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The Chandra//ACIS-S observation (performed on 2002 March 03, when the source was rather bright) has been already published in Sanwaletal. (2002),, the XMM-Newton//EPIC (2005 | The /ACIS-S observation (performed on 2002 March 03, when the source was rather bright) has been already published in \citet{sanwal02}, , the /EPIC (2005 |
ens and hundreds of vears. depending on the physical size of the wind launching region and the SMDLIL mass. | tens and hundreds of years, depending on the physical size of the wind launching region and the SMBH mass. |
Hlowever. observations (see €3 below) suggest that he structure responsible for AGN obscuration should. be optically thick over a Large (e.g. a third) fraction of the whole sky. | However, observations (see 3 below) suggest that the structure responsible for AGN obscuration should be optically thick over a large (e.g., a third) fraction of the whole sky. |
Our moclels cannot achieve such a high average ine of sight column depth unless the mass loss rate in he wind is super-Edeington. | Our models cannot achieve such a high average line of sight column depth unless the mass loss rate in the wind is super-Eddington. |
In fact. we show via simple analvtical estimates that this conclusion is valid. whatever he nature of the wind driving mechanism is. | In fact, we show via simple analytical estimates that this conclusion is valid whatever the nature of the wind driving mechanism is. |
We feel it would oe very hard. to produce highly. super-IExldington winds in a plausible quasi steady-state AGN model for a typical (very sub-Eddington luminositv) GN. | We feel it would be very hard to produce highly super-Eddington winds in a plausible quasi steady-state AGN model for a typical (very sub-Eddington luminosity) AGN. |
Thus we suggest hat winds. while very important for the AGN phenomenon. cannot replace the "torus? in its role in the unification schemes of AGN. | Thus we suggest that winds, while very important for the AGN phenomenon, cannot replace the “torus” in its role in the unification schemes of AGN. |
Either the torus does exist. despite the dillieulties faced by the models. or another optically thick structure. e... a warped accretion disce plavs its role in ACN. | Either the torus does exist despite the difficulties faced by the models, or another optically thick structure, e.g., a warped accretion disc plays its role in AGN. |
AGN clises are expected to be massive ancl thusself-eravitating at large enough distances (7?20.010.1 pc) from the super-massive black holes (SMBID. | AGN discs are expected to be massive and thusself-gravitating at large enough distances $R \simgt 0.01-0.1$ pc) from the super-massive black holes (SMBH). |
Lo the disc cooling time is short. gravitational collapse ancl formation of stars or even planets is predicted. (e.g...2???).. | If the disc cooling time is short, gravitational collapse and formation of stars or even planets is predicted \citep[e.g.,][]{Paczynski78,Kolykhalov80,Shlosman89,Gammie01}. |
Young massive stars in the central parsee of our own Galaxy were most likely created in this way (2). with an apparent significant over-abundance of high mass stars (7). over their raction in “normal galactic” star formation. | Young massive stars in the central parsec of our own Galaxy were most likely created in this way \citep{Paumard06}, with an apparent significant over-abundance of high mass stars \citep{NS05} over their fraction in “normal galactic” star formation. |
As stars are »orn from the accretion disc. they launch powerful radiation ields and stellar winds. some of which will break through he disc. | As stars are born from the accretion disc, they launch powerful radiation fields and stellar winds, some of which will break through the disc. |
Low mass proto-stars may prove equally effective in aunching winds in these circumstances as the rates at which gas is captured from the disk into the Lill (or capture) zones of the protostars are super-IExdcdington (7).. | Low mass proto-stars may prove equally effective in launching winds in these circumstances as the rates at which gas is captured from the disk into the Hill (or capture) zones of the protostars are super-Eddington \citep{Nayakshin06a}. |
In addition. when he gas in the disk is depleted to about half its initial surface density. stellar velocity. dispersion starts to grow. and soon he stellar disc becomes geometrically thicker than the gas disc (2).. | In addition, when the gas in the disk is depleted to about half its initial surface density, stellar velocity dispersion starts to grow, and soon the stellar disc becomes geometrically thicker than the gas disc \citep{Nayakshin06a}. |
Stellar winds then escape from above the disc clirectLv. | Stellar winds then escape from above the disc directly. |
We attempt. to simulate this complex situation in a simplified. setting. concentrating only on the wind part of it. | We attempt to simulate this complex situation in a simplified setting, concentrating only on the wind part of it. |
For this reason we do not include the gas in the negligibly thin and Dat accretion disc from which the stars are born. | For this reason we do not include the gas in the negligibly thin and flat accretion disc from which the stars are born. |
To capture a degree of the expected. diversity of the stellar populations ancl conditions in this problem. the stars are divided into two groups. | To capture a degree of the expected diversity of the stellar populations and conditions in this problem, the stars are divided into two groups. |
The first. produces winds with velocities ον=300 Km/sec. whereas the second has (y=TOO km/sec. | The first produces winds with velocities $v_{\rm w} = 300$ km/sec, whereas the second has $v_{\rm w} = 700$ km/sec. |
Both types of stars have mass loss rates of 2.510M. vear and are situated in a Lat circularly rotating Ixeplerian disk of geometrical thickness H(R)=OAR (the unit of length. used here is 1 aresecond at the distance of S kpe. which is about 1.2041 em or 0.04 pc). | Both types of stars have mass loss rates of $2.5 \times 10^{-4} \msun$ $^{-1}$, and are situated in a flat circularly rotating Keplerian disk of geometrical thickness $H(R) = 0.1 R$ (the unit of length used here is 1 arcsecond at the distance of 8 kpc, which is about $1.2\times 10^{17}$ cm or 0.04 pc). |
In total. we have 200 mass shredcing stars. thus amounting to the nass loss rate of 0.1. M. +. which is à factor of few. super-IExldington. for a SALBLE mass of Moy=3510"M... | In total, we have 200 mass shredding stars, thus amounting to the mass loss rate of 0.1 $\msun$ $^{-1}$, which is a factor of few super-Eddington for a SMBH mass of $\mbh = 3.5 \times 10^6 \msun$. |
Phe disc inner ancl outer radii are Rin=15 and Raw=8. respectively, | The disc inner and outer radii are $R_{\rm in} = 1.5$ and $R_{\rm out} = 8$, respectively. |
The stellar surface clensity follows thelaw “i.(2)xFR and the stars are in the Keplerian circular rotation around the SALBIL. | The stellar surface density follows thelaw $\Sigma_*(R) \propto R^{-1}$, and the stars are in the Keplerian circular rotation around the SMBH. |
The initial phase (i.e. the ó-coordinate) of the stars in the clisk is eenerated randomly. as is the vertical coordinate within the df to L4 limits. | The initial phase (i.e. the $\phi$ -coordinate) of the stars in the disk is generated randomly, as is the vertical coordinate within the $-H$ to $+H$ limits. |
Initially. the computational domain is Glled with hot tenuous gas with velocity greater than the escape velocity. | Initially, the computational domain is filled with hot tenuous gas with velocity greater than the escape velocity. |
This gas quickly outflows from the region. | This gas quickly outflows from the region. |
At the end of the simulation stellar winds exceed the initial eas mass by a factor of few tens. | At the end of the simulation stellar winds exceed the initial gas mass by a factor of few tens. |
‘To perform simulations. we use the Gadget-2 code (?) moclified to model stellar winds as new SPLL particles ejected from the stellar wind sources. | To perform simulations, we use the Gadget-2 code \citep{Springel05} modified to model stellar winds as new SPH particles ejected from the stellar wind sources. |
These methods were described and thoroughly tested in ?7).. | These methods were described and thoroughly tested in \cite{Cuadra05,Cuadra06}. . |
A snapshot of the simulation is shown in Figure 1. at time /&2200 vears alter the beginning of the simulation. which corresponds to 36 dynamical times at 2=1. | A snapshot of the simulation is shown in Figure \ref{fig:fig1} at time $t
\approx 2200$ years after the beginning of the simulation, which corresponds to $\sim 36$ dynamical times at $R=1$. |
While there is no true steady state in this finite number of moving stars system. the snapshot is fairly typical of the morphology of the stellar wind. | While there is no true steady state in this finite number of moving stars system, the snapshot is fairly typical of the morphology of the stellar wind. |
The face-on view (left panel) shows that some of the shocked wind managed to cool down and formed a small-scale clisc. | The face-on view (left panel) shows that some of the shocked wind managed to cool down and formed a small-scale disc. |
Ehe inner part of the simulation domain is the most likely. place for the dise to form since stellar wind density is highest there. leading to shocks. thermalization and rapid cooling: the frequent. collisions of gaseous clumps "waste" momentum of the winds. and finally. the escape velocity from the region is around. GOO km/sec. Le. higher than wind. velocities of the slower winds. | The inner part of the simulation domain is the most likely place for the disc to form since stellar wind density is highest there, leading to shocks, thermalization and rapid cooling; the frequent collisions of gaseous clumps “waste” momentum of the winds, and finally, the escape velocity from the region is around 600 km/sec, i.e. higher than wind velocities of the slower winds. |
With time the disc erows racdiallv to both smaller and larger racii. | With time the disc grows radially to both smaller and larger radii. |
However. the formation of the dise (οἱ.2). is physically significant only in the case when the gaseous accretion disc from which the stars were born iscafiredy absent. as is the case presently in the Galactic Centre. | However, the formation of the disc \citep[cf.][]{Cuadra06} is physically significant only in the case when the gaseous accretion disc from which the stars were born is absent, as is the case presently in the Galactic Centre. |
In the opposite case the cooled disc and the “spiral armis" seen at intermediate raclit would simply blend in with the much more massive underlving eas accretion disc. | In the opposite case the cooled disc and the “spiral arms” seen at intermediate radii would simply blend in with the much more massive underlying gas accretion disc. |
Contrary to the frequent collisions and. high escape velocity in the inner region of the computational domain. 10 conditions in the outer regions allow direct. escape of both the faster cilluse and the slower cooler clumpy winds. | Contrary to the frequent collisions and high escape velocity in the inner region of the computational domain, the conditions in the outer regions allow direct escape of both the faster diffuse and the slower cooler clumpy winds. |
Due to the final extent of the stellar disc and the projection ellect. the wind morphology reminds and 7X-shape (Figure l.. right panel). | Due to the final extent of the stellar disc and the projection effect, the wind morphology reminds and “X”-shape (Figure \ref{fig:fig1}, right panel). |
The edge-on view (right panel) of the stellar disc shows that the plane of the gas disc at /?X2 is very 4ighthy tilted with respect to the orbital plane of the stellar disc. | The edge-on view (right panel) of the stellar disc shows that the plane of the gas disc at $R\simlt 2$ is very slightly tilted with respect to the orbital plane of the stellar disc. |
This is an artefact of the initial conditions. | This is an artefact of the initial conditions. |
The left panel of Figure 2. shows the obscuring column epth of the winds as seen from the SMDII. | The left panel of Figure \ref{fig:fig2} shows the obscuring column depth of the winds as seen from the SMBH. |
Since the orientation. of the tiltecl inner disc is influencecl by the initial conditions. we eliminate its obscuration. only plotting eas with radial distances #21.6 [rom the SMDII. which excludes all of the disc. | Since the orientation of the tilted inner disc is influenced by the initial conditions, we eliminate its obscuration, only plotting gas with radial distances $R \simgt 1.6$ from the SMBH, which excludes all of the disc. |
Besides. hack we included. in the simulations the more massive gas accretion disk whose midplane would coincide with that of the stars. the inner disc seen in Figure 1. would have completely blended in with the accretion disce and would not present much obscuration at all. | Besides, had we included in the simulations the more massive gas accretion disk whose midplane would coincide with that of the stars, the inner disc seen in Figure \ref{fig:fig1} would have completely blended in with the accretion disc and would not present much obscuration at all. |
Note the very irregular patchy structure of the (vellow) optically thicker regions. | Note the very irregular patchy structure of the (yellow) optically thicker regions. |
The contrast. between those and neighbouring less dense patches of sky is frequently a factor of 10 or more. | The contrast between those and neighbouring less dense patches of sky is frequently a factor of 10 or more. |
Since the winds are rotating at a fraction of the local angular frequency. ©. this implies that the columndepth sampled. by the observer will vary on time scales as short as ~10 to 107 of 1/O. Le. | Since the winds are rotating at a fraction of the local angular frequency, $\Omega$ , this implies that the columndepth sampled by the observer will vary on time scales as short as $\sim 10^{-3}$ to $10^{-2}$ of $1/\Omega$ , i.e. |
We have shown that there is a source of hard. power-Iaw-like. A-ray emission in the nuclear regions of Arp 220. | We have shown that there is a source of hard, power-law-like, X-ray emission in the nuclear regions of Arp 220. |
This source is extended EW. consistent with the emission coming [rom both the radio/IR. nuclei. | This source is extended EW, consistent with the emission coming from both the radio/IR nuclei. |
The central concentration of hard X-ray emission in Arp 220 is in contrast to other interacting galaxies. where hard emission comes from chunps distributed across niuch larger physical distances. | The central concentration of hard X-ray emission in Arp 220 is in contrast to other interacting galaxies, where hard emission comes from clumps distributed across much larger physical distances. |
This difference may be associated with the merger in Arp 220 being older than in (he Antennae or NGCC3256. and that compact objects have sunk lo its core (Tremaine. Ostriker Spitzer. 1975) but may also be associated with Arp 2205 ereater Iuminosity. | This difference may be associated with the merger in Arp 220 being older than in the Antennae or NGC3256, and that compact objects have sunk to its core (Tremaine, Ostriker Spitzer, 1975) but may also be associated with Arp 220's greater luminosity. |
The origin of the hard emission is unclear. | The origin of the hard emission is unclear. |
Its spectrum is unlikely to be produced by voung supernovae. but inverse Compton emission. albeit of very. low efficiency. accretion onto ultra-Iuminous X-ray binaries or onto an AGN are all possible. | Its spectrum is unlikely to be produced by young supernovae, but inverse Compton emission, albeit of very low efficiency, accretion onto ultra-luminous X-ray binaries or onto an AGN are all possible. |
If there is an AGN contribution. it has too low a luminosity for it to play a significant role in (he energetics | If there is an AGN contribution, it has too low a luminosity for it to play a significant role in the energetics |
the break time for a collimated jet is where /j is the fiducial break time cefinecl in Rhoads (1999). | the break time for a collimated jet is where $t_b$ is the fiducial break time defined in Rhoads (1999). |
The break in this preclictecl light curve isH extremely broad. auc would giveH 47> little. better than a single power law in fitting the observed break in either A" or 2 band. | The break in this predicted light curve is extremely broad, and would give $\chi^2$ little better than a single power law in fitting the observed break in either $K'$ or $R$ band. |
The model curve is based on numerical integration of the remuzaut's dvuaumieal equations. aud ignores differences iu light travel time between the center aud edge of the remuant. which will ouly smooth the break further (e.g:.. Mocderski et al 2000: Panaitescu Meszaros 1990). | The model curve is based on numerical integration of the remnant's dynamical equations, and ignores differences in light travel time between the center and edge of the remnant, which will only smooth the break further (e.g., Moderski et al 2000; Panaitescu Meszaros 1999). |
If we ignore tlie issue of break sharpness aud fit a collimated jet moclel to the observed R baud light curve. we cau infer the opening auele from the measured yeak time. | If we ignore the issue of break sharpness and fit a collimated jet model to the observed R band light curve, we can infer the opening angle from the measured break time. |
To do so. we need a reasonable measurement of /; and crude estimates of O/E aud n. since the inferred opening augle scales as (nOE)* (Rhoads 1999). | To do so, we need a reasonable measurement of $t_b$ and crude estimates of $\Omega/E$ and $n$, since the inferred opening angle scales as $(t_b^3 n \Omega/E )^{1/8}$ (Rhoads 1999). |
We use /j=5.1 days and E=3x107ergO/(Ix). | We use $t_b = 5.1$ days and $E = 3 \times 10^{53} \erg \times \Omega / (4
\pi)$ . |
To estimate i lore precisely. we use the column deusity WCHL)z-1022-5>D)— Doinferred+ from" Lyman a absorption (Jensen et al 2000) and estimate the linear size of tle source as Z50.2" based on its uoudetection in late HST images. | To estimate $n$ more precisely, we use the column density $N(HI) \approx 10^{21.2 \pm 0.5}$ inferred from Lyman $\alpha$ absorption (Jensen et al 2000) and estimate the linear size of the source as $\la 0.2''$ based on its nondetection in late HST images. |
This implies a number cleusity 7120.L consistent with our earlier estimate. | This implies a number density $n \ga 0.4$ , consistent with our earlier estimate. |
Usiug »zz|. the iulerred opening augle becomes 2.5°. or LO? of the sky if the jet is bipolar. | Using $n \approx 1$, the inferred opening angle becomes $2.5^\circ$, or $10^{-3}$ of the sky if the jet is bipolar. |
The trausitiou to the nonrelativistic regime lias been proposed as auotlier mechauisin for light curve breaks both in this burst (Dai Lu 2000) and others (Dai Lu 1999). | The transition to the nonrelativistic regime has been proposed as another mechanism for light curve breaks both in this burst (Dai Lu 2000) and others (Dai Lu 1999). |
However. we clo uot know of a detailed calculation of the sharpuess of this break. makiug a fair evaluation of this possibility diffieult. | However, we do not know of a detailed calculation of the sharpness of this break, making a fair evaluation of this possibility difficult. |
Light travel time effects seem likely to broaden this feature to 07//1. as with most other features. | Light travel time effects seem likely to broaden this feature to $\delta t / t \sim 1$, as with most other features. |
A final possible cause for sharp breaks in CRB alterelow light curves is discontinuities iu the ambient density distribution. | A final possible cause for sharp breaks in GRB afterglow light curves is discontinuities in the ambient density distribution. |
Assuming that the density is a fuiction of radius alone. a 1ninimun timescale for breaks due to such discoutinuities is A/o=/. wlvere / is the time elapsed in the observers frame siuce the burst aud .M the characteristic duraion of a light curve feature. | Assuming that the density is a function of radius alone, a minimum timescale for breaks due to such discontinuities is $\Delta t \ga t$, where $t$ is the time elapsed in the observer's frame since the burst and $\Delta t$ the characteristic duration of a light curve feature. |
This duratiou is set by differential light. travel time effects between uaterial moving alongthe line of sight and off-axismaterial moving iu direction 1/E. | This duration is set by differential light travel time effects between material moving alongthe line of sight and off-axismaterial moving in direction $1/\Gamma$ |
into account the neighbours of the objects Iving at least 5 degrees far from the boundary (we called this sample). | into account the neighbours of the objects lying at least 5 degrees far from the boundary (we called this sample). |
Fie. | Fig. |
7 shows the surface number density. of neighbours dd=DDπιλσ(ᾳσ|Ao)] in the initial sample (circles) and per—RRwAo(20|Ao) in the random one (triangles) correspondinglv. where we choose the bin size Aq1. | \ref{fig_3_4} shows the surface number density of neighbours $dd=DD/[\pi\Delta\sigma(2\sigma+\Delta\sigma)]$ in the initial sample (circles) and $rr=RR/[\pi \Delta\sigma(2\sigma+\Delta\sigma)]$ in the random one (triangles) correspondingly, where we choose the bin size $\Delta\sigma=1$. |
Values PAR are the mean numbers over 50 realizations of random sample. error bars represent La jack-knife errors. | Values $RR$ are the mean numbers over 50 realizations of random sample, error bars represent $\sigma$ jack-knife errors. |
As can be seen from Fig. 7.. | As can be seen from Fig. \ref{fig_3_4}, |
in. practice the surface density of neighbours for the random. catalogue does. not depend on the distance (except. for the distance range Ss12 Alpe). which is expected for a random clistribution of objects. | in practice the surface density of neighbours for the random catalogue does not depend on the distance (except for the distance range $\lesssim 1-2$ Mpc), which is expected for a random distribution of objects. |
The gap at ~2 Alpe can he explained. by a deficit. of close pairs in the initial catalogue. | The gap at $\sim 1-2$ Mpc can be explained by a deficit of close pairs in the initial catalogue. |
Note that angular distribution of objects in our random samples is the same as in: the initialuu one (see Sec.] PT ??)). "p | Note that angular distribution of objects in our random samples is the same as in the initial one (see Sec. \ref{sec:2.7}) ). |
p:This deficit ⋅⋠⋠is a result of⋅ the fibre⋅ D. collisionellect⋅ (Sec.↴ oe ??)). | This deficit is a result of the fibre collision effect (Sec. \ref{sec:2.6}) ). |
One can see that the number densities for Επ...initial and random samples coinsideos up to 40.50 Alpe. thus we use L=50 Mpe (or 35h.+ Ape). | One can see that the number densities for initial and random samples coinside up to $40-50$ Mpc, thus we use $L=50$ Mpc (or $35\,h^{-1}$ Mpc). |
And we also put Ly=Lh| Alpe because. of the fiber collision clfect. | And we also put $L_0=1\,h^{-1}$ Mpc because of the fiber collision effect. |
“Phe quasar projected. 2pCk a,(a) divided on c in logarithmic scale for the sample is presented in Fig. | The quasar projected 2pCF $w_{p}(\sigma)$ divided on $\sigma$ in logarithmic scale for the sample is presented in Fig. |
S with the the best fit single power-low over the range Leao<35h+ Alpe. | \ref{fig_proj_in} with the the best fit single power-low over the range $1<\sigma<35\,h^{-1}$ Mpc. |
The same calculations were provided for 50 test samples (see Sec. 22)). | The same calculations were provided for 50 test samples (see Sec. \ref{sec:2.2}) ). |
Phe scattering of the resulting parameters is presented in Fig. | The scattering of the resulting parameters is presented in Fig. |
9 with the mean values presented in the | \ref{fig_pp} with the mean values presented in the |
1) The "outer shell’ of low ionization knots is not expanding racially for knots can be found with a large range of radial velocities on its perimeter = 1200 kms+ to 1300. kmsL ). | 1) The `outer shell' of low ionization knots is not expanding radially for knots can be found with a large range of radial velocities on its perimeter = $-$ 1200 $\kms$ to +300 $\kms$ ). |
Even4 some knots. emitting. only. may have = 1450 kms.+. | Even some knots, emitting only, may have = $-$ 1450 $\kms$. |
Some association with a ‘polar blowout” is favoured. | Some association with a `polar blowout' is favoured. |
2) Changes in both racial velocity anc direction of the ‘spike’ favour its interpretation as a continuous "jet with a Es»ed of 1080 kms.+. | 2) Changes in both radial velocity and direction of the `spike' favour its interpretation as a continuous `jet' with a speed of 1080 $\kms$. |
3) The observations of the ‘arc’ remain ambiguous. | 3) The observations of the `arc' remain ambiguous. |
A coherent loop in the py array of prolile suggests ju it could be a narrow cone with an outflow speed. οἱ 1500 kms+. | A coherent loop in the pv array of profile suggests that it could be a narrow cone with an outflow speed of 1500 $\kms$. |
The absence of the most negative velocity features in the corresponding profiles. casts oubt on this interpretation. | The absence of the most negative velocity features in the corresponding profiles casts doubt on this interpretation. |
Somewhat artificial means can be introduced to overcome this problem. | Somewhat artificial means can be introduced to overcome this problem. |
Alternatively. the the loop is composed. of profiles at extreme negative radial velocities ancd profiles at more modest racial velocities. | Alternatively, the the loop is composed of profiles at extreme negative radial velocities and profiles at more modest radial velocities. |
In this case à most fortuitious alignment of velocity components must occur to give the coherent appearance of the loop. | In this case a most fortuitious alignment of velocity components must occur to give the coherent appearance of the loop. |
Observations of S. i] profiles are needed to resolve the issue. | Observations of [S ] profiles are needed to resolve the issue. |
4) X cgrape-shot model is proposed to explain the knottiness of much of the emission seen in HIST images. | 4) A `grape-shot' model is proposed to explain the knottiness of much of the emission seen in HST images. |
We wish to thank the stalfat the New Technology Telescope. La Silla. and the remote observing facility. Garching for their excellent. assistance during these observations. | We wish to thank the staff at the New Technology Telescope, La Silla, and the remote observing facility, Garching for their excellent assistance during these observations. |
WS and RIRW are grateful to PPARC for receipt of the Post. Doctoral Research Associateships. AJLE a Post Graduate Research Studentship and MB a University. of Alanchester Research Fellowship. | WS and RJRW are grateful to PPARC for receipt of the Post Doctoral Research Associateships, AJH a Post Graduate Research Studentship and MB a University of Manchester Research Fellowship. |
The five slit positions. 1-5 are shown against a sketch of the principal features of the 5j. Carinac nebulositv. | The five slit positions, 1-5 are shown against a sketch of the principal features of the $\eta$ Carinae nebulosity. |
. The spatial extents of the positionvelocity (py) arrays shown in Figs. | The spatial extents of the position–velocity (pv) arrays shown in Figs. |
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