACS bands., PEP was designed to provide data in all three PACS bands.26 Since PACS can only observe in two banels simultaneously al tun (Cred)andeitherTO(b ue)or100 pun (ereen) we Use wo sets of observations to produce maps at all three wavelengths., Since PACS can only observe in two bands simultaneously — at $\micro$ m (red) and either 70 (blue) or $\micro$ m (green) — we use two sets of observations to produce maps at all three wavelengths.27 We combine the available deep observations using the standard PACS pipeline. choosing a high-pass filter xuwameter of 20 for the blue and green. bands. ancl 30 for he red band?2).," We combine the available deep observations using the standard PACS pipeline, choosing a high-pass filter parameter of 20 for the blue and green bands, and 30 for the red band."28. In order o prevent ringing effects around. bright sources caused. by he high-pass filter the pipeline performs an initial crude reduction and automatically masks out the brightest sources in the subsequent iterations of de-elitching and filtering., In order to prevent ringing effects around bright sources caused by the high-pass filter the pipeline performs an initial crude reduction and automatically masks out the brightest sources in the subsequent iterations of de-glitching and filtering.29 The depths of the final maps are 0.31. 0.44. and mmy at TO. 100. and. pun. respectively.," The depths of the final maps are 0.31, 0.44, and mJy at 70, 100, and $\micro$ m, respectively."30 As reported by2.. the relatively strong high-pass filter adopted along with the masking of the bright sources may attenuate the final photometry of faint sources.," As reported by, the relatively strong high-pass filter adopted along with the masking of the bright sources may attenuate the final photometry of faint sources."31" To account for these effects. we produce maps of a Lew. isolated. unmasked. faint point sources. of. cillerent [lux density, using the same parameters as were used in the reduction. of the GOODS-South maps: we then mask these sources and create new maps."," To account for these effects, we produce maps of a few, isolated, unmasked, faint point sources of different flux density, using the same parameters as were used in the reduction of the GOODS-South maps; we then mask these sources and create new maps."32 We use the average ratio of the Lux densities of the same sources in the two maps as our estimate of the attenuation factor due to the high-pass filter., We use the average ratio of the flux densities of the same sources in the two maps as our estimate of the attenuation factor due to the high-pass filter.33 We find that the magnitude of the attenuation müldlv increases for increasing wavelengths. as expected given. the shape of the 1/f noise over the relevant frequency range7," We find that the magnitude of the attenuation mildly increases for increasing wavelengths, as expected given the shape of the 1/f noise over the relevant frequency range."34) Phe estimated attenuation [actors are 0:50. 0.78. and 0.75 at το. 100. and tun. respectively.," The estimated attenuation factors are 0.80, 0.78, and 0.75 at 70, 100, and $\micro$ m, respectively."35 Nolethatastightlydi fferentapproach was followedby?. fillering: they findthatthe filleringmod? ficsthefluresbyl6'A forvery," Note that a slightly different approach was followed by, who perform tests on the red band by adding simulated sources to the timelines before masking and high-pass filtering; they find that the filtering modifies the fluxes by for very faint unmasked point sources."36 {e ane because of the lack of an estimate for the blue and ereen bands from the PEP team. we choose to adopt our three estimated factors for consistency.," Despite the slight disagreement with our finding at $\micro$ m, and because of the lack of an estimate for the blue and green bands from the PEP team, we choose to adopt our three estimated factors for consistency."37 The BLAST maps in consist. of a deop region covering 0.9) ddeg? which completely encompasses the southern sources in the catalog (Figure. 1)). ancl have depths of 11. 9. and GomJ.. at 250. 350. and pun. respecbiveblgC," The BLAST maps in consist of a deep region covering $\sim 0.9$ $^2$ which completely encompasses the southern sources in the catalog (Figure \ref{fig:blast_map}) ), and have depths of 11, 9, and mJy, at 250, 350, and $\micro$ m, respectively."38"?).Ductolargeinstrimentalbeams(36 42. and60"" Jane ) source confusion contributes substantially to the noise in these maps. and is estimated to be eougneia7 21. 17. and mindy in the three bands(2)."," Due to large instrumental beams (36, 42, and ) and steep source counts, source confusion contributes substantially to the noise in these maps, and is estimated to be $\sigma_{\rm confusion} \approx$ 21, 17, and mJy in the three bands."39. Phe BLAST maps using à naive mapniaker(7)., The BLAST maps were made using a naive mapmaker.40. Further details on the REXEinstrumentBayle may be found in2.. while Dight performance and calibration are provided in?.," Further details on the instrument may be found in, while flight performance and calibration are provided in."41. The LABOCA E-CDES Submum Survey provides deep Lundata.wilhanrain.s.deplhlobellerthanl.2 mind yacros sthef ullàül «30l field. with an elective resolution of FEWILIAL.," The LABOCA E-CDFS Submm Survey provides deep $\micro$ m data, with an depth to better than mJy across the full $30\arcmin \times 30\arcmin$ field, with an effective resolution of FWHM."42 For a detailed. description of the instrument see, For a detailed description of the instrument see.43 Stacking is a well established technique for. finding the average properties of objects which individually are uncletectable by using external knowledge of their positions in a mapTT)., Stacking is a well established technique for finding the average properties of objects which individually are undetectable by using external knowledge of their positions in a map.44 We follow the formalism ofMOO.. to which we refer to for a full description of the stacking method.," We follow the formalism of, to which we refer to for a full description of the stacking method."45 Here we summarize he salient features of the technique., Here we summarize the salient features of the technique.46 AIO9 showed that the mean lux density of an external catalog is simply the covariance of the mean-subtracted map, M09 showed that the mean flux density of an external catalog is simply the covariance of the mean-subtracted map47ὃς GGA. In both figures black points correspond to observed data. red points to the spectra reconstructed with the F1.0 EBL model. aud blue points to the spectra reconstructed with ihe F1.6 EBL model.,"3C 66A. In both figures black points correspond to observed data, red points to the spectra reconstructed with the F1.0 EBL model, and blue points to the spectra reconstructed with the F1.6 EBL model."48 The reconstructed spectra are significantly harder compared to the observed ones., The reconstructed spectra are significantly harder compared to the observed ones.49 In particular. in case of high EBL flux (F1.6). the spectra have Dj1 and 0.2 for LES 0229-200 and 3C 664. respectively: i.e. they would be significantly harder than the conventional value of 1.5.," In particular, in case of high EBL flux (F1.6), the spectra have $\Gamma_{\rm int}\simeq 1$ and $0.2$ for 1ES 0229+200 and 3C 66A, respectively; i.e. they would be significantly harder than the conventional value of $1.5$."50 To study the case of LES 0229-200. we have combined the reconstructed. VIE data with archive X-ray and optical data from SWIFT (Taveechioetal.2009) and BeppoSAX (Costamanteetal.2002).. together with Fermi LAT observations (Abdoetal.2009).," To study the case of 1ES 0229+200, we have combined the reconstructed VHE data with archive X-ray and optical data from SWIFT \citep{tavecchio09} and BeppoSAX \citep{costamante02}, together with Fermi LAT observations \citep{abdo09}."51. The observational data are summarized in Fig. 3.., The observational data are summarized in Fig. \ref{fig:0229}.52 We have applied the internal absorption scenario as described in Section ?? (0 reproduce the VIIE spectrum together with X-ray spectrum. for both levels of intergalactic absorption In the case of a soft energy. distribution of protons (pZ 2). The exact location of this cutoff depends on different model parameters. e.g. intrinsic optical depth. Doppler boosting factorele. but to satislv the Fermi upper limits the resulting 5-rav spectral break should occur close to ~100 GeV. Thus. the cutoff in the proton spectrum should be located roughly at," We have applied the internal absorption scenario as described in Section \ref{sec:model} to reproduce the VHE spectrum together with X-ray spectrum, for both levels of intergalactic absorption In the case of a soft energy distribution of protons $p\gtrsim2$ ), The exact location of this cutoff depends on different model parameters, e.g. intrinsic optical depth, Doppler boosting factor, but to satisfy the Fermi upper limits the resulting $\gamma$ -ray spectral break should occur close to $\sim100$ GeV. Thus, the cutoff in the proton spectrum should be located roughly at"53produce a 0.1 M. object that entirely depletes its lithium content within about 10. Myr. whereas for the non accreting counterpart. complete lithium depletion takes more than 50 Myr.,"produce a 0.1 $\msol$ object that entirely depletes its lithium content within about 10 Myr, whereas for the non accreting counterpart, complete lithium depletion takes more than 50 Myr."54 As clearly illustrated in Fig. 3..," As clearly illustrated in Fig. \ref{li_mf01},"55 depending on the initial mass and the burst accretion rate. different episodic acceretion histories can produce objects with the same mass. say 0.1 M... ata same age. say x 50 Myr. exhibitingdepletion.," depending on the initial mass and the burst accretion rate, different episodic acccretion histories can produce objects with the same mass, say 0.1 $\msol$, at a same age, say $\simle$ 50 Myr, exhibiting."56 Similar effects are found for episodic aceretion sequences producing final objects within the entire characteristic domain of fully convective stars. M. «0.35M...," Similar effects are found for episodic accretion sequences producing final objects within the entire characteristic domain of fully convective stars, M $\simle \, 0.35 \msol$."57" If the central temperature exceeds ~ 2-3 10° K. a radiative core develops. because of the opacitiy decrease after the last ""k-bump"" due to metals (C. O. Ne and Fe. see Rogers Iglesias 1992). as explained in Chabrier Baratfe (1997. see their 33.2 and their Fig."," If the central temperature exceeds $\sim$ 2-3 $^6$ K, a radiative core develops, because of the opacitiy decrease after the last $\kappa$ -bump"" due to metals (C, O, Ne and Fe, see Rogers Iglesias 1992), as explained in Chabrier Baraffe (1997, see their 3.2 and their Fig."58 9)., 9).59 The exact temperature at which this occurs depends on the density. given the sensitivity of the opacities to density in this temperature range (see Fig.," The exact temperature at which this occurs depends on the density, given the sensitivity of the opacities to density in this temperature range (see Fig."60 2a of Rogers Iglesias 1992)., 2a of Rogers Iglesias 1992).61 The higher the density. the higher the temperature required for the radiative core to develop.," The higher the density, the higher the temperature required for the radiative core to develop."62 Since this range of temperatures is also characteristic of the temperature required for Li nuclear fusion. the central temperature and density at which the radiative core develops will determine the temperature at the bottom of the convective envelope and thus the level of Li depletion in the convective envelope.," Since this range of temperatures is also characteristic of the temperature required for Li nuclear fusion, the central temperature and density at which the radiative core develops will determine the temperature at the bottom of the convective envelope and thus the level of Li depletion in the convective envelope."63 For sequences producing | M. stars. Fig.," For sequences producing 1 $\msol$ stars, Fig."64 4 shows that aceretion history has a strong impact on (1) the age for the onset of a radiative core. (il) the nass of the convective envelope at ages « 30 Myr. and (iii) the Li abundance in the convective envelope.," \ref{li_mf1}65 shows that accretion history has a strong impact on (i) the age for the onset of a radiative core, (ii) the mass of the convective envelope at ages $\simle$ 30 Myr, and (iii) the Li abundance in the convective envelope."66 The higher central temperature reached by sequences with strong burst aceretion rates (see Fig. 2)), The higher central temperature reached by sequences with strong burst accretion rates (see Fig. \ref{mrtc_mf1}) )67 results in a radiative core that develops earlier in time., results in a radiative core that develops earlier in time.68 For the sequence starting from Mini = 10 Myyp. its more compact structure yields significantly higher temperatures at the bottom of the convective envelope. with a maximum of ~ 7 10° Κ. resulting in complete Li depletion in the convective envelope at ages < | Myr.," For the sequence starting from $\minit$ = 10 $\mjup$, its more compact structure yields significantly higher temperatures at the bottom of the convective envelope, with a maximum of $\sim$ 7 $^6$ K, resulting in complete Li depletion in the convective envelope at ages $<$ 1 Myr."69 In comparison. for the accreting sequence starting with Mi; = O.1 M... the maximum temperature reached at the bottom of the convective envelope is ~ 3.5 10° K. which is comparable to the value found in the non accreting sequence (see Fig. 4)).," In comparison, for the accreting sequence starting with $\minit$ = 0.1 $\msol$, the maximum temperature reached at the bottom of the convective envelope is $\sim$ 3.5 $^6$ K, which is comparable to the value found in the non accreting sequence (see Fig. \ref{li_mf1}) )."70 These differences in temperature are crucial for Li depletion. yielding very different final lithium abundances in the convective envelope. as illustrated in Fig. 4..," These differences in temperature are crucial for Li depletion, yielding very different final lithium abundances in the convective envelope, as illustrated in Fig. \ref{li_mf1}."71 The main results of this section can be summarized as follows., The main results of this section can be summarized as follows.72 Depending on the initial seed mass and the burst rates. differentepisodic accretion histories can produce | M.. star models with different sizes of the convective envelope at the same age. for ages x 30 Myr.," Depending on the initial seed mass and the burst rates, different episodic accretion histories can produce 1 $\msol$ star models with different sizes of the convective envelope at the same age, for ages $\simle$ 30 Myr."73 The models can have different surface Li abundancesstructure. 1.9. after ~ 30 Myr in the particular case portrayed in Fig. 4..," The models can have different surface Li abundances, i.e. after $\sim$ 30 Myr in the particular case portrayed in Fig. \ref{li_mf1}."74 Similar effects are found for models with episodic accretion producing partly convective objects in the mass range 0.35 -1 Me., Similar effects are found for models with episodic accretion producing partly convective objects in the mass range 0.35 -1 $\msol$.75 We have shown in the present work that an early protostar/BD accretion history based on episodes of short. intense bursts of accretion with typical accretion rates Mua=ο510¢Mayr!. as obtained in 2D hydrodynamical simulations of gravitationally unstable aceretion disks (Vorobyov Basu 2005). can affect the internal structure of low-mass objects even after several Myr. up to a few times 10 Myr.," We have shown in the present work that an early protostar/BD accretion history based on episodes of short, intense bursts of accretion with typical accretion rates $\mdotburst = 10^{-4} \, - \, 5 \, 10^{-4} \msolyr$, as obtained in 2D hydrodynamical simulations of gravitationally unstable accretion disks (Vorobyov Basu 2005), can affect the internal structure of low-mass objects even after several Myr, up to a few times 10 Myr."76 The main results obtained in the present analysis can be summarized as follows., The main results obtained in the present analysis can be summarized as follows.77find it to He in the range 1.25—1.3 from galaxies to massive clusters. narrowed down to 1.27—1.3 from poor (overall mass M— a few 107M.) to very rich (M— a few I0?M.) clusters.,"find it to lie in the range $1.25 - 1.3$ from galaxies to massive clusters, narrowed down to $1.27 - 1.3$ from poor (overall mass $M \sim$ a few $10^{14}\,M_{\odot}$) to very rich $M \sim$ a few $10^{15}\,M_{\odot}$ ) clusters."78" Two basic features of these ""o-profiles' entering the SM frame are briefly recalled in Appendix A. We stress that our treatment of the a-profiles (see Lap Cavaliere 2009a.b) agrees with recent numerical simulations (Zhao et al."," Two basic features of these $\alpha$ -profiles' entering the SM frame are briefly recalled in Appendix A. We stress that our treatment of the $\alpha$ -profiles (see Lapi Cavaliere 2009a,b) agrees with recent numerical simulations (Zhao et al."79 2003: Diemand et al., 2003; Diemand et al.80 2001: Wang White 2008) in picturing the DM halo formation as a development., 2007; Wang White 2008) in picturing the DM halo formation as a development.81" The epoch z, ranging from 0.5 to 2 marks the halo's transition from the stage of fast violent collapse scarred by major merging events. to one of progressively smoother and slower accretion building up the outskirts with little body perturbation."," The epoch $z_t$ ranging from $0.5$ to $2$ marks the halo's transition from the stage of fast violent collapse scarred by major merging events, to one of progressively smoother and slower accretion building up the outskirts with little body perturbation."82 During the latter. the DM profiles develop an increasing concentration c=R/r.». that actually measures the current outer extension out to the virial radius RX. relative to the region inner to 2 where p(r) is flatter than 77 (see CLFFO9).," During the latter, the DM profiles develop an increasing concentration $c \equiv R/r_{-2}$, that actually measures the current outer extension out to the virial radius $R$, relative to the region inner to $r_{-2}$ where $\rho(r)$ is flatter than $r^{-2}$ (see CLFF09)."83" Starting with values ο~3.5 at the transition. c increases to current values ezz3.5(1 zjJ. to attain values ez10 or more for the 10% fraction of rich clusters with a transition as early as zc1.5 (see Lapi Cavaliere 2009b): thus c measures the age z, of the cluster."," Starting with values $c\approx 3.5$ at the transition, $c$ increases to current values $c\approx 3.5\,84(1+z_t)$, to attain values $c\approx 10$ or more for the $10\%$ fraction of rich clusters with a transition as early as $z_t85\approx 1.5$ (see Lapi Cavaliere 2009b); thus $c$ measures the age $z_t$ of the cluster."86" Since R is typically of order 2 Mpe. values οz5—10 imply r_>~100—200 kpe: in the following. we will refer to ‘outskirts’ for the region r7-5. and to ""inner body for the region r=+1."," Since $R$ is typically of order $2$ Mpc, values $c\approx 5-10$ imply $r_{-2}\approx 100-200$ kpc; in the following, we will refer to `outskirts' for the region $r>r_{-2}$, and to `inner body' for the region $r\la r_{-2}$."87 Based on the SM. we will discuss in particular how the outer distribution of the ICP entropy relates to the DM concentration.," Based on the SM, we will discuss in particular how the outer distribution of the ICP entropy relates to the DM concentration."88 The SM describes the equilibrium of the ICP in the DM potential well. as we recall in 2: there we briefly recap the main features of the ICP description in the form of the SM as developed by CLFFO9. with the addition of Eq. (," The SM describes the equilibrium of the ICP in the DM potential well, as we recall in 2; there we briefly recap the main features of the ICP description in the form of the SM as developed by CLFF09, with the addition of Eq. ("898) that enables computing the total mass.,8) that enables computing the total mass.90 Then in 3 we describe our analysis procedure based on the SM., Then in 3 we describe our analysis procedure based on the SM.91 In 4 we apply such à procedure to the data for a set of six clusters with diverse profiles of X-ray brightness and temperature., In 4 we apply such a procedure to the data for a set of six clusters with diverse profiles of X-ray brightness and temperature.92 Finally. in 5 we discuss the specific physical information we extract from our analysis of the X-ray observations.," Finally, in 5 we discuss the specific physical information we extract from our analysis of the X-ray observations."93 In our treatment we adopt a standard flat cosmology with normalized matter density O4;=0.27. dark energy density Q420.73 and Hubble constant Ho=70 km s! Mpc!. except when comparing with data otherwise scaled.," In our treatment we adopt a standard flat cosmology with normalized matter density $\Omega_M = 0.27$, dark energy density $\Omega_{\Lambda} = 0.73$ and Hubble constant $H_0 =9470$ km $^{-1}$ $^{-1}$, except when comparing with data otherwise scaled."95 We concentrate first on DM halos close to self-gravitational equilibrium under smooth and slow accretion. past the fast collapse stage and after any residual violent mergers. às discussed in 5.," We concentrate first on DM halos close to self-gravitational equilibrium under smooth and slow accretion, past the fast collapse stage and after any residual violent mergers, as discussed in 5."96 Within these DM gravitational wells. the ICP approaches hydrostatic conditions after outgoing blastwaves and shocks driven by central mergers or AGN outbursts have subsided.," Within these DM gravitational wells, the ICP approaches hydrostatic conditions after outgoing blastwaves and shocks driven by central mergers or AGN outbursts have subsided."97 In such conditions. the ICP density is governed by the balance between the gravitational force and the gradient of the pressure p2nkpTpi; with the latter expressed as pxKj? in terms of ICP entropy./ the balance reads ∖∖⇁∣↴⊜∣⋪⊜∕∣≈↭∙↻≣⋋⋔⊜⋯⊖⋅≏⋯⋯⋂∣⊖∁∐∣⋅≏∐⋪∖∖⇁⊜≣∶↔⊺∣⇈⋅∣⊔∕≽⇪⋋⋔⊖ proton mass. and G ts the gravitational constant.," In such conditions, the ICP density is governed by the balance between the gravitational force and the gradient of the pressure $p=n\,k_BT/\mu$; with the latter expressed as $p98\propto k\, n^{5/3}$ in terms of ICP entropy, the balance reads where $\mu\approx 0.6$ is the mean molecular weight, $m_p$ is the proton mass, and $G$ is the gravitational constant."99" Once the radial entropy run K(r) is given as discussed below. the solution of this 1° order differential equation allowed us (CLFFO9) to write the profiles of the gas temperature and density in the form Here barred variables are normalized to their boundary value at r=ΑΔ: the squared circular velocity V)-M(Gnjr is taken from the o-profiles with their weak dependence on à. and is expressed in Appendix A in terms of hypergeometric functions: finally. we define b,=pw)kpT(R)."," Once the radial entropy run $k(r)$ is given as discussed below, the solution of this $1^{\rm st}$ order differential equation allowed us (CLFF09) to write the profiles of the gas temperature and density in the form Here barred variables are normalized to their boundary value at $r = R$; the squared circular velocity $\bar{v}^2_c(\bar{r}) =100\bar{M} (< \bar{r})/\bar{r}$ is taken from the $\alpha$ -profiles with their weak dependence on $\alpha$, and is expressed in Appendix A in terms of hypergeometric functions; finally, we define $b_R \equiv \mu m_p101v^2_c(R)/k_BT(R)$."102 The latter incorporates the boundary condition. required for solving Eq. (, The latter incorporates the boundary condition required for solving Eq. (1031).,1).104 It is physically convenient to fix such a reference value at the virial radius 7R rather than at the center. where Vr) vanishes steeply (see Lapi Cavaliere 20093). while the ICP is often affected by violent stochastic events such as mergers and AGN outbursts.," It is physically convenient to fix such a reference value at the virial radius $r\approx R$ rather than at the center, where $v_c^2(r)$ vanishes steeply (see Lapi Cavaliere 2009a) while the ICP is often affected by violent stochastic events such as mergers and AGN outbursts."105 Moreover. at r~R closely universal accretion of external InterGalactic Medium (IGM) prevails for both Cool Core (CC) and Non Cool Core (NCC) clusters: this holds to lowest order. although the related conversion of infall kinetic into thermal energy may differ somewhat related to cluster age and preheating conditions. as discussed in 3 and 5.," Moreover, at $r\sim R$ closely universal accretion of external InterGalactic Medium (IGM) prevails for both Cool Core (CC) and Non Cool Core (NCC) clusters; this holds to lowest order, although the related conversion of infall kinetic into thermal energy may differ somewhat related to cluster age and preheating conditions, as discussed in 3 and 5."106 In fact. in the absence of substantial preheating in the IGM larger than some 1/2 keV per particle (see Lapi et al.," In fact, in the absence of substantial preheating in the IGM larger than some $1/2$ keV per particle (see Lapi et al."107 2005: MeCarthy et al., 2005; McCarthy et al.108 2008). the energy conversion takes place in accretion shocks that linger at the virial radius (see Tozzi Norman 2001: Lapi et al.," 2008), the energy conversion takes place in accretion shocks that linger at the virial radius (see Tozzi Norman 2001; Lapi et al."109 2005; Voit 2005). to imply that takes on values around 2.5.," 2005; Voit 2005), to imply that takes on values around $2.5$."110" These obtain from maximal conversion at a strong shock of the gravitational infall energy. that yields ApTp=ium,Ay: this quantity. involves the specific kinetic energy 2,Avr gained by the IGM that starts fromthe turning point its free fall toward R across the potential drop Ad,=VOR (see Lapi et al."," These obtain from maximal conversion at a strong shock of the gravitational infall energy, that yields $k_B T_R = \frac{2}{3}\mu m_p \,\Delta111\Phi_{R}$; this quantity involves the specific kinetic energy $2\, v^2_R\Delta \phi_{R}$ gained by the IGM that starts fromthe turning point its free fall toward $R$ across the potential drop $\Delta \Phi_{R}=v_R^2\,\Delta\phi_R$ (see Lapi et al."112 2005)., 2005).113 The latter is provided by the DM oa-profiles. and is conveniently normalized to the related v— vz(R): whence Eq. (," The latter is provided by the DM $\alpha$ -profiles, and is conveniently normalized to the related $v_R^2\equiv114v_c^2(R)$ ; whence Eq. ("1153) follows.,3) follows.116 We stress that &o7g lowers when Ao and the proportional infall kinetic energy are smaller owing to a large concentration c; e.g. à value οz10 holding for a cluster with an early transition (see 1). in place of cz4 holding for clusters with a recent transition. implies Ao to decrease from 0.57 to 0.47 (see CLFFO9).," We stress that $k_B T_R$ lowers when $\Delta\phi$ and the proportional infall kinetic energy are smaller owing to a large concentration $c$; e.g., a value $c\approx 10$ holding for a cluster with an early transition (see 1), in place of $c\approx 4$ holding for clusters with a recent transition, implies $\Delta \phi$ to decrease from $0.57$ to $0.47$ (see CLFF09)."117 On the other hand. &p»7y« also lowers when a high preheating level of the IGM the shock and impairs the conversion efficiency as discussed by Lapi et al. (," On the other hand, $k_B T_R$ also lowers when a high preheating level of the IGM the shock and impairs the conversion efficiency as discussed by Lapi et al. ("1182005) and Voit (2005).,2005) and Voit (2005).119 To obtain the [CP temperature and density profiles the full radial run &(r) of the entropy is required., To obtain the ICP temperature and density profiles the full radial run $k(r)$ of the entropy is required.120 [ts physical modeling is based upon the notions that entropy is eroded by radiative cooling on the timescale f£.—OS(keT/SkeV)(1/107.emmy! Gyr (Sarazin 1988). whilst it is enhanced by sfocks:: at rzR it is deposited by standing aceretion shocks. and in the central region is injected by outbound blastwaves terminating in a shock. as are driven by supersonic outflows from AGN outbursts or head- mergers.," Its physical modeling is based upon the notions that entropy is eroded by radiative cooling on the timescale $t_c \approx 65\, (k_B T/5~121\mathrm{keV})^{1/2}\, (n/10^{-3}~\mathrm{cm}^{-3})^{-1}$ Gyr (Sarazin 1988), whilst it is enhanced by ; at $r\approx R$ it is deposited by standing accretion shocks, and in the central region is injected by outbound blastwaves terminating in a shock, as are driven by supersonic outflows from AGN outbursts or head-on mergers."122 In the vicinity of krzR the powerlaw approximation kX7 always applies (see CLFFO9) with boundary value varying around 1.1.," In the vicinity of $r \approx R$ the powerlaw approximation $k123\propto r^a$ always applies (see CLFF09) with boundary value varying around $1.1$ ."124 Note that after Eq. (, Note that after Eq. (1253) theupper bound de=45/19~2.4 obtains for large Ao as then bp tends to,3) theupper bound $a_R = 45/19 \approx 2.4$ obtains for large $\Delta \phi$ as then $b_R$ tends to126The bottom-up growth of structure is a hallmark of hierarchical cosmological models such as the A cold. clark matter CXCDM) model.,The bottom-up growth of structure is a hallmark of hierarchical cosmological models such as the $\Lambda$ cold dark matter $\Lambda$ CDM) model.127 ln these universes. dark matter haloes of lower mass are expected to form earlier. on average. than massive haloes.," In these universes, dark matter haloes of lower mass are expected to form earlier, on average, than massive haloes."128 la observations. theoretical studies. ancl semi-analytical modelling of galaxy. formation (see ? for a review). the mass of a halo is the key variable upon which many properties of galaxies and their host. haloes depend. e.g. formation redshift. galaxy occupation number. colour. morphology. star formation rate. and stellar feedback processes.," In observations, theoretical studies, and semi-analytical modelling of galaxy formation (see \citealt{BaughReview} for a review), the mass of a halo is the key variable upon which many properties of galaxies and their host haloes depend, e.g., formation redshift, galaxy occupation number, colour, morphology, star formation rate, and stellar feedback processes."129 Recent work has shown that in addition to the mass. various properties of halo formation and evolution. also depend on the environment within whieh the haloes reside.," Recent work has shown that in addition to the mass, various properties of halo formation and evolution also depend on the environment within which the haloes reside."130 For instance. at a [fixed halo mass. older haloes have been found to be more clustered than vounger haloes. ancl the correlation between clustering strength and formation redshift is stronger for lower mass haloes (???????) 1n ," For instance, at a fixed halo mass, older haloes have been found to be more clustered than younger haloes, and the correlation between clustering strength and formation redshift is stronger for lower mass haloes \citep{Gottlober02, ShethTormen04, Gao05, Harker06, GaoWhite07, JingSutoMo07, Wechsler06, WangMoJing07, Hahn08, MFM09}."131many of these studies. halo environment is characterised via the halo bias parameter. which is determined. from the relative clustering strengths of haloes to the underlving clark matter clüistribution.," In many of these studies, halo environment is characterised via the halo bias parameter, which is determined from the relative clustering strengths of haloes to the underlying dark matter distribution."132 In Part L of this series (?:: henceforth EMO9). we have instead chosen to use local overdensities as à more cirect measure of halo environment.," In Part I of this series \citealt{FM09}; henceforth FM09), we have instead chosen to use local overdensities as a more direct measure of halo environment."133 The focus of the study was to quantify the environmental dependence. of the merger rate of haloes., The focus of the study was to quantify the environmental dependence of the merger rate of haloes.134 We compared a number of local overdensity variables (both including and excluding the halo mass itself). some of which hac been used in earlier studies. (?*77) ," We compared a number of local overdensity variables (both including and excluding the halo mass itself), some of which had been used in earlier studies \citep{LemsonKauffmann, Harker06, WangMoJing07, Maulbetsch07, Hahn08}."135The kev finding in EMOO was that halo-halo mergers occur more frequently in denser regions than in voids. ane that his environmental dependence is similar regardless of the merger mass ratio (o... minor vs major) or the descendant ido mass (galaxv- vs cluster-sized): we found. mergers to occur about 2.5 times more frequently in the densest regions han in the emptiest regions.," The key finding in FM09 was that halo-halo mergers occur more frequently in denser regions than in voids, and that this environmental dependence is similar regardless of the merger mass ratio (e.g., minor vs major) or the descendant halo mass (galaxy- vs cluster-sized): we found mergers to occur about 2.5 times more frequently in the densest regions than in the emptiest regions."136 We provided. an analytical ormula as a function of local density to approximate this environmental trend., We provided an analytical formula as a function of local density to approximate this environmental trend.137 This expression can be used with the it lor the global mean rate of 2. (henceforth. EFMOS) to xediet halo merger rates as a function of descendant mass. »ogenitor mass ratio. redshift. and environment over a wide range of parameter space.," This expression can be used with the fit for the global mean rate of \cite{FM08} (henceforth FM08) to predict halo merger rates as a function of descendant mass, progenitor mass ratio, redshift, and environment over a wide range of parameter space."138" Using assumption 3 that My—iMac. and substituting (vpical values. this condition reads The accretion rate M, is the accretion rate onto the SMDII. and the inflow rate of the surrounding gas is assumed to form stars in the bulee (if it is not expelled bv the jets)."," Using assumption 3 that $\dot M_f=\eta \dot M_{\rm acc}$, and substituting typical values, this condition reads The accretion rate $\dot M_{\rm acc}$ is the accretion rate onto the SMBH, and the inflow rate of the surrounding gas is assumed to form stars in the bulge (if it is not expelled by the jets)."139 If the inflow rate is above the value 5givenin equation (8)). the deposition of enerevο. bv the jets is efficient enough to expel the mass back (ο laree distances and heat it (Soker 2008).," If the inflow rate is above the value givenin equation \ref{eq:mscluster}) ), the deposition of energy by the jets is efficient enough to expel the mass back to large distances and heat it (Soker 2008)."140 Following assumption 4 above. I take (he SAIBIT to bulge mass ratio to be equal to Mà44.," Following assumption 4 above, I take the SMBH to bulge mass ratio to be equal to $\dot M_{\rm acc}/{\dot M_s}$."141 Equation (3)) vields The last equation has the form of equation (2)). and by using twpical properties of AGN jets and bulges ο ey: 0). the numerical value of equation (2)) is reproduced with X=1.," Equation \ref{eq:mscluster}) ) yields The last equation has the form of equation \ref{eq:corr2}) ), and by using typical properties of AGN jets and bulges $\eta$; $v_f$; $\sigma$ ), the numerical value of equation \ref{eq:corr2}) ) is reproduced with $\chi=1$."142 For a more detailed comparison to observations. füture work is needed to incorporate the proposed feedback mechanism into numerical simulations of galaxy formation and evolution. e.g.. such as those described by Di Matteo et al. (," For a more detailed comparison to observations, future work is needed to incorporate the proposed feedback mechanism into numerical simulations of galaxy formation and evolution, e.g., such as those described by Di Matteo et al. ("1432005).,2005).144 ] built a simple feedback mechanism where jets launched by the central supermassive black hole (SAIBID interact with the surrounding inflowing gas in a voung galaxy., I built a simple feedback mechanism where jets launched by the central supermassive black hole (SMBH) interact with the surrounding inflowing gas in a young galaxy.145 Dv using several verv plausible assumptions. most of which are well supported by observations. I derived a relation between the mass (hat is converted to stars and (he mass accreted by the central SMDBII (eqs.," By using several very plausible assumptions, most of which are well supported by observations, I derived a relation between the mass that is converted to stars and the mass accreted by the central SMBH (eqs."146 8. and 9))., \ref{eq:mscluster} and \ref{eq:mBH2}) ).147 The basic process to suppress star formation requires (hat the jets do not penetrate the surrounding gas: if thev do. their energy is deposited at large distances. rather than in the gas that potentially can form stars in the bulge.," The basic process to suppress star formation requires that the jets do not penetrate the surrounding gas; if they do, their energy is deposited at large distances, rather than in the gas that potentially can form stars in the bulge."148 For that. the jets should encounter new material before they escape.," For that, the jets should encounter new material before they escape."149" Namely. the (vpical Gime for the jets Lo cross ilsown width at radius r; should be shorter (than the penetration time at radius ος, uS~NS fy "," Namely, the typical time for the jets to cross itsown width at radius $r_s$ should be shorter than the penetration time at radius $r_s$ , $\tau_s \la t_p$ ."150This leads to equation (3))., This leads to equation \ref{eq:mscluster}) ).151C. N. O. Ne. Ar. S. Cl and Fe. rather less than the 7100 lines seen in our spectra of knots J] and J3 in Abell 30 (Wesson et al.,"C, N, O, Ne, Ar, S, Cl and Fe, rather less than the $>$ 100 lines seen in our spectra of knots J1 and J3 in Abell 30 (Wesson et al."152 2003)., 2003).153 None the less. these spectra are the deepest that have been obtained of this object. and are of high enough quality to carry out a thorough analysis of the properties of the knot.," None the less, these spectra are the deepest that have been obtained of this object, and are of high enough quality to carry out a thorough analysis of the properties of the knot."154 We find typical errors on our flux measurements of for lines with FCA)-:0.2. where F(H.3)== 11. forlines with 0.2 FCA) 3.0. and <5% for lines with FCA)-3.0.," We find typical errors on our flux measurements of for lines with $\lambda$ $<$ 0.2, where $\beta$ 1, forlines with $\le$ $\lambda$ $\le$ 3.0, and $<$ for lines with $\lambda$ $>$ 3.0."155 We detect only the brightest lines from the outer nebula of Abell 58., We detect only the brightest lines from the outer nebula of Abell 58.156 We derive a logarithmic extinction at H. e(11.7) of 1.04 from the observed Hao/H:? ratio.," We derive a logarithmic extinction at $\beta$, $c({\rm H}\beta)$ of 1.04 from the observed $\alpha$ $\beta$ ratio."157 This is similar to the value obtained using the reddening maps of Schlegel et al. (, This is similar to the value obtained using the reddening maps of Schlegel et al. (1581998). which for 558's galactic longitude and latitude of 37.60° and -5.16° respectively give E(B-V)=0.544.,"1998), which for 58's galactic longitude and latitude of $^\circ$ and $^\circ$ respectively give E(B-V)=0.544."159" This is equivalent to À,.—1.69. and c(1.3)21.16."," This is equivalent to $_v$ =1.69, and $c({\rm H}\beta)$ =1.16."160 No temperature diagnosties are available. and both the [O 11] and [S 11] density diagnostic line ratios fall below the low density limit.," No temperature diagnostics are available, and both the [O ] and [S ] density diagnostic line ratios fall below the low density limit."161 We calculate ionic abundances for three temperatures — 7.5. 10 and ISkKKK — assuming that the density is 7. the value found by GM96.," We calculate ionic abundances for three temperatures – 7.5, 10 and kK – assuming that the density is $^{-3}$, the value found by GM96."162 Total abundances are calculated using the ionisation correction scheme of Kingsburgh Barlow (199-4)., Total abundances are calculated using the ionisation correction scheme of Kingsburgh Barlow (1994).163 The results are shown in Table 2.., The results are shown in Table \ref{outerabunds}. .164 The spectrum of the outer nebula is of low excitation. with I([O 11] 37274-3729) — IO 1] 49594-5007). so a T. between 7.5 and IOKK seems most likely.," The spectrum of the outer nebula is of low excitation, with I([O ] 3727+3729) $\sim$ I([O ] 4959+5007), so a $_{\rm e}$ between 7.5 and 10kK seems most likely."165 We derive a Ne/O ratio between 0.49 and 0.74. depending on the temperature adopted.," We derive a Ne/O ratio between 0.49 and 0.74, depending on the temperature adopted."166 This is rather higher than the average value for PNe of about 0.25., This is rather higher than the average value for PNe of about 0.25.167 At 7.5kK. the derived oxygen abundance (ton a logarithmic scale where NCH)212.0) of 8.74 is close to the average value of 8.68 found for galactic PNe by Kingsburgh Barlow (1994). while the neon abundance of 8.61 is more than a factor of three higher than their average.," At 7.5kK, the derived oxygen abundance (on a logarithmic scale where N(H)=12.0) of 8.74 is close to the average value of 8.68 found for galactic PNe by Kingsburgh Barlow (1994), while the neon abundance of 8.61 is more than a factor of three higher than their average."168 At LOKKK. on the other hand. the derived neon abundance is about 30 per cent lower than the averagefound by Kingsburgh Barlow (1994). while the oxygen abundance is a factor of three lower.," At kK, on the other hand, the derived neon abundance is about 30 per cent lower than the averagefound by Kingsburgh Barlow (1994), while the oxygen abundance is a factor of three lower."169 Assuming T.=7.5 KKK gives N/O = 0.38. while T.ZIOKKK gives N/O = 0.68.," Assuming $_{\rm e}$ kK gives N/O = 0.38, while $_{\rm e}$ kK gives N/O = 0.68."170 Adopting a temperature of KKK gives N/O = 1.20. which would classify 558 as a Type I planetary nebula as detined by Kingsburgh Barlow (1994). with N/O>-0.8.," Adopting a temperature of kK gives N/O = 1.20, which would classify 58 as a Type I planetary nebula as defined by Kingsburgh Barlow (1994), with $>$ 0.8."171 The criteria of Peimbert Torres-Peimbert (1983) require N/O>0.5 and He/H0.125. and so the outer nebula of 558 would be classitied as a Peimbert Type I nebula if its temperature is higher than about KKK. The knot of 558 is expanding at about 4. and is moving at-LOOKkmss ο relative to the systemic velocity (Pollacco et al.," The criteria of Peimbert Torres-Peimbert (1983) require $>$ 0.5 and $>$ 0.125, and so the outer nebula of 58 would be classified as a Peimbert Type I nebula if its temperature is higher than about kK. The knot of 58 is expanding at about $^{-1}$, and is moving at $^{-1}$ relative to the systemic velocity (Pollacco et al."172 1992)., 1992).173 Estimates of the distance to the system have ranged from ~3 kkpe (Cahn Kaler 1971)to kKpe (Pollacco et al., Estimates of the distance to the system have ranged from $\sim$ kpc (Cahn Kaler 1971) to kpc (Pollacco et al.174 1992)., 1992).175 The knot is ionised by a star which is completely obscured from view. although stellar emission lines are visible in scattered light.," The knot is ionised by a star which is completely obscured from view, although stellar emission lines are visible in scattered light."176 Clayton et al. (, Clayton et al. (1772006) determined a stellar temperature of kKKK and a luminosity of IO! L.; . assuming a distance of kkpe.,"2006) determined a stellar temperature of kK and a luminosity of $^4$ $_{\sun}$ , assuming a distance of kpc."178 We estimate the mass of the knot in two ways., We estimate the mass of the knot in two ways.179" First. we use the following relation: where D is the distance to the knot. He A5876) is the dereddened flux of the He line at5876À. sry, is the mean ionic mass per jon. n, is the electron density. o,rprizszii is the effective recombination coefficient of the AS5876 line and Ένατη is the energy of each photon at5876À."," First, we use the following relation: where D is the distance to the knot, I(He $\lambda$ 5876) is the dereddened flux of the He line at, $\mu_{He}$ is the mean ionic mass per $^+$ ion, $n_{\rm e}$ is the electron density, $\alpha_{eff(5876)}$ is the effective recombination coefficient of the $\lambda$ 5876 line and $_{\lambda5876}$ is the energy of each photon at."180 Adopting a distance of kkpe from Clayton et al. (, Adopting a distance of kpc from Clayton et al. (1812006). an electron density of 2100cem7 (Section 4.3)) and a mean ionic mass per ion of 9.2. based on ionic abundances derived from ORLs (Section 4.6). we derive a mass of M.; for the Knot.,"2006), an electron density of $^{-3}$ (Section \ref{diagnostics}) ) and a mean ionic mass per $^+$ ion of 9.2, based on ionic abundances derived from ORLs (Section \ref{abundances}) ), we derive a mass of $\times$ $^{-5}$ $_{\sun}$ for the knot."182" Second. we derive a mass based on the angular size of the knot: where O is the angular radius of the knot. D its distance. jie the mean ionic mass per electron. ¢ the volume filling factor. and 5», the electron density."," Second, we derive a mass based on the angular size of the knot: where $\Theta$ is the angular radius of the knot, D its distance, $\mu_e$ the mean ionic mass per electron, $\epsilon$ the volume filling factor, and $n_{\rm e}$ the electron density."183 We measure the angular radius of the knot from Hubble Space Telescope images which we downloaded from the STSci archive., We measure the angular radius of the knot from Hubble Space Telescope images which we downloaded from the STSci archive.184 The images were taken in 2001 as part of GO program 9092. in narrow-band filters isolating [O ΠΠ] and [N 11] emission.," The images were taken in 2001 as part of GO program 9092, in narrow-band filters isolating [O ] and [N ] emission."185 We derive an angular radius of 0.38 aresec., We derive an angular radius of 0.38 arcsec.186 Taking D to be KKpe. and yy. to be 6.16. based on abundances derived from ORLs (Section 4.63). we derive a mass of 322.10. te M. .," Taking D to be kpc, and $\mu_e$ to be 6.16, based on abundances derived from ORLs (Section \ref{abundances}) ), we derive a mass of $\times$ $^{-4}$$\epsilon\,$ $_{\sun}$ ."187" Equating these two masses we find that« — 0.16. if n, =2100cem"," Equating these two masses we find that $\epsilon$ = 0.16, if $n_{\rm e}$ =$^{-3}$ ."188 GM96 found a value of e(LI:2) of 0.29 forthe outer nebula of 558. but with a large uncertainty due to the weakness of the hydrogen line intensities.," GM96 found a value of $c({\rm H}\beta)$ of 0.29 forthe outer nebula of 58, but with a large uncertainty due to the weakness of the hydrogen line intensities."189 They suggested that the extinction in the central, They suggested that the extinction in the central190variable light curve shape (see Pavlenko et al.,variable light curve shape (see Pavlenko et al.191 2003 and the scatter of the phase diagrams in reflc))., 2003 and the scatter of the phase diagrams in \\ref{lc}) ).192 A tentative spectral type has been assigned to the companion based on medium-resolution spectra with the AAT., A tentative spectral type has been assigned to the companion based on medium-resolution spectra with the AAT.193 Adopting E(B—V)=0006. d=1.6 kpe and assuming the mean corrected R magnitude (~1673) of the nova is close to the quiescent brightness. the progenitor’s absolute magnitude was about Mj=39. which leaves unchanged our conclusion on it (Kiss Thomson 2000).," Adopting $E(B-V)=0\fm6$, $d=1.6$ kpc and assuming the mean corrected $R$ magnitude $\sim16\fm3$ ) of the nova is close to the quiescent brightness, the progenitor's absolute magnitude was about $M_{\rm R}=3\fm9$, which leaves unchanged our conclusion on it (Kiss Thomson 2000)."194 However. further high-resolution imaging. preferably made under sub-aresecond seeing. is needed in order to allow accurate photometric corrections and light curve modelling.," However, further high-resolution imaging, preferably made under sub-arcsecond seeing, is needed in order to allow accurate photometric corrections and light curve modelling."195"actual Dux from the 7blank fields"" is higher than that fron some patches of the Galactic place free of strong sources.",actual flux from the “blank fields” is higher than that from some patches of the Galactic place free of strong sources.196 The morphology of the slice in the 508514 keV containing the 511 keV line (Fig. nmi," The morphology of the slice in the 508–514 keV, containing the 511 keV line (Fig. \ref{fig:mslice},"197üdcdle panel). is markedly dillerent.," middle panel), is markedly different."198. With the bins along the Galactic plane. the only. prominent feature is the peak at the Galactic Centre. with the extent along / roughly similar to the extent of the peak near Cvg N-1 at low energies.," With the bins along the Galactic plane, the only prominent feature is the peak at the Galactic Centre, with the extent along $l$ roughly similar to the extent of the peak near Cyg X-1 at low energies."199 The peak Lux is 190phots|em7., The peak flux is $\sim 10^{-3}~\pscm2$.200 As already emphasized. this value is not a precise measure of the true 511 keV flux from the Galactic Centre region. but it should be accurate to a factor of better than 2 if the source size does not exceed. several degrees.," As already emphasized, this value is not a precise measure of the true 511 keV flux from the Galactic Centre region, but it should be accurate to a factor of better than 2 if the source size does not exceed several degrees."201 No strong asvmmetry of the positive ancl negative ongitude wings of the central peak is apparent in the ight bucket. profiles (sec. however. Weidenspointneret.al.2008.whoreportanexcess[luxfornegativelongitudes and Bouehetctal.2008.. who do not find strong evidence Or asvmmetrv).," No strong asymmetry of the positive and negative longitude wings of the central peak is apparent in the light bucket profiles (see, however, \citealt[][who report an excess flux for negative202 longitudes]{2008Natur.451..159W} and \citealt{2008ApJ...679.1315B}, who do not find strong evidence for asymmetry)."203 An excess al neealive longituces lo is visible in the light bucket. profile in Teegarden.et(2005). (their Fig., An excess at negative longitudes $l\sim$ is visible in the light bucket profile in \citet{2005ApJ...621..296T} (their Fig.204 2)., 2).205 Renormalization of their. Duxes o the units used in Fig., Renormalization of their fluxes to the units used in Fig.206 2 suggests an excess at the level 1510phots+em;," \ref{fig:mslice} suggests an excess at the level $\sim207 1.5~10^{-4}~\pscm2$."208 In our analysis the Lux at bo is essentially consistent with zero., In our analysis the flux at $l\sim$ is essentially consistent with zero.209 If anything. he profiles shown inFig.," If anything, the profiles shown inFig."210 2. suggest an excess in the 508r514 keV Ilux at. positive longituces /.., \ref{fig:mslice} suggest an excess in the 508–514 keV flux at positive longitudes $l\sim$.211 However. our experiments with “alternative” background models mentioned in refsec:ibackground showed that a weak spurious asvmmetrv (corresponding to a flux excessfclelicit of ~10photslom at |~20 )) can appear at either negative or positive longituces.," However, our experiments with “alternative” background models mentioned in \\ref{sec:background} showed that a weak spurious asymmetry (corresponding to a flux excess/deficit of $\sim 10^{-4}~\pscm2$ at $|l|\sim$ ) can appear at either negative or positive longitudes."212 We therefore conclude that. with the present. knowledge of the SPL background. light bucket. profiles do not. provide compelling evidence for asvmmetry in the 508514 keV lus along the Galactic plane.," We therefore conclude that, with the present knowledge of the SPI background, light bucket profiles do not provide compelling evidence for asymmetry in the 508–514 keV flux along the Galactic plane."213 Finally in the bottom panel of Fig. 2..," Finally in the bottom panel of Fig. \ref{fig:mslice},"214" the slice of the Galactic plane in the 1804.1813 keV. band. containing the 1.8 MeV line of ""PAL is shown."," the slice of the Galactic plane in the 1804–1813 keV band, containing the 1.8 MeV line of $^{26}$ Al, is shown."215 Unlike the 511 keV emission. there is a broad. peak (much broader than the SPL response to à point source). centred at the GC.," Unlike the 511 keV emission, there is a broad peak (much broader than the SPI response to a point source), centred at the GC."216 This distribution is qualitatively consistent with the Comptel (Plüschkeοἱal.2001) and earlier results (Waneetal.2009)., This distribution is qualitatively consistent with the Comptel \citep{2001ESASP.459...55P} and earlier results \citep{2009A&A...496..713W}.217 A similar approach of light bucket mapping can be applied to the special scans of the Galactic plane along 6 performed in 2007. 2008 ancl 2009.," A similar approach of light bucket mapping can be applied to the special scans of the Galactic plane along $b$ performed in 2007, 2008 and 2009."218 The positions of the scans in / are denoted with the grav. vertical lines in Fig. 2.., The positions of the scans in $l$ are denoted with the gray vertical lines in Fig. \ref{fig:mslice}.219 Phree out of four scans were mace as a sequence of pointings starting and ending 30° away from the plane on either side of it., Three out of four scans were made as a sequence of pointings starting and ending $\sim 30$ away from the plane on either side of it.220 The pointing direction was changing in ~2” steps., The pointing direction was changing in $\sim$ steps.221 The remaining ssean at f=0 had been done earlier as à set of pointings directly at the plane and —25 above and below the plane., The remaining “scan” at $l=0$ had been done earlier as a set of pointings directly at the plane and $\sim$ above and below the plane.222 The advantage of making “fast” (taking 10.15 hours) scans over a selected region of the sky is that the quality. of the background. modeling can be verifiec/improvecl if typical variations of the detector background occur on longer time scales., The advantage of making “fast” (taking 10–15 hours) scans over a selected region of the sky is that the quality of the background modeling can be verified/improved if typical variations of the detector background occur on longer time scales.223 The scans in the 508514 keV band in the direction perpendicular to the Galactic plane at /—Οι22. are shown in Fig. 3..," The scans in the 508–514 keV band in the direction perpendicular to the Galactic plane at $l\sim$, are shown in Fig. \ref{fig:vslice}."224 Phe signal is clearly visible at /— and is close to zero for /~, The signal is clearly visible at $l\sim$ and is close to zero for $l\sim$.225 ‘To guide the eve. we have plotted with the black curves the expected. light bucket. profiles for the case where the spatial distribution of the Dux in the 508514 keV. band follows the “Bulge” moclel described later in ," To guide the eye, we have plotted with the black curves the expected light bucket profiles for the case where the spatial distribution of the flux in the 508–514 keV band follows the “Bulge” model described later in \\ref{sec:templates}."226The expected profiles were cruclely estimated assuming that the axis of the instrument was moving along b at fixed /. while keeping the position angle fixed.," The expected profiles were crudely estimated assuming that the axis of the instrument was moving along $b$ at fixed $l$, while keeping the position angle fixed."227 This causes visible asymmetry. in the peak of the LUN profile. relecting the particular orientation of the SPL coding mask relative to the source.," This causes visible asymmetry in the peak of the $l\sim$ profile, reflecting the particular orientation of the SPI coding mask relative to the source."228 In reality. the observed. profiles. are combinations of observations with varving position angle ancl varving / and hence they may dilfer from the simulated. profiles in subtle details.," In reality, the observed profiles are combinations of observations with varying position angle and varying $l$ and hence they may differ from the simulated profiles in subtle details."229 Broacly. the light bucket imaging shows the overall consistency. with a bulge-dominated distribution of the 508514 keV tux and indicates that the disk emission is weak at /," Broadly, the light bucket imaging shows the overall consistency with a bulge-dominated distribution of the 508–514 keV flux and indicates that the disk emission is weak at $l\sim$."230 We now use a simple function. two-dimensional Gaussian to model the spatial distribution of the annihilation line emission near the Galactic Centre., We now use a simple function – two-dimensional Gaussian – to model the spatial distribution of the annihilation line emission near the Galactic Centre.231 Phe data collected when the pointing direction was within of the GC were used., The data collected when the pointing direction was within of the GC were used.232 Two mocdoels were considered: where Gf.) is the surface brightness distribution as à function of Galactic longitude and latitude. £ is the Dux and Wy Ms ave the full widths at half maximum along / and b. respectively.," Two models were considered: where $G(l,b)$ is the surface brightness distribution as a function of Galactic longitude and latitude, $F$ is the flux and $W_l$, $W_b$ are the full widths at half maximum along $l$ and $b$, respectively."233 Lhe second model has an additional component that is aimed to aecount. for emission. elongated: over the Galactic plane., The second model has an additional component that is aimed to account for emission elongated over the Galactic plane.234 Since only the data within of the GC were used corresponds to the deviation of the SPL axis from the GC). the “infinite” extent of this component over f means that the surface brightness of this component does not decrease much at a distance of ~45° from the GC.," Since only the data within of the GC were used corresponds to the deviation of the SPI axis from the GC), the “infinite” extent of this component over $l$ means that the surface brightness of this component does not decrease much at a distance of $\sim$ from the GC."235 To verily the sensitivity of the results to a particular funetional form. we also used an exponential law insteacl of the Gaussian: For a given pair of M; and VW». the model is convolvedwith the simulated SPL response (Sturneretal.2003). and compared with the count rate in the 508514 keV band inindividual detectors during individual observations.," To verify the sensitivity of the results to a particular functional form, we also used an exponential law instead of the Gaussian: For a given pair of $W_l$ and $W_b$ , the model is convolvedwith the simulated SPI response \citep{2003A&A...411L..81S} and compared with the count rate in the 508–514 keV band inindividual detectors during individual observations."236 The best-fitting value of £4 for the one-component model or the pair of values £4 and £5 for the two-component model is, The best-fitting value of $F_1$ for the one-component model or the pair of values $F_1$ and $F_2$ for the two-component model is237was interpreted in terms of an eclipse of the source by the companion star (Bozzoetal..2008.2009;Jain2009).,"was interpreted in terms of an eclipse of the source by the companion star \citep{bozzo08,bozzo09,jain09}."238. This observation revealed that in at least one case the X-ray variability of a SEXT. was due to the obscuration by the companion star., This observation revealed that in at least one case the X-ray variability of a SFXT was due to the obscuration by the companion star.239 In the case of JJ18483-0311. hhelped identify pulsations in the quiescent X-ray flux. of this source (Giuntaetal..2009).. and thus provided strong support for the idea that SFXTs also acerete matter during their quiescent states (seee.g..Sidolietal..2007:Bozzo 2008).," In the case of J18483-0311, helped identify pulsations in the quiescent X-ray flux of this source \citep{giunta09}, and thus provided strong support for the idea that SFXTs also accrete matter during their quiescent states \citep[see e.g.,][]{sidoli07,bozzo08}."240. To study the low level emission of SEXT sources. we present in this paper quiescent state oobservations of the prototypical SFXTs XTEJJ1739-302 and JJ08408-4503.," To study the low level emission of SFXT sources, we present in this paper quiescent state observations of the prototypical SFXTs J1739-302 and J08408-4503."241 In Sect. 2..," In Sect. \ref{sec:sources},"242 we summarize previous observations of these sources. and in Sect.," we summarize previous observations of these sources, and in Sect."243 3 and 4+ we present our data analysis and results., \ref{sec:data} and \ref{sec:results} we present our data analysis and results.244 In particular. we find that the quiescent spectra of these sources contain a soft component below —2 keV. We diseuss some possible interpretations of this component in Sect. 5.," In particular, we find that the quiescent spectra of these sources contain a soft component below $\sim$ 2 keV. We discuss some possible interpretations of this component in Sect. \ref{sec:discussion}."245 A comparison is also carried out between the quiescent and outburst spectral properties of παπά OOur conclusions are summarized in Sect. 6.., A comparison is also carried out between the quiescent and outburst spectral properties of and Our conclusions are summarized in Sect. \ref{sec:conclusion}.246 iis à SEXT prototype. and was discovered with dduring a bright outburst in. 1997. (Smithetal...1998).," is a SFXT prototype, and was discovered with during a bright outburst in 1997 \citep{smith98}."247. The identification of its supergiant companion led to the determination of the source distance at 2.7 kpe (Rahouietal.. 2008)., The identification of its supergiant companion led to the determination of the source distance at 2.7 kpc \citep{rahoui08}.248. Several outbursts from this source were detected later with RXTE (Smithetal.. 2006).. and citeplutovinovOS.sguera05.sguera06.blay08..," Several outbursts from this source were detected later with RXTE \citep{smith06}, , and \\citep{lutovinov05,sguera05,sguera06,blay08}."249 wwas observed in outburst with //BAT on three occasions. on 2008 April 8 (Sidolietal..2009).. on 2008 August 13 (Sidolietal..2009b).. and on 2009 March 10 (Romanoetal..2009).," was observed in outburst with /BAT on three occasions, on 2008 April 8 \citep{sidoli09}, on 2008 August 13 \citep{sidoli09b}, and on 2009 March 10 \citep{romano09}."250. In only the first two cases. sslewed to the source and observations with the X-ray Telescope. XRT. were carried out.," In only the first two cases, slewed to the source and observations with the X-ray Telescope, XRT, were carried out."251 During the 2008 April 8 outburst. XRT observed ~387 s after the BAT trigger.," During the 2008 April 8 outburst, XRT observed $\sim$ 387 s after the BAT trigger."252 These data showed that the source was rapidly (~ 1000 s) decreasing in intensity. and the X-ray spectrum (0.3-10 keV) could be reproduced well by using an absorbed (Nyj=13 «1077 2) power law (hereafter. PL) model (photon index Dz1.5).," These data showed that the source was rapidly $\sim$ 1000 s) decreasing in intensity, and the X-ray spectrum (0.3-10 keV) could be reproduced well by using an absorbed $N_{\rm H}$ $\times$ $^{22}$ $^{-2}$ ) power law (hereafter, PL) model (photon index $\Gamma$ =1.5)."253 The 0.5-100 keV X-ray luminosity was 73.0 «1076 ere/s. Sidolietal.(2009) also performed an analysis of the bbroad band (0.3-60 keV) spectrum of dduring this outburst. and found that this spectrum could be reasonably well described by using either a power law with a cutoff at high energy (~13 keV). or a Comptonizing plasma model in XSPEC).," The 0.5-100 keV X-ray luminosity was $\sim$ $\times$ $^{36}$ erg/s. \citet{sidoli09} also performed an analysis of the broad band (0.3-60 keV) spectrum of during this outburst, and found that this spectrum could be reasonably well described by using either a power law with a cutoff at high energy $\sim$ 13 keV), or a Comptonizing plasma model in )."254 For the outburst of 2008 August 13. XRT data were obtained starting from ~390 s after the BAT trigger. and revealed a more complex behavior than that observed during the previous event (Sidolietal..2009b).," For the outburst of 2008 August 13, XRT data were obtained starting from $\sim$ 390 s after the BAT trigger, and revealed a more complex behavior than that observed during the previous event \citep{sidoli09b}."255. A time-resolved analysis showed that the source X- spectrum could be fit equally well by using an absorbed PL or a black-body (BB) model with constant photon index or temperature (T~1.2. kTpp-—1.8 keV) and a varying absorption column density (in the range «107? 7).," A time-resolved analysis showed that the source X-ray spectrum could be fit equally well by using an absorbed PL or a black-body (BB) model with constant photon index or temperature $\Gamma$$\sim$ 1.2, $kT_{\rm BB}$$\sim$ 1.8 keV), and a varying absorption column density (in the range $\times$ $^{22}$ $^{-2}$ )."256 The combined XRT+BAT broad band (0.3-60 keV) spectrun= could be well fit by using either a model of Comptonization of seed photons in a hot plasma in XSPEC) or a model., The combined XRT+BAT broad band (0.3-60 keV) spectrum could be well fit by using either a model of Comptonization of seed photons in a hot plasma in ) or a model.257 The BMC comprises a BB component and a component accounting for the Comptonization of the BB due to thermal and/or dynamical (bulk) Comptonization., The BMC comprises a BB component and a component accounting for the Comptonization of the BB due to thermal and/or dynamical (bulk) Comptonization.258 The 0.1-100 keV X-ray luminosity derived from the simultaneous XRT+BAT spectrum was ~3.8 « 1000 erg/s. On 2009 March 10. aagan triggered BAT (Romanoetal..2009).," The 0.1-100 keV X-ray luminosity derived from the simultaneous XRT+BAT spectrum was $\sim$ $\times$ $^{36}$ erg/s. On 2009 March 10, again triggered BAT \citep{romano09}."259. On this occasion. ddid not perform any quick slew towards the source and XRT data were accumulated only ~1.5 h after the BAT trigger.," On this occasion, did not perform any quick slew towards the source and XRT data were accumulated only $\sim$ 1.5 h after the BAT trigger."260 At this time. the source was already much fainter (X-rayluminosity ~7 « 10°! erg/s. 2-10 keV). and the XRT spectrum could be reproduced well by using an absorbed power-law model (\yj=4 \ 10°? 2: ΓΞΙ 2).," At this time, the source was already much fainter (X-rayluminosity $\sim$ $\times$ $^{34}$ erg/s, 2-10 keV), and the XRT spectrum could be reproduced well by using an absorbed power-law model $N_{\rm H}$ $\times$ $^{22}$ $^{-2}$, $\Gamma$ =1.2)."261 Little is known about the quiescent emission of AAn oobservation in 1999 did not detect the source and placed à 3.6 upper limit on its X-ray luminosity of « 107? erg/s 2002)..," Little is known about the quiescent emission of An observation in 1999 did not detect the source and placed a 3 $\sigma$ upper limit on its X-ray luminosity of $\times$ $^{32}$ erg/s \citep[exposure time $\sim$13~ks,][]{sakano02}. ."262 A —5 ks oobservation in. 2001 caught the source in a relatively lowluminosity state « 107 erg/s) and the X-ray spectrum was fit well by using an absorbed power-law model (Ny , A $\sim$ 5 ks observation in 2001 caught the source in a relatively lowluminosity state $\times$ $^{33}$ erg/s) and the X-ray spectrum was fit well by using an absorbed power-law model \citep[N$_{\rm H}$ 263echuiques based ou the streneth of the mid-IR cussion with respect to the optical and X-ray cussion have also cen developed aud applied to select large populations of candidate Conptou-thick ACN up to. =23 (DidiDaueretal. 2010).,"techniques based on the strength of the mid-IR emission with respect to the optical and X-ray emission have also been developed and applied to select large populations of candidate Compton-thick AGN up to $z=2-3$ \citep{daddi07,alex08,fiore09,bauer10}."264. Once more. however. the lack of X- spectra prevents an unambiguous determination of he absorbing column density. οπιο the measuremoeuts x these works hugely uncertain.," Once more, however, the lack of X-ray spectra prevents an unambiguous determination of the absorbing column density, making the measurements by these works largely uncertain."265" Iu this paper we report the discovery of a ""bona-fide Comptou-thick ACN at :=L76 in the | Ms CDES.", In this paper we report the discovery of a “bona-fide” Compton-thick AGN at $z$ =4.76 in the 4 Ms CDFS.266" A concordance cosnologv with My=το kins | |l Q,,=0.27. Oy=0.73 is adopted throughout this paper."," A concordance cosmology with $H_0=70$ km $^{-1}$ $^{-1}$, $\Omega_m=0.27$, $\Omega_{\Lambda}=0.73$ is adopted throughout this paper."267 We searched for Wo band dropout objects in the UST/ACS v2.0 data of COODS-South (Ciüavaliscoetal.POOL)— associated. with XN-rav ciuission in the [| Mx Chandra image., We searched for $V-$ band dropout objects in the HST/ACS v2.0 data of GOODS-South \citep{java04} associated with X-ray emission in the 4 Ms $Chandra$ image.268 We used the Voyg-dropout selection criteria from Oeschetal.(2007). which cftectively pick sources at LTci:<5g.," We used the $V_{606}$ -dropout selection criteria from \citet{oesch07}, which effectively pick sources at $4.7<z<5.7$."269 Details on the production of the V-dropout catalog are given iu Suetal. (2011)., Details on the production of the $V$ -dropout catalog are given in \citet{su11}.270 Additionally. we required a stellarity piriuueter (CLASS_SSTAR} ereater than 0.9 in the κου baud to choose poiut-like sources.," Additionally, we required a stellarity parameter STAR) greater than 0.9 in the $z_{850}$ band to choose point-like sources."271 This led to 21 star-like Voyg-dropouts with κου<25.5. among which there are four z5 galaxies. eleven stars. three lower-vedshitt ealaxies. and one z—5 ACN. which is the ouly object detected in N-vavs (ΝΕΟ102 in the 1. MS CDFS catalog of Nueetal. 20011)).," This led to 21 star-like $V_{606}$ -dropouts with $z_{850}<25.5$, among which there are four $\sim$ 5 galaxies, eleven stars, three lower-redshift galaxies, and one $\sim$ 5 AGN, which is the only object detected in X-rays (XID403 in the 4 Ms CDFS catalog of \citealt{xue11}) )."272 The renuünius two candidates have not been ideutified spectroscopically., The remaining two candidates have not been identified spectroscopically.273 The measured AD inaguitudes of XID103 in ACS images are: συ=26.5LEO.fers=25.21EOOL.κου25.05 0.01. The 5o detection limit in the D i55-baud is 28.1 AB mae.," The measured AB magnitudes of XID403 in ACS images are: $V_{606}=26.84\pm0.10, i_{775}=25.21\pm0.04, z_{850}= 25.05\pm0.04$ The $\sigma$ detection limit in the $B_{435}$ -band is 28.4 AB mag."274" NID103. (02990503:32:29.209. 8,59997-27:56:19.5)x was recognized as an ACN at 2=LT6 based ou FORS-2 spectroscopy (Vauzellaetal.2006)."," XID403 $\alpha_{J2000}$ =03:32:29.29, $\delta_{J2000}$ =-27:56:19.5) was recognized as an AGN at $z=4.76$ based on FORS-2 spectroscopy \citep{vanzella06}."275. Its optical spectrum exhibits a narrow (EWIIMEE1000 uis 1) Lye emission line aud a broader (ΕΠΙ2000 Jun +) NVA1210 cluission liue. with an iutegrated flux simular to Lye.," Its optical spectrum exhibits a narrow $\lesssim1000$ km $^{-1}$ ) $\alpha$ emission line and a broader $\sim2000$ km $^{-1}$ ) $\lambda$ 1240 emission line, with an integrated flux similar to $\alpha$."276 A more recent spectrun with DEIMOS/RKeck confirmsboth features (Coppinetal.2009)., A more recent spectrum with DEIMOS/Keck confirmsboth features \citep{coppin09}.277.. The Spectral Euergy Distiibution of NID103 was published by Coppinctal. (20093., The Spectral Energy Distribution of XID403 was published by \citet{coppin09}.278. Based on a LABOCA detection at STOfan. they showed that this source is a bright subiuillimeter galaxy with SFR~1000 AZ. vr+.," Based on a LABOCA detection at $\mu$ m, they showed that this source is a bright submillimeter galaxy with $\sim$ 1000 $M_{\odot}$ $^{-1}$."279 A large reservoir of molecular gas (~1.6.1019 AL) was also identified through observations (Coppinetal.2010).," A large reservoir of molecular gas $\sim1.6\times10^{10}\,M_{\odot}$ ) was also identified through CO(2-1) observations \citep{coppin10}."280. We considered the same optical to mid-IR datasets used by Coppinetal.(2009) aud improved ou the SED bv adding the detection at L1 nuu by AzTec/ASTE ΕΕτν=3340 οι Scott et al., We considered the same optical to mid-IR datasets used by \citet{coppin09} and improved on the SED by adding the detection at 1.1 mm by AzTec/ASTE $f_{1.1mm}=3.3\pm0.5$ mJy; Scott et al.281" 2010) aud the Y. J aud As magnitudes from the deep NIR inagiug bx ILAWRK-I/VLT (Vy,=2h5640.12. Jyp=21374011. Ίντο=5102020: Castellanoetal. 20101)."," 2010) and the $Y$, $J$ and $Ks$ magnitudes from the deep NIR imaging by HAWK-I/VLT $Y_{AB}=24.56\pm0.12$, $J_{AB}=24.37\pm0.14$, $K_{AB}=24.03\pm0.20$; \citealt{castellano10}) )."282 This object is also detected (at ~ 30) at 1.1 GITz with a peak flux of 22.3 Jv (N. Miller priv., This object is also detected (at $\sim3\sigma$ ) at 1.4 GHz with a peak flux of 22.3 $\mu$ Jy (N. Miller priv.283 comin)., comm).284 Unfortuuately. it falls just outside the areas covered by the 167221 Spitzer/IRS mosaic Cleplitzetal.2011) and GOODS-Ierschel (PI D. Elbaz).," Unfortunately, it falls just outside the areas covered by the $16\mu$ m Spitzer/IRS mosaic \citep{teplitz11} and GOODS-Herschel (PI D. Elbaz)."285 NID103 is not detected in the 3Ms NALAL inage of the CDES., XID403 is not detected in the 3Ms $XMM$ image of the CDFS.286 A total exposure of =[Ms has Όσοι acciniulated on the CDFS as a result of 51 individual observations with ACTS-I performed during three different time periods: (kN Ms in 2000. ~1 Ms in 2007 aud ~2 Ms in 2010.," A total exposure of $\lesssim 4$ Ms has been accumulated on the CDFS as a result of 54 individual observations with ACIS-I performed during three different time periods: $\sim0.8$ Ms in 2000, $\sim1$ Ms in 2007 and $\sim 2$ Ms in 2010."287 X-ray data products. iucliding eveut files for cach observation aud also for the mereed dataset are publiclyavailable!.," X-ray data products, including event files for each observation and also for the merged dataset are publicly."288. In this paper we use the data products by Auectal.(2011) who derived N-rav source catalogs roni a full reprocessing aud astrometric recalibration of he event files., In this paper we use the data products by \citet{xue11} who derived X-ray source catalogs from a full reprocessing and astrometric recalibration of the event files.289 We used CIAO vill aud the Fuutools to perform N-ray aperture photometry at the xositiou of NID103., We used CIAO v4.1 and the Funtools to perform X-ray aperture photometry at the position of XID403.290 The separation between the optical aud N-rav centroids is ~OL. which is well within he lo XN-rav source positional wucertainty (70.177).," The separation between the optical and X-ray centroids is $\sim 0.4""$, which is well within the $1\sigma$ X-ray source positional uncertainty $\sim$ 0.47”)."291 To uaxiuize the S/N ratio. we measured the source counts in different bands within a πα aperture of 37 radius. which encloses ~50% of the PSF at 1.5 keV at the source ocation (~8 απο off-axis).," To maximize the S/N ratio, we measured the source counts in different bands within a small aperture of 3” radius, which encloses $\sim$ of the PSF at 1.5 keV at the source location $\sim8$ arcmin off-axis)."292 We measured 37.048.7 iet counts in the (.9-b keV band. correspoucding to a ~L2«detection!!.," We measured $37.0\pm 8.7$ net counts in the 0.9-4 keV band, corresponding to a $\sim4.2\sigma$."293. We verified that similar results are obtained when using differcut local background regions., We verified that similar results are obtained when using different local background regions.294" The larduess ratio. defined as IIR. = ο)|S). where S and II are the net counts observed iu the 0.5-2 keV and 2-7 keV bands. respectively, is 20.23+0.21 (not corrected for vignettiug)."," The hardness ratio, defined as HR = (H-S)/(H+S), where S and H are the net counts observed in the 0.5-2 keV and 2-7 keV bands, respectively, is $0.23\pm0.24$ (not corrected for vignetting)."295 This value. for an ACN at 2~Db with a standard intrinsic spectrum (1... TP= La) is highly sugeestive of heavy obscuration.," This value, for an AGN at $z\sim5$ with a standard intrinsic spectrum (i.e., $\Gamma=1.8$ ) is highly suggestive of heavy obscuration."296" For an. ACN at: ; ~5with Ny: z5l07'«m7 comparison, is expected to have TRS0.35 at the source position."," For comparison, an AGN at $z\sim5$ with $N_H\lesssim10^{23}$ $^{-2}$ is expected to have $\lesssim-0.35$ at the source position."297 We extracted the A-rav spectrun in the 0.5-7 keV baud using the same 3° radius aperture aud verified that below 0.9 keV aud above L keV the source enussion is indistinemishable from the background., We extracted the X-ray spectrum in the 0.5-7 keV band using the same 3” radius aperture and verified that below 0.9 keV and above 4 keV the source emission is indistinguishable from the background.298" The προσ and response files were created using the script in CIAO,", The spectrum and response files were created using the script in CIAO.299 Since docs not account for he PSF fraction when building up the effective. area fle. we multiplied the 0.5-2 keV aud 2-10 keV. fluxes as obtained from the spectral fit by a factor of 2 aud 2.5. respectively. to recover the full aperture-correctecd X-ray fixes.," Since does not account for the PSF fraction when building up the effective area file, we multiplied the 0.5-2 keV and 2-10 keV fluxes as obtained from the spectral fit by a factor of 2 and 2.5, respectively, to recover the full aperture-corrected X-ray fluxes."300 We fouud consistent results either usiug spectral responses extracted from individual observations. or hose obtained as an exposurewveiehted mean over all individual responses.," We found consistent results either using spectral responses extracted from individual observations, or those obtained as an exposure-weighted mean over all individual responses."301 To double check the reliability of lis procedure we also built spectral response fles for one of the Chandra exposures (ObsID-8591) using the[um ACTIS-Extract software (Broosetal.2010). which allows xoper construction of effective area files at auv PSF Traction., To double check the reliability of this procedure we also built spectral response files for one of the $Chandra$ exposures (ObsID=8594) using the ACIS-Extract software \citep{broos10} which allows proper construction of effective area files at any PSF fraction.302Again. consistent results are found when usiug he ACIS-Extract responses.,"Again, consistent results are found when using the ACIS-Extract responses."303 We analyzed the X-ray spectrin with NSPEC v1.3.2 using the Cash statistic (Cash1979) to estimate the best-fit paramcters., We analyzed the X-ray spectrum with XSPEC v11.3.2 using the Cash statistic \citep{cash79} to estimate the best-fit parameters.304 Errors are quoted at lo confidence level., Errors are quoted at $1\sigma$ confidence level.305 We first fitted the data using a powerlaw spectrum modified by ealactic absorption. which returns T=0.617 m," We first fitted the data using a powerlaw spectrum modified by galactic absorption, which returns $\Gamma=-0.64^{+1.15}_{-0.73}$ ."306 We then used the model (Yaqoob 1997).. which follows the propagation of X-ray photous within a uniform. spherical," We then used the model \citep{yaqoob97}, , which follows the propagation of X-ray photons within a uniform, spherical"307resultant probabilities that the frequency distributions of the type 1 aud type 2 AGNs come from the same underlying populations are 2.0&10.9 and 5.9«10© for INe v|A3126/|O. 111]|A5007 and [Fe vir|AG0sT/[O niA5007. respectively.,"resultant probabilities that the frequency distributions of the type 1 and type 2 AGNs come from the same underlying populations are $2.0 \times 10^{-6}$ and $5.9 \times 10^{-6}$ for [Ne $\lambda$ 3426/[O $\lambda$ 5007 and [Fe $\lambda$ 6087/[O $\lambda$ 5007, respectively."308 It is thus confirmed that the two emissiou-liue flux ratios are significantly larger iu the type 1 ACNs than in the tvpe 2 ACNs., It is thus confirmed that the two emission-line flux ratios are significantly larger in the type 1 AGNs than in the type 2 AGNs.309 This result is consistent with the previous reports by Muraviuna Taniguchi (1998a). Nagao et al. (," This result is consistent with the previous reports by Murayama Taniguchi (1998a), Nagao et al. ("31020005). acd Nagao et al. (,"2000b), and Nagao et al. ("3112001b).,2001b).312 They concluded. that these differences in the emissiou-liue fiux ratios are due to the oricutation effect: ic.. the DINER is ocated very close to the nucleus aud thus can be hidden w dusty tori if it is observed from an edge-on view toward he tori.," They concluded that these differences in the emission-line flux ratios are due to the orientation effect; i.e., the HINER is located very close to the nucleus and thus can be hidden by dusty tori if it is observed from an edge-on view toward the tori."313 Coutrary to the above results. the difference in the fiux ratio of [Fe v11]AGOST/[Ne v|A3126 between the type 1 aud vpe 2 AGNs is fairly πα. as shown in Figure 2.," Contrary to the above results, the difference in the flux ratio of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 between the type 1 and type 2 AGNs is fairly small, as shown in Figure 2."314 The average and the l-o standard deviation of this flux ratio are 0.523 4 0.161 for the type 1 AGNs and 0.335 4 0.315 or the type 2 ones., The average and the $\sigma$ standard deviation of this flux ratio are 0.523 $\pm$ 0.461 for the type 1 AGNs and 0.335 $\pm$ 0.315 for the type 2 ones.315 The average aud the l-o stanclard deviation for the all sample (i.c.. 58 ACNs) are 0.115 + 115. aud the median value is 0.309.," The average and the $\sigma$ standard deviation for the all sample (i.e., 58 AGNs) are 0.445 $\pm$ 0.415, and the median value is 0.309."316 Although the flix ratios of the tvpe 2 AGNs sccm to be somewhat smaller han the type 2 ACGNs. the difference is not. significant.," Although the flux ratios of the type 2 AGNs seem to be somewhat smaller than the type 2 AGNs, the difference is not significant."317 The K-S statistical test results in the Ίντο probability of l.l « 7. which means that the tuferred difference is uarginal.," The K-S statistical test results in the K-S probability of 1.1 $\times$ $^{-2}$, which means that the inferred difference is marginal."318 Iu order to exanune whether or not the IIINETRs are dusty. we investigate how the cuussiou-line flux ratio of [Fe VIAGOST/[Ne v]A3126 depends on some physical properties bv performing calculations of photoionization models with aud without dust erains.," In order to examine whether or not the HINERs are dusty, we investigate how the emission-line flux ratio of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 depends on some physical properties by performing calculations of photoionization models with and without dust grains."319 Our method aud results of the photoionization model caleulatious are preseuted below., Our method and results of the photoionization model calculations are presented below.320 We carry out several photoionization model calculations by using the publicly available code Cloudy version 91.00 (Ferland 1997. 2000).," We carry out several photoionization model calculations by using the publicly available code $Cloudy$ version 94.00 (Ferland 1997, 2000)."321 Tere we assume uniforii deusitv gas clouds with a plane-parallel ecometry., Here we assume uniform density gas clouds with a plane-parallel geometry.322 The parameters for the calculations are: (1) the spectral cucrey distribution (SED) of the input continuum radiation: (2) the lvdrogen density of a cloud (ay): (3) the ionization parameter (C). Le. the ratio of the ionizing photon density to the ivdrogeu density at the imadiated surface of a cloud: (1) he colin density of a cloud (Ny): aud (5) the elemental composition aud the dust abundance of the gas.," The parameters for the calculations are: (1) the spectral energy distribution (SED) of the input continuum radiation; (2) the hydrogen density of a cloud $n_{\rm H}$ ); (3) the ionization parameter $U$ ), i.e., the ratio of the ionizing photon density to the hydrogen density at the irradiated surface of a cloud; (4) the column density of a cloud $N_{\rm H}$ ); and (5) the elemental composition and the dust abundance of the gas."323" We adopt the SED in the form of f£,~v' with a=2.5 or A> Lyre a=1 between 10 pau and 50 keV, and a=2.0 for fv>50 keV. taking account of he SEDs actually observed in ACNs (eg... Ioski 1978: Storchi-Bergimaun Pastoriza 1989. 1990: Iiunev ct al."," We adopt the SED in the form of $f_{\nu} \propto \nu^{\alpha}$ with $\alpha = 2.5$ for $\lambda > 10$ $\mu$ m, $\alpha = -1.5$ between 10 $\mu$ m and 50 keV, and $\alpha = -2.0$ for $h\nu > 50$ keV, taking account of the SEDs actually observed in AGNs (e.g., Koski 1978; Storchi-Bergmann Pastoriza 1989, 1990; Kinney et al."324 1991: see also To. Shields. Filippeuko 1993).," 1991; see also Ho, Shields, Filippenko 1993)."325 Since orbidden lines arise at eas clouds with a density near heir critical densities most effectively. the gas deusitv of he [Fe vijJAGOST aud the [Ne v|A3126 ciitting regions are expected to be a few «107 cm3.," Since forbidden lines arise at gas clouds with a density near their critical densities most effectively, the gas density of the [Fe $\lambda$ 6087 and the [Ne $\lambda$ 3426 emitting regions are expected to be a few $\times 10^7$ $^{-3}$."326" We thus perform uodel ruus with vy=10%.106"". aud Lae 3."," We thus perform model runs with $n_{\rm H} = 10^{6.5}, 10^{7.0}$, and $10^{7.5}$ $^{-3}$."327 As or the ionization parameter. we investigate models with - Πω”... ≀∶↓∩−⊳⊽↾∙↓∩−⊳∪⋜⋯≼↧↓∩↓⊳⊽↾∙⋜↧↕↑∐∪∏∶↴∙⊾∐↑↕∐∖↕∪⋯∑⋜↧↑↕∪∐.- ⋅⋅⋅ ≻⋜∐⋅⋜⋯∐∖↑↸∖↥⋅∪↕≯∐↕⋀∖⊽⊏↕⊰↴∖↴↕↴∖↴↑∐∪∏∶↴∙⊾∐↑↑∪↴⋝↸∖↥⋅⋜↧∐↸∖↥⋅↕∐∶↴⋁∐∙↑∐⋜↧↑ ↕↴∖↴↑∪↴∖↴⋜↧⋅↖↽∙≀⊽≧∐∣⇉≺↸∖∙∶↴∙⊾∙∙⋀∖↕↿∐⋅⋜↧⋅↖↽⋜∐⊔⋜↧∙∖↽⊺⋜⋯↕∶↴⋁⋯⊳↕∐↕⋂∩≺∖↴⋝∶ ⋀∖⊽," As for the ionization parameter, we investigate models with $U = 10^{-2.5}$, $10^{-2.0}$ and $10^{-1.5}$, although the ionization parameter of HINERs is thought to be rather high, that is to say, $U \gtrsim 10^{-2}$ (e.g., Murayama Taniguchi 1998b; Nagao et al."328⋜↧∶↴∙⊾⋜⋯, 2001b).329↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖," Because the the column density of a cloud may be much different from object to object, we"330↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐," Because the the column density of a cloud may be much different from object to object, we"331↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑," Because the the column density of a cloud may be much different from object to object, we"332↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕," Because the the column density of a cloud may be much different from object to object, we"333↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟," Because the the column density of a cloud may be much different from object to object, we"334↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥," Because the the column density of a cloud may be much different from object to object, we"335↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅," Because the the column density of a cloud may be much different from object to object, we"336↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪," Because the the column density of a cloud may be much different from object to object, we"337↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕," Because the the column density of a cloud may be much different from object to object, we"338↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔," Because the the column density of a cloud may be much different from object to object, we"339↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪," Because the the column density of a cloud may be much different from object to object, we"340↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴," Because the the column density of a cloud may be much different from object to object, we"341↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝," Because the the column density of a cloud may be much different from object to object, we"342↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙," Because the the column density of a cloud may be much different from object to object, we"343↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣," Because the the column density of a cloud may be much different from object to object, we"344↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡," Because the the column density of a cloud may be much different from object to object, we"345↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸," Because the the column density of a cloud may be much different from object to object, we"346↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖," Because the the column density of a cloud may be much different from object to object, we"347↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸," Because the the column density of a cloud may be much different from object to object, we"348↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳," Because the the column density of a cloud may be much different from object to object, we"349↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑," Because the the column density of a cloud may be much different from object to object, we"350↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪," Because the the column density of a cloud may be much different from object to object, we"351↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪," Because the the column density of a cloud may be much different from object to object, we"352↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴," Because the the column density of a cloud may be much different from object to object, we"353↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝," Because the the column density of a cloud may be much different from object to object, we"354↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙," Because the the column density of a cloud may be much different from object to object, we"355↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿," Because the the column density of a cloud may be much different from object to object, we"356↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡," Because the the column density of a cloud may be much different from object to object, we"357↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸," Because the the column density of a cloud may be much different from object to object, we"358↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖," Because the the column density of a cloud may be much different from object to object, we"359↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼," Because the the column density of a cloud may be much different from object to object, we"360↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳," Because the the column density of a cloud may be much different from object to object, we"361↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑," Because the the column density of a cloud may be much different from object to object, we"362↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅," Because the the column density of a cloud may be much different from object to object, we"363↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅↖," Because the the column density of a cloud may be much different from object to object, we"364↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅↖↖," Because the the column density of a cloud may be much different from object to object, we"365↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅↖↖⇁," Because the the column density of a cloud may be much different from object to object, we"366↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅↖↖⇁↸," Because the the column density of a cloud may be much different from object to object, we"367↸∖↑⋜↧↕∙⊇∩∩∐⋝⋟∙↕≧↸∖↸⊳⋜⋯↴∖↴↸∖↑∐↸∖↑∐↸∖⋯↕⋯⊔∐≼∐∖∐↴∖↴↕↑⋅↖↽∪↕⋟ ⋜↧↸⊳↕∪∏≼⊔⊔⋜↧⋅↖↽↴⋈∖↕⊔⋯⊳∐≼∐↕−↥⋅↸∖↥⋅↸∖∐↑↕⋟↥⋅∪↕⊔∪↴⋝∙∣⋡↸∖↸⊳↑∪∪↴⋝∙↿⋡↸∖≼⊳↑⋅↖↖⇁↸∖," Because the the column density of a cloud may be much different from object to object, we"368Q0957|56LA. D. In Section 3. we measure and interprete a large collection of infrared/optical/ultraviolet Dux ratios.,"Q0957+561A, B. In Section 3, we measure and interprete a large collection of infrared/optical/ultraviolet flux ratios."369 Finally. Section 4 summarizes our main conclusions.," Finally, Section 4 summarizes our main conclusions."370 Spectra of Q0957|561X and. Q09571561B were obtained. in 1999. April 15 and 2000 June 23. respectively. with the Space Telescope Imaging Spectrograph (STIS) on board the LIST.," Spectra of Q0957+561A and Q0957+561B were obtained in 1999 April 15 and 2000 June 2–3, respectively, with the Space Telescope Imaging Spectrograph (STIS) on board the HST."371 Phe spectra of each image cover a wide range of wavelengths. from. the nearinfrared. (NI) to the ultraviolet (UV). since the €G230L. G430L and G750L eratings were used in the experiment.," The spectra of each image cover a wide range of wavelengths, from the near–infrared (NIR) to the ultraviolet (UV), since the G230L, G430L and G750L gratings were used in the experiment."372 The central wavelengths of these gratings are 0.24 sam (6230L). 0.43. fam (6G430L) and 0.77 pam (6GT50L).," The central wavelengths of these gratings are 0.24 $\mu$ m (G230L), 0.43 $\mu$ m (G430L) and 0.77 $\mu$ m (G750L)."373 A 52 0.2 arCsec slit was also used in each observation., A 52 $\times$ 0.2 $^2$ slit was also used in each observation.374" The final data (Hux vs. wavelength) are calibrated with the last version of the CALSTIS pipeline software. so the corrections for timedependent. sensitivity of the ALAALA detectors (which is relevant for the (2201, observations) as well as for chargetransfer elliciency of the CCD detector (relevant for the G430L ancl G750L observations) are incorporated."," The final data (flux vs. wavelength) are calibrated with the last version of the CALSTIS pipeline software, so the corrections for time–dependent sensitivity of the MAMA detectors (which is relevant for the G230L observations) as well as for charge–transfer efficiency of the CCD detector (relevant for the G430L and G750L observations) are incorporated."375 We also note that the LISTSTIS data of the two images do not incorporate any correction for crosscontamination of the spectra ancl contamination by the lens galaxy., We also note that the HST–STIS data of the two images do not incorporate any correction for cross–contamination of the spectra and contamination by the lens galaxy.376 However. the crosscontamination of the spectra anc the contamination by the galaxy light at the bluest wavelengths are expected to be negligible.," However, the cross–contamination of the spectra and the contamination by the galaxy light at the bluest wavelengths are expected to be negligible."377 Phe possible contamination of Q0957|561B (due to its proximity to the galaxy) only must be checked at the reciclest wavelengths., The possible contamination of Q0957+561B (due to its proximity to the galaxy) only must be checked at the reddest wavelengths.378 In the next section. using complementary photometric observations. we test the LISTSTIS continuum flux ratios at 0.540.65 jum. We firstly focus on the continuum in the wavelength range 0.22-1 yam. The data cover the interval 0.09 <Aye OAL jun in the rest [rame of the emitted radiation. νου. blue/UVY emission.," In the next section, using complementary photometric observations, we test the HST–STIS continuum flux ratios at 0.54–0.65 $\mu$ m. We firstly focus on the continuum in the wavelength range 0.22-1 $\mu$ m. The data cover the interval 0.09 $< \lambda_{qso} <$ 0.41 $\mu$ m in the rest frame of the emitted radiation, i.e., blue/UV emission."379 This emitted radiation comes to the lens galaxy at 0.16 Agar Ο.Ε pim (opticalUV)., This emitted radiation comes to the lens galaxy at 0.16 $< \lambda_{gal} <$ 0.74 $\mu$ m (optical/UV).380 We average the continuum Lux over independent intervals of 100 ((0.01 jum) avoiding the prominent emission/absorption lines and. bad. pixels. and then infer 32 flux ratios 7:4.," We average the continuum flux over independent intervals of 100 (0.01 $\mu$ m) avoiding the prominent emission/absorption lines and bad pixels, and then infer 32 flux ratios $B/A$."381 The ratios appear in Figure 1 (blue. green. ancl red. circles).," The ratios appear in Figure 1 (blue, green and red circles)."382 For comparison to previous work. two additional black open circles are depicted in Fig.," For comparison to previous work, two additional black open circles are depicted in Fig."383 1., 1.384 These two optical continuum flux ratios (black open circles) present no contamination by the lens galaxy licht. since the contamination in image D. was accurately subtracted (CGolicoechea et al.," These two optical continuum flux ratios (black open circles) present no contamination by the lens galaxy light, since the contamination in image B was accurately subtracted (Goicoechea et al."385 2005: see also Ovaldsen et al., 2005; see also Ovaldsen et al.386 2003: Ullánn et al., 2003; Ullánn et al.387 2003)., 2003).388 The corresponding frames in the V and. 77 bands were obtained with the Nordic Optical Telescope in 2000 February/March. ancl 2001 April., The corresponding frames in the $V$ and $R$ bands were obtained with the Nordic Optical Telescope in 2000 February/March and 2001 April.389 In Fig., In Fig.390 1 there is an apparent agreement. between the old. results ancl the new measurements. so the previous data validate all the LISTSTIS flux ratios at ezc 1.5 Cr αλ. A in jim).," 1 there is an apparent agreement between the old results and the new measurements, so the previous data validate all the HST–STIS flux ratios at $x \ge$ 1.5 $x$ = $\lambda$, $\lambda$ in $\mu$ m)."391 However. as the contamination could. be important at redder wavelengths (e.g. see the frames in the ff anc £ bands at the ratios ate« 1.5 are not considered from here on.," However, as the contamination could be important at redder wavelengths (e.g., see the frames in the $H$ and $I$ bands at http://cfa-www.harvard.edu/glensdata/Individual/Q0957.html), the ratios at $x <$ 1.5 are not considered from here on."392 Using the results abor> 1.5. we reach two conclusions: (a) the opticalUV. ratios are not achromatic ancl (b) there is a correlation between Df£ and ic including5 a bump close to. = 3.5. that resembles extinction laws for 5galaxies in the Local Group (Gordon οἱ al.," Using the results at $x \ge$ 1.5, we reach two conclusions: (a) the optical/UV ratios are not achromatic and (b) there is a correlation between $B/A$ and $x$, including a bump close to $x$ = 3.5, that resembles extinction laws for galaxies in the Local Group (Gordon et al."393 2003)., 2003).394 Are our results consistent with only extinction and no Cgravitational microlensinge (i.c. no ogravitational ellects due to collapsed objects in the lens galaxy)?.," Are our results consistent with only extinction and no gravitational microlensing (i.e., no gravitational effects due to collapsed objects in the lens galaxy)?."395 Raclio signals come from regions much larger than the optical/UV. continuum source. so microlensing is not expected to alfect those signals.," Radio signals come from regions much larger than the optical/UV continuum source, so microlensing is not expected to affect those signals."396 The extinction by dust is also irrelevant for radio Utixes., The extinction by dust is also irrelevant for radio fluxes.397 The lux ratio at racio wavelengths is thus assumed to be à ratio free of perturbations caused by microlensing and extinction. ancl we take DA = 0.55 (radio ratio) as the macrolens ratio (e.g.. Garrett ct al.," The flux ratio at radio wavelengths is thus assumed to be a ratio free of perturbations caused by microlensing and extinction, and we take $B/A$ = 0.75 (radio ratio) as the macrolens ratio (e.g., Garrett et al."398 1994)., 1994).399 We then fit the 20 data at c> 1.5 to the Milky Ways (MW)like extinction law (Cardelli. Clavton Mathis 1989: Falco et al.," We then fit the 20 data at $x 400\ge$ 1.5 to the Milky Way (MW)–like extinction law (Cardelli, Clayton Mathis 1989; Falco et al."401" 1999). varving the possible redshift of the dust systeni tans, = 0 (MW). 0.36 (lens galaxy). 1.125 (Lyman limit svstem). 14 (damped Lya system and quasars host galaxy)."," 1999), varying the possible redshift of the dust system: $z_{dust}$ = 0 (MW), 0.36 (lens galaxy), 1.125 (Lyman limit system), 1.4 (damped $\alpha$ system and quasar's host galaxy)."402 Details on the two Lyman systems can be found in Michalitsianos et al. (, Details on the two Lyman systems can be found in Michalitsianos et al. (4031997).,1997).404 The best fit is plotted in Pig., The best fit is plotted in Fig.405 1 (dashed line) and it corresponds to a dust system in the lens galaxy., 1 (dashed line) and it corresponds to a dust system in the lens galaxy.406 Phere is no surprise: the observed bump is placed at Ayer& 0.21 jm. i.c. à wavelength close to the centre of the well-known extinction feature for local galaxies (2175 Aj).," There is no surprise: the observed bump is placed at $\lambda_{gal} \approx$ 0.21 $\mu$ m, i.e., a wavelength close to the centre of the well-known extinction feature for local galaxies (2175 )."407 As our best solution does not accurately trace the observed trend (47 = 1.8). we use a more general extinction model.," As our best solution does not accurately trace the observed trend $\chi^2$ = 1.8), we use a more general extinction model."408 At wavelengths larger than 3000A.. the extinction curves for lines of sight in the Local Ciroup (MW. and Magellanie Clouds) are well described by a AIW-like extinction law. whereas at wavelengths shorter than 3000A.. the extinction curves follow a model including a linear background term. a Drude profile and a far-UV curvature term (Cordon et al.," At wavelengths larger than 3000, the extinction curves for lines of sight in the Local Group (MW and Magellanic Clouds) are well described by a MW-like extinction law, whereas at wavelengths shorter than 3000, the extinction curves follow a model including a linear background term, a Drude profile and a far-UV curvature term (Gordon et al."409 2003)., 2003).410 Therefore. we fit our 13 measurements ad Agar> 0.3 pam to a MW. model. and derive a differential extinction (D.V) — 67.5 d 5.0 mmag and a ratio of total to selective extinction in the V optical band Ay = 4.4 d 0.5 confidence intervals).," Therefore, we fit our 13 measurements at $\lambda_{gal} >$ 0.3 $\mu$ m to a MW model, and derive a differential extinction $\Delta E(B - V)$ = 67.5 $\pm$ 5.0 mmag and a ratio of total to selective extinction in the $V$ optical band $R_V$ = 4.4 $\pm$ 0.5 confidence intervals)."411 We also fit the ratios at A;« 0.3 pim to a background. || Drude model (we do not need to include the fa-UV curvature). setting the Drude parameters to ο = 4.6 and  = 1.," We also fit the ratios at $\lambda_{gal} <$ 0.3 $\mu$ m to a background + Drude model (we do not need to include the far-UV curvature), setting the Drude parameters to $x_0$ = 4.6 and $\gamma$ = 1."412 The two new fits are depicted in Eig., The two new fits are depicted in Fig.413 1 (solid lines)., 1 (solid lines).414 They are characterized bv 47 values of 0.7.0.8. which suggests," They are characterized by $\chi^2$ values of 0.7–0.8, which suggests"415In Newtonian limit. we can approximate i6taτὸ because Co1l for a non-relativistic shock.,"In Newtonian limit, we can approximate $ \langle G^{s-3} \rangle \sim \langle G \rangle ^{s-3}$ because $G-1\ll1$ for a non-relativistic shock."416 Thus the energy.o spectral index is ὃνgiven by the similar [orm as leq. (24)., Thus the energy spectral index is given by the similar form as $Eq.~(24)$ .417 Vietri insisted that (67 can not be approximated bv £657 in general for a relativistic shock. contrary to the Peacock's argument.," Vietri insisted that $ \langle G^{s-3} \rangle $ can not be approximated by $\langle G \rangle ^{s-3}$ in general for a relativistic shock, contrary to the Peacock's argument."418 In this section. we calculate. a power-law index using Peacock's approximation and. Vietri's formulation., In this section we calculate a power-law index using Peacock's approximation and Vietri's formulation.419 Phen we examine the characteristic of the power-law index and how the power-law index derived from Peacock’s approximation is different from the one derived. from. Vietri's method to examine the elfect of the variance., Then we examine the characteristic of the power-law index and how the power-law index derived from Peacock's approximation is different from the one derived from Vietri's method to examine the effect of the variance.420 We consider moclel A (where the large-angle scattering occurs in both the upstream anc the downstream). and model Bo (where he larec-angle scattering occurs in the downstream and »uticles are dellected by large-scale magnetic fields in the upstream).," We consider model A (where the large-angle scattering occurs in both the upstream and the downstream), and model B (where the large-angle scattering occurs in the downstream and particles are deflected by large-scale magnetic fields in the upstream)."421 We need the averaged chance probability for a xuticle to return from the downstream to the upstream. per a crossing cycle Py (dq. (, We need the averaged chance probability for a particle to return from the downstream to the upstream per a crossing cycle $\langle P_R \rangle$ (Eq. (42219)) and the averaged energv-gain factor tn€ (Iq. (,19)) and the averaged energy-gain factor $\langle \ln G \rangle$ (Eq. (423213) to caleulate the power-law index by Peacock's approximation.,21)) to calculate the power-law index by Peacock's approximation.424 Ln Fig., In Fig.425 l we show the averaged return oobability., 1 we show the averaged return probability.426 In the both cases of model A and mocdel D. the return probability converges along the shock velocity.," In the both cases of model A and model B, the return probability converges along the shock velocity."427 This is because the shock velocity converges to one third of the speed of light in the down stream when we use the Svnge equation as the equation of the state., This is because the shock velocity converges to one third of the speed of light in the down stream when we use the Synge equation as the equation of the state.428 So the probability for particles to catch up with the shock front converges., So the probability for particles to catch up with the shock front converges.429 In Fig., In Fig.430 2 we show the averaged energy gain factor., 2 we show the averaged energy gain factor.431 The value increases monotonically in model A (Fig., The value increases monotonically in model A (Fig.432 2 top)., 2 top).433 On the other hand. the value converges in model D (Fig.," On the other hand, the value converges in model B (Fig."434 2 bottom)., 2 bottom).435 Particles eet more energv bv the shock-acceleration when he shock velocity becomes faster and the scattering angle xconies larger., Particles get more energy by the shock-acceleration when the shock velocity becomes faster and the scattering angle becomes larger.436 In model X. particles are scattered at à large angle and get large energy as the shock moves with the large velocity.," In model A, particles are scattered at a large angle and get large energy as the shock moves with the large velocity."437 In model B. the dellection angle becomes smaller as he shock velocity becomes Large.," In model B, the deflection angle becomes smaller as the shock velocity becomes large."438 Both of the two ellects (a arge velocity and a small scattering angle) cancel out cach other and the energy amplification converges., Both of the two effects (a large velocity and a small scattering angle) cancel out each other and the energy amplification converges.439 In Fig., In Fig.440 3. we show the power-law index.," 3, we show the power-law index."441" Phe power-law index becomes murder as a shock moves fast in model AX. Ehis is because the averaged energv-gain factor increases monotonically,", The power-law index becomes harder as a shock moves fast in model A. This is because the averaged energy-gain factor increases monotonically.442 On the other hand. the power-law index converges in mocel B. Γι» is due to the convergence of the averaged return probability and the averaged energv-gain factor.," On the other hand, the power-law index converges in model B. This is due to the convergence of the averaged return probability and the averaged energy-gain factor."443 This convergence is the same behavior as Dednarz&Ostrowski(1998)., This convergence is the same behavior as \citet{b3}.444. They studied about the case of the small-angle scattering and the power-law index converges as the shock velocity becomes larec., They studied about the case of the small-angle scattering and the power-law index converges as the shock velocity becomes large.445 The power-law index converges to 4 for a non-relativistic case which is consistent with the result of the diffusive shock-acceleration., The power-law index converges to 4 for a non-relativistic case which is consistent with the result of the diffusive shock-acceleration.446 Next. we consider the dillerence between the power-law index derived from. Peacock’s approximation and. Vietri's formulation.," Next, we consider the difference between the power-law index derived from Peacock's approximation and Vietri's formulation."447 “Phe power-law index converges in model D. but the converged values are dillerent between Peacock’s approximation and Vietris formulation.," The power-law index converges in model B, but the converged values are different between Peacock's approximation and Vietri's formulation."448 So Peacock's approximation is not suitable for relativistic shocks., So Peacock's approximation is not suitable for relativistic shocks.449 On the other hand. in the case of model A. the dillerence of the power-law index is the largest when Py. is about 3. and the index becomes nearly equal when Vy. is greater than 10.," On the other hand, in the case of model A, the difference of the power-law index is the largest when $\Gamma_s \beta_s$ is about 3, and the index becomes nearly equal when $\Gamma_s \beta_s$ is greater than 10."450 We can use Peacock’s approximation for highlv-relativistic shocks., We can use Peacock's approximation for highly-relativistic shocks.451 We also show that the power-law index caleulated by Vietri's formulation is harder than the one calculated by Peacock’s approximation in both cases of model A and model DB. We explain why the power-law spectrum becomes hard inthe discussion., We also show that the power-law index calculated by Vietri's formulation is harder than the one calculated by Peacock's approximation in both cases of model A and model B. We explain why the power-law spectrum becomes hard inthe discussion.452 Next we see what makes the dillerence between, Next we see what makes the difference between453ime interval. the mean gas fraction of all detected subclusters decreases by 0.63 dex. from ifatty=0.238 to (fat)=0.056.,"time interval, the mean gas fraction of all detected subclusters decreases by 0.63 dex, from $\langle f_{\rm gas}(t_1)\rangle=0.238$ to $\langle f_{\rm gas}(t_2)\rangle=0.056$."454" The mean half-mass radius of the decreases from m=0.020 pe to intr),=0.013 pe. which is a decrease of 0.18 dex."," The mean half-mass radius of the decreases from $\langle r_{\rm h}(t_1)\rangle=0.020$ pc to $\langle r_{\rm h}(t_2)\rangle=0.013$ pc, which is a decrease of 0.18 dex."455" Even though the shrinking of subclusters is a second order effect caused by gas accretion, it is interesting to ask which of both mechanisms contributes most to the decrease of the gas fraction."," Even though the shrinking of subclusters is a second order effect caused by gas accretion, it is interesting to ask which of both mechanisms contributes most to the decrease of the gas fraction."456 Is it mainly driven by the increasing mean stellar density of the subelusters or by the ongoing gas accretion onto the sink particles?, Is it mainly driven by the increasing mean stellar density of the subclusters or by the ongoing gas accretion onto the sink particles?457 To assess the relative contributions to gas depletion by üceretion. and. subcluster shrinking. we consider the spatial distribution of the sink particles and the gas.," To assess the relative contributions to gas depletion by accretion and subcluster shrinking, we consider the spatial distribution of the sink particles and the gas."458 Due to the relatively small numbers of stars in individual subelusters. it is best to examine the mean density profiles of the populations of subclusters in the two snapshots at fj and tf.," Due to the relatively small numbers of stars in individual subclusters, it is best to examine the mean density profiles of the populations of subclusters in the two snapshots at $t_1$ and $t_2$."459 Such a combination of the different density profiles decreases the influence of low-number statistics on the result., Such a combination of the different density profiles decreases the influence of low-number statistics on the result.460 In Fig. 8..," In Fig. \ref{fig:gas2},"461 we show the subcluster mass-weighted. mean cumulative mass distributions of gas. sink particles. and both combined.," we show the subcluster mass-weighted, mean cumulative mass distributions of gas, sink particles, and both combined."462 The distributions represent the enclosed mass fractions p. normalised to the sum of the subcluster mass M; and the enclosed gas mass within threestellar half-mass radii Mea: with i=[stars.gas.all) and Mit£) the enclosed mass at £=rín. which is the radial distance in units of thestellar half- radius.," The distributions represent the enclosed mass fractions $\mu$, normalised to the sum of the subcluster mass $M_{\rm cl}$ and the enclosed gas mass within three half-mass radii $M_{\rm gas}$: with $i=\{{\rm stars, gas, all}\}$ and $M_i(\xi)$ the enclosed mass at $\xi\equiv r/r_{\rm h}$, which is the radial distance in units of the half-mass radius."463 The mean distributions shown in Fig., The mean distributions shown in Fig.464 8 are weighted by subcluster mass to emphasise those subclusters with better statistics., \ref{fig:gas2} are weighted by subcluster mass to emphasise those subclusters with better statistics.465 A first comparison of both panels in Fig., A first comparison of both panels in Fig.466 8 shows that the gus fraction indeed decreases between fj and tf., \ref{fig:gas2} shows that the gas fraction indeed decreases between $t_1$ and $t_2$.467 The contribution to this decrease by subcluster shrinking can be estimated by a simple thought experiment. in which the gas distribution is kept fixed and the distribution of stellar mass is compressed by the appropriate amount.," The contribution to this decrease by subcluster shrinking can be estimated by a simple thought experiment, in which the gas distribution is kept fixed and the distribution of stellar mass is compressed by the appropriate amount."468 In the top panel of Fig. 8..," In the top panel of Fig. \ref{fig:gas2},"469 the half-mass radii imt)? and IC)? correspond to £=| and £=0.66. between which the enclosed gas mass differs by 0.26 dex.," the half-mass radii $\langle r_{\rm h}(t_1)\rangle$ and $\langle r_{\rm h}(t_2)\rangle$ correspond to $\xi=1$ and $\xi=0.66$, between which the enclosed gas mass differs by 0.26 dex."470 In other words. if the gas distribution was held fixed and the stellar distribution was shrunk appropriately. then the gas fraction within the new half-mass radius would have declined by 0.26 dex.," In other words, if the gas distribution was held fixed and the stellar distribution was shrunk appropriately, then the gas fraction within the new half-mass radius would have declined by 0.26 dex."471 This is a probe for the decrease of the gas fraction that is solely caused by the shrinking of the subclusters., This is a probe for the decrease of the gas fraction that is solely caused by the shrinking of the subclusters.472 Comparing it with the actual decrease of the mean gas fraction between { and t of 0.63 dex. we see that it covers about half of the decrease. with the remaining 0.37 dex covered by gas accretion itself — not only by adding to the mass in stars. but also by decreasing the gas mass.," Comparing it with the actual decrease of the mean gas fraction between $t_1$ and $t_2$ of 0.63 dex, we see that it covers about half of the decrease, with the remaining 0.37 dex covered by gas accretion itself – not only by adding to the mass in stars, but also by decreasing the gas mass."473 We conclude that the evacuation of the gas due to ongoing gas accretion is about equally important for the gas depletion as the shrinking of the subclusters., We conclude that the evacuation of the gas due to ongoing gas accretion is about equally important for the gas depletion as the shrinking of the subclusters.474 Apart from enabling a quantitative comparison of the effect of gas aecretion and cluster shrinking. Fig.," Apart from enabling a quantitative comparison of the effect of gas accretion and cluster shrinking, Fig."475 8. also demonstrates he spatial variation of the gas fraction in the subclusters., \ref{fig:gas2} also demonstrates the spatial variation of the gas fraction in the subclusters.476 At early imes. the gas is still prevalent in the outskirts of the subclusters. contributing of the enclosed mass at €=3.," At early times, the gas is still prevalent in the outskirts of the subclusters, contributing of the enclosed mass at $\xi=3$."477 At the end of the simulation this gas has mostly vanished. leaving only a few xreent of the mass within the half-mass radius as gas. and ypically at £=3.," At the end of the simulation this gas has mostly vanished, leaving only a few percent of the mass within the half-mass radius as gas, and typically at $\xi=3$."478 It is interesting to note that the relative increase of the enclosed gas mass Traction with respect to the enclosed sink particle mass fraction only occurs at radii where the latter flattens. i.e. the subclusters only become gas-rich at radii where very little stellar mass is esent.," It is interesting to note that the relative increase of the enclosed gas mass fraction with respect to the enclosed sink particle mass fraction only occurs at radii where the latter flattens, i.e. the subclusters only become gas-rich at radii where very little stellar mass is present."479 The influence of the gas on the subcluster dynamics is herefore best evaluated at radii smaller than where the flattening of fia Occurs., The influence of the gas on the subcluster dynamics is therefore best evaluated at radii smaller than where the flattening of $\mu_{\rm stars}$ occurs.480 At fi. the ratio between the enclosed stellar mass and gas mass just before the flattening is about 4:1. while at t» it has increased to [1:1.," At $t_1$, the ratio between the enclosed stellar mass and gas mass just before the flattening is about 4:1, while at $t_2$ it has increased to 11:1."481 This suggests that if feedback starts at a time t> b. the resulting gas expulsion will not strongly affect the subcluster dynamics. and that their virialised state (see Sect. 3.2))," This suggests that if feedback starts at a time $t>t_2$ , the resulting gas expulsion will not strongly affect the subcluster dynamics, and that their virialised state (see Sect. \ref{sec:virial}) )"482 will be largely retained., will be largely retained.483 Motivated by the low gas fractions found in Sect. 3..," Motivated by the low gas fractions found in Sect. \ref{sec:results},"484 we now address the response of the subclusters to gas expulsion in more detail., we now address the response of the subclusters to gas expulsion in more detail.485 The long-term response of the subclusters to gas expulsion can be evaluated by once again omitting the gas from the simulations and considering only the identified stellar subelusters and their evolution towards virial equilibrium., The long-term response of the subclusters to gas expulsion can be evaluated by once again omitting the gas from the simulations and considering only the identified stellar subclusters and their evolution towards virial equilibrium.486 Given a certain virial ratio and bound mass fraction. does a subcluster expand or shrink after gas removal?," Given a certain virial ratio and bound mass fraction, does a subcluster expand or shrink after gas removal?"487 We combine the data from the simulationswith a, We combine the data from the simulationswith a488and ZI) should be identical to those measured along the line of sight.,and $T_{\rm e}$ (H ) should be identical to those measured along the line of sight.489" However. Table 1 shows that the values of £2 ave. cousisteutly lower than those of £2 and the values of Tj, are conusisteutlv hieher than those of Jy. ποσο] favoring the presence of extremely-cold nebular components where the |O ΠΠ lines cannot be excltec but the ID brecombination spectrin is significantly: strenetened."," However, Table \ref{tab1} shows that the values of $t^2_v$ are consistently lower than those of $t^2$ and the values of $T_{0,v}$ are consistently higher than those of $T_0$, seemingly favoring the presence of extremely-cold nebular components where the [O ] lines cannot be excited but the H recombination spectrum is significantly strengthened."490 This conclusion. however. needs to be treated with ⋅ ↸⊳⋜⋯⊓∪," This conclusion, however, needs to be treated with caution."491↕∙↽∕∏∐∖↖↽⋜↧↕⋯∖∪↕∤∣−⋅≺∐∖∏∐↸∖≺∏⋝∙↖⇁⊏≺∣∙≺, The value of $t^2_v$ defined by Eq. (492⊔↥⋅↸∖↻↥⋅↸∖↴∖↴↸∖∐↑↴∖↴↑⊔∖↴ ⊳⋅≻ ⋅ ↸∖∐∩↸∖↥⋅⋜↧⊓∐⋅↸∖↖⇁⋜∐⋅↕⋜↧↑↕∪∐↴∖↴↕∐↖↽↸∖↕∪↸,4) represents the temperature variations in velocity space.493⊳↕↑⋅↖⇁↴∖↴↻⋯⊳↸∖∙↖↖⊽↸∖↴∖↴∐∪∏↕≺∏⋝↸∖⋜∐⋅ ⋯∐⊔∐≼⊔∐⋜↧↑↕∐∖↖↽⋜↧∐≼↕↑⋅↖↽∪↕↑∐↸∖⋜∐⋅∶↴∙↴∐⊔↸∖∐↑↑∐⋜↧↑⊺∠↾−⋅↕↴∖↴↸∖≺∏⋜↧↕ ∪↑∐↸∖⋯↸∖⋜⋯↴∖↴≺∣∏⋜⋯∖↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖↖↽⋜∐⋅↕⋜↧↑↕∪∐↻⋜∐⋅⋜↧⋯↸∖↑↸∖↥⋅⋜↧↕∪↕∩⊾ he line of sight requires two couditious to be satisfied: a) iebular expausion-velocitv dominates the broadeniug of iue profiles: b) the expansiou-velocity is monotonic along he nebular radius. as fouud iu iiost of PNe.," We should bear in mind that the validity of the argument that $t^2_v$ is equal to the mean square temperature variation parameter along the line of sight requires two conditions to be satisfied: a) nebular expansion-velocity dominates the broadening of line profiles; b) the expansion-velocity is monotonic along the nebular radius, as found in most of PNe."494 Otherwise. τρ provides only a lower indt to the temperature fluctuations along the liue⋅ of ⋟⋅sight.," Otherwise, $t^2_v$ provides only a lower limit to the temperature fluctuations along the line of sight."495 Tn any case. the comparison. of D)fZ aud f£n provides: an opportunity: to .investigate the preseuce of extremely-cold conmponeuts.," In any case, the comparison of $t^2_v$ and $t^2$ provides an opportunity to investigate the presence of extremely-cold components."496 This paper is begiuuiug such a project., This paper is beginning such a project.497 Iieheraesolutiou spectroscopic data of PNe with a high expausion velocity will be invaluable for this study., Higher-resolution spectroscopic data of PNe with a high expansion velocity will be invaluable for this study.498 Figure 2 also depicts that if temperature variations exist. he O lines show a profile that differs from that of the O ul lines. which cau be attributed to the fact that CELs and ORLs weight lieh- and low-temperature regions. respectively.," Figure \ref{t_flu} also depicts that if temperature variations exist, the O lines show a profile that differs from that of the [O ] lines, which can be attributed to the fact that CELs and ORLs weight high- and low-temperature regions, respectively."499 Iu the low velocity regious. the clectrou cluperature is relatively hieh.," In the low velocity regions, the electron temperature is relatively high."500 As a result. tle O lines rave a deeper “gap” between the two peaks than the [O ines.," As a result, the O lines have a deeper “gap” between the two peaks than the [O ] lines."501 In the lieh velocity regions which correspoud to the ine wines. the electron temperature is sheltly higher. depressing recombination lines.," In the high velocity regions which correspond to the line wings, the electron temperature is slightly higher, depressing recombination lines."502 Consequeuth. the ο ines are narrower than the JO lines," Consequently, the O lines are narrower than the [O ] lines."503 The O line xofiles differ most significantly from those of the 1Ο A1363 line than of the [O mt} A5007 line. since the |O ΑΡΟΟΤ line has a lower excitation cluperature than the O U1 aurora line.," The O line profiles differ most significantly from those of the [O ] $\lambda4363$ line than of the [O ] $\lambda5007$ line, since the [O ] $\lambda5007$ line has a lower excitation temperature than the [O ] auroral line."504 Ilieh spectralvesohition observations of PN NCC 7009 show that the |O 1) A1363 CEL is broader than the O ORLs bv a Actor of abou 1.5 Barlow et al. 2006))., High spectral-resolution observations of PN NGC 7009 show that the [O ] $\lambda4363$ CEL is broader than the O ORLs by a factor of about 1.5 (Barlow et al. \cite{barlow06}) ).505" For their observatious. a 0.9”L0"". nuaesc-slicer was placed on a bright-edee region located 5.6%it nortlavest of the central star."," For their observations, a $0.9\arcsec\times0.9\arcsec$ image-slicer was placed on a bright-edge region located $5.6\arcsec$ northwest of the central star."506 We attempted to model the observed CEL/ORL width discrepancy., We attempted to model the observed CEL/ORL width discrepancy.507 However. we fud that chemicallv-hoimogeueous nebula models are unable to reproduce à CEL/ORL width ratio that is ereater than 1.2 because our modecllings are unable to achieve such laree temperature variations.," However, we find that chemically-homogeneous nebula models are unable to reproduce a CEL/ORL width ratio that is greater than 1.2 because our modellings are unable to achieve such large temperature variations."508 The modelled CEL/ORL width ratio is csseutially independent of the adopted velocity feld and is mseusitive to the assumed cleusity distribution since the |O 11] A1363 CEL aud the O ORLs both have high critical densities., The modelled CEL/ORL width ratio is essentially independent of the adopted velocity field and is insensitive to the assumed density distribution since the [O ] $\lambda4363$ CEL and the O ORLs both have high critical densities.509 To match the observed CEL/ORL widths properly. the electron temperature is required to increase sharply outwards. which is difficult to achieve using pure photoionization modelling of gaseous webulac.," To match the observed CEL/ORL widths properly, the electron temperature is required to increase sharply outwards, which is difficult to achieve using pure photoionization modelling of gaseous nebulae."510 The most direct explanation for the observed aree CEL/ORL width discrepancy is that the O?!. CELs and ORLs originate in differcu uecbular components which jwe separate kinematic fields., The most direct explanation for the observed large CEL/ORL width discrepancy is that the $^{2+}$ CELs and ORLs originate in different nebular components which have separate kinematic fields.511 Liu (2003)) proposed lat evaporating plauetesimals within the nebulae might xoduce. cold. U-deficient inclusions where ORLs are ereatly enhanced aud CELs aro depressed., Liu \cite{liu03}) ) proposed that evaporating planetesimals within the nebulae might produce cold H-deficient inclusions where ORLs are greatly enhanced and CELs are depressed.512 If this is ie case. these IEdeficieut. components should move im relatively-stable orbits. aud thus have a lower velocity dispersion than the diffuse uebular gas;," If this is the case, these H-deficient components should move in relatively-stable orbits, and thus have a lower velocity dispersion than the diffuse nebular gas."513 As an alternative. --- relmaius possible that heating bw shock waves du he outer regions may be significant and leads to a sharp temperature increase in the liel-velocity regions.," As an alternative, it remains possible that heating by shock waves in the outer regions may be significant and leads to a sharp temperature increase in the high-velocity regions."514 Furthermore. Stasinsska Szezerba (2001)) showed that Nhotoelectrie heating bv dust eraius can cause arec cluperature fluctuations du the presence of density iuhoimogeneities.," Furthermore, Stasińsska Szczerba \cite{stasinska01}) ) showed that photoelectric heating by dust grains can cause large temperature fluctuations in the presence of density inhomogeneities."515 However. iu the case of a constant dust-o-gas ratio iu the eutire nebula. the temperature is wore efficieutlvy cuhauced iu the iuuer regions.," However, in the case of a constant dust-to-gas ratio in the entire nebula, the temperature is more efficiently enhanced in the inner regions."516 If the CEL/ORL width discrepancy is caused by heating by dust. these dust erains nist be predominantly located iu the outer zones of the ionized. gas.," If the CEL/ORL width discrepancy is caused by heating by dust, these dust grains must be predominantly located in the outer zones of the ionized gas."517 The determination of nebular density structure is critical to the investigation of ORL/CEL abundauce discrepancies., The determination of nebular density structure is critical to the investigation of ORL/CEL abundance discrepancies.518 Viegas Cleeeee (199 1)) poimted out that the presence of high density regions could lead to an overestimate of {ο mn and consequentlv au nuderestimate of CEL abundances since the ΑΡΟΟΤ nebular line has a far lower critical deusitv than the ο mu A 1363 auroral line aud is nore siguificantly suppressed by collisional de-excitation in-high density regions., Viegas Clegg \cite{viegas94}) ) pointed out that the presence of high density regions could lead to an overestimate of $T_{\rm e}$ ([O ]) and consequently an underestimate of CEL abundances since the $\lambda$ 5007 nebular line has a far lower critical density than the [O ] $\lambda$ 4363 auroral line and is more significantly suppressed by collisional de-excitation in-high density regions.519 However. Liu et al. (20003) ," However, Liu et al. \cite{liubarlow00}) )"520claimed that density inhomogeucitics can be ruled out as the Cause of the ORL/CEL abuudance discrepancies because no correlation between the abundance discrepancies and critical densities was found., claimed that density inhomogeneities can be ruled out as the cause of the ORL/CEL abundance discrepancies because no correlation between the abundance discrepancies and critical densities was found.521 Line profiles cau be used to xobe density variations along the Ime of sight., Line profiles can be used to probe density variations along the line of sight.522 The |S uf AAGT31.671G double lues are classical clectron-deusity diagnostics.," The [S ] $\lambda\lambda6731,6716$ doublet lines are classical electron-density diagnostics."523 We therefore simulate the profiles of he [S r1] lines uuder the assmuptiou that density decreases outwards aud increases outwards. as slow iu Figs," We therefore simulate the profiles of the [S ] lines under the assumption that density decreases outwards and increases outwards, as shown in Figs."524 { and 5.. respectively.," \ref{n_neg}525 and \ref{n_pos}, respectively."526 We asstuue that the ionizing star has a eniperature of 191. For the nebulae that we model S! isthe dominant ionizatiou-state ofsulplur.," We assume that the ionizing star has a temperature of K. For the nebulae that we model, $^+$ isthe dominant ionization-state ofsulphur."527 Figures { and 5. indicate that the the profiles of the |S 1] AAGT3BL.6716 doublet lines differ if deusitv variations are," Figures \ref{n_neg} and \ref{n_pos} indicate that the the profiles of the [S ] $\lambda\lambda6731,6716$ doublet lines differ if density variations are"528Polars are magnetic cataclysmic binaries consisting of a late-type main-sequence star and a strongly magnetic white dwarf locked in synchronous rotation.,Polars are magnetic cataclysmic binaries consisting of a late-type main-sequence star and a strongly magnetic white dwarf locked in synchronous rotation.529 wwas one of the 11 polars known in the pre-ROSAT era. it was the second brightest at optical and at X-ray wavelengths after the prototypical system AM Herculis.," was one of the 11 polars known in the pre-ROSAT era, it was the second brightest at optical and at X-ray wavelengths after the prototypical system AM Herculis."530 It was studied with all major X-ray observatories (EINSTEIN. EXOSAT. GINGA. ROSAT) in the past and was always found in a high aceretion state.," It was studied with all major X-ray observatories (EINSTEIN, EXOSAT, GINGA, ROSAT) in the past and was always found in a high accretion state."531 EINSTEIN observations revealed the presence of uncorrelated soft and hard X-ray emission and were used to observationally establish the standard picture of magnetic accretion onto white dwarfs in the high /i-regime dominated by a shock-heated accretion column and cooling by free-free radiation (Beuermann. Stella Patterson 1987. henceforth BSP87).," EINSTEIN observations revealed the presence of uncorrelated soft and hard X-ray emission and were used to observationally establish the standard picture of magnetic accretion onto white dwarfs in the high $\dot{m}$ -regime dominated by a shock-heated accretion column and cooling by free-free radiation (Beuermann, Stella Patterson 1987, henceforth BSP87)."532 The absence of a pronounced soft X-ray excess made BSP87 to coin tthe textbook example of AM Herculis-type systems., The absence of a pronounced soft X-ray excess made BSP87 to coin the textbook example of AM Herculis-type systems.533 The shape of the X-ray light curves. in particular the presence of a soft X-ray absorption dip. was used to uncover the aceretion geometry.," The shape of the X-ray light curves, in particular the presence of a soft X-ray absorption dip, was used to uncover the accretion geometry."534 hhad a main accretion pole which was continuously in view. the observer has a moderate inclination with respect to the orbital plane. so that the line of sight crosses the aceretion stream on its way through the magnetosphere.," had a main accretion pole which was continuously in view, the observer has a moderate inclination with respect to the orbital plane, so that the line of sight crosses the accretion stream on its way through the magnetosphere."535 This special geometry allowed detailed stream-density diagnostics with GINGA and EXOSAT (Watson et al., This special geometry allowed detailed stream-density diagnostics with GINGA and EXOSAT (Watson et al.536 1989)., 1989).537 The accretion geometry was intensively studied using photo- and spectro-polarimetric data (e.g. Bailey et al., The accretion geometry was intensively studied using photo- and spectro-polarimetric data (e.g. Bailey et al.538 1982. Cropper 1985. Pitrola et al.," 1982, Cropper 1985, Piirola et al."539 1987. Meggitt Wickramasinghe 1989. Beuermann et al.," 1987, Meggitt Wickramasinghe 1989, Beuermann et al."540 2007)., 2007).541 The latter three papers agree that the white dwarf's magnetic field is probably more complex than that of a centered dipole., The latter three papers agree that the white dwarf's magnetic field is probably more complex than that of a centered dipole.542 The zero point of Bailey's ephemeris (Bailey et al., The zero point of Bailey's ephemeris (Bailey et al.543 1982) centered on the IR (X-ray) absorption dip is widely used in the literature., 1982) centered on the IR (X-ray) absorption dip is widely used in the literature.544 Pitrola et al. (, Piirola et al. (5451987) determined an updated orbital period based on a linear regression of the arrival times of linear polarisation pulses.,1987) determined an updated orbital period based on a linear regression of the arrival times of linear polarisation pulses.546 Beuermann et al. (, Beuermann et al. (5472007) derived a slightly revised ephemeris by including the ROSAT PSPC X-ray dip timings from July 1990.,2007) derived a slightly revised ephemeris by including the ROSAT PSPC X-ray dip timings from July 1990.548 Phases in this paper refer to Bailey’s phase zero and Piirola’s period., Phases in this paper refer to Bailey's phase zero and Piirola's period.549 tturned into a deep low state at V.=18 in 1997 (Wheatley Ramsay 1998) and remains therein since then., turned into a deep low state at $V\simeq 18$ in 1997 (Wheatley Ramsay 1998) and remains therein since then.550 À re-brightening was reported in VSNET on March 5. 2006 (ERI EF 20060305.724 at 14.2 unfiltered CCD based on the Henden-Sumner sequence). but the system returned to the low state shortly thereafter.," A re-brightening was reported in VSNET on March 5, 2006 (ERI EF 20060305.724 at 14.2 unfiltered CCD based on the Henden-Sumner sequence), but the system returned to the low state shortly thereafter."551 While in the high state the stellar photospheres are outshone by aceretion radiation. the low state offers the opportunity to investigate the stars. at least in principle.," While in the high state the stellar photospheres are outshone by accretion radiation, the low state offers the opportunity to investigate the stars, at least in principle."552 Since iis the polar with the shortest orbital period. P4=81 mmin. just a few minutes above the CV minimum period. low state observations are of utmost importance to test current scenarios of CV evolution and to search for the cool secondary.," Since is the polar with the shortest orbital period, $P_{\rm orb} =81$ min, just a few minutes above the CV minimum period, low state observations are of utmost importance to test current scenarios of CV evolution and to search for the cool secondary."553 Indeed. following the more indirect. conclusion by Beuermann et al. (," Indeed, following the more indirect conclusion by Beuermann et al. ("5542000) of a substellar secondary in ffrom the non-detection of any spectral signature of the companion in optical spectra. Howell Ciardi (2001) claimec the detection of the secondary in near-infrared spectra.,"2000) of a substellar secondary in from the non-detection of any spectral signature of the companion in optical spectra, Howell Ciardi (2001) claimed the detection of the secondary in near-infrared spectra."555 A more likely explanation of the observed infrared humps was given i terms of cyclotron radiation (Beuermann et al., A more likely explanation of the observed infrared humps was given in terms of cyclotron radiation (Beuermann et al.556 2000. Harriso et al.," 2000, Harrison et al."557 2004)., 2004).558 Beuermann et al. (, Beuermann et al. (5592000) also estimated the photospheriο temperature of the white dwarf from their low-resolutio optical spectra. Twp=9500€500 KK. one of the coldest WDs among all CVs.,"2000) also estimated the photospheric temperature of the white dwarf from their low-resolution optical spectra, $T_{\rm WD}560=9500\pm500$ K, one of the coldest WDs among all CVs."561 This allowed to draw some conclusions o the likely evolutionary state of the object., This allowed to draw some conclusions on the likely evolutionary state of the object.562 Recently. Szkody et al. (," Recently, Szkody et al. ("5632006) report on phase-resolved observations with the puzzling result of a distinct source of ultraviolet flux much larger than the underlying KK white dwarf.,2006) report on phase-resolved observations with the puzzling result of a distinct source of ultraviolet flux much larger than the underlying K white dwarf.564 Here we report on archival XMM-Newton observations of oobtained in a low aceretion state., Here we report on archival XMM-Newton observations of obtained in a low accretion state.565 We search for remaining X-ray emission in the low state either originating from the white dwarf or the secondary and analyse the data from the optical monitor taken through two different filters., We search for remaining X-ray emission in the low state either originating from the white dwarf or the secondary and analyse the data from the optical monitor taken through two different filters.566We thank R. Schneider and L. Fornatore for collaborative support.,We thank R. Schneider and L. Tornatore for collaborative support.567 Discussions and the stimulating environment atIV. held at OXXrcetri. Florence is kindly acknowledged.," Discussions and the stimulating environment at, held at OAArcetri, Florence is kindly acknowledged."568the sunuued spectra will have a higher sigual-to-noise ratio than in the individual spectra.,the summed spectrum will have a higher signal-to-noise ratio than in the individual spectra.569" Tf. ou the other hand. oue uses the incorrect orbital clements. then the absorption lines from the companion will be at different waveleugths in different individual spectra. and hence will be ""averaged out” in the resulting sunmnunued spectiun."," If, on the other hand, one uses the incorrect orbital elements, then the absorption lines from the companion will be at different wavelengths in different individual spectra, and hence will be “averaged out” in the resulting summed spectrum."570 Tn ourimplementation of the restframe analysis. we used a FORTRAN program that writes IRAF scripts to do the Doppler slitting aud the spectiuu suunation.," In our implementation of the restframe analysis, we used a FORTRAN program that writes IRAF scripts to do the Doppler shifting and the spectrum summation."571 The fxcor task was used to provide aoineasure of how well a restframe spectra matched the template spectrum., The fxcor task was used to provide a measure of how well a restframe spectrum matched the template spectrum.572 Iu particular. we used the peak cross-corrclation value as the measure of the goodness-of-fit of the template spectrin to the restframe spectrum," In particular, we used the peak cross-correlation value as the measure of the goodness-of-fit of the template spectrum to the restframe spectrum."573 The wavelength region used for the cross-correlation analysis was5907-6261À..1-651234...G610-G8OOA. and TOO00-7239A.," The wavelength region used for the cross-correlation analysis was, and ."574. This region excludes telluric lines. id the Πο cussion line.," This region excludes telluric lines, and the $\alpha$ emission line."575 We are interested in three orbitalelements: the period P. the semuümuplitude A» aud the time 6 πιακα velocity Ty.," We are interested in three orbitalelements: the period $P$, the semiamplitude $K_2$ and the time of maximum velocity $T_0$."576 Tlis three climensional paralucter space was searched by making two diamcusional erids of restframe spectra in the AoTy plane at several differcut values of the period P., This three dimensional parameter space was searched by making two dimensional grids of restframe spectra in the $K_2-T_0$ plane at several different values of the period $P$.577" Specifically, we used periods in- the range 0.1500 to 0.5000 days iu steps of 0.0020 davs."," Specifically, we used periods in the range 0.1500 to 0.5000 days in steps of 0.0020 days."578 The restframe spectra for cach trial period were stored iu a separate subdirectory on disk., The restframe spectra for each trial period were stored in a separate subdirectory on disk.579 For cach period. we coustructed a evid of restframe spectra iu the A5Jy plane as follows.," For each period, we constructed a grid of restframe spectra in the $K_{2}-T_{0}$ plane as follows."580 The range considered for ου was 300GOO kn 1 in steps of 2 lan |., The range considered for $K_{2}$ was 300–600 km $^{-1}$ in steps of 2 km $^{-1}$.581 For Ty the range was centered. on Π.Ο) 2.152.136.6 aud. the step size was 0.005 davs.," For $T_{0}$ the range was centered on HJD 2,452,436.6 and the step size was 0.005 days."582 The extent of the range for To was adapted to acconunodate the trial value of he period., The extent of the range for $T_{0}$ was adapted to accommodate the trial value of the period.583 Au individual restframe spectrun was generated for cach set of orbital clemeuts., An individual restframe spectrum was generated for each set of orbital elements.584 In total we eenerated 1.166.116. restirame spectra sine lis same number of scripts stored as separate tiles.," In total we generated 1,166,146 restframe spectra using this same number of scripts stored as separate files."585 Next. we did a cross-correlation analysis ou the restfriune spectra.," Next, we did a cross-correlation analysis on the restframe spectra."586 It was therefore necessary to fud a good template spectrum., It was therefore necessary to find a good template spectrum.587 To do this. we adopted a period of 0.320 days found from the xXhotometrv aud considered the restfine spectra inthe AeTy plane at this period.," To do this, we adopted a period of 0.320 days found from the photometry and considered the restframe spectra in the $K_2-T_0$ plane at this period."588 We then cross-correlated these restframe spectra against all six of he template spectra in turn and determined that the IV. star BS5568 provided the best match since it eave the strongest cross-correlation peaks., We then cross-correlated these restframe spectra against all six of the template spectra in turn and determined that the K4V star BS5568 provided the best match since it gave the strongest cross-correlation peaks.589 Next. all of the restframe spectra at all periods were crosscorrelated iu batch mode against the IWiV template.," Next, all of the restframe spectra at all periods were cross-correlated in batch mode against the K4V template."590 We made a periodoeram by parsing the fxcor log files to eet the peak cross-correlation value for cach restframe spectrum., We made a periodogram by parsing the fxcor log files to get the peak cross-correlation value for each restframe spectrum.591 The anit peak cross-correlation value within cach period subdirectory was saved. resulting in the periodoeraim shown at the bottom of veffiel..," The maximum peak cross-correlation value within each period subdirectory was saved, resulting in the periodogram shown at the bottom of \\ref{fig1}."592 The mnaxiuun cross-correlation value occurs for a trial period of P=0.320 davs. iu aerecuent with the photometric results. which hereby establishes. the orbital period.," The maximum cross-correlation value occurs for a trial period of $P=0.320$ days, in agreement with the photometric results, which thereby establishes the orbital period."593 Thus. we rule out the alias photometric period near WSs davs and again rule out the 0.212 dav spectroscopic period reported by SF2002 rofügel)).," Thus, we rule out the alias photometric period near 0.38 days and again rule out the 0.212 day spectroscopic period reported by SF2002 \\ref{fig1}) )."594 We also made a periodogram using a eniplate with spectral type 65V. Overall. the cross-correlation values were uch lower aud the oxiodoerani was πιο. noisier.," We also made a periodogram using a template with spectral type G5V. Overall, the cross-correlation values were much lower and the periodogram was much noisier."595 However. there still was a peak near P=0.320 davs. and we conclude our adopted spectroscopic period is not sensitive to the choice of the template.," However, there still was a peak near $P=0.320$ days, and we conclude our adopted spectroscopic period is not sensitive to the choice of the template."596 Adopting a period of P=0.320 davs. we performed a finer search of the A»Jy planc. with step sizes of 1 kins tin AS and 0.001 days in Zy.," Adopting a period of $P=0.320$ days, we performed a finer search of the $K_2-T_0$ plane, with step sizes of 1 km $^{-1}$ in $K_2$ and 0.001 days in $T_0$."597 The results are shown in reffie3.., The results are shown in \\ref{fig3}.598 The cross-correlation values have a fairly well-defined peak value along the Zi axis. where the mania cross-correlation is for Ti=2.152.136.600 (IIJD).," The cross-correlation values have a fairly well-defined peak value along the $T_0$ axis, where the maximum cross-correlation is for $T_0=5992,452,436.600$ (HJD)."600 On the other haud. the peak along the [Xo axis is broad with a muxiunun that occurs iu the range [30<Ae=139 km !.," On the other hand, the peak along the $K_2$ axis is broad with a maximum that occurs in the range $430 \lesssim K_2 \lesssim 439$ km $^{-1}$."601" Judeiug frou the width of the peal aud on our experience with 1705-250 (Remillard et 11996). we adopt hte=35 430lans ο,"," Judging from the width of the peak and on our experience with H1705-250 (Remillard et 1996), we adopt $K_2=435 \pm 30$ km $^{-1}$."602 The optical mass function is then The best resttrame spectrum found for the finer erid is shown in refplotspect.., The optical mass function is then The best restframe spectrum found for the finer grid is shown in \\ref{plotspect}. .603 Although few obvious K-star absorption features are apparent. the cross-correlation of this restfraie spectrum using the KIV star BS 5568 as a template docs vield a significant peak near zero velocity refplotspect)).," Although few obvious K-star absorption features are apparent, the cross-correlation of this restframe spectrum using the K4V star BS 5568 as a template does yield a significant peak near zero velocity \\ref{plotspect}) )."604 Since the spectroscopic period agrees with the, Since the spectroscopic period agrees with the60511994. 2003) are limited to lines of sight with extinctions (ly<2) well below the level ab which ice is detected.,"1994, 2003) are limited to lines of sight with extinctions $\av\la2$ ) well below the level at which ice is detected."606 The gas-phase lines within the infrared. vibration-rotation bands are narrow and suffer [rom significant telluric contamination in ground-based observations. requiring both high spectral resolving power (2=10+) and excellent signal-to-noise [or effeclive observation.," The gas-phase lines within the infrared vibration-rotation bands are narrow and suffer from significant telluric contamination in ground-based observations, requiring both high spectral resolving power $R\ga 10^4$ ) and excellent signal-to-noise for effective observation."607 Data of sullicient quality have been acquired to date only for a small number of relatively bright sources (Mitchell 11938. 1990: Shuping 22001: Rettig 22005). and these turn out to be exclusively YSOs of high or intermediate mass.," Data of sufficient quality have been acquired to date only for a small number of relatively bright sources (Mitchell 1988, 1990; Shuping 2001; Rettig 2005), and these turn out to be exclusively YSOs of high or intermediate mass."608 The further possibility to observe pure rotational transitions in absorption proves to be infeasible because ol the lack of sufficiently strong background sources at millimeter wavelengths and the need to observe multiple transitions to account for the total column density. which is spread over many J-states.," The further possibility to observe pure rotational transitions in absorption proves to be infeasible because of the lack of sufficiently strong background sources at millimeter wavelengths and the need to observe multiple transitions to account for the total column density, which is spread over many $J$ -states."609 In summary. to date. no convincing detection of interstellar gaseous CO has has been made in absorption toward any [fiekl star located behind a sienilicant column. of dense molecular-cloud. material. whieh is unfortunate. as it is the field stars that vield the most reliable information on CO depletion in quiescent regions of the clouds. whereas embedded YSOs may drive sublimation of ices in their local environment.," In summary, to date, no convincing detection of interstellar gaseous CO has has been made in absorption toward any field star located behind a significant column of dense molecular-cloud material, which is unfortunate, as it is the field stars that yield the most reliable information on CO depletion in quiescent regions of the clouds, whereas embedded YSOs may drive sublimation of ices in their local environment."610 To measure CO depletion in quiescent clouds it is necessary (o estimate gas phase column densities Irom millimeter-wave observations of CO in lines of sight toward background [eld stars. for comparison with the infrared absorption-line data for solid CO.," To measure CO depletion in quiescent clouds it is necessary to estimate gas phase column densities from millimeter-wave observations of CO in lines of sight toward background field stars, for comparison with the infrared absorption-line data for solid CO."611 This raises an obvious concern that the two tvpes οἱ observation might not sample identical regions of space. because of dillerences in effective beam size ancl the possibilitv of material behind the star. which contributes to the millimeter-wave intensity but not to the ice absorption.," This raises an obvious concern that the two types of observation might not sample identical regions of space, because of differences in effective beam size and the possibility of material behind the star, which contributes to the millimeter-wave intensity but not to the ice absorption."612 The well-known problem of saburalion in the gas-phase CO emission lines is a further concern. generally. necessitating observation of raver isotopologues such as PCO. CAO. CAO and C0 (e.g. Frerking 11932: Bensch 22001a: Larjunpaa 22004).," The well-known problem of saturation in the gas-phase CO emission lines is a further concern, generally necessitating observation of rarer isotopologues such as $^{13}$ CO, $^{18}$ O, $^{17}$ O and $^{13}{\rm C}^{18}$ O (e.g., Frerking 1982; Bensch 2001a; Harjunpää 2004)."613 This paper compares gas-phase CO colunn densities extracted from the Five College Radio Astronomy Observatory (FCRAQO) CO Mapping Survey of the Taurus molecular eloud (Goldsmith 22008: Naravanan 22008: see Section 2) with extinction data lor some 292 background field stars in the same Galactic region (Shenov 22008: Whittet 22001)., This paper compares gas-phase CO column densities extracted from the Five College Radio Astronomy Observatory (FCRAO) CO Mapping Survey of the Taurus molecular cloud (Goldsmith 2008; Narayanan 2008; see Section 2) with extinction data for some 292 background field stars in the same Galactic region (Shenoy 2008; Whittet 2001).614 The spatial resolution of the gas-phase CO observations from the FCRAO survey, The spatial resolution of the gas-phase CO observations from the FCRAO survey615of the true nature of the viscosity as being due to the maenetorotational instability Balbus&Lawley(1991). did not simply set the issue.,of the true nature of the viscosity as being due to the magnetorotational instability \cite{balb91} did not simply set the issue.616 Phe limit evele cannot be seen in 3-D simulations of this instability since the authors cannot follow the elobal evolution of the disc in à viscous timescale and the radial propagation of heating ancl cooling fronts is here. neglected., The limit cycle cannot be seen in 3-D simulations of this instability since the authors cannot follow the global evolution of the disc in a viscous timescale and the radial propagation of heating and cooling fronts is here neglected.617 However. recent. computations. indicate that radiation. pressure contributes to the viscous torque and the instability may be there Hiroseetal.(2009b).," However, recent computations indicate that radiation pressure contributes to the viscous torque and the instability may be there \cite{hirose09b}."618. The comparison of the precictions of this instability with observational data was done only in a [ον papers so far., The comparison of the predictions of this instability with observational data was done only in a few papers so far.619 In this work we propose to study. further the models of the accretion. cise instabilities in a global. picture. as well as to better support. them observationallv.," In this work we propose to study further the models of the accretion disc instabilities in a global picture, as well as to better support them observationally."620 In. case of the radiation. pressure instability. we use two models: one is based on the assumption of the viscoscous torque proportional to the total pressure and. the other is. based on the assumption of the viscosous torque proportional to geometrical mean of the gas and radiation pressure.," In case of the radiation pressure instability, we use two models: one is based on the assumption of the viscoscous torque proportional to the total pressure and the other is based on the assumption of the viscosous torque proportional to geometrical mean of the gas and radiation pressure."621 We also include the cooling term due to the outflow., We also include the cooling term due to the outflow.622 We mark the instability strips in the disk radius - accerction rate plane and we compare them with the observed. properties of the X-ray lightcurves of accreting black holes in binary systems. taken from the literature.," We mark the instability strips in the disk radius - acccretion rate plane and we compare them with the observed properties of the X-ray lightcurves of accreting black holes in binary systems, taken from the literature."623 We claim that the observational constraints for the thermal disc instability and the limit-evele behaviour is much more than the one famous example of the microquasar GRS 1915|105., We claim that the observational constraints for the thermal disc instability and the limit-cycle behaviour is much more than the one famous example of the microquasar GRS 1915+105.624 We give some examples of other sources and. discuss further need. for the detailed observational studies lor both Galactic individual X-ray sources anc ACNs., We give some examples of other sources and discuss further need for the detailed observational studies for both Galactic individual X-ray sources and AGNs.625 This article is organized as follows., This article is organized as follows.626 In Section 2. we discuss the theoretical background for the radiation pressure ancl partial hyerogen ionization instabilities., In Section \ref{sec:results} we discuss the theoretical background for the radiation pressure and partial hydrogen ionization instabilities.627 We also present the results for the extension and overlapping of the unstable regions in the aceretion dises., We also present the results for the extension and overlapping of the unstable regions in the accretion discs.628 In particular. we focus on the constraints for the accretion rate and jet cllicicney which would be adequate for the astrophysical black hole. clises o become unstable for one or both types of instabilities.," In particular, we focus on the constraints for the accretion rate and jet efficiency which would be adequate for the astrophysical black hole discs to become unstable for one or both types of instabilities."629 These results are based on the numerical codes developed »w ourselves and. discussed in detail in à series of previous works., These results are based on the numerical codes developed by ourselves and discussed in detail in a series of previous works.630 In section 3. we present the observational constraints or the dise instabilities found for a number of Galactic Mack hole binaries., In section \ref{sec:obs} we present the observational constraints for the disc instabilities found for a number of Galactic black hole binaries.631 We also discuss the supermassive black role ACGINs and some observational constraints found in the iterature., We also discuss the supermassive black hole AGNs and some observational constraints found in the literature.632 In Section 4 we give a summary and conclusions., In Section \ref{sec:diss} we give a summary and conclusions.633 The black bole accretion cise with the classical. heating terni proportional to the pressure with the a coellicient (Shakura&Sunvaey (1973))) is subject to the thermal and viscous instability when the radiation pressure dominates over the gas pressure., The black hole accretion disc with the classical heating term proportional to the pressure with the $\alpha$ coefficient \cite{ss73}) ) is subject to the thermal and viscous instability when the radiation pressure dominates over the gas pressure.634 Εις occurs in the innermost raclii of the accretion disc around a compact object (in case of a white dwarl such a region cannot be present)., This occurs in the innermost radii of the accretion disc around a compact object (in case of a white dwarf such a region cannot be present).635 Itacdiation pressure instability. of classical alpha models of Sunvaev(LO73) was noticed very carly (Lightman&Earel-lev (1974).. Pringleetal. (1974))) and i was fully analyzecl by Shakura&Sunvacy(10τ0).," Radiation pressure instability of classical alpha models of \cite{ss73} was noticed very early \cite{lightman74}, \cite{pringle74}) ) and it was fully analyzed by \cite{ss76}."636. The time evolution of the svstem ancl its stability is eoverned by the accretion rate outside the unstable region (Le. the mean aceretion rate).," The time evolution of the system and its stability is governed by the accretion rate outside the unstable region (i.e., the mean accretion rate)."637 Lf the aceretion rate is low. hen the disc remains cold and stable. with a constant low tumiinosity.," If the accretion rate is low, then the disc remains cold and stable, with a constant low luminosity."638 HE the accretion rate is large. then the whole disc comes hot and is stabilized by advection. ic. enters a slim disc solution (Abramowiezetal.CD988))).," If the accretion rate is large, then the whole disc becomes hot and is stabilized by advection, i.e. enters a slim disc solution \cite{abr88}) )."639 However. for the intermittent accretion rates. [larger than some critical value. he unstable mode activates.," However, for the intermittent accretion rates, larger than some critical value, the unstable mode activates."640 In this case the source enters a evcle of bright. hot states. separated by the cold. low uminositv states.," In this case the source enters a cycle of bright, hot states, separated by the cold, low luminosity states."641 Phe outburst amplitudes anc cluratious are sensitive to black hole mass. viscosity parameter and he mean accretion rate.," The outburst amplitudes and durations are sensitive to black hole mass, viscosity parameter and the mean accretion rate."642 Also. the heating prescription is important here.," Also, the heating prescription is important here."643 LE the viscous heating is proportional to he total pressure. the outburst amplitudes are very. large. jowever. they can be reduced if the heating is proportional ο the square root of the gas times the total pressure.," If the viscous heating is proportional to the total pressure, the outburst amplitudes are very large, however, they can be reduced if the heating is proportional to the square root of the gas times the total pressure."644 IH we assume the heating proportional only to the eas pressure. he instability disappears.," If we assume the heating proportional only to the gas pressure, the instability disappears."645 In case of the geometrical mean of the two components the parameter space of the instability is greatly reduced., In case of the geometrical mean of the two components the parameter space of the instability is greatly reduced.646 Preliminary shearine-box 391) simulations replacing the alpha viscosity with a physical (magnetic) viscosity mechanism indicated that there is no thermal runaway even when the radiation pressure is 10 times larger than the gas pressure (Lliroseοἱal. (2009a)))., Preliminary shearing-box 3D simulations replacing the alpha viscosity with a physical (magnetic) viscosity mechanism indicated that there is no thermal runaway even when the radiation pressure is 10 times larger than the gas pressure \cite{hirose09a}) ).647 However. the same authors (Lliroseetal. (2009b))) in their follow-up work alreacky sugeested the possibility of radiation pressure instability in some of their calculations. as the unstable solutions may be seen on the surface density - elfective temperature plot.," However, the same authors \cite{hirose09b}) ) in their follow-up work already suggested the possibility of radiation pressure instability in some of their calculations, as the unstable solutions may be seen on the surface density - effective temperature plot."648 The limit evele cannot be seen in those simulations since they do not follow the global evolution of the disc in à viscous timescale and the radial propagation of heating and cooling fronts is here neglected., The limit cycle cannot be seen in those simulations since they do not follow the global evolution of the disc in a viscous timescale and the radial propagation of heating and cooling fronts is here neglected.649 Full time-dependent computations of the global evolution can be only performed with a simple viscosity parameterization. and in such computations a limit-cvcle behaviour is seen. with disc alternating between the hot and cold states (c.g. etal. (2000).. Janiuketal. (2002)... Janiuk&σον (2007).. Czernyetal... (2009))).," Full time-dependent computations of the global evolution can be only performed with a simple viscosity parameterization, and in such computations a limit-cycle behaviour is seen, with disc alternating between the hot and cold states (e.g. \cite{nayak00}, \cite{janiuk02}, \cite{janiuk07}, \cite{czerny09}) )."650 Observationallv. the situation is far from clear (see e.g. the review by Donectal. (20011).," Observationally, the situation is far from clear (see e.g. the review by \cite{done07}) )."651 Several authors sugeested that the radiation. pressure instability is an attractive explanation of the regular outbursts lasting a few huncreds of seconds observed in the microquasar GRS 1915|105 (e.g. Taametal.(1997).. Deeganetal. (2009))).," Several authors suggested that the radiation pressure instability is an attractive explanation of the regular outbursts lasting a few hundreds of seconds observed in the microquasar GRS 1915+105 (e.g. \cite{taam97}, \cite{deegan09}) )."652 The radiation pressure instability is the only model which explains the absence of the direct transitions from the state C to the state D in this source., The radiation pressure instability is the only model which explains the absence of the direct transitions from the state C to the state B in this source.653 The lehteurves of some other objects also show the fluctuations in a form of a limit evele on the appropriate timescales., The lightcurves of some other objects also show the fluctuations in a form of a limit cycle on the appropriate timescales.654 An interesting example is N-ray. pulsar GRO J1744-28 (Cannizzo(1996).. Cannizzo (1998))) with periods of very high aceretion rate and low magnetic field which allows for the presence of the inner racliation-pressure dominated. part. of the disk.," An interesting example is X-ray pulsar GRO J1744-28 \cite{can96}, \cite{can97}) ) with periods of very high accretion rate and low magnetic field which allows for the presence of the inner radiation-pressure dominated part of the disk."655 This instability operates when the radiation. pressure is important. so it is expected only in high Exldington ratio Galactic sources and AGN.," This instability operates when the radiation pressure is important, so it is expected only in high Eddington ratio Galactic sources and AGN."656 However. for many sources accreting at very high rates the limit excle oscillations have not been reported.," However, for many sources accreting at very high rates the limit cycle oscillations have not been reported."657 Comparison of the Lcleington ratios of the stable sources and. those showing fast (100 - 1000, Comparison of the Eddington ratios of the stable sources and those showing fast (100 - 1000658Aeronautics and Space Administration. the National Science Foundation. the U.S. Department of Energy. the Japanese Monbukagalm and the Max Planck Society.,"Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society."659 The SDSS Web site ts http://www.sdss. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions., The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.660 The Participating Institutions| are The University of Chicago. Fermilab. the Institute for Advanced Study. the Japan Participation Group. The Johns Hopkins University. Los Alamos National Laboratory. the Max-Planck-Institute for Astronomy (MPIA). the Max-Planck-Institute for Astrophysics (MPÀ). New Mexico State University. University of Pittsburgh. Princeton University. the United States Naval Observatory. and the University of Washington.," The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."661the total emission for these trausitious of high critical density. leading to larger ceisity estimates.,"the total emission for these transitions of high critical density, leading to larger density estimates."662 However. another possibility also exists that the low-lying aud high-Ivyiug lines may actually trace dilfereut components ο ‘the shock.," However, another possibility also exists that the low-lying and high-lying lines may actually trace different components of the shock."663 Reach et ((2005) proposed that He 5(9) cau arise lareely from the dissociative part of shock. while H» $(3) is attributed almost entirely to the 1ou-dissociative shock.," Reach et (2005) proposed that $_2$ S(9) can arise largely from the dissociative part of shock, while $_2$ S(3) is attributed almost entirely to the non-dissociative shock."664 In reality. these two situations may both exist when part of the higlily-excited CO and Ho rotational lines come rom warner regious where the shock is partially clissociative aud atomic hydrogen becomes an important collision:il partner. au ellect neglected in our inodel.," In reality, these two situations may both exist when part of the highly-excited CO and $_2$ rotational lines come from warmer regions where the shock is partially dissociative and atomic hydrogen becomes an important collisional partner, an effect neglected in our model."665 The best-fit power-law index 6. whic represents the gas temperature distribution aloug the line-of-ight. is in the rauge 2.3 — 3.1 accor(ing to our fits to the IRS H» emissious.," The best-fit power-law index $b$, which represents the gas temperature distribution along the line-of-sight, is in the range 2.3 – 3.1 according to our fits to the IRS $_2$ emissions."666 If H» (0) is also considered. b is enliauced by ~ 0.2. and a trther increase of 0.3 — 0.6 is needed if the IRAC L5 µι 1 band fux or CO lines are included.," If $_2$ S(9) is also considered, $b$ is enhanced by $\sim$ 0.2, and a further increase of 0.3 – 0.6 is needed if the IRAC 4.5 $\mu$ m band flux or CO lines are included."667 The alove ellects are probably caused by the cdegereracy of the two parameters — n(H») aud b — as meuloned in Section [: au increase in the best-[i density ca be comipeusated for by a larger 6 (ie. by assuming the presence of less gas at hieli tenperatu05)., The above effects are probably caused by the degeneracy of the two parameters | $n$ $_2$ ) and $b$ | as mentioned in Section 4; an increase in the best-fit density can be compensated for by a larger $b$ (i.e. by assuming the presence of less gas at high temperatures).668 The best-lit b index for all six sources is 5aller than preclictious from a classical bow shock whose shape cau be approximated as parabolic., The best-fit $b$ index for all six sources is smaller than predictions from a classical bow shock whose shape can be approximated as parabolic.669" According to Smith Brane (1990). the elective shocX surlace area cLd with a perpendicular shock velocity Vy is proportiona to V,Ad owuch leacs 10 a power-law iidex b ~3.8 if the relationship between tle column density and shock velocity [n]0give1 by equation. B6 in NOG is adopted: NV (He paVP (NYOS)."," According to Smith Brand (1990), the effective shock surface area $A$ with a perpendicular shock velocity $V_{s}$ is proportional to $V_{s}^{-4}670dV_{s}$, which leads to a power-law index $b$ $\sim 3.8$ if the relationship between the column density and shock velocity given by equation B6 in N06 is adopted: $N$ $_2)$ $\varpropto V_{s}^{-0.75}671$ (NY08)."672 A b ineex stnaller tha1 3.8 can be caused by a AÀ which drops less steeply with velocity Vi than does a parabolic shock., A $b$ index smaller than 3.8 can be caused by a $A$ which drops less steeply with velocity $V_{s}$ than does a parabolic shock.673 This would require that the curvature of the shock Lrot be sinaller than that of a yarabola. or more probably. that there exists an acinixture of shocks with different geometries whose shapes vary [rom planar to bow.," This would require that the curvature of the shock front be smaller than that of a parabola, or more probably, that there exists an admixture of shocks with different geometries whose shapes vary from planar to bow."674" The average column deusity of the shocked Πο at T> 100 lx within the rectaugular areas marked iu Figures 1 — 6 vajes"" [romH 2x 4107>emn 2400 4x 1001 7. with the two Herbig-Haro objects the weakest sources of the total H» emissions."," The average column density of the shocked $_2$ at $T>$ 100 K within the rectangular areas marked in Figures 1 – 6 varies from $\times$ $10^{20}~$ $^{-2}$ to 4 $\times$ $10^{21}~$ $^{-2}$, with the two Herbig-Haro objects the weakest sources of the total $_2$ emissions."675 The length scale defined by .NCHa)/ (H2). which should be equal to tlie product of shock thickuess aud covering facor within the regious. is iu the range of |LOM cm — 10 ‘cm.," The length scale defined by $N$ $_2)$ $n$ $_2)$, which should be equal to the product of shock thickness and covering factor within the regions, is in the range of $10^{16}$ cm – $10^{18}$ cm."676 All sources except HH51 show .N(Ha)/n (Ho) above LO! cm., All sources except HH54 show $N$ $_2)$ $n$ $_2)$ above $10^{17}~$ cm.677 The thickness of the shocks in these regions. olXalned [rom expressious (B6) arda (B7) from NOG. should be less than 10*=T cm.," The thickness of the shocks in these regions, obtained from expressions (B6) and (B7) from N06, should be less than $10^{17}~$ cm."678 The alalysis above unplies that the covering [actor for all six sources except HH» is larger than uiitv: for W28 it is even as large as ~6., The analysis above implies that the covering factor for all six sources except HH54 is larger than unity; for W28 it is even as large as $\sim 6$.679 The high covering [actors are uot surprisnug becaise we are probably uot observing these shocks [ace-on., The high covering factors are not surprising because we are probably not observing these shocks face-on.680 The filamentary structures appeariug in part of the maps of W28 aud WIL imply the existence of incividual shock [routs seen close to edge-on. as sugeestedMOD by Reach et ((2005).," The filamentary structures appearing in part of the maps of W28 and W44 imply the existence of individual shock fronts seen close to edge-on, as suggested by Reach et (2005)."681 Actually. for most of the sources except HHT.," Actually, for most of the sources except HH7,"682telescope (GBT) and the Arecibo radio telescope.,telescope (GBT) and the Arecibo radio telescope.683 Four pulsars with no detectable radio emission are being observed with the Rossi X-ray Timing Explorer satellite (RXTE)., Four pulsars with no detectable radio emission are being observed with the Rossi X-ray Timing Explorer satellite (RXTE).684 The Urumqi Observatory (?) is using a 25 meter antenna to monitor 38 of the brighter radio pulsars.," The Urumqi Observatory \citep{WMZ01}685 is using a 25 meter antenna to monitor 38 of the brighter radio pulsars."686 The goal is to build a database of rotation parameters that will allow folding of the gamma-rays as they are accumulated over the 5 to 10 year lifetime of the LAT., The goal is to build a database of rotation parameters that will allow folding of the gamma-rays as they are accumulated over the 5 to 10 year lifetime of the LAT.687 This work is similar in spirit to what was done for CGRO (?:9.. 252)., This work is similar in spirit to what was done for CGRO \citealp{ZA94}; ;\citealp{SJ95}; \citealp{DAGM96}; \citealp{KaspiThesis}) ).688 The LAT is described by ?.., The LAT is described by \citet{WA08}.689 In brief. gamma-rays convert to electron-positron pairs in tungsten foil interleaved with layers of silicon microstrip detectors in the tracker. yielding direction information.," In brief, gamma-rays convert to electron-positron pairs in tungsten foil interleaved with layers of silicon microstrip detectors in the tracker, yielding direction information."690 The particle cascade continues in the cesium iodide crystals of the calorimeter. providing energy information.," The particle cascade continues in the cesium iodide crystals of the calorimeter, providing energy information."691 Seintillators surrounding the tracker aid rejection of the charged cosmic ray background., Scintillators surrounding the tracker aid rejection of the charged cosmic ray background.692 The seintillators are segmented to reduce the “backsplash”: a self-veto effect that reduced EGRET's sensitivity to high energy photons., The scintillators are segmented to reduce the “backsplash”: a self-veto effect that reduced EGRET's sensitivity to high energy photons.693 The LAT is à 4-by-4 array of detector “modules” covering an area of roughly 1.7 meters on a side., The LAT is a 4-by-4 array of detector “modules” covering an area of roughly $1.7$ meters on a side.694 It is sensitive to photons with energies between 20 MeV and 300 GeV. whereas EGRET's sensitivity fell off significantly above 10 GeV. After event reconstruction and background rejection. the effective area for gamma-rays above | GeV is >8000 em? at normal incidence. as compared to 1200 em for EGRET.," It is sensitive to photons with energies between $20$ MeV and $300$ GeV, whereas EGRET's sensitivity fell off significantly above 10 GeV. After event reconstruction and background rejection, the effective area for gamma-rays above 1 GeV is $> 8000$ $^2$ at normal incidence, as compared to $1200$ $^2$ for EGRET."695 The angular resolution Is. also better than EGRET's. such that source localisation for typical sources will be of order of 0.17'.," The angular resolution is also better than EGRET's, such that source localisation for typical sources will be of order of $0.1$."696. The height-to-width aspect ratio of the LAT is 0.4. for a of 2.5 sr. or nearly 20% of the sky at a given time.," The height-to-width aspect ratio of the LAT is $0.4$, for a field-of-view of $2.5$ sr, or nearly $20$ of the sky at a given time."697 Combined with the large effective area. this makes a sky survey observation strategy possible: on a given orbit. the LAT will sweep the sky 35° away from the orbital plane. covering of the sky.," Combined with the large effective area, this makes a sky survey observation strategy possible: on a given orbit, the LAT will sweep the sky $35^\circ$ away from the orbital plane, covering of the sky."698 At the end of the orbit. will rock to 35° on the other side of the orbital plane. and continue to scan.," At the end of the orbit, will rock to $35^\circ$ on the other side of the orbital plane, and continue to scan."699 Thus. the entire sky is covered with good uniformity every three hours. and no time 1s lost to earth occultation.," Thus, the entire sky is covered with good uniformity every three hours, and no time is lost to earth occultation."700 Survey mode. large effective area. and good localisation together give the LAT an overall steady point-source sensitivity 30 times better than EGRETs.," Survey mode, large effective area, and good localisation together give the LAT an overall steady point-source sensitivity 30 times better than EGRET's."701 Gamma-ray events recorded with the LAT have timestamps that derive from. a GPS clock on the satellite., Gamma-ray events recorded with the LAT have timestamps that derive from a GPS clock on the satellite.702 Ground tests using cosmic ray muons demonstrated that the LAT measures event times with precision relative to UTC significantly better than a microsecond (?).., Ground tests using cosmic ray muons demonstrated that the LAT measures event times with precision relative to UTC significantly better than a microsecond \citep{DAS08b}.703 On orbit. satellite telemetry indicates comparable accuracy.," On orbit, satellite telemetry indicates comparable accuracy."704 The contribution to the barycentered time resolution from uncertainty in the LAT’s position is negligible., The contribution to the barycentered time resolution from uncertainty in the LAT's position is negligible.705 The EGRET pulsars showed a variety of pulse profiles and emission spectra and raised as many questions às they answered (?).., The EGRET pulsars showed a variety of pulse profiles and emission spectra and raised as many questions as they answered \citep{DJT04}.706 The high-energy emission is thought to arise from basic electromagnetic interactions of highly relativistic particles. namely synchrotron emission. curvature emission and inverse Compton emission.," The high-energy emission is thought to arise from basic electromagnetic interactions of highly relativistic particles, namely synchrotron emission, curvature emission and inverse Compton emission."707 In the two main categories of models describing high-energy emission by pulsars. charged particles are accelerated along the magnetic dipole field lines by parallel electric fields.," In the two main categories of models describing high-energy emission by pulsars, charged particles are accelerated along the magnetic dipole field lines by parallel electric fields."708" The ""polar cap"" model (??) argues that the acceleration begins above the stellar magnetic pole. but can extend to the outer magnetosphere."," The “polar cap” model \citep{STU71,RS75} argues that the acceleration begins above the stellar magnetic pole, but can extend to the outer magnetosphere."709" In the ""outer gap” model (??) particles are thought to be accelerated to high energies only in the outer magnetosphere. in vacuum gaps between a null-charge surface and the light cylinder."," In the “outer gap” model \citep{CHR86a,CHR86b} particles are thought to be accelerated to high energies only in the outer magnetosphere, in vacuum gaps between a null-charge surface and the light cylinder."710 The models predict different high-energy emission features such as spectra and profiles. that LAT observations may elucidate. through a hierarchy of observables.," The models predict different high-energy emission features such as spectra and profiles, that LAT observations may elucidate, through a hierarchy of observables."711 First. the different models have very different predictions of which and how many pulsars emit gamma-rays.," First, the different models have very different predictions of which and how many pulsars emit gamma-rays."712 Along with detections of radio-quiet pulsars in. gamma-rays using blind search techniques. the LAT analysis using this timing program will constrain the ratio of radio-loud to radio-quiet pulsars.," Along with detections of radio-quiet pulsars in gamma-rays using blind search techniques, the LAT analysis using this timing program will constrain the ratio of radio-loud to radio-quiet pulsars."713 This ratio is different for the two emission models. with outer gap models predicting a much lower ratio (??)..," This ratio is different for the two emission models, with outer gap models predicting a much lower ratio \citep{HGG07,PG04}."714 Reliable flux upper limits in the absence of gamma-ray pulsations are useful in this context (?) and also require good timing solutions., Reliable flux upper limits in the absence of gamma-ray pulsations are useful in this context \citep{NAB96} and also require good timing solutions.715 The second observable 1s the emission profile., The second observable is the emission profile.716 Its shape. as the beam sweeps the Earth. provides a cross-section of the regions in the pulsar magnetosphere where the emission originates.," Its shape, as the beam sweeps the Earth, provides a cross-section of the regions in the pulsar magnetosphere where the emission originates."717 Coupled with radio intensity and polarization profile studies. as well as absolute phase. the gamma-ray light curve provides information on the emission geometry. which differs significantly from one model to another (??)..," Coupled with radio intensity and polarization profile studies, as well as absolute phase, the gamma-ray light curve provides information on the emission geometry, which differs significantly from one model to another \citep{GOB02,JCRWR94}."718 The EGRET pulsars typically have two peaks. with the first one slightly offset in phase relative to the single radio peak.," The EGRET pulsars typically have two peaks, with the first one slightly offset in phase relative to the single radio peak."719 Although the pulsar breaks this trend. LAT observations will study the prevalence of this behaviour as a function of pulsar age or other parameters.," Although the pulsar breaks this trend, LAT observations will study the prevalence of this behaviour as a function of pulsar age or other parameters."720 Pulsar detections and emission profiles can only be achieved through solid knowledge of the pulsar’s rotation and good absolute time precision., Pulsar detections and emission profiles can only be achieved through solid knowledge of the pulsar's rotation and good absolute time precision.721 The timing precision will allow finely binned profiles over many years even for millisecond pulsars., The timing precision will allow finely binned profiles over many years even for millisecond pulsars.722 The large energy range covered by the LAT will enable measurements of pulsar spectral cut-offs., The large energy range covered by the LAT will enable measurements of pulsar spectral cut-offs.723 Although EGRET observed high-energy cut-offs in pulsar spectra around a few GeV. it did not have the sensitivity to measure the exact energy or shape of the turnovers.," Although EGRET observed high-energy cut-offs in pulsar spectra around a few GeV, it did not have the sensitivity to measure the exact energy or shape of the turnovers."724 For instance. the LAT should provide a determination of the pulsar’s spectral cut-off energy. known only to be less than a few tens of GeV (?2).. where EGRET lost sensitivity due to the backsplash effect.," For instance, the LAT should provide a determination of the pulsar's spectral cut-off energy, known only to be less than a few tens of GeV \citep{MT08,CELESTE02}, , where EGRET lost sensitivity due to the backsplash effect."725 The on-axis LAT energy resolution ts better than 15% above 100 MeV and is better than 10% in the range between roughly 500 MeV and 50 GeV. and improves somewhat off-axis.," The on-axis LAT energy resolution is better than $15$ above 100 MeV and is better than $10$ in the range between roughly $500$ MeV and $50$ GeV, and improves somewhat off-axis."726 The LAT should quickly measure the shape of the pulsar spectral cut-off expected to be around 4 GeV. a powerful discriminator between polar cap and outer gap models and a potential diagnosticof high-energy emission altitude (2).. Finally. a subset of the pulsars detected by the LAT will have sufficient photon," The LAT should quickly measure the shape of the pulsar spectral cut-off expected to be around 4 GeV, a powerful discriminator between polar cap and outer gap models and a potential diagnosticof high-energy emission altitude \citep{AH07}.. Finally, a subset of the pulsars detected by the LAT will have sufficient photon"727gas and dark matter radial profiles) and the magnitude of the ram pressure.,gas and dark matter radial profiles) and the magnitude of the ram pressure.728" In contrast to the uniform medium runs, however, the ram pressure is not constant with time."," In contrast to the uniform medium runs, however, the ram pressure is not constant with time."729" Using the orbit from the simulations, along with the density profile of the group, Prax(f) is calculated and passed to the analytic model."," Using the orbit from the simulations, along with the density profile of the group, $P_{\rm ram}(t)$ is calculated and passed to the analytic model."730 The analytic model can then predict M(t) once the values of α and 8 have been selected., The analytic model can then predict $M(t)$ once the values of $\alpha$ and $\beta$ have been selected.731 'The mass loss curves for the default 2-system run are plotted in the top panel of 66., The mass loss curves for the default 2-system run are plotted in the top panel of 6.732" Overall, the simple analytic model with a= 2, 0.5«80.7 (shown is B= 2/3) and tram=$tsouna reproduces the mass loss seen in the default 2-system run very well."," Overall, the simple analytic model with $\alpha = 2$ , $0.5 < \beta < 0.7$ (shown is $\beta = 7332/3$ ) and $t_{\rm ram} = \beta 734t_{\rm sound}$ reproduces the mass loss seen in the default 2-system run very well."735" For example, both the simulations and the model show evidence for near convergence in M(t) at t=1.5 Gyr, which corresponds to the (first) pericentric passage and, therefore, to the maximum ram pressure which the galaxy experiences along its orbit (see the bottom panel of 6)."," For example, both the simulations and the model show evidence for near convergence in $M(t)$ at $t 736\ga 1.5$ Gyr, which corresponds to the (first) pericentric passage and, therefore, to the maximum ram pressure which the galaxy experiences along its orbit (see the bottom panel of 6)."737 The analytic model slightly underestimates the mass loss seen in the simulations at early times., The analytic model slightly underestimates the mass loss seen in the simulations at early times.738 This is a result of the fact that the hot halo of the galaxy is initially slightly over-pressurised with respect to the surrounding hot halo of the group. (, This is a result of the fact that the hot halo of the galaxy is initially slightly over-pressurised with respect to the surrounding hot halo of the group. (739Note that this was not the case for the uniform medium simulations plotted in 44.),Note that this was not the case for the uniform medium simulations plotted in 4.)740" This leads to some expansion of the outer gas which, in turn, makes it more susceptible to stripping."," This leads to some expansion of the outer gas which, in turn, makes it more susceptible to stripping."741" Since this effect is in general small and is an artifact of our idealised setup, we do not attempt to model it."," Since this effect is in general small and is an artifact of our idealised setup, we do not attempt to model it."742" While the analytic model with a time delay factor matches the simulations well, an instantaneous stripping model with a~4 (represented by the second dotted curve from the bottom) also performs well."," While the analytic model with a time delay factor matches the simulations well, an instantaneous stripping model with $\alpha 743\approx 4$ (represented by the second dotted curve from the bottom) also performs well."744" However, even if the agreement is reasonable, this model is without physical justification and should not be expected to apply in situations that differ significantly from those of the default run."," However, even if the agreement is reasonable, this model is without physical justification and should not be expected to apply in situations that differ significantly from those of the default run."745" Indeed, this is indicated by the results presented later in the paper (c.f."," Indeed, this is indicated by the results presented later in the paper (c.f."746 88)., 8).747" We also note that a significant fraction of the dark matter halo is also stripped, particularly near the first pericentric passage."," We also note that a significant fraction of the dark matter halo is also stripped, particularly near the first pericentric passage."748 This is not unexpected andis due to, This is not unexpected andis due to749lt is well known from the recent observations that the standard Bie Bang model of cosmology fails to describe the present accelerating phase of the universe.,It is well known from the recent observations that the standard Big Bang model of cosmology fails to describe the present accelerating phase of the universe.750 The nioclel is also pleagued by a time like singularity in the past., The model is also pleagued by a time like singularity in the past.751 Accelerating phase of the universe can. however. be incorporated in a number of wavs.," Accelerating phase of the universe can, however, be incorporated in a number of ways."752 A number of models e.g... models with mocified theory of gravity. (Sotiriou.2007).. models. with unusual matters like Chaplvein. gas (Bento.Bertolami&Sen2002:Bilic.Tupper&Viollier 2001).. scalar and tachyon ficlds (Lyth2003) are taken into account to accommodate present phase of acceleration.," A number of models e.g., models with modified theory of gravity \citep{b19}, models with unusual matters like Chaplygin gas \citep{b3,b4}, scalar and tachyon fields \citep{b18} are taken into account to accommodate present phase of acceleration."753 There are other models based on mostly non-equilibrium thermocnamics and Boltzmann formulation which do not require. dark. energy. (Zimcdahl 2008)..," There are other models based on mostly non-equilibrium thermodynamics and Boltzmann formulation which do not require dark energy \citep{b15,b16,b17}."754 Only very. recently some other models. appeared in the literature which cliseusses cold. dark matter (CDM) and CDAL interactions as alternative to the ACDAL mocel (Lima.Jesus&Oliveira2009:BasilakosPlionis2009).," Only very recently some other models appeared in the literature which discusses cold dark matter (CDM) and CDM interactions as alternative to the $\Lambda$ CDM model \citep {n1,n2}."755. Llowever. exploring singularity free cosmological mioclels is an interesting area in cosmology and Emerecnt Universe scenario (EW) is one of the well known choices.," However, exploring singularity free cosmological models is an interesting area in cosmology and Emergent Universe scenario (EU) is one of the well known choices."756 A number of literature appeared. which discussed EU mocel as it was free from initial singularitv and the size of the universe was large enough so that quantum gravity cllects were not important (Llarrison1967:Ellis&Alaartens2004)..," A number of literature appeared which discussed EU model as it was free from initial singularity and the size of the universe was large enough so that quantum gravity effects were not important \citep{b8,b7}."757 These models evolve from a static phase in the infinite past into an inflationary phase., These models evolve from a static phase in the infinite past into an inflationary phase.758 Ehe idea is in conformity with Lemaitre-Eelelington concepts from. early. days. of modern cosmology., The idea is in conformity with Lemaitre-Eddington concepts from early days of modern cosmology.759 M developed in a consistent manner an emergent universe model is capable of solving some of the well known conceptual problems not understood in the Big-Dang model., If developed in a consistent manner an emergent universe model is capable of solving some of the well known conceptual problems not understood in the Big-Bang model.760 A model of an ever-existing universe. which eventually. enters into the standard. Big Dang epoch at some stage and consistent with features known to us today. is worth considering.," A model of an ever-existing universe, which eventually enters into the standard Big Bang epoch at some stage and consistent with features known to us today is worth considering."761 Recently an interesting EWU model has been proposed by Mukherjeeetal.(2006) which requires some exotic matter in addition to normal matter as cosmic Iuid., Recently an interesting EU model has been proposed by \citet{b12} which requires some exotic matter in addition to normal matter as cosmic fluid.762 The model has been explored in a [lat universe. as such universe is supported by recent observations., The model has been explored in a flat universe as such universe is supported by recent observations.763 SubsequentIv the EU model was taken up to examine the suitability of, Subsequently the EU model was taken up to examine the suitability of764for. and it did not address the under-sampling of the PSE as Cdrizzle does). this output frame had. the advantage that the noise properties were preserved and adjacent. pixels were uncorrelated.,"for, and it did not address the under-sampling of the PSF as “drizzle"" does), this output frame had the advantage that the noise properties were preserved and adjacent pixels were uncorrelated."765 We measured. the standard: deviation of the counts in. blank areas of sky in this shift-and-acd mosaic. and we verified. that the noise (normalized. per unit time) decreased as the square root. of the number of frames combined.," We measured the standard deviation of the counts in blank areas of sky in this shift-and-add mosaic, and we verified that the noise (normalized per unit time) decreased as the square root of the number of frames combined."766" The limiting magnitudes found using these uncorrelated ""true-nolse [rames are in good agreement with the STSel HST/WECS Exposure Time Calculator (PC) Vable 2) presents our 26 limits in a 0 6-diameter aperture. with the aperture correction applied."," The limiting magnitudes found using these uncorrelated “true-noise frames"" are in good agreement with the STScI /WFC3 Exposure Time Calculator (ETC) – Table \ref{tab:obssum} presents our $2\,\sigma$ limits in a $0\farcs6$ -diameter aperture, with the aperture correction applied."767 We also produced a noise model based on the detector gain. reaclout noise and Poisson counts of the measured. background. (including the instrument dark current). and. verified that our sensitivity was well within LO per cent of the expected. noise.," We also produced a noise model based on the detector gain, readout noise and Poisson counts of the measured background (including the instrument dark current), and verified that our sensitivity was well within 10 per cent of the expected noise."768" Finally. we measure the correlated noise (the standard deviation of he background counts) in the drizzled image mosaics which we use for our source detection and. photometry. ancl use he relations in equation ALS of Casertano et ((2000) to introduce a correction factor which depends on the output pixel seale ancl the size of the ""droplet in the drizzling oocedure (7pixfrac)."," Finally, we measure the correlated noise (the standard deviation of the background counts) in the drizzled image mosaics which we use for our source detection and photometry, and use the relations in equation A13 of Casertano et (2000) to introduce a correction factor which depends on the output pixel scale and the size of the “droplet"" in the drizzling procedure (“pixfrac"")."769 We generally found good agreement (at the mmasg level) with our sensitivity measurements using the true-noise frames. except for the LIUDE data where he corrected. clrizzle noise underestimated the truce noise o» mmaeg. perhaps because of the laree number of rames combined with small sub-pixel shifts.," We generally found good agreement (at the mag level) with our sensitivity measurements using the true-noise frames, except for the HUDF data where the corrected drizzle noise underestimated the true noise by mag, perhaps because of the large number of frames combined with small sub-pixel shifts."770 We adopted the sensitivity measurements from the truc-noise frame. having checked that consistent results were produced by the LPC. he noise model. and the noise in the drizzle frame corrected or pixel correlations.," We adopted the sensitivity measurements from the true-noise frame, having checked that consistent results were produced by the ETC, the noise model, and the noise in the drizzle frame corrected for pixel correlations."771 Our measured noise in the LIUDE is in good agreement. with Bouwens et ((2010a). but we note that MeLure ct ((2010) appear to be z0.3 mmag ess sensitive (although we note that their 5o magnitude ⋠⋠⋠ ∕∕⋠ ↓⊔↓↓∣↓⊔⋜↧∪⋅≟⊣⊔⋜⋯↓⋖⋅⊓⋅↓⋅⋜↧↓≻⋖⋅↓⋅⋯↓⋅≺⋅⋜↧↓≻↓≻⋖⋅⋜⊔⋅⊳∖⊔∪↿∣∪⇂↥⋜↧∖⇁≺⋅∣⋈⋅≺⋅⊔ Corrected to total magnitudes with an aperture correction. unlike in Bouwens et 220104).," Our measured noise in the HUDF is in good agreement with Bouwens et (2010a), but we note that McLure et (2010) appear to be $\approx 0.3$ mag less sensitive (although we note that their $5\,\sigma$ magnitude limit in a $0\farcs4$ -diameter aperture appears not to have been corrected to total magnitudes with an aperture correction, unlike in Bouwens et 2010a)."772 The optical41S ACS imaging comes from the Ultra Deep Field (Beckwith et 22006). ancl we used the publicl-available 0.7.2 reductions of Uanking field. UDE-D12 provided bv the UDEOS team (Oesch et 22007).," The optical ACS imaging comes from the Ultra Deep Field (Beckwith et 2006), and we used the publicly-available $v,i,z$ reductions of flanking field UDF-P12 provided by the UDF05 team (Oesch et 2007)."773 We reduced the 0.7.2 ACS data for UDE-D34 from the5T archive. using MULTVIDIIZZLIS to combine a large subset ofthe data comprising blocks of data taken close in time with similar telescope roll angles. again using an output 0703 pixel scale.," We reduced the $v,i,z$ ACS data for UDF-P34 from the archive, using MULTIDRIZZLE to combine a large subset of the data comprising blocks of data taken close in time with similar telescope roll angles, again using an output $0\farcs03$ pixel scale."774 Ehese subsets of drizzled images were then registered and combined withimcombine., These subsets of drizzled images were then registered and combined with.775 Our combined images were kksec in e-band. kksec in band and kksec in z-band.," Our combined images were ksec in $v$ -band, ksec in $i$ -band and ksec in $z$ -band."776 All the ACS images were then block-averaged 2.2 and registered with our drizzled WECS frames., All the ACS images were then block-averaged $2\times 2$ and registered with our drizzled WFC3 frames.777 ‘Lo perform the candidate selection we used the SIZxtractor photometry package (Bertin Arnouts 1996). version 2.5.0.," To perform the candidate selection we used the SExtractor photometry package (Bertin Arnouts 1996), version 2.5.0."778 Since we are searching for 3-drops (objects clearly detected in the WECS3 J-band but with minimal flux in the Y-band and ACS images). fixed circular apertures 076 in ciameter were ‘trained in the Z-image. and running Slxtractor in dual-image mode those apertures were used to measure the Hux in the same locations in the Y-band image.," Since we are searching for $Y$ -drops (objects clearly detected in the WFC3 $J$ -band but with minimal flux in the $Y$ -band and ACS images), fixed circular apertures $0\farcs6$ in diameter were `trained' in the $J$ -image, and running SExtractor in dual-image mode those apertures were used to measure the flux in the same locations in the $Y$ -band image."779 The same procedure was repeated. between Z-band image and. all the other ACS and. WECS3 images with cdillerent. filters., The same procedure was repeated between $J$ -band image and all the other ACS and WFC3 images with different filters.780 For each waveband we used a weight image derived. from. the exposure map., For each waveband we used a weight image derived from the exposure map.781 To identify the objects we set the SIExtractor xwameters to have a lower limit of 5 contiguous pixels above a threshold of 20 per pixel (data were drizzled to a scale of Y'00G pixel ly, To identify the objects we set the SExtractor parameters to have a lower limit of 5 contiguous pixels above a threshold of $2\sigma$ per pixel (data were drizzled to a scale of 06 $^{-1}$ ).782 We corrected. the aperture magnitudes to approximate total magnitudes with the aperture correction appropriate for that filter., We corrected the aperture magnitudes to approximate total magnitudes with the aperture correction appropriate for that filter.783 With this cut we were able o detect all significant sources. along with some spurious detections just above the noise limit or due to diffraction spikes from stars.," With this cut we were able to detect all significant sources, along with some spurious detections just above the noise limit or due to diffraction spikes from stars."784 We also impose a Ga limit on the J- magnitude for all fields. with the Jiz magnitude limit isted in the last column of Table 1.," We also impose a $\sigma$ limit on the $J$ -band magnitude for all fields, with the $J_{AB}$ magnitude limit listed in the last column of Table \ref{tab:exptimes}."785" Table presents our photometry of Y -drops from SlZxtractor. where we have corrected the magnitude errors returned by SExtractor for he elfects of correlated noise in the drizzled images. using our “true noise frames"" to determine the sealing factor (tvpically SIExtractor underestimated the magnitude errors w à [factor of zz1.5 for pixfrac-0.6 used in most of our data. ancl a factor of zz2.6 for pixfracz1.0 as used in the ERS and he df-bane of P12)."," Table \ref{tab:objects} presents our photometry of $Y$ -drops from SExtractor, where we have corrected the magnitude errors returned by SExtractor for the effects of correlated noise in the drizzled images, using our “true noise frames"" to determine the scaling factor (typically SExtractor underestimated the magnitude errors by a factor of $\approx1.5$ for pixfrac=0.6 used in most of our data, and a factor of $\approx2.6$ for pixfrac=1.0 as used in the ERS and the $H$ -band of P12)."786 Iclentification of candidates is achieved using the Lyman break technique (e.g. Steidel ct al., Identification of candidates is achieved using the Lyman break technique (e.g. Steidel et al.787 1996). where a large colour decrement is observed. between filters either side of Lyman-a in the rest-frame of the galaxy.," 1996), where a large colour decrement is observed between filters either side of $\alpha$ in the rest-frame of the galaxy."788 At 2> 6. the Hux decrement comes principally from the larec integrated optical depth of the intervening absorbers (the. Lyman-a forest).," At $z>6$ , the flux decrement comes principally from the large integrated optical depth of the intervening absorbers (the $\alpha$ forest)."789" AtaeS9 the location of the Lyman-a break is redshiftecd to ~Ldyan the WEC3 Yossuos, ancl Joss are suitably located such that a 7.6<z«0δ star forming galaxywill experience a significant Dux clecrement between", At $z\approx 8-9$ the location of the $\alpha$ break is redshifted to $\sim 1.1\mu m$ – the WFC3 $Y_{105w/098m}$ and $J_{125w}$ are suitably located such that a $7.6<z<9.8$ star forming galaxywill experience a significant flux decrement between790If the right-hand side of (he magnetic induction equation (14)) is negligible. the field is well coupled to the whole fIuid.,"If the right-hand side of the magnetic induction equation \ref{ebmag}) ) is negligible, the field is well coupled to the whole fluid."791 As a basis for small-perturbation analvsis. we consider an axisvmmetrie evlindrieal backeround (hat is contracting in a quasi-lydrostatic balance between the sell-gravity aud the repulsive forces (i.e.. thermal and magnetic pressure forces).," As a basis for small-perturbation analysis, we consider an axisymmetric cylindrical background that is contracting in a quasi-hydrostatic balance between the self-gravity and the repulsive forces (i.e., thermal and magnetic pressure forces)."792 The backeround quantities will be denoted with the subscript ο).," The background quantities will be denoted with the subscript ""0""."793 We consider a similarity solution for the bulk unperturbed {nicl so that its velocity field. depends linearly on the axial distance r as vi=dsdlr where s; is a non-dimensional parameter which presents axial contraction in the co-moving Lagrangian coordinate (5< 0).," We consider a similarity solution for the bulk unperturbed fluid so that its velocity field depends linearly on the axial distance $r$ as $\mathbf{v}_{(\mathbf{r},t)} = \frac{ds/dt}{s} \mathbf{r}$, where $s_{(t)}$ is a non-dimensional parameter which presents axial contraction in the co-moving Lagrangian coordinate $\dot{s}<0$ )."794 The basic equations (8))-(14)) lead to where p.. 7. and D. are respectively initial central density. temperature. and magnetic field at the inner radius rj. (>=) is only a function of radius. and so; follows the equation ∖∖⊽∐≼↲↕⋅≼↲∕∣∣↕⊳∖⇁≀↧↴∐∐∐↲−⊳∖⇁≺∢≀↧↴↥≼↲∐⋅≼↲≼↲↕↽≻≀↧↴↕⋅≀↧↴∐∐↲∥↲," The basic equations \ref{drift}) \ref{ebmag}) ) lead to where $\rho_c$, $T_c$, and $B_c$ are respectively initial central density, temperature, and magnetic field at the inner radius $r_{in}$, $f(\frac{r}{r_{in}})$ is only a function of radius, and $s_{(t)}$ follows the equation where $t_0$ is a time-scale free parameter $t \leq t_0$ )."795↕⋅∩≤∕∣∣⇄⋝⋅↴⊺↥∐↲≼↲≺⇂∏≀↧↴∐∪∐⊳∖⇁⊔↱≻⇄⋝↕⋝−⊔⊤⇄⋝⇄⋝≀↧↴↕⋅≼↲≀↧↴≺∢≺∢∏↕⋅≀↧↴∥↲ ↓⋡∪↕⋅∣⋮∕∖∕∖∣⋮∣⋅∣∣⋅≀↧↴↕∐⇂∖∖↽≼↲≀↧↪∖⊽⊳∖⇁∏∐∐↲⊔∐↲≺⇂≼↲∐⋝∖⊽∐⋡∖↽↥⊳∖⊽∐∪∐↓∪≸↽↔↴≼↲∐∪∏⊳," The equations \ref{backgden}) \ref{backgmag}) ) are accurate for $r >796r_{in}$, and we assume the density is homogenous for $r \leq797r_{in}$ ."798∖⇁↓⋟∪↕⋅∣⋮≤∣⋮∣⋅∣∣⋅⊟≻↕⋅≼⇂≼↲↥≀↧∐⊳∖⇁⋅⊔∐↲↕⋅≼↲≀↧↴≼⇂≼↲↕⋅ is referred (o the appendix.," For details, the reader is referred to the appendix."799 At the end of this contraction process in which / is equal to fy. the minimum value of si; takes place. thus the maximum effect of this contraction is an approximately tenfold over density in its profile.," At the end of this contraction process in which $t$ is equal to $t_0$, the minimum value of $s_{(t)}$ takes place, thus the maximum effect of this contraction is an approximately tenfold over density in its profile."800 For obtaininge a linearized svstem of equations. we split each variable into unperturbed and perturbed components. the latter is indicated by subscript 1.," For obtaining a linearized system of equations, we split each variable into unperturbed and perturbed components, the latter is indicated by subscript ""1""."801 Then. we carry out a spatial Fourier analvsis withcomponents proportional to exp(/kr) where f is Che component," Then, we carry out a spatial Fourier analysis withcomponents proportional to $\exp(ikr)$ where $k$ is the component"802Let us assume that the jet plasma is movine with velocity ο towards the exterual medium.,Let us assume that the jet plasma is moving with velocity $\beta_w$ towards the external medium.803" We found (Lyutikov. subinitted) exact self-similar solution of relativistic Riemann problem for the expansion of cold plasma with density. po aud Bo (1naguetizatiou parameter o=Bg/ pu: iis normalized by V/Ix). moving initially with velocity c. towards the vacuum interface where the Doppler factors 6,=YL+3,)/(1—8,) are delined in terms of the plasma velocity 3. local velocity 24. sell-5imilar parameters=z/f. initial wind velocity μαμα the velocity in the undisturbed plasma 244= Yo/(lt+o)."," We found (Lyutikov, submitted) exact self-similar solution of relativistic Riemann problem for the expansion of cold plasma with density $\rho_{0}$ and $B_0$ (magnetization parameter $\sigma=B_0^2/\rho_{0}$ ; is normalized by $\sqrt{4 \pi} $), moving initially with velocity $v_w$ towards the vacuum interface = _w _A where the Doppler factors $\delta _a= \sqrt{(1+\beta_a)/ ( 1-\beta_a)} $ are defined in terms of the plasma velocity $\beta$, local velocity $ \beta_{A}$, self-similar parameter $\eta= z/t$, initial wind velocity $\beta_w$ and the velocity in the undisturbed plasma $ \beta_{A,0}= \sqrt{ \sigma/(1+\sigma)}$ ."804 These equations give the velocity 3. deusity p=UFpolo (U4=Dfve— 31)and properfiekl.. B.—(p/py)By as a function of the self-sinilar variable 7=z/! (expansion of plasina starts at /=0.20 and proceeds into positive direction z> 0).," These equations give the velocity $\beta$ , density $\rho = U_A^2 \rho_0/\sigma$ $U_A= \beta_{A} /\sqrt{1-\beta_{A} ^2}$ )and proper, $B =(\rho/\rho_0) B_0$ as a function of the self-similar variable $\eta= z/t$ (expansion of plasma starts at $t=0, z=0$ and proceeds into positive direction $z>0$ )."805 We stress that these solutions are exact. no assumptions about the value ol the parameter σ and velocity ty were made.," We stress that these solutions are exact, no assumptions about the value of the parameter $\sigma$ and velocity $v_w$ were made."806 Particularly simple relatious are obtained for plasiua initially at rest expaucliug into vacuum Oyc0.04=) (Lyutikov. submitted).," Particularly simple relations are obtained for plasma initially at rest expanding into vacuum $\beta_w=0, \, \delta _\beta=1$ (Lyutikov, submitted)."807 The flow accelerates from rest towards the vacuum interlace., The flow accelerates from rest towards the vacuum interface.808 The bulk of the How is moving with Lorentz factor 561705., The bulk of the flow is moving with Lorentz factor $\gamma' \sim \sigma^{1/3} $.809 The flow becomes supersonic at jj=0. at which point +!=(0/2)!.," The flow becomes supersonic at $\eta =0$, at which point $\gamma'= (\sigma/2)^{1/3}$."810" The vacuum interface moves with Lorentz [actor 2/,.=14-26.", The vacuum interface moves with Lorentz factor $\gamma_{vac} '= 1+2 \sigma$.811" In the observer [rame the vacuum interface is moving withà,=0449,,. which in the limit o.>.>>1 this gives"," In the observer frame the vacuum interface is moving with$\delta _\eta= \delta_{A,0}^{2} \delta_w$, which in the limit $\sigma, \, \gamma_w \gg1$ this gives = 4 _w"812"AGN, Lx(2—10keV)=104!10**ergs~' and do not apply to QSOs with luminosities Lx>1044ergs~!, as the volumes of all three samples used here are too small to identify large numbers of those sources.","AGN, $L_X \rm813(2-10\,keV)=10^{41}-10^{44} \, erg \, s^{-1}$ and do not apply to QSOs with luminosities $\rm L_X>10^{44} \rm \, erg \, s^{-1}$, as the volumes of all three samples used here are too small to identify large numbers of those sources."814" Scattered AGN light may also contaminate the observed colours of galaxies, even in the absence of broad optical emission lines or a dominant nuclear point sources."," Scattered AGN light may also contaminate the observed colours of galaxies, even in the absence of broad optical emission lines or a dominant nuclear point sources."815" Nevertheless, the work of suggests that this effect is small."," Nevertheless, the work of suggests that this effect is small."816 Reddening of the intrinsic galaxy colours by dust is also expected to affect our conclusions., Reddening of the intrinsic galaxy colours by dust is also expected to affect our conclusions.817" If there is little or no change with redshift in the relative importance of SMBH fueling modes, how is the strong evolution of the AGN population explained?"," If there is little or no change with redshift in the relative importance of SMBH fueling modes, how is the strong evolution of the AGN population explained?"818 An important piece of evidence for addressing this question is the nearly constant fraction of ray AGN relative to the OLF of galaxies as a function of redshift., An important piece of evidence for addressing this question is the nearly constant fraction of X-ray AGN relative to the OLF of galaxies as a function of redshift.819" This suggest that the decline of the accretion power of the Universe since z&1 is linked, at least to the first approximation, to the evolution of the galaxy OLF, which in turn is the result of secular evolution processes and the decrease with time of the star-formation rate density of the Universe."," This suggest that the decline of the accretion power of the Universe since $z\approx1$ is linked, at least to the first approximation, to the evolution of the galaxy OLF, which in turn is the result of secular evolution processes and the decrease with time of the star-formation rate density of the Universe."820 Our results are therefore consistent with a fixed accretion mode (or combination of accretion modes) superimposed on an evolving galaxy population., Our results are therefore consistent with a fixed accretion mode (or combination of accretion modes) superimposed on an evolving galaxy population.821 The stellar mass function of AGN relative to galaxies provides complementary information on the evolution of active SMBHs., The stellar mass function of AGN relative to galaxies provides complementary information on the evolution of active SMBHs.822" It isfound that X-ray AGN hosts out to zzz0.8 span a wide range of stellar masses, although their fraction increases with stellar mass in agreement with previous studies(?????)."," It isfound that X-ray AGN hosts out to $z\approx0.8$ span a wide range of stellar masses, although their fraction increases with stellar mass in agreement with previous studies."823". This trend is more pronounced at lower redshift, where the distribution of the fraction of galaxies with X-ray AGN increases nearly monotonically with stellar mass."," This trend is more pronounced at lower redshift, where the distribution of the fraction of galaxies with X-ray AGN increases nearly monotonically with stellar mass."824" At higher redshifts however, the distribution levels off at intermediate masses, 71015Mo."," At higher redshifts however, the distribution levels off at intermediate masses, $\rm \ga 10^{10}\,M_{\odot}$."825" Similar trends are reported by and?,, although their X-ray AGN samples are selected at brighter luminosities, Lx>10*?ergs~*, compared to our limit of Lx(2—10keV)>10?!ergs !."," Similar trends are reported by and, although their X-ray AGN samples are selected at brighter luminosities, $L_X> \rm 10^{42}826\, erg \, s^{-1}$, compared to our limit of $L_X \rm (2-10\,keV) > \rm82710^{41} \, erg \, s^{-1}$ ."828 It is also found that the fraction, It is also found that the fraction829We have presented deep and wide-field photometry for Trumpler 20. a rich open star cluster. heavily contaminated by field stars. which lies inside the solar ring ancl in the infer-arm region between Carina and Scutum-Cryrux.,"We have presented deep and wide-field photometry for Trumpler 20, a rich open star cluster, heavily contaminated by field stars, which lies inside the solar ring and in the inter-arm region between Carina and Scutum-Crux."830 We have exploited our dataset. aiming to improve our knowledge of the cluster basic parameters., We have exploited our dataset aiming to improve our knowledge of the cluster basic parameters.831" ILaving repeatedly stressed the crucial role in the interpretation plaved by high contamination due to field stars. we conclude that Trumnpler 20 has an age of 1.1$2:0.2 ντ, making it a twin of the better-known open cluster NGC As anticipated in the Introduction. Galactic open clusters are ideal laboratories to test theories of stellar evolution. ancl to probe Galactic structure."," Having repeatedly stressed the crucial role in the interpretation played by high contamination due to field stars, we conclude that Trumpler 20 has an age of $1.4\pm0.2$ Gyr, making it a twin of the better-known open cluster NGC As anticipated in the Introduction, Galactic open clusters are ideal laboratories to test theories of stellar evolution, and to probe Galactic structure."832 Trumpler 20 appears to be quite a promising confirmation of On the stellar evolution sile. we have shown that Trumnpler 20 [alls in the age range where (he chump of He burning stus exhibits a peculiar morphology. most possibly due {ο mass-loss variation during the RGB evolutionary phase.," Trumpler 20 appears to be quite a promising confirmation of On the stellar evolution side, we have shown that Trumpler 20 falls in the age range where the clump of He burning stars exhibits a peculiar morphology, most possibly due to mass-loss variation during the RGB evolutionary phase."833 Other clusters of this age. like NGC 2660. NGC 2204 and NGC 7789. are known to have a chump with the same On the Galactic structure side. we position Trumpler 20 in the inter-arm region between the Carina and Seutium-Crux arms.," Other clusters of this age, like NGC 2660, NGC 2204 and NGC 7789, are known to have a clump with the same On the Galactic structure side, we position Trumpler 20 in the inter-arm region between the Carina and Scutum-Crux arms."834 We remind the reader that not many clusters of this age are present in (he inner disk. possibly because of environmental effects. which prevents survival of open clusters lor a long time (Carraro et al.," We remind the reader that not many clusters of this age are present in the inner disk, possibly because of environmental effects, which prevents survival of open clusters for a long time (Carraro et al."835As before. enlarging the parameter space to make more complicated assumptions about the stellar distribution function is not warranted given the other simplifications we adopt here.,"As before, enlarging the parameter space to make more complicated assumptions about the stellar distribution function is not warranted given the other simplifications we adopt here."836 In any case. the results are not likely to be stronely allected unless the stellar distribution function is extremely anisotropic.," In any case, the results are not likely to be strongly affected unless the stellar distribution function is extremely anisotropic."837 For our disc galaxy. we take the parameters use bv Sofue (1996) for the Alilky Way which assumes: (i) a spherically svmmetric (@= 0) nuclear star cluster of mass 510°M... and scale b=120pe (i) a spherically svnumetric bulge of mass 102AL. and scale 6=750pe: (il) an aNxisvinmetric disc of mass 1.61011 ALL. radia scale @=Gkpe and vertical scale b.=500pe: and (iv) a spherically symmetric dark halo of mass 3101M. ance scale b=10kpe.," For our disc galaxy, we take the parameters used by Sofue (1996) for the Milky Way which assumes: (i) a spherically symmetric $a=0$ ) nuclear star cluster of mass $5 \times 10^9 \ \msun$, and scale $b =838120\pc$ (ii) a spherically symmetric bulge of mass $10^{10} \msun$ and scale $b = 750\pc$; (iii) an axisymmetric disc of mass $1.6 \times83910^{11} \msun$ , radial scale $a = 6\kpc$ and vertical scale $b = 500840\pc$; and (iv) a spherically symmetric dark halo of mass $3 \times84110^{11} \msun$ and scale $b = 15\kpc$."842 As above. we model the satellite galaxy in this case as a Αμμος sphere. subject to tidal stripping but with the addition ofa central point mass to represent a nuclear black role ancl surrounding star cluster.," As above, we model the satellite galaxy in this case as a Plummer sphere, subject to tidal stripping but with the addition of a central point mass to represent a nuclear black hole and surrounding star cluster."843" We increase the mass of he satellite to be 410""M. but retain the same scale ength of b=OAkpe."," We increase the mass of the satellite to be $4 \times 10^9844\msun$ but retain the same scale length of $b = 0.4 \kpc$."845 This is consistent with the moclels of Sagittarius of Lelmi ancl White (2001)., This is consistent with the models of Sagittarius of Helmi and White (2001).846 This represents he inclusion of clark matter in the model for the cwart., This represents the inclusion of dark matter in the model for the dwarf.847 This could have be added: into the previous simulations. rut would have had Little effect as the satellite was very rapidly destroved by tidal stripping.," This could have be added into the previous simulations, but would have had little effect as the satellite was very rapidly destroyed by tidal stripping."848 We set the central point mass to be 1.10À. although we shall see later that our simulations are basically independent of this.," We set the central point mass to be $1849\times 10^7 \msun$, although we shall see later that our simulations are basically independent of this."850 We model the orbital dynamics ofthe satellite in a manner using basically the same ideas as used above (Section 2.3))., We model the orbital dynamics of the satellite in a manner using basically the same ideas as used above (Section \ref{tidalstripping}) ).851 The satellite galaxy is subject to the gravitational inlluence of the main galaxy. to dynamical friction and to tidal stripping.," The satellite galaxy is subject to the gravitational influence of the main galaxy, to dynamical friction and to tidal stripping."852 The gravitational force due to the main galaxy is garuehtforware to caleulate., The gravitational force due to the main galaxy is straightforward to calculate.853 Due to the axisvmumetry of our model. there are only (wo independent components to consider.," Due to the axisymmetry of our model, there are only two independent components to consider."854 “Phe radial force component (in a evlindrical sense. that parallel to the plane of the disc) is and the vertical component (perpendicular to the disc). The force in both directions is calculated for cach of the four components of our model galaxy ancl these forces sunmimed to produce the resultant force on the satellite due to the main galaxy.," The radial force component (in a cylindrical sense, that parallel to the plane of the disc) is and the vertical component (perpendicular to the disc), The force in both directions is calculated for each of the four components of our model galaxy and these forces summed to produce the resultant force on the satellite due to the main galaxy."855 Yo caleulate the dynamical friction on the satellite. we first calculate the mean streaming velocity. v. of the disc component of the galaxy at the satellite’s current location using Equation 16 (note that all other components have zero mean streaming velocity since they are spherically. svmumetric) ane from this we caleulate the relative velocity of the satellite which respect to the disc.," To calculate the dynamical friction on the satellite, we first calculate the mean streaming velocity, $\overline{v}_\phi$, of the disc component of the galaxy at the satellite's current location using Equation 16 (note that all other components have zero mean streaming velocity since they are spherically symmetric) and from this we calculate the relative velocity of the satellite which respect to the disc."856 Second we compute he local velocity clispersion of each component of the galaxy., Second we compute the local velocity dispersion of each component of the galaxy.857 Then. finally. we use the standard gravitational drag formula (Equation 7)) to calculate the dynamical friction force on he satellite.," Then, finally, we use the standard gravitational drag formula (Equation \ref{chandra}) ) to calculate the dynamical friction force on the satellite."858 Once again. the contribution from. all four components of the galaxy is then summed. ancl this added ο the gravitational force to produce the total force on the satellite.," Once again, the contribution from all four components of the galaxy is then summed, and this added to the gravitational force to produce the total force on the satellite."859 After cach timestep in our integration. we apply simple idal stripping to the satellite.," After each timestep in our integration, we apply simple tidal stripping to the satellite."860 We calculate the mass internal o the satellites current position. via and then use this mass to calculate the current tidal radius of the satellite. as given in Equation 10..," We calculate the mass internal to the satellite's current position, via and then use this mass to calculate the current tidal radius of the satellite, as given in Equation \ref{tidalradius}."861 Whilst this is not strictly accurate. it is an adequate approximation for our current purposes.," Whilst this is not strictly accurate, it is an adequate approximation for our current purposes."862 As we shall see below. the stripping of the satellite is not the most important. factor in. these siniulations.," As we shall see below, the stripping of the satellite is not the most important factor in these simulations."863" Again. we use a variant of the ""odeint routine [roni Numerical Recipes (1992). to control the timestep in the simulation. so that we use small timesteps when necessary. but largerὃν ones when possible to speed up our simulation."," Again, we use a variant of the ' routine from Numerical Recipes \shortcite{numrec} to control the timestep in the simulation, so that we use small timesteps when necessary, but larger ones when possible to speed up our simulation."864 Since we know from our previous set of calculations that the only satellite trajectories which are likely to have a chance of reaching the centre of the host galaxy. are those which are almost radial. we assume that the satellite galaxy is initially approaching the host galaxy from a Large distance (2100 κρο).," Since we know from our previous set of calculations that the only satellite trajectories which are likely to have a chance of reaching the centre of the host galaxy are those which are almost radial, we assume that the satellite galaxy is initially approaching the host galaxy from a large distance $\geq 100\kpc$ )."865 To be specilic. we celine the centre of the Sevíert as the origin of our coordinate svstem. and then consider the trajectories of small. incoming galaxies which start on the plane which has a normal vector at an angle 006 from the axis of rotation of the Seyfert and. has a point. of closest approach of LOOkpe to the origin.," To be specific, we define the centre of the Seyfert as the origin of our coordinate system, and then consider the trajectories of small, incoming galaxies which start on the plane which has a normal vector at an angle $90 - \alpha$ from the axis of rotation of the Seyfert and has a point of closest approach of $100 \kpc$ to the origin."866 We consider. satellite orbits which start in this plane. and which have their initial velocity being perpendicular to the plane.," We consider satellite orbits which start in this plane, and which have their initial velocity being perpendicular to the plane."867" We shall call the angle à the ""approach angle’. as it corresponds to the angle between the initial velocity. vector and the plane of the host galaxy."," We shall call the angle $\alpha$ the `approach angle', as it corresponds to the angle between the initial velocity vector and the plane of the host galaxy."868 Thus à=90° corresponds to an orbit which is initially directed at right anglesto the cise of the host galaxy., Thus $\alpha = 90\degr$ corresponds to an orbit which is initially directed at right anglesto the disc of the host galaxy.869 For this paper. we investigated orbits with approach," For this paper, we investigated orbits with approach"870We make use of two different data sets in this study to maximize the available information and constrain the effect of any possible svstematics.,We make use of two different data sets in this study to maximize the available information and constrain the effect of any possible systematics.871 For the cluster center. we use photometry obtained with USTs Advanced Camera for," For the cluster center, we use photometry obtained with 's Advanced Camera for"872including those selected in Section 2.1.,including those selected in Section 2.1.873" Thus, we have been able to triple the original list, which was largely based on what was found at the Galaxy Zoo forum of ring galaxies."," Thus, we have been able to triple the original list, which was largely based on what was found at the Galaxy Zoo forum of ring galaxies."874 Fig., Fig.875 2 shows the distribution of catalogue objects on the celestial sphere., \ref{fig_sky} shows the distribution of catalogue objects on the celestial sphere.876" We can see that they uniformly fill the region, covered by the Galaxy Zoo, except for a small fragment around (a=140°,6 7°) close to the Galacticplane?."," We can see that they uniformly fill the region, covered by the Galaxy Zoo, except for a small fragment around $\alpha=140^\circ,\,\delta=7^\circ$ ) close to the Galactic."877". By analogy with the PRC, we named our list SPRC — Sloan-Based Polar Rings Catalogue."," By analogy with the PRC, we named our list SPRC = Sloan-Based Polar Rings Catalogue."878" The catalogue is divided into four unequal groups (best candidates, good candidates, related objects and possible face-on rings)."," The catalogue is divided into four unequal groups (best candidates, good candidates, related objects and possible face-on rings)."879" Since the division into these types is often ambiguous, in contrast to Whitmoreetal.(1990) we did not assign each group with its own index, but used a continuous numeration of objects instead."," Since the division into these types is often ambiguous, in contrast to \citet{Whitmore1990} we did not assign each group with its own index, but used a continuous numeration of objects instead."880" However, for convenience sake, the objects in each group are ordered by their R.A. The lists of galaxies themselves are given in Table 2,, where in addition to numbers in our catalogue and equatorial coordinates, we give the total apparent magnitude in the r filter (corrected for the Galaxy extinction, as listed in the SDSS/DRS8table), redshift and the name of the galaxy from well- catalogues (NGC, UGC, CGCG, etc)."," However, for convenience sake, the objects in each group are ordered by their R.A. The lists of galaxies themselves are given in Table \ref{tab_ABCD}, where in addition to numbers in our catalogue and equatorial coordinates, we give the total apparent magnitude in the $r$ filter (corrected for the Galaxy extinction, as listed in the SDSS/DR8, redshift and the name of the galaxy from well-known catalogues (NGC, UGC, CGCG, etc)."881 Redshifts are given according to the NED database., Redshifts are given according to the NED database.882" If the NED data are missing, we use the SDSS data."," If the NED data are missing, we use the SDSS data."883 For several galaxies no SDSS spectra exist., For several galaxies no SDSS spectra exist.884 Fig., Fig.885 3 demonstrates some of the most typical examples of different types of catalogue objects., \ref{fig_example} demonstrates some of the most typical examples of different types of catalogue objects.886 Images of all the catalogue galaxies are given in Fig., Images of all the catalogue galaxies are given in Fig.887 4 (presented in the electronic version of the paper).," \ref{fig_atlas}888 (presented in the electronic version of the paper)."889" Combined colour JPEG images are provided by the SDSS/DRS, which was released in the early 2011 (Eisensteinetal.2011),, and only in a few cases of mosaic creation problems the data were taken from the SDSS/DR7."," Combined colour JPEG images are provided by the SDSS/DR8, which was released in the early 2011 \citep{sdssdr8}, and only in a few cases of mosaic creation problems the data were taken from the SDSS/DR7."890" The division of candidates into types has a significant uncertainty, since it is based on the appearance of galaxies in the optical images, not employing explicit numerical criteria."," The division of candidates into types has a significant uncertainty, since it is based on the appearance of galaxies in the optical images, not employing explicit numerical criteria."891" Taken that, we tried not to discard any “suspicious” objects, which seemed to be related to PRGs."," Taken that, we tried not to discard any “suspicious” objects, which seemed to be related to PRGs."892" Nevertheless, like in the old PRC catalogue, the introduction of several categories is useful in terms of selecting objects for a further detailed study."," Nevertheless, like in the old PRC catalogue, the introduction of several categories is useful in terms of selecting objects for a further detailed study."893 Keep in mind that only a detailed analysis of internal kinematics can confidently confirm the presence of decoupled components in a galaxy., Keep in mind that only a detailed analysis of internal kinematics can confidently confirm the presence of decoupled components in a galaxy.894" For some galaxies, such a confirmation is already available."," For some galaxies, such a confirmation is already available."895" In the cases when there was additional data provided on individual galaxies in the publications, we have used this information to refine the classification."," In the cases when there was additional data provided on individual galaxies in the publications, we have used this information to refine the classification."896" Namely, we classified the galaxy SPRC-1 in which the polar component is vaguely expressed (the dust lane is mostly visible) to the group of best candidates, as Reshetnikovetal.(2011) have shown the presence of an extended ring in the optical images (see as well the ESO Press Release 14/98)."," Namely, we classified the galaxy SPRC-1 in which the polar component is vaguely expressed (the dust lane is mostly visible) to the group of best candidates, as \citet{Resh2011} have shown the presence of an extended ring in the optical images (see as well the ESO Press Release 14/98)."897" The belonging of the SPRC-7 galaxy to PRGs was proved by spectral observations by Broschetal.(2010),, hence we did not classify it into a group of candidates with nearly face-on rings."," The belonging of the SPRC-7 galaxy to PRGs was proved by spectral observations by \citet{Brosch2010}, hence we did not classify it into a group of candidates with nearly face-on rings."898" The polar ring of SPRC-33 is not noticeable in the SDSS images, however, we included the galaxy in the list, since its outer ring reveals both HI and a young stellar population (see Section 4.1))."," The polar ring of SPRC-33 is not noticeable in the SDSS images, however, we included the galaxy in the list, since its outer ring reveals both HI and a young stellar population (see Section \ref{sect_known}) )."899" At the same time, our catalogue does not include the NGC 6822 and SDSS J102819.244-623502.6 galaxies, mentioned in the introduction, since their HI polar discs revealed no detected stellar population."," At the same time, our catalogue does not include the NGC 6822 and SDSS J102819.24+623502.6 galaxies, mentioned in the introduction, since their HI polar discs revealed no detected stellar population."900" We classified 70 galaxies into this type, the belonging of which to the PRG class is almost undoubted, and their images are similar to those, observed in the “classical” long-known PRGs, like NGC4650A, NGC 2685 or UGC 7576."," We classified 70 galaxies into this type, the belonging of which to the PRG class is almost undoubted, and their images are similar to those, observed in the “classical” long-known PRGs, like NGC4650A, NGC 2685 or UGC 7576."901" In the galaxies of this group, the outer component is extended, homogeneous, often distinguished by a bluer colour in comparison to the central galaxy, indicating the presence of a young stellar population."," In the galaxies of this group, the outer component is extended, homogeneous, often distinguished by a bluer colour in comparison to the central galaxy, indicating the presence of a young stellar population."902" Sometimes the rings are accompanied by a dust belt, crossing the central galaxy along the minor axis: SPRC-1, SPRC-25, SPRC-27, SPRC-42, SPRC-48, SPRC-54, SPRC-66, SPRC-69."," Sometimes the rings are accompanied by a dust belt, crossing the central galaxy along the minor axis: SPRC-1, SPRC-25, SPRC-27, SPRC-42, SPRC-48, SPRC-54, SPRC-66, SPRC-69."903" In most candidates the outer ring is close to being viewed edge-on, which is obviously due to the selection effect during the visual selection of objects."," In most candidates the outer ring is close to being viewed edge-on, which is obviously due to the selection effect during the visual selection of objects."904" In some cases, we can clearly see how the outer ring, inclined to the line of sight covers the central body: SPRC-39, SPRC-47, SPRC-51, SPRC-56, SPRC-65, SPRC-65, SPRC-69."," In some cases, we can clearly see how the outer ring, inclined to the line of sight covers the central body: SPRC-39, SPRC-47, SPRC-51, SPRC-56, SPRC-65, SPRC-65, SPRC-69."905" Detailed photometric studies of the SPRC-1 and SPRC- galaxies, which argue for their membership in the PRG class are presented in the papers by Reshetnikovetal.(2011) and Finkelmanetal.(2011),, respectively."," Detailed photometric studies of the SPRC-1 and SPRC-41 galaxies, which argue for their membership in the PRG class are presented in the papers by \citet{Resh2011} and \citet{Finkelman2011}, respectively."906" In the latter case, the authors used the SUBARU Deep Field images, where the outer component looks considerably more extended than in"," In the latter case, the authors used the SUBARU Deep Field images, where the outer component looks considerably more extended than in"907to the Alfven speed by a lavee factor.,to the Alfven speed by a large factor.908" This factor in the textbook: Sweet-Parker model of recouucction is $77, where S=(LVaA£g) is the Lunquist uuuber. where £L is the leneth of the current sheet (Sweet 1958. Parker 1957. see also Parker 1979)."," This factor in the textbook Sweet-Parker model of reconnection is $S^{1/2}$, where $S\equiv (LV_A/\eta)$ is the Lunquist number, where $L$ is the length of the current sheet (Sweet 1958, Parker 1957, see also Parker 1979)."909 Iu general. satisfviug the conservation of lass condition dictates that V7/VaAGL/ZAN).," In general, satisfying the conservation of mass condition dictates that $V_{rec}\sim V_A (L/\Delta)$."910 Observations require a speed close to Wy. so this expression iuplies that L./—A. ie. that the region of over which magnetic flux tubes intersect should be comparable with the outflow region.," Observations require a speed close to $V_A$, so this expression implies that $L\sim \Delta$, i.e. that the region of over which magnetic flux tubes intersect should be comparable with the outflow region."911" This can be achieved either via making £ as small as the Olunic diffusion region ie. that the magnetic field lines reconnect iu an ""X point” (Petscheck 196 or the outflow region should ect lucreased dramatically 0).bevoud the size that is predicted in the Sweet-Parker model."," This can be achieved either via making $L$ as small as the Ohmic diffusion region i.e. that the magnetic field lines reconnect in an “X point” (Petscheck 1964), or the outflow region should get increased dramatically beyond the size that is predicted in the Sweet-Parker model."912 While for vears the problem of fast reconnection was viewed as connected with xoviug of the stability of X point. the situation has changed recently with the second wav of dramatically increasing the thickness of the outflow is becomine nore popular and getting observational support (sec Ciaravella Ravinoud 2008).," While for years the problem of fast reconnection was viewed as connected with proving of the stability of X point, the situation has changed recently with the second way of dramatically increasing the thickness of the outflow is becoming more popular and getting observational support (see Ciaravella Raymond 2008)."913 The first model of X point reconnection was proposed x Petschek (1961)., The first model of X point reconnection was proposed by Petschek (1964).914 Iu this case the recounection speed uay have little or uo dependence on theresistivity., In this case the reconnection speed may have little or no dependence on the.915 The X point configuration is known to be uustable to collapse iuto a sheet iu the MIID reguue (see Diskuup 1996). but in a collisionless plasma it cau be niüinutained through coupling to a dispersive plasma mode (Sturrock. 1966).," The X point configuration is known to be unstable to collapse into a sheet in the MHD regime (see Biskamp 1996), but in a collisionless plasma it can be maintained through coupling to a dispersive plasma mode (Sturrock 1966)."916 Recent vears. have beeu marked by the progress in understanding some of the key processes of reconnection in astroplvsical plasmas.," Recent years, have been marked by the progress in understanding some of the key processes of reconnection in astrophysical plasmas."917" In particular. a substantial progress has been obtained by considering reconnection in the presence of Tall effect. which is described.by the J«B ter in Olinis Iw: where e is clectron charge aud à, is concentration of electrons."," In particular, a substantial progress has been obtained by considering reconnection in the presence of Hall effect, which is describedby the ${\bf J}\times{\bf B}$ term in Ohm's law: where $e$ is electron charge and $n_e$ is concentration of electrons."918 Nuerical experiments showed that ITa1I-MITD reconnection is capable of supporting X-poiuts aud thus can make the reconnection fast. ic. comparable to the Alfven speed (Shay ct al.," Numerical experiments showed that Hall-MHD reconnection is capable of supporting X-points and thus can make the reconnection fast, i.e. comparable to the Alfven speed (Shay et al."919 1998. 200," 1998, 2004)."920" The condition at which ITall-MITD 1).term ects iuportaut for the reconnection is that the ion skin depth 65, is comparable with the Sweet-Parker diffusion scale A.", The condition at which Hall-MHD term gets important for the reconnection is that the ion skin depth $\delta_{ion}$ is comparable with the Sweet-Parker diffusion scale $\Delta$.921" The jon skin depth is a microscopic characteristic aud it can be viewed at the evroradius of an ion moving at the Alfven speed. ie. Ajo,Vyfey Where wey; ds the cyclotron frequency of an jon."," The ion skin depth is a microscopic characteristic and it can be viewed at the gyroradius of an ion moving at the Alfven speed, i.e. $\delta_{ion}=V_A/\omega_{ci}$, where $\omega_{ci}$ is the cyclotron frequency of an ion."922" Iu the leliotail for a proton we find that 3;,,~10? kan.", In the heliotail for a proton we find that $\delta_{ion}\sim 10^3$ km.923" Thus one can get the constraint on the scale £ for which ITiAll-MITD effects should dominate the recounectiou: where A, is the electron mea free path. where Jj is the ratio of thermal pressure to magnetic pressure (see mniore discussion in Yanada et al."," Thus one can get the constraint on the scale $L$ for which Hall-MHD effects should dominate the reconnection: where $\lambda_{mfp}$ is the electron mean free path, where $\beta_{pl}$ is the ratio of thermal pressure to magnetic pressure (see more discussion in Yamada et al."924 2006)., 2006).925 We aerue in Appendix A that in realistic situations in turbulent media the scales A|=Lent (soe Figure 5)) over which the mucroscale Sweet-Parker reconnection of iudividual turbuleut patches may fake place are much siualler thui the scale. of the system aud therefore the collisiouless effects take place within the heliotail., We agrue in Appendix A that in realistic situations in turbulent media the scales $\lambda_{\|}=L_{turb}$ (see Figure \ref{fig:recon1}) ) over which the microscale Sweet-Parker reconnection of individual turbulent patches may take place are much smaller than the scale of the system and therefore the collisionless effects take place within the heliotail.926 This. as we argue below. does not change the overall rates of magnetic reconuectiou.," This, as we argue below, does not change the overall rates of magnetic reconnection."927 A shortcoming of many discussions of maguctic reconnection is that the traditional setup does uot include ubiquitous pre-existing astrophysical (κος. Armstrone. Rickett Spaneler 1991. Ehucerecu Sealo 2001. Mel&oe Ostriker 2007. Lazarian 2009. Chepurnov Lazari 2010).," A shortcoming of many discussions of magnetic reconnection is that the traditional setup does not include ubiquitous pre-existing astrophysical (see Armstrong, Rickett Spangler 1994, Elmegreen Scalo 2004, McKee Ostriker 2007, Lazarian 2009, Chepurnov Lazarian 2010)."928 As turbulence radically changes mauv astrophwsical processes. the mfuence of turbulence ou reconnection has attracted the attention of researchers for a long time (see Speizer 1970. Straus 1988).," As turbulence radically changes many astrophysical processes, the influence of turbulence on reconnection has attracted the attention of researchers for a long time (see Speizer 1970, Straus 1988)."929 An exteuded discussion of turbulence role in accelerating reconnection can be found in Mathaeus Lamilsin (1985. 1986).," An extended discussion of turbulence role in accelerating reconnection can be found in Mathaeus Lamkin (1985, 1986)."930 However. there tle X-poiuts created by turbulence together with the effects of compressibility aud heating were ideutified as the means of accelerating reconnection.," However, there the X-points created by turbulence together with the effects of compressibility and heating were identified as the means of accelerating reconnection."931 A very different approach to the effects of turbulence was adopted in Lazarvian Vishuiac (1999. heuceforth LV99).," A very different approach to the effects of turbulence was adopted in Lazarian Vishniac (1999, henceforth LV99)."932 Their model does not appeal to auv of the effects of turbulence-created X-poiuts compressibility or heating and it is applicable to a wide rauge of astroplivsical conditions.," Their model does not appeal to any of the effects of turbulence-created X-points, compressibility or heating and it is applicable to a wide range of astrophysical conditions."933 Fortunately. this approach provides a robust way of accelerating reconnection.," Fortunately, this approach provides a robust way of accelerating reconnection."934 Tudeed. as we mentioned above. the approach in LVY99 is to consider wavs to decouple the width of the plasia outflow region from the scale determined by Oluic The plasma is constrained to move along maguctic feld lines. but uot necessarily in the cirection of the can magnetic field.," Indeed, as we mentioned above, the approach in LV99 is to consider ways to decouple the width of the plasma outflow region from the scale determined by Ohmic The plasma is constrained to move along magnetic field lines, but not necessarily in the direction of the mean magnetic field."935 Ina turbulent mediu the two are decoupled. aud fluid elements that have some small initial separation will be separated by a large eddy scale or more after moving the leneth of the current sheet.," In a turbulent medium the two are decoupled, and fluid elements that have some small initial separation will be separated by a large eddy scale or more after moving the length of the current sheet."936 As long as lis separation is larger than the width of the current sheet. the result will not depend on yj.," As long as this separation is larger than the width of the current sheet, the result will not depend on $\eta$."937 The mental picture xeseuted in LV99 is that the fluid follows magnetic field ines. Which are uot straight. but wider (see Figure 5 and also Lazarian ct al.," The mental picture presented in LV99 is that the fluid follows magnetic field lines, which are not straight, but wander (see Figure \ref{fig:recon1} and also Lazarian et al."938 200 where this wandering was calculated πισαν aud compared1. with the analytical xedictiouns in LV99).," 2004, where this wandering was calculated numerically and compared with the analytical predictions in LV99)."939 As a result. the thickness A ofthe Huid outflow is determined not by microphysical Olunic diffusivity. but magnetic field wanderme which for the injection velocity V; of the order of V4 is of the order of the turbulence injection scale ἐν ie. Axd.," As a result, the thickness $\Delta$ of the fluid outflow is determined not by microphysical Ohmic diffusivity, but magnetic field wandering which for the injection velocity $V_l$ of the order of $V_A$ is of the order of the turbulence injection scale $l$, i.e. $\Delta \approx l$."940 Tf the leugth of the current sheet L is of the order of / it is clear from Figure 5. that VW... can be comparable with Va., If the length of the current sheet $L$ is of the order of $l$ it is clear from Figure \ref{fig:recon1} that $V_{rec}$ can be comparable with $V_A$ .941 Note that LV99 cousiders generic 3D configurations of magnetic fluxes with non-zero magnetic eue. ie. shared. magnetic field.," Note that LV99 considers generic 3D configurations of magnetic fluxes with non-zero magnetic guide, i.e. shared, magnetic field."942 The shared maguctic field is, The shared magnetic field is943Planetary nebulae (PNe) represent the terminal breath of 90% of the stars in the Universe.,Planetary nebulae (PNe) represent the terminal breath of $\%$ of the stars in the Universe.944 However. their shaping mechanism is still poorly understood.," However, their shaping mechanism is still poorly understood."945 Bipolar PNe are undoubtedly the most challenging case., Bipolar PNe are undoubtedly the most challenging case.946 Several attempts have been made to explain their shaping (see the review by ?)). breaking spherical symmetry by invoking elements which fall in two distinct categories: rapid stellar rotation. and/or magnetic fields (e.g.??).. and a close interacting companion to the star (e.g. ?.. for a review see 2)).," Several attempts have been made to explain their shaping (see the review by \citealp{balick02}) ), breaking spherical symmetry by invoking elements which fall in two distinct categories: rapid stellar rotation and/or magnetic fields \citep[e.g.][]{garciasegura99,blackman01a}, and a close interacting companion to the star (e.g. \citealt{nordhaus06}, for a review see \citealt{demarco09}) )."947 This latter hypothesis seems to be gaining some ground as close binary systems are progressively being found (e.g. ?:: 2)) at the cores of bipolar PNe., This latter hypothesis seems to be gaining some ground as close binary systems are progressively being found (e.g. \citealp{miszalski09a}; \citealp{miszalski10}) ) at the cores of bipolar PNe.948 Spatio-kinematical modelling of Pe constitutes an excellent tool to test theoretical models., Spatio-kinematical modelling of PNe constitutes an excellent tool to test theoretical models.949 It provides us with important parameters to be matched by the different models of formation. such as the 3-D morphologies and velocity fields of the outflows. their kinematical age (once disentangled from the distance to the nebula) and their orientation to the line of sight.," It provides us with important parameters to be matched by the different models of formation, such as the 3-D morphologies and velocity fields of the outflows, their kinematical age (once disentangled from the distance to the nebula) and their orientation to the line of sight."950 M 1-75 (PN G068.8-00.0. a = 20 04 44.086 ο = +31 27 24.42 J2000) is a good example of a complex nebula.," M 1–75 (PN G068.8-00.0, $\alpha$ = 20 04 44.086 $\delta$ = +31 27 24.42 J2000) is a good example of a complex nebula."951 It displays a seemingly irregular horseshoe-like central region. out of which two systems of faint lobes emerge.," It displays a seemingly irregular horseshoe-like central region, out of which two systems of faint lobes emerge."952 It was first classified as quadrupolar by ?.. and a tentative attempt to recover its kinematic parameters was done by ?..," It was first classified as quadrupolar by \citet{manchado96b}, and a tentative attempt to recover its kinematic parameters was done by \citet{dobrincic08}."953 In this paper we present Fabry-Perot interferometry of1—75.. from which we derive a detailed spatio-kinematical model (section. 3).," In this paper we present Fabry-Perot interferometry of, from which we derive a detailed spatio-kinematical model (section 3)."954 We also report the first imaging and spectroscopic detection of its central star (section 4)., We also report the first imaging and spectroscopic detection of its central star (section 4).955 We then discuss both results and their implications in the shaping of the nebula in section 5., We then discuss both results and their implications in the shaping of the nebula in section 5.956 The [Nu] 658.3 nm emission of M 1-75 was scanned with GHeFAS (Galaxy Ha Fabry-Perot System) on the 4.2 m WHT (William Herschel Telescope) on July 6. 2007. as part of its commissioning. programme.," The ] 658.3 nm emission of M 1-75 was scanned with $\alpha$ FAS (Galaxy $\alpha$ Fabry-Perot System) on the 4.2 m WHT (William Herschel Telescope) on July 6, 2007, as part of its commissioning programme."957" The nebula was observed in high-resolution mode with the OMA etalon (resolving power R~ 18000. effective finesse 9,224) and a plate scale of 0.2 pixel!."," The nebula was observed in high-resolution mode with the OM4 etalon (resolving power $\sim$ 18000, effective finesse $\Im_\mathrm{e}$ =24) and a plate scale of $''$ .2 $^{-1}$."958 The free spectral range was 8.62 A or 392 km s split into 48 channels. thus leading to a velocity step of 8.16 km s! per channel.," The free spectral range was 8.62 $\AA$ or 392 km $^{-1}$ split into 48 channels, thus leading to a velocity step of 8.16 km $^{-1}$ per channel."959 The total exposure time of the scanning was 1.9 hr. and the seeing 07.8.," The total exposure time of the scanning was 1.9 hr, and the seeing $''$ .8."960 The instrumental response function (IRF) was measured by fitting a Lorentzian to the profile of a Neon lamp line and resulted in an instrumental width (FWHM) of 18.6 πε..., The instrumental response function (IRF) was measured by fitting a Lorentzian to the profile of a Neon lamp line and resulted in an instrumental width (FWHM) of 18.6 km $^{-1}$.961 The data were reduced following the standard procedure for GHaFAS data. which are deseribed in ?..," The data were reduced following the standard procedure for $\alpha$ FAS data, which are described in \citet{hernandez08}."962 Several artifacts persisted. through the data reduction. process., Several artifacts persisted through the data reduction process.963 These include slight contamination by Πα emission from adjacent orders (specially in the first and last channels of the datacube). a ghost of the inner region of the nebula. and an are-shaped artifact which runs across several channels. at different locations (see Fig. 1)).," These include slight contamination by $\alpha$ emission from adjacent orders (specially in the first and last channels of the datacube), a ghost of the inner region of the nebula, and an arc-shaped artifact which runs across several channels, at different locations (see Fig. \ref{F1}) )."964 Several images of M 1-75 in the light of different filters (U. B. V. 1. Πα and Strómmgren Y) were taken with ACAM (Auxiliary-port Camera) on the WHT and with the WFC (Wide Field Camera) on the 2.5 m INT (Isaac Newton Telescope).," Several images of M 1–75 in the light of different filters (U, B, V, I, $\alpha$ and Strömmgren Y) were taken with ACAM (Auxiliary-port Camera) on the WHT and with the WFC (Wide Field Camera) on the 2.5 m INT (Isaac Newton Telescope)."965 The log of the observations can be found in Table I., The log of the observations can be found in Table 1.966 All these data were reduced following standard procedures., All these data were reduced following standard procedures.967" An 3600 s spectrum with the slit at parallactic angle (P.A.=284""). crossing the centre of the inner nebula. was taken with IDS (Intermediate Dispersion Spectrograph) on the INT"," An 3600 s spectrum with the slit at parallactic angle $^\circ$ ), crossing the centre of the inner nebula, was taken with IDS (Intermediate Dispersion Spectrograph) on the INT"968 We consider two continuous random scalar fields. Fy and Fy. in a space with d dimensions and any topology.," We consider two continuous random scalar fields, $F_A$ and $F_B$, in a space with $d$ dimensions and any topology."969 Without loss of generality. we shall assume that the mean values (over realization) of both fields is zero and they are not correlated: (F4)=(F5)(FyFx)0.," Without loss of generality, we shall assume that the mean values (over realization) of both fields is zero and they are not correlated: $\langle F_A\rangle=\langle F_B\rangle =\langle F_AF_B\rangle=0$."970 On the other hand. the values of each field at two different points x4. x» are not independent random variables: coμιX?){ΕνμίαΕνμία). where the average is over realizations.," On the other hand, the values of each field at two different points $\vec{x}_1$, $\vec{x}_2$ are not independent random variables: $\omega _{A/B}(\vec{x}_1,\vec{x}_2)=971\langle F_{A/B}(\vec{x}_1)F_{A/B}(\vec{x}_2)\rangle$, where the average is over realizations."972" In practice. in most interesting cases the fields are statistically homogeneous and ergodic. so that co depends only on x,—x». and the correlation may be defined as spatial averages. which is the useful definition since in most cases only one realization is available."," In practice, in most interesting cases the fields are statistically homogeneous and ergodic, so that $\omega $ depends only on $\vec{x}_1-\vec{x}_2$, and the correlation may be defined as spatial averages, which is the useful definition since in most cases only one realization is available."973" If the fields are also statistically isotropic. «o depends only on r=|x,—x2]."," If the fields are also statistically isotropic, $\omega $ depends only on $r\equiv |\vec{x}_1-\vec{x}_2|$."974 For the following derivation. we shall assume homogeneity and isotropy: the full expression might easily be recovered if needed.," For the following derivation, we shall assume homogeneity and isotropy; the full expression might easily be recovered if needed."975" We first derive the ""field-to-field error"" (i.e.. the true error)for the zero lag estimator: where Fypf)=Fog."," We first derive the “field-to-field error” (i.e., the true error)for the zero lag estimator: where $F_{A/B}(i)\equiv F_{A/B}(\vec{x}_i)$."976 We have replaced the d-dimensional volume integral over the sample with a sum over V equal volume cell indexed by / and centered on x;., We have replaced the $d$ -dimensional volume integral over the sample with a sum over $N$ equal volume cell indexed by $i$ and centered on $\vec{x}_i$.977 We would have to multiply the contribution of the field for each pixel by a weight equal to the volume of the pixel in the case of non-equal volume cells., We would have to multiply the contribution of the field for each pixel by a weight equal to the volume of the pixel in the case of non-equal volume cells.978 For the variance in Eq. CAI. ," For the variance in Eq. \ref{bet1}) ),"979we have: Asince. by construction. the mean value of E over realizations. cip. is assumed to be zero.," we have: since, by construction, the mean value of $E$ over realizations, $\omega _{AB}$, is assumed to be zero."980 Developing the square of expression (ΑΟ). and taking its average. we have: In principle. the factor of 2 should not be there in the case i=j. but this will be negligible in the limit of arbitrarily small cells.," Developing the square of expression \ref{bet1}) ), and taking its average, we have: In principle, the factor of 2 should not be there in the case $i=j$ , but this will be negligible in the limit of arbitrarily small cells."981 Now. since the fields Εν and Fy are uncorrelated whereOF r;;=CAFx;RL)—xj].," Now, since the fields $F_A$ and $F_B$ are uncorrelated where $r_{ij}\equiv |\vec{x}_i-\vec{x}_j|$."982" Thus. we have where V, represents the volume of theNunple sample."," Thus, we have where $V_s$ represents the volume of the sample."983 The correlation estimator for any non-zero lag is where rp—Ar/2«rry+ Ar/2., The correlation estimator for any non-zero lag is where $r_0-\Delta r/2<r<r_0+\Delta r/2$ .984 Following Eg. (A6)), Following Eq. \ref{bet6}) )985 using the same procedure asfor Eg. CAL9. ," using the same procedure asfor Eq. \ref{bet1}) ),"986one obtains (AppendixA: CAS). (AppendixΑ:)), one obtains \ref{bet7}) \ref{bet5}) \ref{bet7})987still obtains a very sinall place at the 1991. Marcel Crossimaun mecting |72]..,still obtains a very small place at the 1991 Marcel Grossmann meeting \cite{MG91}.988 But. increasingly. the research in supereravity and higher derivative theories has merged iuto strings. and strime theory is iucreasinely viewed as a stroug colmpcting candidate for the quantum theory of the gravitational field.," But, increasingly, the research in supergravity and higher derivative theories has merged into strings, and string theory is increasingly viewed as a strong competing candidate for the quantum theory of the gravitational field."989 As a side product. nauy particle plivsicists beein to study ecucral relativity. or at least sole bits of it.," As a side product, many particle physicists begin to study general relativity, or at least some bits of it."990 Strings provide a cousisteut perturbative theory., Strings provide a consistent perturbative theory.991 The covariaut program is fully re-boru., The covariant program is fully re-born.992 The problem becomes uuderstaudius why the world described by the theory appears so different frou ours., The problem becomes understanding why the world described by the theory appears so different from ours.993 Ted Jacobson aud Lee Sinolin find loop-like solutions to the Wheeler-DeWitt equation formulated iu the connection formulation |73].. opening the way to loop quantuni eravity.," Ted Jacobson and Lee Smolin find loop-like solutions to the Wheeler-DeWitt equation formulated in the connection formulation \cite{TL}, opening the way to loop quantum gravity."994" The ""loop represeutation of quantum general relativity” is introduced iu [r4]..."," The “loop representation of quantum general relativity"" is introduced in \cite{rovellismolin88}."995 For an carly review. sce [75]..," For an early review, see \cite{report}."996 Tt is based on the new connection formulation of CR [69].. on the Jacobson-Simolin solutions [73].. and on Chris Tshau’s ideas ou the need of non-Gaussian. or non-Fock represcutations im quantum eravity [r6]..," It is based on the new connection formulation of GR \cite{ashtekar}, , on the Jacobson-Smolin solutions \cite{TL}, and on Chris Isham's ideas on the need of non-Gaussian, or non-Fock representations in quantum gravity \cite{chris}."997 Loop quantization had been previously aud independently developed by Rodolfo Canibini and his collaborators for Yang Mills theories |77].., Loop quantization had been previously and independently developed by Rodolfo Gambini and his collaborators for Yang Mills theories \cite{gambini}.998 Iu. the eravitational context. the loop representation leads imuuediately to two surprising results: an infinite family of exact solutions of the Whoeoler-DeWitt equation is found. aud knot theory controls the plivsical quanti states ofthe eravitational field.," In the gravitational context, the loop representation leads immediately to two surprising results: an infinite family of exact solutions of the Wheeler-DeWitt equation is found, and knot theory controls the physical quantum states of the gravitational field."999 Classical knot theory. with its extensions. becomes a brauch of matlematics relevant to describe the diff iuvariaut states of quautum spacetime [78]..," Classical knot theory, with its extensions, becomes a branch of mathematics relevant to describe the diff invariant states of quantum spacetime \cite{pullin}."1000 The theory traustorms the old Wheeler-DeWitt theory in a formali that can be concretely used to compute plivsical quautities in quantum gravity., The theory transforms the old Wheeler-DeWitt theory in a formalism that can be concretely used to compute physical quantities in quantum gravity.1001 The canonical program is fully re-bornu., The canonical program is fully re-born.1002 Today. the theory is called “loop απιαλται For a review. complete references; andan account of the development of thetheory. κος |79]..," Today, the theory is called “loop quantum For a review, complete references, andan account of the development of thetheory, see \cite{lqg}. ."1003Figure l.. aud performed the spectral analysis separately for each shell.,"Figure \ref{images}, and performed the spectral analysis separately for each shell."1004 We excluded. all of the detected point sources (see rofsccesultsinuge)) froiicaclshell, We excluded all of the detected point sources (see \\ref{sec_results_image}) ) from each shell.1005 Thebackyroundrcy ο μμ allied γι freeregionsinthesanmee CDandobsercationsasthoscofthesha]eucut, The background regions were also selected separately for each shell from the source-free regions in the same CCD and observations as those of the shells.1006 The fields covered the whole diffuse structure., The fields covered the whole diffuse structure.1007 Towever. for cousisteney with theChandra analysis. we divided the diffuse structure into shells AD. the same as iu theChandra case.," However, for consistency with the analysis, we divided the diffuse structure into shells A–D, the same as in the case."1008 The background. reeious were selected from the source-free regions im the same observation., The background regions were selected from the source-free regions in the same observation.1009 The spectra of cach shell taken from all the available observations with and were analyzed simultancously., The spectra of each shell taken from all the available observations with and were analyzed simultaneously.1010 However. for brevity. we show only the spectra (backerouud-subtracted) aud the fitting results in Figure 2..," However, for brevity, we show only the spectra (background-subtracted) and the fitting results in Figure \ref{shell_spec}."1011" For spectral analyses. we used NSPEC v11.00 (Arnaud1996) The spectrum of shell A shows many line-like structures with the center energies at 0.58. 0.68. 0.92. and 1.35 keV. which correspond to the emission lines of Πομκο ο Ko and Kj. Te-like Ne. and Heike Ale. respectively,"," For spectral analyses, we used XSPEC v11.00 \citep{arnaud}1012 The spectrum of shell A shows many line-like structures with the center energies at 0.58, 0.68, 0.92, and 1.35 keV, which correspond to the emission lines of He-like O $\alpha$ and $\beta$, He-like Ne, and He-like Mg, respectively."1013 We therefore ft the spectrmu with a thin-thenual plasma model in nou-equilibriuni ionization CNET) (NEI: Borkowski ct 22001a) with the mean LAIC abundauces (Russel&Dopita1992:Hughes.Tavashi.," We therefore fit the spectrum with a thin-thermal plasma model in non-equilibrium ionization (NEI) ; Borkowski et 2001a) with the mean LMC abundances \citep{russel,hughes}."1014"&Ἱνοναιπα 1998).. The interstellar absorption in our Galaxy aud LMC were treated separately,", The interstellar absorption in our Galaxy and LMC were treated separately.1015 The Calactic absorption colunui was estimated using the III data by Dickey&Loclanan(L990) as NIL=οον1029 Cni2| Arvahady," The Galactic absorption column was estimated using the HI data by \citet{dicky} as $N_{\rm H, HI} = 6.35\times 10^{20}$ $^{-2}$."1016jisMBreeman(1999) reported that the value of Iq. neasured im the N-rav bid. is twice that of ΑΠΠι in the case of |b]>25° and Nyy>5«102 7.," \citet{arabadjis} reported that the value of $N_{\rm H}$, measured in the X-ray band, is twice that of $N_{\rm H, HI}$ in the case of $|b|>25^\circ$ and $N_{\rm H, HI}>5\times 10^{20}$ $^{-2}$."1017 Therefore. we fixed the galactic absorption colin o be Ny=1.27<loeb 7: we used the cross sections by Morrison&Al[cCamunon(1983). and he solar abundances (Anders&Crevesse1989).," Therefore, we fixed the galactic absorption column to be $N_{\rm H} = 1.27\times 10^{21}$ $^{-2}$; we used the cross sections by \citet{morrison}1018 and the solar abundances \citep{anders}."1019 The absorption colum iu the LMC was. on the other haud. treated as a free parameter with the nean LMC abundance (Russel&Dopita1992:IIughes.Hayashi.," The absorption column in the LMC was, on the other hand, treated as a free parameter with the mean LMC abundance \citep{russel,hughes}."1020&Ἱνοναπια 1998).. This thin-thermal plaswa model was rejected with a u ‘cleeree of freedom (d.o.£.), This thin-thermal plasma model was rejected with a $\chi^2$ /degree of freedom (d.o.f.)1021 = 331.5/211. even if we allow the abundances to be free. leaving a systematic data residual at the high-energv baud.," = 331.8/211, even if we allow the abundances to be free, leaving a systematic data residual at the high-energy band."1022 We hence added a power-law component ou the thin-thermal model., We hence added a power-law component on the thin-thermal model.1023" Since the two-component model still leaves large residuals at about 0.5 keV and 1.3 keV. we allowed the abuudauce of Fe ix Me in the thermal plasiua (""NET componcut) to be free."," Since the two-component model still leaves large residuals at about 0.8 keV and 1.3 keV, we allowed the abundance of Fe and Mg in the thermal plasma (“NEI” component) to be free."1024 The fitting was then ereath improvec ran sb sed, The fitting was then greatly improved with a $\chi^2$ /d.o.f = 264.6/209.1025 te model is still rejected in a statistica point of view. further flue tunius ou the mode is bevoud the scope of this paper.," Although this two-component model is still rejected in a statistical point of view, further fine tuning on the model is beyond the scope of this paper."1026 Figure 2. iux Table 3. show the best-fit models aud parameters. respectively.," Figure \ref{shell_spec} and Table \ref{spec_A} show the best-fit models and parameters, respectively."1027 Uulike shell A. the N-ray spectra of shells D re hard aud featureless. sugeesting nou-henual origin.," Unlike shell A, the X-ray spectra of shells B--D are hard and featureless, suggesting non-thermal origin."1028 Iu fact. a thin thermal model fittine requires an unrealistically high temperature (> 2 keW) and low abundances (2<0.3).," In fact, a thin thermal model fitting requires an unrealistically high temperature $>$ 2 keV) and low abundances $z < 0.3$ )."1029 We herefore fitted the spectra with a power-law uodel with absorption. which was caleulated iu he same wav as that for shell A. aud fouud acceptable fits for all of the spectra.," We therefore fitted the spectra with a power-law model with absorption, which was calculated in the same way as that for shell A, and found acceptable fits for all of the spectra."1030 The best-ft models aud parameters are shown in Figure 2 aud Table L.. respectively.," The best-fit models and parameters are shown in Figure \ref{shell_spec} and Table \ref{spec_NT}, respectively."1031 It is couceivable that the spectra of shells D may include a small fraction of the thiu-hermal component found in shell A. We therefore added the same thin-thermal spectrum as that or shell A. aud fitted with this composite model (thin-thermal plus power-law).," It is conceivable that the spectra of shells B–D may include a small fraction of the thin-thermal component found in shell A. We therefore added the same thin-thermal spectrum as that for shell A, and fitted with this composite model (thin-thermal plus power-law)."1032 The free parameUS are norlualizations ofthe two components: powcr-aw index and Ny values, The free parameters are normalizations of the two components: power-law index and $N_{\rm H}$ value.1033 However. uo siguificaut raction of the thLB.i thermal conmiponeut is fouud roni shells D. For all of the poiut sources. the X-ray. photous are collected from an ellipse with the radii of the »omt spread function (PSF). as listed in Table 2..," However, no significant fraction of the thin thermal component is found from shells B–D. For all of the point sources, the X-ray photons are collected from an ellipse with the radii of the point spread function (PSF), as listed in Table \ref{point}."1034" We note that all of the sources are located far roni the on-axis position of the N-rav 1iinror. aud rence the PSEs are larecr than the best value of PSF (~ 0.5""on the anu point)."," We note that all of the sources are located far from the on-axis position of the X-ray mirror, and hence the PSFs are larger than the best value of PSF $\sim$ on the aim point)."1035 The vackerouud regious were selected from source-free regions in the same wav as the diffuse cussions., The background regions were selected from source-free regions in the same way as the diffuse emissions.1036 We first fitted the spectra with a thin-tlermal asma inodel in collisional equilibrium(MEKAL: Alewe. Groucuschild. van den 11985: Iaastra 1992) with an absorption caleulated iu he same wav as diffuse enissiou.," We first fitted the spectra with a thin-thermal plasma model in collisional equilibrium; Mewe, Gronenschild, van den 1985; Kaastra 1992) with an absorption calculated in the same way as diffuse emission."1037 The abundances are fixed to be 0.3 solar. the average value of interstellar medium in the LMC.," The abundances are fixed to be 0.3 solar, the average value of interstellar medium in the LMC."1038 The fittings are acceptable for two sources ll aud 1) with reasonable temperature (2.1 and 1.0 keV). but for the spectra of the other Lt sources. the models areeither rejected or require an unreasonablv high," The fittings are acceptable for two sources 1 and 4) with reasonable temperature (2.1 and 1.0 keV), but for the spectra of the other 4 sources, the models areeither rejected or require an unreasonably high"1039distinguish four populations of different ages aud star onuation history. which we now describe.,"distinguish four populations of different ages and star formation history, which we now describe."1040 The oldest population is the splieroid., The oldest population is the spheroid.1041 For this »opulation we assume a single-burst star formation istorv ocurnug carly in the lifetime of the Galaxy. around Ll Cor ago. from eas still iu a spherokal distribution.," For this population we assume a single-burst star formation history ocurring early in the lifetime of the Galaxy, around 14 Gyr ago, from gas still in a spheroidal distribution."1042 The iuitial mass function (IME) and the deusity distribution of lis population is characterised by a ]rower law index. au axis ratio anda local normalisation. wuch are constrained o» reniofe star counts at lieh ane medium Calactic atitudes (Robin ct al. 3 )003)..," The initial mass function (IMF) and the density distribution of this population is characterised by a power law index, an axis ratio and a local normalisation, which are constrained by remote star counts at high and medium Galactic latitudes (Robin et al. \cite{Robin2000}) )."1043 The sInelatics are also deduced. from iu situ vekcity deterinations., The kinematics are also deduced from in situ velocity determinations.1044 The mean uetallicitv of the spheroid is assumed obe ddex with an intrinsic dispersiou of ddex., The mean metallicity of the spheroid is assumed to be dex with an intrinsic dispersion of dex.1045 SecoIv. a population. called the hick disc. is formed of stars born about 11-12 Gar ago ina short period of time as Προς by recent iueallicity. deteYnünatious for this population.," Secondly, a population, called the thick disc, is formed of stars born about 11-12 Gyr ago in a short period of time as implied by recent metallicity determinations for this population."1046 We also assuue a single must for simplicity., We also assume a single burst for simplicity.1047 For the thick disk. star formation occurred from the gas already setled iu the disk.," For the thick disk, star formation occurred from the gas already settled in the disk."1048 The kinematics. dediCO( frou observatiowal constraints (Ojlia et al. 1996.. ," The kinematics, deduced from observational constraints (Ojha et al. \cite{Ojha96}, ,"1049Ojha 1 999)). nuples that if has uudergone a imereie event sjortlv after he disc formation (Robin et al. 1996)).," Ojha \cite{Ojha99}) ), implies that it has undergone a merging event shortly after the disc formation (Robin et al. \cite{Robin96}) ),"1050 1icreasiues he disk thickuess and eiviug a higher velocity «Ispersion aud scale height., increasing the disk thickness and giving a higher velocity dispersion and scale height.1051 The IME. density distribution aux local iornalisation were constramed from star counts {Revlé Robin 20013).," The IMF, density distribution and local normalisation were constrained from star counts (Reylé Robin \cite{Reyle2001}) )."1052 The mean metallicity of the thick disc IS assu| to be dadex with an intrinsic dispersion of ddex., The mean metallicity of the thick disc is assumed to be dex with an intrinsic dispersion of dex.1053 Thirdly. a bulge population is picent in the center of he Calaxyv and extends to about 2 Kkpe.," Thirdly, a bulge population is present in the center of the Galaxy and extends to about 2 kpc."1054 Its age is of the order of 1) Cyr., Its age is of the order of 10 Gyr.1055 This value is however poorly coustrainecd., This value is however poorly constrained.1056 This popuation has a triaxial distribution. as a bar as determined by DPicaud et al. (20013) ," This population has a triaxial distribution, as a bar as determined by Picaud et al. \cite{Picaud2004}) )"1057from near-infrared star couns., from near-infrared star counts.1058 Velocity dispersions are large. similar to the spheroid.," Velocity dispersions are large, similar to the spheroid."1059 The mean metallicity is assunied solar with a dispersion of ddex., The mean metallicity is assumed solar with a dispersion of dex.1060 A staidard evolution model is used to compute the disc poptlation. based on a typical set of paraicters: al IMIF. asar fornation rate (SFR). a set of evolutionary tracks (see Tavwood et al.," A standard evolution model is used to compute the disc population, based on a typical set of parameters: an IMF, a star formation rate (SFR), a set of evolutionary tracks (see Haywood et al.,"1061 L997a aud references therein)., 1997a and references therein).1062" The cise populaloli Is assned to evolve during 10 Cor,", The disc population is assumed to evolve during 10 Gyr.1063 A set of IME slop sand SER aue tentatively assuiied aud tested agalust sar counts., A set of IMF slopes and SFR's are tentatively assumed and tested against star counts.1064 The tuning of disc parameters agalust relevant observatioval data las been described in Tavwood et al. (19972... 1997b)).," The tuning of disc parameters against relevant observational data has been described in Haywood et al. \cite{Haywood97}, \cite{Haywood97a}) )."1065 A revised IAF has been used as a starting poit in the prescut analysis. adjusted ο agree With the most recent Ilipparcos resuls: the age-velocity dispersion relation is from CóΠιο ο al. (19973).," A revised IMF has been used as a starting point in the present analysis, adjusted to agree with the most recent Hipparcos results: the age-velocity dispersion relation is from Gómmez et al. \cite{Gomez97}) ),"1066 the local luminosity function is from Wiclen (1997)) and an IME is adjusted to it. eiviug au IMF slope a = 1.5 in the low mass range [0.5-0.08 ALJ]. in good agrecment with Iroupa (2001)).," the local luminosity function is from Wielen \cite{Jahreiss97}) ) and an IMF is adjusted to it, giving an IMF slope $\alpha$ = 1.5 in the low mass range [0.5-0.08 ], in good agreement with Kroupa \cite{Kroupa2001}) )."1067 The scale height has been computed selfcousisteutlv. using the potential via tιο Boltzmann equation., The scale height has been computed self-consistently using the potential via the Boltzmann equation.1068 The local dynamical Lass was Τaken from Crézzé et al. (1998) )., The local dynamical mass was taken from Crézzé et al. \cite{Creze1998}) ).1069 T1e evolutionary model fixes the distribution of stars witlii the parameter space of effective teniperaure. eyavlY. absolute uaenitude. mass and age.," The evolutionary model fixes the distribution of stars within the parameter space of effective temperature, gravity, absolute magnitude, mass and age."1070 These paraltors are COLNrorted iuto colours nmi various svsclus throug1 stellar atic»pliere models., These parameters are converted into colours in various systems through stellar atmosphere models.1071 Iu he case of tre MEGACAMNM. photometric svete. we have used the ο]otics. CCD aud filter definition of the passhaids. aud appied thei to the spectral libraries.," In the case of the MEGACAM photometric system, we have used the optics, CCD and filter definition of the passbands, and applied them to the spectral libraries."1072 As SCCLL 111 Fig., As seen in Fig.1073 1. the Dasel3.1 hbrarv is more suitable for hot stars. οwing better predicted colours. especially &4’. while ιο NextCen ibrary is more realistic for cool stars.," 1, the Basel3.1 library is more suitable for hot stars, giving better predicted colours, especially $u^{*}-g'$, while the NextGen library is more realistic for cool stars."1074 IIeuce we have adoted: a combivation of both: Dasel3.1 at Tor>LOOOTS and NextCen or cooler stars.," Hence we have adopted a combination of both: Basel3.1 at $\rm T_{eff} >10754000\,K$ and NextGen for cooler stars."1076 As the cooler stars are mostly located 1i the dise. we used ouly [Fe/TT=0.0 aud |Fe/TI|21.0 fex Tog<10001.," As the cooler stars are mostly located in the disc, we used only [Fe/H]=0.0 and $=-1.0$ for $\rm1077T_{eff} < 4000\,K$."1078 The inodel sinulationus aIso include a nodel of extiuction aud ACCOMit for observational errors., The model simulations also include a model of extinction and account for observational errors.1079 The Besancoon model predictions in the MECACAM photometric svsteni can be fouxd at http:vww.obs-besaicon.fr/iodele., The Besançoon model predictions in the MEGACAM photometric system can be found at http://www.obs-besancon.fr/modele.1080 Simlatious from the Besangoon model in the MEGACAAL photometric svstem have been performed and compared with the CFITTLS iu three fields: Dl. D2 aud D3»," Simulations from the Besançoon model in the MEGACAM photometric system have been performed and compared with the CFHTLS in three fields: D1, D2 and D3."1081 We uxedthe saue maguitude Imt iu / of 21.0 niposed by the stir/ealaxvno separation and took the photometric ΤΟΥΣ in each baud as a function of magnitude iuto account., We used the same magnitude limit in $i'$ of 21.0 imposed by the star/galaxy separation and took the photometric errors in each band as a function of magnitude into account.1082 m he discussion below we indicate also t1ο different coniponents of the €alactic model such as ju disc. thick dise aud the spieroid.," In the discussion below we indicate also the different components of the Galactic model such as thin disc, thick disc and the spheroid."1083 Figure Lt shows 1 ονpecte colour-colour cdiagraie predicted by fie Desancoon model fo rthe D1 fied., Figure \ref{ccd-components} shows the expected colour-colour diagram predicted by the Besançoon model for the D1 field.1084 Indicated are the three lifferent com2010111s. thin disc. tick disc aud halo.," Indicated are the three different components, thin disc, thick disc and halo."1085" While 1e blue part of the colour-colour ciaeranm is populated "" uainlv spheroid stars. the thick cse stars are concentraed wound 0Gκήi'«0.9 aud the thin disc stars cover je red part of he «liagrai."," While the blue part of the colour-colour diagram is populated by mainly spheroid stars, the thick disc stars are concentrated around $0.6 <r'-i'< 0.9$ and the thin disc stars cover the red part of the diagram."1086 Figure 5Ὁ shows the colour-colour diagrams of the CFIITLS for the Dl. D2 and D3 fields respectively as well as the 1uoccl predictiois.," Figure \ref{ccd-all} shows the colour-colour diagrams of the CFHTLS for the D1, D2 and D3 fields respectively as well as the model predictions."1087 Note the excellent overall agreement between observe and predicted colours for the three. Galactic components., Note the excellent overall agreement between observed and predicted colours for the three Galactic components.1088 Figue 5 also shows that the stellar ?opuations of the threedeep fields are simiar. although o16 notices severa differences:," Figure \ref{ccd-all} also shows that the stellar populations of the threedeep fields are similar, although one notices several differences:"1089The measured brightness of ILDG9830 al 22 pam is 270.443.) mJy. with a dispersion of between April 2004 and January 2009.,"The measured brightness of HD69830 at 22 $\mu$ m is $\pm$ 3.1 mJy, with a dispersion of between April 2004 and January 2009."1090 For epochs with observations of both 11D69330 and IIED63146. the [Inx measurements of the stars have dispersions of 1.2 aud1.154... respectively.," For epochs with observations of both HD69830 and HD68146, the flux measurements of the stars have dispersions of 1.2 and, respectively."1091 Although these levels of variability are higher than Hipparcos limits on their visible light dispersion1997).. the similarity between the (wo stars’ dispersions suggests an origin (hat is instrumental. not physical.," Although these levels of variability are higher than Hipparcos limits on their visible light dispersion, the similarity between the two stars' dispersions suggests an origin that is instrumental, not physical."1092 As further evidence of instrumental effects. (he variations between epochs are correlated between the {wo stars.," As further evidence of instrumental effects, the variations between epochs are correlated between the two stars."1093" The strongest deviations come during ""Repeat 3b (2008 Apr 30) when stars are about higher than at other times.", The strongest deviations come during “Repeat 3b” (2008 Apr 30) when stars are about higher than at other times.1094 Hence. for measurements of IID69830 with contemporary observations of IID68146. some additional photometric calibration is possible.," Hence, for measurements of HD69830 with contemporary observations of HD68146, some additional photometric calibration is possible."1095 Nonnalizing the ILD69830 data using the ILD68146 observations where available leads to the corrected values for ILD69530 in Table 3. and Figure 3.., Normalizing the HD69830 data using the HD68146 observations where available leads to the corrected values for HD69830 in Table \ref{IRSPeakup} and Figure \ref{IRSPeakupFig}.1096 The dispersion in (he renormalized data for 1D69330 (total of star + disk) is just (10) which should be taken as an upper limit on the mic-IR variability of the svstem as a whole., The dispersion in the renormalized data for HD69830 (total of star $+$ disk) is just (1 $\sigma$ ) which should be taken as an upper limit on the mid-IR variability of the system as a whole.1097 We discuss the possible variability of the disk emission in Section ??.., We discuss the possible variability of the disk emission in Section \ref{outburstlimits}.1098 Three different epochs of Hilies data of both I1LD69330. and 11ID68416 were collected on 2007 December 20. 2008 December 4. and 2008 December 13.," Three different epochs of HiRes data of both HD69830 and HD68416 were collected on 2007 December 20, 2008 December 4, and 2008 December 13."1099 The first two epochs each included (svo observations at (wo Nod positions with the final epoch including one observation ol two Nod positions for a total of ten spectra of each star., The first two epochs each included two observations at two Nod positions with the final epoch including one observation of two Nod positions for a total of ten spectra of each star.1100 An observation on 2008 April 30 had downlink problems and is not considered further., An observation on 2008 April 30 had downlink problems and is not considered further.1101 Each observation also included an exposure offset from the star to be used for background. ancl bad pixel removal., Each observation also included an exposure offset from the star to be used for background and bad pixel removal.1102 All images were run through to flag and repair bad pixels using the bad pixel masks provided by the Spitzer Science Center (SSC) lor each observational campaign. as well as bad pixels flagged by hand usingSelean.," All images were run through to flag and repair bad pixels using the bad pixel masks provided by the Spitzer Science Center (SSC) for each observational campaign, as well as bad pixels flagged by hand using."1103 After bad pixel repair. the background images were subtracted [rom their corresponding target images.," After bad pixel repair, the background images were subtracted from their corresponding target images."1104 The 1-D spectra were (hen extracted from the reduced images using and its default settings., The 1-D spectra were then extracted from the reduced images using and its default settings.1105 Prior to stitehing together the ten orders within the Long and Short wavelength modules. the 1-D spectra are put through. to remove the low level of interference fringes inherent to the data.," Prior to stitching together the ten orders within the Long and Short wavelength modules, the 1-D spectra are put through to remove the low level of interference fringes inherent to the data."1106 At this point. there were still spikes in the spectra cue to missed bad pixels ancl noise at the red end of each module.," At this point, there were still spikes in the spectra due to missed bad pixels and noise at the red end of each module."1107 To remove (he remaining errant pixels in, To remove the remaining errant pixels in1108we obtain c=MT tl] which leads to ο.ο= 2 | Sor = BAGeV.,we obtain c= ] which leads to 0)= 2m_c + r - = 3.1.1109" With τοι=1.3 (ιο, o0=0.2GeV. and a= 7/12. this results in e=1.56 aud for the radius of the the value ο=0)0.12 fin. which agrees very well with that obtained in the correspondiug (πιαταmechanical study [15].."," With $m_c=1.3$ GeV, $\sigma=0.2$$^2$, and $\alpha=\pi/12$ , this results in $c=1.56$ and for the radius of the the value $r(T=0)\simeq 0.42$ fm, which agrees very well with that obtained in the corresponding quantum-mechanical study \cite{KMS}. ."1110 The imiuimuizatiou requirement (18)) for finite T. leads to SI500 =3f BOT. where apteT) is the effective coupling determined above (see rofal-FU)).," The minimization requirement \ref{min}) ) for finite $T$ leads to = r^2 = _U(r,T), where $\alpha_U(r,T)$ is the effective coupling determined above (see \\ref{al-FU}) )."1111 In5.. we solve this equation eraphically.," In, we solve this equation graphically."1112 It is seen that wp to some teniperature value Tucm15Tho areT) attains a peak large enough to intersect the kinetic term ¢fier. so that there is a nui in the enerev and hence a bound state.," It is seen that up to some temperature value $T_{dis} \simeq 1.5 ~T_c$, $\alpha_U(r,T)$ attains a peak large enough to intersect the kinetic term $c /m_c r$, so that there is a minimum in the energy and hence a bound state."1113 For T2T; this is uo longer the case. the energy decreases monotonically with rs and there is no more bound state.," For $T > T_{dis}$, this is no longer the case, the energy decreases monotonically with $r$ and there is no more bound state."1114 The radius of the surviving is seen to vary very little withtempcrature: it remaius in the rauge 0.5.—0.15 fin up to Tiss. Where the bound state clisappears.," The radius of the surviving is seen to vary very little withtemperature; it remains in the range $0.3 - 0.45$ fm up to $T_{diss}$, where the bound state disappears."1115 The potential C(r.D) used inthe Schrodcinger equation has the correct P-indepeucent form for small r. andcuhauced attractionin the intermediate r range up to Z;;;;.," The potential $U(r,T)$ used inthe Schröddinger equation has the correct $T$ -independent form for small $r$ , andenhanced attractionin the intermediate $r$ range up to $T_{diss}$ ."1116 At large r. it leads to the +-indepenudeut," At large $r$ , it leads to the $r$ -independent"1117the Galactocentric radius of the Sun. and. ey. the circular velocity. of the LSR.,"the Galactocentric radius of the Sun, and $v_0$, the circular velocity of the LSR."1118 We have found that the best-fitting values. considered. separately. are strongly dependent. on the Galaxy model used. to interpret. the data. but that the ratio ryfy is consistently found to lie in the range 208 31L5kms+kpe," We have found that the best-fitting values, considered separately, are strongly dependent on the Galaxy model used to interpret the data, but that the ratio $v_0/R_0$ is consistently found to lie in the range $29.8-31.5\kms\,\kpc^{-1}$."1119 We have also used these data to explore the value of the Sun's peculiar velocitw v.. in light of the recent argument of D09 that the canonical DBOS value is incorrect.," We have also used these data to explore the value of the Sun's peculiar velocity $\vsol$, in light of the recent argument of B09 that the canonical DB98 value is incorrect."1120 We find that these cata support the conclusion that V. is significantly higher than the DBOs value., We find that these data support the conclusion that $V_\odot$ is significantly higher than the DB98 value.1121 By a small but. significant marein. the cata prefer models. with the BOO revision of V. t0 ΣαινI over models in which v. is left as a free parameter (equation 10)).," By a small but significant margin, the data prefer models with the B09 revision of $V_\odot$ to $\sim 11\kms$ over models in which $\vsol$ is left as a free parameter (equation \ref{eq:PB09}) )."1122 The best-fitting models have € and WV. near to the DBOS values. and the even larger value οἱ Ἐν—16.20kms," The best-fitting models have $U_\odot$ and $W_\odot$ near to the DB98 values, and the even larger value of $V_\odot \sim 16 - 20\kms$."1123 7. suggested. that USERS significantly lag circular rotation., \cite{Reidetal2009} suggested that s significantly lag circular rotation.1124 We have investigated. the possibility that the HAISFRs have a twpical peculiar velocity veer and find that the models only constrain the velocity dillerence Veerν.., We have investigated the possibility that the s have a typical peculiar velocity $\vsfr$ and find that the models only constrain the velocity difference $\vsfr-\vsol$.1125 We have argued that models in which the UMsSrRs have large peculiar velocities in the opposite direction to Galactic rotation are neither needed nor plausible., We have argued that models in which the s have large peculiar velocities in the opposite direction to Galactic rotation are neither needed nor plausible.1126 This work. in conjunction with BOO. casts severe doubt on the accuracy of the widely used DB9S value for V...," This work, in conjunction with B09, casts severe doubt on the accuracy of the widely used DB98 value for $V_\odot$."1127 This must be a concern for anvone interested. in the dynamics within the Milky Way because velocities. are. inevitably measured with respect to the Sun., This must be a concern for anyone interested in the dynamics within the Milky Way because velocities are inevitably measured with respect to the Sun.1128 Both observations of more masers and the development of a detailed: dvnanical model of the Galaxy's spiral structure would contribute to establishing more securely what the true value of V. 7. , Both observations of more masers and the development of a detailed dynamical model of the Galaxy's spiral structure would contribute to establishing more securely what the true value of $V_\odot$ \cite{Rygletal2009} 11292009).,2009).1130 Phe (0.0) component is static and does not contribute to tidal dissipation.," The $(0,0)$ component is static and does not contribute to tidal dissipation."1131" Phe (0m.m)-component is physically identical. to the (m.m, )-component."," The $(-m,-m')$ -component is physically identical to the $(m,m')$ -component."1132" Each Gnai')-component of the tidal potential drives. [uid »erturbation inside the star. which can be specified by the Lagrangian displacement. £,,,,,(v./) and Eulerian density »erturbation 9p,(αι£D)."," Each $(mm')$ -component of the tidal potential drives fluid perturbation inside the star, which can be specified by the Lagrangian displacement $\bxi_\mmp(\br,t)$ and Eulerian density perturbation $\delta\rho_\mmp(\br,t)$."1133 In the absence ofdissipation. these »erturbations are proportional to (Crfios)Daundial4E (where wy—COMIP is the dynamical frequency of the star). exactly in phase with the tidal potential.," In the absence of dissipation, these perturbations are proportional to $(U_\mmp/\omega_0^2) e^{im\phi-im'\Omega t}$ (where $\omega_0\equiv1134\sqrt{GM/R^3}$ is the dynamical frequency of the star), exactly in phase with the tidal potential."1135 When tuicl dissipation is present. there will be phase shift. between he fluid. perturbation and the tidal potential.," When fluid dissipation is present, there will be phase shift between the fluid perturbation and the tidal potential."1136 This phase shift. in general. depends on mm. the forcing frequency. (in he rotating frame) ρω. as well as the intrinsic property (including the rotation rate) of the star.," This phase shift, in general, depends on $m$, the forcing frequency (in the rotating frame) $\tomega_\mmp$, as well as the intrinsic property (including the rotation rate) of the star."1137" We write this phase shift as(AS) where /,,,,7"" is: the ""lag time"".- associated. with. the ; (mnm)-perturbation."," We write this phase shift as, where $t_\mmp$ is the “lag time” associated with the $(mm')$ -perturbation."1138 Thus. without loss of generality. we can write ; ⋅ ↿∖⊔⋯∣⊐⊣∼∪⊔↓↓≻∪⊔∢⊾⊔↿∪⇂⇂↓↥∢⊾∐," Thus, without loss of generality, we can write $(mm')$ -component of the fluid displacement as )."1139⋯∠⇂∠⊔⊳∖↓≻↓⋯∼⋖⋅⊔↓⋖⋅⊔↿⋜↧⊳∖⋠⋠ ad (53) with ," The corresponding density perturbation is ), with )."1140"Note that £,,,,(r) and op,,,,;(v) are proportional to «HE τ except for this factor. they are real functions."," Note that $\bar\bxi_\mmp(\br)$ and $\delta\bar\rho_\mmp(\br)$ are proportional to $e^{im\phi}$ ; except for this factor, they are real functions."1141 Also note that with Mouου Th," Also note that =, with $\cU_\mmp\equiv W_{2m'}\cD_\mmp(\Theta)$ ."1142"us £09). and δρ) specify the amplitudes (and. shapes) the Iuid perturbations after factoring out the ""equilibrium dimensionless tidal distortion. (ALου."," Thus, $\bar\bxi_\mmp(\br)$ and $\delta\bar\rho_\mmp(\br)$ specify the amplitudes (and shapes) the fluid perturbations after factoring out the “equilibrium” dimensionless tidal distortion, $(M'/M)(R/a)^3$."1143" Note that although our ansatz of the tidal responses is motivated by the weakfriction theory of equilibrium tides. it is actually quite general. provided that one keeps in mind that the lag time /,,,,; depends on m. c,"" and the intrinsic property of the star."," Note that although our ansatz of the tidal responses is motivated by the weakfriction theory of equilibrium tides, it is actually quite general, provided that one keeps in mind that the lag time $t_\mmp$ depends on $m$, $\tomega_\mmp$ and the intrinsic property of the star."1144 The tidal torque on the star is (2.1)) and. (2.2)). the z- (along the spin axis) of the tidal torque reduceslo (73) where PR and 5nennt is (45)the “Love coellicient”:," The tidal torque on the star is ) and \ref{eq:deltarho}) ), the $z$ -component (along the spin axis) of the tidal torque reducesto^2 , where )^2R^5, and $\kappa_\mmp$ is the “Love coefficient”:"1145ancl others).,and others).1146 The unique dispersion. measures identified them: as astronomical signals. and. the peak (ux densities (ranging from mmy to 3.6JJ.) made these transients among the brightest radio sources in the universe.," The unique dispersion measures identified them as astronomical signals, and the peak flux densities (ranging from mJy to Jy) made these transients among the brightest radio sources in the universe."1147 The most natural interpretation of these sources is that they are rotating neutron stars (NSs) which emit radio bursts sporaclically: hence the name Rotating RAcio Transients (RATS)., The most natural interpretation of these sources is that they are rotating neutron stars (NSs) which emit radio bursts sporadically: hence the name Rotating RAdio Transients (RRATs).1148 Based on their ephemeral nature. the total number of ULNIS in the Galaxy (~ 107: 2008)) might. exceed that of the radio pulsars. hus representing a large fraction of the population of neutron stars with age smaller than 10* vr.," Based on their ephemeral nature, the total number of RRATs in the Galaxy $\sim 10^5$ ; ) might exceed that of the radio pulsars, thus representing a large fraction of the population of neutron stars with age smaller than $\sim 10^7$ yr."1149 That calls fora deep investigation of these objects. aimed at understanding heir emission properties. environments. ancl evolutionary inks (i£ any) with other classes of neutron stars.," That calls for a deep investigation of these objects, aimed at understanding their emission properties, environments, and evolutionary links (if any) with other classes of neutron stars."1150 Lt has been ooposed Chat RIGVES could be an extreme manifestation of orocesses already observed in at least some ordinary pulsars: c.g. giant pulses 2006).. nulling2007). or the fading of the radio signal of an old pulsar while approaching the pair-production death line2007).," It has been proposed that RRATs could be an extreme manifestation of processes already observed in at least some ordinary pulsars: e.g., giant pulses , nulling, or the fading of the radio signal of an old pulsar while approaching the pair-production death line."1151.. A particularly interesting hypothesis is that RRATs may be distant midcle-aged. (spin-down ages 15.LO? vvr) ordinary. pulsars (such as PSR. BOGHSG|14). whose emission is only detectable when particularly bright pulses oceur2006).," A particularly interesting hypothesis is that RRATs may be distant middle-aged (spin-down ages $1-5\times 10^5$ yr) ordinary pulsars (such as PSR B0656+14), whose emission is only detectable when particularly bright pulses occur."1152.. Aceorcing to other models. the RRATPS co not belong to the population of the rotation-powered. pulsars: instead. associations with the steadilv-emitting magnetars 2006).. transient magnetars2009). or isolated neutron stars (INS: 2006)) were suggested.," According to other models, the RRATs do not belong to the population of the rotation-powered pulsars: instead associations with the steadily-emitting magnetars , transient magnetars, or thermally-emitting isolated neutron stars (INS; ) were suggested."1153 Even more exotic hypotheses explain the bursting behavior of the Αν with the sporadic ellects of an asteroid or radiation belt surrounding the pulsar)., Even more exotic hypotheses explain the bursting behavior of the RRATs with the sporadic effects of an asteroid or radiation belt surrounding the pulsar.1154. The observation of the counterpart to a RAT in the X- band could be crucial for discriminating between many of the models reported above. since they predict distinct N-rav spectral signatures for the source.," The observation of the counterpart to a RRAT in the X-ray band could be crucial for discriminating between many of the models reported above, since they predict distinct X-ray spectral signatures for the source."1155" Unfortunately. only ((with spin period 2= 4.26ss. dipole magnetic field B=5. ντος. iage τ=1.210"" vr. and spin- luminosity characteristicE=3.10732meres+: for definitions of these parameters see 2004)) has been detected at high energies."," Unfortunately, only (with spin period $P=4.26$ s, dipole magnetic field $B= 5\times 10^{13}$ G, characteristic age $1156\tau=1.2\times 10^5\,$ yr, and spin-down luminosity $\dot E= 3\times115710^{32}\,\ergsec$; for definitions of these parameters see ) has been detected at high energies."1158 Ht was serendipitously found in oobservations of the (almost certainly unrelated) supernova remnant 15.9|0.2., It was serendipitously found in observations of the (almost certainly unrelated) supernova remnant G15.9+0.2.1159 Additional long pointed. integrations revealed that hhas a thermal spectrum with AZ= O14kkeV and an unabsorbed N-ray luminosity Ly=d105eres 2007).. which is larger than the spin- luminosity.," Additional long pointed integrations revealed that has a thermal spectrum with $kT=0.14$ keV and an unabsorbed X-ray luminosity $L_X \approx11604\times10^{33}\,\ergsec$ , which is larger than the spin-down luminosity."1161 The spectrum. dillers from that of the maenetars. which are hotter2006).. although it may be similar (at. times) to that of the transient magnetar NTE 1510 197 0.15.—Q.IS kkeV: 2005).," The spectrum differs from that of the magnetars, which are hotter, although it may be similar (at times) to that of the transient magnetar XTE $-$ 197 $kT \sim 0.15-0.18$ keV; )."1162 Instead. the spectrum of lis comparable in temperature and X-ray luminosity to the similarlv-aged. radiopulsar DB0656|14 (which itself is consistent with being a nearby RRAL: see above)," Instead, the spectrum of is comparable in temperature and X-ray luminosity to the similarly-aged radiopulsar B0656+14 (which itself is consistent with being a nearby RRAT; see above)."1163 It is also similar to the slightly older radio-quiet INS (AZ—0.050.1 keV at ages of zz0.5 Myr: see 2008))., It is also similar to the slightly older radio-quiet INS $kT=0.05-0.1$ keV at ages of $\approx 0.5$ Myr; see ).1164 In addition. the spectrum displays à broad. absorption line resembling the lines detected in many of the INS2007).," In addition, the spectrum displays a broad absorption line resembling the lines detected in many of the INS."1165. The case of illustrates the potentiality of X-ray observations in unveiling possible connections between RRATS and other NS populations., The case of illustrates the potentiality of X-ray observations in unveiling possible connections between RRATs and other NS populations.1166 Llowever. until recently these studies have been significantly limited by the lack of precise positions for most of the sources.," However, until recently these studies have been significantly limited by the lack of precise positions for most of the sources."1167 wwas one of the only three RRATLS whose celestial coordinates hack been precisely. determined. through pulse timing (e.g.. as discussed in 2004)).," was one of the only three RRATs whose celestial coordinates had been precisely determined through pulse timing (e.g., as discussed in )."1168 ALL the other IIGXTS were only localised to within the primary beam of the Parkes telescope: z1d aremin diameter for the observing frequency of Εαν.," All the other RRATs were only localised to within the primary beam of the Parkes telescope: $\approx 14\,$ arcmin diameter for the observing frequency of GHz."1169 That. prevented follow-up multi-wavelength observations and. in turn. the possibility of constraining the spectrum of the RATS and their primary CHerev stores.," That prevented follow-up multi-wavelength observations and, in turn, the possibility of constraining the spectrum of the RRATs and their primary energy stores."1170 Dedicated timing observations presented in have more than doubled the sub-sample of the RRAPS with phase-connected. timing solutions ancl hence high precision rotational cid astrometriepeusunoeters?., Dedicated timing observations presented in have more than doubled the sub-sample of the RRATs with phase-connected timing solutions and hence high precision rotational and astrometric.1171. ‘Pwo among the RRATs with à new timing solution aand 02573) are Located in the P?P diagram in a region almost devoid of ordinary pulsars (see Section 2. for timing parameters) but close to the area occupied. by. the INS. lending support to the hypothesis2008: 20003) that they could be transition objects between ordinary pulsars and. INS.," Two among the RRATs with a new timing solution and ) are located in the $P-\dot{P}$ diagram in a region almost devoid of ordinary pulsars (see Section \ref{sec:obs} for timing parameters) but close to the area occupied by the INS, lending support to the hypothesis; ) that they could be transition objects between ordinary pulsars and INS."1172 Inspecting the P?p diagram: (see e.g.. Figure 4 in 20093). aand aalso appear the closest RRATS to the. location of1458.. the only RRAT detected in X-ray. so far (see above).," Inspecting the $P-\dot{P}$ diagram (see e.g., Figure 4 in ), and also appear the closest RRATs to the location of, the only RRAT detected in X-ray so far (see above)."1173 Here. we present searches for. X-ray emission. [rom these two particularly promising sources. based on targeted and archivalChandra observations.," Here we present searches for X-ray emission from these two particularly promising sources, based on targeted and archival observations."1174 In. what. follows. all Iuminosities are corrected. for interstellar absorption.," In what follows, all luminosities are corrected for interstellar absorption."1175 We searched for an X-ray counterpart to (P— 5918s. D—310 GG. 7=US Myr. and=2 +) with the 2000).. as summamarised. in “Table," We searched for an X-ray counterpart to $P=5.978\,$ s, $B=\expnt{3}{13}\,$ G, $\tau=0.8\,$ Myr, $\dot1176E=\expnt{2}{31}\,\ergsec$ ) with the , as summarised in Table \ref{tab:obs}. ."1177 , 1178one.,one.1179 As {ος Increases. however. the expected spectral index trends get flatter and flatter and a best fit is found for 0.1.," As $t_{\rm OFF}/t_{\rm s}$ increases, however, the expected spectral index trends get flatter and flatter and a best fit is found for $t_{\rm OFF}/t_{\rm s}=0.8\pm0.1$ ."1180 The behavior of the spectral index seen in B2 0120433 and B2 1610-29 could be typical of many dying sources., The behavior of the spectral index seen in B2 0120+33 and B2 1610+29 could be typical of many dying sources.1181 In practice. as a result of the jets switch-off. any pre-existent spectral index gradient along the lobes is quickly canceled since the break frequency reaches roughly the same value in each part of the source.," In practice, as a result of the jets switch-off, any pre-existent spectral index gradient along the lobes is quickly canceled since the break frequency reaches roughly the same value in each part of the source."1182 We can therefore expect that extreme dying sources. those for which {ος=1. are characterized by very uniform spectral index distributions along the fading lobes.," We can therefore expect that extreme dying sources, those for which $t_{\rm OFF}/t_{\rm s}\simeq 1$, are characterized by very uniform spectral index distributions along the fading lobes."1183 As the dying source gets older. the spectral index increases systematically but with small variations from point to point.," As the dying source gets older, the spectral index increases systematically but with small variations from point to point."1184 We performed the minimum energy calculation for the five dying galaxies by considering the source power at 151 MHz. where the energy losses of the synchrotron electrons are less dramatic. and the volumes as measured from the VLA images at arcsecond resolution.," We performed the minimum energy calculation for the five dying galaxies by considering the source power at 151 MHz, where the energy losses of the synchrotron electrons are less dramatic, and the volumes as measured from the VLA images at arcsecond resolution."1185 In the calculation of the sources power we made use of the 151 MHz flux density after the deconvolution of confusing sources., In the calculation of the sources power we made use of the 151 MHz flux density after the deconvolution of confusing sources.1186 The resulting monochromatic powers in the sources rest frame. Ls). are listed in 66 and place them below the FRI-FRII division (Fanaroff Riley 1974).," The resulting monochromatic powers in the sources rest frame, $L_{151}$, are listed in 6 and place them below the FRI-FRII division (Fanaroff Riley 1974)."1187 We assume that the radio sources contain relativistic particles and magnetic fields uniformly distributed and in energy equipartition conditions., We assume that the radio sources contain relativistic particles and magnetic fields uniformly distributed and in energy equipartition conditions.1188" The equipartition parameters (magnetic field B; and energy density myin) are generally computed assuming that the relativistic particle energies are confined between a minimum e, and a maximum €hivh- corresponding to the observable radio frequency range. typically 10 MHz - 100 GHz (see. e.g.. Pacholezyk 1970)."," The equipartition parameters (magnetic field $B_{\rm min}$ and energy density $u_{\rm min}$ ) are generally computed assuming that the relativistic particle energies are confined between a minimum $\epsilon_{\rm low}$ and a maximum $\epsilon_{\rm high}$, corresponding to the observable radio frequency range, typically 10 MHz - 100 GHz (see, e.g., Pacholczyk 1970)."1189 This choice minimizes the source energetics required by the observed radiation in the radio band., This choice minimizes the source energetics required by the observed radiation in the radio band.1190 However. a fixed frequency range corresponds to an energy range that depends on the source magnetic field. which may change from source to source.," However, a fixed frequency range corresponds to an energy range that depends on the source magnetic field, which may change from source to source."1191 A fixed frequency range computation would miss the contribution from lower energy electrons. since the ej. corresponding to 10 MHz ts larger than 200 MeV for Bj<30uG (Brunetti et al.," A fixed frequency range computation would miss the contribution from lower energy electrons, since the $\epsilon_{low}$ corresponding to 10 MHz is larger than 200 MeV for $B_{\rm min} \le 30~\mu G$ (Brunetti et al."1192 1997 and Beck Krause 2005)., 1997 and Beck Krause 2005).1193" Because of this we have computed the equipartition parameters assuming a fixed low energy cutoff e, = 10 MeV. The high energy ey, cut-off is chosen to match the spectral break in the emission spectrum. vp. since no electrons are present in the source beyond this limit. see Parma et al. ("," Because of this we have computed the equipartition parameters assuming a fixed low energy cutoff $\epsilon_{\rm low}$ = 10 MeV. The high energy $\epsilon_{\rm high}$ cut-off is chosen to match the spectral break in the emission spectrum, $\nu_{\rm b}$, since no electrons are present in the source beyond this limit, see Parma et al. ("11942007) for further details.,2007) for further details.1195 The equipartition parameters computed in this way are reported in Tab., The equipartition parameters computed in this way are reported in Tab.1196 6., 6.1197" A comparison with fixed frequency range calculations shows that the values of our B, are larger by up to a factor two.", A comparison with fixed frequency range calculations shows that the values of our $B_{\rm min}$ are larger by up to a factor two.1198 Assuming a constant magnetic field and neglecting expansion losses. the total source age can be calculated from the break frequency. vp: where the synchrotron age ἐς is in Myr. the magnetic field in μα. the break frequency vy in GHz. while the inverse Compton equivalent field (see Sect.," Assuming a constant magnetic field and neglecting expansion losses, the total source age can be calculated from the break frequency, $\nu_{b}$: where the synchrotron age $t_{\rm s}$ is in Myr, the magnetic field in $\mu$ G, the break frequency $\nu_{\rm b}$ in GHz, while the inverse Compton equivalent field (see Sect."1199 5.2)., 5.2).1200" By adopting the equipartition value B,,, for the magnetic field strength we can thus derive the synchrotron age £,.", By adopting the equipartition value $B_{\rm min}$ for the magnetic field strength we can thus derive the synchrotron age $t_{\rm s}$.

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